{ "1003/1003.3836_arXiv.txt": { "abstract": "{Tycho Brahe completed his catalogue with the positions and magnitudes of 1004 fixed stars in 1598. This catalogue circulated in manuscript form. Brahe edited a shorter version with 777 stars, printed in 1602, and Kepler edited the full catalogue of 1004 stars, printed in 1627. We provide machine-readable versions\\thanks{The full Tables \\keplere\\ and \\variants\\ (see Table\\,\\ref{t:machine}) and the Table with the latin descriptions of the stars are available in electronic from only at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/} of the three versions of the catalogue, describe the differences between them and briefly discuss their accuracy on the basis of comparison with modern data from the Hipparcos Catalogue. We also compare our results with earlier analyses by Dreyer (1916) and Rawlins (1993), finding good overall agreement. The magnitudes given by Brahe correlate well with modern values, his longitudes and latitudes have error distributions with widths of 2\\arcmin, with excess numbers of stars with larger errors (as compared to Gaussian distributions), in particular for the faintest stars. Errors in positions larger than $\\simeq$10\\arcmin, which comprise about 15\\%\\ of the entries, are likely due to computing or copying errors. ", "introduction": "The astronomical observations of Tycho Brahe improved by an order of magnitude on the positional accuracy achieved by his predecessors. His measurements of the positions of stars on the celestial sphere resulted in a manuscript catalogue in 1598 (Brahe 1598).\\nocite{brahe98} Astronomers and mapmakers throughout Europe used handwritten copies of this catalogue. Brahe edited a shorter version, with 777 stars, which was printed in 1602 as part of \\textit{Astronomiae Instauratae Progymnasmata} (Brahe 1602).\\nocite{brahe02} The full list of 1004 entries with some modifications was published by Johannes Kepler in 1627 as part of the \\textit{Tabulae Rudolphinae} (Kepler 1627).\\nocite{kepler} These catalogues are a monument in the history of astronomy, and as such have been studied repeatedly (e.g. Baily 1843, Dreyer 1890, Rawlins 1993). \\nocite{baily}\\nocite{dreyer} In this paper we describe machine-readable versions of the catalogues. In addition to the numbers given by Brahe (and Kepler) the machine-readable tables provide cross-references between the catalogues, identifications with stars from the (modern!) \\textit{Hipparcos Catalogue} (ESA 1997) and on the basis of these the accuracy of the positions and magnitudes tabulated in the old catalogues. \\nocite{esa} The accuracy of the measurements of Brahe is best studied by reference to his observational logs rather than by reference to his reduced data, and Wesley (1978) has shown that the measurement accuracy varies between the instruments used by Brahe. For the mural quadrant, for example, the average error is 34\\farcs6. In his comprehensive study of Brahe's star catalogue, Rawlins (1993) also refers to the observational logs to correct errors that Brahe made in producing his star catalogue. An important conclusion drawn by Rawlins is that errors exceeding 6\\arcmin\\ are usually the consequence of errors in the reduction of the measurements, rather than in the measurements themselves. Such errors include repetition of stars in several entries, mixing data for different stars for one entry, and spurious entries. Rawlins (1993) produces an improved version of the star catalogue of Brahe, the best version that Brahe could have published on the basis of his measurements. The goal of {\\em our} edition is to present the star catalogue of Brahe as it was available for use to a 17th century astronomer or map maker, i.e.\\ the versions as given by Brahe and Kepler. Our analysis improves on earlier ones in three ways. First and foremost, our analysis is based on the \\textit{Hipparcos Catalogue} (ESA 1997), which is more accurate, complete and homogeneous than the stellar catalogues used in earlier analyses. Second, we grade each identification, discriminating between secure, probable and merely possible. Third, we provide images of each constellation comparing the positions and magnitudes from the Brahe catalogue with the posistion and magnitudes of {\\em all stars in the field}, thus illustrating which stars were selected by Brahe for measurements and which ones not. In describing the different versions of the Catalogues we use the following notation: \\manuscript\\ refers to the Manuscript version (Brahe 1598), \\progym\\ refers to the version edited and printed by Brahe (1602), \\kepler\\ refers to the version edited by Kepler (1627), and \\keplere\\ to our emended version of the latter catalogue. Individual entries are numbered according to the order in which they appear in the different versions. For example, the third entry in \\kepler\\ corresponds to entry 338 in \\manuscript, which we denote as K\\,3 and M\\,338, respectively. (Thus, our M numbers correspond to D numbers of Rawlins 1993.) The sequence number within a constellation is indicated by a number following the abbreviated constellation name: Oph\\,14 is the 14th star of Ophiuchus. ", "conclusions": "\\begin{figure} \\includegraphics[angle=270,width=\\columnwidth]{dvariants.pdf} \\caption{Comparison of the accuracies of positions that are different in different versions of the Star Catalogue of Brahe. $\\Delta_E$ is the positional accuracy of \\keplere, $\\Delta$ the positional accuracy of \\kepler, \\progym, or \\manuscript. Note that symbols may be superposed when a variant position is common to several versions of the catalogue. Stars identified in one but not in the other catalogue are omitted.\\label{f:variants}} \\end{figure} It is beyond the purpose of this paper to provide a full analysis of the Star Catalogue by Brahe and it different editions, but a few general remarks may be made. \\subsection{Identifications, emendations and the three versions} In Table\\,\\ref{t:dreyer}, in the columns labelled {\\em all}, we list the numbers of our identifications of the entries in \\keplere\\ and in \\progym. It is seen that only 14 entries remain unidentified in \\keplere, and 5 in \\progym. A large majority of stars is securely identified. \\begin{table} \\caption{Frequency of flags D of Dreyer (1916) identifications as a function of the flags I of our identifications \\label{t:dreyer}} \\begin{tabular}{l|rrrrr|rrrrr} & \\multicolumn{5}{c}{KeplerE} & \\multicolumn{5}{c}{Progymnasmata} \\\\ I\\verb+\\+D & 0 & 1 & 2 & 3 & all & 0 & 1 & 2 & 3 & all \\\\ \\hline 1 & 11 & 896 & 0 & 22 & 929 & 2 & 735 & 0 & 15 & 752 \\\\ 2 & 0 & 12 & 0 & 2 & 14 & 0 & 8 & 0 & 1 & 9 \\\\ 3 & 11 & 5 & 1 & 1 & 18 & 1 & 2 & 0 & 1 & 4 \\\\ 4 & 3 & 4 & 8 & 2 & 17 & 0 & 4 & 3 & 0 & 7 \\\\ 5 & 11 & 0 & 0 & 3 & 14 & 3 & 0 & 0 & 2 & 5 \\\\ 6 & 2 & 8 & 0 & 2 & 12 & 0 & 0 & 0 & 0 & 0 \\\\ all & 38 & 925 & 9 & 32 & 1004 & 6 & 749 & 3 & 19 & 777 \\end{tabular} \\end{table} The sixteen emendations that we apply to \\kepler\\ affect the number of identifications. In particular, eight lead to secure identifications (all with $\\Delta\\le3\\farcm2$) of previously unidentified stars. One other emendation leaves an unidentified star (K\\,67) unidentified, one gives a probable identification of a previously unidentified star (K\\,68), one improves the positional correspondence with its {\\em Hipparcos} counterpart (K\\,801), and five lead to different identifications with better positions. Note that with one exception (K\\,801) all the emendations that we make to \\kepler\\ are taken from \\progym\\ and/or \\manuscript. In Figure\\,\\ref{f:variants} we show the change in position caused by our emendations to \\kepler, and by different positions in \\progym\\ or \\manuscript\\ with respect to \\keplere. Not surprisingly, all emendations lead to better positions. Figure\\,\\ref{f:variants} shows that most differences between \\progym\\ and \\keplere\\ are small, as are most differences between \\manuscript\\ and \\keplere. Remarkably, the Figure also shows that in all cases where the positions differ strongly between \\kepler\\ and \\progym\\ and/or \\manuscript, the positions in the older catalogues are better. 25 of the 26 entries from \\progym\\ that differ from the corresponding entries in \\keplere\\ have the same identification in both versions, but K\\,547 is identified in \\keplere\\ but unidentified in \\progym. Similarly, 39 of 42 entries from \\manuscript\\ that differ from the corresponding entries in \\keplere\\ have the same identification in both versions, one (K\\,251) is identified in \\keplere\\ but not in \\manuscript, and two (K\\,64 and K\\,300) have a different identification in \\keplere\\ than in \\manuscript. \\subsection{Comparison with Dreyer (1916) and Rawlins (1993)} In Table\\,\\ref{t:dreyer} we compare the identifications as found by us with those given by Dreyer (1916). For both the Brahe (1602) and the emended Kepler (1627) versions, we find that our identifications agree with the earlier ones by Dreyer in most cases. We have identified a number of stars not identified by Dreyer, in some cases prefer another one from several plausible possibilities, and in some cases reject an identification by Dreyer. The numbers in Table\\,\\ref{t:dreyer} should be read as indicative rather than exact, due to unavoidable arbitrariness in some classifications. The pair K\\,146/K\\,147 is an example (see Sect.\\,\\ref{s:notes}): we flagged our identifications as secure (I=1) and Dreyer's as wrong (D=3), but could have chosen ours as one of several possibilites (I=4) and Dreyer's as an alternative to our choice (D=2). Another example is the case of K\\,218 in Cygnus. We have assigned HIP\\,106062 as its counterpart, because the closer and brighter counterpart has been assigned to K\\,440, a star in Pegasus. This is a reasonable choice {\\em if} we assume that Kepler was aware of the proximity of K\\,440 to K\\,218. If such was not the case, we could follow Dreyer and consider K\\,218 as a repeated entry for K\\,440, and our identification flag would be I=6 rather than I=2; and the flag for Dreyer's identification D=1 rather than D=3. The sixteen emendations that we apply to \\kepler\\ also affect the numbers in Table\\,\\ref{t:dreyer}. Comparison of our identifications with those by Rawlins (1993) must be made with some care, because his identifications refer to an ideal version of the catalogue, that Brahe might have produced given the time, whereas our identifications refer to the catalogue in the versions edited by Brahe and Kepler. Thus Rawlins identifies the stars that Brahe actually observed, whereas we identify the stars closest to the catalogue positions. In Table\\,\\ref{t:rawlins} we compare the identifications as found by us with those given by Rawlins for the emended Kepler edition. The three entries in \\keplere\\ that do not occur in \\manuscript\\ are not discussed by Rawlins; all other entries are identified. In 937 cases our identification agrees with the one by Rawlins. We include in this four identifications (of K\\,583, K\\,718, K\\,120, and K\\,411) given by Rawlins (1993) that refer to one of a close pair of stars, whereas our identification refers to the other star of the pair. In each of these cases the pair is not separable with the naked eye, with a separation $<2$\\arcmin, and our identification refers to the star that is brighter in the {\\em Hipparcos Catalogue}. In 911 cases the identifications given by Dreyer (1916), Rawlins (1993) and us all agree. In one case the identification given by Rawlins refers to one of two possibilities considered by us. In 53 cases Rawlins finds a different identification because he has corrected the catalogue position. For 4 entries, all in Ophiuchus, Rawlins concludes that Brahe invented positions without having observed them: they are `utter fakes' (see Fig.\\,\\ref{f:ophdetail}). \\begin{table} \\caption{Frequency of flags R of Rawlins (1993) identifications as a function of the flags I of our identifications for \\keplere. \\label{t:rawlins}} \\begin{tabular}{l|rrrrrrrr} I\\verb+\\+R & 0 & 1 & 2 & 3 & 4 & 5 & all \\\\ \\hline 1 & 2 & 900 & 0 & 5 & 22 & 0 & 929 \\\\ 2 & 0 & 11 & 0 & 0 & 3 & 0 & 14 \\\\ 3 & 0 & 6 & 0 & 0 & 8 & 4 & 18 \\\\ 4 & 0 & 9 & 1 & 0 & 7 & 0 & 17 \\\\ 5 & 0 & 0 & 0 & 1 & 13 & 0 & 14 \\\\ 6 & 1 & 11 & 0 & 0 & 0 & 0 & 12 \\\\ all & 3 & 937 & 1 & 6 & 53 & 4 & 1004 \\end{tabular} \\end{table} This leaves 6 entries where our identification is different from that given by Rawlins. Five of these concern pairs of stars, with separations varying from 3\\farcm8 to 10\\farcm5, in which our suggested counterpart is closer to the catalogue position than the counterpart given by Rawlins. An example is shown in Fig.\\,\\ref{f:oridetail}. In four cases (K\\,183, K\\,209, K\\,671, K\\,870) our counterpart is the brighter star of the pair, in one case (K\\,804) only slightly fainter than the other star. In some cases, e.g. K\\,175, K\\,183 and K\\,209, Rawlins combines two stars separated by 8-11\\arcmin\\ into one counterpart; in such cases we may choose the brighter and/or closer star as the counterpart, or leave the entry unidentified. The corrections applied by Rawlins (1993) also affect the number of repeated entries, as indicated in Table\\,\\ref{t:doubles}. \\subsection{Accuracy} Table\\,\\ref{t:dreyer} shows that 14 stars remain unidentified in our emended Kepler catalogue. Two entries in \\keplere\\ have a secure identification, but no counterpart from the {\\em Hipparcos Catalogue}: K\\,267 is SN\\,1572 and K\\,577 is Praesepe. To identify the fourteen unidentified stars would require that one accepts either a fainter counterpart, or a larger positional offset. It is necessary to note that such acceptance increases the probability of chance coincidences, i.e.\\ of spurious identifications. That this is a serious problem may be concluded from the fact that we classified as `secure' four identifications of entries in Kepler (1627) that were identified with {\\em other} counterparts after our emendation was applied. Thirteen of our unidentified entries are corrected by Rawlins (1993) to new positions, that allow him to identify them. One of our unidentified entries (K\\,175) is identified by Rawlins as the combined light of HIP\\,87045 ($V$=6.47) and HIP\\,87119 ($V$=6.83) two stars separated by 8\\farcm8. \\begin{figure} \\includegraphics[angle=270,width=\\columnwidth]{progymv.pdf} \\includegraphics[angle=270,width=\\columnwidth]{keplerev.pdf} \\caption{Distribution of the magnitudes in \\progym\\ (above) and \\keplere\\ (below). In the large frames the histograms indicate the magnitudes according to Brahe for stars which we have securely identified (red; flags 1-2 in Table\\,\\ref{t:idflags}) or not securely identified (blue, flags 3-5), and the magnitudes from the {\\em Hipparcos} catalogue for securely identified stars (black). The numbers of securely and not-securely identified stars are indicated. The small frames give the {\\em Hipparcos} magnitude distributions for securely identified stars for each magnitude according to Brahe separately. The number of securely identified stars at each (Brahe) magnitude is indicated. \\label{f:magnitudes}} \\end{figure} In Figure\\,\\ref{f:magnitudes} we compare the magnitude distributions of the stars in \\keplere\\ to those of their securely identified counterparts. In our opinion, a difference between an {\\em Hipparcos} magnitude and the magnitude assigned by Brahe cannot be called an error, since the magnitudes for Brahe correspond to a classification rather than a measurement. It is striking that Brahe's classification in general correlates well with the modern magnitude; only his magnitude 6 corresponds to mostly brighter stars in {\\em Hipparcos}. The number of securely identified stars peaks at modern magnitudes 4 and 5, and rapidly drops for magnitude 6. This lends support to our hesitance in accepting stars with modern magnitudes $V>6$ as feasible counterparts. \\begin{figure} \\includegraphics[angle=270,width=0.48\\columnwidth]{lamdlam.pdf} \\includegraphics[angle=270,width=0.48\\columnwidth]{lamdbet.pdf} \\includegraphics[angle=270,width=\\columnwidth]{dlondlat.pdf} \\caption{Above: Correlations of the differences in longitude $\\Delta\\lambda\\equiv(\\lambda_\\mathrm{H}-\\lambda)\\cos\\beta$ and latitude $\\Delta\\beta\\equiv\\beta_\\mathrm{H}-\\beta$ of the entries in \\keplere\\ and their secure {\\em Hipparcos} counterparts (converted to 1601) with longitude. Below left: Distributions of $\\Delta\\lambda$ and $\\Delta\\beta$. The numbers of sources with $\\Delta\\lambda,\\Delta\\beta<-10\\arcmin$, of sources included in the histogram ($-10\\arcmin<\\Delta\\lambda,\\Delta\\beta<10\\arcmin$), and of sources with $\\Delta\\lambda,\\Delta\\beta>10\\arcmin$ are indicated. Below right: correlation between $\\Delta\\lambda$ and $\\Delta\\beta$. \\label{f:dlongdlat}} \\end{figure} In Figure\\,\\ref{f:dlongdlat} we show the offsets between the position in \\keplere\\ and the position derived from the {\\em Hipparcos} data, for longitudes and latitudes separately. We showed in Sect.\\,\\ref{s:conversion} that errors in conversion of the modern data to the positions in 1601 are negligible, so that the offsets describe the errors in the position as given by Brahe (or Kepler). If the errors were fully random, we might expect their distributions to be Gaussian, but this is not the case: gaussians that fit the peak of the distribution ($|\\Delta\\lambda|,|\\Delta\\beta|<5\\arcmin$, say) have a width $\\sigma\\simeq2\\arcmin$ and predict much smaller numbers at $|\\Delta\\lambda|,|\\Delta\\beta|>5\\arcmin$ than observed. The excess in the wings of the distributions with respect to a gaussian description is presumably due to the correlations and buildup of errors when the position of a star is determined by measurement with respect to another star which already has a positional error. However, for large errors, $>10\\arcmin$ say, the possibility of a copying error must be considered, as illustrated by the differences between the different versions of the catalogue (see Fig.\\,\\ref{f:variants}), and as proven for many cases by Rawlins (1993). \\begin{figure} \\includegraphics[angle=270,width=\\columnwidth]{alphadb.pdf} \\includegraphics[angle=270,width=\\columnwidth]{brahrefr.pdf} \\caption{Above: Latitude errors $\\Delta\\beta$ in \\keplere\\ as a function of right ascension. The nine standard stars of Brahe are indicated with a red star; the red line indicates 1\\farcm9 $\\sin\\alpha$ (see Eq.\\,\\ref{e:alphadb}). Below: the difference between modern values for diffraction and the values used by Brahe (zero at $h>20$) for stars, as a function of altitude $h$ (Saemundsson, T. 1986; Meeus 1998; Brahe 1602, p.287). \\label{f:alphadb}} \\end{figure} The average offsets $\\Delta\\lambda$ and $\\Delta\\beta$ are not zero, but for Gaussian fits to the central part of the distributions are both around $-0\\farcm5$. Together with the small but systematic dependence of the average offsets on longitude, also shown in Figure\\,\\ref{f:dlongdlat}, this suggests that a small part of the errors may be due to small errors in the position of the zero point in longitude and in the value of the obliquity that Brahe used. Brahe used an obliquity $\\epsilon_B$=23\\fdg525 (Brahe 1602, p.18 and p.208) whereas the correct value for 1601 according to modern determinations (see Sect.\\,\\ref{s:conversion}) was $\\epsilon$=23\\fdg491. For small declinations $\\delta$, the resulting error in latitude $\\Delta\\beta\\equiv\\beta-\\beta_B$ due to the error $\\Delta\\epsilon\\equiv\\epsilon-\\epsilon_B$ after converting equatorial to ecliptic coordinates may be written \\begin{equation} \\cos\\beta\\,\\Delta\\beta \\simeq -sin\\alpha\\cos\\epsilon\\,\\Delta\\epsilon= 1\\farcm9\\sin\\alpha \\label{e:alphadb}\\end{equation} This dependence on right ascension is clearly seen in Brahe's data, as illustrated in Fig.\\,\\ref{f:alphadb}. Brahe also assumed that refraction at altitudes above 20$^\\circ$ is negligible, whereas a modern estimate would give 0\\farcm7 at the altitude of the equatorial pole for Hven (latitude almost 56$^\\circ$; the actual refraction depends somewhat on weather circumstances). Declination measurements with respect to this pole would thus all be offset by $-$0\\farcm7, and this would lead to a systematic offset in $\\beta$. Even though the real situation would be more complicated, involving differences in refraction errors between stars measured at different altitudes (see Fig.\\,\\ref{f:alphadb}), we think that this offset largely explains the overall offset in $\\beta$ seen in Figs.\\,\\ref{f:dlongdlat} and \\ref{f:alphadb}, which averages $-$0\\farcm5, and is already present in the positions of Brahe's nine standard stars (Fig.\\,\\ref{f:alphadb}; see also Dreyer 1890, p.387). \\begin{figure} \\includegraphics[angle=270,width=\\columnwidth]{Delta.pdf} \\caption{Distribution of the position errors $\\Delta$ in \\progym\\ (solid histograms) and in \\keplere\\ (open histograms) for all securely identified stars (left) and for the securely identified stars at each Brahe magnitude separately (right). The numbers indicate the number of stars included in the plot (i.e.\\ with $\\Delta<10\\arcmin$) and those excluded ($\\Delta\\ge10\\arcmin$) \\label{f:delta}} \\end{figure} If $\\Delta\\lambda$ and $\\Delta\\beta$ were distributed as Gaussians centered on zero, the distribution of the total position errors $\\Delta\\equiv\\sqrt{\\Delta\\lambda^2+\\Delta\\beta^2}$ could be written \\begin{eqnarray} N f(\\Delta\\lambda)f(\\Delta\\beta)d\\Delta\\lambda d\\Delta\\beta & = & {N\\over2\\pi\\sigma^2}e^{-0.5\\Delta\\lambda^2/\\sigma^2} e^{-0.5\\Delta\\beta^2/\\sigma^2}d\\Delta\\lambda d\\Delta\\beta \\nonumber \\\\ \\equiv N f(\\Delta)d\\Delta & = & N{\\Delta\\over\\sigma}e^{-0.5(\\Delta/\\sigma)^2}d{\\Delta\\over\\sigma} \\label{e:delta} \\end{eqnarray} The maximum of this function is at $\\Delta=\\sigma$. In Figure\\,\\ref{f:delta} we show the distribution of $\\Delta$. The observed error distribution is similar to Eq.\\,\\ref{e:delta}, but has more measurements at large $\\Delta$ than expected from Eq.\\,\\ref{e:delta} with $\\sigma=2$ because the observed distributions of $\\Delta\\lambda$ and $\\Delta\\beta$ are not centered on zero (a small effect) and show excesses at higher offsets with respect to a Gaussian (the dominant effect). Using Kolmogorow-Smirnov tests to compute the probability that the distributions of the errors $\\Delta$ in the range $\\Delta<10\\arcmin$ for $V=n$ is the same as for $V=n+1$ we find no significant differences. However, the distribution of $\\Delta$ for $V\\le2$ is significantly different from that for $3\\le V\\le5$ (probability for being identical is 0.003). Comparing the differences between the distributions of $\\Delta$ over the full range of $\\Delta$, i.e.\\ including large offsets, we find a significant difference between $V=5$ and $V=6$ (probability for being identical is 0.001). We conclude that the positions of stars with Brahe magnitudes 1 and 2 are better than those of fainter stars; and that a larger fraction of the positions of stars with magnitude 6 is wrong by more than 10\\arcmin\\ than for positions of fainter stars." }, "1003/1003.1144_arXiv.txt": { "abstract": "Positive results of dark matter searches in DAMA/NaI and DAMA/LIBRA experiments, being put together with the results of other groups, can imply nontrivial particle physics solutions for cosmological dark matter. Stable particles with charge -2, bound with primordial helium in O-helium \"atoms\" (OHe), represent a specific Warmer than Cold nuclear-interacting form of dark matter. Slowed down in the terrestrial matter, OHe is elusive for direct methods of underground Dark matter detection used in cryogenic experiments. However radiative capture of OHe by Na and I nuclei can lead to annual variations of energy release in the interval of energy 2-5 keV in DAMA/NaI and DAMA/LIBRA experiments. ", "introduction": "Following Refs.~\\refcite{I}--\\refcite{Khlopov:2006dk},\\refcite{KK} and \\refcite{I2} consider charge asymmetric case, when excess of $X^{--}$ provides effective suppression of positively charged species. In the period $100\\s \\le t \\le 300\\s$ at $100 \\keV\\ge T \\ge T_o= I_{o}/27 \\approx 60 \\keV$, $^4He$ has already been formed in the SBBN and virtually all free $X^{--}$ are trapped by $^4He$ in O-helium ``atoms\" $(^4He^{++} X^{--})$. Here the O-helium ionization potential is\\footnote{The account for charge distribution in $He$ nucleus leads to smaller value $I_o \\approx 1.3 \\MeV$ \\cite{Pospelov}.} \\beq I_{o} = Z_{x}^2 Z_{He}^2 \\alpha^2 m_{He}/2 \\approx 1.6 \\MeV,\\label{IO}\\eeq where $\\alpha$ is the fine structure constant,$Z_{He}= 2$ and $Z_{x}= 2$ stands for the absolute value of electric charge of $X^{--}$. The size of these ``atoms\" is \\cite{I,FKS} \\beq R_{o} \\sim 1/(Z_{x} Z_{He}\\alpha m_{He}) \\approx 2 \\cdot 10^{-13} \\cm \\label{REHe} \\eeq Here and further, if not specified, we use the system of units $\\hbar=c=k=1$. Due to nuclear interactions with nuclei of cosmic plasma, the O-helium gas is in thermal equilibrium with plasma and radiation on the Radiation Dominance (RD) stage, while the energy and momentum transfer from plasma is effective. The radiation pressure acting on the plasma is then transferred to density fluctuations of the O-helium gas and transforms them in acoustic waves at scales up to the size of the horizon. At temperature $T < T_{od} \\approx 200 S^{2/3}_3\\eV$ the energy and momentum transfer from baryons to O-helium is not effective \\cite{I,KK} because $n_B \\sv (m_p/m_o) t < 1,$ where $m_o$ is the mass of the $OHe$ atom, $S_3= m_o/(1 \\TeV)$, $m_p$ is the mass of proton, $\\sigma \\approx \\sigma_{o} \\sim \\pi R_{o}^2 \\approx 10^{-25}\\cm^2$ and $v = \\sqrt{3T/m_p}$ is the baryon thermal velocity. Then O-helium gas decouples from plasma. It starts to dominate in the Universe after $t \\sim 10^{12}\\s$ at $T \\le T_{RM} \\approx 1 \\eV$ and O-helium ``atoms\" play the main dynamical role in the development of gravitational instability, triggering the large scale structure formation. The composite nature of O-helium determines the specifics of the corresponding dark matter scenario, which has qualitative feature of a Warmer Than Cold Dark Matter model \\cite{unesco}. Being decoupled from baryonic matter, the $OHe$ gas does not follow the formation of baryonic astrophysical objects (stars, planets, molecular clouds...) and forms dark matter halos of galaxies. It can be easily seen that O-helium gas is collisionless for its number density, saturating galactic dark matter. Taking the average density of baryonic matter one can also find that the Galaxy as a whole is transparent for O-helium in spite of its nuclear interaction. Only individual baryonic objects like stars and planets are opaque for it. O-helium atoms can be destroyed in astrophysical processes, giving rise to acceleration of free $X^{--}$ in the Galaxy. If the mechanisms of $X^{--}$ acceleration are effective, the anomalous low $Z/A$ component of $-2$ charged $X^{--}$ can be present in cosmic rays at the level \\cite{unesco,Mayorov} $X/p \\sim n_{X}/n_g \\sim 10^{-9}S_3^{-1},$ and be within the reach for PAMELA and AMS02 cosmic ray experiments. In the framework of Walking Technicolor model the excess of both stable $X^{--}$ and $Y^{++}$ is possible \\cite{KK2}, the latter being two-three orders of magnitude smaller, than the former. It leads to the two-component composite dark matter scenario with the dominant OHe accompanied by a subdominant WIMP-like component of $(X^{--}Y^{++})$ bound systems. Technibaryons and technileptons can be metastable and decays of $X^{--}$ and $Y^{++}$ can provide explanation for anomalies, observed in high energy cosmic positron spectrum by PAMELA and in high energy electron spectrum by FERMI and ATIC. O-helium collisions in the galactic bulge can lead to excitation of O-helium. If 2S level is excited, pair production dominates over two-photon channel in the de-excitation by $E0$ transition and positron production with the rate $3 \\cdot 10^{42}S_3^{-2} \\s^{-1}$ is not accompanied by strong gamma signal. According to Ref.~\\refcite{Finkbeiner:2007kk} this rate of positron production for $S_3 \\sim 1$ is sufficient to explain the excess in positron annihilation line from bulge, measured by INTEGRAL (see Ref.~\\refcite{integral} for review and references). If $OHe$ levels with nonzero orbital momentum are excited, gamma lines should be observed from transitions ($ n>m$) $E_{nm}= 1.598 \\MeV (1/m^2 -1/n^2)$ (or from the similar transitions corresponding to the case $I_o = 1.287 \\MeV $) at the level of $3 \\cdot 10^{-4}S_3^{-2}(\\cm^2 \\s \\MeV sr)^{-1}$. ", "conclusions": "To conclude, the existence of heavy stable charged particles may not only be compatible with the experimental constraints but even lead to composite dark matter scenario of nuclear interacting Warmer than Cold Dark Matter. This new form of dark matter can provide explanation of excess of positron annihilation line radiation, observed by INTEGRAL in the galactic bulge. The search for stable -2 charge component of cosmic rays is challenging for PAMELA and AMS02 experiments. Decays of heavy charged constituents of composite dark matter can provide explanation for anomalies in spectra of cosmic high energy positrons and electrons, observed by PAMELA, FERMI and ATIC. In the context of our approach search for heavy stable charged quarks and leptons at LHC acquires the significance of experimental probe for components of cosmological composite dark matter. The results of dark matter search in experiments DAMA/NaI and DAMA/LIBRA can be explained in the framework of our scenario without contradiction with negative results of other groups. Our approach contains distinct features, by which the present explanation can be distinguished from other recent approaches to this problem \\cite{Edward} (see also review and more references in Ref.~\\refcite{Gelmini}). The proposed explanation is based on the mechanism of low energy binding of OHe with nuclei. Within the uncertainty of nuclear physics parameters there exists a range at which OHe binding energy with sodium and/or iodine is in the interval 2-6 keV. Radiative capture of OHe to this bound state leads to the corresponding energy release observed as an ionization signal in DAMA detector. OHe concentration in the matter of underground detectors is determined by the equilibrium between the incoming cosmic flux of OHe and diffusion towards the center of Earth. It is rapidly adjusted and follows the change in this flux with the relaxation time of few minutes. Therefore the rate of radiative capture of OHe should experience annual modulations reflected in annual modulations of the ionization signal from these reactions. An inevitable consequence of the proposed explanation is appearance in the matter of DAMA/NaI or DAMA/LIBRA detector anomalous superheavy isotopes of sodium and/or iodine, having the mass roughly by $m_o$ larger, than ordinary isotopes of these elements. Our results show that the ionization signal, detected by DAMA, is proportional to the temperature and should be suppressed in cryogenic detectors. Therefore test of results of DAMA/NaI and DAMA/LIBRA experiments by other experimental groups can become a very nontrivial task, especially, in view of their rejection of electromagnetic part of counting rate in the absence of nuclear recoil. The presented approach sheds new light on the physical nature of dark matter. Specific properties of composite dark matter and its constituents are challenging for their experimental search. OHe interaction with matter is an important aspect of these studies. In this context positive result of DAMA/NaI and DAMA/LIBRA experiments may be a signature for exciting phenomena of O-helium nuclear physics." }, "1003/1003.3141_arXiv.txt": { "abstract": "{} {We investigate the molecular gas properties of the deeply obscured luminous infrared galaxy NGC~4418. We address the excitation of the complex molecule HC$_3$N to determine whether its unusually luminous emission is related to the nature of the buried nuclear source.} {We use IRAM 30m and JCMT observations of rotational and vibrational lines of HC$_3$N to model the excitation of the molecule by means of rotational diagrams.} {We report the first confirmed extragalactic detection of vibrational lines of HC$_3$N. We detect 6 different rotational transitions ranging from $J$=10--9 to $J$=30--29 in the ground vibrational state and obtain a tentative detection of the $J$=38--37 line. We also detect 7 rotational transitions of the vibrationally excited states $v_6$ and $v_7$, with angular momenta ranging from $J$=10--9 to 28--27. The energies of the upper states of the observed transitions range from 20 to 850 K. In the optically thin regime, we find that the rotational transitions of the vibrational ground state can be fitted for two temperatures, 30 K and 260 K, while the vibrationally excited levels can be fitted for a rotational temperature of 90 K and a vibrational temperature of 500 K. In the inner 300 pc of NGC~4418, we estimate a high HC$_3$N abundance, of the order of 10$^{-7}$. } {The excitation of the HC$_3$N molecule responds strongly to the intense radiation field and the presence of warm, dense gas and dust at the center of NGC~4418. The intense HC$_3$N line emission is a result of both high abundances and excitation. The properties of the HC$_3$N emitting gas are similar to those found for hot cores in Sgr~B2, which implies that the nucleus ($<$ 300 pc) of NGC~4418 is reminiscent of a hot core. The potential presence of a compact, hot component (T=500 K) is also discussed.} ", "introduction": "Luminous infrared galaxies (LIRGs) are intriguing challenges to modern astronomy. They emit most of their radiation in the infrared (IR) region of the spectrum in the form of dust thermal continuum, and have typical IR luminosities of L$_{IR}>$10$^{10}$ L$_\\odot$. In many of these objects, the central power source responsible for the total energy output is buried deep inside the dusty central region and has an origin that remains unclear. For instance, \\citet{spoon07,nascent} suggest that dusty, compact LIRGs may represent the early obscured stages of either active galactic nuclei (AGNs) or starbursts and thus play a fundamental role in galaxy formation and evolution.\\\\ Owing to the large column of intervening material, observations at IR or shorter wavelengths only detect the surface of the optically thick nuclear regions, where molecular gas column densities can reach values that exceed $N({\\rm H}_2)=10^{24}$ $\\cmmt$. Thus optical emission becomes completely absorbed by the gas and dust, unless the source geometry allows the emission to escape. These high masses of gas and dust may also cause an AGN to become Compton thick and hard X-rays to become absorbed by the intervening material, adding to problems in identifying the nuclear source. At radio wavelengths, free-free emission from intervening ionized material may also obscure an AGN \\citep[e.g., ][]{sanders_96}. This leads to great observational ambiguities that are reflected by the classification of LIRGs varying considerably with the observed frequency band. To produce the high central luminosity observed in many LIRGs requires either an AGN or a nuclear starburst (or a combination of the two) to heat a large volume of dust. Dust temperatures in the inner hundred to a few hundred parsec vary among galaxies and can range from 20-30 K to in excess of 170 K \\citep[e.g., ][]{evans03,downes07}. Dissipation of turbulence, which acts on the gas phase \\citep[e.g., ][]{guesten93} may also indirectly contribute to the heating of dust. Because of the aforementioned large amount of obscuring material, the interplay between the possible energy sources remains difficult to ascertain by direct investigation.\\\\ Millimetre observations of molecular lines provide a potentially valuable tool for determining the effects of the central power source on the interstellar medium (ISM) of LIRGs and indirectly constraining some of their key properties, such as gas temperature, density, and chemistry. Chemical models \\citep{meijerink07} show how X-ray dominated regions (XDRs), generally expected in the case of accretion onto a central compact object, and photodissociation regions (PDRs), generated by large UV fluxes from young stars, leave different imprints in the ISM composition. The peculiar molecular chemistry of dense hot cores around nascent stars was also described by \\citet{viti08}. The gas temperature structure is also expected to be different around an AGN compared to a starburst. For example, in an XDR model bulk gas temperatures can be as high as 200 K, while in a starburst the temperatures should be around 20-50 K \\citep{meijerink07}.\\\\ Molecular lines surveys of LIRGs, focused mainly on high density tracers, such as HCN, HCO$^+$, HNC, and CS, have been carried out by several groups \\citep[e.g., ][]{krips08,gracia08,baan08,imanishi07,gao04,aalto02}. However, even if these studies provide unprecedented insights into the ISM of active galaxies, the interpretation of their results remains debated and more sensitive tracers need to be found. \\\\In surveys of external galaxies (see Sect. \\ref{sec:othergal}), bright HC$_3$N\\ $J$=10--9 line emission is found in a subset of IR luminous galaxies. The HC$_3$N\\ line emission is a useful tracer of warm and dense regions and is extremely sensitive to a strong IR-field because of its multiple, mid-IR bending modes. The molecule is often found to be abundant in Galactic hot cores \\citep[e.g., ][]{devicente00} and can be easily destroyed by intense UV and particle radiation. Thus HC$_3$N\\ serves as a tracer of the gas (and dust) properties of galaxies with intense IR fields but where the dense gas is not too exposed to destructive radiation. The edge-on, Sa-type galaxy NGC~4418 has one of the highest luminosities in HC$_3$N (relative to HCN) found for an external galaxy so far \\citep{nascent,monje08}. Its unusual HC$_3$N emission was first reported by \\citet{nascent}. The inner region of NGC~4418 is deeply dust-enshrouded \\citep{spoon_01} with mid-IR intensities that are indicative of dust temperatures of 85 K \\citep{evans03} inside a radius of 50 pc. The IR luminosity-to-molecular gas mass ratio is high for a non-ULIRG galaxy indicating that intense, compact activity is hidden behind the dust. Interferometric observations of HCN $J$=1--0 by \\citet{imanishi04} show that the bulk of the dense gas is contained in the inner 2$''$, corresponding to a region of 300 pc in diameter. NGC~4418 is a FIR-excess galaxy with a logarithmic IR-to-radio continuum ratio ($q$) of 3 \\citep{roussel03}. This excess may be caused by either a young pre-supernova starburst or a buried AGN \\citep{nascent,roussel03, imanishi04}. The luminous HC$_3$N signature was interpreted as young starburst activity \\citep{nascent}. However we need to perform more detailed observations and modeling of HC$_3$N emission of NGC~4418 to interpret its origin accurately .A tentative detection of a mm vibrational HC$_3$N\\ line \\citep{nascent} led us to explore the excitation and abundance of HC$_3$N\\ in NGC~4418 and search for a model of the dense gas properties in the inner 300 pc of the galaxy. Throughout this paper we assume that the observed HC$_3$N\\ line emission emerges from a region smaller than or equal to the scale of the interferometric HCN observations by \\citet{imanishi04}. \\\\ In this paper, we report the first confirmed extragalactic detections of vibrationally excited HC$_3$N in the LIRG NGC~4418 with the IRAM 30m telescope. We report in total 13 different transitions of HC$_3$N - including vibrationally excited levels - allowing us to compile a first model of the excitation and abundance of HC$_3$N. We also report a tentative detection of the 345 GHz rotational $J$=38--37 line with the JCMT telescope. In Sect. 2, we present the observations and results in terms of line intensities. In Sect. 3, we present the results in terms of HC$_3$N rotational diagrams and, in Sect. 4, we briefly discuss the interpretation of our results and future aims. ", "conclusions": "We have confirmed the first extragalactic detection of vibrational lines of HC$_3$N. We have detected 6 different rotational transitions ranging from $J$=10--9 to $J$=30--29 in the ground vibrational state, plus a tentative detection of the $J$=38--37 line. We have also detected 7 rotational transitions of the vibrationally excited states $v_6$ and $v_7$, with angular momenta ranging from $J$=10--9 to 28--27. In the optically thin regime, we find that the $v=0$ transitions can be reproduced by models of two temperatures, 29 K and 265 K, while the $v_7$ lines can be fitted by a temperature of 91 K. The vibrational temperature, fitted to the J=25-24 transitions is, 519 K. By allowing the column density to vary between the different temperature components, we inferred HC$_3$N column densities of 8$\\times10^{14}$ cm$^{-2}$ for the vibrational ground level, and 9$\\times10^{15}$ cm$^{-2}$ for the vibrational $v_7$=1 transitions.\\\\ The excitation of the HC$_3$N molecule responds strongly to the intense radiation field and the presence of warm, dense gas and dust at the center of NGC~4418. The intense HC$_3$N line emission is a result of both high abundances and excitation. The HC$_3$N excitation and abundances seem similar to those found for hot cores in Sgr B2 in the Galactic Center. This implies that the nucleus of NGC~4418 has properties in common with Galactic hot cores. It cannot be excluded that the hot (500 K) component may be associated with a buried AGN." }, "1003/1003.4282_arXiv.txt": { "abstract": "The anisotropy of clustering in redshift space provides a direct measure of the growth rate of large scale structure in the Universe. Future galaxy redshift surveys will make high precision measurements of these distortions, and will potentially allow us to distinguish between different scenarios for the accelerating expansion of the Universe. Accurate predictions are needed in order to distinguish between competing cosmological models. We study the distortions in the redshift space power spectrum in $\\Lambda$CDM and quintessence dark energy models, using large volume N-body simulations, and test predictions for the form of the redshift space distortions. We find that the linear perturbation theory prediction by Kaiser (1987) is a poor fit to the measured distortions, even on surprisingly large scales $k \\ge 0.05 h$Mpc$^{-1}$. An improved model for the redshift space power spectrum, including the non-linear velocity divergence power spectrum, is presented and agrees with the power spectra measured from the simulations up to $k \\sim 0.2 h$Mpc$^{-1}$. We have found a density-velocity relation which is cosmology independent and which relates the non-linear velocity divergence spectrum to the non-linear matter power spectrum. We provide a formula which generates the non-linear velocity divergence $P(k)$ at any redshift, using only the non-linear matter power spectrum and the linear growth factor at the desired redshift. This formula is accurate to better than $5\\%$ on scales $k<0.2 h $Mpc$^{-1}$ for all the cosmological models discussed in this paper. Our results will extend the statistical power of future galaxy surveys. ", "introduction": "The rate at which cosmic structures grow is set by a competition between gravitational instability and the rate of expansion of the Universe. The growth of structure can be measured by analysing the distortions in the galaxy clustering pattern, when viewed in redshift space (i.e. when a galaxy's redshift is used to infer its radial position). Proof of concept of this approach came recently from \\citet{Guzzo:2008ac} who used spectroscopic data for 10,000 galaxies from the VIMOS-VLT Deep Survey \\citep{LeFevre:2004hv} to measure the growth rate of structure at redshift $z=0.77$ to an accuracy of $\\sim 40\\%$ \\citep[see also][]{2001Natur.410..169P}. To distinguish between competing explanations for the accelerating expansion of the Universe, we need to measure the growth of structure to an accuracy of a few percent over a wide redshift interval. The next generation of galaxy redshift surveys, such as ESA's Euclid mission \\citep{2009ExA....23...39C}, will be able to achieve this precision. These redshift space distortions are commonly modelled using a linear perturbation theory expression. We test the validity of this approximation using large volume N-body simulations to model the redshift space distortions in $\\Lambda$CDM and quintessence dark energy models, to see if it works at the level required to take advantage of the information in forthcoming surveys. The large volume of our simulations means that we are able to find the limits of perturbation theory models. We can also study the impact of non-linearities on large scales in cosmologies with different expansion histories from $\\Lambda$CDM, such as quintessence dark energy. One explanation of the accelerating expansion of the Universe is that a negative pressure dark energy component makes up approximately 70\\% of the present density of the Universe \\citep{Komatsu:2008hk, 2009MNRAS.400.1643S}. Examples of dark energy models include the cosmological constant and a dynamical scalar field such as quintessence \\citep[ see e.g.][for a review]{Copeland:2006wr}. Other possible solutions require modifications to general relativity and include extensions to the Einstein-Hilbert action, such as $f(R)$ theories or braneworld cosmologies \\citep[see e.g.][]{Dvali:2000hr,Oyaizu:2008sr}. The expansion history of the Universe is described by the scale factor, $a(t)$. Dark energy and modified gravity models can produce similar expansion histories for the Universe, which can be derived from the Hubble parameter measured, for example, using type Ia supernovae. As both dark energy and modified gravity models can be described using an effective equation of state which specifies the expansion history, it is not possible to distinguish between these two possibilities using measurements of the expansion history alone. The growth rate is a measure of how rapidly structures are forming in the Universe. Dark energy or modified gravity models predict different growth rates for the large scale structure of the Universe, which can be measured using redshift space distortions of clustering. As noted by \\citet{Linder:2005in}, in the case of general relativity, the second order differential equation for the growth of density perturbations depends only on the expansion history through the Hubble parameter, $H(a)$, or the equation of state, $w(a)$. This is not the case for modified gravity theories. By comparing the cosmic expansion history with the growth of structure, it is possible to distinguish the physical origin of the accelerating expansion of the Universe as being due either to dark energy or modified gravity \\citep*{PhysRevD.69.124015,Linder:2005in}. If there is no discrepancy between the observed growth rate and the theoretical prediction assuming general relativity, this implies that a dark energy component alone can explain the accelerated expansion. Galaxy redshift surveys allow us to study the 3D spatial distribution of galaxies and clusters. In a homogeneous universe, redshift measurements would probe only the Hubble flow and would provide accurate radial distances for galaxies. In reality, peculiar velocities are gravitationally induced by inhomogeneous structure and distort the measured distances. \\citet{Kaiser:1987qv} described the anisotropy of the clustering pattern in redshift space but restricted his calculation to large scales where linear perturbation theory should be applicable. In the linear regime, the matter power spectrum in redshift space is a function of the power spectrum in real space and the parameter $\\beta = f/b$ where $f$ is the linear growth rate. The linear bias factor, $b$, characterises the clustering of galaxies with respect to the underlying mass distribution \\citep*[e.g.][]{Kaiser:1987qv}. \\citet{Scoccimarro:2004tg} extended the analysis of \\citet{Kaiser:1987qv} into the non-linear regime, including the contribution of peculiar velocities on small scales. We test this model in this paper. Perturbations in bulk flows converge more slowly then perturbations in density, and so very large volume simulations are needed to model these flows, and hence the redshift space distortion of clustering, accurately. Our simulation boxes are 125 times the volume of those used by \\citet{Cole:1993kh} and $\\sim 30$ times the volume of the N-body results interpreted by \\citet{Scoccimarro:2004tg}. \\citet{2009MNRAS.393..297P} used a single $1h^{-1}$Gpc box to study redshift space distortions in a $\\Lambda$CDM model. Their simulation is over three time smaller than the one we consider. This paper is organised as follows: In Section \\ref{RSD} we discuss the linear growth rate and review the theory of redshift space distortions on linear and non-linear scales. In Section \\ref{2} we present the quintessence models considered and the details of our N-body simulations. The main results of the paper are presented in Sections \\ref{PS12} and \\ref{RT12}. The linear theory redshift space distortion, as well as models for the redshift space power spectrum which include non-linear effects are examined in Section \\ref{PS12} for various dark energy cosmologies. In Section \\ref{RT12} we present the density-velocity relation measured from the simulations. Using this relation the non-linear models used in the previous section can be made cosmology independent. We present a prescription for obtaining the non-linear velocity divergence power spectrum from the non-linear matter power spectrum at an arbitrary redshift in Section \\ref{rt2}. Our conclusions are presented in Section \\ref{5.1}. ", "conclusions": "} One of the primary goals of future galaxy redshift surveys is to determine the physics behind the accelerating expansion of the Universe by making an accurate measurement of the growth rate, $f$, of large scale structure \\citep{2009ExA....23...39C}. Measuring the growth rate with an error of less than $10\\%$ is one of the main science goals of Euclid, as this will allow us to distinguish modified gravity from dark energy models. With an independent measurement of the expansion history, the predicted growth rate for a dark energy model would agree with the observed value of $f$ if general relativity holds. We use simulations of three quintessence dark energy models which have different expansion histories, linear growth rates and power spectra compared to $\\Lambda$CDM. In a previous paper, \\citet{2010MNRAS.401.2181J}, we carried out the first fully consistent N-body simulations of quintessence dark energy, taking into account different expansion histories, linear theory power spectra and best fitting cosmological parameters $\\Omega_{\\rm m}$, $\\Omega_{\\rm b}$ and $H_0$, for each model. In this paper we examine the redshift space distortions in the SUGRA, CNR and 2EXP quintessence models. These models are representative of a broader class of quintessence models which have different growth histories and dark energy densities at early times compared to $\\Lambda$CDM. In particular the SUGRA model has a linear growth rate that differs from $\\Lambda$CDM by $\\sim 20\\%$ at $z=5$ and the CNR model has high levels of dark energy at early times, $\\Omega_{\\rm \\tiny{DE}} \\sim 0.03$ at $z \\sim 200$. The 2EXP model has a similar expansion history to $\\Lambda$CDM at low redshifts, $z<5$, despite having a dynamical equation of state for the dark energy component. For more details on each of the dark energy models see \\citet{2010MNRAS.401.2181J}. Redshift space distortions observed in galaxy surveys are the result of peculiar velocities which are coherent on large scales, leading to a boost in the observed redshift space power spectrum compared to the real space power spectrum \\citep{Kaiser:1987qv}. On small scales these peculiar velocities are incoherent and give rise to a damping in the ratio of the redshift to real space power spectrum. The Kaiser formula is a prediction of the boost in this ratio on very large scales, where the growth is assumed to be linear, and can be expressed as a function of the linear growth rate and bias, neglecting all non-linear contributions. In previous work, using N-body simulations in a periodic cube of $300 h ^{-1}$Mpc on a side, \\citet{Cole:1993kh} found that the measured value of $\\beta =f/b$, where $b$ is the linear bias, deviates from the Kaiser formula on wavelengths of $50 h^{-1}$ Mpc or more as a result of these non-linearities. \\citet{Hatton:1997xs} extended this analysis to slightly larger scales using the Zel'dovich approximation combined with a dispersion model where non-linear velocities are treated as random perturbations to the linear theory velocity. These previous studies do not provide an accurate description of the non-linearities in the velocity field for two reasons. Firstly, the Zel'dovich approximation does not model the velocities correctly, as it only treats part of the bulk motions. Secondly, in a computational box of length $300 h ^{-1}$Mpc, the power which determines the bulk flows has not converged. In this work we use a large computational box of side $1500 h^{-1}$Mpc, which allows us to measure redshift space distortions on large scales to far greater accuracy than in previous work. In this paper we find that the ratio of the monopole of the redshift space power spectrum to the real space power spectrum agrees with the linear theory Kaiser formula only on extremely large scales $k<0.03 h$Mpc$^{-1}$ in both $\\Lambda$CDM and the quintessence dark energy models. We still find significant scatter between choosing different axes as the line of sight, even though we have used a much larger simulation box than that employed in previous studies. As a result we average over the three power spectra, assuming the distortions lie along the $x$, $y$ and $z$ directions in turn, for the redshift space power spectrum in this paper. Instead of using the measured matter power spectrum in real space, we find that the estimator suggested by \\citet{Cole:1993kh}, involving the ratio of the quadrupole to monopole redshift space power spectrum, works better than using the monopole and agrees with the expected linear theory on slightly smaller scales $k< 0.07h$Mpc$^{-1}$ at $z=0$ for both $\\Lambda$CDM and the quintessence models. As the measured redshift space distortions only agree with the Kaiser formula on scales $k<0.07h$Mpc$^{-1}$, it is clear that the linear approximation is not correct on scales which are normally considered to be in the \\lq linear regime\\rq, $k<0.2 h$Mpc$^{-1}$. In linear theory, the velocity divergence power spectrum is simply a product of the matter power spectrum and the square of the linear growth rate. In this work we have demonstrated that non-linear terms in the velocity divergence power spectrum persist on scales $0.04< k (h$Mpc$^{-1})<0.2$. These results agree with \\citet{Scoccimarro:2004tg} who also found significant non-linear corrections due to the evolution of the velocity fields on large scales, assuming a $\\Lambda$CDM cosmology. We have shown that including the non-linear velocity divergence auto and cross power spectrum in the expression for the redshift space $P(k)$ leads to a significant improvement when trying to match the measured quadrupole to monopole ratio for both $\\Lambda$CDM and quintessence dark energy models. Including the non-linear velocity divergence cross and auto power spectra in the expression for the redshift space power spectrum increases the number of parameters needed and depends on the cosmological model that is used. Using the non-linear matter and velocity divergence power spectra we have found a density velocity relation which is model independent over a range of redshifts. Using this relation it is possible to write the non-linear velocity divergence auto or cross power spectrum at a given redshift, $z'$, in terms of the non-linear matter power spectrum and linear growth factor at $z=0$ and $z=z'$. This formula is given in Eq. \\ref{fullmodel} in Section \\ref{rt2}. We find that this formula accurately reproduces the non-linear velocity divergence $P(k)$ to within $10\\%$ for $k<0.3 h $Mpc$^{-1}$ and to better than $5\\%$ for $k<0.2 h$Mpc$^{-1}$ for both $\\Lambda$CDM and the dark energy models used in this paper. It is clear that including the non-linear velocity divergence terms results in an improved model for redshift space distortions on scales $k<0.2 h$Mpc$^{-1}$ for different cosmological models. Current galaxy redshift surveys can provide only very weak constraints on $P_{\\delta \\theta}$ and $P_{\\theta \\theta}$ \\citep{2002MNRAS.335..887T}. The relation given in this paper between the non-linear velocity divergence and matter power spectra will be useful for analysing redshift space distortions in future galaxy surveys as it removes the need to use noiser and sparser velocity data." }, "1003/1003.3694_arXiv.txt": { "abstract": "We have assembled new {\\em Spitzer Space Telescope} Infrared Array Camera observations of the mysterious binary star $\\epsilon$~Aurigae, along with archival far-ultraviolet to mid-infrared data, to form an unprecedented spectral energy distribution spanning three orders of magnitude in wavelength from 0.1 $\\mu$m to 100 $\\mu$m. The observed spectral energy distribution can be reproduced using a three component model consisting of a 2.2$^{+0.9}_{-0.8}$ $M_{\\odot}$ F type post-asymptotic giant branch star, and a 5.9$\\pm$0.8 $M_{\\odot}$ B5$\\pm$1 type main sequence star that is surrounded by a geometrically thick, but partially transparent, disk of gas and dust. At the nominal HIPPARCOS parallax distance of 625 pc, the model normalization yields a radius of $135\\pm5$ $R_{\\odot}$ for the F star, consistent with published interferometric observations. The dusty disk is constrained to be viewed at an inclination of $i\\gtrsim87^{\\circ}$, and has effective temperature of $550\\pm50$ K with an outer radius of 3.8 AU and a thickness of 0.95 AU. The dust content of the disk must be largely confined to grains larger than $\\sim10$ $\\mu$m in order to produce the observed gray optical--infrared eclipses and the lack of broad dust emission features in the archival {\\em Spitzer} mid-infrared spectra. The total mass of the disk, even considering a potential gaseous contribution in addition to the dust that produces the observed infrared excess, is $\\ll1$ $M_{\\odot}$. We discuss evolutionary scenarios for this system that could lead to the current status of the stellar components and suggests possibilities for its future evolution, as well as potential observational tests of our model.\\\\ ", "introduction": "\\label{s:intro} The bright star $\\epsilon$~Aurigae (HD 31964) is a single-lined spectroscopic binary that is famous for its long orbital period (27.1 yr), which is punctuated by an almost 2 yr long eclipse caused by an essentially invisible object \\citep{1991ApJ...367..278C}. The central problem posed by this system is that if the F star component, which dominates the light from the system over a wide range of wavelength and is the eclipsed object, is a massive supergiant (as its spectrum implies), then the invisible companion is surprisingly under-luminous for its mass. Exotic solutions for this mass conundrum involving, for example, a black hole \\citep{1971Natur.229..178C} are not viable because of the train of ever more complicated additional requirements that observational constraints impose on such a model. For example, the lack of significant X-ray emission from the system precludes a black hole {\\em unless} there is no accretion from the disk, which is not possible {\\em unless} there is yet another unseen body (a massive planet, perhaps?) that clears out the space around the black hole, and so on (see the discussions in \\citealt{1986PASP...98..637B}, \\citealt{1991ApJ...367..278C}, and \\citealt{wolk10}). By examining the optical spectra of $\\epsilon$~Aur near the end of its 1954--1956 eclipse, \\citet{1959ApJ...129..291H} was able to deduce the electron density and develop the hypothesis of a Be star-like hot object at the center of a large disk of occulting material \\citep{1961MmSAI..32..351H}. \\citet{1973ApJ...185..229W} reported pioneering infrared (IR) observations that revealed the presence of an excess consistent with the disk being a cloud of partially ionized gas, with a total projected area comparable to that of the F star. An estimate of electron density from the IR excess was consistent with the optical--ultraviolet (UV) estimates of $10^{11}$ cm$^{-3}$. As IR detector technology advanced, \\citet{1984ApJ...284..799B} reported ground-based IR photometry obtained during the 1982--1984 eclipse that demonstrated that the excess could be characterized as a $T=500\\pm150$ K source subtending $8\\times10^{-16}$ sr. \\citet{1985ApJ...299L..99B} refined this result with {\\em Infrared Astronomy Explorer} (IRAS) satellite photometry during the eclipse, extending the wavelength coverage well into the thermal IR, and yielding a revised temperature of $475\\pm50$ K and angular extent of $8.6\\pm1.0\\times10^{-16}$ sr. \\citet{1985Obs...105...90S} examined the same {\\em IRAS} data and was led to conclude that the disk temperature could be better characterized as a $750\\pm100$ K source with a projected area about six times that of the F star photosphere. Regardless of its exact characteristics, the transiting disk in the $\\epsilon$~Aur binary offers a valuable opportunity to study its longitudinal structure in ways not possible with circumstellar disks around single stars. We report here on the results from new mid-IR observations of $\\epsilon$~Aur made with the {\\em Spitzer Space Telescope} \\citep{werner04}, which provide a more precise characterization of the occulting body, as well as a new look at archival data at shorter wavelengths, which better constrain the stellar components. The importance of this exercise lies in the fact that the spectral energy distribution (SED) now can be much more precisely defined, thanks to the availability of new and recalibrated data spanning the far-UV to the mid-IR. These results strongly imply that the putative F supergiant star in $\\epsilon$~Aur is more likely a lower mass, unstable post-AGB object that previously transferred matter to a B(e)-like star companion -- as proposed by \\citet{webbink85} -- resulting in a complex ``dark matter'' disk \\citep{2008ApJ...685..418H} that causes the eclipses. ", "conclusions": "We have analyzed an unprecedented SED of $\\epsilon$~Aur, assembled from observational data spanning three orders of magnitude in wavelength. In conjunction with constraints provided by other published information about this enigmatic binary star, such as its orbital dynamics, we can conclude that:\\ (1) the F star component is a low mass post-AGB object, (2) the B star component is a B5$\\pm$1~{\\rm V} star, and (3) the dusty disk is a partially transparent, low mass disk of predominantly 10 $\\mu$m or larger grains. The identification of the F star in $\\epsilon$~Aur as a normal high mass F supergiant is simply no longer tenable as a plausible scenario. The requisite mass of the B star plus disk cannot be made to satisfy the well-known mass function for this system (without invoking exotic scenarios involving compact, non-luminous sources of 5--10 $M_{\\odot}$ of additional mass that do not produce X-rays -- see Section \\ref{s:intro}), while at the same time satisfying the constraints on the luminosity and SED by the data we have assembled here. From an evolutionary standpoint, the F star must have been initially the more massive of the stellar components in the binary, and some or all of the mass in the dusty disk around the B star may have been transferred from the F star precursor. The bulk of the mass is likely to have escaped from the binary during the evolution of the F star precursor, and might be visible in sensitive observations at far-IR wavelengths. As a post-AGB object, the F star is in a relatively rapid transitional phase of stellar evolution, and we should expect significant changes in the appearance of $\\epsilon$~Aur after the next few thousand to tens of thousands of years. In the meantime, observational studies of the out-of-eclipse disk at wavelengths greater than 30 $\\mu$m and via interferometric imaging are to be encouraged." }, "1003/1003.5551_arXiv.txt": { "abstract": "{}{We aim to interpret the photometric and spectroscopic variability of the luminous blue variable supergiant HD\\,50064 ($V=8.21$).} {CoRoT space photometry and follow-up high-resolution spectroscopy with a time base of 137\\,d and 169\\,d, respectively, was gathered, analysed, and interpreted using standard time series analysis and light curve modelling methods, as well as spectral line diagnostics.} {The space photometry reveals one period of 37\\,d, which undergoes a sudden amplitude change with a factor 1.6. The pulsation period is confirmed in the spectroscopy, which additionally reveals metal line radial velocity values differing by $\\sim 30\\,$km\\,s$^{-1}$ depending on the spectral line and on the epoch. We estimate \\teff$\\sim$13\\,500\\,K, \\logg$\\sim$1.5 from the equivalent width of Si lines. The Balmer lines reveal that the star undergoes episodes of changing mass loss on a time scale similar to the changes in the photometric and spectroscopic variability, with an average value of $\\log\\dot{\\rm M}\\simeq-5$ (in M$_\\odot$\\,yr$^{-1}$). We {tentatively interpret the 37\\,d period as the result of a strange mode oscillation.}}{} ", "introduction": "One of the goals of the asteroseismology programme (Michel et al.\\ 2006) of the CoRoT satellite (Baglin et al.\\ 2006) is to explore the Hertzsprung-Russell diagram (HRD) through uninterrupted time series of white-light photometry of unprecedented precision. In this context, numerous non-radial pulsators of various kind have been observed and analysed, among which massive stars on the main sequence (e.g., Degroote et al.\\ 2009; Neiner et al.\\ 2009). With the goals of mapping the uppermost part of the HRD and understanding the role of oscillations in the mass loss of evolved massive stars, a hot supergiant was observed in the seismology programme of the satellite. \\begin{figure*}[ht!] \\begin{center} \\rotatebox{270}{\\resizebox{5.cm}{!}{\\includegraphics{lichtkromme.eps}}} \\end{center} \\caption{The CoRoT light curve of HD\\,50064 corrected for a linear downward trend (full line). For an explanation of the fits in different linestyles, see text. The $y$-axis of the lower panel is reduced by a factor 2 compared with the one in the upper panel.} \\label{lc} \\end{figure*} \\begin{figure*} \\begin{center} \\rotatebox{270}{\\resizebox{3.65cm}{!}{\\includegraphics{gemha.eps}}} \\rotatebox{270}{\\resizebox{3.65cm}{!}{\\includegraphics{gemhb.eps}}} \\rotatebox{270}{\\resizebox{3.65cm}{!}{\\includegraphics{gemhg.eps}}} \\rotatebox{270}{\\resizebox{3.65cm}{!}{\\includegraphics{gemhe.eps}}} \\rotatebox{270}{\\resizebox{3.65cm}{!}{\\includegraphics{gemmg.eps}}} \\end{center} \\caption{Average profiles at three epochs of H, He, and an Mg line of HD\\,50064 (from left to right: H$\\alpha$, H$\\beta$, H$\\gamma$, \\ion{He}{I} at 6678\\AA, and \\ion{He}{I} at 4471\\AA, as well as \\ion{Mg}{II} at 4481\\AA).} \\label{spectra} \\end{figure*} The B-type supergiant HD\\,50064 ($V$ mag of 8.21) has not been studied in detail. Its spectral type assignments range from B1Ia (Jacoby \\& Hunter 1984) to B6Ia (Blanco et al.\\ 1970). It was monitored by CoRoT, whose performance not only delivers two orders of magnitude better precision than any ground-based photometry, but, even more importantly for supergiant stars, also guarantees uninterrupted data during several months. This is essential for progress in understanding massive evolved stars, because ground-based data for supergiants have so far suffered severely from very low duty cycles. The oscillations of evolved massive stars known so far essentially come in two flavours. Classical gravity mode oscillations with periods of a few days, excited by the $\\kappa\\,$mechanism, have recently been found from space photometry (Saio et al.\\ 2006; Lefever et al.\\ 2007a). On the other hand, theory predicts so-called strange modes with periods between roughly 10 and 100\\,d in stars with masses above 40\\,M$_\\odot$. These strange modes, which can be both radial and non-radial in nature, are modes trapped in a cavity caused by a density inversion in the very outer, highly non-adiabatic envelope of stars with a high L/M ratio whose radiation pressure dominates over the gas pressure (Glatzel \\& Kiriakidis 1993; Saio et al.\\ 1998; Glatzel et al.\\ 1999; Dorfi \\& Gautschy 2000). This type of oscillation has been claimed to be responsible for the mass-loss episodes of luminous stars, such as luminous blue variables (LBVs), e.g.\\ Glatzel \\& Kiriakidis (1993), but observational proof of the occurrence of strange modes has not been established so far. Our data of HD\\,50064 { suggests that the star undergoes a strange mode oscillation.} ", "conclusions": "" }, "1003/1003.0232_arXiv.txt": { "abstract": "The existence of a sizable, $\\mathcal{O}\\left(10^{-10}\\text{--}10^{-9}\\mathrm{G}\\right)$, cosmological magnetic field in the early Universe has been postulated as a necessary step in certain formation scenarios for the large scale $\\mathcal{O}(\\mu\\mathrm{G})$ magnetic fields found in galaxies and galaxy clusters. If this field exists then it may induce significant mixing between photons and axion-like particles (ALPs) in the early Universe. The resonant conversion of photons into ALPs in a primordial magnetic field has been studied elsewhere by Mirizzi, Redondo and Sigl (2009). Here we consider the non-resonant mixing between photons and scalar ALPs with masses much less than the plasma frequency along the path, with specific reference to the chameleon scalar field model. The mixing would alter the intensity and polarization state of the cosmic microwave background (CMB) radiation. We find that the average modification to the CMB polarization modes is negligible. However the average modification to the CMB intensity spectrum is more significant and we compare this to high precision measurements of the CMB monopole made by the far infrared absolute spectrophotometer (FIRAS) on board the COBE satellite. The resulting 95\\% confidence limit on the scalar-photon conversion probability in the primordial field (at $100\\,\\mathrm{GHz}$) is $\\Pbar<2.6\\times10^{-2}$. This corresponds to a degenerate constraint on the photon-scalar coupling strength, $g_{\\mathrm{eff}}$, and the magnitude of the primordial magnetic field. Taking the upper bound on the strength of the primordial magnetic field derived from the CMB power spectra, $B_{\\lambda}\\leq 5.0\\times 10^{-9}\\mathrm{G}$, this would imply an upper bound on the photon-scalar coupling strength in the range $g_{\\mathrm{eff}}\\lesssim 7.14\\times10^{-13}\\mathrm{GeV}^{-1}$ to $g_{\\mathrm{eff}}\\lesssim 9.20\\times10^{-14}\\mathrm{GeV}^{-1}$, depending on the power spectrum of the primordial magnetic field. ", "introduction": "} The origin of the observed large-scale magnetic fields of order $\\mu\\mathrm{G}$ found in nearly all galaxies and galaxy clusters is still largely unknown. It is generally believed that the galactic magnetic field develops by some form of amplification from a pre-galactic cosmological magnetic field. The two popular formation scenarios that are considered are either some exponential dynamo mechanism which amplifies a very small seed field of order $10^{-30}\\mathrm{G}$ as the galaxy evolves, or the adiabatic collapse of a larger existing cosmological field of order $(10^{-10}\\text{--}10^{-9})\\mathrm{G}$. There are a number of pros and cons to both scenarios. See \\cite{Kronberg94} for reviews on the subject. Theories explaining the origin of this primordial magnetic (PMF) field are still highly speculative. It has been suggested that a large-scale magnetic field could be produced during inflation if the conformal invariance of the electromagnetic field is broken; see for example \\cite{Ratra92} or more recently \\cite{Kunze10}. To date, there is no astrophysical evidence for the existence of a large-scale cosmological magnetic field, and only upper bounds on its magnitude have been derived. So far the strongest constraints have come from measurements of the cosmic microwave background (CMB) and big-bang nucleosynthesis \\cite{Kronberg94,Tashiro06,Kosowsky09,Yamazaki10,Paoletti10}. The CMB bounds are derived by considering the effects of Faraday rotation induced by the PMF on the CMB power spectra. In \\cite {Kosowsky09} an upper limit in the range $6\\times10^{-8}$ to $2\\times10^{-6}\\mathrm{G}$ was derived for the mean-field amplitude of the PMF at a comoving length scale of $1\\,\\mathrm{Mpc}$ by comparison to the WMAP 5-year data. More recent work \\cite{Yamazaki10,Paoletti10} analysing the WMAP data in combination with other CMB experiments such as ACBAR, CBI and QUAD have placed tighter constraints on the amplitude of the PMF with an upper bound on the mean-field amplitude at $1\\,\\mathrm{Mpc}$ of $\\sim 5\\times10^{-9}\\mathrm{G}$. The existence of a primordial magnetic field would induce mixing between CMB photons and axion-like particles (ALPs). ALPs refer collectively to any very light scalar or pseudo-scalar with a linear coupling to $F_{\\mu\\nu}F^{\\mu\\nu}$ or $\\epsilon_{\\mu\\nu\\rho\\sigma}F^{\\mu\\nu}F^{\\rho\\sigma}$ respectively. In this paper we consider non-resonant mixing of scalar ALPs and CMB photons, with specific reference to the chameleon scalar field model \\cite{Khoury04,Brax07,Burrage08}. Standard ALPs have constant mass and photon-scalar coupling everywhere, while the chameleon model has a density dependent mass. In sparse environments, such as the primordial plasma, the chameleon acts as a very light scalar field and would be indistinguishable from a standard scalar ALP. However in dense environments the chameleon is very heavy and evades the standard ALP constraints \\cite{Brax09}. The current best constraints on the coupling strength between the chameleon and electromagnetic fields are: $g_{\\mathrm{eff}} \\lesssim 9.1\\times 10^{-10} \\mathrm{GeV}^{-1}$ \\cite{Burrage08} and $g_{\\mathrm{eff}} \\lesssim \\left(0.72\\sim22\\right)\\times 10^{-9} \\mathrm{GeV}^{-1}$ \\cite{Davis09}. Other chameleon-like theories exist such as the Olive-Pospelov model \\cite{Davis09,Olive08} which have a density-dependent coupling strength, but we do not discuss these further here. The resonant mixing of photons with scalar and pseudo-scalar ALPs in a primordial magnetic field has been analysed by Mirizzi, Redondo and Sigl \\cite{Mirizzi09}. They find a constraint on the combined magnetic field strength, $B$, and ALP-photon coupling strength, $g$: $g\\langle B^2 \\rangle ^{1/2}\\lesssim 10^{-13}\\sim10^{-11}\\mathrm{GeV}^{-1}\\mathrm{nG}$, for ALP masses between $10^{-14}\\mathrm{eV}$ and $10^{-4}\\mathrm{eV}$. We believe that these results will not necessarily be applicable to the chameleon because its mass evolves as the density of the Universe decreases. In the following analysis we assume a stochastic primordial magnetic field with a power-law power spectrum, similar to the treatment in \\cite{Kosowsky09}. We assume fluctuations in the magnetic field are damped on small scales due to Alfv$\\mathrm{\\acute{e}}$n wave dissipation \\cite{Barrow98}, and subdivide the magnetic field into multiple domains of length $L$ of a comparable size to just above the Alfv$\\mathrm{\\acute{e}}$n wave damping scale $k_{D}^{-1}$. The magnetic field in each domain is assumed to be approximately constant, and correlated to the other domains according to the magnetic power spectrum. A similar method was applied to correlations in quasar polarization spectra by Agarwal, Kamal and Jain \\cite{Agarwal09}. The degree of conversion between chameleons and photons in a magnetic field is inversely proportional to the electron density in the plasma \\cite{Davis09}. Hence the dominant contribution to photon-scalar mixing will take place in the region after recombination when the ionization fraction drops significantly and before reionization. This greatly simplifies the mixing equations because we do not need to evolve the photon-scalar mixing equations through the last scattering surface, nor include the density inhomogeneities present after reionization. We model a scenario in which the primary CMB is formed at the last scattering surface and then evolves through a primordial magnetic field extending from recombination ($z\\sim1100$) to the epoch of reionization ($z\\sim20$). This paper is organized as follows: in section \\ref{chameleon_model} the chameleon model is introduced and the calculations describing photon-scalar mixing in a magnetic field, living in a Friedman-Robertson-Walker (FRW) spacetime, are presented. The power spectrum of the primordial magnetic field is discussed in more detail in section \\ref{magnetic_field}. In section \\ref{evolution} we analyse the evolution of the photon and chameleon states as they propagate through the multiple magnetic domains, and predict the average modification to the CMB intensity and polarization. In section \\ref{experiment}, our predictions are compared to precision measurements of the CMB monopole made by the far infrared absolute spectrophotometer (FIRAS) on board the cosmic background explorer (COBE) satellite. We present a summary of the work and our conclusions in section \\ref{conclusion}. The appendix contains details of the equations governing the evolution of the CMB Stokes parameters. ", "conclusions": "} The existence of a large-scale cosmological field in the early Universe is so far unconfirmed. It would need to be in the region of $\\mathcal{O}\\left(10^{-10}\\sim10^{-9}\\mathrm{G}\\right)$ if it is to explain the formation of the $\\mathcal{O}(\\mu\\mathrm{G})$ magnetic fields in galaxies and galaxy clusters, through adiabatic collapse. If this primordial magnetic field exists it would induce mixing between CMB photons and axion-like particles (ALPs) as they propagate from the last-scattering surface to Earth. In this paper we have studied the case of non-resonant mixing between scalar-ALPs and photons in a primordial magnetic field, with specific reference to the chameleon scalar field model. The chameleon model is a promising candidate for the dark energy scalar field since it can have a gravitational strength (or stronger) coupling to normal matter while at the same time evading fifth-force constraints. To date, the strongest bounds on the chameleon coupling strength come from chameleon-photon mixing in local astrophysical environments, such as starlight propagating through the galactic magnetic field: $g_{\\mathrm{eff}}\\lesssim 9.1\\times 10^{-10}$ \\cite{Burrage08}. Should there be a detection of a primordial magnetic field of order $\\gtrsim 10^{-10}\\mathrm{G}$, our results would place far greater constraints on the coupling strength. We have considered a stochastic primordial magnetic field described by a power-law power spectrum, $P(k)\\propto k^{n_B}$, up to some cut-off damping scale $k_D$. Photon-scalar mixing in this field was solved by dividing the path length into multiple magnetic domains in which the field is approximated as constant. The length of the domains was taken to be of a comparable size to $k_D ^{-1}$ since the magnetic power spectrum will be approximately flat, but non-zero, just above the damping scale. Correlations between the magnetic field strength and direction in each domain are determined by the magnetic power spectrum. In addition, we approximated the ionization fraction as being constant in the region from redshift $\\sim750$ to $20$, and held at its post-recombination freeze-out value of $X_e\\sim 5\\times 10^{-4}$. The dominant contribution to photon-scalar mixing in the CMB will occur in this region of low electron density, and we neglected contributions from other elements of the path length. A more sophisticated approach to modelling the electron density from recombination to the present day may result in small changes to the predictions, but would require a numerical rather than analytical approach to the mixing equations. We have compared our predictions of the average modification of the CMB intensity over the whole sky, to precision measurements of the CMB monopole by the FIRAS instrument on board the COBE satellite \\cite{Fixsen96,Fixsen02}. This constrains the probability of photon-scalar mixing over the path length, to be \\[\\bar{\\mathcal{P}}_{\\gamma\\rightarrow\\phi}\\lesssim 0.026\\;\\;(95\\%\\:\\mathrm{CL})\\] at $100\\,\\mathrm{GHz}$. The corresponding bounds on the magnitude of the magnetic field and photon-scalar coupling strength are plotted in Fig. \\ref{fig:MBbounds} for different values of the magnetic spectral index. Until a detection of the primordial magnetic field is made, to break the degeneneracy of this constraint, we cannot place limits on the chameleon-photon coupling strength, $g_{\\rm eff}$. For the largest magnetic field allowed by the constraints in \\cite{Yamazaki10,Paoletti10} we would find the strongest possible constraint, \\[ g_{\\mathrm{eff}}\\lesssim \\left(0.92 \\sim 7.14\\right)\\times10^{-13}\\mathrm{GeV}^{-1},\\] depending on the slope of the magnetic field power spectrum. The results in this paper apply to any scalar ALP with a mass less than $\\sim 10^{-14}\\mathrm{eV}$, since we require the mass to be much smaller than the plasma frequency along the path. These nicely complement the bounds derived in \\cite{Mirizzi09} for resonant conversion between photons and ALPs in a primordial magnetic field, $g_{\\rm eff}\\langle B^2 \\rangle^{1/2}\\lesssim 10^{-13}\\sim 10^{-11}\\,{\\rm GeV^{-1}\\,nG}$, which apply to ALP masses in the range $10^{-14}{\\rm eV}$ to $10^{-4}{\\rm eV}$. In addition to the average modification to the CMB intensity, the formalism developed in this paper can straightforwardly be extended to calculate the change to correlations between the CMB Stokes parameters along different lines of sight. In section \\ref{multiple_evolution}, an example was given of how a correlation between the $U$ and $V$ polarization modes can arise from photon-scalar mixing in the primordial magnetic field, given a non-zero $\\langle IQ\\rangle$ correlation. If this effect is present in the CMB cross-correlations, it would lead to a significant chameleon signature in the $\\langle EB\\rangle$ power spectrum. A full analysis of this effect is kept for a separate publication \\cite{Schelpe10}. \\vspace{0.5cm} \\noindent{\\bf Acknowledgments:} I am funded by STFC. I am grateful to Douglas Shaw and my supervisor Anne Davis for their support, and to Anthony Challinor for interesting and helpful discussions. \\appendix" }, "1003/1003.0835_arXiv.txt": { "abstract": "We employ a bias-corrected abundance matching technique to investigate the coevolution of the $\\LCDM$ dark halo mass function (HMF), the observationally derived velocity dispersion and stellar mass functions (VDF, SMF) of galaxies between $z=1$ and $0$. We use for the first time the evolution of the VDF constrained through strong lensing statistics by Chae (2010) for galaxy-halo abundance matching studies. As a local benchmark we use a couple of $z \\sim 0$ VDFs (a Monte-Carlo realised VDF based on SDSS DR5 and a directly measured VDF based on SDSS DR6). We then focus on connecting the VDF evolution to the HMF evolution predicted by $N$-body simulations and the SMF evolution constrained by galaxy surveys. On the VDF-HMF connection, we find that the local dark halo virial mass-central stellar velocity dispersion ($\\Mvir$-$\\sigma$) relation is in good agreement with the individual properties of well-studied low-redshift dark haloes, and the VDF evolution closely parallels the HMF evolution meaning little evolution in the $\\Mvir$-$\\sigma$ relation. On the VDF-SMF connection, it is also likely that the stellar mass-stellar velocity dispersion ($\\Mstars$-$\\sigma$) relation evolves little taking the abundance matching results together with other independent observational results and hydrodynamic simulation results. Our results support the simple picture that as the halo grows hierarchically, the stellar mass and the central stellar velocity dispersion grow in parallel. We discuss possible implications of this parallel coevolution for galaxy formation and evolution under the $\\LCDM$ paradigm. ", "introduction": "The current $\\LCDM$ hierarchical structure formation theory predicts robustly the evolution of the dark halo mass\\footnote{Throughout this refers to the total mass within the virial radius of the halo. Accordingly, it includes the stellar mass once the galaxy is formed.} function (HMF) over cosmic time (e.g., \\citealt{Spr05,War06,Ree07,Luk07,Tin08,Kly10}). Because visible galaxies are believed to form and reside within the haloes under the $\\LCDM$ paradigm, the statistical functions of galaxies, such as the luminosity function (LF), the stellar mass function (SMF), and the stellar velocity (dispersion) function (VF, VDF), are also expected to evolve. Connection of these statistical functions of galaxies with the theoretical HMF is not straightforward due to, and mirrors, the complex processes of galaxy formation and evolution including star formations, supernovae explosions, AGN activities and galaxy merging. Some recent works in the literature are focused on the connection of the HMF with the broadly measured stellar mass function (SMF) of galaxies (e.g., \\citealt{CW09,Mos09,Guo09}). Notice that the SMF as well as the LF have mainly to do with the star formation history of galaxies. Galaxy formation in the halo has dynamical consequences as well. As stars are formed in the inner halo, the halo responds and the inner halo dark matter distribution is modified (e.g., \\citealt{Blu86,Gne04,Rud08,Aba09,Tis10}). Consequently, not only the total (i.e.\\ dark plus stellar) mass distribution but also the dark matter distribution may become different from the pure dark matter distribution predicted by the $\\LCDM$. This dynamical aspect of galaxy formation is a crucial part of cosmological studies. Ultimately, the theory of galaxy formation should predict successfully the dynamical evolution as well as the star formation history of galaxies. The statistical property of the dynamics of galaxies is encoded in the VDF of galaxies. The local total VDF is carefully reconstructed by \\citet{Cha10} using SDSS DR5 galaxy counts and intrinsic correlations between luminosities and velocities of galaxies. More recently, \\citet{Ber10} estimates directly the local total VDF based on SDSS DR6 measurements (of a DR4 sample). \\citet{Cha10} then constrains the evolution of the VDF up to $z \\sim 1$ through the statistical properties of strong lensing galaxies based on the empirical result that the average total (luminous plus dark) mass profile of galaxies is isothermal in the optical region. \\citet{Cha10} notices that the {\\it differential} evolution of the derived VDF is qualitatively similar to the evolution of the theoretical HMF under the current $\\LCDM$ paradigm. In this work we make a detailed quantitative comparison between the VDF evolution constrained from strong lensing statistics through the method of \\citet{Cha10} and the evolutions of mass functions (i.e.\\ the HMF and the SMF) from the literature. In doing so, we investigate the local (statistical) correlations of the velocity dispersion ($\\sigma$) of a galaxy with the virial mass ($\\Mvir$) of the surrounding halo and the stellar mass ($\\Mstars$) of the galaxy, i.e.\\ $\\sigma(\\Mvir)$ and $\\sigma(\\Mstars)$, and their evolutions out to $z \\sim 1$. These empirical correlations will provide independent {\\it statistical} constraints on the structures of galaxies and haloes and their evolutions. We investigate the implications of the local correlations for the baryon-modified dark halo structures in a following work. In this work we focus on the evolutions of the correlations. We find that the halo mass, the stellar mass and the stellar velocity dispersion are coevolving in a parallel way for $0 \\la z \\la 1$. We discuss the implications of this result for galaxy formation and evolution under the $\\LCDM$ paradigm. This paper is organised as follows. In \u00a72, we describe the method of analysis and the statistical functions to be used in this work; some details of the analysis are given in the Appendix~A. In \\S 3, we investigate the connection between the (evolving) VDF of galaxies from the SDSS and strong lensing statistics and the HMF from N-body simulations. We obtain the relation $\\sigma(\\Mvir)$ and its evolution. We also examine the compatibility of the evolutions of the VDF and the HMF. In \\S 4, we compare the VDF evolution with the SMF evolution from galaxy surveys. We investigate the evolution of $\\sigma(\\Mstars)$ and the compatibility of the current observationally constrained VDF and SMF. In \\S 5, we discuss the implications of the results for galaxy formation and evolution in the context of the current $\\LCDM$ structure formation paradigm and cosmological observations. We conclude in \\S 6. Unless specified otherwise, we assume a WMAP 5 year $\\LCDM$ cosmology (\\citealt{Dun09}) with $(\\Omega_{{\\rm m}0},\\Omega_{{\\Lambda}0})=(0.25,0.75)$ and $H_0 =100h\\kms$~Mpc$^{-1}$. When parameter $h$ does not appear explicitly, $h=0.7$ is assumed. In Appendix~B we compare the results of this work with the nearly concurrent results by \\citet{Dut10}. \\citet{Dut10} focus on the connection between the circular velocity in the optical region (at about the projected half-light radius) $\\vopt$ and that at the virial radius $\\vvir$. While \\citet{Dut10} use various estimates of the halo mass including satellite kinematics, weak lensing and abundance matching, their results are confined to $z \\sim 0$. Another important difference between \\citet{Dut10} and this work is that \\citet{Dut10} use observed stellar mass-velocity relations while this work uses the observationally derived velocity dispersion functions for abundance matching. ", "conclusions": "Through an abundance matching analysis of the lensing constrained VDF evolution along with the theoretical HMF and the observed SMF from galaxy surveys, we find the following. \\begin{enumerate} \\item The dark halo virial mass-central stellar velocity dispersion ($\\Mvir$-$\\sigma$) relation at $z=0$ is in excellent agreement with the observed properties of low-redshift individual haloes. \\item The stellar mass-central stellar velocity dispersion ($\\Mstars$-$\\sigma$) relation at $z=0$ is consistent with the local scaling relations of galaxies in the literature. \\item The $\\Mvir$-$\\sigma$ relation does not evolve between $z=1$ and $0$ independent of current observation and simulation data. \\item The $\\Mstars$-$\\sigma$ relation does not evolve between $z=1$ and $0$ for the COSMOS SMF. This is well in line with the observed non-evolution of $\\sigma$ with $z$ at $\\Mstars = 10^{11} \\Msun$. This is also consistent with the predicted little or mild evolution of $\\sigma$ with $z$ insensitive to $\\Mstars$ from cosmological simulations. However, the Spitzer SMF (a typical downsizing SMF) requires the $\\Mstars$-$\\sigma$ relation to evolve in a differential way that is not supported by independent observational results on the structural evolutions of galaxies in the literature. \\item The non-evolution in the $\\Mvir$-$\\sigma$ and the $\\Mstars$-$\\sigma$ relations imply a parallel coevolution of $\\Mvir$, $\\Mstars$ and $\\sigma$ between $z=1$ and $0$. This is corroborated by the little evolution in the abundance matching $\\Mvir$-$\\Mstars$ relation between $z=1$ and $0$ for $\\Mvir \\ga 10^{12}\\Msun$. \\item The parallel coevolution of $\\Mvir$, $\\sigma$ and $\\Mstars$ with $z$ may imply a universality and regularity in galaxy formation and evolution despite complex baryonic physics processes. \\end{enumerate} \\vspace{1cm} The author would like to thank Mariangela Bernardi, Nacho Trujillo and Michele Cappellari for useful communications and Andrey Kravtsov, Robert Feldmann, Joshua Frieman, Nick Gnedin and Steve Kent for helpful discussions and conversations. The author also gratefully acknowledges the referee comments that were helpful in clarifying and improving the manuscript significantly." }, "1003/1003.0004_arXiv.txt": { "abstract": "The frequency and effects of multiple weak deflections in galaxy-galaxy lensing are investigated via Monte Carlo simulations. The lenses in the simulations are galaxies with known redshifts and known rest-frame blue luminosities. The frequency of multiple deflections above a given threshold shear value is quantified for discrete source redshifts, as well as for a set of sources that are broadly distributed in redshift space. In general, the closest lens in projection on the sky is not the only lens for a given source. In addition, $\\sim 50$\\% of the time the closest lens is not the most important lens for a given source. Compared to a naive single-deflection calculation in which only the lensing due to the closest weak lens is considered, a full multiple-deflection calculation yields a higher net shear for individual sources, as well as a higher mean tangential shear around the lens centers. The full multiple-deflection calculation also shows that galaxy-galaxy lensing may contribute a substantial amount to cosmic shear on small angular scales. The degree to which galaxy-galaxy lensing contributes to the small-scale cosmic shear is, however, quite sensitive to the mass adopted for the halos of $L_B^\\ast$ galaxies. Changing the halo mass by a factor of $\\sim 2.5$ changes the contribution of galaxy-galaxy lensing to the cosmic shear by a factor of $\\sim 3$ on scales of $\\theta \\sim 1$~arcmin. The contribution of galaxy-galaxy lensing to cosmic shear decreases rapidly with angular scale and extrapolates to zero at $\\theta \\sim 5$~arcmin. This last result is roughly independent of the halo mass and suggests that for scales $\\theta \\gtrsim 5$~arcmin, cosmic shear is insensitive to the details of the gravitational potentials of large galaxies. ", "introduction": "Galaxy-galaxy lensing is the systematic weak gravitational lensing of background galaxies by foreground galaxies. Unlike weak lensing by massive galaxy clusters, where the only important lens in the problem is the cluster itself, galaxy-galaxy lensing involves multiple weak deflections. For example, it is common for a distant source galaxy at redshift $z_s$ to be weakly lensed by a more nearby galaxy at redshift $z_{l1}$, and for both of these galaxies to then be lensed by another (even more nearby) galaxy at redshift $z_{l2}$. Thus, the galaxy with redshift $z_{l1}$ serves simultaneously as a lens for the galaxy at $z_s$ and a source for the galaxy at $z_{l2}$. In addition, the galaxy at $z_s$ is lensed by two independent foreground galaxies. The importance of such multiple deflections in galaxy-galaxy lensing was first noted by Brainerd et al.\\ (1996; hereafter BBS). Since the work of BBS, galaxy-galaxy lensing has been detected with impressively high statistical significance by a number of different groups. This has enabled constraints to be placed on the nature of the dark matter halos that surround the lens galaxies as well as the bias between mass and light in the universe (see, e.g., Fischer et al.\\ 2000; Guzik \\& Seljak 2002; Hoekstra et al.\\ 2004, 2005; Sheldon et al.\\ 2004; Heymans et al.\\ 2006; Kleinheinrich et al.\\ 2006; Mandelbaum et al.\\ 2006ab; Limousin et al.\\ 2007; Parker et al.\\ 2007; Natarajan et al.\\ 2009; Tian et al.\\ 2009 ). The purpose of the present investigation is to: (i) quantify the frequency of multiple weak lensing deflections in a relatively deep galaxy-galaxy lensing data set, (ii) determine the effect of multiple deflections on the net shear for distant source galaxies that have been weakly lensed by foreground galaxies, and (iii) demonstrate that galaxy-galaxy lensing alone may contribute a substantial amount to the ``cosmic shear'' signal on small angular scales. To do this, theoretical shear fields are constructed using a set of observed galaxies with known redshifts and known rest-frame blue luminosities. A simple halo model is used to assign masses to the observed galaxies and Monte Carlo simulations are then used to lens various theoretical source galaxy distributions by the observed galaxies. Theoretical shear fields for full, multiple-deflection calculations are computed; i.e., each source galaxy in the simulation is lensed by all foreground galaxies. In addition, theoretical shear fields for naive, single-deflection calculations (where the closest lens on the sky is assumed to be the only lens) are also computed. The results of the single-deflection calculations are compared to those of the full, multiple-deflection calculations in order to assess the effects of multiple deflections in galaxy-galaxy lensing. Throughout, the weak lensing of an entire source galaxy by a single foreground lens galaxy will be referred to as a ``deflection''. The paper is organized as follows. The Monte Carlo simulations of galaxy-galaxy lensing are described in Section 2, the frequency of multiple weak deflections in galaxy-galaxy lensing is computed in Section 3, the effects of multiple weak deflections on the galaxy-galaxy lensing shear are computed in Section 4, the contribution of galaxy-galaxy lensing to cosmic shear is computed in Section 5, and a discussion of cosmic variance in relation to the field size in presented in Section 6. The conclusions are summarized in Section 7. ", "conclusions": "The frequency and effects of multiple weak lensing deflections in galaxy-galaxy lensing have been investigated using Monte Carlo simulations. The lenses in the simulations are modeled using observed galaxies with magnitudes $R \\le 23$, contained within a circle of radius of 4~arcminutes, centered on the HDF-N. The lenses have known redshifts and known rest-frame B-band luminosities. By adopting a simple halo mass model it is possible to determine the relative strengths of each of the lenses using scaling relations. The Monte Carlo simulations reveal a number of expected results: (i) the frequency of multiple deflections depends upon the minimum value of the shear (i.e., the lower is the minimum value, the more likely it is that multiple deflections will be experienced by a given source), (ii) the frequency of multiple deflections depends upon the source redshift (i.e., the higher is the source redshift, the more likely it is that will experience multiple deflections) and (iii) the higher are the masses of the lenses, the more likely it is that multiple deflections will occur. For a deep galaxy-galaxy lensing data set in which the sources have a median redshift $z_s \\sim 1$ and the lenses have a median redshift $z_l \\sim 0.6$, the probability that a given source galaxy will have experienced more than one weak lens that induces a ``typical'' shear of $\\gamma = 0.005$ ranges from 26\\% to 59\\%, depending upon the masses adopted for the lenses. The Monte Carlo simulations also reveal a number of results that may seem counter-intuitive at first glance: (i) of order 50\\% of the time, the closest lens in projection on the sky is not the most important weak lens for a given source, (ii) for a given source, the net shear due to all foreground lenses generally exceeds the shear due to the strongest individual weak lens, and (iii) multiple deflections give rise to a larger tangential shear around the lens galaxies than a simple, single-deflection calculation in which the closest lens is assumed to be the only lens. This emphasizes the importance of using full, multiple-deflection calculations when using observations of galaxy-galaxy lensing to constrain the parameters of the dark matter halos of the lens galaxies. If multiple deflections are not incorporated into the calculation, this will result in halo masses that are systematically too large. Lastly, the Monte Carlo simulations reveal that galaxy-galaxy lensing alone can give rise to a cosmic shear signal on small angular scales. This is unsurprising because cosmic shear occurs when photons from distant galaxies are deflected by all mass along the line of sight. In the case of galaxy-galaxy lensing, it is the very large $k$ end of the power spectrum of density fluctuations that contributes to the cosmic shear by inducing correlated image shapes for the distant galaxies. On scales $\\theta \\sim 1$~arcmin, the degree to which galaxy-galaxy lensing contributes to cosmic shear is quite sensitive to the masses of the lens galaxies. Changing the mass of the halo of a fiducial $L_B^\\ast$ galaxy by a factor of $\\sim 2.5$ changes the contribution to the top hat shear variance, $\\left< \\gamma^2 \\right>$, by a factor of $\\sim 3$. Comparing the theoretical values of $\\left< \\gamma^2 \\right>$ at $\\theta = 1$~arcmin to the value observed by Fu et al.\\ (2008) for sources with a similar redshift distribution, galaxy-galaxy lensing alone could account for as little as $\\sim 5$\\% or as much as $\\sim 58$\\% of the observed value, depending upon the halo mass for $L_B^\\ast$ galaxies. While the small-scale contribution of galaxy-galaxy lensing to cosmic shear is quite sensitive to the masses of the lenses, the scale at which galaxy-galaxy lensing becomes unimportant to cosmic shear is relatively independent of the lens masses. If the results for the galaxy-galaxy lensing contribution to cosmic shear are extrapolated to large scales, the contribution of galaxy-galaxy lensing to cosmic shear should vanish for scales $\\theta \\gtrsim 5$~arcmin, largely independent of the lens masses." }, "1003/1003.4231_arXiv.txt": { "abstract": "A possible solution to the dark energy problem is that Einstein's theory of general relativity is modified. A suite of models have been proposed that, in general, are unable to predict the correct amount of large scale structure in the distribution of galaxies or anisotropies in the Cosmic Microwave Background. It has been argued, however, that it should be possible to constrain a {\\it general} class of theories of modified gravity by focusing on properties such as the growing mode, gravitational slip and the effective, time varying Newton's constant. We show that assuming certain physical requirements such as stability, metricity and gauge invariance, it is possible to come up with consistency conditions between these various parameters. In this paper we focus on theories which have, at most, 2$^{\\rm nd}$ derivatives in the metric variables and find restrictions that shed light on current and future experimental constraints without having to resort to a (as yet unknown) complete theory of modified gravity. We claim that future measurements of the growth of structure on small scales (i.e. from 1-200 $h^{-1}$ Mpc) may lead to tight constraints on both dark energy and modified theories of gravity. ", "introduction": "The {\\it dark energy} problem, i.e. the possibility that 70$\\%$ of the Universe seems to be permeated by an invisible fluid which behaves repulsively under gravity and does not cluster, has been the focus of research in cosmology for over decade. There are a host of proposals \\cite{CopelandSameTsujikawa2006} and a battery of experiments are under way, or on the drawing board, to characterize the nature of this elusive source of energy \\cite{SKA,RefregierEtAl2010,DES,LSST}. In recent years, an alternative possibility has emerged, that Einstein's General theory of relativity is incorrect on cosmological scales and must be modified. Although the idea that General Relativity is incomplete has been around since the early 1960s \\cite{Dirac1938a,Jordan1949,BransDicke1961,Sakharov1968a}, there are now a number of proposals for what this theory of modified gravity might be \\cite{FerreiraStarkman2009}. The Einstein-Hilbert action, $S_g\\propto\\int d^4x \\sqrt{-g}R$ (where $g$ is the metric determinant and $R$ is the scalar curvature of a metric $g_{ab}$) can be replaced by a more general form $S_g\\propto\\int d^4x \\sqrt{-g}F(R)$ where $F$ is an appropriately chosen function of $R$ \\cite{Amendola07a,SotiriouFaraoni2010a}; the dynamics of the gravitational field can emerge from a theory in higher dimensions such as one might encounter in brane worlds \\cite{DvaliGabadadzePorrati2000d}; a preferred reference frame may emerge from the spontaneous symmetry breaking of local Lorentz symmetry \\cite{ZlosnikFerreiraStarkman2007,JimenezMaroto2008,DimopoulosKarciauskasWagstaff2010a,KoivistoMota2008}; the metric which satisfies the Einstein equation is not necessarily the one that defines geodesic motion \\cite{Bekenstein1993} but is related to a second metric via additional fields~\\cite{Bekenstein2004a,Skordis2008a,Skordis2009a} or connections~\\cite{Banados2008,BanadosEtAl2008,Milgrom2009}; the Einstein-Hilbert action may be deformed by choosing as fundamental variables of gravity, $SU(2)$ connections~\\cite{Plebanski1977,Krasnov2008,KrasnovShtanov2010}. Many of these models have been successful in reproducing, for example, the observed relation between redshift and luminosity distances from distant supernovae. They have, however, generally failed to reproduce the observed clustering of galaxies on large scales as well as the anisotropies in the Cosmic Microwave Background (CMB) unless the modified theory becomes effectively equivalent to general relativity (i.e. the Einstein-Hilbert action and a cosmological constant), e.g.~\\cite{BanadosFerreiraSkordis2008,ZuntzFerreiraZlosnik2008,ZuntzEtAl2010}. The general problem that seems to plague most theories is an excess of power on the very largest scales which manifests itself through the Integrated Sachs Wolfe (ISW) effect and a mismatch between the normalization of the power spectrum of fluctuations on the largest and smallest scales. As yet, a truly compelling and viable model of modified theory of gravity has yet to be but forward which may resolve the dark energy problem. All is not lost, however, and progress can be made in learning about potential modifications to gravity by extracting phenomenogical properties that can be compared to observation- the \"Parametrized Post Friedmannian\" approach \\cite{HuSawicki2007a}. In this paper we focus on a key observable characterizing the evolution of large scale structure: the growing mode of gravitational collapse. The time evolution of the density field can be a sensitive probe of not only the expansion rate of the the Universe but also its matter content. In a flat, matter dominated universe we have that $\\delta_M$, the density contrast of matter, evolves as $\\delta_M\\propto a$ where $a$ is the scale factor of the Universe. We can parametrize deviations from this behaviour in terms of $\\gamma$ \\cite{Peebles1980,Linder2005a,Lee09a,Lee09b} through \\begin{eqnarray} \\gamma\\equiv\\frac{\\ln\\left[\\frac{{\\dot \\delta}_M}{ \\adotoa \\delta_M}\\right] }{\\ln \\Omega_M} \\end{eqnarray} where $\\Omega_M$ is the fractional density of matter, $\\adotoa=\\frac{\\dot a}{a}$ and overdots are derivatives with regards to conformal time, $\\tau$. For standard growth in the presence of a cosmological constant, one has that $\\gamma\\approx6/11$ to a very good approximation. This is not true over a wide range of values for $\\Omega_M$. In fact, in Figure \\ref{lambda} we can see that $\\gamma$ deviates from its early-universe asymptotic value as $\\Omega_M\\rightarrow 0$. A natural question to ask is how $\\gamma$ depends on different aspects of the Universe and how might use it to constrain dark energy and modifications to gravity. In this paper we will focus on a few of these properties. \\begin{figure}[htbp] \\begin{flushleft} \\vspace{-15pt} \\epsfig{figure=purelambda.ps,width=9cm} \\end{flushleft} \\vspace{-100pt} \\caption{The solid line is the growth parameter, $\\gamma$, for a $\\Lambda$CDM universe, as a function of $\\Omega_M$. For small values of $\\Omega_\\Lambda$, $\\gamma$ is well approximated by $6/11$ (dashed line) but there are deviations as $\\Omega_\\Lambda$ grows; we find errors of $0.7\\%$, $3.3\\%$ and $4.2\\%$ when $\\Omega_M=0.7$, $0.3$ and $0.05$.} \\label{lambda} \\vspace{-10pt} \\end{figure} One important property of the Universe is the {\\it equation of state} of dark energy, characterized by the constant (or function of time), $w$: \\begin{eqnarray} P_E=w\\rho_E. \\label{fluid_eq_state} \\end{eqnarray} $P_E$ and $\\rho_E$ are the pressure and energy densities of the dark energy. The function $w$ may be time varying and is related to the adiabatic speed of sound $c_a^2$ as \\begin{equation} c_a^2 = w - \\frac{\\dot{w}}{3\\adotoa(1+w)} \\end{equation} Another important property is {\\it gravitational slip}, $\\zeta$, which is normally defined to be \\begin{eqnarray} \\Phi-\\Psi\\equiv\\zeta\\Phi \\end{eqnarray} where we are taking a linearly perturbed metric in the conformal Newtonian gauge, \\begin{equation} ds^2=-a^2(1+2\\Psi)d\\tau^2+a^2(1-2\\Phi)d\\vec{x}^2. \\end{equation} Such a parametrization has been advocated in a number of papers on modified gravity \\cite{Bertschinger2006,BertschingerZukin2008,DanielEtAl2009,DanielEtAl2010,PogosianEtAl2010,Bean2009,ReyesEtAl2010} and it has been shown that it can lead to a number of observational effects. Albeit simple, and appealing, such a parametrization of slip is not necessarily general and, as we shall see in the next section, necessarily implies other non-trivial modifications to the gravitational sector. Such modifications are, in general, not explicitely acknowledged but may correspond to unexpected assumptions about any putative, underlying theory. Hence a more general assumption (at least within the context of 2$^{\\rm nd}$ order theories) would be that gravitational slip would depend on $\\Phi$ and ${\\dot \\Phi}$ (this is explained in more detail in section-\\ref{sec_consistency} and in ~\\cite{Skordis2008b}) Finally, we can define an {\\it effective Newton's constant} in the relativistic Newton-Poisson equation \\begin{eqnarray} \\nabla^2\\Phi=4\\pi a^2G_{eff}\\sum_X\\rho_X[\\delta_X+3(1+w_X)\\frac{\\dot a}{a}\\theta_X] \\nonumber \\label{eq_eff_Poisson} \\end{eqnarray} where $\\delta_X$ is the density contrast and $\\theta_X$ is the momentum of the cosmological fluid $X$ which has an equation of state $w_X$. We can define the dimensionless function \\begin{eqnarray} \\mu^2\\equiv\\frac{G}{G_{eff}} \\end{eqnarray} where $G$ is the \"bare\" Newton constant. It then makes sense to try to constrain ($\\gamma$, $w$, $\\zeta$ and $\\mu$) in the hope that it may be possible to shed light on a possible theory of modified gravity. Although there alternative proposals \\cite{KoivistoMota2006,Mota2007}, a number of groups have have pioneered the use of this simple parametrization of modified gravity (in terms of $\\zeta$, $\\mu^2$ or both): in \\cite{Bertschinger2006,BertschingerZukin2008} it was argued that gravitational slip might be a generic prediction for modified theories of gravity, in \\cite{DanielEtAl2009,DanielEtAl2010} it was shown that it would be possible to constrain it through cross correlations of the CMB with galaxy surveys and in \\cite{Hu2008} from the ISW effect; weak lensing has been proposed as a possible route for constraining these parameters \\cite{UzanBernardeau2001,ZhangEtAl2007,Schmidt2009a} with a tentative detection of modification being proposed in \\cite{Bean2009,BeanTangmatitham2010}. Much is expected from applying these methods to future ambitious experiments that will map out the large scale structure of the Universe. Indeed constraints of General Relativity are a core element of the science that could be extracted from the Euclid experiment~\\cite{RefregierEtAl2010}. Given that such an approach is phenomenological, the general attitude has been to leave these parameters completely free. There is merit to such an approach in that one isn't restricting oneself to a particular theory and hence constraints will be general. It is true however that is is possible to idenitfy (reasonably general) consistency conditions for ($\\gamma$, $w$, $\\zeta$ and $\\mu$), contingent on specific physical assumptions. In this paper we state these assumptions and present restrictions on ($\\gamma$, $w$,$\\zeta$ and $\\mu$). We shall use the formalism first proposed by one of us\\cite{Skordis2008b} which spells out how to build consistent modifications to gravity. This paper is structured as follows. In Section \\ref{formalism} we recap the formalism presented in \\cite{Skordis2008b} and relate it to the parameters we wish to study phenomenologically. We discuss how the consistency conditions reduce the freedom to choose arbitrary ($\\gamma$, $\\zeta$, $\\mu^2$). In Section \\ref{small} we implement the consistency conditions and find a relationship between the parameters by looking at the evolution equation for the density contrast in matter for small wavelengths. In doing so, we find find analytic expressions for the relationships and briefly assess the range of scale to which they are applicable. In Section \\ref{General} we find analytic expressions for $\\gamma$ to 2$^{\\rm nd}$ order for a general parametrization which is consistent with the Parametrized Post Newtonian approximation on small scales.In Section \\ref{Discussion} we discuss the generality of the results and how they may be extended to other, more exotic models. ", "conclusions": "\\label{Discussion} Let us briefly recap what we have done. The main point of this paper is that, when introducing modifications to gravity in linear perturbation theory, one must take into account the consistency conditions in the field equations. These necessarily lead to restrictions in the form of the modifications that can be introduced. Most notably, and within the context of second order theories, this means that if one wishes to include modifications to the Newton-Poisson equation, then one {\\it cannot consider the simplified gravitational slip}, $\\Psi=\\zeta\\Phi$, and must include an extra term such that $\\Psi=\\zeta\\Phi+(g/k){\\dot \\Phi}$ where $\\frac{g}{k}\\adotoa=1-G_{eff}/G_0$. If we wish to construct a proper \"Parametrized Post Friedmanian\" approach to modified gravity, any parameter we introduce must be independent of the environment or initial conditions in the perturbations. The only way to do this is to use the parametrization we are advocating. To our knowledge, all attempts at studying cosmological deviations from general relativity have ignored this and hence it is unclear what class of theories they map onto and which types of theories are being constrained. Having taken this point on board, we have found the expression for the growth parameter on small scales in terms of both the gravitational slip, $\\zeta$ and the modified Newton constant, $\\mu^2=1- \\frac{g}{k} \\adotoa$. Given a set of cosmological constraints on $\\gamma$ and its dependence on $\\Omega_M$, it is now straightforward to calculate constraints on $\\zeta$ and $g$. The growth parameter is given by equation (\\ref{gamma_param}) which can be seen as a Taylor expansion in terms of $ 1-\\Omega_M$. The coefficients in this expansion, $\\gamma_0$, $\\gamma_1$ and $\\gamma_2$ can be expressed in terms of the equation of state, $w$ (see Equation \\ref{w_param}), $\\zeta$ (see Equation \\ref{zeta_param}) and $g$ (see Equation \\ref{g_param}) by using equations (\\ref{UV_param}), followed by equations (\\ref{Ytot_param}) and finally equations (\\ref{gammatot_param}). With these relations in hand, it is now possible to use cosmological observations to place constraints on theories of modified gravity. In this paper we have focused on small scales (by which we mean between 1 and 200h$^{-1}$Mpc), scales that should be probed by redshift space distortion measurements, galaxy power spectra and weak lensing. Furthermore, we can now do this consistently, relating modifications in the growth rate with changes in the gravitational slip. This is of particular importance when considering weak lensing where observations probe $\\Phi+\\Psi$. It is also clear from our analysis that we have come up against the limitations of the $\\gamma$ parametrization: it is useful and effective on very small scales but not on scales comparable to the cosmological horizon. On those scales, one should be using the full set of field equations. We therefore do not advocate using our fitting formula to the growth on larger scales such as would be probed by the Integrated Sach-Wolfe. How general is this method? We have declared from the outset, the class of theories that we are considering. They must be metric, with 2$^{\\rm nd}$ order equations and satisfy gauge form-invariance. From what we have learnt about modifications of gravity, these seem a reasonable set of conditions to apply- they lead to theories which are less likely to be marred by gross instabilities either at the classical or quantum level. We should point out that all other attempts at developing such a parametrization have {\\it implicitely} made these assumptions although have not necessarily done so self-consistently. It is possible to extend this analysis beyond the scope of these theories. If we are to go beyond 2$^{\\rm nd}$ order, one must include terms in ${\\ddot \\Phi}$ or even higher. The Bianchi conditions will, again, impose a set of constraints on the coefficients of these terms and should allow a similar type of analysis. Two well studied theores are worth mentioning. $F(R)$, $F(R^{\\mu\\nu} R_{\\mu\\nu})$, etc, theories come with up-to four time derivatives in the field equations. Thus they do not fall directly within the methods of this paper but do under the general scheme outlined in~\\cite{Skordis2008b}. In this case one would have to include terms involving $\\dot{\\Phi}$, $\\ddot{\\Phi}$, $\\Psi$ and $\\dot{\\Psi}$ in to the $G_{00}$ and $G_{0i}$ Einstein equations, while the $G_{ij}$ equations would need $\\dddot{\\Phi}$ and $\\ddot{\\Psi}$ in addition. Theories with higher derivatives are a subject that warrants further investigation and have yet to be properly incorporated in any parametrized modifications of standard general relativity. One other theory, studied extensively is the DGP theory \\cite{DvaliGabadadzePorrati2000d}. In this case, only two time derivatives are present in the field equations and just like our frame, DGP contains two non-metric dynamical degrees of freedom, which can be effectively written ad $\\delta_E$ and $\\theta_E$. However, our framework cannot encompass DGP because DGP cannot be written as a generalized fluid as we have assumed of dark energy in this work (hence it is not a failure of our use of $\\delta U_{ab}$). Nevertheless, our $\\gamma$ parameterization in powers of $\\Omega_E$ is still valid, and indeed needed. In the case of DGP we find that \\begin{equation} \\gamma = \\frac{11}{16} + \\frac{7}{5632} \\Omega_E - \\frac{93}{4096} \\Omega_E^2 \\end{equation} gives an error on $\\gamma$ around $5\\%$ at $\\Omega_M <0.1$, dropping to $2\\%$ at $\\Omega_M\\sim 0.2$ and $<1\\%$ for larger values of $\\Omega_M$. The error on the corresponding density contrast at those values of $\\Omega_M$ is $<2\\%$, $\\sim 1\\%$ and $<0.5\\%$ respectively. Notice how the coefficients are entirely fixed and {\\emph do not} depend on the only free parameter of the theory, namely the scale $r_c$. Rather $r_c$ comes to play a role only through $\\Omega_E = \\frac{1}{H r_c}$. \\vspace{-10pt}" }, "1003/1003.1078_arXiv.txt": { "abstract": "{We review new results on strong and electroweak interactions, flavour physics, cosmic rays and cosmology, which were presented at this conference, focussing on physics beyond the Standard Models. Special emphasis is given to the Higgs sector of the Standard Model of Particle Physics and recent results on high-energy cosmic rays and their implications for dark matter.} \\FullConference{European Physical Society Europhysics Conference on High Energy Physics\\\\ July 16-22, 2009\\\\ Krakow, Poland} \\begin{document} ", "introduction": "At this conference, the Standard Models of Particle Physics and Cosmology have again been impressively confirmed. In the experimental talks on strong interactions, electroweak precision tests, flavour and neutrino physics and searches for `new physics' no significant deviations from Standard Model predictions have been reported. Also in Astrophysics, where unexpected results in high-energy cosmic rays were found, conventional astrophysical explanations of the new data appear to be sufficient. In Cosmology, we have entered an era of precision physics with theory lagging far behind. Given this situation, one faces the question: What are the theoretical and experimental hints for physics beyond the Standard Models, and what discoveries can we hope for at the LHC, in non-accelerator experiments, and in astrophysical and cosmological observations? In the following I shall summarize some results of this conference using this question as a guideline. Particular emphasis will therefore be given to the Higgs sector of the Standard Model, the ``topic number one'' at the LHC, and the recent results in high-energy cosmic rays, which caused tremendous excitement during the past year because of the possible connection to dark matter. The Standard Model of Particle Physics is a relativistic quantum field theory, a non-Abelian gauge theory with symmetry group \\begin{equation} G_{\\mathrm{SM}} = SU(3)\\times SU(2)\\times U(1) \\end{equation} for the strong and electroweak interactions, respectively. Three generations of quarks and leptons with chiral gauge interactions describe all features of matter. The current focus is on \\begin{itemize} \\item Precision measurements and calculations in QCD \\item Heavy ions and nonperturbative field theory \\item Electroweak symmetry breaking, with the key elements: top-quark, W-boson and Higgs bosons \\item Flavour physics and neutrinos. \\end{itemize} The cosmological Standard Model is also based on a gauge theory, Einstein's theory of gravity. Together with the Robertson-Walker metric this leads to Friedmann's equations. Within current errors, the universe is known to be spatially flat, and its expansion rate is increasing. Most remarkably, its energy density is dominated by `dark matter' and `dark energy'. The desire to disentangle the nature of dark matter and dark energy, and to understand their possible connection to particle physics is the main driving force in observational cosmology today. On the theoretical frontier, string theory is the main theme, despite the fact that after more than thirty years of research it still has not become a falsifiable theory. Nevertheless, string theory has inspired many extensions of the Standard Model, which will be tested at the LHC and it has stimulated interesting models for the early universe which can be probed by cosmological observations. String theory goes beyond field theory by replacing point-interactions of particles by nonlocal interactions of strings. In this way it has also become a valuable tool to analyze strongly interacting systems of particles at high energies and high densities. ", "conclusions": "" }, "1003/1003.3885_arXiv.txt": { "abstract": "{} % { \\textit{Fermi} can measure the spectral properties of gamma-ray bursts over a very large energy range and is opening a new window on the prompt emission of these energetic events. Localizations by the instruments on \\textit{Fermi} in combination with follow-up by \\textit{Swift} provide accurate positions for observations at longer wavelengths leading to the determination of redshifts, the true energy budget, host galaxy properties and facilitate comparison with pre-Fermi bursts. } {Multi-wavelength follow-up observations were performed on the afterglows of four bursts with high energy emission detected by {\\textit{Fermi}/LAT} : GRB\\,090323, GRB\\,090328, GRB\\,090510 and GRB\\,090902B. They were obtained in the optical/near-infrared bands with GROND mounted at the MPG/ESO 2.2\\ m telescope and additionally of GRB\\,090323 in the optical with the 2\\ m telescope in Tautenburg, Germany. Three of the events are classified as long bursts while GRB\\,090510 is a well localized short GRB with GeV emission. In addition, host galaxies were detected for three of the four bursts. Spectroscopic follow-up was initiated with the VLT for GRB\\,090328 and GRB\\,090510. } { The afterglow observations in 7 bands are presented for all bursts and their host galaxies are investigated. Knowledge of the distance and the local dust extinction enables comparison of the afterglows of LAT-detected GRBs with the general sample. The spectroscopic redshifts of GRB\\,090328 and GRB\\,090510 were determined to be $z=0.7354\\pm0.0003$ and $z=0.903\\pm0.001$ and dust corrected star-formation rates of $4.8$ M$_\\odot$ yr$^{-1}$ and $0.60$ M$_\\odot$ yr$^{-1}$ were derived for their host galaxies, respectively. } { The afterglows of long bursts exhibit power-law decay indices ($\\alpha$) from less than 1 to $\\sim$2.3 and spectral indices ($\\beta_{\\rm opt}$) values from 0.65 to $\\sim$1.2 which are fairly standard for GRB afterglows. Constraints are placed on the jet half opening angles of $\\lesssim 2.1^\\circ$ to $\\gtrsim 6.4^\\circ$, which allows limits to be placed on the beaming corrected energies. These range from $\\lesssim $ $5 \\times 10^{50}$ erg to the one of the highest values ever recorded, $\\gtrsim 2.2 \\times 10^{52}$ erg for GRB~090902B, and are not consistent with a standard candle. The extremely energetic long \\textit{Fermi} bursts have optical afterglows which lie in the top half of the brightness distribution of all optical afterglows detected in the \\textit{Swift} era or even in the top 5 \\% if incompleteness is considered. The properties of the host galaxies of these LAT detected bursts in terms of extinction, star formation rates and masses do not appear to differ from previous samples. } ", "introduction": "The follow-up of gamma-ray bursts (GRBs) detected by the \\textit{Swift} satellite \\citep{2004ApJ...611.1005G} has led to the determination of the distance scale for a large sample of bursts. The Burst Alert Telescope \\citep[BAT,][]{2005SSRv..120..143B} is sensitive in the energy range 15$-$150~keV and has good localization capabilities with typical uncertainties in the arcminute range. Rapid follow-up by \\textit{Swift's} narrow field instruments in the X-rays \\citep[XRT,][]{2005SSRv..120..165B} and optical/UV \\citep[UVOT,][]{2005SSRv..120...95R} have lead to the arcsecond localizations required for ground-based observers and in turn to spectroscopic redshift measurements of a large sample of GRBs \\citep[e.g.,][]{2009arXiv0907.3449F} and investigation of their host galaxies \\citep[e.g.,][]{2009ApJ...691..182S,2009AJ....138.1690P}. To date distances to $\\sim$ 200 GRB sources have been established with redshifts ranging from $z = 0.0085$ \\citep[GRB 980425:][]{1998IAUC.6896....3W,1998Natur.395..670G} to $z\\sim8.2$ \\citep[GRB090423:][]{2009arXiv0906.1577T,2009arXiv0906.1578S}. The BAT has a narrow spectral range and is not able to determine the spectral parameters of the prompt emission of GRBs over a broad energy range. Since the launch of the \\textit{Fermi} Gamma-Ray Space Telescope \\citep{1994NIMPA.342..302A,1996SPIE.2806...31M} there is now the possibility to investigate the spectrum over seven decades in energy. These events can be localized by instruments from \\textit{Fermi} and followed up by \\textit{Swift} and ground-based obervatories, enabling investigation of their afterglows and host properties and facilitating comparison to a general sample. In some fortuitous cases instruments on the \\textit{Fermi} and \\textit{Swift} satellites may trigger on the same event with high-energy emission enabling both rapid follow-up and broadband prompt emission coverage (e.g. GRB~090510, see Section 2.3). The high-energy spectral properties of GRBs were investigated previously by instruments on-board the \\textit{Compton Gamma-ray Observatory}, however the redshifts of these events are unknown. The spectra of these GRBs were described by an extrapolation of the low energy spectra to energies $>$100 MeV \\citep[e.g.,][]{1994A&A...285..161H}, and in some cases an additional component at high-energies \\citep{2003Natur.424..749G,2008ApJ...677.1168K} and long-lived GeV emission \\citep{1995Natur.374...94H}. Recently a photometric redshift of $z=1.8^{+0.4}_{-0.3}$ was reported by \\citet{2008A&A...491L..29R} for GRB~080514B which was detected at energies up to 300~MeV by instruments on the AGILE satellite \\citep{2008A&A...491L..25G}. \\textit{Fermi} hosts two instruments that detect GRBs: the Gamma-ray Burst Monitor \\citep[GBM:][]{2009arXiv0908.0450M} sensitive to photons from $\\sim$8 keV to $\\sim$40 MeV, and the Large Area Telescope with a spectral coverage from $\\sim$ 30 MeV to $\\sim$100 GeV \\citep[LAT:][]{2009ApJ...697.1071A}. Together they provide valuable information on the spectral properties of the prompt emission \\citep[e.g.][]{2009arXiv0908.1832F,2009arXiv0910.4192F,090902B_PAPER,2009Sci...323.1688A}, the presence or absence of intrinsic spectral cut-offs or those due to the optical depth of the universe to high-energy $\\gamma$-rays due to pair production on infrared diffuse Extragalactic Background Light \\citep[e.g.,][]{2002ApJ...566..738D,2003A&A...407..791M,2004A&A...413..807K,2005Natur.438...45K,2006ApJ...648..774S,2008A&A...487..837F,2009arXiv0905.1115F,090902B_PAPER} and test for quantum gravity effects \\citep[e.g.,][]{1998Natur.393..763A,2008ApJ...673..972S,2005LRR.....8....5M,2009PhRvD..80h4017A,2009Sci...323.1688A,2009arXiv0908.1832F}. Up to the end of January 2010, fourteen GRBs have been detected by both instruments and localized by the LAT : GRB 080825C \\citep[][]{LAT_080825C,2009arXiv0910.4192F}, GRB 080916C \\citep[][]{2009Sci...323.1688A}, GRB 081024B \\citep[][]{2008GCN..8407....1O}, GRB 081214 \\citep[][]{2008GCN..8723....1W}, GRB 090217 \\citep[][]{LAT_090217}, GRB 090323 \\citep[][]{2009GCN..9021....1O}, GRB 090328 \\citep[][]{2009GCN..9044....1M}, GRB 090510 \\citep[][]{2009GCN..9334....1O,2009arXiv0908.1832F}, GRB 090626 \\citep[][]{090626_LAT}, GRB 090902B \\citep[][]{LAT_090902B,090902B_PAPER}, GRB 090926A \\citep[][]{GBM_090926,LAT_090926}, GRB 091003 \\citep[][]{McEnery09_GCN9985}, GRB 091031 \\citep[][]{Palma09_GCN10163}, and GRB 100116A \\citep[][]{McEnery10_GCN10333}. Twelve are long duration GRBs and two have reported durations compatible with the short burst class (GRBs~081024B and 090510). The redshifts of five of these bursts have been determined and range from $z=0.736$ to $z=4.35\\pm0.15$ \\citep{2009A&A...498...89G,2009GCN..9026....1U,2009GCN..9028....1C,2009GCN..9053....1C,2009_arne,Z_090902B,Xshooter_090926}. Among the most impressive of these events are GRB 080916C \\citep{2009Sci...323.1688A} and GRB 090510 \\citep{2009arXiv0908.1832F}. GRB~080916C is a long bright GRB for which the prompt emission spectrum could be fit over six decades in energy by the empirical Band function \\citep[{see}][{for detailed results}]{2009Sci...323.1688A}. The burst was found to have a high photometric redshift of $z=4.35\\pm0.15$ via ground-based follow up observations \\citep{2009A&A...498...89G}. The distance information enabled the determination of the rest frame properties, energetics, and the placing of lower limits on the bulk Lorentz factor of the outflow ($\\gtrsim 900-1100$) \\citep{2009Sci...323.1688A,2009A&A...498...89G}. GRB 090510 is a short burst with a T$_{\\rm 90}$ of 2.1 s and from which a 31~GeV photon was detected \\citep{2009arXiv0908.1832F}. The distance determination \\citep{2009_arne} and high energy emission allow a limit to be placed on photon dispersion \\citep{2009arXiv0908.1832F}. Moreover, very recently a 33 GeV photon was detected from GRB~090902B \\citep{090902B_PAPER} and this burst is at redshift $z = 1.822$ \\citep{Z_090902B}. Clearly, the distance determination via optical spectroscopy is crucial to the interpretation of these events. We report on the optical and near-infrared (NIR) observations of the afterglow of four bursts with high energy emission, GRB\\,090323, GRB\\,090328, GRB\\,090510 and GRB\\,090902B and compare them to the sample of GRB afterglows to date. Furthermore we also report on the spectroscopic redshift determination for GRB~090328 and GRB~090510. Throughout the paper, we adopt concordance $\\Lambda$CDM cosmology ($\\Omega_M=0.27$, $\\Omega_{\\Lambda}=0.73$, $H_0=71$~km/s/Mpc), and the convention that the flux density of the GRB afterglow can be described as $F_\\nu (t)\\propto \\nu^{-\\beta}t^{-\\alpha}$ ", "conclusions": "We have presented the multi-colour afterglow observations of four bursts with high energy emission including three long and one short burst. In addition, we present spectroscopic observations and redshift determinations of two of these bursts. It is now possible for the first time to combine the high energy information from the prompt emission with the afterglow, host galaxy and redshifts of these sources. Follow-up of GRBs detected by \\textit{Fermi} with the seven band imager, GROND, allows simultaneous determination of the temporal, $\\alpha$, and spectral, $\\beta_{\\rm opt}$, indices of the burst afterglows. The long bursts exhibit power-law decay indices ($\\alpha$) from less than 1 to $\\sim$2.3 and spectral indices ($\\beta_{\\rm opt}$) from 0.65 to $\\sim$1.2 which are fairly standard for GRB afterglows. Moreover, an estimation of the jet break time is vital to determine the broadband properties and energetics of these events. The redshifts of the long bursts span the range from 0.7354 to 3.57 and the beaming corrected energies differ by a factor of at least $\\sim$30 with a value of $E_{\\rm{\\gamma}}$ $\\gtrsim 2.2 \\times 10^{52}$ {erg} for GRB~090902B and are not compatible with a standard candle. Interestingly, the higher redshift bursts detected by the LAT are very luminous in the keV-MeV region of the prompt spectrum, have high $E_{\\gamma, \\rm{iso}}$ and $E_{\\rm{\\gamma}}$ and also exceptionally bright afterglows. The lower redshift GRB~090328 is the exception in this respect, with a more standard $E_{\\gamma, \\rm{iso}}$ as compared to \\textit{Swift} GRBs. The host galaxies of GRB~090328 and GRB~090510 fit well within the distributions of the long and short burst host galaxies respectively. Future photometric and spectroscopic follow-up of these rare \\textit{Fermi} bursts with high-energy emission is crucial to further investigate the energetics and host galaxies of a larger sample." }, "1003/1003.2417_arXiv.txt": { "abstract": "In this paper, we present and analyse optical photometry and spectra of the extremely luminous and slowly evolving Type~Ia supernova (SN~Ia) 2009dc, and offer evidence that it is a super-Chandrasekhar mass (SC) SN~Ia and thus had a SC white dwarf (WD) progenitor. Optical spectra of SN~2007if, a similar object, are also shown. \\dc\\ had one of the most slowly evolving light curves ever observed for a SN~Ia, with a rise time of \\about23~days and $\\Delta m_{15} (B)=0.72$~mag. We calculate a lower limit to the peak bolometric luminosity of \\about$2.4\\times10^{43}$~\\ergps, though the actual value is likely almost 40\\% larger. Optical spectra of \\dc\\ and \\snif\\ obtained near maximum brightness exhibit strong \\ion{C}{II} features (indicative of a significant amount of unburned material), and the post-maximum spectra are dominated by iron-group elements. All of our spectra of \\dc\\ and \\snif\\ also show low expansion velocities. However, we see no strong evidence in \\dc\\ for a velocity ``plateau'' near maximum light like the one seen in \\snif\\ \\citep{Scalzo10}. The high luminosity and low expansion velocities of \\dc\\ lead us to derive a possible WD progenitor mass of more than 2~\\msun\\ and a \\nic\\ mass of about 1.4--1.7~\\msun. We propose that the host galaxy of \\dc\\ underwent a gravitational interaction with a neighboring galaxy in the relatively recent past. This may have led to a sudden burst of star formation which could have produced the SC~WD progenitor of \\dc\\ and likely turned the neighboring galaxy into a ``post-starburst galaxy.'' No published model seems to match the extreme values observed in \\dc, but simulations do show that such massive progenitors can exist (likely as a result of the merger of two WDs) and can possibly explode as SC~SNe~Ia. ", "introduction": "\\label{s:intro} Type~Ia supernovae (SNe~Ia) are differentiated from other types of SNe by the absence of hydrogen and the presence of broad absorption from \\ion{Si}{II} $\\lambda$6355 in their optical spectra \\citep[for a review see][]{Filippenko97}. SNe~Ia have been used to measure cosmological parameters to high precision \\citep[e.g.,][]{Kowalski08,Hicken09,Kessler09,Amanullah10}, as well as to discover the accelerating expansion of the Universe \\citep{Riess98,Perlmutter99}. Broadly speaking, SNe~Ia are the result of thermonuclear explosions of C-O white dwarfs (WDs) resulting from either the accretion of matter from a nondegenerate companion star \\citep[e.g.,][]{Whelan73} or the merger of two degenerate objects \\citep[e.g.,][]{Iben84,Webbink84}. However, the nature of the companion and the details of the explosion itself are both still quite uncertain. The cosmological utility of SNe~Ia comes from the fact that they are standardizable candles (i.e., their luminosities at peak can be calibrated). This naively seems reasonable since SNe~Ia should all have the same amount of fuel and the same trigger point: they should all explode when a WD nearly reaches the Chandrasekhar mass of \\about1.4~\\msun. In 2003, SNLS-03D3bb (also known as \\fg) was discovered \\citep{Howell06} and was shown to be a SN~Ia that was overluminous by about a factor of 2, and had a slowly declining light curve, quite low expansion velocities, and unburned material present in near-maximum light spectra. This last observation implies that a layer of carbon and oxygen from the progenitor existed on top of the burned silicon layer. In SNe~Ia, carbon is usually extremely weak or completely absent (even at very early times) in optical and infrared (IR) spectra, although it has been (sometimes tentatively) identified in a few other cases \\citep[e.g.,][]{Branch03,Marion06,Thomas07,Foley10}. \\citet{Howell06} suggest that all of the oddities seen in \\fg\\ could be explained if its progenitor WD had a mass greater then the canonical upper limit for WDs of \\about1.4~\\msun, a so-called ``super-Chandrasekhar mass'' (SC) WD. \\citet{Hicken07} then presented data on \\gz\\ which shared some of the strange properties of \\fg, and they concluded that \\gz\\ must have come from a WD merger leading to a SC~SN~Ia. Recently, \\citet{Scalzo10} published observations of \\snif\\ which is yet another example of this emerging class of possible SC~SNe~Ia. They not only observed low expansion velocities, but they also saw a plateau in the expansion velocity near maximum-brightness which they interpret as evidence for the SN ejecta running into a shell of material. The accurately determined total mass of the \\snif\\ system is well above the Chandrasekhar mass. Recently, it was pointed out that there might be another member of this SC~SN~Ia class, \\dc\\ \\citep{Harutyunyan09,Marion09,Yamanaka09}. \\dc\\ was discovered $15\\farcs8$ west and $20\\farcs8$ north of the nucleus of the S0 galaxy UGC~10064 by \\citet{Puckett09} on 2009~Apr.~9.31 (UT dates are used throughout this paper), though in \\S\\ref{ss:lc} we will show a detection of the SN \\about5~days earlier. It is located at $\\alpha_{\\rm J2000} = 15^{\\rm h}51^{\\rm m}12\\fs12$ and $\\delta_{\\rm J2000} = +25\\degr42\\arcmin28\\farcs0$; the SN and its host are shown in Figure~\\ref{f:galaxy}. No object was visible in our data at the position of the SN on 2009~Mar.~28 to a limiting magnitude of \\about19.5, so the actual explosion date was almost certainly between Mar.~28 and Apr.~4. \\begin{figure}% \\centering \\includegraphics[width=8.5cm]{sn2009dc_kait} \\caption{KAIT image of \\dc\\ and its host galaxy, UGC~10064. The field of view is 6.7\\arcmin$\\times$6.7\\arcmin. \\dc\\ is labeled along with the comparison stars used for differential photometry. The scale of the image is marked in the top right; north is up and east is to left. The SN is sufficiently far from its host galaxy that template subtraction was not performed when conducting photometry.}\\label{f:galaxy} \\end{figure} \\citet{Harutyunyan09} obtained a spectrum of \\dc\\ one week after the announced discovery date and showed it to be a SN~Ia before maximum light. They also noted that \\dc\\ spectroscopically resembled the SC~SN~Ia candidate SN~2006gz \\citep{Hicken07} at this time, but with a much lower expansion velocity as derived from the \\ion{Si}{II} $\\lambda$6355 absorption feature. Three days later, nearly simultaneous optical and IR spectra were obtained by \\citet{Marion09}, covering a wavelength range of 0.36 to 1.3 $\\mu$m. They again note some similarities to (as well as differences from) SN~2006gz, as well as similarities with the possible SC~SN~Ia 2003fg \\citep{Howell06}. \\citet{Yamanaka09} presented early-time optical and near-IR observations of \\dc\\ and showed that it did indeed share many of the properties seen in both \\fg\\ and \\gz, suggesting that it too is a possible SC~SN~Ia. Furthermore, \\citet{Tanaka09} published spectropolarimetry of \\dc\\ which indicated that the explosion was quite spherically symmetric. In this paper we present and analyse our own optical photometric and spectroscopic data for \\dc\\ (as well as spectra of \\snif) with the goal of more definitively answering the question of whether \\dc\\ and \\snif\\ were truly SC~SNe~Ia. We show some of the most precise data on the rise time of a possible SC~SN~Ia ever published, as well as some of the latest-time photometric and spectral observations of a member of this class. In \\S\\ref{s:obs} we describe our observations and data reduction, and in \\S\\ref{s:analysis} we discuss our analysis of the photometry and spectra of \\dc\\ (and its host galaxy) and \\snif. We attempt to robustly calculate physical parameters of \\dc\\ using a variety of methods and compare them to theoretical predictions in \\S\\ref{s:discussion}. Finally, in \\S\\ref{s:conclusions}, we summarize our conclusions and ruminate about the future. ", "conclusions": "\\label{s:conclusions} In this paper we have presented and analysed optical photometry and spectra of \\dc\\ and \\snif, both of which are possibly SC~SNe~Ia. Our photometric and spectral data on \\dc\\ constitute one of the richest datasets ever published on a SC~SN~Ia candidate. Our well-sampled light curve follows \\dc\\ from about 1 week before maximum brightness until about 5~months past maximum, and shows that \\dc\\ is one of the slowest photometrically evolving SNe~Ia ever observed. We derive a rise time of 23~days and $\\Delta m_{15} (B)=0.72$~mag, which are two of the most extreme values for these parameters ever seen in a SN~Ia. Assuming no host-galaxy reddening, we derive a peak bolometric luminosity of about $2.4\\times10^{43}$~\\ergps, though this is almost certainly an underestimate since we observe strong evidence for at lease {\\it some} host reddening. Using our nonzero values for $E(\\bv)_\\textrm{host}$, the peak bolometric luminosity increases by about 40\\%--200\\%. Spectroscopically, \\dc\\ also evolves relatively slowly. Strong \\ion{C}{II} absorption features (which are rarely observed in SNe~Ia) are seen in the spectra near maximum brightness, implying a significant amount of unburned fuel from the progenitor WD in the outer layers of the SN ejecta. \\ion{Si}{II} absorption also appears in our spectra of \\dc\\ and remains visible even 2~months past maximum. Our post-maximum spectra are dominated by a forest of IGE features and, interestingly, resemble spectra of the peculiar SN~Ia 2002cx. Finally, the spectra of \\dc\\ all show very low expansion velocities at all layers (i.e., unburned carbon, IMEs, and IGEs) as compared to other SNe~Ia. This may be explained by a massive WD progenitor which consequently has a large binding energy. Even though the expansion velocities are small, we see no strong evidence in \\dc\\ for a velocity ``plateau'' near maximum light like the one seen in \\snif\\ \\citep{Scalzo10}. Using various luminosity and energy arguments, we calculate that the progenitor of \\dc\\ is possibly a SC~WD with a mass greater than \\about2~\\msun, and that at least \\about1~\\msun\\ of \\nic\\ was likely formed in the explosion (though the most probable value is in the range 1.4--1.7~\\msun). These values are larger than (or about as large as) those calculated for any other SN~Ia ever observed. We propose that the host galaxy of \\dc\\ underwent a gravitational interaction with a nearby galaxy (UGC~10063) in the relatively recent past, and that this could have induced a sudden burst of star formation which may have given rise to the progenitor of \\dc\\ and turned UGC~10063 into the ``post-starburst'' galaxy that we observe today. We also compare our observed quantities for \\dc\\ to theoretical models, and while no model seems to match or explain every aspect of \\dc, simulations show that SC~WDs with masses near what we calculate for the progenitor of \\dc\\ can possibly form, likely from the merger of two WDs. Furthermore, models of extremely luminous SNe~Ia which employ a Chandrasekhar-mass WD progenitor cannot explain our observations of \\dc. Thus, taking all of these extreme values into account, we conclude that \\dc\\ is very likely a SC~SN~Ia. As mentioned previously, many of the observed peculiarities of \\dc\\ are also seen in \\fg\\ and \\snif. Therefore, we concur with \\citet{Howell06} and \\citet{Scalzo10} that both \\fg\\ and \\snif\\ (respectively) are also probably SC~SNe Ia. However, given their fairly normal expansion velocities and relative weakness (or even absence) of \\ion{C}{II} features near maximum brightness, it seems that \\gz\\ and SN~2004gu are less likely to be SC~SNe~Ia. New large transient searches such as Pan-STARRS \\citep{Kaiser02} and the Palomar Transient Factory \\citep[][]{Rau09,Law09} will probably find many SC or other super-luminous SNe~Ia in the near future. Since it seems that they cannot be standardized in the same way as most SNe~Ia, they will need to be handled separately or ignored in cosmological surveys which will use large numbers of SNe~Ia. However, the simulations of \\citet{Chen09} show that donor stars with lower metallicities (e.g., Population II stars) are less likely to form WDs with masses greater than 1.7~\\msun\\ than higher metallicity stars. Thus, it is possible that contamination levels from SC~SNe~Ia, which are already rare at low redshifts (i.e., average metallicity), may be relatively small in medium or high-redshift surveys." }, "1003/1003.5001_arXiv.txt": { "abstract": "{ We present the results of a deep spectral analysis of all Swift observations of Mrk 421 between April 2006 and July 2006, when it reached its highest X-ray flux recorded until the end of 2006. We completed this data set with other historical X-ray observations. We used the full data set to investigate the correlation between the spectral parameters.\\\\ We found a signature of stochastic acceleration in the anticorrelation between the peak energy ($E_p$) of the spectral energy distribution (SED) and the spectral curvature parameter ($b$). We found signature of energetic budget of the jet in the correlation between the peak flux of the SED ($S_p$) and $E_p$. Moreover, using simultaneous \\swf~ UVOT/XRT/BAT data, we demonstrated, that during the strongest flares, the UV-to-X-ray emission from \\mrk~ requires that the curved electron distribution develops a low energy power-law tail.\\\\ The observed spectral curvature and its anticorrelation with $E_p$ is consistent with both stochastic acceleration or energy-dependent acceleration probability mechanisms, whereas the power-law slope of \\xrt-\\uvt~ data is close to that inferred from the GRBs X-ray afterglow and in agreement with the \\textit{universal} first-order relativistic shock acceleration models. This scenario implies that magnetic turbulence may play a twofold role: spatial diffusion relevant to the first order process and momentum diffusion relevant to the second order process.\\\\ } \\FullConference{The Extreme sky: Sampling the Universe above 10 keV - extremesky2009,\\\\ October 13-17, 2009\\\\ Otranto (Lecce) Italy} \\begin{document} \\begin{figure}[t] \\begin{center} \\includegraphics[angle=0,width=9cm]{Ep-b.pdf} \\end{center} \\caption{Scatter plot of the curvature ($b$) vs. $E_p$. Solid red circles represent data from \\cite{Trama2007b}. Black boxes represent Swift data from \\cite{Tramacere2009}, without the cooling-dominated events. Empty circles represent the whole XRT data set presented in \\cite{Tramacere2009}.} \\label{fig:Epb} \\end{figure} ", "introduction": "BL Lac objects are Active Galactic Nuclei (AGNs) characterized by a polarised and highly variable nonthermal continuum emission extending from radio to $\\gamma$-rays. In the most accepted scenario, this radiation is produced within a relativistic jet that originates in the central engine and points close to our line of sight. The relativistic outflow has a typical bulk Lorentz factor of $\\Gamma \\approx 10$, hence the emitted fluxes, observed at an angle $\\theta$, are affected by a beaming factor $ \\delta = 1/(\\Gamma (1 - \\beta \\cos\\theta ))$. \\\\ The Spectral Energy Distribution (SED) of these objects has a typical two-bump shape. According to current models, the lower-frequency bump is interpreted as synchrotron (S) emission from highly relativistic electrons with Lorentz factors $\\gamma$ in excess of $10^2$. This component peaks at frequencies ranging from the IR to the X-ray band. In the framework of the Synchrotron Self Compton (SSC) emission mechanism, the higher-frequency bump can be attributed to inverse Compton scattering of synchrotron photons by the same population of relativistic electrons that produce the synchrotron emission \\citep{Jones1974}.\\\\ \\mrk~ is classified as a High energy peaked BL Lac (HBL) \\citep{Padovani1995} because its synchrotron emission peak ranges from a fraction of a keV to several keV. With its redshift $z$ = 0.031, it is among the closest and most well studied HBLs. In spring/summer 2006, \\mrk~ reached its highest X-ray flux recorded until that time. The peak flux was about 85 milli-Crab in the 2.0-10.0 keV band, and corresponded to a peak energy of the spectral energy distribution (SED) that was often at energies higher than 10 keV.\\\\ In this paper we use \\swf~ UV and X-ray data \\citep{Tramacere2009} of the 2006 flaring activity, completed with historical X-ray observations, to interpret the correlation between the spectral parameters in terms of acceleration processes and energetic budget of the jet. We remand the reader to \\citet{Tramacere2009} that paper for a more complete picture. \\begin{figure}[t] \\begin{center} \\includegraphics[angle=0,width=9cm]{Sp-Ep.pdf} \\end{center} \\caption{Scatter plot of $L_p$ vs $E_p$. Solid red circles represent data from \\cite{Trama2007b} spanning from 1997 to 2005. Solid blue boxes represent Swift data from \\cite{Tramacere2009}. The green dashed dotted line represents the best fit using a power-law and the black dashed line represents the best fit using a broken power law.} \\label{fig:SpEp} \\end{figure} ", "conclusions": "We have shown that the X-ray flaring activity of Mrk 421 results in a complex spectral evolutions due to drastic changes in the electron energy distribution probably related to a complex acceleration scenario. Indeed, in our analysis we found both signatures of first and second order acceleration processes acting at the same time. The $E_p-b$ trend is consistent with a SA scenario with the X-ray spectral curvature related to the acceleration rather than to the cooling process. The presence of a power-law low-energy tail, found during the brightest X-ray flares in 2006, and the corresponding values of the electron distribution index ($s\\simeq[2.2-2.4]$) are consistent with the predictions of relativistic Fermi first-order acceleration models ($s\\simeq 2.3$). Our findings hint for a simultaneous role of the first and second order processes both related to the magnetic field turbulence. The stochastic acceleration hence is related to a momentum diffusion coefficient which drives the curvature consistently with the $E_p-b$ observed trend. The first order acceleration, observationally signed by the low-energy power-law spectral index, is linked to the spatial diffusion coefficient. Interestingly, recent particle-in-cell simulations by \\cite{Spit2008} obtain electron distributions that are compatible with this scenario.\\\\ The $S_p-E_p$ trend demonstrates the connection between the average energy of the particle distribution and the power output of the source. The observed values of the expected power-law dependence ($S_p \\propto E_p^{\\alpha}$) exclude $ B$, and $ \\delta $, indicating $ \\gamma_p $ as the main driver of the $S_p-E_p$ trend. The break in the $S_p-E_p$ scatter plot (see Fig. \\ref{fig:SpEp}) at about 1 keV, where the typical source luminosity is about $L_p \\simeq 10^{45} $ erg/s, can be interpreted as an indicator of the energetic content of the jet." }, "1003/1003.5223_arXiv.txt": { "abstract": "A Rotating Modulator (RM) is one of a class of techniques for indirect imaging of an object scene by modulation and detection of incident photons. Comparison of the RM to more common imaging techniques, the Rotating Modulation Collimator and the coded aperture, reveals trade-offs in instrument weight and complexity, sensitivity, angular resolution, and image fidelity. In the case of a high-energy (hundreds of keV to MeV), wide field-of-view, satellite or balloon-borne astrophysical survey mission, the RM is shown to be an attractive option when coupled with a reconstruction algorithm that can simultaneously achieve super-resolution and suppress fluctuations arising from statistical noise. We describe the Noise-Compensating Algebraic Reconstruction (NCAR) algorithm, which is shown to perform better than traditional deconvolution techniques for most object scene distributions. Results from Monte Carlo simulations demonstrate that NCAR achieves super-resolution, can resolve multiple point sources and complex distributions, and manifests noise as fuzzy sidelobes about the true source location, rather than spurious peaks elsewhere in the image as seen with other techniques. ", "introduction": "\\IEEEPARstart{I}{maging} hard x-ray and gamma-ray photons (hundreds of keV to MeV) cannot be accomplished using focusing techniques, as with lower energy photons. A concise overview of several techniques is given in \\cite{Caroli1987}. In one class of methods, incident photons are spatially or temporally modulated before detection. The recorded data are not a direct representation of the object scene, and so additional steps are required to deconvolve this information with a pre-determined instrument response function (i.e., system matrix). A variety of deconvolution techniques have been developed across a wide range of applications, each employed for its demonstration of computational speed, noise suppression, resolving power, or fidelity; classes include statistical, algebraic, and ``ad-hoc'' algorithms. Algebraic techniques attempt to solve directly for the unknown image. Since this class requires convergence of the reconstruction to the data, noise in the data can be amplified and cause spurious peaks to arise. As we will show, however, an algebraic technique has a distinct advantage (the ability to achieve ``super-resolution'') that makes it entirely suitable for the rotating modulator, and with an appropriate non-linear step, statistical noise may be adequately suppressed. ", "conclusions": "A rotating modulator is an instrument capable of imaging at hard x-ray and gamma-ray energies. The instrument is composed of a single grid of slats above a collection of circular non-imaging detector elements. While the instrument response for the RM is non-ideal (the sensitivity is approximately 62\\% of that of an equivalent coded aperture), the near-continuous nature of the observed count profiles enables an analysis which can go beyond the instrinsic resolution as defined by the geometry, whereas the coded aperture cannot. When compared to the RMC, the RM is able to achieve higher and more uniform sensitivity, does not suffer from mechanical collimation at high energies, and has lower overall weight. We have found that algebraic solutions are the only reconstruction techniques that can both achieve super-resolution and perform relatively fast. Algebraic methods, however, have in the past suffered from poor reconstructions due to the requirement of strict agreement with the noisy data. We have developed a novel technique, NCAR, based on an algebraic solution and non-linear physical constraints. A noise-compensation factor derived directly from the Poisson uncertainty in the data is added to a Gauss-Seidel iteration. By performing a running average on the successive results after a specific agreement criterion has been met, the spurious peaks arising from noise in the data are effectively suppressed, while the true sources remain. The uncertainty of the true source location due to noise in the data is manifest as a blurring of the scene rather than spurious peaks elsewhere in the image. When compared to other common deconvolution algorithms (MEM, CLEAN, and DDM), NCAR provides higher fidelity in most cases. Furthermore, the technique is general enough that it could be used in imaging and other applications where deconvolution is required in the presence of large background." }, "1003/1003.0892_arXiv.txt": { "abstract": "We fit every emission line in the high-resolution \\chandra\\/ grating spectrum of \\zpup\\/ with an empirical line profile model that accounts for the effects of Doppler broadening and attenuation by the bulk wind. For each of sixteen lines or line complexes that can be reliably measured, we determine a best-fitting fiducial optical depth, $\\taustar \\equiv \\kappa\\Mdot/4{\\pi}\\Rstar\\vinf$, and place confidence limits on this parameter. These sixteen lines include seven that have not previously been reported on in the literature. The extended wavelength range of these lines allows us to infer, for the first time, a clear increase in \\taustar\\/ with line wavelength, as expected from the wavelength increase of bound-free absorption opacity. The small overall values of \\taustar, reflected in the rather modest asymmetry in the line profiles, can moreover all be fit simultaneously by simply assuming a moderate mass-loss rate of $3.5 \\pm 0.3 \\times 10^{-6}$ \\Msunyr, without any need to invoke porosity effects in the wind. The quoted uncertainty is statistical, but the largest source of uncertainty in the derived mass-loss rate is due to the uncertainty in the elemental abundances of \\zpup, which affects the continuum opacity of the wind, and which we estimate to be a factor of two. Even so, the mass-loss rate we find is significantly below the most recent smooth-wind \\Ha\\/ mass-loss rate determinations for \\zpup, but is in line with newer determinations that account for small-scale wind clumping. If \\zpup\\/ is representative of other massive stars, these results will have important implications for stellar and galactic evolution. ", "introduction": "\\label{sec:intro} Massive stars can lose a significant fraction of their original mass during their short lifetimes due to their strong, radiation-driven stellar winds. Accurate determinations of these stars' mass-loss rates are therefore important from an evolutionary point of view, as well as for understanding the radiative driving process itself. Massive star winds are also an important source of energy, momentum, and (chemically enriched) matter deposition into the interstellar medium, making accurate mass-loss rate determinations important from a galactic perspective. A consensus appeared to be reached by the late 1990s that the mass-loss rates of O stars were accurately known observationally and theoretically, using the modified \\citep{Pauldrach1986} CAK \\citep{cak1975} theory of line-driven stellar winds. This understanding was thought to be good enough that UV observations of spectral signatures of their winds could be used to determine their luminosities with sufficient accuracy to make extragalactic O stars standard candles \\citep{Puls1996}. This consensus has unraveled in the last few years, mostly from the observational side, where a growing appreciation of wind clumping -- an effect whose importance has long been recognized \\citep{elm1998,hm1999,hk1999} -- has led to a re-evaluation of mass-loss rate diagnostics, including \\Ha\\/ emission, radio and IR free-free emission, and UV absorption \\citep{blh2005,Puls2006,fmp2006}. Accounting for small-scale clumping that affects density squared emission diagnostics -- and also ionization balance and thus ionic column density diagnostics like UV resonance lines -- leads to a downward revision of mass-loss rates by a factor of several, with a fair amount of controversy over the actual factor \\citep{hfo2008,pvn2008}. X-ray emission line profile analysis provides a good and independent way to measure the mass-loss rates of O stars. Like the UV absorption line diagnostics, X-ray emission profile diagnostics are sensitive to the wind column density and thus are not directly affected by clumping in the way density-squared diagnostics are. Unlike the UV absorption line diagnostics, however, X-ray profile analysis is not very sensitive to the ionization balance; moreover, as it relies on continuum opacity rather than line opacity, it is not subject to the uncertainty associated with saturated absorption lines that hamper the interpretation of the UV diagnostics. In this paper, we apply a quantitative line profile analysis to the \\chandra\\/ grating spectrum of the early O supergiant, \\zpup, one of the nearest O stars to the Earth and a star that has long been used as a canonical example of an early O star with a strong radiation-driven wind. Previous analysis of the same \\chandra\\/ data has established that the kinematics of the X-ray emitting plasma, as diagnosed by the line widths, are in good agreement with wind-shock theory, and that there are modest signatures of attenuation of the X-rays by the dominant cold wind component in which the shock-heated X-ray emitting plasma is embedded \\citep{kco2003}. The work presented here goes beyond the profile analysis reported in that paper in several respects. We analyze many lines left out of the original study that are weak, but which carry a significant amount of information. We better account for line blends and are careful to exclude those lines where blending cannot be adequately modelled. We model the continuum emission underlying each line separately from the line itself. We use a realistic model of the spectrometers' responses and the telescope and detector effective area. And we include the High Energy Grating (HEG) spectral data, where appropriate, to augment the higher signal-to-noise Medium Energy Grating (MEG) data that \\citet{kco2003} reported on. Implementing all of these improvements enables us to derive highly reliable values of the fiducial wind optical depth parameter, $\\taustar \\equiv \\kappa\\Mdot/4{\\pi}\\Rstar\\vinf$, for each of sixteen emission lines or line complexes in the \\chandra\\/ grating spectrum of \\zpup. Using a model of the wavelength-dependent wind opacity, $\\kappa$, and values for the star's radius, \\Rstar, and wind terminal velocity, \\vinf, derived from UV and optical observations, we can fit a value of the mass-loss rate, \\Mdot, to the ensemble of \\taustar\\/ values, and thereby determine the mass-loss rate of \\zpup\\/ based on the observed X-ray emission line profiles. In doing this, we also can verify that the wavelength-dependence of the optical depth values -- derived separately for each individual line -- is consistent with that of the atomic opacity of the bulk wind, rather than the gray opacity that would, for example, be obtained from an extremely porous wind \\citep{ofh2006,oc2006}. While a moderate porosity might reduce somewhat the effective absorption while still retaining some wavelength dependence, for simplicity our analysis here assumes a purely atomic opacity set by photoelectric absorption, with no reduction from porosity. This assumption is justified by the large porosity lengths required for any appreciable porosity effect on line profile shapes \\citep{oc2006} and the very small-scale clumping in state-of-the-art two-dimensional radiation hydrodynamics simulations \\citep{do2003}. Furthermore, preliminary results indicate that profile models that explicitly include porosity are not favored over ones that do not \\citep{clt2008}. We will extend this result in a forthcoming paper but do not address the effect of porosity on individual line profile shapes directly in the current work. The paper is organized as follows: We begin by describing the \\chandra\\/ data set and defining a sample of well behaved emission lines for our analysis in \\S 2. We briefly evaluate the stellar and wind properties of \\zpup\\/ in \\S 3. In \\S 4 we describe the empirical profile model for X-ray emission lines and report on the fits to the sixteen usable lines and line complexes in the spectrum. We discuss the implications of the profile model fitting results in \\S 5, and summarize our conclusions in \\S 6. \\begin{figure*} \\includegraphics[angle=0,width=160mm]{fig1.eps} \\caption{The entire usable portions of the MEG (top) and HEG (bottom) first order (negative and positive orders coadded) spectra of \\zpup. The binning is native (5 m\\AA\\/ for the MEG and 2.5 m\\AA\\/ for the HEG). Vertical dashed lines in the data panels themselves represent the laboratory rest wavelengths of all the 21 lines and line complexes we fit with the profile model. Solid (red) vertical lines between the two spectral plots indicate the lines we successfully fit with profile models and lines we attempted to fit but which were too blended to extract meaningful model parameters are indicated by dashed (green) lines. For all blended emission lines we show only one of these solid or dashed lines between the panels, and align it with the bluest line in the blend. } \\label{fig:bigplot_both} \\end{figure*} ", "conclusions": "By quantitatively analyzing all the X-ray line profiles in the \\chandra\\/ spectrum, we have determined a mass-loss rate of $3.5 \\times 10^{-6}$ \\Msunyr, with a confidence range of $2$ to $4 \\times 10^{-6}$ \\Msunyr. Within the context of the simple, spherically symmetric wind emission and absorption model we employ, the largest uncertainty arises from the abundances used in the atomic opacity model. This method of mass-loss rate determination from X-ray profiles is a potentially powerful tool for addressing the important issue of the actual mass-loss rates of O stars. Care must be taken in the profile analysis, however, as well as in the interpretation of the trends found in the derived \\taustar\\/ values. It is especially important to use a realistic model of the wind opacity. And for O stars with weaker winds, especially, it will be important to verify that the X-ray profiles are consistent with the overall paradigm of embedded wind shocks. Here, an independent determination of the terminal velocity of the X-ray emitting plasma by analyzing the widths and profiles of the observed X-ray lines themselves will be crucial. In the case of \\zpup, we have shown that the X-ray profiles are in fact consistent with the same wind kinematics seen in UV absorption line spectra of the bulk wind. And the profile analysis also strongly constrains the onset radius of X-ray production to be about $r = 1.5$ \\Rstar. An additional conclusion from the profile analysis is that there is no need to invoke large scale porosity to explain individual line profiles, as the overall wavelength trend is completely consistent (within the measurement errors) with the wavelength-dependence of the atomic opacity. The lower-than-expected wind optical depths are simply due to a reduction in the wind mass-loss rate. This modest reduction is consistent with other recent determinations that account for the effect of small-scale clumping on density-squared diagnostics and ionization corrections \\citep{Puls2006}." }, "1003/1003.3858_arXiv.txt": { "abstract": "{We introduce and carefully define an entire class of field theories based on non-standard spinors. Their dominant interaction is via the gravitational field which makes them naturally dark; we refer to them as \\emph{Dark Spinors}. We provide a critical analysis of previous proposals for dark spinors noting that they violate Lorentz invariance. As a working assumption we restrict our analysis to non-standard spinors which preserve Lorentz invariance, whilst being non-local and explicitly construct such a theory. We construct the complete energy-momentum tensor and derive its components explicitly by assuming a specific projection operator. It is natural to next consider dark spinors in a cosmological setting. We find various interesting solutions where the spinor field leads to slow roll and fast roll de Sitter solutions. We also analyse models where the spinor is coupled conformally to gravity, and consider the perturbations and stability of the spinor.} \\preprint{} ", "introduction": "\\label{sect:Introduction} In recent years, our understanding of the universe has become greatly improved thanks to the high precision cosmological observations that we have available today. According to the Standard Model of Cosmology, which assumes General Relativity as the theory describing the gravitational interaction, our universe is composed by about $4\\%$ of baryons, $23\\%$ of dark matter and $73\\%$ of dark energy. Moreover, in addition to these components, we need to assume an early inflationary epoch in order to explain the current state of our universe. Although this budget enables us to successfully account for the current cosmological data, it needs to assume the existence of three unknown components from a particle physics point of view, namely: dark matter, dark energy and inflaton field. Thus, we find that predictions based on General Relativity plus the Standard Model of particle physics are at odds with current astronomical observations, not only on cosmological scales, but also on galactic scales where dark matter plays a crucial role. This indicates failures either in particle physics or in general relativity (or both) and, in particular, it might be indicating the existence of new particles/fields as candidates to dark matter, dark energy and the inflaton which could arise in high energy physics \\cite{Li4,Li5,Bi1,Li1,Li2,Ci1,Ci2,Ci3,Li3,Bi2,Ci4,Ci5,Bi3,Watanabe:2009nc,Poplawski2010,Mota:2007sz,alpha}. Spinors have played an important role in mathematics and physics throughout the last 80 years. They theoretically model particles with half integer spin, like the electron in the massive case or the neutrino (massive or massless). The spin structure of manifolds has played an important part in modern mathematics, while in mathematical physics this structure motivated the twistor program. In the framework of particle physics all spinors used are either Dirac, Weyl (massless Dirac spinors) or Majorana spinors, $\\psi$. Such spinors obey a field equation which is first order in the derivatives (momenta) of $\\psi$. Cosmologically, this first order field equation implies that the average value of both $\\Phi = \\bar{\\psi}\\psi$ and the spinor energy density of a free spinor field evolves like the energy density of pressure-less dust i.e.~proportional to $(1+z)^{3}$, where $z$ is the redshift. Additionally, the first order nature of the field equation results in a quantum propagator, $G_{F}$, which, for large momenta $p$, behaves as $G_{F} \\propto p^{-1}$. This limits the form of perturbatively renormalizable spinor self-interaction terms in the action to be no more than quadratic in $\\psi$ e.g.~$\\bar{\\psi}\\psi$ and $\\bar{\\psi}\\gamma_{\\mu}A^{\\mu}\\psi$. The momentum drop-off of $G_{F}$ also results in $\\psi$ having a canonical mass dimension of $3/2$. A wider range of renormalizable self-interaction terms and cosmological behavior would be allowed if one could construct a viable spinor field theory where $G_F \\propto p^{-2}$, for large $p$, resulting in a $\\psi$ with a canonical mass dimension of unity. We refer to this entire class of spinor field theories with such properties as Non-Standard Spinors (NSS). This class of spinors is closely related to Wigner's non-standard classes~\\cite{Wigner:1939cj}. Weinberg showed that, under the assumptions of Lorentz invariance and locality, the only spin-$1/2$ quantum field theory is that which describes standard spinors (Dirac, Weyl, Majorana). NSS will therefore violate either locality or Lorentz invariance, or possibly both. Our working assumption is that reasonable NSS models preserve Lorentz invariance, while being non-local. Along these lines of reasoning, Ahluwalia-Khalilova and Grumiller~\\cite{jcap,prd} constructed a NSS model using momentum space eigen-spinors of the charge conjugation operator \\emph{Eigenspinoren des LadungsKonjugationsOperators} (ELKO) to build a quantum field. They showed that such spinors belong to a non-standard Wigner class, and to exhibit non-locality~\\cite{Wigner:1939cj}. They satisfy $(CPT)^2=-\\mathbb{I}$ while Dirac spinors satisfy $(CPT)^2=\\mathbb{I}$. In more mathematical terms, they belong to a wider class of spinorial fields, so-called flagpole spinor fields~\\cite{daRocha:2005ti}. The spinors correspond to the class 5, according to Lounesto's classification which is based on bilinear covariants, similar to Majarona spinors, see also~\\cite{daRocha:2008we,daRocha:2009gb,HoffdaSilva:2009is}. Locality issues and Lorentz invariance were further investigated in~\\cite{Ahluwalia:2008xi,Ahluwalia:2009rh} resulting in results along the lines of the current work. Causality has been analyzed in~\\cite{Fabbri:2009ka,Fabbri:2009aj}. The construction of ELKOs using momentum space eigenspinors, $\\lambda(\\mathbf{p},h,e)$, of the charge conjugation operator leads to a spinor field with a double helicity structure. The left-handed and the right-handed spinor have opposite helicities which in turn requires a careful construction of the resulting field theory. These spinors have received quite some attention recently~\\cite{Boehmer:2006qq,Boehmer:2007dh,n6} and their effects in cosmology have been investigated~\\cite{Boehmer:2007ut,Boehmer:2008ah,Boehmer:2008rz,Gredat:2008qf,Boehmer:2009aw,n1,n3,n4,n5,Shankaranarayanan:2009sz,Shankaranarayanan:2010st,Boehmer:2010tv,Wei:2010ad}. However, as we will show in \\S \\ref{sec:sNSS}, ELKOs spinors, defined in the way described above, are not Lorentz invariant. We demonstrate using our construction of NSS where this Lorentz violation appears, thus confirming~\\cite{Ahluwalia:2008xi,Ahluwalia:2009rh}. The original analyses defined the field structure entirely in terms of momentum space basis spinors rather than say starting with an action whose minimization would imply that structure. This led to the violation of Lorentz invariance being hidden in the mathematical structure of the model. In the present work, on the other hand, we start with a general action principle for NSS. When applied to the ELKOs and an alternative model also based eigenspinors of the charge conjugation operator, the violation of Lorentz invariance and other issues with their construction are explicit at the level of the action. The original ELKO definition is seen to require a preferred space-like direction and is ill-defined when the momentum points along that direction. We offer a new NSS field theory which is also based on the eigenspinors of the charge conjugation operator (i.e.~using the basis $\\lambda(\\mathbf{p},h,e)$) which respects the rotational group $SO(3)$ but is not invariant under boosts. We shall see that the general construction of NSS models can be seen as the choice of some operator $\\opp{P}$ satisfying $\\opp{P}^2 = \\mathbb{I}$ which acts on $\\psi$ to project out what states that would otherwise give an inconsistent Hamiltonian density. In this article we provide a general treatment of class of NSS models based on an action principle and choice of operator $\\opp{P}$. We show that there is one, potentially unique, choice of $P$ which results in a Lorentz invariant, ghost-free but non-local spinor field theory with canonical mass dimension one. We are also interested in the cosmological behavior of general NSS models and construct the energy momentum tensor, $T_{\\mu\\nu}$. For ELKO spinors it appears that, at present, no one has obtained the full $T_{\\mu\\nu}$ as all previous works in the literature, including ours, have overlooked contributions to $T_{\\mu\\nu}$ from the variation of spin connection. This article is organized as follows: we define our notation, general spinors and exactly what a non-standard spinor is in \\S \\ref{sec:spinors}, then in \\S\\ref{sec:sNSS} we look specifically at the original ELKO definition, offer a modified version, finishing the section by examining the possibility of a Lorentz invariant non-standard spinor. In \\S \\ref{sec:emt} we examine the energy momentum tensor both with and without the projection operator, this then leads us nicely into sections \\S \\ref{sec:ELKOCos} and \\S \\ref{sec:NSSCos}, where we examine the cosmological applications of both the original ELKO and the Lorentz invariant NSS respectively and in each case note the existence of non-trivial de Sitter solutions. We make our final remarks in \\S \\ref{sec:conc}, followed by three appendices showing explicit calculations of the variation of the spin connection with respect to the metric for the general case, the Dirac spinor and finally the ELKO spinor. ", "conclusions": "\\label{sec:conc} In this article, we have constructed a new class of theories of non-standard spinors (NSS). Their dynamics is more general that than that of Dirac or Majorana spinors, even when self-interactions are taken into account. In contrast to standard spinors, the dynamics of NSS is \\emph{not} described by a first order equations of motion like the Dirac equation. This leads to a more general and thus more interesting cosmological behavior than that exhibited by normal spinors, including for instance the existence of non-trivial de Sitter solutions. It is therefore possible to invoke NSS, as an alternative to scalar fields, as one possible explanation of the early and late time acceleration of our Universe. As example of a NSS theory is that of the eigenspinors of $C$, originally proposed by Ahluwalia-Khalilova and Grumiller in Ref.~\\cite{jcap,prd}. We have constructed a general action for NSS. We began by considering a Klein-Gordon action for spinors and noted that because there is no positive definite Lorentz invariant norm for spinors, such an action would have propagating negative energy ghost modes. These would lead to instabilities in the quantum theory. These negative energy modes can, however, be eliminated by including an additional term in the action. This term depends on an operator $\\opp{P}$, which must have the property that $\\opp{P} \\psi = \\psi$ on positive energy modes, and $\\opp{P}^2 \\psi = -\\psi$ on negative energy ones. Hence $\\opp{P}^2 = \\mathbb{I}$. We also noted that in momentum space $P$ must be an odd function of momentum i.e. $\\opp{P}(\\mathbf{p}) = -\\opp{P}(-\\mathbf{p})$. The original ELKO model as well as Dirac and Majorana spinors correspond simply to specific choices of the projection operator $\\opp{P}$. By constructing the NSS action in this way, we found that ELKO spinors require a choice of $\\opp{P}$ that is \\emph{not} Lorentz invariant but instead includes a preferred axis. Previous works on this field have effectively made a specific choice of frame so that this preferred direction and hence the violation of Lorentz invariance was not manifest explicitly. Using our general definition of NSS theories, however, the violation of Lorentz invariance which is required to define ELKOs is clear at the level of the action. We note that an alternative definition of the eigenspinors of $\\opp{C}$ replaces the spatial direction with a preferred time-like direction. A truly Lorentz invariant NSS theory requires a Lorentz invariant projection operator $\\opp{P}$. For most such choices of the operator, $\\opp{P} \\psi = \\psi$ is essentially equivalent to the Dirac equation with self-interaction terms. The projection condition then effectively reduces the dynamics from second to first order equations of motion. We found that there was only one suitable, Lorentz invariant choice of $P$ which preserves the second order dynamics. In momentum space and by assuming a flat background this operator is given by $\\opp{P} = p_{\\mu}\\gamma^{\\mu}/\\sqrt{p_{\\mu}p^{\\mu}}$, and so $\\opp{P}$ is a non-local operator. In the absence of self-interactions, i.e.~we have $V = V_0 + m^2 \\bar{\\psi}\\psi$, the field equations then reduce to the Dirac equation, but for more complicated choices of $V$ this is no longer the case. Having provided a general definition of NSS we then constructed and examined the full energy-momentum tensor. In the case of ELKOs we noted that even if one ignores the additional contributions to the energy-momentum tensor from the variation of $P$ which respect to the metric, the energy momentum tensor differs from that which has previously appeared in the literature~\\cite{Boehmer:2006qq,Boehmer:2007dh,n6,Boehmer:2007ut,Boehmer:2008ah,Boehmer:2008rz,Gredat:2008qf,Boehmer:2009aw,n1,n3,n4,n5,Shankaranarayanan:2009sz,Shankaranarayanan:2010st,Boehmer:2010tv,Wei:2010ad}. We show explicitly in Appendix~\\ref{appelko} where the additional terms come from, namely from the variation of the spin connection with respect to the metric. In case of Dirac spinors this contribution identically vanishes, see Appendix~\\ref{appdirac}, and therefore, we believe, this has been neglected in the past for ELKO spinors. The presence of additional terms even in the ELKO energy-momentum tensor led us to re-address the cosmology of such models. We defined $\\psi = \\varphi(t)\\xi$, where $\\xi$ is a constant spinor, and so were able to treat $\\varphi$ as the only dynamical variable cosmological. In the simplest case, the flat FLRW background, it produces an effective gravitational coupling $G$ which places a simple limit on the maximum value of $\\Phi$, namely $\\Phi < 1 / \\sqrt{\\pi G} = 2 \\sqrt{2}M_{Pl}$. When we examine de Sitter type solutions, this ansatz produces a potential very similar to those discussed in previous work. Its form is surprisingly similar to the case where the variation of the spin connection with respect to the metric was neglected. However, as we noted, it is not clear if a Lorentz invariant NSS model can be found which has these dynamics. We also considered the dynamics of the only Lorentz invariant NSS model we were able to construct. We found that cosmologically the dynamics can be written in terms of $\\Phi = \\bar{\\psi}\\psi$ and $\\Psi = \\dot{\\bar{\\psi}}\\dot{\\psi}$. The potential describing self-interactions depends only on $\\Phi$: $V=V(\\Phi)$. When $V_{,\\Phi\\Phi} =0$, we had previously noted that the theory should be equivalent to that of a Dirac spinor, and so $\\rho_{\\psi}\\propto a^{-3}$. We confirmed that this was indeed the case. When $V_{,\\Phi \\Phi} \\neq 0$ and $V_{,\\Phi} < 0$ we found that there stable de Sitter solutions; stability of these solutions is ensured when $V_{,\\Phi \\Phi} > 0$. In contrast to the situation with scalar fields, de Sitter solutions do \\emph{not} required $\\vert V_{,\\Phi}/\\kappa V(\\Phi)\\vert \\ll 1$ (i.e. slow-roll). Instead with NSS we generally have $-V_{,\\Phi}/\\kappa V(\\Phi) \\sim O(1)$. With NSS spinors, the expansion of the Universe acts as a brake to prevent the effective scalar field $\\Phi$ rolling down the potential. The main results of the paper lie in our discussion of the definition and dynamics of the entire class of NSS, and their cosmology. We laid the foundations of an in depth analysis of the dynamics of this field in an arbitrary spacetime with a focus on cosmological dynamics. Importantly we also constructed what is, to the best of our knowledge, the only Lorentz invariant, ghost free proposal for a theory of non-standard spinors. The cosmological dynamics of the effective scalar degree of freedom in both ELKO and Lorentz invariant NSS cosmology show a large number of very interesting properties, mainly due to their more complicated couplings to the gravitational sector when compared to the scalar field. The cosmological evolution of the NSS energy density exhibits a much wider range of behavior than that seen with Dirac spinors where it always scales as $a^{-3}$. The existence of stable de Sitter solutions means that NSS could represent an alternative to scalar field inflation / dark energy." }, "1003/1003.3405_arXiv.txt": { "abstract": "The European Pulsar Timing Array (EPTA) is a multi-institutional, multi-telescope collaboration, with the goal of using high-precision pulsar timing to directly detect gravitational waves. In this article we discuss the EPTA member telescopes, current achieved timing precision, and near-future goals. We report a preliminary upper limit to the amplitude of a gravitational wave background. We also discuss the Large European Array for Pulsars, in which the five major European telescopes involved in pulsar timing will be combined to provide a coherent array that will give similar sensitivity to the Arecibo radio telescope, and larger sky coverage. ", "introduction": "\\label{sec:intro} Pulsars are highly magnetized neutron stars (NS), the remnants of the supernovae of massive stellar progenitors \\citep{gol68,pac68}. The majority of known pulsars rotate approximately once every second, and most are visible via their radio emission. It is generally accepted that radio pulsars radiate from above their magnetic poles; with each rotation of the NS, we observe a cross-section through the emission region that cuts our line of sight at Earth \\citep{gol68,pac68,gj69,stu71}. This gives rise to the ``pulses'' that lend pulsars their name. The intrinsic periodicity of most pulsars is remarkably stable, with typical period derivatives of $\\sim 10^{-15}$ s\\,s$^{-1}$ \\citep[see e.g.,][]{lk05}. There is a population of so-called ``recycled'' pulsars with rotational periods of $\\sim 1-100$ ms. They are the result of a mass-accretion episode from a stellar binary companion onto the NS. The accreted matter carries angular momentum, resulting in the ``spin-up'' of the NS \\citep{sb76,acrs82,rs82,sv82}. These millisecond pulsars (MSPs) are observed to be more stable as a population than the mainstream pulsars, with rotational period derivatives of $\\sim 10^{-20}$ s\\,s$^{-1}$ \\citep[see e.g.,][]{lk05}. They also show fewer instances of intrinsic, often low-frequency noise, corresponding to instabilities which affect the timing of mainstream pulsars. This is frequently referred to as ``timing'' or ``red'' noise, and can be detrimental to long-term timing precision \\citep[e.g.,][]{antt94, vbc+09, hlml09}. For these reasons, in addition to their $2-3$ orders of magnitude shorter rotation rate, MSPs are thus the pulsars of choice when the highest timing precision is required, as is the case for the pursuit of gravitational wave (GW) detection. Through measurement of the pulse arrival times at Earth, we can continually improve our model ephemeris for the rotation of a pulsar. This is done by properly accounting for any deviations to the expected arrival times due to, for example, the motion of the Earth, the proper motion of the pulsar, or the binary parameters of the pulsar. This pulsar ``timing'' aims to successfully account for every rotation of the NS, which is referred to as ``phase connection''; this is what provides the high precision for which radio pulsar astronomy is known. The differences between the observed and expected times of arrival (TOAs) are referred to as ``timing residuals''. An ideal timing model will generally give residuals that are normally distributed about a mean of zero \\citep[see e.g.,][]{lk05}. Pulsar timing has been very successful for the measurement pulsar properties, and the study of the environments around NSs, for example. Perhaps its most celebrated results, however, are those that have tested general relativity (GR) and other theories of gravity, which we briefly outline in the next section. The potential that pulsar timing also holds for the direct detection of GWs is discussed in section \\ref{sec:gw_detection}. \\subsection{Testing gravitational theories with pulsars} The very first test of GR using pulsar observations involved timing of the first-discovered binary pulsar, PSR~B1913+16. Through high-precision timing of this double-neutron-star system, the measured decay of its orbit was found to agree spectacularly well with the orbital energy loss due to quadrupolar GW radiation from that system, as predicted by GR \\citep{ht75a,tw89}. More recently, observations of the double-pulsar system PSR~J0737$-$3039A/B, in which both NSs are observed as radio pulsars, have provided the most stringent test to date of GR in the strong-field regime \\citep{ksm+06}. In addition, several pulsars have been, and continue to be, used to test various aspects of GR and other theories of gravity. These include, for example, probing Strong Equivalence Principle violation by setting limits on phenomena predicted to exist by alternative theories of gravity. Among these are gravitational dipole radiation, secular evolution of binary parameters due to preferred frame effects, and a variation in Newton's gravitational constant \\citep[e.g.,][]{nor90, dt91, wex00, arz03, vbs+08, lwj+09}. \\subsection{Direct detection of gravitational waves} \\label{sec:gw_detection} Although some of the tests described in the preceding subsection compare predictions of GR and other theories of gravity to pulsar observations in order to \\emph{infer} the existence of GW radiation as prescribed by GR, they fall short of being able to claim a \\emph{direct} detection of GWs. Pulsars can, however, be used to directly detect a stochastic GW background (GWB), presumably produced by supermassive black hole binaries at the centers of distant galaxies \\citep[][]{rr95a, jb03} or perhaps by exotic phenomena such as cosmic strings \\citep[e.g.,][]{dv05}. This is done by using an ensemble of pulsars as arms of a very large GW detector. The perturbing effect of a GWB passing by the Earth would be correlated amongst observations of all pulsars. This would manifest itself as a distinct signature in the timing residuals, once all model parameters particular to the pulsar system were properly accounted for. The cross-correlation of these residuals is predicted to show a distinct dependence on the separations between the various pulsars observed, and so a distribution of pulsars evenly spread over as much of the sky as possible is desired for such a study \\citep[e.g.,][]{det79, hd83}. A distribution of sources in a ``pulsar timing array'' (PTA) thus provides a unique laboratory for performing this experiment. The PTA experiment is sensitive to GWs with frequencies of $\\sim 10^{-9}$ Hz, since the limiting frequency corresponds to the timespan of the data sets used, typically $1-10$ years in length. Such a detection will be complementary to those that will be made by terrestrial and space-based interferometers such as LIGO \\citep{aad+92} and LISA \\citep{dan00}, which are most sensitive to wave frequencies of $\\sim 10-10^3$ Hz and $\\sim 10^{-4}-10^{-2}$ Hz, respectively. The expected correlated signal is very subtle, and its detection requires the highest-precision timing, with rms of timing residuals well below $1\\ \\mu$s. It is also crucial that the pulsars used for this undertaking are stable over long timescales \\citep[e.g.,][]{vbc+09}. As stated in section \\ref{sec:intro}, most MSPs fit this description. Ideally, observations of $20-40$ MSPs are required to be observed over $5-10$ years, with consistent $\\sim 100$ ns rms for the resulting timing residuals, in order to be able to significantly distinguish a GWB signal in the measured cross-correlation spectrum \\citep[][]{jhlm05}. The direct detection of GWs using a PTA is certainly a very challenging task. There now exist three major research groups dedicated to this effort: The Parkes Pulsar Timing Array (PPTA) in Australia, which uses the Parkes radio telescope \\citep{vbb+10}; the North American Nanohertz Observatory for Gravitational Waves (NANOGrav) collaboration in North America, which takes its data from the Arecibo and Green Bank telescopes \\citep{jfl+09}; and the European Pulsar Timing Array (EPTA). These groups have begun to develop a partnership, with the goal of sharing resources in an International Pulsar Timing Array (IPTA) effort, which will be necessary to achieve the objective of GW detection \\citep{haa+10}. In the following section, we focus on the instruments used by the EPTA and its current progress; in section \\ref{sec:leap}, we describe the Large European Array for Pulsars (LEAP), the next step for the EPTA toward GW detection; finally, we conclude in section \\ref{sec:summary} with a short summary. ", "conclusions": "\\label{sec:summary} We have outlined in this article some of the challenges that are faced by the pulsar astronomy community in detecting a stochastic GWB, and have described the role of the EPTA within this international effort. The EPTA, with its access to five large radio telescopes, already has the ability to provide high-quality pulsar timing data that is comparable to the other existing PTA groups. It is also in the process of moving a step further with the construction of the LEAP telescope array, which will coherently combine the time-series data from the five major European observatories. This will form a telescope that will be equivalent in sensitivity to the Arecibo radio telescope, but with much larger sky coverage. This project will not only provide very high-precision pulsar data, but will also be a vital testing ground for future multi-telescope arrays, most important of which is the SKA." }, "1003/1003.1400_arXiv.txt": { "abstract": "Motivated by a considerable scatter in the observationally inferred lifetimes of the embedded phase of star formation, we study the duration of the Class~0 and Class~I phases in upper-mass brown dwarfs and low-mass stars using numerical hydrodynamics simulations of the gravitational collapse of a large sample of cloud cores. We resolve the formation of a star/disk/envelope system and extend our numerical simulations to the late accretion phase when the envelope is nearly totally depleted of matter. We adopted a classification scheme of Andr\\'e et al. and calculate the lifetimes of the Class~0 and Class~I phases ($\\tau_{\\rm C0}$ and $\\tau_{\\rm CI}$, respectively) based on the mass remaining in the envelope. When cloud cores with various rotation rates, masses, and sizes (but identical otherwise) are considered, our modeling reveals a sub-linear correlation between the Class~0 lifetimes and stellar masses in the Class~0 phase with the least-squares fit exponent $m=0.8\\pm0.05$. The corresponding correlation between the Class~I lifetimes and stellar masses in the Class~I is super-linear with $m=1.2\\pm0.05$. If a wider sample of cloud cores is considered, which includes possible variations in the initial gas temperature, cloud core truncation radii, density enhancement amplitudes, initial gas density and angular velocity profiles, and magnetic fields, then the corresponding exponents may decrease by as much as 0.3. The duration of the Class~I phase is found to be longer than that of the Class~0 phase in most models, with a mean ratio $\\tau_{\\rm CI}/\\tau_{\\rm C0}\\approx$~1.5--2. A notable exception are YSOs that form from cloud cores with large initial density enhancements, in which case $\\tau_{\\rm C0}$ may be greater than $\\tau_{\\rm CI}$. Moreover, the upper-mass ($\\ga 1.0~M_\\odot$) cloud cores with frozen-in magnetic fields and high cloud core rotation rates may have the $\\tau_{\\rm CI}/\\tau_{\\rm C0}$ ratios as large as 3.0--4.0. We calculate the rate of mass accretion from the envelope onto the star/disk system and provide an approximation formula that can be used in semi-analytic models of cloud core collapse. ", "introduction": "Constraining the lifetimes associated with different phases of low-mass star formation has been traditionally one of the goals of research. In spite of much effort in this field, yet there is a considerable scatter among observationally estimated lifetimes of the embedded phase. Most previous observational studies suggested the duration of the Class~0 and Class~I phases to be within the 0.1--0.7~Myr range \\citep{Wilking89,Greene94,Kenyon95,Visser02,Hatchell07}, with similar relative lifetimes \\citep{Visser02,Hatchell07}, while some argue in favour of a significantly shorter lifetime for the Class 0 phase, $\\tau_{\\rm C0}$=0.01--0.06~Myr \\citep{Andre94}. In the most recent work involving a large set of YSOs from different star forming clouds, \\citet{Enoch09} and \\citet{Evans09} advocated for the lifetimes of 0.1--0.2~Myr for the Class~0 sources and 0.44--0.54~Myr for the Class~I ones. The aforementioned lifetimes may vary in part due to different estimation techniques and employed wavelengths and in part due to observations being confined to different molecular clouds. The cornerstone assumption of steady star formation, adopted in most observational studies, may break down if star formation relies heavily on external triggering, as may be the case in $\\rho$~Ophiuchi \\citep{Visser02}. Furthermore, the environment and local initial conditions determine how long the Class~0 and Class~I phases last. It is therefore interesting to compare the observationally inferred lifetimes with those predicted from numerical modeling of the cloud core collapse. For instance, \\citet{Masunaga00} performed radiation hydrodynamics simulations of the gravitational collapse of spherically symmetric cloud cores and found, based on the synthetic spectral energy distribution, that the Class~0 phase lasts around $2\\times10^4$~yr. However, their modeling was limited to only two cloud cores. Some work in this direction was also done by \\citet{Froebrich06}, who employed smooth particle hydrodynamics simulations of the fragmentation and collapse of turbulent, self-gravitating clouds and found Class 0 lifetimes to be of the order of $(2-6) \\times 10^{4}$~yr, significantly lower than those inferred from most observations. However, they did not provide estimates for the Class I lifetime, probably due to an enormous computational cost, and their lifetimes were inferred from the mass infall rates at $\\sim$~300~AU rather than from more customary diagnostics such as envelope masses or spectral properties. In this paper, we perform a comprehensive numerical study of the duration of the embedded phase of star formation in upper-mass brown dwarfs and low-mass stars. Using numerical hydrodynamics simulations, we compute the gravitational collapse of a large sample of gravitationally unstable cloud cores with various initial masses, rotation rates, gas temperatures, density enhancement amplitudes, truncation radii, and strengths of frozen-in magnetic fields. The paper is organized as follows. In Sections~\\ref{model} and \\ref{init}, we briefly review the numerical model and initial conditions. In Section~\\ref{schemes}, we discuss the adopted classification scheme of YSOs. In Section~\\ref{lifetimes}, we search for possible correlations between the Class~0 and Class~I lifetimes, from one hand, and stellar and cloud core masses, from the other hand. The effect of varying initial conditions is considered in Section~\\ref{initcond}. The model envelope depletion rates are tested against several empirical functions in Section~\\ref{depletion} and the main results are summarized in Section~\\ref{summary}. ", "conclusions": "\\label{summary} Using numerical hydrodynamics simulations, we have computed the gravitational collapse of a large set of cloud cores. We start from a gravitationally unstable, pre-stellar phase and terminate the simulations in the Class~II phase when most of the cloud core has been accreted by the forming star/disk system. We have considered cloud cores with various initial masses ($M_{\\rm cl}$=0.06--3.9~$M_\\odot$), ratios of the rotational to the gravitational energy ($\\beta$=(0.2--2.2)$\\times10^{-2}$), initial gas temperatures ($T_0$=10--18~K), truncation radii ($r_{\\rm out}/r_0$=6--12), strengths of frozen-in magnetic fields ($\\alpha_{\\rm m}$=0--0.3), initial gas density enhancements ($A$=2--8), and initial radial profiles of the gas surface density and angular velocity ($\\Sigma,\\Omega \\propto r^{-1}$ and $\\Sigma, \\Omega =const$). We employ a sophisticated method for distinguishing between the infalling envelope and the forming disk in our numerical simulations and determine the duration of the embedded phase of star formation, adopting a classification scheme based on the remaining mass in the envelope \\citep{Andre93}. We also calculate the envelope depletion rate $\\dot{M}_{\\rm env}$ (or, equivalently, the rate of mass accretion onto the star/disk system) and test the utility of two empirical functions for $\\dot{M}_{\\rm env}$ provided by \\citet{Bontemps96} (BATC function) and \\citet{Rice09} (RMA function). We find the following. \\begin{enumerate} \\item Class~0 ($\\tau_{\\rm C0}$) and Class~I ($\\tau_{\\rm CI}$) lifetimes correlate with the corresponding stellar mass $M_{\\rm \\ast,C0}$ and $M_{\\rm \\ast,CI}$, however the scaling is more complex than can be expected from simple theoretical grounds based on the linear correlation between the free-fall time and the cloud core mass. In particular, the form of the correlation depends on the cloud core properties such as the initial ratio of rotational to gravitational energy, initial gas temperature, density enhancement amplitude, and the initial radial distribution of gas surface density and angular velocity. In addition, frozen-in magnetic fields with constant flux-to-mass ratio can substantially influence the Class~I lifetimes. The correlation is further complicated due to the fact that the stellar mass varies considerably during the Class~0 phase. This makes the $\\tau_{\\rm C0}$--$M_{\\rm \\ast,C0}$ correlation dependent on the stellar age as well. \\item When cloud cores with varying rotation rates, masses and sizes (but otherwise identical) are considered, Class~0 and Class~I lifetimes have sub- and super-linear correlations with the corresponding stellar masses, respectively. These correlations extend from sub-stellar masses ($\\sim 0.01~M_\\odot$) to at least solar masses ($\\sim 1.5~M_\\odot$) and obey the following scaling laws $\\tau_{\\rm C0}=0.18~M_{\\rm \\ast,C0}^{0.8\\pm0.05}$ and $\\tau_{\\rm CI}=0.3~M_{\\rm \\ast,CI}^{1.2 \\pm 0.05}$ for Class~0/I stars of {\\it all possible ages}. This makes the observational determination of the mean lifetimes sensitive to the form of the initial mass function and/or to instrumental biases toward a particular mass band. For instance, our modeling predicts a mean Class~0 lifetime of $\\langle \\tau_{\\rm C0}\\rangle =0.044 $~Myr for stars in the 0.008--0.85~$M_\\odot$ mass range and $\\langle \\tau_{\\rm C0,tr}\\rangle =0.086 $~Myr for stars in the 0.02--0.85~$M_\\odot$ range. In the case of the Class I objects, we obtain $\\langle \\tau_{\\rm CI}\\rangle=0.09$~Myr and $\\langle \\tau_{\\rm CI}\\rangle_{\\rm tr}=0.15$~Myr for stars in the 0.03--1.24~$M_\\odot$ and 0.2--1.24~$M_\\odot$ mass ranges, respectively. It is evident that the neglect of objects at the lower mass end results in almost a factor of 2 overestimate of the mean lifetimes. \\item An increase in the initial cloud core temperature and density enhancement amplitude tend to lower the Class~0 and Class~I lifetimes, whereas cloud cores with initially constant gas surface density and angular velocity distributions (as compared to those with $\\Sigma,\\Omega \\propto r^{-1}$) have longer Class~0 and especially Class~I lifetimes. In addition, frozen-in magnetic fields may increase the Class~I lifetimes, particularly for the upper-mass cloud cores, thus steepening the corresponding $\\tau_{\\rm CI}$--$M_{\\rm \\ast,CI}$ relation. The net effect of varying initial condition is to {\\it weaken} the aforementioned sub-linear (Class~0) and super-linear (Class~I) correlations between the lifetimes and stellar masses by decreasing the corresponding exponents by as much as 0.3. However, more accurate modeling with ambipolar diffusion and magnetic braking taken into account is needed to accurately assess the magnitude of this effect. The outer truncation radius of a cloud core has little effect on the resulting lifetimes. \\item Most cloud cores give birth to YSOs whose Class~I lifetimes are longer than those of the Class~0 phase by roughly a factor of 1.5--2.0. A notable exception are YSOs formed from cloud cores a with large initial density enhancement. In the latter case, the duration of the Class~I phase may actually be shorter than that of the Class~0. In addition, the upper-mass models $M_{\\rm cl}>1~M_\\odot$ with frozen-in magnetic fields and high cloud core rotation rates $\\beta>10^{-2}$ may have the $\\tau_{\\rm CI}/\\tau_{\\rm C0}$ ratios as large as 3.0--4.0. \\item The time evolution of $\\dot{M}_{\\rm env}$ reveals two distinct modes: a shorter period of near-constant depletion rate and a longer period of gradual (and terminal) decline of $\\dot{M}_{\\rm env}$. The boundary between these two modes lies near the end of the Class~0 phase and the beginning of Class I phase. In the later mode, $\\dot{M}_{\\rm env}$ may show short-term fluctuations due to episodic disk expansions and contractions. The BATC function may provide an adequate fit to our model envelope depletion rates if the characteristic time of decline of $\\dot{M}_{\\rm env}$ is chosen properly. The RMA function fails to provide an acceptable fit to our model data, irrespective of the free parameters. \\end{enumerate} We emphasize that our lifetimes have been derived based on the AWTB classification scheme \\citep{Andre93} and may change by a factor of unity if other schemes are used. For instance, if we re-define the boundary between the Class I and Class~II phases in the AWTB scheme and assume that the Class~II phase begins when the envelope mass drops below 5\\% of the initial cloud core mass (in contrast to 10\\% adopted in our paper), the resulting Class~I lifetimes in model set~2 increase by a factor of 1.5. However, the derived trends and correlations between the lifetime of the embedded phase, from one hand, and the stellar masses, from the other hand, are expect to stay valid irrespective of the classification scheme used." }, "1003/1003.5918_arXiv.txt": { "abstract": "Cosmological probes are steadily reducing the total neutrino mass window, resulting in constraints on the neutrino-mass degeneracy as the most significant outcome. In this work we explore the discovery potential of cosmological probes to constrain the neutrino hierarchy, and point out some subtleties that could yield spurious claims of detection. This has an important implication for next generation of double beta decay experiments, that will be able to achieve a positive signal in the case of degenerate or inverted hierarchy of Majorana neutrinos. We find that cosmological experiments that nearly cover the whole sky could in principle distinguish the neutrino hierarchy by yielding `substantial' evidence for one scenario over the another, via precise measurements of the shape of the matter power spectrum from large scale structure and weak gravitational lensing. ", "introduction": "In the past decade, there has been great progress in neutrino physics. It has been shown that atmospheric neutrinos exhibit a large up-down asymmetry in the SuperKamiokande (SK) experiment. This was the first significant evidence for a finite neutrino mass \\cite{SuperK} and hence the incompleteness of the Standard Model of particle physics. Accelerator experiments \\cite{K2K, MINOS} have confirmed this evidence and improved the determination of the neutrino mass splitting required to explain the observations. The Sudbury Neutrino Observatory (SNO) experiment has shown that the solar neutrinos change their flavors from the electron type to other active types (muon and tau neutrinos)\\cite{SNO}. Finally, the KamLAND reactor anti-neutrino oscillation experiments reported a significant deficit in reactor anti-neutrino flux over approximately 180~km of propagation \\cite{KamLAND}. Combining results from the pioneering Homestake experiment \\cite{Homestake} and Gallium-based experiments \\cite{Gallium}, the decades-long solar neutrino problem \\cite{solarproblem} has been solved by the electron neutrinos produced at Sun's core propagating adiabatically to a heavier mass eigenstate due to the matter effect \\cite{MSW}. As a summary, two hierarchical neutrino mass splittings and two large mixing angles have been measured, while only a bound on a third mixing angle has been established. Furthermore the standard model has three neutrinos and the motivation for considering deviations from the standard model in the form of extra neutrino species has now disappeared \\cite{mena,miniboone}. New neutrino experiments aim to determine the remaining parameters of the neutrino mass matrix and the nature of the neutrino mass. Meanwhile, relic neutrinos produced in the early universe are hardly detectable by weak interactions but new cosmological probes offer the opportunity to detect relic neutrinos and determine neutrino mass parameters. It is very relevant that the maximal mixing of the solar mixing angle is excluded at a high significance. The exclusion of the maximal mixing by SNO \\cite{SNO} has an important impact on a deep question in neutrino physics: ``are neutrinos their own anti-particle?\". If the answer is yes, then neutrinos are Majorana fermions; if not, they are Dirac. If neutrinos and anti-neutrinos are identical, there could have been a process in the Early Universe that affected the balance between particles and anti-particles, leading to the matter anti-matter asymmetry we need to exist \\cite{leptogenesis}. This question can, in principle, be resolved if neutrinoless double beta decay is observed. Because such a phenomenon will violate the lepton number by two units, it cannot be caused if the neutrino is different from the anti-neutrino (see \\cite{murayama} and references therein). Many experimental proposals exist that will increase the sensitivity to such a phenomenon dramatically over the next ten years (e.g., \\cite{0nbb} and references therein). The crucial question we want to address is if a negative result from such experiments can lead to a definitive statement about the nature of neutrinos. Within three generations of neutrinos, and given all neutrino oscillation data, there are three possible mass spectra: a) degenerate, with mass splitting smaller than the neutrino masses, and two non-degenerate cases, b) normal hierarchy, with the larger mass splitting between the two more massive neutrinos and c) inverted hierarchy, with the smaller spitting between the two higher mass neutrinos. For the inverted hierarchy, a lower bound can be derived on the effective neutrino mass \\cite{murayama}. The bound for the degenerate spectrum is stronger than for inverted hierarchy. Unfortunately, for the normal hierarchy, one cannot obtain a similar rigorous lower limit. Neutrino oscillation data have measured the neutrino squared mass differences, which are hierarchical. Given the smallness of neutrino masses and the hierarchy in mass splittings, we can characterize the impact of neutrino masses on cosmological observables and in particular on the \\begin{center} \\begin{figure*}[!t] \\includegraphics[width=6.8cm]{standardplot.eps} \\includegraphics[width=6.8cm]{delta_definition.eps} \\caption{Left: constraints from neutrino oscillations and from cosmology in the $m$-$\\Sigma$ plane. Right: constraints from neutrino oscillations (shaded regions) and from cosmology in the $\\Sigma$-$\\Delta$ plane. In this parameterization the sign of $\\Delta$ specifies the hierarchy.} \\label{fig:0} \\end{figure*} \\end{center} the matter power spectrum by two parameters: the total mass $\\Sigma$ and the ratio of the largest mass splitting to the total mass, $\\Delta$. As we will show, one can safely neglect the impact of the solar mass splitting in cosmology. In this approach, two masses characterize the neutrino mass spectrum, the lightest one, $m$, and the heaviest one, $M$. Neutrino oscillation data are unable to resolve whether the mass spectrum consists in two light states with mass $m$ and a heavy one with mass $M$, named normal hierarchy (NH) or two heavy states with mass $M$ and a light one with mass $m$, named inverted hierarchy (IH). Near future neutrino oscillation data may resolve the neutrino mass hierarchy if one of the still unknown parameters that relates flavor with mass states is not too small. On the contrary, if that mixing angle is too small, oscillation data may be unable to solve this issue. Analogously, a total neutrino mass determination from cosmology will be able to determine the hierarchy only if the underlying model is normal hierarchy and $\\Sigma<0.1$ eV (see e.g., Fig~\\ref{fig:0}). If neutrinos exist in either an inverted hierarchy or are denegerate, (and if the neutrinoless double beta decay signal is not seen within the bounds determined by neutrino oscillation data), then the three light neutrino mass eigenstates (only) will be found to be Dirac particles. In this paper, we investigate whether cosmological data may positively establish the degenerate spectrum from the inverted hierarchy (or vice versa). Our approach is to take cosmic variance limited surveys, rather than specifically planned experiments, so that we can determine if (even in the ideal case) cosmology can make any impact on this question. ", "conclusions": "The shape of the matter power spectrum contains information, in order of decreasing sensitivity, about the sum of neutrino masses, the amplitude of the mass splitting and the hierarchy (i.e., the mass splitting order). We have introduced a novel parameterization of the neutrino mass hierarchy, $\\Delta$, that has the advantage of changing continuously between normal, degenerate and inverted hierarchies and whose sign changes between normal and inverted. The absolute value of $\\Delta$ describes the maximum mass difference between the eigenstates. We stress that, current constraints from neutrino oscillations have ruled out large part of the parameter space given by the sum of the masses and the $\\Delta$ parameter, leaving two narrow regions: for a fixed value of the total mass, the value of $\\Delta$ for the normal hierarchy is related to that of the inverted one and $\\Delta_{NH}\\simeq -\\Delta_{IH}$ (but, in detail, $\\Delta_{NH} \\not \\equiv |\\Delta_{IH}|$). It is the allowed region that cosmology should explore. We found that the information about $\\Delta$ accessible from the power spectrum shape yields a degeneracy: parameters values $\\Delta$ and $-\\Delta$ yield nearly identical power spectra and therefore that the likelihood surface in $\\Delta$ is bimodal. This was not noted in the literature before and not taking this into account when using the Fisher matrix-approach to forecast future surveys performance may lead to spurious indications of a surveys ability to determine the hierarchy. Detecting the signature of the hierarchy in the sky is therefore extremely challenging, and therefore we asked: ``can cosmology in the cosmic-variance limit, and for an ideal experiment, distinguish the neutrino heirarchy?\" or in other words, ``is there enough information in the sky to measure the neutrino hierarchy?\" To address these questions we have considered ideal, full-sky, cosmic variance-limited surveys and found that substantial Bayesian evidence ($\\ln B\\geq 1$) can be achieved. Are such a surveys feasible in the next $5$-$10$ years? There are two candidates for such surveys : a full extragalactic survey in the optical/infrared like Euclid\\footnote{http://sci.esa.int/euclid} \\cite{Euclid} and a full 21cm survey by the SKA\\footnote{http://www.skatelescope.org}. Each of these surveys is scheduled to start operations by 2018. Euclid will make an all sky Hubble-quality map for weak lensing and will directly trace the dark matter using this technique; whilst the cosmic variance limited survey we consider here is ambitous with respect to this survey these result serve as a qualitative measure of this surveys expected performance (costraints should be only a factor $\\leq 1.5$ larger at worst). Euclid will also target emission line galaxies up to $z \\sim 3$ (therefore these galaxies will have bias parameter close to $1$) however $nP$, quantifying the the ratio of the signal to shot noise, will be only slightly above $1$. The 21cm surveys provide the most un-biased indirect tracer of the dark matter distribution in the Universe and have negligible shot noise. For the degenerate and inverted mass spectra, the next generation neutrinoless double beta decay experiments can determine if neutrinos are their own anti-particle. For the normal hierarchy, the effective electron-neutrino mass may even vanish. However, if the large-scale structure cosmological data, improved data on the tritium beta decay, or the long-baseline neutrino oscillation experiments establish the degenerate or inverted mass spectrum, the null result from such double-beta decay experiments will lead to a definitive result pointing to the Dirac nature of the neutrino mass. This is summarized in figure \\ref{fig:flowchart}. If the small mixing in the neutrino mixing matrix is negligible, cosmology might be the most promising arena to help in this puzzle. Our work shows that depending on the total neutrino mass, there might be substantial evidence by cosmological data to infer the neutrino hierarchy." }, "1003/1003.5773_arXiv.txt": { "abstract": "We present a multi-wavelength, UV-to-radio analysis for a sample of massive (M$_{\\ast}$ $\\sim$ 10$^{10}$ M$_\\odot$) IRAC- and MIPS 24$\\mu$m-detected Lyman Break Galaxies (LBGs) with spectroscopic redshifts z$\\sim$3 in the GOODS-North field (L$_{\\rm UV}$$>$1.8$\\times$L$^{\\ast}_{z=3}$). For LBGs without individual 24$\\mu$m detections, we employ stacking techniques at 24$\\mu$m, 1.1mm and 1.4GHz, to construct the average UV-to-radio spectral energy distribution and find it to be consistent with that of a Luminous Infrared Galaxy (LIRG) with L$\\rm_{IR}$=4.5$^{+1.1}_{-2.3}$$\\times$10$^{11}$ L$_{\\odot}$ and a specific star formation rate (SSFR) of 4.3 Gyr$^{-1}$ that corresponds to a mass doubling time $\\sim$230 Myrs. On the other hand, when considering the 24$\\mu$m-detected LBGs we find among them galaxies with L$\\rm_{IR}>$10$^{12}$ L$_{\\odot}$, indicating that the space density of $z\\sim$3 UV-selected Ultra-luminous Infrared Galaxies (ULIRGs) is $\\sim$(1.5$\\pm$0.5)$\\times$10$^{-5}$ Mpc$^{-3}$. We compare measurements of star formation rates (SFRs) from data at different wavelengths and find that there is tight correlation (Kendall's $\\tau >$ 99.7\\%) and excellent agreement between the values derived from dust-corrected UV, mid-IR, mm and radio data for the whole range of L$\\rm_{IR}$ up to L$\\rm_{IR}$ $\\sim$ 10$^{13}$ L$_{\\odot}$. This range is greater than that for which the correlation is known to hold at z$\\sim$2, possibly due to the lack of significant contribution from PAHs to the 24$\\mu$m flux at $z\\sim$3. The fact that this agreement is observed for galaxies with L$\\rm_{IR}$ $>$ 10$^{12}$ L$_{\\odot}$ suggests that star-formation in UV-selected ULIRGs, as well as the bulk of star-formation activity at this redshift, is not embedded in optically thick regions as seen in local ULIRGs and submillimeter-selected galaxies at $z=2$. ", "introduction": "One of the most fundamental quantities needed for understanding the nature and evolution of galaxies is the star formation rate (SFR). To get reliable and meaningful estimates of the SFR for galaxies at high redshift, one needs a well defined sample of objects, coupled with multi-wavelength data that can provide a thorough and comprehensive investigation. One of the most successful methods of detecting high--z star--forming galaxies is the Lyman-break technique, pioneered by Steidel et al.\\ (1996,2003). This technique has revealed a wealth of Lyman Break Galaxies (LBGs) at $z\\sim$3, now comprising an impressive catalogue of thousands star--forming galaxies at this redshift. Multi-wavelength studies of LBGs have provided extensive information on various physical properties of these objects. In particular, measurements at near-infrared wavelengths and at 3.6-8$\\mu$m from the Spitzer Space Telescope IRAC instrument indicate that their stellar masses are typically 10$^{9}$-10$^{11}$ M$_{\\odot}$ (e.g., Shapley et al.\\ 2001, Papovich et al.\\ 2001, Magdis et al.\\ 2010). The dust content and the SFR of LBGs at $z \\approx 3$ are still poorly constrained. For their siblings at lower redshift, $z\\sim$2, Reddy \\& Steidel (2004) and Reddy et al.\\ (2006), using multi-wavelength data ranging from X-rays to radio, have reported that UV can be a reliable SFR indicator if corrected for dust attenuation by an average factor between 4.4 and 5.1. The validity of the UV as a robust SFR indicator has also been presented by Daddi et al.\\ (2005, 2007) for a sample of near-infrared selected galaxies at 1.5 $<$ $z$ $<$ 2.5 identified using the BzK technique (see also Dannerbauer et al.\\ 2006). A similar multi-wavelength study for the $z\\sim$3 LBGs, though, is still needed. In this letter we make use of the unique compilation of multi-wavelength data on the Great Observatories Origins Deep Survey North field (GOODS-N) to explore the SFR and the infrared luminosities (L$\\rm_{IR}$) of $z\\sim$3 LBGs. Our aim is to fully characterize the spectral energy distribution (SED) of a typical LBG from rest-frame UV to radio wavelengths, to compare different tracers of star formation, and to test whether the UV can provide a reliable measurement of star formation at $z\\sim$3. For this letter we adopt a $\\Lambda$ cold dark matter ($\\Lambda$CDM) cosmology with H$_{0}$= 71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m}$= 0.27 and $\\Omega_{\\Lambda}$= 0.73, while the magnitudes presented in this work are all in the AB magnitude system. ", "conclusions": "Considering the median fluxes of the IRAC-LBGs for the rest--frame UV to NIR (i.e U$_{n}$ to 8.0$\\mu$m) and stacked fluxes at 24$\\mu$m, 1.1mm and 1.4GHz, we construct the average SED of a typical (24$\\mu$m faint) IRAC-LBG. We fit the rest--frame UV to NIR with model SEDs generated using the CB07 code, and the mid-IR to radio with CE01 templates. The photometric points along with the best fit model are shown in Figure 4. The best-fit CB07 model indicates an average stellar mass of M$_{\\ast}$ $\\sim$2.2$\\times$10$^{10}$M$_{\\odot}$ and an average SFR $\\sim$85 M$_{\\odot}$yr$^{-1}$. Using this mass estimate and the SFR derived from our multi-wavelength analysis we derive a specific SFR (SSFR, defined as SFR/M$_{\\odot}$) $\\sim$4.3 Gyr$^{-1}$, corresponding to a mass doubling time of $\\sim$ 230Myrs. This value is very close to the one presented by Magdis et al.\\ (2010) (4.5 Gyr$^{-1}$) and is larger than that found at lower and higher redshifts, reinforcing their argument that the evolution of the SSFR peaks at $z\\sim$3. Based on CE01 models we find that the average IRAC-LBG is a LIRG with L$\\rm_{IR}$=4.5$^{+1.1}_{-2.3}\\times$10$^{11}$ L$_{\\odot}$. We note that L$_{\\rm IR}$ corresponds to a typical dust temperature T$_{\\rm d}$=35K, based on T$_{\\rm d}$ measurements of local LIRGs (Yang et al. 2007). On the other hand, MIPS-LBGs have higher luminosities, indicating that ULIRGs are present among the UV-selected galaxies at $z\\sim$3. For instance, HDFN-M23 which was individually detected in the radio, has L$\\rm_{IR}\\sim$5($\\pm$2)$\\times$10$^{12}$ L$_{\\odot}$. We calculate the comoving volume for the redshift range of our sample (2.5$<$ z $<$3.5) and find that the space density of ULIRGs (based on UV L$\\rm_{IR}$) LBGs is $\\sim$1.5$\\pm$0.5$\\times$10$^{-5}$ Mpc$^{-3}$. This is a factor of $\\sim$10 smaller than the space density of the $z\\sim$2 BzK-selected ULIRGs (Daddi et al.\\ 2007). We note that the UV selection is likely to miss more obscured, UV faint ULIRGs as our sample is limited to R$_{\\rm AB}$ $<$ 25.5 while, for example, an Arp220 like object at $z\\sim$3 would have R$_ {\\rm AB}\\sim$26.9. The good agreement between SFR indicators that are affected by dust extinction (UV) or not (MIR, mm, radio) suggests that the bulk of the star formation activity in massive UV-selected galaxies takes place in optically thin regions. Since by focusing on the massive LBGs we also select those that are most affected by dust extinction (e.g., Magdis et al.\\ 2010), it is reasonable to extend this result to the whole population of UV galaxies including the less massive, less dusty LBGs. Furthermore, the fact that the agreement holds even for the case of the most massive, dusty, 24$\\mu$m-detected LBGs with L$\\rm_{IR}>$10$^{12}$ L$_{\\odot}$, indicates that UV-selected ULIRGs at z=3 are transparent to UV light, meaning that we can estimate its SFR given its rest-frame UV 1500 $\\rm \\AA$ luminosity and UV slope, contrary to the local ULIRGs and z=2 SMGs. Similar results have been reached by Reddy et al. (2006) for UV-selected and by Daddi et al. (2007) for BzK selected z$\\sim$2 LIRGs and ULIRGs. Both studies find that galaxies of a given bolometric luminosity are on average a factor of 8-10 less dust obscured at z$\\sim$2 than at the present epoch. The fact that our study suggests that at z$\\sim$3 ULIRGs are optically thin at even higher L$_{\\rm IR}$ than that at z$\\sim$2, indicates even less obscuration at z$\\sim$3 confirming the trend between galaxies at z = 0, z $\\sim$ 1, and z $\\sim$ 3 (Adelberger $\\&$ Steidel 2000). This could plausibly be explained as a result of increasing dust-to-gas ratios as we move from the high-z to the present universe. As galaxies evolve, they convert gas into stars which in turn enrich the interstellar medium with dust. If the dust distribution becomes more compact with time (assuming mergers that drive dust and gas to the center of the galaxy) the overall result would be an increase of the dust column density toward star-forming regions making ULIRGs progressively optically opaque at later epochs. To put constraints on the census of the cosmic star formation at z=3, we should also consider the missing fraction of star formation embedded in optically thick regions, that takes place in galaxies not selected in UV. On the other hand, there is evidence of decreasing obscuration with increasing redshift for a given L$\\rm_{IR}$ (e.g., Reddy et 2008), pointing towards less optically thick star-formation at higher redshifts. Combining the above with the fact that the bulk of the currently known SMGs are at $z\\sim$2.2 (Chapman et al.\\ 2005), we can assume that at z=3 the contribution of the SMGs to the total SFR density, is not dominant, and hence locate the census of high-z star-formation in optically thin regions. We stress, that such a scenario cannot be confirmed based solely on our sample, as it has been shown that the Lyman Break technique can miss a large fraction of massive (dusty) galaxies at high-z (e.g., Daddi et al. 2004, Van Dokkum et al.\\ 2006). Further insights into the far-IR properties of the LBGs, such as dust temperature and dust mass will be provided by deep surveys with the Photodetector Array Camera \\& Spectrometer (PACS, 70-,100-,160$\\mu$m) and the Photometric Imaging REceiver (SPIRE, 250-,350-500$\\mu$m) on board the Herschel Space Observatory (HSO). Based on the average spectrum of Figure 4, the predicted flux densities in the HSO bands of a typical IRAC-detected LBG are $f_{100}$=0.13 mJy, $f_{160}$=0.47 mJy, $f_{250}$=1.18 mJy, $f_{350}$=1.52 mJy and $f_{500}$=1.22 mJy." }, "1003/1003.5034_arXiv.txt": { "abstract": "We suggest a possible explanation for the high frequency quasi-periodic oscillations (QPOs) in black hole low mass X-ray binaries. By solving the perturbation general relativistic magnetohydrodynamic equations, we find two stable modes of the Alf\\'ven wave in the the accretion disks with toroidal magnetic fields. We suggest that these two modes may lead to the double high frequency QPOs if they are produced in the transition region between the inner advection dominated accretion flow and the outer thin disk. This model naturally accounts for the $3:2$ relation for the upper and lower frequencies of the QPOs, and the relation between the black hole mass and QPO frequency. ", "introduction": "Low-mass X-ray binaries (LMXBs) are binary systems consisting of a neutron star (NS) or black hole (BH) accreting from a low-mass ($\\lesssim 1 M_{\\sun}$) companion star. X-ray emission of LMXBs often shows fast X-ray variability in the form of high frequency quasi-periodic oscillations (HFQPOs), which frequently appear in pairs in certain state simultaneously (van der Klis 2006). Abramowicz \\& Klu\\'zniak (2001) pointed out that the frequency ratio of the twin-peak HFQPOs in the BH source GRO J1655-40 equals 3/2, and that this commensurability of frequencies may be a signature of a non-linear resonance. Later, Abramowicz et al. (2003) found a signature of the same commensurable ratio in the twin-peak HFQPOs observed in an NS-LMXB, Sco X-1. Based on this observational evidence, Klu{\\'z}niak and Abramowicz argued in several papers (Klu{\\'z}niak \\& Abramowicz 2001; Klu{\\'z}niak, Abramowicz, \\& Lee 2004; Klu{\\'z}niak et al. 2004) that the twin-peak HFQPOs in both BH and NS sources are due to the same physical mechanism --- a non-linear parametric resonance in accretion disk global oscillations. However, while for the BH sources the commensurable ratio 3/2 was quickly confirmed and generally accepted (e.g. Remillard \\& McClintock 2006), the presence of the same commensurability in the NS sources is denied by several experts (e.g. Boutelier et al., 2009). There is no consensus whether the nature of the twin-peak HFQPOs in the two types of LMXBs is the same. We have proposed a mechanism for the twin kilohertz QPOs in NS-LMXBs using the magnetohydrodynamic (MHD) Alf\\'ven wave oscillations, and the results seem to fit the observation well (Li \\& Zhang 2005; Shi \\& Li 2009). In this paper we focus on an MHD explanation of the HFQPOs in BH-LMXBs. Barret et al. (2005) measured the quality factor Q for the HFQPOs measured in the NS-LMXB 4U 1608-52, and found that Q $\\thicksim$ 200. They argued that such high coherency is impossible to achieve from kinematic effects in orbital motion of hot spots, clumps or other similar features located at the accretion disk surface, because these features are too quickly sheared out by the differential rotation. Although orbital motion cannot explain the the HFQPOs in LMXBs, the frequencies of several fluid oscillatory modes are expressed by the three characteristic orbital frequencies: the ``Keplerian\" frequency, the ``radial\" epicyclic frequency, and the ``vertical\" epicyclic frequency. In the Kerr metric, these three orbital frequencies and the Lense-Thirring ``frame-dragging\" frequency have been listed (e.g. Perez et al. 1997). Several HFQPOs models use their ratios (in various combinations) to explain the observed 3/2 commensurability. Cui, Zhang, \\& Chen (1998) suggested the Lense-Thirring nodal precession frequency near the inner stable circular orbit (ISCO) radius as the lower HFQPO frequency, such as the 300 Hz QPOs in GRO J1655-40. The relativistic precession model of Stella et al. (1999) applies to both BH and NS sources; the pariastron precession frequency and the Keplerian frequency were taken as the lower and upper frequencies of the twin HFQPOs, respectively, whereas the QPOs at much lower frequencies were interpreted in terms of the Lense-Thirring nodal precession ferquency. Wang et al. (2003, 2005) suggested that a non-axisymmetric magnetic coupling of a rotating BH with its surrounding accretion disk coexists with the Blandford-Znajek process. The two frequencies were supposed as the Keplerian frequencies of two hotspots, one near the inner edge of the disk and the other somewhere outside, respectively. Wagoner et al. (2001) considered the modes of the diskoseismic wave, such as g-modes, p-modes and c-modes as the explanation of the HFQPOs. They estimated the masses and angular momenta of some BHs with the measured frequencies of the HFQPOs when the g-modes or c-modes were selected. Rezzolla et al. (2003) discussed the inertial-acoustic modes in a small-size torus very close to the horizon of the BH while the centrifugal and pressure gradients were selected as the only restoring forces. In this model the black hole spin had to be very close to the maximal value to produce the $3:2$ ratio. Tassive \\& Bertschinger (2008) investigated the kinematic density waves in the accretion disks when nothing but the gravity was considered as the restoring force, and several discrete radii were adopted. Several modes in pairs close the ratio ($3:2$) could be got, but the correct frequencies couldn't be reproduced. Similar to the parametric resonance models, it is difficult to explain why other modes such as the fundamental frequencies weren't observed except the two modes in pairs. In modeling HFQPOs in BH-LMXBs two points need to be mentioned. First, in most models the $3:2$ ratio was often overemphasized and substituted into those models directly. In fact, the $3:2$ ratio of the twin HFQPOs in BH-LMXBs isn't rigorous but approximate. Second, the HFQPO frequencies were often considered invariable so that these frequencies (168, 113 Hz and 67, 41Hz in GRS 1915+105, Remillard (2004); 450, 300 Hz in GRO J1655-40, Strohmayer (2001a); 276, 184 Hz in XTE J1550-564, Miller et al. (2001); 240, 165 Hz in H1743-322, Homan et al. (2005)) were used to estimate the parameters of the BHs such as the masses and the spins. In fact, the frequencies are stable, i.e. there are small varitions in them, rather than invariable (Miller et al. 2001; Remillard et al. 1999; Morgan et al. 1997; Strohmayer 2001b). This paper is organized as follows. In Section 2 we suggest the basic model and get two stable modes of the Alf\\'ven wave by solving the general relativistic magnetohydrodynamic (GRMHD) equations of the perturbed plasma in BH accretion disks. In Section 3 we compare the results with the observations and discuss their possible implications. ", "conclusions": "Equations (63), (64), and (72) indicate the existence of Alf\\'ven waves with two frequencies in the accretion disk. The ratio of the upper and lower frequencies is close to $3:2$, suggesting that these waves may account for the HFQPO pairs. Given the mass and spin of a BH, one can determine the radius where the QPOs are produced. In Table 1 we present the inferred radius by comparing Eq.~(72) with either the upper or the lower centroid QPO frequency of several BH-LMXBs. We adopt the averaged masses for GRO J1655$-$40 ($6.0-6.6M_\\odot$, McClintock \\& Remillard 2006), GRS 1915+105 ($10-18M_\\odot$, Greiner et al. 2001), and XTE J1550$-$564 ($8.4-10.8 M_\\odot$, McClintock \\& Remillard 2006) and averaged dimensionless spins of GRO J1655$-$40 ($0.65-0.75$, Shafee et al. 2006) and GRS 1915$+$105 ($0.98-1.0$, McClintock \\& Shafee et al. 2006). Because the spin of XTE J1550$-$564 hasn't been measured, we take it to be 1, 0.5, and 0. It is interesting to see that in most cases the radii are $\\sim 70 r_{\\rm g}$, consistent with the transition radii between an ADAF and a thin disk discussed above. The frequencies of the HFQPOs mainly depends on the transition radius ($r_{\\rm tr}$) and the ratio of $l'=(v_{\\rm r}/v_{\\varphi})/(B_{\\rm r}^{'}/B_{\\varphi}^{'})$. According to the discussion in \\S2.3 the parameter $l'$ changes little with the $\\alpha_*$, $r$, $M$, and $\\dot{M}$, this may explain why the HFQPO frequencies are relatively stable during the VHS. Considering the similarities in BH and NS accretion disks, it is interesting to ask why the 1.5 frequency ratio is not evident in the HFQPOs in NS-LMXBs. The reasons may lie in the differences in the configuration of the magnetic fields and the structure of the accretion disks in BH- and NS-LMXBs. It is well known that NSs generally hold dipolar magnetic fields, and the toroidal field component, induced by difference in the angular velocity between the disk and the NS magnetosphere, can never be stronger than the poloidal component, otherwise the field configuration becomes unstable leading to inflation of the field lines (e.g. Aly 1985). On the other hand, both theories and simulations (Ruzmaikin et al., 1979; Tout et al., 1992; Ruediger et al., 1995; Hawley et al., 2000; Moss et al, 2004; Hirose et al., 2004) show that in BH accretion disks the toroidal magnetic field can be much stronger than the poloidal one. Additionally, recent observations suggest that the inner disk around an NS may not be an ADAF even in low state, as the thermal emission from the surface of the NS would tend to cool an ADAF, making such a flow more difficult than in BH systems (Cackett et al., 2009). Finally, a correlation between the lower HFQPO frequencies and the BH masses has been suggested, i.e., $\\nu_{\\rm l}\\propto M^{-1}$ (McClintock \\& Remillard 2006). This correlation can be naturally reproduced in our model. In Fig.~1 we plot the predicted relation between $\\nu_{\\rm l}$ and $m$ when $r_{\\rm tr}$ changes from $66 r_{ g}$ to $76 r_{ g}$, which fit reasonably with the measured data." }, "1003/1003.2422_arXiv.txt": { "abstract": "{We present a suite of full hydrodynamical cosmological simulations that quantitatively address the impact of neutrinos on the (mildly non-linear) spatial distribution of matter and in particular on the neutral hydrogen distribution in the Intergalactic Medium (IGM), which is responsible for the intervening \\lya absorption in quasar spectra. The free-streaming of neutrinos results in a (non-linear) scale-dependent suppression of power spectrum of the total matter distribution at scales probed by Lyman-$\\alpha$ forest data which is larger than the linear theory prediction by about 25~\\% and strongly redshift dependent. By extracting a set of realistic mock quasar spectra, we quantify the effect of neutrinos on the flux probability distribution function and flux power spectrum. The differences in the matter power spectra translate into a $\\sim2.5\\%$ ($5\\%$) difference in the flux power spectrum for neutrino masses with $\\Sigma m_{\\nu} = 0.3$ eV (0.6 eV). This rather small effect is difficult to detect from present Lyman-$\\alpha$ forest data and nearly perfectly degenerate with the overall amplitude of the matter power spectrum as characterised by $\\sigma_8$. If the results of the numerical simulations are normalized to have the same $\\sigma_8$ in the initial conditions, then neutrinos produce a smaller suppression in the flux power of about 3\\% (5\\%) for $\\Sigma m_{\\nu} = 0.6$ eV (1.2 eV) when compared to a simulation without neutrinos. We present constraints on neutrino masses using the Sloan Digital Sky Survey flux power spectrum alone and find an upper limit of $\\Sigma m_{\\nu} < 0.9$ eV (2$\\sigma$ C.L.), comparable to constraints obtained from the cosmic microwave background data or other large scale structure probes.} \\begin{document} ", "introduction": "One of the most exciting results in particle physics in the last decade has been that neutrinos have been established to be massive particles. Solar, atmospheric, reactor and accelerator neutrino experiments have confirmed the existence of flavour oscillations of active neutrinos, implying that neutrinos have non-zero mass (see Ref.~\\cite{lespast} and references therein). This is generally considered as definite evidence for new physics beyond the Standard Model. The neutrino oscillation experiments do, however, not pin down the absolute neutrino masses. The experiments instead provide a lower limit for the sum of the neutrino masses of $0.05-0.1$ eV. Current measurement of the matter power spectrum from Cosmic Microwave Background (CMB) data extrapolated to smaller scales alone already give an upper limit on the sum of the neutrino masses of about $1.5$ eV well below what has been reached with particle physics experiments leaving an allowed range of only a factor about twenty for the sum of the neutrino masses. There is thus very strong motivation to push hard for an actual measurement of neutrino masses. The tritium $\\beta$-decay experiment {\\small{KATRIN}}\\footnote{http://www-ik.fzk.de/$\\sim$katrin} is the most ambitious current direct detection experiment and is expected to probe an electron neutrino mass of $\\sim 0.2$ eV in the near future (see \\cite{fogli} for a recent review). The matter distribution in the Universe is sensitive to the free-streaming of cosmological neutrinos. Astrophysical constraints are therefore a very competitive alternative method to measure/constrain the masses of neutrinos. Measurements of the matter power spectrum can in principle probe neutrino masses significantly smaller than the upper limit from CMB experiments. Early on the neutrinos are relativistic and travel at the speed of light with a free-streaming length equal to the Hubble radius. Neutrinos in the mass range $0.05$ eV $\\le \\Sigma m_{\\nu} \\le 1.5$ eV, become non-relativistic in the redshift range $3000\\ge z \\ge 100$. In the mass range of degenerate neutrino masses the thermal velocities can be approximated as, \\begin{equation} v_{\\rm th}\\sim 150\\, (1+z)\\,\\left[\\frac{1\\,\\rm{eV}}{\\Sigma m_{\\nu}}\\right] \\rm {km/s} \\,. \\label{eqvel} \\end{equation} As a result present-day velocities (of the most massive neutrino species) range between 100 km/s for the upper and and 3000 km/s for the lower end of the still allowed range of the sum of the neutrinos masses. Dark matter particles with such a high velocity dispersion are usually called hot dark matter. A dominant contribution of hot dark matter to the total dark matter content would be at odds with current observations. Neutrinos in the still allowed mass range instead constitute a sub-dominant contribution complementing cold dark matter comprised of some other elementary particle, such as neutralinos or axions. The effect of cosmological neutrinos on the evolution of density perturbation in the linear regime is well understood. Neutrinos affect both the cosmic expansion rate and the growth of structure (\\cite{maber95,lespast}). The neutrino contribution in terms of energy density can be expressed as: \\begin{equation} f_{\\nu}=\\Omega_{0\\nu}/\\Omega_{\\rm 0m}\\,,\\,\\,\\,\\,\\,\\,\\,\\,\\Omega_{0\\nu}=\\frac{\\Sigma\\, m_{\\nu}}{93.8\\,h^2\\rm{eV}}, \\end{equation} where $h$ is the present value of the Hubble constant in units of 100 km/s/Mpc and $\\Omega_{\\rm 0m}$ is the matter energy density in terms of the critical density. When neutrinos become non relativistic in the matter dominated era, there is a minimum wavenumber \\begin{equation} k_{\\rm nr}\\sim 0.018\\,\\Omega_{\\rm 0m}^{1/2} \\left[\\frac{\\Sigma m_{\\nu}}{1\\, {\\rm eV}}\\right]^{1/2} h/{\\rm Mpc}\\, , \\end{equation} above which the physical effect produced by neutrino free-streaming damps small-scale neutrino density fluctuations, while modes with $k0.1 h/$Mpc the suppression is constant while at scales $0.01 $10MG) white dwarfs have been termed LARPS (Low Accretion Rate Polars; Schwope et al. 2002). Since the low mass transfer means the accretion luminosity is low, the individual stellar components can be viewed. Thus, these systems provide a unique comparison to widely separated binaries for the study of how angular momentum losses and heating of the white dwarf affect the binary evolutionary scenario. While normal polars can reach mass accretion rates close to those of LARPS when they undergo long periods of low mass transfer, the main difference between LARPS and normal polars in low states appears to be the temperature of the white dwarf. Ultraviolet observations of normal polars have shown temperatures of 11-14,000K for their white dwarfs (Araujo-Betancor et al. 1995). In contrast, the white dwarfs in the LARPS are less than 10,000K (Schmidt et al. 2005, 2007, hereafter S05, S07; Vogel et al. 2007, hereafter V07; Schwope et al. 2009, hereafter Sw09). These low temperatures imply that compressional heating has not taken place and so these objects have been suggested to be pre-Polars (S05) or PREPs (Sw09). Figure 5 of Sw09 shows a nice plot of the temperatures of 9 PREPs compared to normal polars and accreting non-magnetic white dwarfs that illustrates this difference. However, the terminology and behavior is not always so cleanly separated. The system SDSS121209.31+013627.7 (Schmidt et al. 2005b; Burleigh et al. 2006; hereafter B06) also has a low accretion rate and a low white dwarf temperature but has a brown dwarf secondary. Its evolutionary path has been suggested to be either a polar in a low state or a PREP. In addition, the recent observation of enhanced activity levels in the LARP SDSSJ204827.9+005008.9 (Honeycutt et al. 2010) suggest that it may not be a LARP/PREP or that even PREPs can have different levels of activity. A secondary criteria for a LARP/PREP is the secondary underfilling its Roche lobe so that Roche lobe overflow does not occur (S05; S07; Sw09). This is generally determined from the system having an spectral type too late to fill the lobe if it has a normal size for its type, as well as the absence of narrow components in the emission lines that can be traced to a stream. However, the Honeycutt et al. (2010) study found evidence for a stream in SDSSJ204827.9+005008.9, although they could not determine if the system had entered an increased state of mass transfer at the time. While the low accretion rates and sizes of the secondaries imply no Roche lobe overflow, there could likely still be some accretion from the stellar wind of the secondary that is funnelled by the high field of the white dwarf (Li et al. 2004). The work of Webbink \\& Wickramasinghe (2005) using considerations of energy, magnetic fields, separations and lifetimes supports a wind model for the LARPS. In this study, we are not so concerned with the correct classification of our objects as LARPS, PREPs or normal polars in low states, or how the accretion might occur, but rather with the effects of low levels of accretion on the white dwarf. These effects are visible through time-resolved UV observations. Accretion spots with temperatures of 30,000-70,000K are easily visible on normal polars even during their sporadic low states of accretion with $\\dot{M}\\sim$10$^{-12}$M$_{\\odot}$ yr$^{-1}$ (Araujo-Betancor et al. 2005; G\\\"ansicke et al. 2006) as large amplitude modulations of the UV light. However, it was somewhat surprising to find that similar modulations were apparent for EF Eri (Szkody et al. 2006) which had been in a low state for seven years with $\\dot{M}$ of 10$^{-13}$ M$_{\\odot}$ yr$^{-1}$ and the LARP MQ Dra with an even lower $\\dot{M}$ of 10$^{-14}$ M$_{\\odot}$ yr$^{-1}$ (Szkody et al. 2008). While the UV light curves could be approximately matched with either hot spots or cyclotron components (Campbell et al. 2008), each interpretation has its problems: the hot spots require different sizes and geometries while the cyclotron origin requires much higher magnetic fields in the UV than are apparent in the optical. In order to further study the effects of this low level accretion, we used {\\it GALEX} to obtain UV observations of two additional LARPS (WX LMi and SDSSJ103100.55+202832.2) and the system SDSSSJ121209.31+013627.7. For convenience, we will refer to the SDSS objects as SDSS1031 and SDSS1212. The known parameters for the three systems are listed in Table 1. While all three of these objects have comparable low mass accretion rates, they present differences in magnetic field strength, white dwarf temperature, spectral type of secondary and orbital period. ", "conclusions": "Our UV observations of three polars with extremely low accretion rates enforce the results found from observations of EF Eri and MQ Dra i.e. that all the white dwarfs have areas of enhanced emission even with these low rates, indicating some accretion is still occuring. Despite a large range in magnetic field (7-70MG), the field is apparently strong enough in all three systems to funnel the accreting material to the magnetic pole(s) of the white dwarf. This appears to happen even in the case of SDSS1212, which likely has a brown dwarf secondary. Thus, if the accretion is via a stellar wind from the secondary that is trapped by the field of the white dwarf, it is difficult to provide a wind of this type for SDSS1212. Time-resolved spectra that would enable Doppler tomography might provide a resolution of this issue. During high states of accretion, the stream flow is usually visible as a narrow component in the Balmer emission lines which is mapped to the stream. In normal polars during their low states, high velocity components in the lines have been interpreted as structures similar to prominences close to the secondary star (Kafka et al. 2007, 2008). However, the Balmer emission in LARPS is weak to non-existent due to the low the accretion rates low so it would be difficult to construct the map. If the increased emission evident at some phases is due to hotter temperature, simple spots of 10-000-14,000K covering a few percent of the white dwarf surface can approximate the UV light curves, although the geometries of the spots require a more complex model shape than simply circular. Alternatively, if there is a large range in field strengths in the white dwarfs, there could be contributions to the UV from higher cyclotron harmonics. Improved UV cyclotron models will be needed to test this possibility. Since these objects are too faint for {\\it GALEX} spectra or high S/N time-resolved UV light curves that would merit detailed models of the shape or of pursuing cyclotron radiation, further work will require UV spectra and polarimetry with larger telescopes." }, "1003/1003.0082_arXiv.txt": { "abstract": "An analysis is made of the manner in which the cosmic ray intensity at Earth has varied over its existence and its possible relevance to both the origin and the evolution of life. Much of the analysis relates to the 'high energy' cosmic rays ($E>10^{14}eV;=0.1PeV$) and their variability due to the changing proximity of the solar system to supernova remnants which are generally believed to be responsible for most cosmic rays up to PeV energies. It is pointed out that, on a statistical basis, there will have been considerable variations in the likely 100 My between the Earth's biosphere reaching reasonable stability and the onset of very elementary life. Interestingly, there is the increasingly strong possibility that PeV cosmic rays are responsible for the initiation of terrestrial lightning strokes and the possibility arises of considerable increases in the frequency of lightnings and thereby the formation of some of the complex molecules which are the 'building blocks of life'. Attention is also given to the well known generation of the oxides of nitrogen by lightning strokes which are poisonous to animal life but helpful to plant growth; here, too, the violent swings of cosmic ray intensities may have had relevance to evolutionary changes. A particular variant of the cosmic ray acceleration model, put forward by us, predicts an increase in lightning rate in the past and this has been sought in Korean historical records. Finally, the time dependence of the overall cosmic ray intensity, which manifests itself mainly at sub-10 GeV energies, has been examined. The relevance of cosmic rays to the 'global electrical circuit' points to the importance of this concept. ", "introduction": "\\subsection{The cosmic radiation and its relation to climate} Over the 4.5By since the formation of the Earth the astronomical environment has been variable and, with it, the cosmic ray (CR) spectrum (~this is in addition to solar irradiation changes caused by the varying Sun-Earth distance and Earth obliquity - the 'Milankovich effect'~). There are three main sources of CR variability: the Geomagnetic field, the Sun (~by way of the solar wind and the occasional 'solar cosmic rays' associated with solar flares~) and the presence of nearby CR sources. The first two relate to variations of the low energy particles, principally below 10 GeV, and the last-mentioned to all energies, with increasing variability as the energy increases, the highest energy recorded being $\\sim 10^{20}$eV (~ie $10^{11}$GeV~). Further remarks on 'cosmic rays' are necessary. Under 'normal circumstances' (~quiescent Sun~) the CR, composed mainly of hydrogen, helium and heavier nuclei, have a power law spectrum from about 1000 GeV to 3 PeV. At energies above 3 PeV the spectrum steepens. Below 1000 GeV there is a progressive 'flattening' of the spectrum (~as one proceeds towards lower energies~) due to both the Geomagnetic field and the solar wind. The magnitude of the flattening depends on the 11 year cycle of the solar wind and is a function of geographic latitude and energy; it is very small above 100 GeV. The energy content of Galactic cosmic rays (~GCR~) is an important parameter. The energy densities at earth above the energies indicated are, in units of $Jm^{-3}$, $\\sim 0.05 (>10^8eV), \\sim 0.03 (>10^{10}eV), 3\\cdot 10^{-3} (>10^{12}eV), 10^{-4} (>10^{14}eV)$ and $10^{-6} (>10^{16}eV)$ (Wolfendale, 1973). By contrast, the total solar irradiance is some $10^8$ times greater than that for all CR. In addition to an interest in its own right, the variation of the GCR energy spectrum with time has relevance to atmospheric properties and thus (perhaps) to `life' on Earth. The first mention of a possible effect of CR on climate seems to be that of Ney (1959), the claimed mechanism being by way of the effect of CR ions on aerosols, leading to enhanced cloud formation. A number of workers have, more recently, followed up the suggestion (eg Svensmark and Friis-Christensen, 1997; Palle Bago and Butler, 2000; Marsh and Svensmark, 2000; Svensmark, 2007). Indeed, Svensmark (2007) has even coined a term for the new discipline: 'Cosmoclimatology' ! Although, a priori, it might be thought that such a mechanism was unlikely, based on the $10^8$ factor referred to above, Voiculescu et al.(2006) have claimed that there is a GCR, low cloud cover correlation over restricted regions of the Globe. Harrison and Ambaum (2008) have claimed that 'the mechanism' exists on the edges of clouds by way of a large reduction in the critical supersaturation needed because of a large degree of droplet charging. This is where it must be pointed out that direct CR may affect the climate. The effect of GCR on the 'Global electrical circuit' (~Williams, 2002; Rycroft et al., 2008~) has been studied by Tinsley (2008) and Aplin et al. (2008) who consider the effect of 'electro-freezing' and 'electro-scavenging'; these processes lead to changes in the density of cloud condensation nuclei. In all these processes it is important to know the manner in which CR ionization varies with atmospheric depth. Such studies have been made by Usoskin and Kovaltsov (2006) and Velinov et al.(2009). Despite doubts about the appreciable effect of GCR on clouds in the lower troposphere, (~Sloan and Wolfendale, 2008; Erlykin et al., 2009b~), in the stratosphere, where GCR intensities are higher and there are the occasional 'solar protons', there are strong signals. That there is an 11-year cycle in the ozone density is beyond doubt, and Lu (2009) has presented strong evidence for a strong causal correlation between GCR and polar ozone loss over Antarctica (~the 'ozone hole'~). The identification of GCR as being responsible, as distinct from solar irradiance, comes from CR being the only source of low energy electrons at the depth in question and the ozone hole being in the Geomagnetic Pole region where the CR intensity is a maximum. Other related effects include the effect of strong Geomagnetic storms and Forbush decreases of GCR intensity on the total ozone content and the lower atmosphere (~troposphere and lower stratosphere~) (~Lastovicka and Krizan, 2005~). It must be mentioned, however, that the claimed effect of Forbush CR decreases on the liquid cloud fraction in the troposphere (~and other atmospheric parameters~) by Svensmark et al. (2009) was not confirmed by Laken et al. (2009). A less certain, but potentially important, process, is that suggested by Shumilov et al. (1996). These workers were impressed by the increase in aerosol concentration after solar proton events (~particularly the 'Ground Level Event' of $16/02/1984$. The increase occured at the altitude range, $\\sim$17km, where energetic solar protons lose their energy in the atmosphere~). The mechanism put forward was CR ionization, as a source of ion nucleation, stratospheric sulphate aerosols forming on the condensation nuclei (~Arnold, 1982; Hofman and Rosen, 1983~). The importance of solar particle events in polar regions has been pointed out by Usoskin et al. (2009). A related argument, which may have relevance to the claimed low cloud cover, GCR correlation, is due to Kudryavtsev and Jungner (2005). These workers argue that the extra CR induced aerosols cause atmospheric transparency changes which in turn affect tropospheric climate. They quote other workers (~eg Starkov and Roldugin, 1994 and Pudovkin et al., 1997) as having also observed transparency decreases during solar proton events. \\subsection{Extensive Air Showers} Our work reported here relates mainly to much higher energies than those concerned with solar effects (~which are mainly below some 10s of GeV~), specifically 0.1PeV and above; these particles being manifest by their production of extensive air showers (EAS). Some remarks about EAS are necessary. When a particle (~often a proton~) of 'high energy', say above 0.1PeV, is incident on the atmosphere it interacts with the air nuclei to produce a cascade of secondary particles (~mainly pions of the three charge states: positive, negative and zero~). The (unstable) charged pions decay into muons ('heavy electrons'), which in turn may decay into electrons, and the neutral pions decay into gamma rays. The process is repeated by the primary proton which survives the interaction with reduced energy and those energetic pions which interact further before they have had chance to decay. The effect of the interactions is to build up an 'electromagnetic cascade' of electrons and gamma rays. To be quantitative, a 1PeV primary will generate a shower having maximum number of charged particles (~mainly electrons~) at an atmospheric depth of $\\sim$560mb, ie height in the atmosphere of $\\sim$5km. The mean number of particles at 'shower maximum' would be $\\sim 5.3\\cdot 10^5$ and the number of particles at ground level for a vertically incident proton would be $\\sim 1.4\\cdot 10^5$ (Ambrosio et al., 1997). The important feature of EAS relevant to the initiation of lightning strokes is the high density of secondary electrons near the axis of the shower. Recent calculation by us (~Erlykin et al.,2009a~) for 100PeV protons and a height of 2km give a particle density at 1m from the shower axis of $\\sim 2\\cdot 10^5 m^{-2}$. For a primary of energy 1PeV the density will be $\\sim 2\\cdot 10^3 m^{-2}$, still a very high value. On the axis itself, the particle density will be some ten times greater. \\subsection{Relevance of EAS to the origin of life} The possible relevance to the atmosphere (~and `life', including humans~), is by way of the very likely role of EAS particles in the initiation of lightning (eg Gurevich and Zybin, 2001, Chubenko et al, 2009, Gurevich et al, 2009 and Chilingarian et al, 2009). The idea is that the leader lightning stroke is initiated by runaway electrons which are generated by particles in EAS. The references quoted include observed coincidences between EAS and lightning and not just the undoubted effect of thunderstorm electric fields on the energies of CR particles (~which are, themselves, not necessarily members of EAS~). The whole question of the electrical conditions of the atmosphere, including its most dramatic manifestation (~lightning~), is tied up with CR insofar as they represent an important source of ions near ground level and the major source at altitudes above a few km. Tinsley et al, (2007), Rycroft et al, (2008), and others have pointed out the great importance of the `global electric circuit' - to which CR contribute considerably - even when the changes considered have been small. The global electric circuit possibly can be influenced also by gigantic red sprites and blue jets. Presumably they are initiated by electrons, which are created by CR and accelerated in their upward movement to runaway energies by thunderstorm electric fields (~Yukhimuk et al. 1998; Tonev and Velinov, 2003~). Effects consequent upon very large changes in CR intensity at Earth could be profound. The question of the global electrical curcuit is considered in more detail later. Lightning has, conceivably, played a role in the evolution of life. Starting with pre-life, the work of Miller and Urey (~Miller, 1953~) involving the passage of electrical discharges through a `pre-biotic soup' of appropriate chemicals (~water, methane, ammonia, etc~) caused quite complex molecules to be generated: amino acids, monomers, RNA etc, which were necessary pre-cursors of elementary life. Lightning could, conceivably, have provided the required discharges. It must be remarked that there are different views about the origin of life; some argue that the initial complex molecules arrived by way of comets instead (~eg Hoyle and Wickrama-singhe, 1993~). Here, we persist with the Miller and Urey hypothesis (MU) since very recent work (~eg Parman, 2009~) concludes that there was little free oxygen in the atmosphere prior to 2.45 Gy BP (~ie $2.45\\cdot 10^9$ years before present~) and the MU hypothesis would have a chance of success. The necessary water oceans were probably present by 4.2 Gy BP (`Bada, 2003; Parman, 2009~). Later, when `life' was advanced, lightning would have had an effect on evolution by virtue of the obnoxious NO$_{x}$ (`NO and NO$_{2}$) produced. Even now, some 20\\% of NO$_{x}$ comes from lightning - much higher lightning rates would have been important. NO$_{x}$ effects include modifications to atmospheric chemistry, with particular relevance to ozone levels and effects on hydroxyl (OH) radicals - thereby increasing the concentration of greenhouse gases. There is a wealth of literature on NO$_x$ production by lightning (~eg Betz et al., 2008~). Allen et al. (2009) estimate that the rate of production of NO$_x$ from lightning is $\\sim 10^{13}kg \\cdot year^{-1}$, to be compared with the total atmospheric mass of $5\\cdot10^{18}kg$ and a mass of ozone of order $5\\cdot10^{11}kg$ (~Allen, 1973~). Although NO$_x$ is damaging to mammals, plants benefit from the nitrates coming from NO$_{x}$ reactions. The interplay between the development of plants and of mammals means that periods of low CR intensity (~low NO$_{x}$~) as well as those of high intensity are important. The likelihood of CR effects here is, no doubt, less contentious and should be put alongside the various meteorological factors, such as temperature and rainfall, which have affected the evolution of life. Even if none of the above effects turn out to be important, a knowledge of the past history of the intensity of high energy GCR (HECR), by which we mean $10^{14}$ eV and above, is of considerable interest because of its relevance to the (~still unsolved~) problem of the origin sites, acceleration mode and propagation characteristics of the primary particles. \\subsection{Scope of the paper} We start with an analysis of the time variation on a statistical basis using results provided by us earlier (~Erlykin and Wolfendale, 2001a~), and based on our supernova remnant (SNR) model of GCR acceleration (~Erlykin and Wolfendale, 2001b~). Later we examine the recent past - some 30,000 y - assuming that our Single Source Model of the `knee' in the spectrum at $\\sim$3PeV (~Erlykin and Wolfendale, 1997, 2003~) is correct. In this model, which is now being increasingly accepted (~eg Hu, 2009~), we argue that the extreme sharpness of the transition region of the energy spectrum is indicative of the presence of a recent, nearby SNR. Finally, some remarks will be made about the possibility of the total intensity of CR, as distinct from just the high energy component, having relevance to the terrestrial climate and thereby to evolutionary mechanisms. This is left to last because we are less convinced by the claims for its relevance to the lightning hypothesis but it is included for its relevance to the other mechanisms. ", "conclusions": "" }, "1003/1003.1751_arXiv.txt": { "abstract": "We present the results of large-area \\co{} emission mapping of three nearby field galaxies, NGC\\,628, NGC\\,3521, and NGC\\,3627, completed at the James Clerk Maxwell Telescope as part of the Nearby Galaxies Legacy Survey. These galaxies all have moderate to strong \\co{} detections over large areas of the fields observed by the survey, showing resolved structure and dynamics in their warm/dense molecular gas disks. All three galaxies were part of the Spitzer Infrared Nearby Galaxies Survey sample, and as such have excellent published multi-wavelength ancillary data. These data sets allow us to examine the star formation properties, gas content, and dynamics of these galaxies on sub-kiloparsec scales. We find that the global gas depletion times for dense/warm molecular gas in these galaxies is consistent with other results for nearby spiral galaxies, indicating this may be independent of galaxy properties such as structures, gas compositions, and environments. Similar to the results from the THINGS \\hi{} survey, we do not see a correlation of the star formation efficiency with the gas surface density consistent with the Schmidt-Kennicutt law. Finally, we find that the star formation efficiency of the dense molecular gas traced by \\co{} is potentially flat or slightly declining as a function of molecular gas density, the \\corat{} ratio (in contrast to the correlation found in a previous study into the starburst galaxy M83), and the fraction of total gas in molecular form. ", "introduction": "\\label{sec:intro} Gas and dust in disk galaxies are intimately linked to almost all aspects of their evolution, dynamics, and appearance. For example, the interstellar medium (ISM) influences star formation and evolution, while these processes in turn have an effect on the distribution and turbulence of the ISM. One of the most important links is that between the molecular clouds, which provide the fuel for star formation, and the location and rate of star formation. However the underlying physical process that drives this relation is still poorly understood \\citep{sch08}. A recent detailed study by \\citet{ler08}, using the high resolution \\hi{} maps from THINGS \\citep{wal08}, found that many of the commonly used prescriptions do not predict the star formation well. Tracing {\\em all} components of the ISM across a wide variety of galaxies can lead to a greater understanding of the physics of the star formation process, its dependence on the properties of the host galaxy's ISM, the evolution of galaxies, and the distribution of dark matter. One of the limiting factors in understanding the physics of star formation is our ability to trace the molecular gas clouds within galaxies. Much work has been done recently to examine the molecular gas within individual and small samples of nearby galaxies using various molecular transition lines that trace different gas densities and temperatures. The most extensive observations thus far have used the millimeter \\cooz{} transition to map molecular gas, including large surveys of nearby galaxies with single-dish telescopes and interferometers \\citep[e.g.][]{you95,hel03,kun07}. This transition is most sensitive to cold, lower density molecular gas, which does not appear to trace star formation on a one-to-one basis \\citep{gre05}. Recent studies such as \\citet{ion09} have found that higher order transitions such as the \\co{} line more closely follow global star formation, nearly linearly over five orders of magnitude. The James Clerk Maxwell Telescope (JCMT) is a 15~meter single dish submillimeter telescope operated by the Joint Astronomy Centre (JAC) on Mauna Kea, Hawai'i, which after a recent major upgrade is undertaking the JCMT Legacy Survey\\footnote{http://www.jach.hawaii.edu/JCMT/surveys/} (JLS) a series of seven surveys from planetary to cosmological scales. One of these, the JCMT Nearby Galaxies Legacy Survey (NGLS), is the first large submillimeter survey of nearby galaxies at $\\sim$15\\arcsec{} spatial resolution. Taking advantage of new instrumentation on the JCMT, the survey will observe a well-selected sample of 155 nearby galaxies (within 25~Mpc) at 850~$\\mu$m and 450~$\\mu$m and in the \\co{} line. The NGLS data set will be a powerful tool for studying the physics of the dusty ISM in galaxies as well as the interplay between star formation and the ISM. With a large, well-selected sample, we will be able to search for variations in the physical properties of the ISM as a function of galaxy type, metallicity, star formation rate, mass, and environment. By having limited our sample galaxies to distances closer than 25~Mpc, we will be able to study spatial scales of 0.2-2~kpc and to search for variations inside a single galaxy as well as between galaxies. \\begin{deluxetable}{llccccccccc} \\tabletypesize{\\scriptsize} \\tablecaption{Summary of Previously Measured Galaxy Properties. \\label{tab:prop}} \\tablewidth{0pt} \\tablehead{\\colhead{Name} & \\colhead{$\\alpha$(J2000.0)} & \\colhead{$\\delta$(J2000.0)} & \\colhead{Type} & \\colhead{\\vsys{}} & \\colhead{$d$} & \\colhead{\\MB{}} & \\colhead{\\MHI{}} & \\colhead{${\\cal M}_{\\rm H_{2}}$} & \\colhead{$\\log L_{{\\rm H}\\alpha}$} & \\colhead{$i$} \\\\ & & & & \\colhead{(\\kkms{})} & \\colhead{(Mpc)} & \\colhead{(mag)} & \\colhead{($10^{8}$\\Msun{})} & \\colhead{($10^{8}$\\Msun{})} & \\colhead{($\\log$\\,erg\\,s$^{-1}$)} & \\colhead{(\\degr{})} \\\\ \\colhead{(1)} & \\colhead{(2)} & \\colhead{(3)} & \\colhead{(4)} & \\colhead{(5)} & \\colhead{(6)} & \\colhead{(7)} & \\colhead{(8)} & \\colhead{(9)} & \\colhead{(10)} & \\colhead{(11)} } \\startdata NGC\\,628 & $01\\,36\\,41.77$ & $+15\\,47\\,00.5$ & SA(s)c & 648 & 7.3 & -19.58 & 38.0 & 6.3 & 40.87 & 7.0 \\\\ NGC\\,3521 & $11\\,05\\,48.88$ & $-00\\,02\\,05.7$ & SAB(rs)bc & 778 & 10.7 & -19.85 & 80.2 & 19.5 & 41.00 & 72.7 \\\\ NGC\\,3627 & $11\\,20\\,15.03$ & $+12\\,59\\,29.6$ & SAB(s)b & 697 & 9.4 & -20.44 & 8.18 & 41.2 & 41.11 & 61.8 \\\\ \\enddata \\tablecomments{Col. (1): Galaxy name. Col. (2) and (3): J2000.0 right ascension and declination as given in RC3 \\citep{dev91}. Units of right ascension are hours, minutes, and seconds, and units of declination are degrees, arcminutes, and arcseconds. Col. (4): Galaxy morphology type from NED. Col. (5): Systemic velocity from the \\hi{} line (Heliocentric). Col. (6): Adopted distance \\citep[][respectively]{kar04,wal08,fre01}. Col. (7): Total absolute {\\em B}-band magnitude \\citep{ken08}. Col. (8): The \\hi{} mass from THINGS \\citep{wal08}. Col. (9): The H$_{2}$ mass from the NRAO 12~m \\cooz{} on-the-fly observations for BIMA SONG \\citep{hel03}. Col. (10): The H$\\alpha$ luminosity from \\citet{ken08}. Col. (11): The adopted inclination \\citep{wal08,deb08}. } \\end{deluxetable} Since November 2007 the first phase of the survey has been complete, observing 57 galaxies of the sample in the \\co{} line; 21 overlap with the Spitzer Infrared Nearby Galaxies Survey \\citep[SINGS,][]{ken03} sample, and 36 are members of the Virgo Cluster. \\co{} observations for the remaining SINGS and field sample galaxies are currently underway. Following on from \\citet[][ hereafter \\pI{}]{wil09}, which investigated four SINGS subsample galaxies that are members of the Virgo Cluster, we present the \\co{} results for three SINGS galaxies outside the cluster environment. For our initial examination of the \\co{} data in non-cluster NGLS galaxies, we have chosen to look at three well known local spiral galaxies that were completed in the early stages of the survey: NGC\\,628, NGC\\,3521, and NGC\\,3627. All three of these galaxies have extensive multi-wavelength ancillary data, most recently from SINGS and the follow-up surveys associated with it. Having been included in SINGS, all three have publicly available infrared observations (seven bands from 3.6 to 160~$\\mu$m, plus some spectroscopy). All three were also included in The \\hi{} Nearby Galaxy Survey \\citep[THINGS,][]{wal08}, the VLA \\hi{} follow up survey to SINGS. Additionally, they were all observed in the \\cooz{} line on the Berkeley-Illinois-Maryland Association (BIMA) millimeter interferometer as part of the BIMA Survey of Nearby Galaxies \\citep[BIMA SONG,][]{hel03}, and two (NGC\\,3521 and NGC\\,3627) have single dish \\cooz{} observations taken with the Nobeyama 45-m telescope by \\citet{kun07}. Comparing some of the abundant ancillary data available with our JCMT \\co{} observations provide us with a powerful tool for exploring star formation within these galaxies. Table~\\ref{tab:prop} outlines some of the previously published properties of these three galaxies. NGC\\,628 (M\\,74) is an isolated, near face-on spiral of type SA(s)c. It is known to have an extended \\hi{} disk, extending to approximately three times the optical Holmberg radius \\citep{kam92}. \\cooz{} observations of NGC\\,628 for BIMA SONG \\citep{hel03} only cover the inner region of the galaxy, but show that the molecular gas is clumpy and distributed mostly along the spiral arms. It is slightly less luminous than the other two galaxies we are examining here, and also has the lowest H$\\alpha$ and \\cooz{} luminosity of the three. We adopt the same distance estimates in this paper as in \\citet{wal08}. For NGC\\,628 the distance we use is 7.3~Mpc, derived by \\citet{kar04} using the luminosity of the brightest stars. NGC\\,3521 is a highly inclined flocculent spiral galaxy (type SAB(rs)bc) that shows a ring-like \\cooz{} structure \\citep{hel03}. It is associated with the nearby Leo Triplet (of which another of our galaxies, NGC\\,3627, is a member). It has a moderately high \\hi{} mass \\citep[$8 \\times 10^9$\\Msun{},][]{wal08}, the highest of these three galaxies, and the \\hi{} disk extends well beyond the stellar disk. The adopted distance to NGC\\,3521 is 10.7~Mpc \\citep[][ Hubble flow distance using the velocity listed on NED]{wal08}. NGC\\,3627 (M\\,66) is an active, asymmetric barred spiral galaxy (type SAB(s)b), a member of the Leo Triplet well known for its unusual kinematics that are influenced by both its bar and external interactions \\citep{zha93,reu96,reg02,che03}. It is the most luminous and most active of these three galaxies, and has the highest molecular gas content of the three as deduced from \\cooz{} observations \\citep[${\\cal M}_{\\rm H_{2}} = (4.1\\pm0.4) \\times 10^{9}$\\Msun{},][]{hel03}. In contrast, the \\hi{} mass is the lowest of the three galaxies, such that the molecular gas dominates the gas content of the galaxy. The adopted distance to NGC\\,3627 is 9.3~Mpc, derived by \\citet{fre01} using Cepheid variables. In this paper we focus on the molecular gas kinematics and detailed star formation properties of these three galaxies. \\S~\\ref{sec:data} describes our observations at the JCMT and data reduction process, and the processing of ancillary data. \\S~\\ref{sec:resmaps} presents the resulting moment maps from our \\co{} observations. \\S~\\ref{sec:resratio} compares our new \\co{} observations to existing \\hi{} and \\cooz{} data. \\S~\\ref{sec:resdyne} deals with the dynamics of the molecular gas, looking at differences in the velocity fields from molecular and atomic gas, and comparing the rotation curves we derive from our \\co{} observations with those from the \\hi{} data of THINGS \\citep{wal08}. In \\S~\\ref{sec:ressf}, we discuss several aspects of star formation within these galaxies. Finally we present our conclusions in \\S~\\ref{sec:conclusions}. ", "conclusions": "\\label{sec:conclusions} We have used the JCMT to map the \\co{} emission from three non-cluster spiral galaxies, completed as part of the ongoing Nearby Galaxies Legacy Survey. These galaxies are all included in the SINGS survey \\citep{ken03}, and as such have excellent published data sets covering their dust and ISM properties. Combined with the new NGLS \\co{} data, these data set also allow us to probe a range of galaxy properties including differences in the distribution of \\co{}, \\cooz{}, and \\hi{} emission, galaxy dynamics, and relationships between star formation rates and warm/dense molecular gas. These galaxies all have moderate to strong \\co{} detections over large areas of the fields observed by the survey, showing resolved structure and dynamics in their warm/dense molecular gas disks. We find some large differences between the \\co{} and \\hi{} velocity fields of some galaxies, in particular NGC\\,3521. These differences appear to be due to the presence of large quantities of low velocity \\hi{} gas in the THINGS data cubes. Global gas depletion times for dense/warm molecular gas in these galaxies, as well as those examined in \\pI{} and other surveys, are relatively consistent with each other (in the range 1 to 2~Gyr). This is despite very different structures, gas compositions, and environments. This is likely more an indication of the transitional nature of such gas in spirals rather than the properties of the galaxy itself. Similar to the results from the \\citet{ler08}, we do not see any correlation of the star formation efficiency with the gas surface density, with at best a slight decline in efficiency in high gas density regions traced by \\co{}. Additionally, we find that the star formation efficiency of the dense/warm molecular gas traced by \\co{} is flat or slightly declining as the \\corat{} ratio increases, in contrast to the correlation found in the study by \\citet{mur07} into the starburst galaxy M83. A similar weak trend is also seen when examining star formation efficiency versus the H$_2$ (CO\\,$J$=1-0) gas mass surface density fraction. All three of these could indicate that in regions of dense and/or warm molecular gas the efficiency is either independent of the gas density, or become slightly less efficient as the density increases." }, "1003/1003.3424_arXiv.txt": { "abstract": "The \\snr\\ is interacting with a protrusion of the cavity wall in its eastern edge (the XA region), where the X-ray emission is very bright. The complexity of the environment and the non-linear physical processes of the shock-cloud interaction make the origin of the \\xray\\ emission still not well understood. Our purpose is to understand the physical origin of the \\xray\\ emission in the XA region, addressing, in particular, the role of thermal conduction in the interaction process. We analyzed two \\xmm\\ data sets, performing image analysis and spatially resolved spectral analysis on a set of homogeneous regions. We applied a recently developed diagnostic tool to compare spectral analysis results with predictions of theoretical models, and to estimate the efficiency of thermal conduction on the X-ray emitting shocked plasma. We found that the inhomogeneous cavity wall contains both large clumps (the protrusion) and small isolated clumps with different densities. A large indentation bent over to the south is detected. The abundance of the surrounding ISM is $\\sim0.2$ times solar value. We confirmed the important role of thermal conduction in the evolution of \\xray\\ emitting plasma during shock-cloud interaction. ", "introduction": "The \\snr\\ is considered to be a proto-typical evolved supernova remnant (SNR). Located at a distance of 540\\witherr{80}{100} pc (Blair \\& Sankrit 2005), the remnant can be studied with high spatial resolution. Its surface brightness is also high, which enables us to attain high signal-to-noise ratio observations. The \\snr\\ is therefore a great laboratory to study the evolution of a SNR at different scales. The X-ray emission of the \\snr\\ is limb-brightened and clumpy, which is thought to be the result of a recent interaction with an inhomogeneous cavity wall (Levenson et al.\\ 1997, Levenson et al.\\ 1999). Several X-ray bright features in the edge of the remnant were studied, in the framework of the scenario of shock-cloud interaction (Miyata \\& Tsunemi 2001, Levenson et al.\\ 2002, Levenson \\& Graham 2005, and Patnaude et al.\\ 2002). However, the issue of the origin of the hot plasma contributing to the \\xray\\ emission in these regions was not addressed in detail yet. One of the difficulty is to take into proper account typical non-linear physical processes occurring in these systems, such as radiative losses and thermal conduction. To better understand the origin of the \\xray\\ emission, a detailed model that could be accurately compared with the observation is required. Miceli et al.\\ (2006) studied an isolated \\xray\\ knot in the Vela SNR (Vela Fi1D) by directly comparing models with observations. They built up a set of hydrodynamic simulations of shock-cloud interaction including thermal conduction and radiative cooling effects. The model setups were determined taking account of the earlier analysis of the X-ray data (Miceli et al.\\ 2005). They synthesized the X-ray emission (maps and spectra) from the models, and compared it directly with the observation for the first time. In this way, a self consistent picture of the hot X-ray emitting plasma was derived, and the importance of thermal conduction was confirmed in this case. \\begin{figure*} \\centerline{ {\\hfil\\hfil \\psfig{figure=fig1.ps,height=14cm,angle=0, clip=} \\hfil\\hfil}} \\caption{EPIC count rate maps (MOS-equivalent) in three energy bands and median photon energy map. The count rate maps were produced by weighted averaging the MOS1, MOS2, and pn images in the soft (0.3--0.5 keV, {\\sl upper-left panel}), medium (0.5--1 keV, {\\sl upper-right panel}), and hard (1--12 keV, {\\sl lower-left panel}) bands. The maps were adaptively smoothed to a signal-to-noise ratio of 25. The peak values in the three maps are: 1.7\\E{-3} ${\\rm cnt}~{\\rm s}^{-1}~{\\rm bin}^{-1}$, 1.6\\E{-3} ${\\rm cnt}~{\\rm s}^{-1}~{\\rm bin}^{-1}$, and 2.1\\E{-4} ${\\rm cnt}~{\\rm s}^{-1}~{\\rm bin}^{-1}$, respectively. The dashed rectangular region in the soft band map indicates the extension of the strip (see text). The median photon energy map ({\\sl lower-right} panel) was produced from a combination of the MOS1 and MOS2 data in 0.3--2 keV band. The image was smoothed with a gaussian distribution with $\\sigma=8.5\\arcsec$, and the minimum number of counts per pixel is 4. The regions chosen for spectra extraction are indicated in red, while the dashed green circles indicate the background regions. The bin size of each image is 4\\arcsec. The white contour marks the $3\\sigma$ confidence level of the \\xray\\ emission in 0.3--2 keV band. } \\label{f:3bandsmpe} \\end{figure*} Orlando et al.\\ (2005) and Orlando et al.\\ (2006) examined the shock-cloud interaction in different physical regimes by using a detailed hydrodynamic model, including thermal conduction and radiative losses, and synthesizing the relevant \\xray\\ emission. Based on these works, Orlando et al.\\ (2010) proposed a diagnostic tool to derive quickly the efficiency of thermal conduction from the results of the \\xray\\ data analysis, without the need of running a time consuming ad-hoc hydrodynamical model. By using such advanced method, the origin of X-rays in the \\snr\\ can be investigated to derive a better understanding of the shock-cloud interaction. In this paper we apply this method to a still not well understood region of the \\snr, an outstanding interaction region named ``XA\" by Hester \\& Cox 1986. Both bright \\xray\\ emission and bright optical emission are present in this region. The scenario of the blast wave interacting with a protrusion from the cavity wall was proposed to explain the properties of optical emission (Danforth et al.\\ 2001, Hester \\& Cox 1986). The proposed scenario is entirely different from the case of interaction with an isolated cloud in the Vela SNR studied by Miceli et al. (2005, 2006): therefore, it is necessary to see whether thermal conduction also plays an important role in this situation. We analyzed the \\xmm\\ European photon imaging camera (EPIC) observations of the XA region to examine the physical processes that account for the observed \\xray\\ emission. The observations and results are described in \\S2 and \\S3. In \\S4, we discuss the interpretation and derive the characteristic parameters of shock-cloud interaction. The conclusions are summarized in \\S5. ", "conclusions": "\\subsection{Global properties of the \\xray\\ emitting plasma} The patchy morphology of the soft band X-ray emission (see Fig.~\\ref{f:3bandsmpe}a) and the filaments observed in the optical band (Danforth et al.\\ 2001) strongly indicate that in the XA region the blast wave of the \\snr\\ is encountering a large scale inhomogeneity of the ISM. In such interacting region, the \\xray\\ emitting plasma can be typically described with two thermal components associated with two phases of ISM swept up by the shock, as shown in theoretical models (Miceli et al.\\ 2006, Orlando et al.\\ 2006, 2010). In this scenario, the two temperature components found in our spectral analysis probably indicate two thermal phases of ISM: i) the shocked dense cloud material and ii) the surrounding corona. Analogous features were also reported in the Vela SNR observed with \\rosat\\ (Bocchino et al.\\ 1999) and \\xmm\\ (Miceli et al.\\ 2005), and in several regions of the \\snr\\ (Levenson et al.\\ 2002, Katsuda et al.\\ 2008a, Miyata et al.\\ 2007, and Uchida et al.\\ 2008). \\begin{table}{} \\caption{Derived parameters for each region.\\label{tbl:dt}} \\medskip \\centering \\renewcommand{\\footnoterule}{} \\begin{minipage}{8.0cm} \\begin{tabular}{lcccc} \\hline region & f$_1$ & $n_1 ({\\rm cm}^{-3})$ & $n_2 ({\\rm cm}^{-3})$ & $t_{int}$ (kyr)\\footnote{the elapsed time of interaction, derived by $\\tau/n_2$ where $\\tau$ is the ionization time scale of the hot component.}\\\\ \\hline 1 & 0.95\\witherr{0.01}{0.02} & 6.4\\err{0.7} & 2.20\\err{0.02} & 5\\witherr{3}{4}\\\\ 2 & 0.87\\err{0.03} & 3.9\\err{0.4} & 1.54\\err{0.03} & 14\\err{8}\\\\ 3 & 0.985\\witherr{0.006}{0.01} & 11\\err{2} & 3.048\\witherr{0.006}{0.01} & 1.3\\witherr{0.6}{1}\\\\ 4 & 0.91\\err{0.02} & 4.4\\witherr{0.4}{0.5} & 1.26\\witherr{0.01}{0.02} & 7\\witherr{2}{3}\\\\ 5 & 0.90\\err{0.03} & 4.2\\witherr{0.4}{0.4} & 1.70\\err{0.03} & $>9$\\\\ 6 & 0.85\\err{0.02} & 3.1\\witherr{0.2}{0.3} & 1.13\\err{0.02} & 7\\witherr{0.8}{0.6}\\\\ 7 & 0.97\\err{0.01} & 8\\err1 & 2.687\\witherr{0.009}{0.01} & $>1$\\\\ 8 & 0.80\\witherr{0.06}{0.09} & 4.9\\witherr{0.9}{1} & 1.94\\witherr{0.05}{0.06} & $>16$\\\\ 9 & 0.76\\err{0.03} & 2.7\\witherr{0.2}{0.3} & 0.90\\err{0.03} & 7.4\\witherr{1}{0.7}\\\\ \\hline \\end{tabular} \\end{minipage} \\end{table} The extension of the line of sight (L) through the nearly spherical remnant (with a radius of $\\sim84'.5$) at the location of each region could be approximated as the length of the chord there. Then the electron density of the $i$-$th$ component can be derived considering $L_i=L\\times f_i$, where $f_i$ is the filling factor. Under the assumption of pressure equilibrium between the two components, the filling factor can be estimated as $f_1=\\left(\\frac{{\\rm EM}_2}{{\\rm EM}_1}\\frac{{\\rm T}_2^2}{{\\rm T}_1^2}+1\\right)^{-1} $ and $f_2=1-f_1$ (see Bocchino et al.\\ 1999). From the ionization time scale and the density of the hot component, we also derived the time elapsed after the interaction with the shock, $t_{int}$. The results of filling factors, densities, and $t_{int}$ (with corresponding uncertainties) are listed in Table~\\ref{tbl:dt}. Since the filling factors of the low temperature component are predominant in all the regions, the dense cloud material always fills most of the volume. The filling factors and the densities of the cold components are always higher in the knot regions than in the corresponding diffuse regions, and this confirms the inhomogeneities of the ISM and indicates that the knot regions mostly subtend the shocked cloud material. The $t_{int}$ is longer in the diffuse regions than in the corresponding knot regions, and this is consistent with the direction of propagation of the blast wave, though uncertainties are large. In regions 5, 7, and 8, the hot component seems close to ionization equilibrium, and $t_{int}$ would be maybe too long considering the \\snr\\ age. The discrepancy may be caused by mistaking the extension along the LoS. As an example, in region 8, if we assume the extension along the LoS to be equal to the size of the region in the plane of the sky, we obtain more reasonable values ($t_{int}>3000$ yr), while the densities become larger (n$_1$=27\\witherr{5}{7}, n$_2$=10.7\\err{0.3}). So the shocked cloud in regions 7 and 8 is probably a small, compact isolated knot which is much denser than the others. The case could be similar in region 5. As explained in Sect. \\ref{Spectral analysis}, the low temperature component can be associated with the shocked ISM, because of its characteristic low abundances. The densities of the \\xray\\ emitting shocked clouds (i.~e. $n_1$ in Table~\\ref{tbl:dt}) are comparable with that obtained in other cloud-shock interaction regions in the \\snr; for example, in the northeastern limb $n=$1.6--4 cm$^{-3}$ (Miyata et al.\\ 2007, the densities of the high temperature component are 1--2 cm$^{-3}$ there), in the southwestern region $n=$0.5--1 cm$^{-3}$ (Patnaude et al.\\ 2002), and in the western limb the pre-shock cloud density is 5 cm$^{-3}$ (Levenson et al.\\ 2002). The densities of FilD, a shocked cloud in the Vela SNR, were found to be 0.32--5 cm$^{-3}$ for the cloud core and 0.14--5 cm$^{-3}$ for the corona (Miceli et al.\\ 2005). The overabundances revealed in the hot components in region 4 and region 9 are consistent with an ejecta origin, and their temperatures are higher than that of the shocked ISM derived in the other regions. Both region 4 and region 9 are bright in the hard band emission. In region 4, a bright knot is visible in the hard band map (Fig.~\\ref{f:3bandsmpe}c), and this may indicate a fragment or protrusion in the ejecta. In region 9, the temperature of the cooler component is higher than that in the other regions, and this is consistent with its lower density. The region is located inside the remnant, where the ejecta is likely to be prominent ($<$~0.85~R$_{\\rm shock}$, Tsunemi et al.\\ 2007, Katsuda et al.\\ 2008b). Indeed, in the large \\xray\\ bright region to the northwest of our FoV (containing region 9) there is also the contribution of some edge-on sheet of shocked ISM (Raymond et al.\\ 1988, Blair et al.\\ 1991), so the \\xray\\ emission in this region is probably associated both with the hot ejecta and the shocked ISM. The long strip (``finger\") that we detected in the soft band (Fig.~\\ref{f:3bandsmpe}a) and in the median photon energy maps (Fig.~\\ref{f:3bandsmpe}b) is likely an extension of the very dense small protrusion observed in the optical band by Danforth et al.\\ (2001). The densities of the cold component in regions 5\\&6 (corresponding to this finger) are high, and this may also be related to the enhancement of 25 $\\mu$m mid-infrared emission (Fig.~\\ref{f:softvsIR}). However, no optical emission from the finger has been detected, thus suggesting that it does not correspond to the densest core of the XA region. The presence of the finger indicates that the protrusion in the cavity wall is much larger than what was thought before. \\subsection{The role of thermal conduction} \\label{The role of thermal conduction} \\begin{figure} \\centerline{ {\\hfil\\hfil \\psfig{figure=fig5.ps,height=6.0cm,angle=0,clip=} \\hfil\\hfil}} \\caption{The figure is the same as Fig.~$7$ in Orlando et al.\\ (2010), but we added (in black) the points with the values derived from the analysis of the EPIC observation of the XA region in the \\snr, by comparing the temperatures and the emission measures per unit area of the cold component derived for the knots and the corresponding diffuse regions. The blue stars are derived from the theoretical model without thermal conduction, the red crosses are derived from the model with thermal conduction. } \\label{f:TvsEM} \\end{figure} To examine the role of thermal conduction in the shock-cloud interaction, we applied the method proposed by Orlando et al.\\ (2010). By adopting this method, we can perform a direct and quick comparison between the observations and the results of the numerical model of shock-cloud interaction, without the need to run ad-hoc (and time-consuming) MHD simulations. In particular, we adopt a robust criterion that consists in comparing the temperature and the emission measure per unit area of the cold component derived from the knot (T$_{\\rm cold}$[knot], em$_{\\rm cold}$[knot]) to that derived from the corresponding diffuse regions (T$_{\\rm cold}$[diffuse], em$_{\\rm cold}$[diffuse]). Our results are shown in Fig.~\\ref{f:TvsEM}. Orlando et al.\\ (2010) showed that when thermal conduction is at work, we expect T$_{\\rm cold}$[diffuse]/T$_{\\rm cold}$[knot]$>1$ and em$_{\\rm cold}$[diffuse]/em$_{\\rm cold}$[knot]$<0.5$. As shown in Fig.~\\ref{f:TvsEM}, we found that thermal conduction is efficient in all the shock-cloud interaction regions. Because of the presence of ejecta in regions 3\\&4, we can not draw a conclusion there. In regions 5 and 6, that are at a relatively large distance from the border of the shell, we expect the emission to be contaminated by the contribution of multiple components along the LoS. This may be the reason of the high value em$_{\\rm cold}$[diffuse]$/$em$_{\\rm cold}$[knot] for these two regions. However, considering the error bars, all regions are consistent with a scenario where the contribution of thermal conduction is relevant. Orlando et al.\\ (2010) also showed that the slope of the $\\overline{E}$ vs $R$ plot can provide important information on the efficiency of thermal conduction and on the physical properties of the shock and the cloud. We presented these plots for our spectral regions in Fig.~\\ref{f:RvsE}. We found that the steepest descending trend of the ($\\overline{E}$) vs $R$ plots (in regions 7\\&8) is about two times flatter than that predicted by the theoretical model without thermal conduction (as shown in Fig.~$6$ in Orlando et al.\\ 2010). Such flat descending trends are instead consistent with the shock-cloud interaction with prominent thermal conduction. As shown in Orlando et al.\\ 2010, the $\\overline{E}$ vs $R$ slope is sensitive to the shock Mach number, and the flatter trends in regions 1\\&2 and 5\\&6 probably indicate higher Mach numbers. In regions 7\\&8, the steeper descending trend and the cold plateau provide convincing evidences for less efficient thermal conduction. This could be caused by slow shocks, presence of magnetic fields, or both. Considering that the magnetic field scale-length should be much larger than the distance between our regions, we do not expect large variations of the magnetic field and we conclude that a slow shock in regions 7\\&8 may be more reasonably account for the steeper trend of the $\\overline{E}$ vs $R$ plot. The ascending trend in regions 3\\&4 is not predicted in the theoretical models of shock-cloud interaction by Orlando et al. (2010). The ejecta origin of the X-ray emission in region 4 could be the reason of such ascending trend, but further investigation is needed to confirm this." }, "1003/1003.2202_arXiv.txt": { "abstract": "The attenuation of starlight by interstellar dust is investigated in a sample of low redshift, disk--dominated star--forming galaxies using photometry from {\\it GALEX} and SDSS. By considering broadband colors as a function of galaxy inclination we are able to confidently separate trends arising from increasing dust opacity from possible differences in stellar populations, since stellar populations do not correlate with inclination. We are thus able to make firm statements regarding the wavelength--dependent attenuation of starlight for disk--dominated galaxies as a function of gas--phase metallicity and stellar mass. All commonly employed dust attenuation curves (such as the Calzetti curve for starbursts, or a power-law curve) provide poor fits to the ultraviolet colors for moderately and highly inclined galaxies. This conclusion rests on the fact that the average FUV-NUV color varies little from face-on to edge-on galaxies, while other colors such as NUV$-u$ and $u-r$ vary strongly with inclination. After considering a number of model variations, we are led to speculate that the presence of the strong dust extinction feature at 2175\\AA\\, seen in the Milky Way (MW) extinction curve is responsible for the observed trends. If the 2175\\AA\\, feature is responsible, these results would constitute the first detection of the feature in the attenuation curves of galaxies at low redshift. Independent of our interpretation, these results imply that the modeling of dust attenuation in the ultraviolet is significantly more complicated than traditionally assumed. These results also imply a very weak dependence of the FUV-NUV color on total FUV attenuation, and we conclude from this that it is extremely difficult to use only the observed UV spectral slope to infer the total UV dust attenuation, as is commonly done. We propose several simple tests that might finally identify the grain population responsible for the 2175\\AA\\, feature. ", "introduction": "\\label{s:intro} Dust is everywhere. Its obscuring and emissive signatures have been observed in environs ranging from nearby dwarf irregulars to giant ellipticals, within our own and nearby galaxies, in intergalactic space, and even in the highest redshift quasars \\citep[e.g.,][]{Trumpler30, Disney89, Knapp89, Goudfrooij94, Zaritsky94b, Xu95, Wang96, Elbaz99, Xilouris99, Calzetti01, Motta02, Wang08, Menard09}. Our understanding of the composition and spatial distribution of dust in a variety of environments has increased markedly in the past decade, thanks largely to the launches of the {\\it Spitzer Space Telescope} and the {\\it Galaxy Evolution Explorer (GALEX)}. For example, the fraction of a galaxy's gas mass that is locked up in dust is now well-characterized as a function of galaxy mass, star formation rate (SFR), and metallicity \\citep[e.g.,][]{Draine07, daCunha08, daCunha10}. The radial distribution of gas and dust has also been extensively studied \\citep{Munoz-Mateos09}. The importance and prevalence of polycyclic aromatic hydrocarbons (PAHs) is now widely acknowledged thanks to moderate resolution spectroscopy in the mid--IR \\citep[e.g.,][]{Leger84, Smith07, ODowd09}. Despite these advances, basic questions remain. It is still not known, for example, if the majority of dust forms in AGB outflows, in supernovae ejecta, or condenses directly out of the ISM \\citep{Galliano08, Draine09}. Computation of the dust extinction curve from first principles relies on a number of uncertain inputs, including the size, shape, and composition of the grains, and the optical constants of the constituent materials \\citep[e.g.,][]{Weingartner01, Draine03, Gordon03, Zubko04}. A further uncertainty with broad implications for our understanding of galaxies is the relation between the stellar and dust content of galaxies; for example, how the grain size distribution varies within and between galaxies. Due to gas inflow and outflow, the stellar and gas--phase metallicities are not simply related, further compounding the problem. Without a solid theoretical understanding of the complex relation between stellar and dust content, these two components need to be modeled flexibly and independently. Despite this widely acknowledged fact, in practice a particular dust obscuration model is often simply assumed in order to then infer the underlying stellar population. Such an assumption can introduce significant biases, as we will discuss below. The net effect of dust obscuration within a galaxy is the result of the combination of the underlying dust extinction curve, geometry, and radiative transfer. Here, geometry is taken to mean not only the large--scale distribution of dust with respect to the stars, but also the local geometry, which includes the clumpiness of the interstellar medium (ISM). Geometrical effects can make the relation between the underlying extinction curve and the resulting attenuation curve (which measures the net loss of photons within a galaxy) arbitrarily complicated \\citep[e.g.,][]{Natta84, Witt92, Calzetti94, Witt96, Varosi99, Witt00, Granato00, Charlot00, Pierini04, Tuffs04, Panuzzo07}. For example, a clumpy ISM, arising for example from turbulence, will result in an attenuation curve that is greyer than the extinction curve. When the optical depth approaches unity, the effects of geometry become especially important. In such cases, the light measured in the blue/ultraviolet may not trace the same regions of the galaxy as the red/near--IR light. Consider a disk galaxy viewed edge--on that is optically thick in the UV and optically thin in the near-IR. The ultraviolet light emitted from the far side of the galaxy will be heavily extinguished, and so the ultraviolet light collected by an observer will be preferentially sampling the near side of the galaxy. In contrast, the near--IR will more faithfully trace the entire stellar population because the galaxy is relatively transparent in this wavelength range. These effects are rarely considered when modeling the SEDs of galaxies. If we knew the underlying, unobscured stellar population, the ratio between the intrinsic and observed light would then provide a direct measure of the attenuation. Unfortunately, it is not possible to know unambiguously the underlying stellar population. Less direct methods are therefore required. For example, the attenuation curve can be estimated in a sample of galaxies with similar stellar populations and variable amounts of dust. In such cases, ratios between more and less heavily attenuated spectra can provide estimates of the net attenuation curve in the sample. Such a technique was utilized by \\citet{Calzetti94} in order to probe the average attenuation curve of starburst galaxies. The Balmer decrement was used as an estimator for the amount of dust attenuation, and, with the assumption that the galaxies in their sample were of similar metallicity and SFR, the average attenuation properties were estimated. Similar techniques for probing the wavelength--dependent attenuation have recently been applied by \\citet{Johnson07b, Johnson07a} to photometry of a large sample of low redshift galaxies, \\citet{Conroy10b} who analyzed the restframe UV photometry of star--forming galaxies at $z\\sim1$, and \\citet{Noll09} who analyzed spectra of star--forming galaxies at $z\\sim2$. In the present work we follow in this vein by considering the colors of disk--dominated galaxies as a function of their inclination. Since inclination will correlate with dust attenuation but not with stellar populations, comparing less to more inclined systems allows us to isolate the effects of increasing dust opacity on the observed properties of galaxies \\citep[e.g.,][]{Giovanelli94, Giovanelli95, Masters03, Driver07, Unterborn08, Maller09, Masters10, Yip10}. Another constraint on the attenuation curve comes from the relation between the ratio of total infrared to UV luminosity and UV spectral slope ($L_{\\rm TIR}/L_{\\rm UV}-\\beta$, or the `IRX--$\\beta$' relation). Star-forming galaxies with redder UV spectra tend to have higher IRX. It is widely believed that this relation is primarily a sequence in dust attenuation \\citep[e.g.,][]{Kong04}. Various recipes have thus been proposed to use the IRX--$\\beta$ relation to estimate the total UV attenuation based on the observed UV spectral slope \\citep[e.g.,][]{Buat05, Burgarella05, Cortese08}. This relation is of particular importance to the study of high--redshift galaxies, where only restframe UV and optical photons can be readily collected for large samples of galaxies \\citep[although significant samples at $z\\sim1-2$ with restframe IR data are growing rapidly, e.g.,][]{Reddy06a, Salim09, Reddy10}. For most studies of high redshift galaxies, the observed UV slope, $\\beta$, is used in conjunction with a locally estimated IRX--$\\beta$ relation to estimate the dust opacity. This approach is essential, for example, to interpret recent observations of galaxies at $z\\approx6-8$ \\citep[e.g.,][]{Bouwens09}. The IRX--$\\beta$ relation is different for starbursts \\citep{Meurer99} and normal star--forming galaxies \\citep{Dale07, Boissier07}, and in addition depends somewhat on the star formation history \\citep{Kong04, Johnson07a, Munoz-Mateos09}, 60$\\mu m$ luminosity \\citep{Takeuchi10}, bolometric luminosity \\citep{Reddy06a}, and the star--dust geometry \\citep{Panuzzo07}. These dependencies result in substantial scatter in the IRX--$\\beta$ plane, which, in conjunction with the fact that the IRX--$\\beta$ is nearly vertical over much of the relevant parameter space (i.e., IRX varies considerably over a relatively narrow range in $\\beta$), calls its utility into question \\citep{Bell02b}. The most prominent feature of the MW extinction curve is the strong, broad dust feature at $2175$\\AA, the `UV bump' \\citep{Stecher65}. This feature is also seen, albeit more weakly, along most sightlines in the LMC and along one of the five sightlines probed in the SMC \\citep{Gordon03}. It is also seen along sightlines in M31 \\citep{Bianchi96}. There are now confident detections of this feature in the extinction curves of more distant galaxies as well, as probed by background gamma ray bursts and quasars, and gravitational lenses \\citep{Ardis09, Motta02, Wang04, Mediavilla05}, Puzzlingly, there are many examples of galaxies that do not show this feature in their extinction curve \\citep[e.g.,][]{York06, Stratta07}. The carrier of this extinction feature is not known, although owing to its strength it must be due to some abundant material, such as carbon. Measurements of the grain albedo suggest that the feature is not due to scattering \\citep[see data compilations in][]{Witt00, Draine03}. Several dust models associate this feature with PAH absorption \\citep[e.g.,][]{Weingartner01} although there are other possibilities \\citep[e.g.,]{Draine93}. Understanding the prevalence of the UV bump in the attenuation curves of other galaxies is essential for broadband photometric studies of galaxies. For example, at $z\\sim0$ the UV bump falls into the {\\it GALEX} NUV-band, at $z\\sim1$ it falls within the $B-$band, and at $z\\sim2$ it redshifts into the $R-$band. If present in the attenuation curves of galaxies, the UV bump would therefore result in substantially more, and likely more uncertain, attenuation in particular filters at particular redshifts. Based on spectra from the {\\it International Ultraviolet Explorer}, Calzetti et al. found an average attenuation curve for starburst galaxies that lacked a UV bump. The absence of a UV bump in the starburst attenuation curve led \\citet{Witt00} to suggest that in such galaxies the underlying dust extinction curve lacks a UV bump. A detailed analysis of the ultraviolet through infrared photometry of M51 revealed little evidence for a UV bump within individual HII regions \\citep{Calzetti05}. \\citet{Conroy10b} also found no evidence for a UV bump as strong as that seen in the MW in star--forming galaxies at $z\\sim1$. In stark contrast, \\citet{Noll09} presented strong evidence for a UV bump in stacked restframe UV spectra of $z\\sim2$ star--forming galaxies. Finally, \\citet{Capak10} presented tantalizing evidence for a strong UV bump in at least one $z\\sim7$ galaxy. Thus, while the UV bump appears to be a ubiquitous feature of the MW and LMC extinction curves, there is little evidence for this feature in the net attenuation curves of other galaxies. However, it is important to recognize that a systematic investigation of the UV attenuation in $z\\sim0$ `normal' star-forming galaxies is currently lacking. While the effects of geometry and scattering can diminish the strength of the UV bump with respect to the underlying extinction curve, it is a generic prediction of radiative transfer calculations that if the UV bump is present in the extinction curve, its presence will be detectable in the resulting attenuation curve \\citep{Witt00, Panuzzo07}. Given the observations of the UV bump in the extinction curves of other galaxies, the attenuation curves of at least some galaxies should show a UV bump, even if the strength of the bump is weak. This expectation serves as motivation for the present study. There is currently no clear picture linking together these various observations. One possibility is that the radiation field modulates the strength of the UV bump, as suggested by \\citet{Gordon03}. The typical radiation field in the SMC and the local starburst galaxy sample is much harsher than the MW, potentially explaining why a UV bump is not seen in such systems. It would be hard to explain the results of Noll et al. in this context, however, given that their sample is dominated by high SFR systems. \\citet{Granato00} explained the result from Calzetti et al. as being due to the fact that the UV energy production in starbursts is dominated by young stars that are heavily embedded within molecular clouds. In their model, attenuation in starbursts is therefore governed by the wavelength-dependent fraction of UV photons produced by young stars (whose light is heavily extinguished by their birth cloud). These authors predicted that the attenuation suffered by `normal' star-forming galaxies would show evidence of a UV bump because in this case significant UV energy is provided by intermediate age stars that have left their birth clouds. Such stars will suffer attenuation primary from the diffuse dust, where significant 2175\\AA\\, absorption may be expected. It is not immediately obvious whether or not this explanation can accommodate the results from Noll et al. \\citet{Draine07} has recently demonstrated a deficiency of PAH emission in the infrared in galaxies with low gas--phase metallicities \\citep[see also][]{Engelbracht05, Smith07}. Metallicity may therefore play an important role. In the present work we investigate the wavelength--dependent attenuation by dust for a sample of disk--dominated galaxies. Our sample is carefully selected to be homogeneous and complete. We then consider the ultraviolet, optical, and near--infrared colors as a function of inclination. Since inclination will only correlate with dust attenuation and not physical parameters such as star formation nor metallicity, the inclination--dependent colors will provide a robust and sensitive probe of the wavelength--dependent attenuation in disk--dominated systems. This technique is therefore similar in spirit to that of \\citet{Calzetti94, Calzetti00}, although for a sample of `normal' star-forming galaxies, and utilizing photometry rather than spectroscopy. Averaging the flux of many galaxies considerably simplifies the modeling of the underlying stellar population because the average star formation history (SFH) of many normal galaxies must be smooth. By normal here we mean galaxies not chosen to have special SFHs, such as starburst or post-starburst galaxies. We will focus especially on the attenuation properties in the NUV band, and will therefore be able to make strong statements regarding the presence of the UV bump in the average attenuation curve of normal star-forming galaxies. Where necessary, a flat $\\Lambda$CDM cosmology with $(\\Omega_m, \\Omega_\\Lambda)=(0.30,0.70)$ is adopted, along with a Hubble constant of $H_0 = 100h$ km s$^{-1}$ Mpc$^{-1}$. All magnitudes are in the $AB$ system \\citep{Oke83}. A \\citet{Chabrier03} initial stellar mass function (IMF) is adopted when quoting stellar masses. ", "conclusions": "\\label{s:disc} \\subsection{An emerging physical picture, and implications} In the previous section we demonstrated that the inclination--dependent UV colors of a mass--selected sample of disk--dominated galaxies are best explained by a MW attenuation curve with $R_V=2.0$ that includes a prominent UV bump at 2175\\AA. The strength of this bump is approximately 80\\% as strong as the average UV bump observed in the MW. This result is summarized in Figure \\ref{fig:best}. In this figure we show our favored attenuation curve and compare this to the standard MW curve ($R_V=3.1$). We also show the actual constraints on the attenuation provided by the FUV, NUV, $u$, and $g$ filters. It is clear from this figure that while our results are consistent with a strong UV bump, we cannot claim that our results {\\it require} a UV bump, since we rely solely on broadband photometry over a narrow redshift range. However, inspection of Figure \\ref{fig:optspec} strongly suggests that the UV bump is responsible, as we know of no other model variation capable of reproducing the depression in the average SED of edge-on galaxies at $\\lambda\\approx2200$\\AA. Our results imply that previous work attempting to measure the SFR in low redshift galaxies has underestimated the amount of attenuation in the near-UV. From inspection of Figure \\ref{fig:best} we can see that typical attenuation in the NUV band is underestimated by approximately the value of the $V-$band optical depth, $\\tau_V$, when comparing standard attenuation curves (power-law, starburst, and average MW) to our favored model. For moderately dusty galaxies the underestimation may therefore be substantial. The bias is such that SFRs based primarily on the near-UV will be lower than the intrinsic SFR. Additional work will be required to understand in detail how the effects of a strong UV bump propagate into the derived physical properties of galaxies. From this figure we can also understand qualitatively why our disk--dominated sample favors an attenuation curve with $R_V=2.0$, as opposed to the canonical value of $R_V=3.1$ preferred in the MW.\\footnote{Of course, as emphasized throughout, attenuation and extinction are conceptually different, and so there is no {\\it a priori} reason for the attenuation curve that best describes our sample to be similar to the average extinction curve measured for the MW.} Decreasing $R_V$ from 3.1 to 2.0 results in substantially more attenuation between the $u$ and NUV filters, and therefore a much redder NUV$-u$ color. The relative attenuation between FUV and NUV, and also between $u$ and $g$ is does not change substantially between the $R_V=3.1$ and $R_V=2.0$ curves, and so these colors change by much less than NUV$-u$. Therefore, it is principally the constraint from the NUV$-u$ color that drives the requirement for $R_V\\approx2.0$. If our interpretation of the observed trends is correct, this would constitute the first detection of the UV bump in the attenuation properties of galaxies at low redshift. The only other detection of the UV bump in the attenuation curves of galaxies was reported by \\citet{Noll09}, who considered restframe UV spectra of star-forming galaxies at $z\\sim2$. In \\citet{Conroy10b} the observed $B-R$ colors of star-forming galaxies at $0.7 5\\times 10^{-6}$. In this paper, we report on our analysis of archival X-ray data for \\etfo\\ collected by {\\it ASCA, Chandra, XMM} and \\Suzaku\\ during the past 17 years, including two observations that were made fortuitously very closely following glitches. We have looked for correlations between the AXP's flux variability and glitch epochs. In section \\ref{sec:obs} we describe the observations and data reduction process. In section \\ref{sec:spec} we describe our spectral analysis and AXP flux extraction method. Our results and conclusions are discussed in section \\ref{sec:sum}. ", "conclusions": "\\label{sec:sum} The goal of this study was to see whether the prolific glitching AXP 1E~1841$-$045 shows phase-averaged flux variability, in spite of showing no evidence for pulsed flux variability. Also, we wished to determine whether any variability is correlated with its glitches as has been seen in AXP \\ttfn\\ in its 2002 major outburst, in 1E~1048.1$-$5937, and also reported for \\seven. As is clear from Panel a of Figure \\ref{fig:flux}, in the 4--10 keV band, the neutron star's flux did not vary by more than $\\sim$30\\% in 13 years. Interestingly the largest variations we find are in the multiple pre-1999 {\\it ASCA} observations: in those seven observations, a fit to a constant flux results in a reduced $\\chi^2$ of 3.7 for 6 degrees of freedom, which has a probability of occurring by chance of $\\sim$0.001. However, during this time, there were certainly no large glitches ($\\Delta \\nu/\\nu < 5\\times 10^{-6}$) \\citep{gvd99}. The \\asca fluxes all appear to be higher than the fluxes measured from other observations. If this difference is accurate, then the AXP's flux must have dropped around 2000, and we might in principle expect some change in the pulsed flux as monitored by \\xte\\ at that epoch. However, the pulsed flux of the pulsar was constant around that time as seen in Panel b of Figure \\ref{fig:flux}. Therefore, we suspect that the relative increase in the flux in the {\\it ASCA} observations compared with those of the other observatories may be due to instrumental calibration issues. The fluxes measured from the last four observations taken by \\chandra, \\xmm and \\Suzaku can be fitted with a constant flux model (reduced $\\chi^2 = 1.9$ for 3 degrees of freedom, corresponding to a probability of having occurred by chance of 0.125). Thus we conclude that the phase-averaged 4--10~keV fluxes of the last four observations were consistent with being constant, and we put an upper limit of 11\\% on long-term variability in this energy band. Importantly, the two \\xmm observations were taken only 88 and 90 days after the first glitch, and the \\Suzaku observation was taken only 27 days after the third glitch. By contrast, the 4--10 keV flux of \\ttfn\\, was 50\\% higher than in quiescence 21 days after its 2002 glitch \\citep{zkd+08}, that of \\tfe was a factor of 6 higher $\\ge$38 days after its 2007 glitches \\citep{tgd+08}, and was 50--70\\% higher for \\seven\\, $\\sim$53 days after its first 2005 glitch as inferred from \\citet{gri+07}\\footnote{This is calculated based on the reported 1--10 keV fluxes of \\seven\\, from its 2003 \\xmm observation and 2005 \\swift observation. We assumed a power-law to blackbody flux ratio of 3 for the 1-10 keV fluxes, and then calculated the 4--10 keV fluxes using webPIMMS (http://heasarc.gsfc.nasa.gov/Tools/w3pimms.html)}. Therefore, we conclude that unlike in \\ttfn, 1E~1048.1$-$5937 and possibly \\seven, there is no evidence for glitch-correlated flux changes in AXP \\etfo. One caveat of our study is that we were limited to the harder part of the neutron star's emission spectrum. The flux from the blackbody component was not well constrained. Therefore, we cannot rule out changes in the neutron star's thermal radiation, only changes in the power-law component in the 4--10~keV band, which we note constitutes $\\sim$0.25 of the stellar flux (BB+POW) in the 1--10 keV band. Thus our results argue against the hypothesis that there exists a generic correlation between the X-ray flux variability and glitch epochs in AXPs, further supporting the argument that glitches in AXPs can be either radiatively loud or radiatively silent \\citep{dkg08}. There is of course precedent for radiatively silent glitches in neutron stars, in that no rotation-powered pulsar glitch has ever been reported to be accompanied with any radiative change, although rapid X-ray follow-up has only been accomplished in one case \\citep{hgh01}. Any physical model of magnetar glitches will have to explain the simultaneous existence of both types. This is true of even a single source, as there is evidence that AXP \\ttfn has both, given that its most recent glitch showed no pulsed flux change \\citep{dkg08b}. Recently \\citet{es10} have argued that AXP glitches are triggered by energy releases at depths below $\\sim$100~m in the crust, with angular momentum vortex unpinning being due to global mechanical motion triggered by the energy release, not by heat as has been proposed in the context of rotation-powered pulsars glitches \\citep{le96,lc02}. If mechanical triggering occurs, then radiatively silent glitches of the same amplitude as radiatively loud glitches are possible since less energy is required to trigger a glitch than to cause a substantial X-ray brightening. If so, then \\citet{es10} predict that all AXP glitches should occur simultaneously with or before the observed X-ray brightening; this can be tested with continuous (daily or better) X-ray monitoring observations. Moreover, although mechanical unpinning of vortices by activity in the lower crust does result in heat release, the latter could take as much as several years to reach the surface. This could help explain the long-term X-ray variability trends that have been reported in some AXPs \\citep[e.g.][]{dkg07}." }, "1003/1003.5119.txt": { "abstract": "Precision cosmology with Type Ia supernovae (SNe Ia) makes use of the fact that SN Ia luminosities depend on their light-curve shapes and colours. Using Supernova Legacy Survey (SNLS) and other data, we show that there is an additional dependence on the global characteristics of their host galaxies: events of the same light-curve shape and colour are, on average, 0.08\\,mag ($\\simeq4.0\\sigma$) brighter in massive host galaxies (presumably metal-rich) and galaxies with low specific star-formation rates (sSFR). These trends do not depend on any assumed cosmological model, and are independent of the SN light-curve width: both fast and slow-declining events show the same trends. SNe Ia in galaxies with a low sSFR also have a smaller slope (``$\\beta$'') between their luminosities and colours with $\\sim$2.7$\\sigma$ significance, and a smaller scatter on SN Ia Hubble diagrams (at 95\\% confidence), though the significance of these effects is dependent on the reddest SNe. SN Ia colours are similar between low-mass and high-mass hosts, leading us to interpret their luminosity differences as an intrinsic property of the SNe and not of some external factor such as dust. If the host stellar mass is interpreted as a metallicity indicator using galaxy mass--metallicity relations, the luminosity trends are in qualitative agreement with theoretical predictions. We show that the average stellar mass, and therefore the average metallicity, of our SN Ia host galaxies decreases with redshift. The SN Ia luminosity differences consequently introduce a systematic error in cosmological analyses, comparable to the current statistical uncertainties on parameters such as $w$, the equation of state of dark energy. We show that the use of two SN Ia absolute magnitudes, one for events in high-mass (metal-rich) galaxies, and one for events in low-mass (metal-poor) galaxies, adequately corrects for the differences. Cosmological fits incorporating these terms give a significant reduction in $\\chi^2$ (3.8--4.5$\\sigma$); linear corrections based on host parameters do not perform as well. We conclude that all future SN Ia cosmological analyses should use a correction of this (or similar) form to control demographic shifts in the underlying galaxy population. ", "introduction": "As calibrateable standard candles, Type Ia supernovae (SNe Ia) provide a direct route to understanding the nature of the dark energy that drives the accelerated expansion of the Universe. Yet, the relationships that allow the calibration of their peak luminosities, and hence permit their cosmological use, remain purely empirical. Relations between the width of the SN Ia light curve and peak luminosity \\citep{1993ApJ...413L.105P} and between the SN Ia optical colours and luminosity \\citep[e.g.][]{1996ApJ...473...88R,1998A&A...331..815T} reduce the scatter in their peak magnitudes to $\\sim$0.15\\,mag \\citep{2007ApJ...659..122J,2007A&A...466...11G,2008ApJ...681..482C}. As the available SN Ia samples increase in both size and quality, and the dark energy constraints they provide become correspondingly more statistically precise, it is increasingly important that the validity of these calibrating relationships is robustly examined. The observed properties of SNe Ia are known to correlate with the physical parameters defining their host galaxy stellar populations. SNe Ia are more than an order of magnitude more common (per unit stellar mass) in actively star-forming or morphologically late-type galaxies than in passive or elliptical systems \\citep{2005A&A...433..807M,2006ApJ...648..868S}. SNe Ia in elliptical or passively evolving systems are also intrinsically fainter, with narrower, faster (or lower ``stretch''), light curves \\citep{1995AJ....109....1H,1996AJ....112.2398H,1999AJ....117..707R,2000AJ....120.1479H,2006ApJ...648..868S}. Though this effect is corrected for by the light-curve shape correction, the amount of star formation activity in the universe increases with redshift, and these differences lead to an observed ``demographic shift'' in mean SN Ia properties. A greater fraction of intrinsically luminous, wider light-curve events in the distant universe are seen compared to that observed locally \\citep{2007ApJ...667L..37H}. These photometric differences are also partially reflected in SN Ia spectra, with SNe Ia in spiral galaxies showing weaker intermediate mass element line strengths than those in elliptical galaxies \\citep{2008A&A...477..717B,2009A&A...507...85B}, and a corresponding evolution in the mean SN Ia spectrum with redshift \\citep{2009ApJ...693L..76S}. There are suggestions that these effects may be the result of multiple astrophysical channels capable of producing SN Ia explosions \\citep[e.g.][]{2005ApJ...629L..85S,2006MNRAS.370..773M}. In particular, delay-time distributions with distinct ``prompt'' and ``delayed'' components, or with a wide range of delay-times, match most observational datasets well \\citep{2006MNRAS.370..773M,2006ApJ...648..868S,2008ApJ...683L..25P,2008PASJ...60.1327T}, though the minimum age for the prompt systems remains controversial \\citep{2008A&A...492..631A,2009ApJ...707...74R} with some evidence that ``prompt'' SNe Ia occur more frequently in metal-poor systems \\citep{2009ApJ...704..687C}. The use of SNe Ia as precision cosmological probes therefore depends on establishing that the demographic shifts, or existence of multiple channels to a SN Ia, do not impact on the light-curve-width/colour/luminosity relationships. If these relations show environmental dependence, then the task of calibrating SNe Ia for cosmology becomes substantially more challenging \\citep[e.g.][]{2008ApJ...684L..13S,2009arXiv0912.0929K}. A second complication arises from the (poorly understood) colour corrections applied to SN Ia luminosities. Redder SNe Ia appear fainter than their bluer counterparts, but the slope of the relationship between SN Ia colour and magnitude is inconsistent with the ratio of total-to-selective absorption appropriate for the diffuse interstellar medium of the Milky Way ($R_V=3.1$). Multiple studies of different SN Ia samples indicate that the effective $R_V$ inferred from normal SNe Ia is smaller than 3.1 \\citep[e.g.,][]{1998A&A...331..815T,2006A&A...447...31A,2006AJ....131.1639K}, and the \\textit{assumption} of $R_V=3.1$, even after light curve shape corrections, leads to serious systematic error problems such as a spurious ``Hubble Bubble'' \\citep{2007ApJ...659..122J,2007ApJ...664L..13C}. The reason for this low effective $R_V$ is not well understood. Although uncorrected intrinsic variations in the SN Ia population could play a role \\citep[e.g.][]{2009Natur.460..869K}, some dust extinction must also affect the SN Ia luminosities and colours, and this may vary by environment. Furthermore, the exact value of $R_V$ obtained is sensitive to method used to determine it, with lower $R_V$ obtained when fitting linear relations between SN Ia luminosities, colours, and light curve widths \\citep{2010AJ....139..120F}, perhaps due to intrinsic variation in SN Ia colours that correlates with luminosity but not light-curve width. Current knowledge of SNe Ia is not sufficient to separate and correct for both intrinsic colour--luminosity and dust-induced colour--luminosity effects in cosmological SN Ia samples. Examining how SN Ia luminosities vary with environment after light curve shape and colour corrections can place constraints on the degree of these possible variations. Early studies showed little evidence that corrected SN Ia luminosities varied with host galaxy morphologies \\citep[e.g.,][]{1999ApJ...517..565P,1999AJ....117..707R,2003MNRAS.340.1057S,2003AJ....126.2608W,2005ApJ...634..210G}, though these tests used relatively small samples of events ($\\la50$), in some cases from the first-generation of SN Ia cosmological samples before dense multi-colour light curves were routinely obtained. More recent analyses, using larger, well-observed samples, have shown tentative evidence for variation. \\citet{2009ApJ...700.1097H} found $\\simeq2\\sigma$ evidence that SNe Ia in morphologically E/S0 galaxies are brighter than those in later-type spirals after light-curve shape and colour corrections. Extending beyond simple host galaxy morphologies to more physically motivated variables gives further tantalising suggestions of variation. \\citet{2008ApJ...685..752G} found evidence for a correlation between Hubble diagram residual and host galaxy stellar metallicity in a sample of 17 local SNe Ia located in E/S0 galaxies, in the sense that fainter SNe Ia after correction were found in metal poor systems (note this is the reverse of the originally published trend due to an error in the original analysis; P. Garnavich, private communication). \\citet{2009ApJ...691..661H} used 55 SNe Ia from the first year of the SNLS and showed no significant correlation between Hubble residual and host galaxy metallicity, albeit using host gas-phase metallicities inferred from average galaxy stellar-mass--metallicity relations, a less direct measure of metallicity. \\citet{2009arXiv0912.0929K} have shown a relation between host galaxy stellar mass and Hubble residual, in the sense that more massive systems host brighter SNe Ia after light curve shape and colour corrections. Under the assumption that more massive galaxies are metal rich, this trend is consistent with the revised \\citet{2008ApJ...685..752G} result. In this paper, we use a sample of 282 high redshift SNe Ia discovered and photometrically monitored by the Canada-France-Hawaii Telescope (CFHT) as part of the Supernova Legacy Survey (SNLS), and which form the SNLS ``three-year'' sample (SNLS3). Using deep optical imaging of their host galaxies taken over the duration of the survey, we place constraints on their recent star-formation activity, stellar masses (and hence inferred metallicity), and compare to the photometric properties of the SNe Ia that they host. In particular, we search for evidence that the corrected SN Ia luminosities correlate with these host properties, indicating possible systematic errors in the light curve fitting framework that underpins their cosmological use. We compare with the properties of a sample of lower-redshift SNe Ia taken from the literature. A plan of the paper follows. In $\\S$~\\ref{sec:data} we introduce the SN Ia sample and the data available on their host galaxies. $\\S$~\\ref{sec:sn-properties-as} investigates how the SN Ia light curve widths and colours of these SNe Ia varies according to their host galaxy properties, and in $\\S$~\\ref{sec:lum-depend-trends} we compare their corrected luminosities to the host properties. We discuss the results, including the cosmological implications, in $\\S$~\\ref{sec:discussion}, and conclude in $\\S$~\\ref{sec:conclusions}. Throughout, where relevant we assume a flat $\\Lambda$CDM cosmological model with $\\omatter=0.256$ (the reason for this non-standard choice is explained in $\\S$~\\ref{sec:lum-depend-trends}) and $\\mathrm{H}_0$=70\\,km\\,s$^{-1}$\\,Mpc$^{-1}$ assumed in all quoted absolute magnitudes. All magnitudes are given on the \\citet{1992AJ....104..340L} photometric system as described in \\citet{2009A&A...506..999R}. ", "conclusions": "" }, "1003/1003.3518_arXiv.txt": { "abstract": "Hoyle \\& Folwler showed that there could be Radiation Pressure Supported Stars (RPSS) even in Newtonian gravity. Much later, Mitra found that one could also conceive of their General Relativistic (GR) version, ``Relativistic Radiation Pressure Supported Stars'' (RRPSSs). While RPSSs have $z\\ll 1$, RRPSSs have $z \\gg 1$, where $z$ is the surface gravitational redshift. Here we elaborate on the formation of RRPSSs during continued gravitational collapse by recalling that a contracting massive star must start trapping radiation as it would enter its {\\em photon sphere}. It is found that, irrespective of the details of the contraction process, the trapped radiation flux should attain the corresponding Eddington value at sufficiently large $z\\gg 1$. This means that continued GR collapse may generate an intermediate RRPSS with $z\\gg 1$ before a true BH state with $z=\\infty$ is formed asymptotically. An exciting consequence of this is that the stellar mass black hole candidates, at present epoch, should be hot balls of quark gluon plasma, as has been discussed by Royzen in a recent article entitled ``{\\it QCD against black holes?}''. ", "introduction": "For any self-gravitating object, the definition of ``compactness'' may be given in terms of the surface gravitational redshift(Weinberg 1972): \\begin{equation} z = \\left(1 - R_s/ R\\right)^{-1/2} -1 \\end{equation} where $R_s = 2 G M/c^2$ is the Schwarzschild radius of the object having a gravitational mass $M$ and radius $R$. Here, $G$ is the Newtonian gravitational constant and $c$ is the speed of light. For the Sun, one has $z \\approx 2 \\times 10^{-6}$ while for a typical neutron star $z \\sim 0.15$. Another important parameter for a self-gravitating object is $x= p_r/p_g$, where $p_r$ is the pure radiation pressure and $p_g$ is the kinetic pressure (Mitra 2009a,b). While the central region of Sun has $x \\approx 0.006$, obviously, a {\\em cold} object at temperature $T=0$ has $x=0$. It is known that strictly static and {\\em cold} stars have an upper mass limit both in Newtonian and Einstein gravity. In the Newtonian case, this is obtained by the marriage of Newtonian gravity and Special Relativity, and is known as ``Chandrasekhar Mass'', $M_{ch}$. On the other hand, in Einstein gravity, such an upper mass limit is known as ``Oppenheimer -Volkoff Mass'' (Weinberg 1972). The precise values of such limits depend on the equation of state (EOS) and other details. However, for a pure {\\em cold}, Helium dwarf, $M_{ch} \\sim 1.4 M_\\odot$, where $M_\\odot$ is the solar mass. On the other hand, for a pure free neutron Fermi-Dirac fluid, $M_{ov} \\sim 0.8 M_\\odot$. But if such upper limits were the full story, we would not have had stars of masses as large as $\\sim 100 M_\\odot$. Further, there are interstellar gas clouds of mass probably as large as $\\sim 10^6 M_\\odot$. The reason behind the existence of non-singular cosmic objects with such higher masses is that they are not supported by {\\em cold} quantum pressure alone. On the other hand, they are supported not only by $p_g$ but partly by $p_r$ too (Mitra 2009a,b). Even when one would work at a purely Newtonian level ($z \\ll 1$), the probable increase in the value of $x$ would support the higher self-gravity of a star $M \\gg 1 M_\\odot$. And this was probably first realized by Hoyle \\& Fowler (1963)and Fowler (1966) who conceived of the Radiation Pressure Supported Stars (RPSS) having $x\\gg 1$. It turns out that, given the self-imposed restriction, $z\\ll 1$, one would require $M > 7200 M_\\odot$ in order to have $x >1$ (Weinberg 1972). With the increase of $x$ and attendant self-gravity, the star tends to be more compact, i.e., $z$ would increase Mitra(2009a,b). For instance even though ``supermassive stars'' are Newtonian (i.e., $z \\ll 1$), they possess $z$ much larger than the solar value of $z_\\odot \\approx 2 \\times 10^{-6}$. Accordingly, it is possible to conceive of a Newtonian RPSS with $z \\sim 0.1$, i.e., one which is almost as compact as a neutron star. A related important concept here is the ``Eddington Luminosity''. In order that a self-gravitating and self -luminous object can remain quasi-static, its luminosity at a given radius $r=r$ must be less than a critical value (Hoyle \\& Fowler 1963, Weinberg 1972): \\begin{equation} L_{ed}(r) = {4 \\pi G M(r) c\\over \\kappa} \\end{equation} where $M(r)$ is the gravitational mass within a given surface, and $\\kappa$ is the appropriate opacity. In case, the star will be composed of pure ionized hydrogen, one will have the lowest value of opacity called Thomson opacity: $ \\kappa = \\sigma_T/ m_p \\approx 0.4$ cm$^2$ g$^{-1}$ where, $\\sigma_T$ is the Thomson crosssection and $m_p$ is the proton rest mass. If one would have the intrinsic luminosity, $L > L_{ed}$, the star would be disrupted by radiation pressure. The Newtonian supermassive stars, conceived by Hoyle \\& Fowler necessarily radiate at the above mentioned Newtonian Eddington rate. As one would tread into Einstein gravity, however, one should be able to conceive of situations with $z \\gg 1$. The definition of Eddington luminosity, in such a case, gets modified (Mitra 1998): \\begin{equation} L_{ed} = {4 \\pi G M c\\over \\kappa } (1+z) \\end{equation} Accordingly, in Einstein gravity, there should be Relativistic RPSSs (RRPSS) which would {\\em locally} radiate at the above mentioned enhanced Edddington rate. In the following, we discuss the basic physics behind the formation of such RRPSSs. \\section {More ~on ~Eddington ~Luminosity} It is known that the concept of Eddington luminosity is relevant not only for the structure of stars but for the accretion process around the stars too. In fact, the accretion luminosity of the star also must be limited by $L_{ed}$ for steady spherical accretion. To appreciate this, let us recall here the basic physics behind the concept of ``Eddington Luminosity'' by considering the fluid to be a fully ionized hydrogen plasma: The average {\\em attractive} gravitational force on one atom is $ F_g = -G M m_p/R^2$. If this would be the only force acting on the plasma, then both the intrinsic luminosity and the accretion luminosity could be infinite. In reality, however, $L$ is finite and has a maximal value, $L_{ed}$, because the ionized H-atom is also subject to the {\\em repulsive} force due to accretion luminosity of the central object: \\begin{equation} F_{rad} = \\sigma_T ~ q/c \\end{equation} where $q = L/ 4 \\pi R^2 $ is the radial energy/heat flux. Thus the effective, \\begin{equation} F_g \\to F_g + F_{rad} = - {G M\\over R^2} [(1-\\alpha) m_p] \\end{equation} where $\\alpha = L/L_{ed}$ and $L_{ed} = 4 \\pi G M c/ \\kappa$. Therefore, as if, the radially outward heat transport in a spherically symmetric isotropic isolated body reduces the Effective Gravity (EG) by a factor of $(1 -\\alpha)$. This may be also seen as a reduction of the inertial mass (IM) by the same factor. A value of $\\alpha >1$, i.e., $L > L_{ed}$ would thus mean EG to be negative, where either the star would be disrupted despite self-gravity or instead of accretion there could be radiation driven winds. The latter indeed happens in the atmosphere of very massive stars with very large $L$. In the context of collapse, if the collapse generated luminosity would approach its Eddington value, it becomes clear then that the collapse process would tend to be stalled. Essentially, $L_{ed}$ corresponds to a critical comoving {\\em outward heat flux} of \\begin{equation} q_{ed} = {L_{ed}\\over 4 \\pi R^2} = {G M\\over \\kappa R^2} (1+z) \\end{equation} Irrespective of the specific mode of reduction of the EG due to heat flow, the very notion of an ``Eddington Luminosity'', both in Newtonian gravity and in GR, implies that {\\em the attainment of} $L=L_{ed}$ {\\em would stop inflow/collapse}. In fact, in a very important study on GR collapse, Herrera \\& Santos(2004) have indeed shown that outward heat flow reduces IM by a factor $ (1 -\\alpha)$. Therefore, in principle, the GR collapse process can certainly {\\em slow down and get stalled} if the collapse generated luminosity would approach its maximal value ($\\alpha \\to 1$)! Further numerical as well as analytical studies of radiative GR collapse have confirmed the above mentioned analytical result (Herrera \\& Santos 2004; Herrera, Prisco \\& Barreto 2006; Herrera, Prisco, \\& Ospino, 2006; Herrera, Prisco, Fuenmayor, \\& Troconis 2009). If such papers (Herrera \\& Santos 2004; Herrera, Prisco \\& Barreto 2006) would be used for weak gravity, one would clearly identify the $\\alpha$ occurring in them as none other than the $\\alpha$ appearing in Eq.(5). While such studies considered enhancement of radiation flux due to matter -radiation interaction and are somewhat non-generic, we would consider here a generic effect: {\\em the unabated enhancement of the gravitationally trapped radiation beyond the photon sphere and consequent attainment of Eddington luminosity}. \\section {Self-Gravitational ~Trapping ~of ~Radiation} General Relativity (GR) predicts that even the trajectories of quanta emitted by a star itself do bend away from the direction of normal to the direction of the tangent of the surface of the body because of the effect of the gravitational field of the star. However, as long as $z < \\sqrt{3} -1$, the emitted quanta nevertheless manage to evade entrapment and move away to infinity. But if the body would be so compact as to lie within its ``photon sphere'', i.e., $R < (3/2) R_s$ or $z > \\sqrt{3} -1$, then only the radiation emitted within a cone defined by a semi-angle $\\theta_c$ (Harrison 2000): \\begin{equation} \\sin\\theta_c = {\\sqrt{27}\\over 2} (1 - R_s/R)^{1/2} (R_s/R) \\end{equation} will be able to escape. Radiation emitted in the rest of the hemisphere would eventually return within the compact object. In the presence of an external radiation, i.e., in the absence of a strict exterior vacuum, even an apparently static supermassive star is not strictly static. This is so because, the spacetime in such a case is described by radiative Vaidya solution (Vaidya 1951) rather than the exact vacuum Schwarzschild solution. In this sense, while the exterior spacetime of a {\\em cold} White Dwarf or a Neutron Star having a fluid at zero tempetature is described by vacuum Schwarzschild metric, the exterior spacetime of any radiative object including the Sun or a supermassive star, {\\em in a strict sense} is described by the Vaidya metric. Accordingly, in a strict GR sense, stars are in quasistatic equilibrium and always evolving. For the Vaidya metric, the value of $z$ could be arbitrary high, and that is the reason a collapsing and radiating object is supposed to attain the black hole (BH) stage having $z=\\infty$. As the collapse/contraction proceeds, both $M$ and $R$ decrease so that $z$ increases and eventually, the collapsing body must approach the exact BH state with $z=\\infty$. The effect of radiation trapping during these {\\sl intermediate} high-$z$ states has {\\em never} been considered, though previously Kembhavi \\& Vishveshwara(1980) observed that: ``If neutrinos are trapped, they will not be able to transport energy to the outside, and this can have serious consequences on the thermal evolution of the star. These considerations might become especially interesting in the case of a collapsing phase which leads to the formation of a compact, dense object.'' At high $z$, $R\\approx R_s$ and from Eq.(7), one can see that, $\\sin \\theta_c \\to \\theta_c \\approx (\\sqrt{27}/2) (1 +z)^{-1}$. Hence the solid angle of escaping radiation is \\begin{equation} \\Omega_c \\approx \\pi \\theta_c^2 \\approx {27 \\pi\\over 4} (1+z)^{-2} \\end{equation} The chance of escape of radiation therefore decreases as $\\Omega_c/2 \\pi \\approx (27/8) (1+z)^{-2}$. This means that if without trapping $10^{10}$ neutrinos/photons would escape a particular spot on the surface, {\\sl with} gravitational trapping, {\\sl only} $1$ out the $10^{10}$ quanta would escape for $z=10^5$. Consequently, as the collapse generates internal heat/radiation, the energy density of trapped radiation $\\rho_r$ and associated outward heat flux {\\em within} the body would increase as \\begin{equation} q_{trap} \\sim R^{-3} (1+z)^2 \\end{equation} The distantly observed luminosity/flux would be lesser by a factor of $(1+z)^2$ than what is indicated by the foregoing Eq. Such a reduction would take into account the fact that, locally, trapped quanta are moving in almost closed orbits. But as far as local flux is concerned, to avoid double counting, one must not introduce any additional factor of $(1+z)$ in Eq.(9). Using Eqs.(6) and(9), we see that, in this regime of $z\\gg 1$, \\begin{equation} \\alpha= {q_{trap}\\over q_{ed}} \\sim {(1 +z)\\over R M}\\end{equation} Initially, of course, $\\alpha \\ll 1$. But during the collapse, both $R$ and $M$ would decrease and Eq.(10) would show that, as $z\\to \\infty$, $\\alpha$ would increase dramatically, $\\alpha \\to 1$, at a sufficiently high $z$. At this stage the collapse would degenerate into a secular quasistatic contraction by the {\\em very definition} of $L_{ed}$. As if, a leaking and contracting balloon {\\it stops contraction as its self gravity fixes the leakage} by forcing the molecules to move in (almost) closed circular orbits. Also, simultaneously, the attendant heat and pressure become large enough to resist further contraction. In a very strict sense, however, the body would still be contracting on extremely long time-scales! This is so because as long as an horizon is not formed, i.e., $z < \\infty$, the body would radiate and $M$ would continue to decrease. Consequently the metric would remain {\\em non-static} and, in response, $R$ too, would decrease. It is this infinitesimal decrease in the value of $R$ and attendant much higher secular increase in the value of $z$ and $q_{trap}$ which would generate just enough energy (at the expense of $Mc^2$) to maintain the GR Eddington luminosity seen by a distant observer: \\begin{equation} L^\\infty_{ed} = { 4 \\pi R^2 q_{ed} \\over (1+z)^2} ={4 \\pi G M c\\over \\kappa (1+z)} \\end{equation} Since $L^\\infty = -c^2 dM/du$, the observed time scale associated with this phase is \\begin{equation} u = { Mc^2 \\over -c^2 dM/du} = {M c^2\\over L^\\infty_{ed}} = { \\kappa c (1+z)\\over 4 \\pi G} \\end{equation} Obviously, $u \\to \\infty$ {\\em irrespective of the value} of $\\kappa$ as the BH stage ($z=\\infty$) would be arrived. Thus the Eddington-limited contracting phase actually becomes eternal. Since for photons, $\\kappa_\\gamma \\approx 0.4$ cm$^2$/g, but, for neutrinos $k_\\nu$ is smaller by an extremely large factor of $\\sim 10^{14-18}$, we will have $u_\\nu \\ll u_\\gamma$. Consequently, initial transition to the RRPSS phase may be dominated by huge $\\nu$-emission with a time scale $u_\\nu$. But as far as eventual secular RRPSS phase is concerned, it should be governed by photon time scale $u_\\gamma$ because it is much easier to maintain a $L_{ed}$ caused by photons than by neutrinos: \\begin{equation} L^\\infty_{ed, \\gamma} = 1.3 \\left(M\\over 1 M_\\odot\\ \\right ) 10^{38} (1+z)^{-1} ~{\\rm erg/s} \\end{equation} Somewhat similar thing happens for the formation of a hot neutron star from a proto neutron star: initial time scale of $\\sim 10$s is dictated by huge $\\nu$-emission, while the hot NS cools for thousands of years by photon emission (Glendenning 2000). For this era of quasi-stability by trapped photons, by using Eq.(13) into Eq.(12), it follows that the observed time scale of an RRPSS at a given $z\\gg 1$ is given by \\begin{equation} u \\approx 1.5 \\times 10^{16} (1+z) ~ {\\rm s} \\approx 4\\times 10^8 (1+z) ~{\\rm yr} \\end{equation} For $z\\gg 1$, the local energy density is almost entirely due to radiation and pairs (Mitra 2006a) so that $\\rho \\approx aT^4/3$ where $a$ is the radiation constant and $T$ is the mean local temperature. Further since $M = (4\\pi/3) \\rho R^3$ and $R\\approx 3 (M/M_\\odot)$ Km for $z\\gg 1$ (see Eq.[1]), we obtain \\begin{equation} T = \\left({3c^4\\over 8\\pi a G}\\right)^{1/4} R_s^{-1/2} \\approx 600 \\left({M\\over M_\\odot}\\right)^{-1/2}~ {\\rm MeV} \\end{equation} Therefore, a RRPSS is an ultrarelativistic fireball of radiation and pairs interspersed with baryons much like the plasma in the very early universe (unless $M$ is too high). In particular, Eq.(15) shows that $T \\sim 200$ MeV for a $10 M_\\odot$ RRPSS. Hence, the stellar mass RRPSSs could be in a Quark Gluon Plasma (QGP) phase. As of now, it is believed that a bulk QGP phase existed only in the very early universe. But now we arrive, through a simple and straight forward analysis, at the exciting possibility that a {\\em bulk and ever lasting QGP phase } may be existing within galaxies. \\section {Analytical \\& Numerical Support for This Scenario} The physical effect described here cannot be obtained by any {\\em exact} analytical solution of Einstein equations simply because the {\\em only exact} solution of GR gravitational collapse is the Oppenheimer - Snyder one (Oppenheimer \\& Snyder 1939): If one would consider the collapse of a homogeneous dust with $p=0$, {\\em and yet assume finite initial density}, $\\rho >0$, the fluid would appear to collapse to a singularity in a flash, $ \\tau \\propto \\rho^{-1/2}$, with no question of slowing down or bounce or oscillation. However, physically, a {\\em strict} $p=0$ EOS should correspond to a fluid mass of $M=0$ (Ivanov 2002, Mitra 2009a) and thus, in a strict sense, a $p=0$ collapse should be eternal: $\\tau =\\infty$. Note, when pressure gradient forces are included, even adiabatic GR collapse {\\em admits bounce and oscillatory behavior} where the fluid need not always plunge inside its Schwarzschild radiu (Nariai 1967; Taub 1968; Bondi 1969). Further Mansouri (1977) showed that a uniform density sphere {\\em cannot undergo any adiabatic collapse at all} if an equation of state would be assumed. Therefore, the effect described here here can be inferred by appropriate numerical studies of radiative physical gravitational collapse or by {\\em generic} physical studies, as carried out here. Indeed realistic gravitational collapse must involve not only pressure gradient but also dissipative processes and radiation emission (Mitra 2006a,b,c; Herrera \\& Santos 2004; Herrera, Prisco \\& Barreto 2006, Herrera, Prisco, \\& Ospino, 2006; Herrera, Prisco, Fuenmayor, \\& Troconis 2009). It is clear that dissipative processes might not only slow down but even stall the collapse by generating a quasistatic state (Herrera \\& Santos 2004; Herrera, Prisco \\& Barreto 2006, Herrera, Prisco, \\& Ospino, 2006; Herrera, Prisco, Fuenmayor, \\& Troconis 2009). And suppose, one is considering the collapse with a certain initial mass $M=M_i$ and the RRPSS state is (first) formed at $M=M_*$ and $z=z_*$. In order to find the values of $M_*$ and $z_*$, one must study the problem numerically by devising a scheme to incorporate the effect of gravitational radiation trapping and matter-radiation interaction. Obviously, the precise value of opacity $\\kappa$ would be enormously different from what has been considered here. And same would be true for all other relevant physical parameters. But at this juncture, we are not claiming to make any such detail study. Indeed, we are interested only in a generic but physically valid picture. Implicitly, we are considering here a modest range of $M_i$ to relate it with associated numerical/analytical works (Herrera \\& Santos 2004; Herrera, Prisco \\& Barreto 2006, Herrera, Prisco, \\& Ospino, 2006; Herrera, Prisco, Fuenmayor, \\& Troconis 2009). On the other hand, for extremely large values of $M$, there would be no $\\nu$-generation and one would be concerned with solely $\\gamma$ and $e^+, e^-$ processes. This generic/qualitative picture becomes strengthened by the fact Goswami \\& Joshi (2005) considered the possibility that trapped surface formation may be avoided because of loss of mass by emission of radiation. Also, there is an {\\em exact} solution of GR collapse which shows that the repulsive effects of heat flow may prevent the formation of an event horizon (Banerjee, Chatterjee \\& Dadhich 2002). Further, Fayos \\& Torres (2008) have shown that in view of the emission of radiation, GR continued collapse may turn out to be singularity free where the {\\it entire mass-energy of the star may be radiated out}. And recently, it has been found that, indeed, for radiative spherical collapse, trapped surfaces are not formed at a finite value of $M$ (Mitra 2009b). Thus, {\\em there are considerable supports from both numerical and analytical studies of Einstein equations}, about the basic feasibility of the picture presented here. ", "conclusions": "The generic, reason why continued radiative GR collapse should (first) result in a radiation supported quasistatic state or Eternally Collapsing Object(ECO), may be better appreciated by recalling Harrison: ``When the contracting body reaches a radius 1.5 times the Schwarzschild radius, all rays emitted tangential to the surface are curved into circular orbits. This is the radius of the {\\em photon sphere}. On further contraction, the emitted rays become more strongly deflected and many now fall back to the surface. Only the rays emitted within an exit cone can escape and this escape cone narrows as contraction continues. When the body reaches the Schwarzschild radius, the exit cone closes completely and no light rays escape. Redshift and deflection conspire to ensure that no radiation escapes from a black hole.'' Just like the case of an event horizon formation is independent of the density of the fluid, {\\em radiation trapping, the basic mechanism} for the scenario presented here, too is independent of density. It simply depends on the fact for continued collapse, the fluid necessarily plunges within the {\\em photon sphere} and the effect of radiation trapping must increase dramatically with increasing, $z$. Essentially we have pursued this {\\em generic} picture by recalling that there is another associated {\\em generic} concept - which is ``Eddington Luminosity''. We showed that the rays falling back inside the contracting object should form an outward flux which would locally grow as $\\sim (1+z)^2$ whereas the definition of Eddington luminosity, in the increased gravity, would increase as $\\sim (1+z)$. And since the former grows more rapidly by a factor of $(1+z)/RM$, it must {\\em catch up} with the latter at some high $z=z_*\\gg 1$ as the journey to the BH stage involves march towards $z\\to \\infty$ and $M=0$ (Mitra 2009a,b). True, there would be many non-generic effects like matter-radiation interaction, asphericity, rotation, magnetic field etc which would affect this collapse process. But all such effects would resist the free collapse scenario and hence they would reduce the value of $z_*$. Thus this generic picture cannot be negated on the plea of complexities of the GR collapse problem. In any case, publication of this paper may prompt researchers to incorporate this effect of gravitational trapping in numerical studies of radiative collapse. While the OS paper is the only {\\em truly exact} study of continued gravitational collapse, it is also the most unrealistic one because it assumes $p=p_r=q=0$ even when the fluid would attain infinite density! And indeed Taub (1968) and Mansouri (1977) showed that a uniform density sphere cannot collapse at all if pressure would be incorporated! Consequently, in the past few years, several alternatives to BH Candidates have been proposed (Mazur \\& Mottola 2004; Chapline 2005). Unlike the OS case, such treatments are inexact, and further they assume that the collapsing matter with positive pressure suddendly udergoes some unspecified phase transition to acquire {\\em negative} pressure. Alternately, they assume that ``Dark Energy'' starts to play dominant role at local level by virtue of mysterios quantum processes. In contrast, though, our teatment too is inexact, we simply considered the natural GR effect that a contracting object must start trapping its own radiation once it is within its {\\em photon sphere}. Note, one defines Event Horizon (EH) as a surface from which ``nothing not even light can escape''. But the surface with $z_c = \\sqrt{3} -1$ is the {\\em precursor} of an EH, because outward movement of everything gets strongly inhibited for $z > z_c$. Further, unless angular momentum is lost, everything may tend to move in closed circular orbits for $z >z_c$. However, there must be equal number of counter rotating orbits in order to conserve angular momentum. In such a case, all stresses must be tangential as radial stresses should almost vanish. In other words, the situation here may approach the idealized case of an ``Einstein Cluster'' (Einstein 1939). And it follows that, when stresses are completely tangential, {\\em there could indeed be static configurations with} $2M/R \\to 1$ or $z \\to \\infty$ Florides (1974). Thus, if one would wish, one might view a RRPSS/ECO as a (quasi) static configuration with $z \\gg 1$. In such a case too, there is no upper limit on the value of $M$! There is no denying the fact that we indeed observe compact objects, as massive as $10^{10}$ $M_\\odot$. It is also certain that compact objects of even $3-4 M_\\odot$s cannot be {\\em cold} neutron stars. But it is fundamentally impossible to prove that these objects are true BHs simply because, by definition, {\\em BH event horizons having $z=\\infty$ cannot be directly detected} (Abramowicz, Kluzniak \\& Lasota 2002). And as mentioned in the introduction, for sufficiently {\\em hot} compact objects there is no upper mass limit even when stresses would be considered to be isotropic. As we found, for stellar mass cases, the quasistatic RRPSS fluid could be in a QGP state with $T\\sim 100$ MeV. Interestingly, in the context of stellar mass BH formation, a recent study entitled {\\it QCD Against Black Holes} concluded that QCD phase transition would ensure that the collapsing object {\\em becomes a QGP fluid rather than a true BH} (Royzen 2009). Note, there are many {\\em observational evidences} (Robertson \\& Leiter 2002; Robertson \\& Leiter, 2003; Robertson \\& Leiter 2004; Schild, Leiter, \\& Robertson 2006; Schild, Leiter \\& Robertson 2008) for the scenario described in this paper which show that the compact objects in some X-ray binaries or quasars could be RRPSSs with $z\\gg 1$ and strong intrinsic magnetic moments rather than exact BHs with $z=\\infty$ and no {\\em intrinsic} magnetic moment. Also, by definition, it is impossible to claim that such compact objects are true BHs with exact EHs (Abramowicz, Kluzniak \\& Lasota 2002). Unlike Mazur \\& Mottola (2004) and Chapline (2005), the scenario considered here {\\em in no way denies that, mathematically}, the final state of continued gravitational collapse is a BH. But, if one would ignore the phenomenon of trapping of own radiation due self-gravity as the body plunges into its photon sphere and also ignore the effect of pressure gradient or radiative transport, surely, one would find text book type prompt formation of BHs. However, the very fact that photon sphere is the precursor of an eventual EH, our scenario actually {\\em fills the missing gap between a photon sphere and a true EH} during continued collapse. Accordingly, this paper may not be rejected on the assumption that it is in conflict with the basic mathematical notion of a BH. As the {\\em quasistatic} hot RRPSSs would become more and more compact, they would asymptotically approach the ultimate state of spherical gravitational collapse; i.e., a BH with $z=\\infty$ \\& $M=0$ (Mitra 2009a,b). Newtonian and Post Newtonian radiation pressure supported stars, during their evolution, may be unstable to radial oscillation (Chandrasekhar 1965). And there is a preliminary numerical computation which suggests that Newtonian supermassive stars may indeed collapse to form RRPSSs rather than true BHs (Cuesta, Salim \\& Santos 2005). For Newtonian supermassive stars, even though, $p_r\\gg p_g$, $p \\ll \\rho c^2$; in contrast the RRPSSs have $p \\approx (1/3) \\rho c^2$. Thus any study of Newtonian or Post Newtonian systems is not relevant for RRPSSs. Further, recently, it has been shown that, the Active Gravitational Mass Density of a quasistatic system is $\\rho_g = \\rho - 3 p/c^2$ (Mitra 2010). Consequently, $\\rho_g \\ll \\rho$ for RRPSSs and this prevents sudden rapid gravitational contraction. However, all RRPSSSs are assymptotically contracting towards the true BH state. To conclude, for continued collapse, we elaborated here on a most natural physical mechanism by which one may have GR version of Radiation Pressure Supported stars first conceived by Hoyle \\& Fowler way back in 1963." }, "1003/1003.3204_arXiv.txt": { "abstract": "In a landscape of compactifications with different numbers of macroscopic dimensions, it is possible that our universe has nucleated from a vacuum where some of our four large dimensions were compact while other, now compact, directions were macroscopic. From our perspective, this shapeshifting can be perceived as an anisotropic background spacetime. As an example, we present a model where our universe emerged from a tunneling event which involves the decompactification of two dimensions compactified on the two-sphere. In this case, our universe is of the Kantowski-Sachs type and therefore homogeneous and anisotropic. We study the deviations from statistical isotropy of the Cosmic Microwave Background induced by the anisotropic curvature, with particular attention to the anomalies. The model predicts a quadrupolar power asymmetry with the same sign and acoustic oscillations as found by WMAP. The amplitude of the effect is however too small given the current estimated bound on anisotropic curvature derived from the quadrupole. ", "introduction": "If inflation lasted sufficiently long, all curvature scales imprinted on our universe by pre-inflationary physics are pushed to undetectable distances beyond our current horizon. If, on the other hand, inflation ended soon after the required amount of accelerated expansion to allow a later epoch of rich structure formation in our local universe, this cosmic amnesia may have been only partial. The largest cosmological scales observable today would then potentially show traces of the initial conditions for our inflating universe, including curvature \\cite{2006JHEP...03..039F, \\openinflation}, anisotropies \\cite{GCP,PPU1,PPU2}, or nontrivial topology \\cite{LachiezeRey:1995kj}. This is the setting in which the scenario we propose can have observational consequences. While it has often been stated that fine-tuning the amount of inflation to this extent is unnatural, the landscape paradigm combined with the difficulty to find long-lasting inflationary solutions in string theory provide sufficiently strong counterarguments to take this possibility seriously \\cite{DeSimone:2009dq}. Like the collision of bubbles in scenarios of false vacuum inflation (see \\cite{Aguirre:2009ug} for a status report), it offers a remote chance to probe the landscape of string theory by observations. What are plausible initial conditions for our local inflationary patch in the landscape? Compactification in string theory is often treated kinematically, as part of the construction of the effective four dimensional theory. With four macroscopic plus a number of microscopic compact dimensions fixed once and for all, the transition between metastable vacua can be described by the spontaneous nucleation of bubbles with open homogeneous and isotropic spatial sections \\cite{CDL,Aguirre:2005nt}. The observable consequences of false vacuum bubble nucleation followed by a brief period of slow-roll inflation have been studied extensively in the context of open inflation \\cite{\\openinflation} and are well understood by now. From the dynamical perspective, however, compactification becomes a problem of string cosmology, for the compactified dimensions can spontaneously open up and become large \\cite{Giddings}. This substantially widens the parameter space for initial conditions of our local universe, since there are now alternative channels to populate the landscape which can be viewed as transitions between vacua with differing numbers of macroscopic dimensions. This was first explicitly spelled out by Carroll, Johnson, and Randall \\cite{CJR} in the context of dynamical compactification from a higher dimensional spacetime to our effectively four-dimensional one. The opposite process, dynamical decompactification, was studied in \\cite{BlancoPillado:2009di}, and in \\cite{GiddingsMyers} in a different context. Although these previous studies were all concerned with the (de-)compactification of higher dimensions, there is no reason to exclude that the three macroscopic dimensions of our present universe may themselves be the result of such a process \\cite{BrandyVafa}. This is the starting point of our work\\footnote{Shortly before completion of our work, two articles (\\cite{GHR} and \\cite{BPS}) were posted which have a significant amount of overlap with ours. We will comment on the similarities and differences at various places in the main text.}. It is intuitively clear that decompactification allows the existence of a preferred direction in the sky if only one or two directions are compact, and therefore gives rise to anisotropic cosmologies. We specifically consider the case where two of our macroscopic dimensions are compact. Before inflation, these dimensions were microscopic, leaving one macroscopic direction which may still play a preferred role for cosmological observations today if inflation was short. As a concrete example, we present a four dimensional model with two dimensions compactified on a sphere by the flux of an Abelian gauge field and a cosmological constant. The solutions of the Euclidean Einstein-Maxwell equations with spherical symmetry are well known. They describe pair creation of charged black holes in de Sitter space \\cite{BoussoHawking1,BoussoHawking2,MannRoss}. We note that the causal patch beyond the cosmological horizon in the Lorentzian spacetime is an anisotropic inflating Kantowski-Sachs (KS) universe. This initial state can be placed in the broader context of the landscape in different ways. It can be viewed as an intermediate phase in a progressive opening up of initially compact dimensions as described in \\cite{GiddingsMyers}, or as the temporarily final step in a sequence of transitions starting from a higher dimensional geometry that triggered the decompactification of two previously compact dimensions. We comment briefly on these scenarios in Sec.\\ \\ref{ref:motivations} but we will leave the discussion of further implications for future work. The remainder of this article is structured as follows. In Sec.\\ \\ref{sec:model}, the model is introduced and its connection to black hole pair production is discussed. Sec.\\ \\ref{sec:perturbations} presents a preliminary exploration of Cosmic Microwave Background (CMB) signatures in our model, which are due to the anisotropy of KS spacetime. Owing to the technical complexity of a full analysis, we only consider the perturbations of a test scalar field and outline the qualitative modifications of CMB temperature anisotropies. With these results we try to assess some of the observed CMB anomalies which seem to indicate a violation of statistical isotropy in our universe. We conclude and discuss some directions for further research along these lines in Sec.\\ \\ref{sec:conclusions}. ", "conclusions": "\\label{sec:conclusions} In the landscape of string theory, vacua with differing numbers of macroscopic dimensions and topologies are connected by tunneling events. In this paradigm, our universe has been produced as the last step in a potentially long chain of events, by tunneling from a parent vacuum with a greater, equal, or smaller number of macroscopic dimensions. Like \\cite{GHR, BPS} which appeared when our work was near completion, we consider the latter case. Our approach differs in the specific model under examination and in its embedding into a cosmological scenario, which we dubbed the ``shapeshifting universe''. Specifically, we take the example of a spacetime where two of our current macroscopic dimensions are compactified on a two-sphere by a magnetic flux. The existence of a long-lived parent vacuum with two macroscopic dimensions, while all others are stably compactified, is not a necessary requirement for this scenario. All we need is a precursor with two of our four large dimensions compactified while other, currently compact, directions may well have been macroscopic. This allows a direct connection of anisotropic cosmologies produced by decompactification with the vast possibilities of transdimensional vacuum transitions described in \\cite{CJR, BSV} and, furthermore, admits histories with greatly varying effective values of $\\Lambda$. In one particular possible scenario, our parent vacuum was a $\\mathrm{dS}_D \\times \\mathrm{S}_2$ spacetime with higher $\\Lambda_{\\rm eff}$ and different compact directions than our universe, connected to ours by a combined spontaneous compactification of $D-2$ dimensions which are now small and decompactification of the $\\mathrm{S}_2$. We have not explicitly constructed the instanton describing this shapeshifting event but we know of no reasons why it shouldn't exist. It is also plausible that models can be constructed where inflation is triggered by the transition, as described in \\cite{CJR}. This line of investigation raises a number of new and interesting questions. In the context of the cosmological measure problem, the contribution of regions with different numbers of macroscopic dimensions and transitions in between is still largely unexplored (see \\cite{CJR,BSV,BPS} for some ideas). We hope to return to this question in future work. From the phenomenological point of view, on the other hand, decompactification offers a concrete framework for parameter studies of anisotropic cosmological models. Our ability to detect statistical anisotropy hinges on the assumption of a sufficiently short period of inflation, which again may find some justification in the string theory landscape. Given the observational constraints, a small but nonzero window for anisotropic curvature remains (see, e.g., \\cite{GHR}). Precisely how strong the constraints already are, and how much can be gained by combining future large-scale galaxy surveys and CMB data, can only be answered by a more detailed investigation of the cosmological signatures of models like the one proposed here. It is also interesting to note that having embedded the anisotropic universe model in a more complete scenario of quantum tunneling, where the KS spacetime corresponds to only a wedge of a large spacetime, solves the ambiguity in the ground state (compare with the discussions of \\cite{GCP,PPU2}). Our preliminary analysis of the deviations from statistical isotropy of the CMB shows the following trends. Some of the off-diagonal correlations $\\langle a_{lm} a_{l'm'}^* \\rangle$, which vanish in the isotropic case, are now non-zero. Because of the invariance under point reflection of the background, the correlations with $l$ and $l'$ of different parities vanish. The nonzero correlations receive contributions from both the inflationary phase and the recent expansion of the universe. Both scale as $\\Omega_{\\rm curv}$ in the observations. The central results are the table \\ref{tab:numerics} and the equations (\\ref{parabola}) and (\\ref{sum rule}). We showed that as a direct consequence of this, and in contrast to Bianchi III models, the Kantowski-Sachs models predict a quadrupolar modulation in the CMB with all the features observed in the data, namely the sign of the bipolar coefficients $A_{ll'}^{LM}$, their acoustic oscillations, and a particular relation between them. To obtain the observed amplitude requires, however, an anisotropic curvature at a level of $\\Omega_{\\rm curv} \\simeq 10^{-2}$, which is incompatible with the bound derived from the value of the quadrupole, namely $\\Omega_{\\rm curv} \\lesssim 10^{-4}$. We briefly argued that this bound could be mitigated. The price to pay is unfortunately an increased complexity of the model, which we leave for future work. We conclude that the model in its present, most simple, form is unable to account for any of the anomalies. As an additional line for future investigation we mention the generation of primordial magnetic fields. Since KS spacetimes are not conformally flat, we indeed expect that long wavelength magnetic fields are produced from vacuum fluctuations during inflation, without recourse to noncanonical couplings of masses. This generation could perhaps be sufficient to explain the primordial magnetic fields necessary to seed galactic dynamos." }, "1003/1003.3026_arXiv.txt": { "abstract": "This work summarises some of the attempts to explain the phenomenon of dark energy as an effective description of complex gravitational physics and the proper interpretation of observations. Cosmological backreaction has been shown to be relevant for observational (precision) cosmology, nevertheless no convincing explanation of dark energy by means of backreaction has been given so far. ", "introduction": "The accelerated expansion of the Universe inferred from the observation of supernovae (SN) Ia \\cite{sn1a} poses one of the biggest puzzles to modern cosmology. While the introduction of a simple term --- a cosmological constant --- provides a phenomenological explanation, it leads to several conceptual problems at the same time. Our minimal model, the inflationary $\\Lambda$ cold dark matter model, is very successful in describing a large number of cosmological observations like the anisotropies of cosmic microwave background (CMB) radiation, the large scale distribution of galaxies, or the cosmic expansion. Thus it is also called concordance model. Surprisingly, $95\\%$ of the cosmic energy density turn out to be in the form of dark components. Despite the success of the model describing the evolution of the Universe, it obviously fails to explain its content --- at least unless we uncover the nature of the dark components. Especially the nature of the phenomenon of ``dark energy'' is completely unclear. There are many routes that have to be explored: observational issues, astrophysics issues, testing models of a new component of dark energy, testing modifications to Einstein's equations --- to name just a few popular directions. Here I focus on the probably most conservative approach to the problem, assuming that there are no unknown observational and astrophysical issues. Accelerated expansion of the Universe seems to start when structures of the size of the Hubble scale at matter-radiation equality (the only scale in hierarchical structure formation) have grown mildly non-linear. This suggests that dark energy and structure formation might be linked in some way. It is also well known that the time evolution of cosmic observables and their volume averages do not commute. One might speculate that a mixture of several effects of non-linear structure formation and the fact that we actually observe (light-cone) averages of local quantities could modify the homogeneous and isotropic ``background'' space-time. One often calls that the idea of cosmological backreaction. In order to assess the status of these ideas and to motivate their study, I first revisit the key assumptions of the concordance model and stress that some pieces of information come entirely from local data sets ($z<1$), while some others come from high redshifts at $z > 1000$. In section 3 the model-independent evidence for cosmic acceleration is discussed, before we briefly turn to some theoretical issues related to the concordance model and a sketch of various solutions proposed so far. For the rest of this work I focus on the issue of cosmological backreaction, stressing some of the important issues of cosmic structure formation in section 5 and discussing the importance of averaging in section 5. Finally, I summarise some of the open problems. ", "conclusions": "Probably the most important open problem is to learn how to treat light-cone averages in realistic ab initio calculations of the Universe. My expectation is that the $100$ Mpc scale is most important for the effects of cosmic backreaction and that we have to design tests to decide the issue by means of observations. The proposed study of the variance of $\\delta_H$ is a first idea in that direction. The preliminary conclusion is that cosmic averaging is important, but not necessarily an explanation of dark energy. However, the hardest challenge for cosmological backreaction is to explain why cosmic backreaction would mimic $\\Lambda$CDM. At the time being, the cosmological constant wins the beauty contest." }, "1003/1003.3356_arXiv.txt": { "abstract": "The dimensionless entropy , ${\\cal S} \\equiv S/k$, of the visible universe, taken as a sphere of radius 50 billion light years with the Earth at its \"center\", is discussed. An upper limit ($10^{112}$), and a lower limit ($10^{102}$), for ${\\cal S}$ are introduced. It is suggested that intermediate-mass black holes (IMBHs) constitute all dark matter, and that they dominate ${\\cal S}$. ", "introduction": "\\bigskip \\noindent Two references useful for further information about the material of this talk are: \\noindent (1) P.H.F. and T.W. Kephart. {\\it Upper and Lower Bounds on Gravitational Entropy}. \\noindent JCAP 06:008 (2008) and \\noindent (2) P.H.F. {\\it Identification of All Dark Matter as Black Holes}. {\\tt arXiv:0905.3632 [hep-th]}. JCAP 0910:016 (2009). ", "conclusions": "" }, "1003/1003.4740_arXiv.txt": { "abstract": "Accurate knowledge of the non-linear dark-matter power spectrum is important for understanding the large-scale structure of the Universe, the statistics of dark-matter haloes and their evolution, and cosmological gravitational lensing. We analytically model the dark-matter power spectrum and its cross-power spectrum with dark-matter haloes. Our model extends the halo-model formalism, including realistic substructure population within individual dark-matter haloes and the scatter of the concentration parameter at fixed halo mass. We consider three prescriptions for the mass-concentration relation and two for the substructure distribution in dark-matter haloes. We show that this extension of the halo model mainly increases the predicted power on the small scales, and is crucial for proper modeling the cosmological weak-lensing signal due to low-mass haloes. Our extended formalism shows how the halo model approach can be improved in accuracy as one increases the number of ingredients that are calibrated from $n$-body simulations. ", "introduction": "Accurate studies of large-scale cosmic structures and galaxy clustering became possible with the advent of large galaxy redshift surveys \\citep{guzzo05,meneux09}. Different analyses have been carried out depending on galaxy luminosity, color and morphology \\citep{norberg02,zehavi05,padmanabhan07}. According to the standard scenario of structure formation, galaxies with dissimilar features reside in different dark matter haloes and have experienced varied formation histories. Galaxies are believed to form and reside in dark-matter haloes that extend much beyond their observable radii. Some of them are located at halo centers while others orbit around it, constituting the satellite population. Dark-matter haloes form by gravitational instability from dark-matter density fluctuations \\citep{bond91,lacey93} and subsequently merge to form increasingly large haloes as cosmic time proceeds. Gas follows the dark-matter density perturbations. Once it reaches sufficiently high densities, dissipative processes, shocks and cooling allow stars to form from this gas \\citep{white78,kauffmann99}. While the main lines of this scenario are widely accepted, its details are still poorly understood. The clustering strength of a given galaxy population is related to that of the dark matter halos which host the galaxies. It is possible to provide an accurate description of clustering in the small-scale non-linear regime even if one has no knowledge of how the haloes themselves are clustered \\citep{sheth97b,smith03}. The halo-model of matter clustering \\citep{scherrer91,peacock00,seljak00,scoccimarro01,cooray02}, which has been the subject of much recent interest, allows a parametrization of clustering even on large scales. To date, almost all analytic work based on the halo-model approach assumes that haloes are spherically symmetric and that the matter density distribution around each halo centre is smooth. However, numerical simulations of hierarchical clustering have shown that haloes are neither spherically symmetric \\citep{jing02,allgood06,hayashi07} nor smooth \\citep{moore98,springel01b,gao04,delucia04,tormen04,giocoli08b}. About $10\\%$ of the mass in cluster-sized haloes is associated with subclumps. A halo model which includes the effects of halo triaxiality on various clustering statistics is developed in \\cite{smith05,smith06}, and formalism to include halo substructures was developed by \\cite{sheth03c}. The main purpose of this work is to include recent advances in our understanding of halo substructures into the halo model approach. This is because an accurate model of substructures is a necessary first step to modeling the small scale weak gravitational lensing convergence and shear signals \\citep{bartelmann01, hagan05} and for the substructure contribution to the gravitational flexion \\citep{bacon06}. The present paper is organized as follows. In Sect.~\\ref{themodel}, we describe the ingredients of our extension of the halo model: the properties of subhaloes and the substructure mass function. In Sect.~\\ref{nonlinearps}, we show how to incorporate these into the halo model. Cross-correlations between haloes and mass as well as clumps and mass are studied in Sect.~\\ref{crosscorrelations}. Sect.~\\ref{sandc} summarizes our methods and conclusions. ", "conclusions": "\\label{sandc} We have shown how the halo-model formalism can be extended to account for scatter in halo concentration at fixed mass, and for the presence of substructures in dark matter haloes. Differences in the mass-concentration relation do affect the predicted non-linear power spectrum. We quantified this by considering three models for the $M$-$c$ relation, as well as a deterministic and a stochastic models for the concentration at fixed mass. Accounting for substructure means that the 1-halo and 2-halo terms should be written as sums of four and three types of pairs, respectively. If one uses realistic models for these different pair-types, then the halo model calculation is in reasonable agreement with the non-linear power spectrum measured in the Millennium Simulation, over a range of redshifts. The cross-correlation between halos and mass can also be estimated using the halo-model formalism. We have shown how, also in this case, substructures can be taken into account and how the cross-correlation signal between clumps and mass can be split into different contributions. The agreement with simulations is not yet at the percent level, so there is room for improvement. The simulations show an excess at around $k\\sim 1$ relative to our model -- this is almost certainly not due problems with our substructure calculation, since it is on larger scales than those on which substructures dominate. The discrepancy is also unlikely to be due to our decision to ignore `assembly bias' effects. Rather it may be due to the fact that our 2-halo term is too crude -- on these smaller scales, one must include the effects of nonlinear bias, perhaps following \\citet{smith07}. \\begin{itemize} \\item Our formulation of the halo model assumes that the mass fraction in substructures is a deterministic function of the halo concentration: in practice, there is a distribution of $f_s$ values at fixed $c$ and $M$. Allowing for this stochasticity simply adds one more integral in most of our expressions, that can matter at the percent level on the small scales of interest here (for essentially the same reason that scatter in $c$ and fixed $M$ matters at the ten percent level on small scales). \\item We only include some of the effects of the known correlation between subhalo concentration and distance from the center of the host. Specifically, our current implementation accounts for the fact that the spatial dependence means subhalo concentrations (at fixed mass) are stochastic, but it does not include the fact that this is correlated with distance from host halo center. Implementing a spatially dependent concentration is difficult in Fourier space, so doing so in configuration is probably the best way forward. \\item We also do not account for the fact that subhalos themselves have substructure. This will matter most for the contribution which comes from pairs which are in the same subhalo (for the same reason that substructure is not important for the 2-halo term). \\item Halo exclusion matters for the cross-correlation signal. There are exclusion effects associated with the subhalos too, which we currently ignore. \\item Although the sum of the smooth and clumpy components should give an NFW profile, our current implementation assumes the smooth component follows and NFW profile, whereas the clumpy one does not. Hence there is no guarantee that the sum of the two actually is NFW at high precision. While this is relatively easy to fix (simply define the profile of the smooth component to be the required NFW minus the clumpy component), we have not done so. \\item And finally, at this level of precision, halo shapes also matter. In particular, the distribution of $f_s$ and $c$ at fixed $M$ may depend on halo shape -- once this has been quantified in simulations, it may become necessary to extend our formalism to include shapes. \\end{itemize} The formalism presented in this paper can be used in subsequent studies to make more realistic predictions for the galaxy-galaxy lensing signals and for the lensing convergence power spectrum. By identifying the different contributions due to the smooth and clumpy components of the dark-matter haloes, this will also allow us to quantify the contributions the convergence power spectrum not just by different halo masses, but also by different substructures." }, "1003/1003.3483_arXiv.txt": { "abstract": "\\vskip1em \\noindent We compute the transition amplitude between coherent quantum-states of geometry peaked on homogeneous isotropic metrics. We use the holomorphic representations of loop quantum gravity and the Kaminski-Kisielowski-Lewandowski generalization of the new vertex, and work at first order in the vertex expansion, second order in the graph (multipole) expansion, and first order in volume${}^{-1}$. We show that the resulting amplitude is in the kernel of a differential operator whose classical limit is the canonical hamiltonian of a Friedmann-Robertson-Walker cosmology. This result is an indication that the dynamics of loop quantum gravity defined by the new vertex yields the Friedmann equation in the appropriate limit. ", "introduction": "The dynamics of loop quantum gravity (LQG) can be given in covariant form by using the {\\em spinfoam} formalism. In this paper we apply this formalism to cosmology. In other words, we introduce a spinfoam formulation of quantum cosmology, or a ``spinfoam cosmology\". We obtain two results. The first is that physical transition amplitudes can be computed, in an appropriate expansion. We compute explicitly the transition amplitude between homogeneous isotropic coherent states, at first order. The second and main result is that this amplitude is in the kernel of an operator $\\hat C$, and the classical limit of $\\hat C$ turns out to be precisely the Hamiltonian constraint of the Friedmann dynamics of homogeneous isotropic cosmology. In other words, we show that LQG yields the Friedmann equation in a suitable limit. LQG has seen momentous developments in the last few years. We make use of several of these developments here, combining them together. The first ingredient we utilize is the ``new\" spinfoam vertex \\cite{Engle:2007uq,Livine:2007vk,Engle:2007qf,Freidel:2007py,Engle:2007wy}. The second is the Kaminski-Kisielowski-Lewandowski extension of this to vertices of arbitrary-valence \\cite{Kaminski:2009fm}. The third ingredient is the coherent state technology \\cite{Hall:2002, Ashtekar:1994nx, Thiemann:2000bw,Thiemann:2000ca,Thiemann:2000bx, Thiemann:2000by,Sahlmann:2001nv,Thiemann:2002vj, Bahr:2007xa,Bahr:2007xn,Flori:2008nw,Flori:2009rw, Freidel:2010aq,Freidel:2010}, and in particular the holomorphic coherent states discussed in detailed in \\cite{Bianchi:2009ky}. These states define a holomorphic representation of LQG \\cite{Ashtekar:1994nx,Bianchi:2010}, and we work here in this representation. Our strategy is the following. We consider the standard Hilbert space of canonical LQG and we assume the dynamics to be given by the new vertex. We consider holomorphic coherent states in this Hilbert space and we work in the holomorphic representation they define. We truncate LQG down to a graph with a finite number of links. In particular, the calculation is based on the ``dipole\" graph formed by two nodes connected by four links \\cite{Rovelli:2008dx}. This choice determines a Hilbert space, which describes a finite number of the degrees of freedom of the gravitational field. These degrees of freedom can be identified as the lowest modes in a multipole expansion of the metric in hyper-spherical harmonics on $S_3$ \\cite{Battisti:2009kp}. That is, they describe a closed cosmology, with anisotropies and a few low-mode inhomogeneities. In particular, we consider coherent states that are peaked on homogeneous isotropic geometries. We emphasize the fact that these states are just peaked on homogeneous and isotropic geometries, but they also include fluctuations of the inhomogeneous and anisotropic degrees of freedom. So, the dynamics of the quantum theory we consider {\\em does} include inhomogeneous and anisotropic degrees of freedom. Homogeneous-isotropic coherent states are labelled by two parameters which capture the scale factor $a$ of standard cosmology and its time derivative $\\dot a$; or, equivalently, the $p$ and $c$ canonical variables used in Loop Quantum Cosmology (LQC). In the holomorphic representation, these two quantities appear in the complex combination $z=\\alpha c+i\\beta p$, and therefore the states we consider are labelled by the complex number $z$. The transition amplitude between two such states is then an analytic function $W(z,z')$ of two complex variables. We write this transition amplitude at first order in a vertex expansion. We view this as the analog of a first order calculation in, say, QED perturbation theory. We compute explicitly $W(z,z')$ in the limit in which the geometry is large compared to the Planck scale. In other words, we compute the transition amplitude between macroscopical homogeneous isotropic cosmological spaces in LQG, taking three approximations from the complete theory: (i) the truncation of the degrees of freedom to those defined on a finite graph; (ii) the restriction to first order in the vertex expansion; (iii) the large volume limit. The validity of these approximation can only be justified a posteriori, from the correctness of the result. Our next step is to notice that the transition amplitude computed solves the equations \\mbox{$H\\, W(z,z')=0$} for a certain operator $H=H(z, {\\scriptstyle\\hbar} \\frac{\\d}{\\d z})$. This fact implies that the amplitude defines a quantum dynamics where the operator constraint $H=0$ holds. The corresponding classical dynamics will be governed by (the $\\hbar\\to 0$ limit of) the classical constraint $H(z, \\bar z)=0$. When written in terms of $p$ and $c$, this turns out to be precisely the Hamiltonian constraint that governs (the gravitational part of) the dynamics of a classical Friedmann cosmology, in the limit of large volume. Therefore LQG yields the Friedmann dynamics in this limit. Several words of caution are necessary. First, we work in the Euclidean theory. Second, the cosmological dynamics that we obtain is the one in the large volume limit and since we do not have any matter present, this has only the solution $a=constant$, which is flat space. With these caveats, our result is that there is an approximation in LQG that leads to classical cosmology. This result can be compared with those of LQC \\cite{Bojowald:2006da,Ashtekar:2008zu}. In LQC, one first reduces the classical theory to a cosmological system with a finite number of degrees of freedom, and then applies a ``loop quantization\" to this symmetry reduced model. Thus, one has a complete quantum theory of a truncation of the classical theory. Here, instead, we start from the full quantum theory and take an approximation. Therefore in LQC one studies exact solution in a truncated system, while here we study approximated solutions in the (hopefully) exact quantum theory. The possibility of introducing a spinfoam-like expansion starting from LQC has been considered in the papers \\cite{Ashtekar:2009dn,Ashtekar:2010ve,Rovelli:2009tp}. These papers and the present work can be seen as two converging attempts to construct a spinfoam version of quantum cosmology. Finally, in our opinion a main reason of interest of the result we present here is that it represents an example of a complete calculation of physical transition amplitude in background independent quantum gravity. It complements the calculation of the two-point function \\cite{Rovelli:2005yj,Bianchi:2006uf,Alesci:2007tx,Alesci:2007tg,Alesci:2008ff}, that has been recently completed \\cite{Bianchi:2009ri}. \\vskip.5cm In Section \\ref{theory}, we briefly recall the definition of the full quantum theory. In Section \\ref{cosmo}, we discuss the approximation that selects a cosmological sector, we compute the resulting transition amplitude. In Section \\ref{limit}, we study the classical limit and recover the Friedmann dynamics. ", "conclusions": "We have introduced a spinfoam formulation of quantum cosmology. We have obtained two results. The first is that it is possible to compute quantum transition amplitudes explicitly in suitable approximations. In detail, we have studied three approximations: (i) cutting the theory to a finite dimensional graph (the dipole), (ii) cutting the spinfoam expansion to just one term with a single vertex and (iii) the large volume limit. The main hypothesis on which this work is based is that the regime of validity in which these approximations are viable includes the semiclassical limit of the dynamics of large wavelengths. ``Large\" means here of the order to the size of the universe itself. This regime includes of course the standard Friedmann cosmology. The second result is that the transition amplitude computed appears to give the correct Friedmann dynamics in the classical limit. This results must be taken with caution, for a number of reasons. First, we have used the Euclidean theory, instead than the physical Lorentzian theory. Second, the dynamics we have obtained is in fact rather trivial. The solution of the constraint equation \\eqref{fine} is either $p=0$ or $c=0$; that is, either the universe has no volume, or it is flat. This is physically correct, since in absence of a matter and in the limit of infinitely large radius one obtain precisely a flat spacetime. But is is only a weak indication that the full Friedmann dynamics is recovered. Wether the result still holds with matter, or a cosmological constant, must still be check. Also, in the derivation of the classical limit, the symplectic structure \\eqref{ss} has been taken as an input. It can be shown that this choice reproduces the symplectic structures of the LQC variables\\footnote{ This fixes the product $\\alpha\\beta=3t/16\\pi G\\g$. Then $\\alpha$ can be determined by noticing that for $c=1$ the connection $A$ is the Cartan connection and its holonomy from the identity to $g$ is $g$ itself. Taking $n_s=\\id$ and $n_t=-\\id$ gives easily $\\alpha=2\\pi$. } $(c,p)$ or the one of the LQG variables $(E_\\ell,U_\\ell)$. Finally, the system that the approximation defines admits obvious improvements. In particular, transitions must be computed on a larger graph, and at the next order in the vertex amplitude, in order to investigate the validity of the approximation. We have noticed that at first order the transition amplitude factorizes (see eq.\\eqref{Wzz}). To this order, the ``projector\" on physical states $P=\\sum_n \\bk{n}{n}$ defined by $\\bek{\\psi_\\fin}{P}{\\psi_\\ini}=\\bk{W}{\\bar\\psi_\\fin \\otimes \\psi_\\ini}$ projects on a single state, say $\\ket0$, which can be identified as the Hartle-Hawking ``wave function of the universe\" defined by the so called ``no boundary proposal\" \\cite{Hawking:1980gf}. We do not expect this factorization to survive higher orders, where the projector can regain its general form. From the point of view of cosmology, the system we have described opens in principle the way to the description of inhomogeneous degrees of freedom at the bounce, circumventing the difficulties of the model given in \\cite{Rovelli:2008dx}. In particular, the covariant dynamics used here can readily be extended to larger graphs. Coherent states have been largely used in loop quantum cosmology (see for instance \\cite{Ashtekar:2006rx,Ashtekar:2006uz,Ashtekar:2006wn,Ashtekar:2006es}) in particular in relation to the problem of finding effective equations or in numerical simulations \\cite{Singh:2005xg,Bojowald:2009jk,Bojowald:2009zz}. Here, however, homogeneous and isotropic states appear naturally as states peaked on homogeneous and isotropic \\emph{mean values} of the quantum states, in the context of a formalism which --we stress-- is not a reduction of the dynamics to homogeneous and isotropic degrees of freedom. In physical terms, these states represent a universe where inhomogeneous and anisotropic degrees of freedom are taken into account but fluctuate around zero. This provides also an elegant solution of the problem of having to choose between coordinate or momenta in imposing a symmetry reduction in cosmology \\cite{Engle:2005ca,Engle:2007zz,Engle:2007qh}. Ideally, this formalism could describe inhomogeneous and anisotropic quantum fluctuations of the geometry at the bounce. \\subparagraph{Acknolegements} A warm thank to Antonino Marcian\\`o, Elena Magliaro and Claudio Perini, for a continuous collaboration that has been essential for the developments of the ideas in this work. We also thank the partecipants to the workshop ``Open problems in Loop Quantum Gravity'' in Zakopane, Poland, for many interesting comments. The work of E.B. is supported by a Marie Curie Intra-European Fellowship within the 7th European Community Framework Programme." }, "1003/1003.2166_arXiv.txt": { "abstract": "\\noindent We present submillimetre and mid-infrared imaging observations of five fields centred on quasi-stellar objects (QSOs) at $1.73$) radio galaxies (HzRGs) concerns very bright SMGs in their fields. Two of the seven HzRG fields imaged with SCUBA contain SMGs with $S_{850}>20$ mJy (Stevens et al. 2003) whereas the brightest object found around the QSOs has $S_{850}\\sim10$~mJy. Although the samples are too small to give a statistically significant result it appears that the SMGs detected around the lower redshift, less powerful, radio-quiet AGN have lower $S_{850}$. A direct comparison is not informative because the sources were extracted with different signal-to-noise criteria and by different methods but comparing flux densities significantly above the catalogue cut-offs yields $8$ SMGs with $S_{850}\\geq6$~mJy in the 7 HzRG fields while only $2$ are detected around the 5 QSOs. Given the effectively flat relationship between $S_{850}$ and redshift for $12.5$ (Archibald et al. 2001; Reuland et al. 2004) which is significantly higher than the average redshift of our QSO sample. Another possibility is that the radio properties of the central AGN are in some way linked to environmental density (e.g. Kauffmann, Heckman \\& Best 2008). Indeed, recent work by Falder et al. (2010) finds evidence for larger over-densities of galaxies around radio-loud objects (based on work carried out at 3.6~$\\mu$m). Perhaps radio synchrotron luminosity is enhanced in regions where the inter-galactic medium is more dense. An obvious next step will be to conduct a statistically meaningful survey of fields around AGN over the full range of redshifts in order to track star-formation activity as a function of cosmic time. The next generation bolometer camera on the JCMT (SCUBA-2) and SPIRE/PACS on-board the {\\em Herschel Space Observatory\\/} are instruments with good enough sensitivity to make such a campaign practical. It would also allow a bigger area to be mapped around each AGN, better matching the size of proto-cluster regions predicted by numerical simulations." }, "1003/1003.0449_arXiv.txt": { "abstract": "We present first results from Galaxy Zoo 2, the second phase of the highly successful Galaxy Zoo project ({\\tt www.galaxyzoo.org}). Using a volume--limited sample of 13665 disk galaxies ($0.01< z < 0.06$ and $M_r<-19.38$), we study the fraction of galaxies with bars as a function of global galaxy properties like colour, luminosity and bulge prominence. Overall, $29.4\\pm0.5\\%$ of galaxies in our sample have a bar, in excellent agreement with previous visually--classified samples of galaxies (although this overall fraction is lower than measured by automated bar--finding methods). We see a clear increase in the bar fraction with redder $(g-r)$ colours, decreased luminosity and in galaxies with more prominent bulges, to the extent that over half of the red, bulge--dominated, disk galaxies in our sample possess a bar. We see evidence for a colour bi-modality for our sample of disk galaxies, with a ``red sequence\" that is both bulge and bar--dominated, and a ``blue cloud\" which has little, or no, evidence for a (classical) bulge or bar. These results are consistent with similar trends for barred galaxies seen recently both locally and at higher redshift, and with early studies using the RC3. We discuss these results in the context of internal (secular) galaxy evolution scenarios and the possible links to the formation of bars and bulges in disk galaxies. ", "introduction": "Bars are common in disk galaxies, and are thought to have an important impact on the evolution of galaxies through their ability to transfer angular momentum in both the baryonic and dark matter components of the galaxy \\citep{CS81,W85,DS00,B06}. Bars are efficient at driving gas inwards, perhaps sparking central star formation \\citep[e.g.][]{H86,K95,J05,S05}, and thus help to grow a central bulge \\citep[e.g. for a review][]{KK04}. Bars may also feed a central black hole (as per \\citealt{S89,S90} and see \\citealt{J06} for a recent review), but so far no correlation has been found between AGN activity and bar fraction in galaxies forming stars (e.g. \\citealt{H97} and as recently discussed by \\citealt {H09}). Early visual inspection of spiral galaxies in catalogues like the RC3 \\citep{RC3} gave an optical bar fraction of $f_{\\rm bar} \\sim 0.25-0.3$, rising to 60\\% if weaker bars or oval distortions were included \\citep[as discussed in e.g.][]{SW93,M95,KSP00,S08} More recent work on the bar fraction of disk galaxies has relied on automated methods of detecting bars in galaxies including elliptical isophote fitting or the Fourier decomposition of CCD images. Such automated studies find optical bar fractions of $\\sim$50\\% for nearby disk galaxies \\citep{Bar08,Ag09}, which is consistent with near infra-red (NIR) studies that find a majority (at least 60\\%) of disk galaxies appear to have a bar \\citep[e.g.][]{MR97,MJ07}. These differences are probably due to a combination of selection effects, wavelength-dependences, differences in the strength of the bar, and small samples sizes. While studies of bars in galaxies have been collectively moving towards automated classifications (and away from possibly subjective visual classifications) in recent years, concerns have arisen about the reliability of a completely automated approach for detecting bars which struggle to distinguish spiral arms and bars in some cases \\citep[see for example the discussion of problems with the Fourier method in][]{Ag09}. For such reasons, especially to provide a larger sample of visually selected barred galaxies, we started a new phase of the successful Galaxy Zoo project\\footnote{{\\tt www.galaxyzoo.org}} by asking the public to provide more detailed visual classifications of galaxies seen in the Sloan Digital Sky Survey (SDSS). This new project is known as ``Galaxy Zoo 2\" (GZ2) throughout this paper. The project and data set will be fully described in a future paper (Lintott et al. in prep.). In this article, we present the first results on the bar properties of 13665 visually classified GZ2 disk galaxies. This sample is nearly an order of magnitude larger than previous studies using SDSS data \\citep{Bar08,Ag09} which facilitates a detailed statistical study of the fraction of barred disk galaxies as a function of other galaxy properties like global optical colour, luminosity, and estimates of the bulge size, or prominence. Where appropriate, we assume a standard cosmological model of $\\Omega_{m}$ = 0.3, $\\Omega_{\\Lambda}$ = 0.7 and $H_{0} = 70$ \\kmsMpc and all photometric quantities are taken from SDSS. ", "conclusions": "\\subsection{Comparison with Other Work} The literature on bar fractions is extensive going back to seminal early work on optical galaxy catalogues. Many studies have attempted to separate the bar fraction into galaxies of different morphological types \\citep[e.g.][]{O96,KSP00,E04} and a picture is emerging which suggests that the bar fraction found in a given sample might depend quite sensitively on the morphological and stellar mass make up of the sample being considered \\citep[as recently discussed by][]{G10,NA10}. One must of course be careful in comparing {\\bf any} galaxy properties across studies using samples at different redshifts and with different morphological and luminosity/mass distributions, but a review of results from different studies is still useful to put our results into context. In this work, we observe a significant increase in the bar fraction of disk galaxies as the galaxies become redder and have more prominent bulges. We also observe a small increase in the very bluest galaxies in our sample. This suggests that the bar fraction is highest in early type disks, has a minimum somewhere in the blue cloud of spiral disks, and increases slightly towards the latest type spirals. In fact there has been a suggestion for some time that bar fraction does not vary monotonically with galaxy type. Both \\citet{O96} and \\citet{E04}, use visual classifications of RC3 galaxies observe that the (strong) bar fraction (i.e. spiral types SB) decreases from around 60\\% in S0/a to 30\\% in Sc, after which it increases again towards very late type disks. Interestingly both studies also show that the weak (or mixed type, SAB) bar fraction is much flatter with Hubble type and if anything shows the opposite trend. \\citet{KSP00} used the same RC3 data to argue that bar fraction remains relatively flat across all types of disk galaxies, however the trend they show in their Figure 1 (for strong bars at least) appears consistent with that used by both \\citet{O96} and \\citet{E04} to argue for variation. Almost concurrently with this work, \\citet{G10} in a multiwavelength study of galaxies in the Virgo cluster, show that the bar fraction depends sensitively on the morphological composition of a sample. In qualitative agreement with our results, they find in a sample of $\\sim 300$ disk galaxies that early type spirals have a higher bar fraction (45-50\\%) than late-type spirals (22-36\\%) and suggest the difference could be explained by the higher baryon fraction of earlier type spirals. Also in very recent work, \\citet{NA10} discuss the bar fraction in 14,043 visually classified galaxies (all classified by PN, \\citealt{NA10a}) and like us find the bar fraction to depend on morphology with a minimum near the division between the blue and red sequences. They suggest that to reconcile the apparently conflicting results on the bar fraction (and in particular the evolution of the bar fraction with redshift) found in the literature one need only consider the different stellar mass ranges of the samples in question. Other recent work on the bar fraction at low redshifts ($z < 0.1$ or so) also support this broad picture of bar fraction having a minimum at around Sc types and rising towards both earlier and later spiral/disk galaxy types. For example, \\citet{Lauri09} study the bar fraction in 127 early type spirals and find it increases from Sas to S0/a (but then drops significantly in S0s - however distinguishing unbarred S0s from elliptical galaxies is notoriously hard and may bias the S0 bar fraction low). At first glance our results appear in conflict with those of \\citet{Bar08, Ag09, We09} all of who argue that bar fraction increases as bulge prominence decreases, however we suggest that sample selection may again be the culprit, with these studies actually picking out significantly later spiral types than are found in our sample, and therefore seeing only the bluest end of the trend we show. For example \\citet{Bar08} and \\citet{Ag09}, who both find larger bar fractions in later/bluer disk galaxies (using a local samples of $\\sim2000$ ``disk\" or spiral galaxies from the SDSS) both use automated techniques to identify a disk/spiral sample of galaxies based on concentration and velocity dispersion \\citep{Ag09} or colour \\citep[][this study also considered Sersic fits, but the final results were for a colour--selected sample of ``spiral\" galaxies]{Bar08}. If we restrict our analysis to the range of colours explored by \\citet{Bar08}, then much of the trend we see in Figure 3 is missed and we would have actually witnessed a mild decrease in bar fraction towards the redder spirals, fully consistent with their findings. In more detail, we mimic the \\citet{Bar08} selection by using their U-V colour cut (from \\citet{B04}) with colour transformations from \\citet{S02}, plus $i<60^\\circ$, and $0.0110^{11}M_\\odot$, but there are almost no such galaxies in our volume--limited sample. While we find a similar overall bar fraction to the STAGES study of barred galaxies in a dense cluster at $z\\sim0.2$ (Marinova et al. 2009) our findings on the trends of bar fraction with other properties differ substantially from that study. They observed that bar fraction (in $\\sim 800$ galaxies found from ellipse fitting to B-band images) rises in brighter galaxies and those which have no significant bulge component and that bar fraction had no dependence on disk galaxy colour. While we do see an increase in bar fraction for brighter spirals (from $26\\pm1$\\% for those with $M_r>-20$, to $37\\pm2$\\% for those with $M_r<-22$), we argue that this is driven by the strong colour dependence of the bar fraction and at a fixed colour we see little dependence of bar fraction on luminosity (Figure 4) - the trend we find is also much smaller than seen by Marinova et al. (2009). Similarly, we agree with the overall bar fraction of 25\\% found for 945 galaxies by \\citet{Bar09} across both field and clusters environments at $z\\sim$0.4--0.8 (observed in rest frame B--V). However we again find opposite trends of bar fraction with bulge-prominence (they find more bars in bluer disk-dominated galaxies). The source of these discrepancies is unclear. Disk galaxies were identified visually in both studies so should be similar to our GZ disks. We do use quite different bar finding techniques (visual versus ellipse fitting), so perhaps this indicates a difference in the trends for strong (visual) and weaker bars (as also hinted at by \\citealt{O96} and \\citealt{E04}). More interestingly it could be pointing to a difference between disk galaxies in high density regions and those elsewhere (as also discussed by \\citet{G10} who find similar results to us in Virgo cluster galaxies). \\citet{Bar09} explored differences between bar fractions in the field and clusters finding hints that high density regions are favourable locations for bars, however this could only be done for a subset of the sample ($N=241$), making the results of limited statistical significance. These possibilities will be explored in future work exploiting the huge number of bar classifications available to us in GZ2. There does remain a difference in our overall conclusions with \\citet{Bar08} and \\citet{Ag09} and, in particular, the total bar fraction we find is lower than either of these studies. It has been argued that bars in early-type spiral galaxies tend to be longer (and thus stronger) than those seen in late--type spirals \\citep{A03}, so it may be that the remaining differences are due to our sensitivity to these longer, stronger bars, while \\citet{Bar08} and \\citet{Ag09} may detect weaker bars using their ellipse--fitting techniques. The cross-over between the sample used here and that in \\citet{Bar08} is small (as we have argued above it is this mismatch which explains the different trends we see in bar fraction with colour and bulge prominence), but in $\\sim$400 galaxies found in both samples, we find GZ2 bars in $\\sim$ 25\\% of the objects while the number of barred objects rises to roughly twice that in the \\citet{Bar08} classifications. The two bar classification methods agree $\\sim$75\\% of the time -- most of the difference is from \\citet{Bar08} identified bars not being found by GZ2. Further studies directly comparing the bars identified from ellipse--fitting methods and the GZ2 identifications will be needed to understand the main reason for this difference. Such a comparison is in progress (Masters et al. in prep.). However, it is interesting that the bar fraction difference between the two techniques is rather similar to the split between strongly barred (SB), weakly barred (SAB) and non-barred (SA) galaxies in both the RC3 \\citep[e.g.][]{SW93,E00} and the de Vaucouleurs Atlas of Galaxies (Buta, Cordwin \\& Odewahn 2007). This seems to further indicate that (as we suggested above) GZ2 bars should be identified with strong bars (SB types) only and so the trends we observe should most likely be considered trends in the fraction of strong bars. \\subsection{The Impact of Bars on Disk Galaxies} Given the trends we have observed, we now focus on the interpretations we can make for the effect of bars on the secular and dynamical evolution of disk galaxies. We observe a significant increase in bar fraction as disk galaxies become redder and have larger (classical) bulges; over half of red, bulge--dominated disk galaxies have a bar. Our observations suggest an important link between the presence of a bulge (perhaps preferentially a classical bulge with a de Vaucouleur profile) and the existence of a bar instability. Bar instabilities are often invoked as a way to form pseudo-bulges (with an exponential profile) by moving material around in the disk of a spiral galaxy (see \\citet{KK04} for a comprehensive review of this subject), but classical bulges are usually thought to have formed during a fast, dissipative process (most likely related to galaxy mergers), which would have likely disrupted any bar. However using a sample of 143 local galaxies with bulge-disk-bar decompositions, \\citet{We09} argue that most bright spirals have (almost) pseudo-bulges and comparing with cosmological simulations argue that most of these must have been made by a combination of minor mergers and secular evolution. Perhaps it is this link we are seeing in the Galaxy Zoo sample. In \\citet{G10} the observation that the barred fraction is higher in early type disk galaxies in their Virgo cluster sample is tentatively explained by a combination of the higher baryon fraction in early type galaxies (which makes their discs heavier and therefore more susceptible to bar instabilities) and environmental effects which might destroy a late-type spiral but would leave an early type spiral with a bar. Further studies of the environmental dependence of the bar fraction are planned with GZ2 data (Skibba et al. in prep.) and will be used to test this scenario. We finish by returning to the suggestion that we see two populations of disk galaxies. Both Figure 3 and 7 suggest a split between disk galaxies on the ``red sequence\", which have large (possibly classical) bulges (maybe formed during merger processes), and disk galaxies in the ``blue cloud\" with either no bulge or a pseudo--bulge. The red sequence population show little change in their bar fraction with luminosity and colour and overall have a high fraction of (strong) bars, approaching $50\\%$. The blue cloud population also show little trend with colour with low bar fractions of 10-20\\%." }, "1003/1003.2957_arXiv.txt": { "abstract": "Heavy Ion Collisions ($HIC$) represent a unique tool to probe the in-medium nuclear interaction in regions away from saturation. In this report we present a selection of new reaction observables in dissipative collisions particularly sensitive to the symmetry term of the nuclear Equation of State ($Iso-EoS$). We will first discuss the Isospin Equilibration Dynamics. At low energies this manifests via the recently observed Dynamical Dipole Radiation, due to a collective neutron-proton oscillation with the symmetry term acting as a restoring force. At higher beam energies Iso-EoS effects will be seen in an Isospin Diffusion mechanism, via Imbalance Ratio Measurements, in particular from correlations to the total kinetic energy loss. For fragmentation reactions in central events we suggest to look at the coupling between isospin distillation and radial flow. In Neck Fragmentation reactions important $Iso-EoS$ information can be obtained from fragment isospin content, velocity and alignement correlations. The high density symmetry term can be probed from isospin effects on heavy ion reactions at relativistic energies (few $AGeV$ range), in particular for high transverse momentum selections of the reaction products. Rather isospin sensitive observables are proposed from nucleon/cluster emissions, collective flows and meson production. The possibility to shed light on the controversial neutron/proton effective mass splitting in asymmetric matter is also suggested. A large symmetry repulsion at high baryon density will also lead to an ``earlier'' hadron-deconfinement transition in n-rich matter. The binodal transition line of the ($T,\\rho_B$) diagram is lowered to a region accessible through heavy ion collisions in the energy range of the new planned facilities, e.g. the $FAIR/NICA$ projects. Some observable effects of the formation of a Mixed Phase are suggested, in particular a {\\it Neutron Trapping} mechanism. The dependence of the results on a suitable treatment of the isovector part of the interaction in effective QCD Lagrangian approaches is critically discussed. We stress the interest of this study in nuclear astrophysics, in particular for supernovae explosions and neutron star structure, where the knowledge of the $Iso-EoS$ is important at low as well as at high baryon density. ", "introduction": "The study of the behavior of nuclear matter under several conditions of density and temperature is of crucial importance for the understanding of a large variety of phenomena, ranging from the structure of nuclei and their decay modes, up to the life and the properties of massive stars. Mechanisms involving an enormous range of scales in size, characteristic time and energy, but all based on nuclear processes at fundamental level, are actually linked by the concept of the nuclear Equation of State ($EoS$) and the associated energy density functional. In particular, the understanding of the properties of exotic nuclei, as well as neutron stars and supernova dynamics, entails constraining the behavior of the nuclear symmetry energy, which describes the difference between the binding energy of symmetric matter (with equal proton and neutron numbers, N=Z), and that of pure neutron matter. Transient states of nuclear matter far from normal conditions can be created in terrestrial laboratories via Heavy Ion Collisions ($HIC$). Many experimental and theoretical efforts have been devoted, over the past 30 years, to the study of nuclear reactions from low to intermediate energies, as a possible tool to learn about the behavior of nuclear matter and its $EoS$. Relevant conclusions have been reached concerning the $EoS$ of symmetric matter for densities up to five time the saturation value \\cite{daniel02}. More recently, the availability of neutron-rich and exotic beams has opened the way to investigate, in laboratory conditions, new aspects of nuclear structure and dynamics up to extreme ratios of neutron to proton numbers. Thus it has become possible to explore the behavior of nuclear matter along a new degree of freedom, the asymmetry $I = (N-Z)/(N+Z)$ (in the rest of the review also defined as $\\beta$ or $\\alpha$), aiming at constraining the density and/or temperature dependence of the symmetry energy ($Iso-EoS$). Here we will review the Isospin Dynamics in $HIC$ from the Coulomb Barrier to Relativistic Energies. The symmetry energy $E_{sym}$ appears in the energy density $\\epsilon(\\rho,\\rho_3)\\equiv \\epsilon(\\rho)+\\rho E_{sym} (\\rho_3/\\rho)^2 + O(\\rho_3/\\rho)^4 +..$, expressed in terms of total ($\\rho=\\rho_p+\\rho_n$) and isospin ($\\rho_3=\\rho_p-\\rho_n$) densities \\cite{baranPR}. The symmetry term gets a kinetic contribution directly from basic Pauli correlations and a potential part from the highly controversial isospin dependence of the effective interactions. Both at sub-saturation and supra-saturation densities, predictions based of the existing many-body techniques diverge rather widely, see \\cite{fuchswci,fantoni08}. We recall that the knowledge of the $EoS$ of asymmetric matter is very important at low densities, in nuclear structure ( neutron skins, pigmy resonances, refs. \\cite{gsi07,pieka06,trippa08,carbone10}, in reactions (neutron distillation in fragmentation \\cite{baran98}, charge equilibration \\cite{tsang04}), in astrophysics (neutron star formation and crust, \\cite{horopieka02, steiner05}) as well as at high densities in relativistic heavy ion reactions (isospin flows \\cite{greco03}, particle production \\cite{baranPR,baoPR08,ferini05,ferini06}), in compact star (neutron star mass-radius relation, cooling, hybrid structure, formation of black holes, \\cite{page06,latpra07,prakash07,baldo07}) and for fundamental properties of strong interacting systems (transition to new phases of the matter, \\cite{muller,ditorodec}). Several observables which are sensitive to the $Iso-EoS$ and testable experimentally, have been suggested \\cite{baranPR,baoPR08,colonnaPRC57,baoIJMPE7,Isospin01,baran2004, wcineck,WCI_betty,sheyen10}. We take advantage of new opportunities in theory (development of rather reliable microscopic transport codes for $HIC$) and in experiments (availability of very asymmetric radioactive beams, improved possibility of measuring event-by-event correlations) to present new results that are constraining the existing effective interaction models. We will discuss dissipative collisions in a wide range of beam energies, from just above the Coulomb barrier up to the $AGeV$ range. Isospin effects on the chiral/deconfinement transition at high baryon density will be also analyzed. Low to Fermi energies will bring information on the symmetry term around (below) normal density, while intermediate energies will probe high density regions. The transport codes are based on mean field theories, with correlations included via hard nucleon-nucleon elastic and inelastic collisions and via stochastic forces, selfconsistently evaluated from the mean phase-space trajectory, see \\cite{baranPR,guarneraPLB373,colonnaNPA642,chomazPR}. Stochasticity is essential in order to get distributions as well as to allow the growth of dynamical instabilities. Relativistic collisions are described via a fully covariant transport approach, related to an effective field exchange model, where the relevant degrees of freedom of the nuclear dynamics are accounted for \\cite{baranPR,ferini05,ferini06,liubo02,theo04,santini05}. We will have a propagation of particles suitably dressed by self-energies that will influence collective flows and in medium nucleon-nucleon inelastic cross sections. The construction of an $Hadron-EoS$ at high baryon and isospin densities will finally allow the possibility of developing a model of a hadron-deconfinement transition at high density for an asymmetric matter \\cite{ditorodec}. The problem of a correct treatment of the isospin in an effective partonic $EoS$ will be stressed. ", "conclusions": "" }, "1003/1003.5600_arXiv.txt": { "abstract": "We investigate the variability of \\ion{C}{4}~$\\lambda$1549 broad absorption line (BAL) troughs over rest-frame time scales of up to $\\approx$7~yr in 14~quasars at redshifts $z \\ga 2.1$. For 9 sources at sufficiently high redshift, we also compare \\ion{C}{4} and \\ion{Si}{4}~$\\lambda$1400 absorption variation. We compare shorter- and longer-term variability using spectra from up to four different epochs per source and find complex patterns of variation in the sample overall. The scatter in the change of absorption equivalent width (EW), $\\Delta$EW, increases with the time between observations. BALs do not, in general, strengthen or weaken monotonically, and variation observed over shorter ($\\lesssim$months) time scales is not predictive of multi-year variation. We find no evidence for asymmetry in the distribution of $\\Delta$EW that would indicate that BALs form and decay on different time scales, and we constrain the typical BAL lifetime to be $\\gtrsim$30~yr. The BAL absorption for one source, LBQS~0022+0150, has weakened and may now be classified as a mini-BAL. Another source, 1235+1453, shows evidence of variable, blue continuum emission that is relatively unabsorbed by the BAL outflow. \\ion{C}{4} and \\ion{Si}{4} BAL shape changes are related in at least some sources. Given their high velocities, BAL outflows apparently traverse large spatial regions and may interact with parsec-scale structures such as an obscuring torus. Assuming BAL outflows are launched from a rotating accretion disk, notable azimuthal symmetry is required in the outflow to explain the relatively small changes observed in velocity structure over times up to 7~yr. ", "introduction": "} High-velocity, structured outflows in the central regions of quasars (QSOs) generate broad absorption lines (BALs) from material along the line of sight to the QSO ultraviolet (UV) emitter. These ``BAL outflows'' are launched, ionized, and accelerated by physical processes in the heart of the QSO, and can provide an important means of carrying material and energy out of the QSO's central region. BAL features can reach high outflow velocities (30,000~km~s$^{-1}$ or more), and the absorption signatures of these outflows are, by definition, at least 2000~km~s$^{-1}$ wide \\citep{wmfh91}, although BALs can be much broader. BALs are observed in up to $\\approx$15\\% of QSOs in optical/UV spectroscopic surveys \\citep[e.g.,][and references therein]{wmfh91, hf03, rrhsvfykb03, thrrsvkafbkn06, gjbhswasvgfy09}, and the fraction of QSOs with BALs rises to 17--23\\% after correcting for optical/UV selection effects \\citep[e.g.,][]{hf03, ksgc08, gjbhswasvgfy09}. The intrinsic fraction may be $\\sim$2 times higher if less-restrictive criteria are used to select BALs \\citep[e.g.,][]{thrrsvkafbkn06, dss08, gb08}; in particular, narrower, mini-BAL absorption features may also be related to BALs \\citep[e.g.,][]{akdjb99, gbcg02, gbgs09}. Some physical models associate BAL outflows with equatorial accretion disk winds that may be ubiquitous in QSOs \\citep[e.g.,][]{mcgv95, psk00}. BAL outflows may also result during evolutionary stages when QSOs expel thick shrouds of gas and dust \\citep[e.g.,][]{vwk93, gbd06}. High-resolution spectroscopic studies have found that BAL absorption is frequently saturated, but does not completely obscure the continuum emitter. \\citet{ablgwbd99} concluded that BAL absorption troughs are generally shaped by the outflow kinematic structure and geometry, with the depth of the absorption feature representing the fraction of the continuum emitter covered by one or more components of the absorbing outflow at a given velocity. In this model, studies of BAL variability primarily reveal ongoing changes in the structure of the outflow. Variability studies have also sought to test the ``covering factor'' hypothesis by identifying correlations between absorption and emission variation that would indicate that absorption is not strongly saturated, but that the absorption strength is related to the photoionization-dependent column density of the absorbing ion. The results of these studies are mixed, but conclusive evidence for photoionization-dominated variation has not been found in the general population of BALs \\citep[e.g.,][]{bjb89, bjbwmk92, lwbhsyvb07, gbsg08G08}. With sufficient spectroscopic observing resources, BAL monitoring studies could greatly improve our knowledge of the composition and evolving structure of QSO outflows, and the physical mechanisms that launch and accelerate outflows. In practice, however, BAL variability studies have been limited by several constraints. Because BAL absorption is most commonly studied in the \\ion{C}{4} $\\lambda$1549 line, BAL observations must be obtained either in the UV or of sources at sufficiently high redshift to place \\ion{C}{4} BALs into a bandpass observable from ground-based telescopes. Because most known BAL QSOs are at $z \\ga 1.5$, it can be difficult to study variation over extended (rest-frame) time scales. With modern AGN surveys such as the Sloan Digital Sky Survey \\citep[SDSS;][]{y+00}, it is now possible to identify thousands of QSOs with observed BAL absorption \\citep[e.g.,][]{thrrsvkafbkn06, gjbhswasvgfy09}. Yet it is necessary to have spectra obtained $\\gtrsim 10$~yr apart in order to span several rest-frame years. In this work, we examine the variation of BAL QSOs observed over rest-frame time scales of up to 7~yr. In order to construct a reasonably representative sample, we have selected all BAL QSOs at redshifts $z \\gtrsim 2.1$ from the catalog of \\citet{hf03} that are visible to the Hobby-Eberly Telescope \\citep[HET;][]{rabbcfggghkklmrrssss98}. Our sample therefore inherits the selection properties of the Large Bright Quasar Survey \\citep[LBQS;][]{hfc95}. We re-observed each of the 14~QSOs matching our selection criteria using the HET Low Resolution Spectrograph \\citep[LRS;][]{hnmtcm98} in order to obtain spectra with adequate signal-to-noise ratio (S/N) and spectral resolution for comparison to the LBQS epoch. Additionally, many of the sources in our study have spectra available from intermediate epochs, obtained either as part of the follow-up study by \\citet{wmfh91}, or during the course of the SDSS. We will refer to the former set of spectra as the ``Palomar'' BAL spectra. The LBQS spectra in our current study and that of \\citet[][hereafter G08]{gbsg08G08} were obtained in the years 1986--1989. With the addition of the Palomar and SDSS spectra, we can sample time scales of months to several years for individual sources. In general, BALs vary in complex ways over multi-year time scales (G08). While studies of individual BAL QSOs are needed to understand how BALs appear \\citep[e.g.,][]{hkrph08, lhcg09}, disappear \\citep[e.g.,][]{jsbbchl01}, and accelerate \\citep[e.g.,][]{hsher07}, it is also important to determine the general characteristics of BAL outflow evolution. In the current work, we take another step in that direction by comparing \\ion{C}{4} BAL absorption profiles in up to four epochs spanning a range of time scales from months to 7~yr. Many of our sources are at sufficiently high redshift to shift the \\ion{Si}{4} absorption region into the optical bandpass, enabling a study of the simultaneous variation of absorption from two different ions (\\ion{Si}{4} and \\ion{C}{4}). \\subsection{Conventions\\label{conventionsSec}} Because conventions vary among studies, we briefly describe the terminology used in this work. We define outflow velocities to be negative (with respect to QSO emission rest frames). Positive velocities would indicate features that are at {\\it longer} wavelengths than the wavelength corresponding to (rest-frame) zero velocity. However, we use the terms ``greater'' and ``smaller'' velocities to refer to the {\\it magnitude} of the velocity, so that an outflow speed of \\mbox{--10,000~km~s$^{-1}$} is ``greater'' than a speed of \\mbox{--5000~km~s$^{-1}$}. We consider absorption equivalent widths (EWs) to be negative, so that a negative change in $\\Delta$EW between two epochs corresponds to a strengthening BAL. Errors in EW are calculated by propagating standard errors, with an estimate of 10\\% error in the continuum. Although we often quote the averaged doublet wavelengths to identify a spectral line, we perform absorption outflow calculations using the wavelength of the red component of the doublet. In effect, we use absorption rest wavelengths of 1550.77~\\AA\\ for \\ion{C}{4} and 1402.77~\\AA\\ for \\ion{Si}{4}. This convention was adopted because the onset of a strong absorption trough from outflowing material would be associated with the red side of the doublet. Unless otherwise noted, wavelengths in this work refer to rest-frame values. We use vacuum wavelengths for all epochs. Sources identifiers that are formatted like ``0009+0219'' refer to LBQS source designations using B1950.0 coordinates. Throughout, we use a cosmology in which $H_0 = 70$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_M = 0.3,$ and $\\Omega_{\\Lambda} = 0.7$. ", "conclusions": "} Here, we briefly summarize results in our study. \\begin{enumerate} \\item{Multi-year BAL variability is complex; there is no clear pattern of \\ion{C}{4} absorption variation over \\mbox{5--7~yr} in our sample of 14 BAL QSOs at $z > 2.1$.} \\item{The scatter in $\\Delta$EW increases with time to $\\sim$1800~km~s$^{-1}$ over 5.8~yr. We find no evidence for asymmetry in the distribution of $\\Delta$EW about 0~km~s$^{-1}$ that would indicate that BALs strengthen and weaken on dramatically different time scales.} \\item{In one source, 1235+1453, we find evidence for a variable, blue continuum component that is relatively unabsorbed by the BAL outflow.} \\item{Although individual BALs may strengthen or weaken monotonically, BALs do not generally evolve monotonically over 3--7~yr. Variation on multi-month time scales does not predict variation over years in an obvious way.} \\item{Based on the fact that we saw at most 1~BAL (in 0022+0150) formally leaving our augmented sample of 19~strong BALs, we estimate that at least stronger BAL features have a lifetime $\\gtrsim$30~yr. The absorption in 0022+0150 did not fully disappear; instead, the absorption weakened to mini-BAL status.} \\item{The evolution of substructures in \\ion{Si}{4} and \\ion{C}{4} BAL trough shapes appears to be related in at least two cases in our sample.} \\item{While $\\Delta$EW is not strongly correlated over 5.8~yr for \\ion{Si}{4} and \\ion{C}{4} BAL trough evolution, there does appear to be a strong correlation in the fractional change in EW for \\ion{Si}{4} and \\ion{C}{4} troughs. This result depends on the removal of one outlier, 1235+1453, which (as we discuss) has unusual properties in several respects.} \\item{Given their apparently long lifetimes and high velocities, BAL outflows can cover large radial and azimuthal distances relative to the accretion disk center. The stability of many BAL troughs (in whole or in specific velocity ranges) over years is an important consideration for models in which BAL outflows are associated with rotating accretion disks.} \\item{We cross-calibrated the wavelengths of spectra between epochs using narrower absorption and emission features and identified two cases (0022+0150 and 1240+1607) of \\ion{C}{4} BALs that showed evidence for possible changes in velocity structure of a magnitude $\\sim$1000~km~s$^{-1}$ over $\\approx$6~yr. Follow-up is needed for these and other candidate sources to determine the nature of possible absorber acceleration.} \\end{enumerate} The current work and previous studies of BAL variability in QSO samples (e.g., \\citealt{b93, lwbhsyvb07}; G08) demonstrate that there is still much to learn about the structure, evolution, and influence of BAL outflows, as well as the physics that launches and shapes them. Improved samples of BAL spectra are clearly needed to examine the nature of BAL variation and acceleration, the ionization structure and evolution of BAL outflows, links with other absorption phenomena such as mini-BALs, and the long-term evolution of BAL outflows. Given the practical issues involved with observing UV spectra over long time frames, we suggest some approaches that would be particularly advantageous for the ongoing study of high-velocity QSO outflows. Firstly, larger samples are needed. The SDSS archive will provide an excellent baseline for studies of BALs at higher resolution and S/N, and the SDSS-III surveys will provide an additional epoch of flux-calibrated spectroscopy for many SDSS BAL QSOs. Secondly, intensive monitoring of a set of variable BALs, ideally at higher resolution, could resolve interesting questions about BAL structure and how BALs evolve over shorter time scales. Among other things, spectra from such a campaign could reveal substructures in the outflow that vary in concert (in the BAL of a single ion or from multiple ions), and would also enable a search for rapid photoionization-driven changes in the outflow. Candidates for BAL acceleration, such as 0022+0150, 1240+1607, and SDSS~J$024221.87+004912.6$ \\citep{hsher07} should also be re-observed over months to years to test whether acceleration has continued. Ideally, we would like to determine a ``structure function for BAL variation'' that would be a useful input to models of quasar outflows. The computation of structure functions and other detailed variability metrics will become possible as large-scale spectroscopic surveys such as SDSS-III re-observe the large samples of BAL quasars identified in earlier SDSS surveys." }, "1003/1003.4210_arXiv.txt": { "abstract": "We produce a 10-day series of simulated Doppler images at a 15-minute cadence that reproduces the spatial and temporal characteristics seen in the SOHO/MDI Doppler data. Our simulated data contains a spectrum of cellular flows with but two necessary components --- a granule component that peaks at wavenumbers of about 4000 and a supergranule component that peaks at wavenumbers of about 110. We include the advection of these cellular components by a differential rotation profile that depends on latitude and wavenumber (depth). We further mimic the evolution of the cellular pattern by introducing random variations to the amplitudes and phases of the spectral components at rates that reproduce the level of cross-correlation as a function of time and latitude. Our simulated data do not include any wave-like characteristics for the supergranules yet can accurately reproduce the rotation characteristics previously attributed to wave-like characteristics. ", "introduction": "Supergranules are cellular flow structures observed in the solar photosphere with typical diameters of about 30 Mm and lifetimes of about one day. They cover the entire surface of the Sun and are intimately involved with the structure and evolution of the magnetic field in the photosphere. The magnetic structures of the chromospheric network form at the boundaries of these cells and magnetic elements are shuffled about the surface as the cells evolve. The diffusion of magnetic elements by the evolving supergranules has long been associated with the evolution of the Sun's magnetic field [\\cite{Leighton64}]. Supergranules were discovered by \\cite{Hart54}. While these cellular flows were quickly identified as convective features [\\cite{LeightonNoyesSimon62}] the difficulty of detecting any associated thermal features consistent with that dentification (i.e. hot cell centers) has made this identification somewhat problematic [\\cite{Worden75}]. The rotation of the supergranules has added further mystery to their nature. \\citet{Duvall80} cross-correlated the equatorial Doppler velocity patterns and found that the supergranules rotated more rapidly than the plasma at the photosphere and that even faster rotation rates were obtained when longer (24-hour vs. 8-hour) time intervals were used. He attributed this behavior to a surface shear layer [proposed by \\cite{FoukalJokipii75} and \\cite{Foukal79} and modeled by \\cite{GilmanFoukal79}] in which larger, longer-lived, cells extend deeper into the more rapidly rotating layers. \\cite{SnodgrassUlrich90} used data from Mount Wilson Observatory to find the rotation rate at different latitudes and noted that the rotation rates for the Doppler pattern were some 4\\% faster than the spectroscopic rate and, more interestingly, some 2\\% faster than the magnetic features and sunspots. More recently \\cite{BeckSchou00} used a 2D Fourier transform method to find that the Doppler pattern rotates more rapidly than the shear layer itself and that larger features do rotate more rapidly than the smaller features. They suggested that supergranules have wave-like characteristics with a preference for prograde propagation. In a previous paper [\\cite{Hathaway_etal06}] we showed that this ``super-rotation'' of the Doppler pattern could be attributed to projection effects associated with the Doppler signal itself. As the velocity pattern rotates across the field of view its line-of-sight component is modulated in a way that essentially adds another half wave and gives a higher rotation rate that is a function of wavenumber. In that paper we took a fixed velocity pattern (which had spatial characteristics that matched the SOHO/MDI data) and rotated it rigidly to show this ``super-rotation'' effect. While this indicated that this Doppler projection effect should be accounted for, the fixed pattern could not account for all the variations reported by \\cite{BeckSchou00}. Furthermore, when \\cite{Schou03} ``divided-out'' the line-of-sight modulation he still saw prograde and retrograde moving components. In this paper we report on our analyses of simulated data in which the supergranules are advected by a differential rotation that varies with both latitude and depth. The data is designed to faithfully mimic the SOHO/MDI data that was analyzed in \\cite{BeckSchou00} and \\cite{Schou03} and the analyses are reproductions of those done in earlier studies. ", "conclusions": "We have produced simulated data in which the cellular structures (supergranules) are advected by differential rotation and evolve by uncorrelated random changes. When we compare results from analyses of this data with those from analyses of the MDI data we find that the simulated data exhibits the same characteristics as the MDI data --- the visual structures, the power spectra, the rotation characteristics, and the evolution rates all match. While some of these characteristics have been attributed to wave-like properties [c.f. \\cite{BeckSchou00} and \\cite{Schou03}] our simulated data is simply advected by a zonal flow (differential rotation) with speeds that never exceed those determined from helioseismology [\\cite{Schou_etal98}]. The differential rotation we impose does, however, have a dependence on wavenumber $\\ell$. If we assume that the rotation rate of cells with diameters, $D = 2 \\pi R_\\odot/\\ell$, reflects the rotation rate at a depth, $d = D/2$, then the surface shear layer indicated by our differential rotation has a thickness of about 20 Mm --- somewhat thinner than the 30 Mm suggested by helioseismic inversions [\\cite{Schou_etal98}]." }, "1003/1003.1745_arXiv.txt": { "abstract": "In these lectures I focus on early universe models which can explain the currently observed structure on large scales. I begin with a survey of inflationary cosmology, the current paradigm for understanding the origin of the universe as we observe it today. I will discuss some progress and problems in inflationary cosmology before moving on to a description of two alternative scenarios - the Matter Bounce and String Gas Cosmology. All early universe models connect to observations via the evolution of cosmological perturbations - a topic which will be discussed in detail in these lectures. ", "introduction": "\\subsection{Overview} Observational cosmology is currently in its ``Golden Age\". A wealth of new observational results are being uncovered. The cosmic microwave background (CMB) has been measured to high precision, the distribution of visible and of dark matter is being mapped out to greater and greater depths. New observational windows to probe the structure of the universe are opening up. To explain these observational results it is necessary to consider processes which happened in the very early universe. The inflationary scenario \\cite{Guth} (see also \\cite{Brout, Starob1, Sato}) is the current paradigm for the evolution of the very early universe. Inflation can explain some of the puzzles which the previous paradigm - Standard Big Bang cosmology - could not address. More importantly, however, inflationary cosmology gave rise to the first explanation for the origin of inhomogeneities in the universe based on causal physics \\cite{Mukh} (see also \\cite{Press,Sato,Starob2})). The theory of structure formation in inflationary cosmology predicted the detailed shape of the angular power spectrum of CMB anisotropies, a prediction which was verified many years later observationally \\cite{WMAP}. In spite of this phenomenological success, inflationary cosmology is not without its conceptual problems. These problems motivate the search for alternative proposals for the evolution of the early universe and for the generation of structure. These alternatives must be consistent with the current observations, and they must make predictions with which they can be observationally distinguished from inflationary cosmology. In these lectures I will discuss two alternative scenarios, the ``Matter Bounce\" (see e.g. \\cite{CosPA08}) and ``String Gas Cosmology\" \\cite{RHBSGrev}. I begin these lectures with a brief survey of the Standard Big Bang (SBB) model and its problems. Possibly the most important problem is the absence of a structure formation scenario based on causal physics. I will then introduce inflationary cosmology and the two alternatives and show how they address the conceptual problems of the SBB model, in particular how structure is formed in these models. Since the information about the early universe is encoded in the spectrum of cosmological fluctuations about the expanding background, it is these fluctuations which provide a link between the physics in the very early universe and current cosmological observations. Section 2 is devoted to an overview of the theory of linearized cosmological perturbations, the main technical tool in modern cosmology. The analysis of this section is applicable to all early universe theories. In Section 3 I give an overview of inflationary cosmology, focusing on the basic principles and emphasizing recent progress and conceptual problems. The conceptual problems of inflationary cosmology motivate the search for alternatives. In Section 4 the ``Matter Bounce\" alternative is discussed, in Section 5 the ``String Gas\" alternative. \\subsection{Standard Big Bang Model} Standard Big Bang cosmology is based on three principles. The first is the assumption that space is homogeneous and isotropic on large scales. The second is the assumption that the dynamics of space-time is described by the Einstein field equations. The third basis of the theory is that matter can be described as a superposition of two classical perfect fluids, a radiation fluid with relativistic equation of state and pressure-less (cold) matter. The first principle of Standard Cosmology implies that the metric takes the following most general form for a homogeneous and isotropic four-dimensional universe \\be \\label{metric} ds^2 \\, = \\, dt^2 - a(t)^2 d{\\bf x}^2 \\, , \\ee where $t$ is physical time, ${\\bf x}$ denote the three co-moving spatial coordinates (points at rest in an expanding space have constant co-moving coordinates), and the scale factor $a(t)$ is proportional to the radiius of space. For simplicity we have assumed that the universe is spatially flat. The expansion rate $H(t)$ of the universe is given by \\be H(t) \\, = \\, \\frac{{\\dot a}}{a} \\, , \\ee where the overdot represents the derivative with respect to time. The cosmological redshift $z(t)$ at time $t$ yields the amount of expansion which space has undergone between the time $t$ and the present time $t_0$: \\be z(t) + 1 \\, = \\, \\frac{a(t_0)}{a(t)} \\, . \\ee In the case of a homogeneous and isotropic metric (\\ref{metric}), the Einstein equations which describe how matter induces curvature of space-time reduce to a set of ordinary differential equations, the Friedmann equations. Written for simplicity in the case of no cosmological constant and no spatial curvature they are \\be H^2 \\, = \\, \\frac{8 \\pi G}{3} \\rho \\, , \\ee where $\\rho$ is the energy density of matter and $G$ is Newton's gravitational constant, and \\be \\frac{\\ddot{a}}{a} \\, = \\, - \\frac{4 \\pi G}{3} \\bigl( \\rho + 3 p \\bigr) \\, \\ee or equivalently \\be {\\dot{\\rho}} \\, = \\, - 3 H \\bigl( \\rho + p \\bigr) \\, , \\ee where $p$ is the pressure density of matter. The third principle of Standard Cosmology says that the energy density and pressure are the sums of the contributions from cold matter (symbols with subscript ``m\") and radiation (subscripts ``r\"): \\bea \\rho \\, &=& \\, \\rho_m + \\rho_r \\, \\\\ p \\, &=& \\, p_m + p_r \\, \\nonumber \\eea with $p_m = 0$ and $p_r = 1/3 \\rho_r$. The first principle of Standard Cosmology, the homogeneity and isotropy of space on large scales, was initially introduced soleley to simplify the mathematics. However, it is now very well confirmed observationally by the near isotropy of the cosmic microwave background and by the convergence to homogeneity in the distribution of galaxies as the length scale on which the universe is probed increases. Standard Cosmology rests on three observational pillars. Firstly, it explains Hubble's redshift-distance relationship. Secondly, and most importantly for the theme of these lectures, it predicted the existence and black body nature of the cosmic microwave background (CMB). The argument is as follows. If we consider regular matter and go back in time, the temperature of matter increases. Eventually, it exceeds the ionization temperature of atoms. Before that time, matter was a plasma, and space was permeated by a thermal gas of photons. A gas of photons will thus remain at the current time. The photons last scatter at a time $t_{rec}$, the ``time of recombination\" which occurs at a redshift of about $10^3$. After that, the gas of photons remains in a thermal distribution with a temperature $T$ which redshifts as the universe expands, i.e. $T \\sim a^{-1}$. The CMB is this remnant gas of photons from the early universe. The precision measurement of the black body nature of the CMB \\cite{CMBblack} can be viewed as the beginning of the ``Golden Age\" of observational cosmology. The third observational pillar of Standard Cosmology is the good agreement between the predicted abundances of light elements and the observed ones. This agreement tells us that Standard Cosmology is a very good description of how the universe evolved back to the time of nucleosynthesis. Any modifications to the time evolution of Standard Cosmology must take place before then. At the present time $t_0$ matter is dominated by the cold pressureless component. Thus, as we go back in time, the universe is (except close to the present time when it appears that a residual cosmological constant is beginning to dominate the cosmological dynamics) initially in a matter-dominated phase during which $a(t) \\sim t^{2/3}$. Since the energy density in cold matter scales as $a^{-3}$ whereas that of relativistic radiation scales as $a^{-4}$, there will be a time before which the energy density in radiation was larger than that of cold matter. The time when the two energy densities are equal is $t_{eq}$ and corresponds to a redshift of around $z = 10^4$. For $t < t_{eq}$ the universe is radiation-dominated and $a(t) \\sim t^{1/2}$. As we go back into the past the density and temperature increase without bound and a singularity is reached at which energy density, temperature and curvature are all infinite. This is the ``Big Bang\" (see Fig. \\ref{fig1} for a sketch of the temperature/time relationship in Standard Cosmology). \\begin{figure} \\includegraphics[height=6cm]{standard.eps} \\caption{A sketch of the temperature (vertical axis) - time (horizontal axis) relation in Standard Cosmology. As the beginning of time is approached, the temperature diverges.} \\label{fig1} \\end{figure} \\subsection{Problems of the Standard Big Bang} Obviously, the assumptions on which Standard Cosmology is based break down long before the singularity is reached. This is the ``singularity problem\" of the model, a problem which also arises in inflationary cosmology. It is wrong to say that Standard Cosmology predicts a Big Bang. Instead, one should say that Standard Cosmology is incomplete and one does not understand the earliest moments of the evolution of the universe. There are, however, more ``practical\" ``problems\" of the Standard Big Bang model - problems in the sense that the model is unable to explain certain key features of the observed universe. The first such problem is the ``horizon problem\" \\cite{Guth}: within Standard Cosmology there is no possible explanation for the observed homogeneity and isotropy of the CMB. Let us consider photons reaching us from opposite angles in the sky. As sketched in Figure \\ref{dfig1}, the source points for these photons on the last scattering surface are separated by a distance greater than the horizon at that time. Hence, there is no possible causal mechanism which can relate the temperatures at the two points. \\begin{figure} \\includegraphics[height=6cm]{dfig1.eps} \\caption{Sketch illustrating the horizon problem of Standard Cosmology: our past horizon at $t_{rec}$ is larger than the causal horizon (forward light cone) at that time, thus making it impossible to causally explain the observed isotropy of the CMB.} \\label{dfig1} \\end{figure} A related problem is the ``size problem\": if the spatial sections of the universe are finite, then the only length scale available at the Planck time $t_{pl}$ is the Planck length. However, it we extrapolate the size of our currently observed horizon back to when the temperature was equal to the Planck temperature, then the corresponding wavelength is larger than $1 {\\rm{\\mu m}}$, many orders of magnitude larger than the Planck length. Standard Cosmology offers no explanation for these initial conditions. The third mystery of Standard Cosmology concerns the observed degree of spatial flatness of the universe. At the current time the observed energy density equals the ``critical\" energy density $\\rho_c$ of a spatially flat universe to within $10\\%$ or better. However, in Standard Cosmology $\\rho = \\rho_c$ is an unstable critical point in an expanding universe. As the universe expands, the relative difference between $\\rho$ and $\\rho_c$ increases. This can be seen by taking the Friedmann equation in the presence of spatial curvature \\be \\label{FRW11} H^2 + \\epsilon T^2 \\, = \\, \\frac{8 \\pi G}{3} \\rho \\, \\ee where \\be \\epsilon \\, = \\, \\frac{k}{(a T)^2} \\, \\ee $T$ being the temperature and $k$ the curvature constant which is $k = \\pm 1$ or $k = 0$ for closed, open or flat spatial sections, and comparing (\\ref{FRW11}) with the corresponding equation in a spatially flat universe ($k = 0$ and $\\rho_c$ replacing $\\rho$. If entropy is conserved (as it is in Standard Cosmology) then $\\epsilon$ is constant and we obtain \\be \\label{flateq} \\frac{\\rho - \\rho_c}{\\rho_c} \\, = \\, \\frac{3}{8 \\pi G} \\frac{\\epsilon T^2}{\\rho_c} \\, \\sim \\, T^{-2} \\, . \\ee Hence, to explain the currently observed degree of spatial flatness, the initial spatial curvature had to have been tuned to a very high accuracy. This is the ``flatness problem\". We observe highly non-random correlations in the distribution of galaxies in the universe. The only force which can act on the relevant distances is gravity. Gravity is a weak force, and therefore the seed fluctuations which develop into the observed structures had to have been non-randomly distributed at the time $t_{eq}$, the time when gravitational clustering begins (see Section 2). However, as illustrated in Figure \\ref{dfig2}, the physical length corresponding to the fluctuations which we observe today on the largest scales (they have constant comoving scale) is larger than the horizon at that time. Hence, there can be no causal generation mechanism for these perturbations \\footnote{This is the usual textbook argument. The student is invited to find (at least) two flaws in this argument.}. This is the ``fluctuation problem\" of Standard Cosmology. \\begin{figure} \\includegraphics[height=6cm]{dfig2.eps} \\caption{Sketch illustrating the formation of structure problem of Standard Cosmology: the physical wavelength of fluctuations on fixed comoving scales which correspond to the large-scale structures observed in the universe is larger than the horizon at the time $t_{eq}$ of equal matter and radiation, the time when matter fluctuations begin to grow. Hence, it is impossible to explain the origin of non-trivial correlations of the seeds for the fluctuations which had to have been present at that time.} \\label{dfig2} \\end{figure} There are also more conceptual problems: Standard Cosmology is based on treating matter as a set of perfect fluids. However, we know that at high energies and temperatures a classical description of matter breaks down. Thus, Standard Cosmology must break down at sufficiently high energies. It does not contain the adequate matter physics to describe the very early universe. Similarly, the singularity of space-time which Standard Cosmology contains corresponds to a breakdown of the assumptions on which General Relativity is based. This is the ``singularity problem\" of Standard Cosmology. All of the early universe scenarios which I will discuss in the following provide solutions to the formation of structure problem. They can successfully explain the wealth of observational data on the distribution of matter in the universe and on the CMB anisotropies. The rest of these lectures will focus on this point. However, the reader should also ask under which circumstances the scenarios to be discussed below address the other problems mentioned above. \\subsection{Inflation as a Solution} The idea of inflationary cosmology is to add a period to the evolution of the very early universe during which the scale factor undergoes accelerated expansion - most often nearly exponential growth. The time line of inflationary cosmology is sketched in Figure \\ref{timeline}. The time $t_i$ is the beginning of the inflationary period, and $t_R$ is its end (the meaning of the subscript $R$ will become clear later). Although inflation is usually associated with physics at very high energy scales, e.g. $E \\sim 10^{16} {\\rm Gev}$, all that is required from the initial basic considerations is that inflation ends before the time of nucleosynthesis. \\begin{figure} \\includegraphics[height=2.2cm]{timeline.eps} \\caption{A sketch showing the time line of inflationary cosmology. The period of accelerated expansion begins at time $t_i$ and end at $t_R$. The time evolution after $t_R$ corresponds to what happens in Standard Cosmology.} \\label{timeline} \\end{figure} During the period of inflation, the density of any pre-existing particles is red-shifted. Hence, if inflation is to be viable, it must contain a mechanism to heat the universe at $t_R$, a ``reheating\" mechanism - hence the subscript $R$ on $t_R$. This mechanism must involve dramatic entropy generation. It is this non-adiabatic evolution which leads to a solution of the flatness problem, as the reader can verify by inspecting equation (\\ref{flateq}) and allowing for entropy generation at the time $t_R$. A space-time sketch of inflationary cosmology is given in Figure \\ref{infl1} . The vertical axis is time, the horizontal axis corresponds to physical distance. Three different distance scales are shown. The solid line labelled by $k$ is the physical length corresponding to a fixed comoving perturbation. The second solid line (blue) is the Hubble radius \\be l_H(t) \\, \\equiv \\, H^{-1}(t) \\, . \\ee The Hubble radius separates scales where microphysics dominates over gravity (sub-Hubble scales) from ones on which the effects of microphysics are negligible (super-Hubble scales) \\footnote{This statement will be demonstrated later in these lectures.}. Hence, a necessary requirement for a causal theory of structure formation is that scales we observe today originate at sub-Hubble lengths in the early universe. The third length is the ``horizon\", the forward light cone of our position at the Big Bang. The horizon is the causality limit. Note that because of the exponential expansion of space during inflation, the horizon is exponentially larger than the Hubble radius. It is important not to confuse these two scales. Hubble radius and horizon are the same in Standard Cosmology, but in all three early universe scenarios which will be discussed in the following they are different. \\begin{figure} \\includegraphics[height=9cm]{infl1.eps} \\caption{Space-time sketch of inflationary cosmology. The vertical axis is time, the horizontal axis corresponds to physical distance. The solid line labelled $k$ is the physical length of a fixed comoving fluctuation scale. The role of the Hubble radius and the horizon are discussed in the text.} \\label{infl1} \\end{figure} {F}rom Fig. \\ref{infl1} it is clear that provided that the period of inflation is sufficiently long, all scales which are currently observed originate as sub-Hubble scales at the beginning of the inflationary phase. Thus, in inflationary cosmology it is possible to have a causal generation mechanism of fluctuations \\cite{Mukh,Press,Sato}. Since matter pre-existing at $t_i$ is redshifted away, we are left with a matter vacuum. The inflationary universe scenario of structure formation is based on the hypothesis that all current structure originated as quantum vacuum fluctuations. From Figure \\ref{infl1} it is also clear that the horizon problem of standard cosmology can be solved provided that the period of inflation lasts sufficiently long. The reader should convince him/herself that the required period of inflation is about $50 H^{-1}$ if inflation takes place at an enegy scale of about $10^{16} {\\rm GeV}$. Inflation thus solve both the horizon and the structure formation problems. To obtain exponential expansion of space in the context of Einstein gravity the energy density must be constant. Thus, during inflation the total energy and size of the universe both increase exponentially. In this way, inflation can solve the size and entropy problems of Standard Cosmology. To summarize the main point concerning the generation of cosmological fluctuations (the main theme of these lectures) in inflationary cosmology: the first crucial criterium which must be satisfied in order to have a successful theory of structure formation is that fluctuation scales originate inside the Hubble radius. In inflationary cosmology it is the accelerated expansion of space during the inflationary phase which provides this success. In the following we will emphasize what is responsible for the corresponding success in the two alternative scenarios which we will discuss. \\subsection{Matter Bounce as a Solution} The first alternative to cosmological inflation as a theory of structure formation is the ``matter bounce\" , an alternative which is not yet well appreciated (for an overview the reader is referred to \\cite{CosPA08}). The scenario is based on a cosmological background in which the scale factor $a(t)$ bounces in a non-singular manner. {F}igure \\ref{bounce} shows a space-time sketch of such a bouncing cosmology. Without loss of generality we can adjust the time axis such that the bounce point (minimal value of the scale factor) occurs at $t = 0$. There are three phases in such a non-singular bounce: the initial contracting phase during which the Hubble radius is decreasing linearly in $|t|$, a bounce phase during which a transition from contraction to expansion takes place, and thirdly the usual expanding phase of Standard Cosmology. There is no prolonged inflationary phase after the bounce, nor is there a time-symmetric deflationary contracting period before the bounce point. As is obvious from the Figure, scales which we observe today started out early in the contracting phase at sub-Hubble lengths. The matter bounce scenario assumes that the contracting phase is matter-dominated at the times when scales we observe today exit the Hubble radius. A model in which the contracting phase is the time reverse of our current expanding phase would obey this condition. The assumption of an initial matter-dominated phase will be seen later in these lectures to be important if we want to obtain a scale-invariant spectrum of cosmological perturbations from initial vacuum fluctuations \\cite{Wands,FB2,Wands2}. \\begin{figure}[htbp] \\includegraphics[height=9cm]{bounce.eps} \\caption{Space-time sketch in the matter bounce scenario. The vertical axis is conformal time $\\eta$, the horizontal axis denotes a co-moving space coordinate. Also, ${\\cal H}^{-1}$ denotes the co-moving Hubble radius.} \\label{bounce} \\end{figure} Let us make a first comparison with inflation. A non-deflationary contracting phase replaces the accelerated expanding phase as a mechanism to bring fixed comoving scales within the Hubble radius as we go back in time, allowing us to consider the possibility of a causal generation mechanism of fluctuations. Starting with vacuum fluctuations, a matter-dominated contracting phase is required in order to obtain a scale-invariant spectrum. This corresponds to the requirement in inflationary cosmology that the accelerated expansion be nearly exponential. With Einstein gravity and matter satisfying the usual energy conditions it is not possible to obtain a non-singular bounce. However, as mentioned before, it is unreasonable to expect that Einstein gravity will provide a good description of the physics near the bounce. There are a large number of ways to obtain a non-singular bouncing cosmology. To mention but a few, it has been shown that a bouncing cosmology results naturally from the special ghost-free higher derivative gravity Lagrangian introduced in \\cite{Biswas1}. Bounces also arise in the higher-derivative non-singular universe construction of \\cite{MBS}, in ``mirage cosmology\" (see e.g. \\cite{BFS}) which is the cosmology on a brane moving through a curved higher-dimensional bulk space-time (the time dependence for a brane observer is induced by the motion through the bulk), or - within the context of Einstein gravity - by making use of ``quintom matter\" (matter consisting of two components, one with regular kinetic term in the action, the other one with opposite sign kinetic action) \\cite{Cai1}. For an in-depth review of ways of obtaining bouncing cosmologies see \\cite{Novello}. Very recently, it has been shown \\cite{HLbounce} that a bouncing cosmology can easily emerge from Horava-Lifshitz gravity \\cite{Horava}, a new approach to quantizing gravity. In the matter bounce scenario the universe begins cold and therefore large. Thus, the size problem of Standard Cosmology does not arise. As is obvious from Figure \\ref{bounce}, there is no horizon problem for the matter bounce scenario as long as the contracting period is long (to be specific, of similar duration as the post-bounce expanding phase until the present time). By the same argument, it is possible to have a causal mechanism for generating the primordial cosmological perturbations which evolve into the structures we observe today. Specifically, as will be discussed in the section of the matter bounce scenario, if the fluctuations originate as vacuum perturbations on sub-Hubble scales in the contracting phase, then the resulting spectrum at late times for scales exiting the Hubble radius in the matter-dominated phase of contraction is scale-invariant \\cite{Wands,FB2,Wands2}. The propagation of infrared (IR) fluctuations through the non-singular bounce was analyzed in the case of the higher derivative gravity model of \\cite{Biswas1} in \\cite{ABB}, in mirage cosmology in \\cite{BFS}, in the case of the quintom bounce in \\cite{Cai2,LWbounce} and for a Horava-Lifshitz bounce in \\cite{HLbounce2}. The result of these studies is that the scale-invariance of the spectrum before the bounce goes persists after the bounce as long as the time period of the bounce phase is short compared to the wavelengths of the modes being considered. Note that if the fluctuations have a thermal origin, then the condition on the background cosmology to yield scale-invariance of the spectrum of fluctuations is different \\cite{Thermalflucts}. \\subsection{String Gas Cosmology as a Solution} String gas cosmology \\cite{BV} (see also \\cite{Perlt}, and see \\cite{RHBSGrev} for a comprehensive review) is a toy model of cosmology which results from coupling a gas of fundamental strings to a background space-time metric. It is assumed that the spatial sections are compact. It is argued that the universe starts in a quasi-static phase during which the temperature of the string gas hovers at the Hagedorn value \\cite{Hagedorn}, the maximal temperature of a gas of closed strings in thermal equilibrium. The string gas in this early phase is dominated by strings winding the compact spatial sections. The annihilation of winding strings will produce string loops and lead to a transition from the early quasi-static phase to the radiation phase of Standard Cosmology. Fig. \\ref{timeevol} shows a sketch of the evolution of the scale factor in string gas cosmology. \\begin{figure} \\includegraphics[height=6cm]{timeevol.eps} \\caption{The dynamics of string gas cosmology. The vertical axis represents the scale factor of the universe, the horizontal axis is time. Along the horizontal axis, the approximate equation of state is also indicated. During the Hagedorn phase the pressure is negligible due to the cancellation between the positive pressure of the momentum modes and the negative pressure of the winding modes, after time $t_R$ the equation of state is that of a radiation-dominated universe.} \\label{timeevol} \\end{figure} In Figure \\ref{spacetimenew} we sketch the space-time diagram in string gas cosmology. Since the early Hagedorn phase is quasi-static, the Hubble radius is infinite. For the same reason, the physical wavelength of fluctuations remains constant in this phase. At the end of the Hagedorn phase, the Hubble radius decreases to a microscopic value and makes a transition to its evolution in Standard Cosmology. \\begin{figure} \\includegraphics[height=10cm]{spacetimenew.eps} \\caption{Space-time diagram (sketch) showing the evolution of fixed co-moving scales in string gas cosmology. The vertical axis is time, the horizontal axis is physical distance. The solid curve represents the Einstein frame Hubble radius $H^{-1}$ which shrinks abruptly to a micro-physical scale at $t_R$ and then increases linearly in time for $t > t_R$. Fixed co-moving scales (the dotted lines labeled by $k_1$ and $k_2$) which are currently probed in cosmological observations have wavelengths which are smaller than the Hubble radius before $t_R$. They exit the Hubble radius at times $t_i(k)$ just prior to $t_R$, and propagate with a wavelength larger than the Hubble radius until they reenter the Hubble radius at times $t_f(k)$.} \\label{spacetimenew} \\end{figure} Once again, we see that fluctuations originate on sub-Hubble scales. In string gas cosmology, it is the existence of a quasi-static phase which leads to this result. As will be discussed in the section on string gas cosmology, the source of perturbations in string gas cosmology is thermal: string thermodynamical fluctuations in a compact space with stable winding modes in fact leads to a scale-invariant spectrum \\cite{NBV}. ", "conclusions": "The main messages of these lectures are the following: \\begin{itemize} \\item{} It is possible to explore physics of the very early universe using current cosmological observations. Given the large amount of new data which is expected over the next decade, early universe cosmology will be a vibrant field of research. The information about the very early universe is transferred to the current time mainly via imprints on the spectrum of cosmological fluctuations. \\item{} The theory of cosmological fluctuations is the key tool of modern cosmology. It links current observations with early universe cosmology. The theory can be applied to any background cosmology, not just inflationary cosmology. \\item{} The Hubble radius plays a key role in the evolution of fluctuations. On sub-Hubble scales the microphysical forces dominate, whereas on super-Hubble scales matter forces freeze out and gravity dominates. In order to have a causal mechanism of structure formation, it is therefore crucial that scales of current interest in cosmology originate at very early times inside the Hubble radius, and that they then propagate over an extended period of time on super-Hubble scales. \\item{} Inflationary cosmology provides a mechanism to explain the origin of structure in the universe via causal physics, but it is not the only one. In these lectures we have presented two alternative scenarios, the ``matter bounce\" and ``string gas cosmology\". \\item{} Inflationary cosmology was predictive: it predicted the detailed shape of the angular power spectrum of CMB anisotropies - more than 15 years before the observations. The challenge for alternative scenarios is to identify clean predictions with which the new scenarios can be distinguished from those of inflation. In the case of the matter bounce the prediction we have identified is a particular shape of the bispectrum, in the case of string gas cosmology the key prediction is a slight blue tilt in the spectrum of gravitational waves. \\end{itemize} Inflationary cosmology has been under development for 30 years. The framework is based on Einstein gravity coupled to classical scalar field matter and is well established. However, when applied to early universe cosmology there are important conceptual problems which arise. In contrast, the ``matter bounce\" and ``string gas\" scenarios of structure formation are much more recent. They both {\\it must} involve physics which goes beyond classical field theory coupled to Einstein gravity. To come up with a consistent model is a major challenge which has not yet been satisfactorily met. There is a lot of interesting work ahead of us." }, "1003/1003.4998_arXiv.txt": { "abstract": "The spatial morphology, spectral characteristics, and time variability of ultracompact \\hii\\ regions provide strong constraints on the process of massive star formation. We have performed simulations of the gravitational collapse of rotating molecular cloud cores, including treatments of the propagation of ionizing and non-ionizing radiation. We here present synthetic radio continuum observations of \\hii\\ regions from our collapse simulations, to investigate how well they agree with observation, and what we can learn about how massive star formation proceeds. We find that intermittent shielding by dense filaments in the gravitationally unstable accretion flow around the massive star leads to highly variable \\hii\\ regions that do not grow monotonically, but rather flicker, growing and shrinking repeatedly. This behavior appears able to resolve the well-known lifetime problem. We find that multiple ionizing sources generally form, resulting in groups of ultracompact \\hii\\ regions, consistent with observations. We confirm that our model reproduces the qualitative \\hii\\ region morphologies found in surveys, with generally consistent relative frequencies. We also find that simulated spectral energy distributions (SEDs) from our model are consistent with the range of observed \\hii\\ region SEDs, including both regions showing a normal transition from optically thick to optically thin emission, and those with intermediate spectral slopes. In our models, anomalous slopes are solely produced by inhomogeneities in the \\hii\\ region, with no contribution from dust emission at millimeter or submillimeter wavelengths. We conclude that many observed characteristics of ultracompact \\hii\\ regions appear consistent with massive star formation in fast, gravitationally unstable, accretion flows. ", "introduction": "Ultracompact (UC) \\hii\\ regions have radii $R< 0.1$~pc and high radio surface brightness \\citep{churchwell02}. They have a characteristic distribution of morphologies \\citep{woodchurch89,kurtzetal94}, and usually are found associated with one another and with compact \\hii\\ regions \\citep[e.g.][]{welchetal87,gaumeclaussen90,mehringeretal93,kimkoo01}. Their spectral energy distributions (SEDs) can reflect a transition from optically thin to optically thick emission, but often show an anomalous intermediate wavelength dependence \\citep{francoetal00,lizano08}. Recent radio continuum observations have suggested that ultracompact \\hii\\ regions can contract, change in shape, or expand anisotropically over intervals of as little as $\\sim$ 10~yr \\citep{francheretal04,rodrigetal07,galvmadetal08}. The observed brightness and size of UC \\hii\\ regions require that they be ionized by massive stars of spectral type earlier than B3. If the regions were to expand at the sound speed of ionized gas, $c_i \\sim 10$~km/s, they would have lifetimes of roughly $10^4$~yr. Less than 1\\% of an OB star's lifetime of a few million years should therefore be spent within such a region, so the same fraction of OB stars should now lie within them. However, surveys find numbers in our Galaxy consistent with over 10\\% of OB stars being surrounded by them \\citep{woodchurch89,depreeetal05}, or equivalently, lifetimes of $\\sim 10^5$~yr if this model is correct. A number of explanations have been proposed for this lifetime problem, including confinement in cloud cores by thermal pressure \\citep{depreeetal95,gsfranco96} or turbulent pressure \\citep{xieetal96}, ram pressure confinement by infall \\citep{yorke86,hollenbachetal94} or bow shocks \\citep{vanburenml90,mlvanburen91,arthurhoare06}, champagne flows \\citep{bodenheimeretal79,gsfranco96,arthurhoare06}, disk evaporation \\citep{hollenbachetal94}, and mass-loaded stellar winds \\citep{dysonetal95,redmanetal96,williamsetal96,lizanoetal96}, but most have been argued to have major flaws \\citep{maclowetal07}. We have modeled accretion on to an ionizing source using three-dimensional simulations \\citep[][hereafter Paper I]{petersetal10}. These calculations suggest that accretion can indeed explain the lifetime problem, but in an unexpected way. \\citet{keto02,keto07} has argued that ultracompact and hypercompact \\hii\\ regions are simply the ionized portion of an accretion flow. However, massive stars require accretion at rates exceeding $10^{-4}$~M$_{\\odot}$~yr$^{-1}$ \\citep{beutheretal02,beltranetal06} to reach their final masses before exhausting their nuclear fuel \\citep{ketoetal06}. The result is gravitational instability during collapse, leading to the formation of dense gas filaments in the rotating, collapsing flow, along with dozens of accompanying stars (Paper I). The strongest sources of ionizing radiation orbit through the dense filaments repeatedly, accreting mass efficiently when they do. The filaments absorb the ionizing radiation locally, though, when this happens, shielding the rest of the \\hii\\ region for long enough for it to recombine. As a result, the size of the observed \\hii\\ region remains independent of the age of the ionizing star until the surrounding secondary star formation cuts off accretion on to the primary and a compact \\hii\\ region begins to grow around it. In this paper we consider in more detail than in Paper~I whether our models of ionization interacting with a gravitationally unstable accretion flow can reproduce the observations of ultracompact \\hii\\ regions summarized at the beginning of this section. We show how \\hii\\ regions fluctuate in size as their central stars pass through density fluctuations, and demonstrate that our models qualitatively reproduce all morphologies observed for ultracompact \\hii\\ regions, even giving general quantitative agreement with the distribution of different morphologies observed by \\citet{woodchurch89} and \\citet{kurtzetal94}. Our models also offer a natural explanation for the observed clustering of ultracompact \\hii\\ regions. We further demonstrate that they reproduce observed SEDs, and provide natural explanations for the anomalous SEDs observed for some ultracompact \\hii\\ regions \\citep{lizano08,beuthetal04,ketoetal08}. In Sect.\\ \\ref{sec:numerics} we describe our methods for modeling ultracompact \\hii\\ regions and simulating observations, while in Sect.\\ \\ref{sec:results} we describe the results of our work relevant for this paper. Finally, in Sect.\\ \\ref{sec:conclusions} we draw conclusions. ", "conclusions": "\\label{sec:conclusions} We have presented synthetic continuum observations of the \\hii\\ regions formed in our collapse simulations of massive star formation described in Paper~I. We find that the \\hii\\ regions are highly variable in size and shape as long as the massive star continues accreting. Dense filaments in the gravitationally unstable accretion flow irregularly shield ionizing radiation from the central star. This shielding effect leads to a large-scale ($\\sim 5000\\,$AU) flickering of the \\hii\\ regions on timescales as short as $\\sim 100\\,$yr. As long as the flickering continues, there is no direct relation between the age of the star and the size of the \\hii\\ region. This result appears able to resolve the \\uchii\\ lifetime problem. Furthermore, we have identified the structures inside the \\hii\\ regions that produce the continuum emission. We find that emission peaks are not necesarily related to the positions of stars. Instead, strong emission is produced by partially ionized gas from the accretion flow that enters the \\hii\\ region. These filaments can be ionized and blown away by thermal pressure. Since the continuum emission is determined solely by the structure of the flow field around the massive star, the appearance of the \\hii\\ region depends strongly on the angle from which it is viewed. For example, we find that the same \\hii\\ region can be classified as shell-like or cometary, depending on the position of the observer. We have evaluated the distribution of apparent \\hii\\ region morphologies. The multiple sink simulation run~B reproduces the observed relative frequencies reported from surveys, including the high fraction of spherical or unresolved regions. The single sink simulation run~A, however, fails to reproduce the morphology statistics. This is because the single star grows so quickly that it is statistically unable to produce a significant fraction of the smallest \\hii\\ regions. The formation of a whole stellar cluster is necessary to get the relative frequencies of the smallest \\hii\\ regions right. This analysis provides strong evidence against models of massive star formation where all or most high-mass stars form in isolation. In addition to the \\hii\\ region morphologies, we can also reproduce the different SEDs characteristic of \\uchii\\ regions. We find that our initial gas mass of $1000\\,M_\\odot$ is too small to produce observable dust emission at VLA or ALMA wavelengths. Nevertheless, ultracompact \\hii\\ regions at different times and places in our models show SEDs with both regular transitions from optically thick to optically thin spectral slopes as well as anomalous scaling with a spectral slope around unity. These slopes are entirely produced by inhomogeneities in the density structure, with no contribution from dust. However, follow-up simulations with more massive initial clumps will be required to completely pin down the role of dust in observed SEDs." }, "1003/1003.5020_arXiv.txt": { "abstract": "Primordial non-Gaussianity is a potentially powerful discriminant of the physical mechanisms that generated the cosmological fluctuations observed today. Any detection of non-Gaussianity would have profound implications for our understanding of cosmic structure formation. In this paper, we review past and current efforts in the search for primordial non-Gaussianity in the large scale structure of the Universe. ", "introduction": "\\label{sec:intro} In generic inflationary models based on the slow roll of a scalar field, primordial curvature perturbations are produced by the inflaton field as it slowly rolls down its potential \\cite{1981JETPL..33..532M,1982PhLB..117..175S, 1982PhLB..115..295H,1982PhRvL..49.1110G}. Most of these scenarios predict a nearly scale-invariant spectrum of adiabatic curvature fluctuations, a relatively small amount of gravity waves and tiny deviations from Gaussianity in the primeval distribution of curvature perturbations \\cite{1987PhLB..197...66A,1992PhRvD..46.4232F,1994ApJ...430..447G}. While the latest measurements of the cosmic microwave background (CMB) anisotropies favor a slightly red power spectrum \\cite{2009ApJS..180..330K}, no significant detection of a $B$-mode or of primordial non-Gaussianity (NG) has thus far been reported from CMB observations. While the presence of a $B$-mode can only be tested with CMB measurements \\cite{1997PhRvL..78.2054S,1997PhRvL..78.2058K}, primordial deviations from Gaussianity can leave a detectable signature in the distribution of CMB anisotropies {\\it and} in the large scale structure (LSS) of the Universe. Until recently, it was widely accepted that measurement of the CMB furnish the best probe of primordial non-Gaussianity \\cite{2000MNRAS.313..141V}. However, these conclusions did not take into account the scale-dependence of the galaxy power spectrum and bispectrum arising for primordial NG of the local $\\floc$ type \\cite{2004PhRvD..69j3513S,2008PhRvD..77l3514D}. These theoretical results, together with rapid developments in observational techniques that will provide large amount of LSS data, will enable us to critically confront predictions of non-gaussian models. In particular, galaxy clustering should provide independent constraints on the magnitude of primordial non-Gaussianity as competitive as those from the CMB and in the long run may even give the best constraints. The purpose of this work is to provide an overview of the search for a primordial non-Gaussian signal in the large scale structure. We will begin by briefly summarizing how non-Gaussianity arises in inflationary models (\\S\\ref{sec:models}). Next, we will discuss the impact of primordial non-Gaussianity on the mass distribution in the low redshift Universe (\\S\\ref{sec:matterng}). The main body of this review is \\S\\ref{sec:lssprobes}, where we describe in detail an number of methods exploiting the abundance and clustering properties of observed tracers of the LSS to constrain the amount of initial non-Gaussianity. We conclude with a discussion of present and forecasted constraints achievable with LSS surveys (\\S\\ref{sec:limits}). ", "conclusions": "" }, "1003/1003.0433_arXiv.txt": { "abstract": "{} {We aim to understand the interplay between non-radial oscillations and stellar magnetic activity and test the feasibility of doing asteroseismology of magnetically active stars. We investigate the active slow rotator EK~Eri which is the likely descendant of an Ap~star.} {We analyze 30 years of photometric time-series data, 3 years of HARPS radial velocity monitoring, and 3 nights of high-cadence HARPS asteroseismic data. We construct a high-S/N HARPS spectrum that we use to determine atmospheric parameters and chemical composition. Spectra observed at different rotation phases are analyzed to search for signs of temperature or abundance variations. An upper limit on the projected rotational velocity is derived from very high-resolution CES spectra.} {We detect oscillations in EK~Eri with a frequency of the maximum power of $\\nu_\\mathrm{max}=320\\pm32$~$\\mu$Hz, and we derive a peak amplitude per radial mode of $\\approx0.15$~m\\,s$^{-1}$, which is a factor of $\\approx 3$ lower than expected. We suggest that the magnetic field may act to suppress low-degree modes. Individual frequencies can not be extracted from the available data. We derive accurate atmospheric parameters, refining our previous analysis, finding $T_\\mathrm{eff} = 5135\\pm80$~K, $\\log g = 3.39\\pm0.12$, and metallicity [M/H] =$+0.02\\pm0.04$. Mass and radius estimates from the seismic analysis are not accurate enough to constrain the position in the HR diagram and the evolutionary state. We confirm that the main light variation is due to cool spots, but that other contributions may need to be taken into account. We tentatively suggest that the rotation period is twice the photometric period, i.e., $P_\\mathrm{rot} = 2P_\\mathrm{phot} = 617.6$~d, and that the star is a dipole-dominated oblique rotator viewed close to equator-on. We conclude from our derived parameters that $v\\sin i < 0.40$~km\\,s$^{-1}$ and we show that the value is too low to be reliably measured. We also link the time series of direct magnetic field measurements available in the literature to our newly derived photometric ephemeris. } {} ", "introduction": "\\label{intro} \\object{EK~Eri} (HR 1362, HD 27536) is a unique case of a slowly rotating ($v\\sin i < 2$~\\kms) G8 giant or sub-giant, which is over-active with respect to its rotation rate and evolutionary state. It is exhibiting brightness variations with a period of more than 300 days, believed to be due to semi-stable star spots being rotated across the projected surface \\citep{strassmeier+1999}. It has been suggested that the associated strong magnetic field is a fossil remnant from its main sequence life as a magnetic Ap star \\citep{stepien1993}. The star has been monitored photometrically since 1978, with first results included in \\citet{strassmeier+1990}. In a subsequent study, \\citet{strassmeier+1999} derived a photometric period of 306.9~d from twenty years of monitoring, but noted that the light curve could be split into two segments with different periods, 311~d (pre-1987) and 294~d (post-1992) respectively, and with a period of relatively small light variations in between, possibly reflecting two distinct magnetic cycles. These photometric data also showed that the star gets redder when it gets fainter, which agrees with cool spots as the cause of the light variation. From high-resolution spectroscopy, \\citet{strassmeier+1999} also determined the fundamental stellar parameters and showed from the radius--$v\\sin i$ constraints that the star must be seen close to equator-on, i.e. $i \\approx 90^\\circ$. \\citet{auriere+2008} published the first direct measurement of the magnetic field of EK~Eri. They find the field to be large scale, rather than a solar-type (small-scale) field, dominated by a poloidal mostly axisymmetric component, resembling a dipole with a strength of $\\approx 270$~G. They also observe some modulation over the rotation period, although their data do not cover the full photometric period. Their results strengthen the interpretation that EK~Eri is a descendant of a slowly rotating magnetic Ap star which is now approaching the giant branch. In \\citet[][hereafter Paper\\,I]{dall+2005} we refined the fundamental parameters of EK~Eri using new HARPS spectra and found radial velocity (RV) variations with peak-to-peak amplitude of $\\approx 100$~\\ms. This variation was shown to correlate extremely well with the calcium H \\& K activity index (\\rhk) as well as with the bisector inverse slope (BIS) of the cross-correlation function. However we found a positive correlation for the BIS rather than the negative correlation expected from spot-induced RV variations \\citep[e.g.,][]{desort+2007}. Such correlations have previously been attributed to fainter stellar companions contributing to the signal, effectively disguising the signal of a planetary companion \\citep[][]{santos+jvc2002,zucker+2004}. In this paper we present the updated results from 30 years of photometric monitoring, as well as from three years of RV monitoring, and three half-nights of high-cadence RV measurements used for an asteroseismic analysis. In Sect.~\\ref{obs} we present our new observations and the data reduction. In Sect.~\\ref{results} we present our results, which we discuss in Sect.~\\ref{discussion}. In Sect.~\\ref{conclusions} we give a summary of our conclusions. ", "conclusions": "\\label{conclusions} We have presented results from an intensive monitoring of the active sub-giant star EK~Eri. We have used photometric data covering 30 years and spectroscopic data probing long-term variation in activity during 3 years and high-cadence radial-velocity monitoring from 3 nights. From the photometry we have refined the rotation period and the ephemeris as listed in Table~\\ref{tab:results}. Also, from 3 half-nights of high-cadence RV monitoring we have detected solar-like oscillations in a late-type spotted sub-giant star for the first time. While oscillations have been detected in a number of late-type giants \\citep[e.g.,][]{hekker+2009,hatzes+zechmeister2008} and for mildly active solar-type stars \\citep[e.g., Pollux:][]{hatzes+zechmeister2007,auriere+2009}, this is the first detection for a sub-giant hosting a strong magnetic field. Unfortunately our data do not allow us to resolve individual frequencies. We measure an amplitude per radial mode of $\\approx 0.15$~\\ms\\ at a position of maximum power $\\nu_\\mathrm{max}=320\\pm32$~$\\mu$Hz. This amplitude is at least a factor of 3 lower than expected and, if confirmed, may mean that the magnetic field has a strong stabilizing effect on the stellar geometry, essentially favoring high-degree oscillation modes in the presence of a near-dipole magnetic field. In that case, the interpretation of the asteroseismic data may become more difficult for this and other sub-giant stars with similar magnetic properties. A longer time series is required in order to obtain quantitative results. For reference purposes, having accurate oscillation data for a non-active star at the position in the HR diagram of EK~Eri would be highly desirable. Unfortunately, neither CoRoT nor Kepler apparently covers this. Based on the roughly sinusoidal shape of the light curve, the likely very high inclination, the field geometry suggested by \\citet{auriere+2008}, and the behavior of the activity indicators as function of RV, we suggest a conceptual model of EK~Eri with two large low-latitude spot covered areas approximately $180^\\circ$ apart on a star viewed equator-on. In this scenario, the rotational period is twice the photometric period, thus $P_\\mathrm{rot} = 2P_\\mathrm{phot} = 617.6$~d. We note however, that a simple two-spot model is not able to account for all the seasonal light variations observed, mostly due to the unknown spot lifetimes, sizes and longitudes. Regardless of the rotation period, the measured values of $v\\sin i$ both from this work and from the literature are inconsistent with the derived radius of the star. Both the radius derived from asteroseismology and from the spectral analysis set strict upper limits on $v\\sin i$ which are lower than previous estimates. We thus conclude that the $v\\sin i$ is too low to be reliably measured with available spectrographs. Based on high-quality HARPS spectra we have derived the atmospheric parameters of EK~Eri to very high precision. The abundance pattern for 17 analysed elements is very similar to the Sun, and we detect no anomalies that could otherwise be attributed to an earlier evolutionary state as a magnetic Ap star. However, in order to argue for or against EK~Eri being a descendant of a magnetic Ap star, stronger constraints on the mass and evolutionary state are needed. Further seismic studies, preferably at varying rotational phases, may deliver such constraints in terms of accurate asteroseismic mass and radius measurements, and are also needed to probe the possible link between solar-like oscillations and the magnetic field." }, "1003/1003.0605_arXiv.txt": { "abstract": "Black holes grow by accreting matter from their surroundings. However, angular momentum provides an efficient natural barrier to accretion and so only the lowest angular momentum material will be available to feed the black holes. The standard sub-grid model for black hole accretion in galaxy formation simulations -- based on the Bondi-Hoyle method -- does not account for the angular momentum of accreting material, and so it is not clear how representative the black hole accretion rate estimated in this way is likely to be. In this paper we introduce a new sub-grid model for black hole accretion that naturally accounts for the angular momentum of accreting material. Both the black hole and its accretion disc are modelled as a composite \\emph{accretion disc particle}. Gas particles are captured by the accretion disc particle if and only if their orbits bring them within its accretion radius $R_{\\rm acc}$, at which point their mass is added to the accretion disc and feeds the black hole on a viscous timescale $t_{\\rm visc}$. The resulting black hole accretion rate $\\dot{M}_{\\rm BH}$ powers the accretion luminosity $L_{\\rm acc} \\propto \\dot{M}_{\\rm BH}$, which drives black hole feedback. Using a series of controlled numerical experiments, we demonstrate that our new accretion disc particle method is more physically self-consistent than the Bondi-Hoyle method. We also discuss the physical implications of the accretion disc particle method for systems with a high degree of rotational support, and we argue that the $M_{\\rm BH}-\\sigma$ relation in these systems should be offset from the relation for classical bulges and ellipticals, as appears to be observed. ", "introduction": "\\label{sec:intro} Understanding how super-massive black holes at the centres of galaxies grow over cosmic time is one of the most important yet challenging problems facing modellers of galaxy formation. Observationally there is clear and compelling evidence that in galaxies that host super-massive black holes the black hole mass $M_{\\rm BH}$ correlates strongly with the stellar mass $M_{\\ast}$ and velocity dispersion $\\sigma$ of the host bulge \\citep[e.g.][]{magorrian.1998,ferrarese.2000,gebhardt.2000,tremaine.2002, haring.rix.2004,gultekin.2009}. Theoretically these correlations are widely interpreted as hallmarks of black hole feedback, which itself is a natural consequence of accretion onto the black hole \\citep[e.g.][]{silk.rees.1998, fabian.1999,king.2003,sazonov.etal.2005,king.2005}. In this picture, feedback acts to regulate the black hole's mass accretion rate $\\dot M_{\\rm BH}$ by modifying the physical and dynamical state of gas in and around its host galaxy -- so the greater $\\dot M_{\\rm BH}$, the stronger the feedback and the greater the impact on $\\dot M_{\\rm BH}$. Therefore, how one estimates $\\dot M_{\\rm BH}$ is crucial because it governs not only the rate at which the black hole grows but also the strength of the black hole feedback. This is a particularly important problem because how black hole feeding and feedback is modelled can have a profound impact on the predictions of how galaxies form \\citep[e.g.][]{bower.etal.2006,croton.etal.2006}.\\\\ The standard approach to estimating $\\dot M_{\\rm BH}$ in galaxy formation simulations is based on the work of \\citet{bondi.hoyle.1944} and \\citet{bondi.1952} \\citep[hereafter the Bondi-Hoyle method; cf.][]{dimatteo.2005,springel.etal.2005}. In the accretion problem as it was originally formulated, a spherically symmetric accretion flow is captured gravitationally by a point-like accretor from a uniform distribution of gas with zero angular momentum. Under these conditions, the accretion rate onto the accretor $\\dot M_{\\rm Bondi}$ is proportional to the square of the black hole mass $M_{\\rm BH}^2$ and the gas density $\\rho$, and inversely proportional to the cube of the sound speed $c_s$. This gives an accretion rate $\\dot{M}_{\\rm Bondi} \\propto M_{\\rm BH}^2 \\rho/c_s^3$. The assumption in galaxy formation simulations is that $\\dot M_{\\rm BH} \\propto \\dot M_{\\rm Bondi}$ \\citep[see, for example, the discussion in][]{booth.schaye.2009}. However, there are good physical reasons to believe that $\\dot M_{\\rm Bondi}$ cannot be representative of the true black hole accretion rate $\\dot M_{\\rm BH}$ in an astrophysically realistic situation \\citep[cf.][]{king.2010}. First, the black hole is embedded in the gravitational potential of a galaxy that is orders of magnitude more massive than it; this means that the gravitational force acting on the accretion flow is dominated by the mass of the galaxy rather than the black hole and so $\\dot M_{\\rm Bondi}$ will be a similar number of orders of magnitude off the true $\\dot M_{\\rm BH}$ (we show this explicitly in Hobbs et al., in preparation). Second, any astrophysically realistic accretion flow will have some angular momentum, violating one of the key assumptions made in calculating $\\dot M_{\\rm Bondi}$. This is important because it implies that infalling material will settle onto a circular orbit whose radius $R_{\\rm circ}$ is set by the angular momentum of the material with respect to the black hole \\citep[cf.][]{hobbs.etal.2010a}. In particular, it means that only the very lowest angular material will be available to feed the black hole because the timescale required for viscous transport of material through the disc is of order a Hubble time on scales of order $R\\sim 1-10\\rm pc$ \\citep[see, for example, ][]{king.2010}. This is a very restrictive condition because it is not straightforward for infalling gas to lose its angular momentum other than by colliding with other gas, which leads to angular momentum cancellation. Therefore, angular momentum provides an efficient natural barrier to accretion by the black hole, and so must be accounted for when estimating $\\dot M_{\\rm BH}$.\\\\ These arguments make clear that the Bondi-Hoyle method cannot provide a reliable estimate of $\\dot M_{\\rm BH}$ in galaxy formation simulations. If feedback from black holes plays as important a role in galaxy formation as we expect it to \\citep[e.g.][]{bower.etal.2006,croton.etal.2006}, then it is crucial that we devise an alternative method for estimating $\\dot M_{\\rm BH}$ in galaxy formation simulations that overcomes the problems that beset the Bondi-Hoyle method. In this short paper, we introduce our new ``accretion disc particle'' method (hereafter the ADP method) for estimating $\\dot M_{\\rm BH}$ in galaxy formation simulations, which accounts naturally for the angular momentum of infalling material. We use a collisionless accretion disc particle (ADP) to model the black hole and its accretion disc. The black hole accretes if and only if gas comes within the accretion radius $R_{\\rm acc}$ of the ADP, at which point it is captured and added to the accretion disc that feeds the black hole on a viscous timescale $t_{\\rm visc}$. In this way the black hole will accrete only the lowest angular momentum material from its surroundings. The layout of this paper is as follows. We describe the main features of the ADP method in \\S\\ref{sec:accretion_models}, showing how the accretion rate $\\dot M_{\\rm acc}$ onto the ADP is linked to the black hole accretion rate $\\dot M_{\\rm BH}$. In \\S\\ref{sec:feedback} we discuss briefly our momentum-driven feedback model \\citep[cf.][]{nayakshin.power.2010} as well as our implementation of the quasar pre-heating model of \\citet{sazonov.etal.2005}. The accretion rate $\\dot M_{\\rm BH}$ estimated using the ADP method is very different from one estimated using the Bondi-Hoyle method. We show this clearly in \\S\\ref{sec:results} using simple idealised numerical simulations, designed to illustrate the key differences between the ADP and Bondi-Hoyle methods for estimating $\\dot M_{\\rm BH}$. These simulations follow the collapse of an initially rotating shell of gas onto a black hole embedded in an isothermal galactic potential. Finally we summarise our results in \\S\\ref{sec:summary} and we discuss the implications for galaxy formation simulations and the $M_{\\rm BH}-\\sigma$ relation in \\S\\ref{sec:conclusions}. ", "conclusions": "\\label{sec:conclusions} We have argued that our new accretion disc particle (ADP) method provides a far more physically motivated and self-consistent approach to modelling black hole accretion than the Bondi-Hoyle method, which is the standard approach in galaxy formation simulations \\citep[cf.][]{springel.etal.2005,dimatteo.2005}. The Bondi-Hoyle method was formulated with a specific astrophysical problem in mind, quite unlike the situations that arise when modelling galaxy formation. It is not applicable to problems in which the accretion flow has non-zero angular momentum (as demonstrated in this paper) and/or in which it is embedded in the potential of a more massive host (as we show in Hobbs et al., in preparation). Therefore it is unsurprising that the Bondi-Hoyle method struggles to capture the behaviour of gas accretion in these kinds of common situations. Our ADP method is similar in spirit to the ``accretion radius'' or ``sink particle'' approaches to modelling accretion that are used extensively in simulating star formation \\citep[e.g.,][]{bate.1995,bate.bonnel.2005} and modelling gas accretion onto the super-massive black hole at the centre of the Milky Way \\citep{cuadra.etal.2006}, and we believe that it is natural to extend this approach into modelling galaxy formation. An important next step in our work, which builds on this accretion disc particle method and our recent momentum-driven wind model for feedback \\citep[cf.][]{nayakshin.power.2010}, is to combine the models in simulations of merging galaxies and ultimately cosmological galaxy formation simulations.\\\\ It is interesting to consider one important astrophysical consequence of our accretion disc particle model and to contrast it with what one would expect using the Bondi-Hoyle model. Recently it has been suggested that there is observationally evidence for separate $M_{\\rm BH}-\\sigma$ relations for elliptical galaxies and classical bulges on the one hand and pseudo-bulges on the other, such that the black holes in pseudo-bulges are underweight \\citep[see, for example,][]{greene.2008,hu.2009}. The properties of pseudo-bulges appear to deviate systematically from those of classical bulges, and in particular they are characterised by a high degree of rotational support. As we have shown, the angular momentum of infalling material provides a natural barrier to black hole growth, and so we would expect that rotationally supported systems to be more likely to be systems in which the central super-massive black hole is malnourished and underweight. Of course, the precise details of a galaxy's assembly history are important but our model would predict a systematic offset between super-massive black hole masses in galaxies that have on average accreted higher angular momentum material than those that have on average accreted lower angular momentum material. In contrast, the Bondi-Hoyle model would predict that the black hole should continue to grow to a critical black hole mass imposed by the depth of the gravitational potential in which it sits, regardless of the angular momentum of infalling material. We shall investigate this question further in future work." }, "1003/1003.0838_arXiv.txt": { "abstract": "Spurred by recent observations of 24$\\micron$ emission within wind-blown bubbles, we study the role that dust can play in such environments, and build an approximate model of a particular wind-blown bubble, `N49.' First, we model the observations with a dusty wind-blown bubble, and then ask whether dust could survive within N49 to its present age (estimated to be $5 \\times 10^5$ to $10^6$ years). We find that dust sputtering and especially dust-gas friction would imply relatively short timescales ($t \\sim 10^4$~years) for dust survival in the wind-shocked region of the bubble. To explain the $24\\micron$ emission, we postulate that the grains are replenished within the wind-blown bubble by destruction of embedded, dense cloudlets of ISM gas that have been over-run by the expanding wind-blown bubble. We calculate the ablation timescales for cloudlets within N49 and find approximate parameters for the embedded cloudlets that can replenish the dust; the parameters for the cloudlets are roughly similar to those observed in other nebula. Such dust will have an important effect on the bubble: including simple dust cooling in a wind-blown bubble model for N49, we find that the luminosity is higher by approximately a factor of six at a bubble age of about $10^4$ years. At ages of $10^7$ years, the energy contained in the bubble is lower by about a factor of eight if dust is included; if dust must be replenished within the bubble, the associated accompanying gas mass will also be very important to wind-blown bubble cooling and evolution. While more detailed models are certainly called for, this work illustrates the possible strong importance of dust in wind-blown bubbles, and is a first step toward models of dusty, wind-blown bubbles. ", "introduction": "In addition to the energy, mass, and momentum input provided by the deaths of massive stars in supernovae, early-type stars also launch hypersonic ($v \\sim 1000 - 3000$\\,km\\,s$^{-1}$, whereas $c_{\\rm s,ISM} \\sim 10$\\,km\\,s$^{-1}$) winds with mass outflow rates of order $10^{-7}$ to $10^{-6}$\\,M$_{\\odot}$\\,yr$^{-1}$ \\citep[although clumping within the wind may imply smaller mass-outflow rates; see][]{Hillier2009}. Such winds drive forward shocks both into the surrounding interstellar medium and reverse shocks back into the stellar wind; these reverse shocks then heat the stellar wind into a pressurized bubble surrounding the young star, resulting in objects known as wind-blown bubbles \\citep[WBBs; for reviews see, e.g.,][]{ZhekovMyasnikov2000, Chu2003, Arthur2007, Chu2008}. WBBs are therefore thought to have four major components: (1) an inner, free-flowing, hypersonic wind surrounded by (2) the hot, post-shocked wind plasma, which is enclosed by (3) a thin shell of $10^4$\\,K gas which resides in (4) the ambient interstellar medium \\citep[see Figure 1 of ][for a sketch of these components; they also form the basis of our Fig.~\\ref{n49Schematic}]{WeaverEtAl1977}. The structure of these WBBs is useful to study in that they help us understand the transformation of gas phases in the interstellar medium, and more generally, the impact of massive stars on their host molecular clouds, the heating of ISM gas, and the input of energy into turbulence in the ISM \\citep[e.g.,][]{Hensler2008}; they may even trigger star formation on their periphery \\citep[][]{ChurchwellEtAl2006}. Observations of WBBs also place constraints on stellar-wind power over the evolution of early-type stars. Finally, WBBs are important for setting the stage for the star's subsequent supernova explosion or perhaps gamma-ray burst. These wind-blown bubbles and their impact on the interstellar medium have been observed for many years \\citep[e.g.,][]{JohnsonHogg1965, Smith1967, ChuLasker1980, NazeEtAl2001, ChurchwellEtAl2006, WatsonEtAl2008}. Theoretical work in understanding these observations has been pursued in parallel by many researchers \\citep[e.g.,][]{JohnsonHogg1965, Mathews66, Pikelner68, Dyson73, CastorEtAl75, Falle75, WeaverEtAl1977, HartquistEtAl1986, ArthurEtAl1993, GarciaSeguraMacLow1995, GarciaSeguraEtAl1996, ArthurEtAl1996, PittardEtAl2001a, PittardEtAl2001b, NazeEtAl2002, FreyerEtAl03, FreyerEtAl2006}. These dynamical studies of wind-blown bubbles have included detailed modelling of the effect of a surrounding inhomogeneous medium \\citep[and especially, the role of clumps in the surrounding medium, which will be become important later in this work; see][for example]{ArthurEtAl1993, PittardEtAl2001b, NazeEtAl2002}. To the present day, however, such models have not included the role of dust. Recent observations by \\citet{ChurchwellEtAl2006} and \\citet{WatsonEtAl2008} present evidence that dust exists within WBBs. What role does this dust play in the structure and the evolution of the bubbles, and how might dust modify the impact of WBBs on the surrounding interstellar medium? This paper aims to study the role of dust in wind-blown bubbles in the context of observations of one WBB: N49. We will start by briefly outlining the observations of N49 in Section~\\ref{bubbleObs}. This paper will then examine the role of dust within N49; in particular, we ask if dust can explain the observed $24\\um$ emission around N49 (examined in Section~\\ref{modelDef}), and then ask if dust can survive within the $10^6$-$10^7$\\,K, post-shocked wind-blown bubble in N49 in Section~\\ref{dustSurvival}. Then, in Section~\\ref{proposedModel}, we offer a simple model to explain the observations and our results on dust survival. Our conclusions are presented in Section~\\ref{conclusions}. ", "conclusions": "We were intrigued by the possibility of dust within the wind-blown bubble 'N49' \\citep{WatsonEtAl2008}. Our goal was to explain the present observations with dust, and to try to understand how dust may survive within that environment. In this paper, we found that both the processes of dust-gas friction and sputtering of small dust grains can lead to the evacuation of dust grains from WBBs on relatively short timescales of approximately $10^4$~years. We also found, though, that if dust is present, it has a very significant effect on the structure of the wind-blown bubble: decreasing the energy within the bubble due to dust cooling, and therefore decreasing the size of the bubble relative to dust-free bubbles. However, in order for dust to be present in N49 today, that dust must be replenished; we hypothesize that high-density cloudlets ($n_{\\rm c} \\sim 10^5$\\,cm$^{-3}$) can be overrun and subsumed by the wind-blown bubble. Those cloudlets are then gradually destroyed by ablation and evaporation in the post-shock wind in the bubble, but on timescales long enough that the cloudlets would continue to supply the required density of dust to N49 over at least $5 \\times 10^5$~years. Such substructure has already been suggested by other researchers \\citep[e.g.,][]{OsterbrockFlather1959, Pittard2007}; what is new here is the consideration of dust. We caution that the dust may be supplied by other processes, too; for instance, the ejection of dust-grains from the debris disks of nearby low-mass Young Stellar Objects might also be important. This model has several implications. First, the external density around N49 must be fairly high to support high-density clouds ($n_{\\rm c} \\sim 10^5$~cm$^{-3}$); our dynamical model also assumes $n_{\\rm external} \\sim 10^4$\\,cm$^{-3}$ in the external medium. This should be tested by constraining the density through ammonia observations (Cyganowski et al., in preparation). Also, as dusty bubbles age, if the cloudlet density is constant with radius within the bubble, dust should gradually be evaporated from the inside out, such that there should be wind-blown bubbles with only an outer-rim of 24$\\um$ emission; this could be tested by comparing free-free and 24$\\um$ observations of WBBs. Finally, a similar structure of 24$\\um$ emission should be visible in other wind-blown bubbles; as long the bubble is young enough (ages of less than $10^6$~years for the cloud parameters outlined herein) that it has not destroyed all of the dusty cloudlets, similar 24$\\um$ emission should be seen; there are hints of this already (C. Watson, personal communication). Finally, the large amount of mass-loading that appears to be necessary to maintain dust in the WBB may lead to large amounts of cooling and a short lifetime for the \\citet{WeaverEtAl1977} pressure-driven phase." }, "1003/1003.2526_arXiv.txt": { "abstract": "The Taiwanese-American Occultation Survey (TAOS) project has collected more than a billion photometric measurements since 2005 January. These sky survey data---covering timescales from a fraction of a second to a few hundred days---are a useful source to study stellar variability. A total of 167 star fields, mostly along the ecliptic plane, have been selected for photometric monitoring with the TAOS telescopes. This paper presents our initial analysis of a search for periodic variable stars from the time-series TAOS data on one particular TAOS field, No.\\,151 (RA = 17$^{\\rm h}$30$^{\\rm m}$6$\\fs$67, Dec = 27\\degr17\\arcmin 30\\arcsec, J2000), which had been observed over 47 epochs in 2005. A total of 81 candidate variables are identified in the 3 square degree field, with magnitudes in the range $8 < R < 16$. On the basis of the periodicity and shape of the lightcurves, 29 variables, 15 of which were previously unknown, are classified as RR Lyrae, Cepheid, $\\delta$ Scuti, SX Phonencis, semi-regular and eclipsing binaries. ", "introduction": "% The Taiwanese-American Occultation Survey (TAOS) project aims to search for stellar occultation by small ($\\sim$1~km diameter) \\emph{Kuiper Belt Objects} (KBOs). The KBO population consists of remnant planetesimals in our Solar System, which typically have low to intermediate (below 30\\degr) inclination orbits and heliocentric distances between 30 and 50~AU \\citep{edg49,kui51,mor03}. The size distribution of large KBOs shows a broken power law with the break occurring $\\sim$30--100~km that indicates a relative deficiency of small KBOs. Such a broken power law is believed to be the consequence of competing processes of agglomeration to form progressively larger bodies versus collisional destruction. The size distribution thus provides critical information of the dynamical history of the Solar System. The stellar occultation technique, namely the dimming of a background star by a passing KBO, is the only technique capable of detecting cometary-sized bodies, which are too faint for direct imaging even with the largest telescopes \\citep{alc03,zha08}. So far, TAOS has collected several billion stellar photometric measurements, and no occultation events have been detected, indicating a significant depletion of small KBOs \\citep{zha08,bia09}. Several projects have discovered numerous variable stars as byproducts, for instance the MACHO \\citep{alc95, alc98}, EROS \\citep{bea95,der02}, OGLE \\citep{cie03, wra04}, and ROTSE-I \\citep{aker00,kin06,hof09}. Such data have enriched our knowledge of stellar variability in the Galactic fields and the Magellanic Clouds, which not only improves the number statistics, but also has helped to shed light on the detailed mechanisms of stellar variability. Knowledge of the variability has been so far still relatively poor for even the bright stars. Recent large-area sky survey projects, however, have started to turn up large numbers of variable stars. These projects include the All Sky Automated Survey \\citep[ASAS]{poj05}, the observations by the Hungarian Automated Telescope \\citep[HAT]{bak01}, the Northern Sky Variability Survey \\citep[NSVS]{woz04}, and ROTSE-I. Variable stars, notably Cepheids, RR Lyrae-type, $\\delta$ Scuti-type, SX Phonenicis-type, semi-regular variables, and eclipsing binaries are shown to be ubiquitous in Galactic fields and in clusters. The next-generation projects like the cyclic all-sky survey by the Panoramic Sky Survey And Rapid Response System (Pan-STARRS) no doubt will provide a much complete variable star census and characterization to enhance vastly our understanding of the cosmos in the time domain. While the main goal of the TAOS project is to conduct a KBO census by detecting stellar occultations, the plethora of time-series stellar photometry renders the opportunity to identify and characterize variable stars spanning a wide range of timescales, from less than a second to a few years. The first paper of the series of the TAOS stellar variability studies deals with detection of low-amplitude $\\delta$ Scuti stars \\citep{kim10}. The current paper, the second in the series, presents the effort to identify variable stars in a targeted star field. ", "conclusions": "% We have identified a total of 81 candidate variable stars in a particular field, No.\\,151 (RA = 17$^{\\rm h}$30$^{\\rm m}$6$\\fs67$, Dec = 27\\degr17\\arcmin30\\arcsec, J2000) in the TAOS survey. Among these, 29 variables can be classified, including 15 previously uncatalogued, as Cepheids, RR Lyrae stars, semiregular variables, eclipsing binaries or $\\delta$ Scuti-type variables. Their lightcurves, derived periods, semi-amplitudes, and hence the variable classification are presented here and the data are avaiable on the TAOS website, {\\verb+ http://taos.asiaa.sinica.edu.tw/demo+}. With the same methodology we expect to produce variable star lists in other TAOS fields, now with observations covering more than 4 years (2005--2009). In addition to stare-mode photometry, the zipper-mode observations provide data sampled at 5~Hz, so may be particularly useful for fast stellar variability \\citep{kim10}. The TAOS database hence has the unique potential to study several thousand stars at timescales from less than a second to a few years." }, "1003/1003.5847_arXiv.txt": { "abstract": " ", "introduction": "The strong CP problem can be solved by the Peccei-Quinn symmetry~\\cite{PQ}, that manifests at low energy as a light axion $a$ with a decay constant $f\\circa{>} 5~10^9\\GeV$~\\cite{Kim}. In supersymmetric models the axion $a$ gets extended into an axion supermultiplet which also contains the scalar saxion $s$ and the fermionic axino $\\tilde{a}$~\\cite{susyaxion}. Depending on the model of supersymmetry breaking, the axino can easily be lighter than all other sparticles, becoming the stable lightest supersymmetric particle (LSP) and consequently a Dark Matter candidate~\\cite{Wil,Ros,Goto}. It is thereby interesting to compute its cosmological abundance. The axino can be produced: i) from decays of the next to lightest sparticle (NLSP), such that $\\Omega_{\\tilde{a}} = m_{\\tilde{a}}\\Omega_{\\rm NLSP}/m_{\\rm NLSP}$~\\cite{Ros}; plus ii) thermally in the early universe when the temperature $T$ was just below the reheating temperature $T_{\\rm RH}$. We here reconsider the thermal axino abundance, improving on previous computations~\\cite{axinoHTL} in the following ways: \\begin{itemize} \\item[a)] in section~\\ref{L} we show that, beyond the well known axino/gluino/gluon interaction, there is a new axino/gluino/squark/squark interaction, unavoidably demanded by superymmetry, that contributes at the same order to the usual $2\\to 2$ scatterings that produce axinos; \\item[b)] in section~\\ref{T} we show that axinos are also thermally produced by $1\\to 2$ decays kinematically allowed by the gluon thermal mass. The gluon $\\to$ gluino + axino process gives the dominant contribution in view of the large value $g_3\\sim 1$ of the strong coupling constant; \\item[c)] in section~\\ref{pheno} we precisely model the reheating process. \\end{itemize} As a result the axino production rate is significantly enhanced. ", "conclusions": "In section~\\ref{L} we derived the axino coupling to the strong sector, finding that it is described by two terms: one is the well known $\\tilde{a} \\tilde{g} G$ coupling, the other $\\tilde{a}\\tilde{g}\\tilde{q}^*\\tilde{q}$ term was missed in previous studies, although they contribute at the same order to the thermal axino production rate. This makes the axino interaction fully analogous to the Goldstino interaction, such that we could infer the axino thermal production rate from the Goldstino rate, computed in~\\cite{RS}. Such result goes beyond the leading order computations~\\cite{axinoHTL} based on the Hard Thermal Loop (HTL) approximation $g_3\\ll 1$, which gives unphysical results at the physical value of $g_3\\sim 1$. As a consequence we find an enhancement, plotted in fig.\\fig{res}, by a factor of 6 (3) at $T_{\\rm RH}= 10^4~(10^7)\\GeV$. The function $F$ determines the axino rate as described by eq.s\\eq{gamma} and\\eq{Omega}. \\small \\paragraph{Acknowledgements} We thank L. Covi, V.S. Rychkov, C. Scrucca and G. Villadoro. \\bigskip \\footnotesize \\begin{multicols}{2}" }, "1003/1003.2656_arXiv.txt": { "abstract": "Magnetic field embedded in a perfectly conducting fluid preserves its topology for all time. Although ionized astrophysical objects, like stars and galactic disks, are almost perfectly conducting, they show indications of changes in topology, `magnetic reconnection\u2019, on dynamical time scales. Reconnection can be observed directly in the solar corona, but can also be inferred from the existence of large scale dynamo activity inside stellar interiors. Solar flares and gamma ray busts are usually associated with magnetic reconnection. Previous work has concentrated on showing how reconnection can be rapid in plasmas with very small collision rates. Here we present numerical evidence, based on three dimensional simulations, that reconnection in a turbulent fluid occurs at a speed comparable to the rms velocity of the turbulence, regardless of the value of the resistivity. In particular, this is true for turbulent pressures much weaker than the magnetic field pressure so that the magnetic field lines are only slightly bent by the turbulence. These results are consistent with the proposal by Lazarian \\& Vishniac (1999) that reconnection is controlled by the stochastic diffusion of magnetic field lines, which produces a broad outflow of plasma from the reconnection zone. This work implies that reconnection in a turbulent fluid typically takes place in approximately a single eddy turnover time, with broad implications for dynamo activity and particle acceleration throughout the universe. In contrast, the reconnection in 2D configurations in the presence of turbulence depends on resistivity, i.e. is slow. ", "introduction": " ", "conclusions": "" }, "1003/1003.0379_arXiv.txt": { "abstract": "Diffusive shock acceleration in supernova remnants is the most widely invoked paradigm to explain the Galactic cosmic ray spectrum. Cosmic rays escaping supernova remnants diffuse in the interstellar medium and collide with the ambient atomic and molecular gas. From such collisions gamma-rays are created, which can possibly provide the first evidence of a parent population of runaway cosmic rays. We present model predictions for the GeV to TeV gamma-ray emission produced by the collisions of runaway cosmic rays with the gas in the environment surrounding the shell-type supernova remnant RX~J1713.7-3946. The spectral and spatial distributions of the emission, which depend upon the source age, the source injection history, the diffusion regime and the distribution of the ambient gas, as mapped by the LAB and NANTEN surveys, are studied in detail. In particular, we find for the region surrounding RX~J1713-3946, that depending on the energy one is observing at, one may observe startlingly different spectra or may not detect any enhanced emission with respect to the diffuse emission contributed by background cosmic rays. This result has important implications for current and future gamma-ray experiments. ", "introduction": "\\label{sec:intro} Cosmic rays (CRs) are the highly energetic protons and nuclei which fill the Galaxy and carry, at least in the vicinity of the Sun, as much energy per unit volume as the energy density of starlight or of the interstellar magnetic fields or the kinetic energy density of the interstellar gas. CRs of energies up to the ``knee'' (${10}^{15}$ eV), or even up to ${10}^{18}$ eV, are believed to originate from sources located within the Galaxy. Galactic CRs are thought to be accelerated via diffusive shock acceleration (DSA) operating in the expanding shells of supernova remnants (SNRs). For references see e.g. \\cite{Blandford,Malkov}. The accelerated CRs interact with ambient gas through inelastic collisions, producing neutral pions, which then decay into \\Grays. If the bulk of Galactic CRs up to at least PeV energies are indeed accelerated in SNRs, then TeV \\Grays are expected to be emitted during the acceleration process CRs undergo within SNRs \\citep{Drury}. Indeed TeV \\Grays have been detected from the shells of SNRs \\citep{Aharonian:nature,Albert}. However, such observations do not constitute a definitive proof that CRs are accelerated in SNRs, since the observed emission could be produced by energetic electrons up scattering low energy photon fields. We note, in fact, that \\Grays are also produced through inverse Compton and bremsstrahlung processes of very highly energetic electrons. \\Grays are also expected to be emitted when the accelerated CRs propagate into the interstellar medium (ISM) \\citep{Montmerle,Issa,Aharonian:1991,Aharonian:1996,Gabici2009}. Before being isotropised by the Galactic magnetic fields, the injected CRs produce $\\gamma$-ray emission, which can significantly differ from the emission of the SNR itself, as well as from the diffuse emission contributed by the background CRs and electrons, because of the hardness of the runaway CR spectrum, which is not yet steepened by diffusion. The extension of such diffuse sources does not generally exceed a few hundred parsecs, the scale at which the spectra of the injected CRs can significantly differ from the spectrum of the CR background \\citep{Aharonian:1996,Gabici,Marrero,Gabici2009}. These diffuse sources are often correlated with dense molecular clouds (MCs), which act as a target for the production of \\Grays due to the enhanced local CR injection spectrum. \\cite{Casse} and \\cite{Montmerle} have pointed out that SNRs are located in star forming regions, which are rich in molecular hydrogen. In other words, CR sources and MCs are often associated and target-accelerator systems are not unusual within the Milky Way. In fact, there have been recently claims of detection of \\Grays in association with dense MCs close to candidate CR sources, both at GeV energies \\citep{Abdo2009d,Abdo2010,Tavani,Castro} and TeV energies \\citep{Albert2007,Aharonian:2008,Aharonian:2008b,Aharonian:2008c}. In order to probe the interaction of the low-energy component of the ambient cosmic ray flux, \\cite{Montmerle:2009} has recently proposed to measure enhanced ionization in TeV-bright molecular clouds, using millimeter observations. On the other hand, high energy CRs (above 1 GeV up to PeV energies) interacting with the ambient gas emit \\Grays. For this reason the high sensitivity, high resolution \\Gray data from current (HESS, Magic, Veritas, Fermi and Agile) and future detectors, such as AGIS, CTA and HAWC \\citep{Sinnis,Hinton}, together with the knowledge of the distribution of the atomic and molecular hydrogen in the Galaxy on sub-degree scales are crucial to explore the flux of high energy CRs close to the candidate CR sources and to pinpoint the long searched-for sites of CR acceleration. The \\Gray radiation from hadronic interactions from regions close to CR sources depends not only on the total power emitted in CRs by the sources and on the distance of the source to us, but also on the ambient interstellar gas density, the local diffusion coefficient and the injection history of the CR source. It is therefore difficult to definitely recognize the sites of CR acceleration from \\Gray observations alone, since very often only qualitative predictions are provided, rather than robust quantitative predictions, especially from a morphological point of view. In order to fully exploit the present and future experimental facilities and to test the standard scenario for CR injection in SNRs and propagation, we present here model predictions of the spectral and morphological features of the hadronic \\Grays emission surrounding the candidate CR source, RX~J1713.7-3946, by constructing as quantitative a model as possible. In particular, we will convey all information concerning the environment, the source age, the acceleration rate and history, which all play a role in the physical process of CR injection and propagation. Building upon the modeling of the broadband emission from MCs close to CR accelerators developed in \\cite{Gabici2009} and upon the analysis of the CR background discussed in \\cite{Casanova2009} (hereafter Paper 1), we compute the expected \\Gray emissivity from hadronic interactions of runaway CRs for the region $340^\\circ4$ keV) characterized by a power law (Yaqoob et al. 1993; Warwick et al. 1995, 1996). The 2--10 keV flux can double on timescale of $\\sim 0.5$ days (Tananbaum et al. 1978; Yaqoob \\& Warwick 1991) and flaring events are seen on timescales of days and weeks (Elvis 1976; Lawrence 1980; Edelson et al. 1996; Markowitz et al. 2003; de Rosa et al. 2007). (2) The photon spectral index ($\\Gamma$) of the hard power law spectrum varies between $\\sim 1.35$ and $\\sim 1.7$, and correlates with the 2--10 keV flux\\footnote{The linear relation between $\\Gamma$ and the absorption corrected 2--10 keV flux in units of $10^{-11}$ erg s$^{-1}$ derived from EXOSAT and Ginga measurements is $\\Gamma=1.18+0.012F_{2-10keV}^c$ (Perola et al. 1986) and $\\Gamma=1.35+0.011F_{2-10keV}^c$ (Yaqoob \\& Warwick 1991), respectively.}. The spectrum becomes softer with increasing flux, and the same variability in the X-ray continuum flux is observed on time scales of days to years in the spectral slope. Note that the range of $\\Gamma$ corresponds to a harder spectrum than the canonical value for Seyferts of $\\Gamma\\sim 1.8-1.9$ (e.g., Mushotzky 1984, Nandra \\& Pounds 1994). (3) A soft ``excess'' ($<2$ keV) is present in the spectra when fitted with the power law measured in the hard-band and a uniform absorption (Holt et al. 1980). Extended soft X-ray emission associated with ionized gas observed in the optical has been detected (Elvis et al. 1983; Morse et al. 1995; Ogle et al. 2000; Yang et al 2001). The soft flux shows stability against the high variations in the UV and hard X-ray continua (Perola et al. 1986; Weaver et al. 1994a,b). (4) A clear emission line is present at $6.39\\pm 0.07$ keV (Matsuoka et al. 1986; Wang et al. 2001; Schurch et al. 2003), consistent with being Fe$K_{\\alpha}$ emission line produced by the fluoresence of cold iron illuminated by the X-ray continuum. The line has a relatively narrow Gaussian profile and the line flux remains constant over time (Warwick et al. 1989; Schurch et al. 2003). Long term monitoring with {\\em RXTE}/PCA over 800 days (Markowitz et al. 2003) shows a decreasing line flux. De Rosa et al. (2007) suggest that the line exhibits some variability related to the reflection component. Particularly interesting results were reported from previous {\\em Chandra} observations. Yang et al. (2001) presented {\\em Chandra} observations of NGC 4151 and found that the 2--9 keV spectrum of the nucleus is described by a heavily absorbed ($N_H\\simeq 3\\times 10^{22}$ cm$^{-2}$), extremely hard power law (photon index $\\Gamma= 0.32^{+0.05}_{-0.12}$). This is consistent with the {\\em Chandra} HETG spectra reported in Ogle et al. (2000), in which the hard continuum emission is characterized by a photon index $\\Gamma=0.4\\pm 0.3$ with a column of $N_H\\simeq 3.7\\times 10^{22}$ cm$^{-2}$. Note that both observations found NGC 4151 in a low state with $F_{2-10keV} \\sim 5.5\\times 10^{-11}$ erg s$^{-1}$ and do not follow the $\\Gamma$--$F_{2-10keV}$ correlation. The rather unusually flat power law slope complicates the understanding of the X-ray spectrum of NGC 4151. We have obtained deep {\\em Chandra} observation aimed to take advantage of {\\em Chandra}'s sub-arcsecond spatial resolution (van Speybroeck et al. 1997) to study the soft circum-nuclear extended emission. These data also present an opportunity to examine the spectra of this puzzling nucleus. The {\\em Chandra} observations and data reduction are briefly described in \\S~\\ref{obs}. We examine the light curves of the nucleus and perform diagnostic analysis of the X-ray spectra in \\S~\\ref{pileup_analysis}. In \\S~\\ref{models} we explain the various ways to model the X-ray spectra. We then analyzed the archival deep {\\em ASCA} observation of NGC 4151 in low state to check the validity of our models (\\S~\\ref{asca_look}). Finally we discuss the results in \\S~\\ref{discussion} and summarize our findings in \\S~\\ref{summary}. ", "conclusions": "\\label{summary} The deep $\\sim$200 ks {\\em Chandra} observation of the Seyfert 1 nucleus in NGC 4151 was analyzed to understand its complex spectral variability, in particular the unusally flat photon index reported from earlier $Chandra$ observations. We find: (1) The 2--10 keV flux of the nucleus varied from $6\\times 10^{-11}$ erg s$^{-1}$ cm$^{-2}$ (ObsID 9218 and ObsID 9217a) to $\\sim 10^{-10}$ erg s$^{-1}$ cm$^{-2}$ (ObsID 9217b), resulting in a non-linear count rate variation of opposite sign because of significant pileup. (2) With pileup corrected spectral fitting, we are able to recover the spectral parameters and find consistency with those derived from unpiled events in the ACIS readout streak and outer region of the bright PSF core. (3) The low flux segment shows a hard photon index $\\Gamma_2\\sim 0.7-0.9$ similar to that seen in the historical low state for a simple power-law fit. More complex, physically meaningful models are attempted and provide good fits to the piled spectra, including a Compton reflection model, the reflection model subjected to a partially covered absorber, and the reflection model subjected to an ionized absorber. (4) The observed flat spectrum and its variability can be interpreted as due to an intrinsically varying continuum with respect to an underlying Compton reflection component, or a variable X-ray absorber partially covering the continuum source. Including the partial covering absorber provides only a marginal improvement over the simpler reflection-only model, nevertheless it gives a continuum photon index that is typical of Seyfert 1s. (5) If the absorption model is correct, the size of X-ray emission region is constrained $r\\simeq 20 r_s$, and the X-ray absorber is located at $r\\leq 1.6\\times 10^4 r_s$ from the nucleus, possibly associated with the BELR clouds. (6) If instead the reflection-only model is correct, the presence of a constant reflection component with respect to the rapidly flaring continuum of NGC 4151 implies the reflector is located far from the nucleus, consistent with the absence of broad FeK$\\alpha$ fluorescence line." }, "1003/1003.0904_arXiv.txt": { "abstract": "The identity of dark matter is a question of central importance in both astrophysics and particle physics. In the past, the leading particle candidates were cold and collisionless, and typically predicted missing energy signals at particle colliders. However, recent progress has greatly expanded the list of well-motivated candidates and the possible signatures of dark matter. This review begins with a brief summary of the standard model of particle physics and its outstanding problems. We then discuss several dark matter candidates motivated by these problems, including WIMPs, superWIMPs, light gravitinos, hidden dark matter, sterile neutrinos, and axions. For each of these, we critically examine the particle physics motivations and present their expected production mechanisms, basic properties, and implications for direct and indirect detection, particle colliders, and astrophysical observations. Upcoming experiments will discover or exclude many of these candidates, and progress may open up an era of unprecedented synergy between studies of the largest and smallest observable length scales. ", "introduction": "\\label{sec:introduction} The evidence that dark matter is required to make sense of our Universe has been building for some time. In 1933 Fritz Zwicky found that the velocity dispersion of galaxies in the Coma cluster of galaxies was far too large to be supported by the luminous matter~\\cite{Zwicky:1933gu}. In the 1970s, Vera Rubin and collaborators~\\cite{Rubin:1970zz,Rubin:1980zd} and Albert Bosma~\\cite{1978PhDT.......195B} measured the rotation curves of individual galaxies and also found evidence for non-luminous matter. This and other ``classic'' evidence for non-luminous matter (see, \\eg, \\citet{Trimble:1987ee}) has now been supplemented by data from weak~\\cite{Refregier:2003ct} and strong~\\cite{Tyson:1998vp} lensing, hot gas in clusters~\\cite{Lewis:2002mfa}, the Bullet Cluster~\\cite{Clowe:2006eq}, Big Bang nucleosynthesis (BBN)~\\cite{Fields:2008}, further constraints from large scale structure~\\cite{Allen:2002eu}, distant supernovae~\\cite{Riess:1998cb,Perlmutter:1998np}, and the cosmic microwave background (CMB)~\\cite{Komatsu:2010fb}. Together, these data now provide overwhelming evidence for the remarkable fact that not only is there non-luminous matter in our Universe, but most of it is not composed of baryons or any of the other known particles. Current data imply that dark matter is five times more prevalent than normal matter and accounts for about a quarter of the Universe. More precisely, these data constrain the energy densities of the Universe in baryons, non-baryonic dark matter (DM), and dark energy $\\Lambda$ to be~\\cite{Komatsu:2010fb} \\begin{eqnarray} \\Omega_{\\text{B}} &\\simeq& 0.0456 \\pm 0.0016 \\\\ \\OmegaDM &\\simeq& 0.227 \\pm 0.014 \\\\ \\Omega_{\\Lambda} &\\simeq& 0.728 \\pm 0.015 \\ . \\label{Lambda} \\end{eqnarray} Despite this progress, all of the evidence for dark matter noted above is based on its gravitational interactions. Given the universality of gravity, this evidence does little to pinpoint what dark matter is. At the same time, the identity of dark matter has far-reaching implications: in astrophysics, the properties of dark matter determine how structure forms and impact the past and future evolution of the Universe; and in particle physics, dark matter is the leading empirical evidence for new particles, and there are striking hints that it may be linked to attempts to understand electroweak symmetry breaking, the leading puzzle in the field today. The identity of dark matter is therefore of central importance in both fields and ties together studies of the Universe at both the largest and smallest observable length scales. In this review, we discuss some of the leading dark matter candidates and their implications for experiments and observatories. The wealth of recent cosmological data does constrain some dark matter properties, such as its self-interactions and its temperature at the time of matter-radiation equality. Nevertheless, it is still not at all difficult to invent new particles that satisfy all the constraints, and there are candidates motivated by minimality, particles motivated by possible experimental anomalies, and exotic possibilities motivated primarily by the desire of clever iconoclasts to highlight how truly ignorant we are about the nature of dark matter. Here we will focus on dark matter candidates that are motivated not only by cosmology, but also by robust problems in particle physics. For this reason, this review begins with a brief summary of the standard model of particle physics, highlighting its basic features and some of its problems. As we will see, particle physics provides strong motivation for new particles, and in many cases, these particles have just the right properties to be dark matter. We will find that many of them predict signals that are within reach of current and near future experiments. We will also find that unusual predictions for astrophysics emerge, and that cold and collisionless dark matter is far from a universal prediction, even for candidates with impeccable particle physics credentials. At the same time, it will become clear that even in favorable cases, a compelling solution to the dark matter problem will not be easy to achieve and will likely rely on synergistic progress along many lines of inquiry. An outline of this review is provided by \\tableref{summary}, which summarizes the dark matter candidates discussed here, along with their basic properties and opportunities for detection. Some of the acronyms and symbols commonly used in this review are defined in \\tableref{definitions}. \\begin{table}[tbp] \\begin{minipage}{\\columnwidth} \\begin{tabular}{lcccccc} \\hline\\hline \\rule[1mm]{0mm}{5mm} & WIMPs & SuperWIMPs & Light $\\tilde{G}$ & Hidden DM & Sterile $\\nu$ & Axions \\vspace*{.1in} \\\\ \\hline Motivation \\rule[3mm]{0mm}{5mm} & GHP & GHP & \\ttwo{GHP}{NPFP} & \\ttwo{GHP}{NPFP} & $\\nu$ Mass & Strong CP \\vspace*{.2in} \\\\ \\ttwo{Naturally}{Correct~$\\Omega$} & Yes & Yes & No & Possible & No & No \\vspace*{.12in} \\\\ \\ttwo{Production}{Mechanism} & Freeze Out & Decay & Thermal & Various & Various & Various \\vspace*{.12in} \\\\ Mass Range & GeV$-$TeV & GeV$-$TeV & eV$-$keV & GeV$-$TeV & keV & $\\mu\\ev - \\text{meV}$ \\vspace*{.12in} \\\\ Temperature & Cold & Cold/Warm & Cold/Warm & Cold/Warm & Warm & Cold \\vspace*{.12in} \\\\ Collisional & & & & $\\surd$ & & \\vspace*{.12in} \\\\ \\ttwo{Early}{Universe} & & $\\surd\\surd$ & & $\\surd$ & & \\vspace*{.12in} \\\\ \\ttwo{Direct}{Detection} & $\\surd\\surd$ & & & $\\surd$ & & $\\surd\\surd$ \\vspace*{.12in} \\\\ \\ttwo{Indirect}{Detection} & $\\surd\\surd$ & $\\surd$ & & $\\surd$ & $\\surd\\surd$ & \\vspace*{.12in} \\\\ \\ttwo{Particle}{Colliders} & $\\surd\\surd$ & $\\surd\\surd$ & $\\surd\\surd$ & $\\surd$ & & \\vspace*{.12in} \\\\ \\hline \\hline \\end{tabular} \\end{minipage} \\caption{Summary of dark matter particle candidates, their properties, and their potential methods of detection. The particle physics motivations are discussed in \\secref{problems}; GHP and NPFP are abbreviations for the gauge hierarchy problem and new physics flavor problem, respectively. In the last five rows, $\\surd\\surd$ denotes detection signals that are generic for this class of dark matter candidate and $\\surd$ denotes signals that are possible, but not generic. ``Early Universe'' includes phenomena such as BBN and the CMB; ``Direct Detection'' implies signatures from dark matter scattering off normal matter in the laboratory; ``Indirect Detection'' implies signatures of late time dark matter annihilation or decay; and ``Particle Colliders'' implies signatures of dark matter or its progenitors produced at colliders, such as the Large Hadron Collider (LHC). See the text for details.} \\vspace{0.2cm} \\label{table:summary} \\end{table} \\begin{table}[tbp] \\begin{minipage}{\\columnwidth} \\begin{tabular}{ll} \\hline\\hline $\\chi \\qquad \\qquad \\qquad $ & lightest neutralino, a supersymmetric dark matter candidate \\\\ $\\gravitino$ & gravitino, a supersymmetric dark matter candidate \\\\ $G_N$ & Newton's gravitational constant \\\\ GMSB & gauge-mediated supersymmetry breaking \\\\ LKP & lightest Kaluza-Klein particle \\\\ LSP & lightest supersymmetric particle \\\\ NLSP & next-to-lightest supersymmetric particle \\\\ $\\mplanck$ & Planck mass $\\simeq 1.2\\times 10^{19}~\\gev$ \\\\ $\\mstar$ & reduced Planck mass $\\simeq 2.4\\times 10^{18}~\\gev$ \\\\ \\ttwo{minimal}{supergravity} & a simple version of the MSSM specified by 5 parameters \\\\ MSSM & supersymmetric standard model with minimal number of extra particles \\\\ SM & standard model of particle physics \\\\ stau & scalar superpartner of the tau lepton \\\\ superWIMP & superweakly-interacting massive particle \\\\ UED & universal extra dimensions \\\\ WIMP & weakly-interacting massive particle \\\\ $X$ & general dark matter candidate \\\\ \\hline \\hline \\end{tabular} \\end{minipage} \\caption{Definitions of acronyms and symbols commonly used in this review.} \\vspace{0.2cm} \\label{table:definitions} \\end{table} ", "conclusions": "Current observations support a remarkably simple model of the Universe consisting of baryons, dark matter, and dark energy, supplemented by initial conditions determined by an early epoch of inflation. If scientific progress is characterized by periods of confusion, which are resolved by neat and tidy models, which are then launched back into confusion by further data, the current era is most definitely of the neat and tidy sort. Dark matter may be the area that launches us back into confusion with further data. The microscopic properties of dark matter are as much of a mystery now as they were in the 1930's. In the next few years, however, searches for dark matter through a variety of means discussed here will discover or exclude many of the most promising candidates. At its core, the dark matter problem is highly interdisciplinary. Rather than attempt a summary of this review, we close with some optimistic, but plausible, scenarios for the future in which experiments from both particle physics and astrophysics are required to identify dark matter. Consider the following examples: \\begin{itemize} \\setlength{\\itemsep}{1pt}\\setlength{\\parskip}{0pt}\\setlength{\\parsep}{0pt} \\item {\\em Scenario 1}: Direct detection experiments see a dark matter signal in spin-independent scattering. This result is confirmed by the LHC, which sees a missing energy signal that is followed up by precision measurements pointing to a 800 GeV Kaluza-Klein gauge boson. Further LHC studies constrain the Kaluza-Klein particle's predicted thermal relic density to be identical at the percent level with $\\OmegaDM$, establishing a new standard cosmology in which the dark matter is composed entirely of Kaluza-Klein dark matter, cosmology is standard back to 1 ns after the Big Bang, and the Universe has extra dimensions. Direct and indirect detection rates are then used to constrain halo profiles and substructure, ushering in a new era of dark matter astronomy. \\item {\\em Scenario 2}: The LHC discovers heavy, charged particles that are apparently stable. Together the LHC and International Linear Collider determine that the new particles are staus, predicted by supersymmetry. Detailed follow-up studies show that, if these staus are absolutely stable, their thermal relic density is larger than the total mass of the Universe! This paradox is resolved by further studies that show that staus decay on time scales of a month to gravitinos. Careful studies of the decays determine that the amount of gravitinos in the Universe is exactly that required to be dark matter, providing strong quantitative evidence that dark matter is entirely in the form of gravitinos, and providing empirical support for supergravity and string theory. \\item {\\em Scenario 3}: An X-ray experiment discovers a line signal. Assuming this results from decaying sterile neutrinos, the photon energy determines the neutrino's mass $m = 2 E_{\\gamma}$, the intensity determines the neutrino mixing angle ($I \\propto \\sin^2\\theta$), and the image morphology determines the dark matter's spatial distribution. {}From the mass and radial distribution, theorists determine the free-streaming length. This favors production from Higgs decays over production by oscillations, leading to predictions of non-standard Higgs phenomenology, which are then confirmed at the LHC. Additional information on neutrino parameters from the LHC strengthens the hypothesis of sterile neutrino dark matter, and the energy distribution of the narrow spectral line is then used to study the expansion history of the Universe and dark energy. \\end{itemize} These scenarios are, of course, highly speculative and idealized, but they illustrate that, even in ideal scenarios that we have studied and understand, close interactions between many subfields will be required. At the same time, if any of the ideas discussed here is correct, there are promising prospects for the combination of detection methods in particle physics and astrophysics to identify dark matter in the not-so-distant future." }, "1003/1003.2901_arXiv.txt": { "abstract": "We investigate the relationship between R Coronae Borealis (RCB) stars and hydrogen-deficient carbon (HdC) stars by measuring precise $^{16}$O/$^{18}$O ratios for five cool RCB stars. The $^{16}$O/$^{18}$O ratios are derived by spectrum synthesis from high-resolution (R$\\sim$50,000) K-band spectra. Lower limits to the $^{16}$O/$^{17}$O and $^{14}$N/$^{15}$N ratios as well as Na and S abundances (when possible) are also given. RCB stars in our sample generally display less $^{18}$O than HdC stars - the derived $^{16}$O/$^{18}$O ratios range from 3 to 20. The only exception is the RCB star WX CrA, which seems to be a HdC-like star with $^{16}$O/$^{18}$O=0.3. Our result of a higher $^{16}$O/$^{18}$O ratio for the RCB stars must be accounted for by a theory of the formation and evolution of HdC and RCB stars. We speculate that a late dredge-up of products of He-burning, principally $^{12}$C and $^{16}$O, may convert a $^{18}$O-rich HdC star into a $^{18}$O-poor RCB star as the H-deficient star begins its final evolution from a cool supergiant to the top of the white dwarf cooling track. ", "introduction": "The R Coronae Borealis (RCB) stars and their likely cousins the hydrogen-deficient carbon (HdC) stars and the extreme helium (EHe) stars have long posed a puzzle as to their origins in terms of stellar evolution. Two leading scenarios have survived decades of theoretical and observational scrutiny. In one, the H-deficient supergiant is formed from the merger of a He white dwarf (WD) with a C-O white dwarf (Webbink 1984; Iben \\& Tutukov 1984; Saio \\& Jeffery 2002). This path is widely referred to as the double-degenerate (DD) scenario. In the other, these H-deficient stars result from a final, post-AGB helium shell flash in the central star of a planetary nebula. The final flash may transform the star on the WD cooling track into a H-deficient supergiant; this the so-called ``born-again\" scenario is discussed by Herwig (2001) and Bl\\\"{o}cker (2001) and often labeled the FF scenario. Open questions remaining include: Is the DD or the FF scenario the dominant mechanism? If both are operative, how does one distinguish a product of the DD from one of the FF scenario? Elemental and isotopic abundances of C, N, and O are potentially powerful agents for testing the different evolutionary scenarios proposed. Previous studies of the CNO abundances support production of the RCBs and the EHes by the DD scenario (Pandey et al. 2006). A dramatic advance was made by Clayton et al.'s (2005, 2007) discovery from medium-resolution infrared spectra of the CO 2.3 $\\mu$m bands that $^{18}$O was very abundant for some cool RCBs and HdCs. Exploratory calculations led Clayton et al. (2007) to propose that $^{18}$O was synthesised from $^{14}$N during the merger in the DD scenario. Extraordinary conditions are required for the FF scenario to lead to abundant $^{18}$O. More recently, we analyzed high-resolution (R=50,000) spectra of a few narrow windows in the K band of the five known HdC stars and a few RCB stars (Garc\\'\\i a-Hern\\'andez et al. 2009). In these spectra, the CO spectrum is resolved and application of spectrum synthesis enables a more precise estimate of the $^{16}$O/$^{18}$O ratio to be obtained. Our recent analysis generally confirmed reports by Clayton and colleagues from R=5,900 (or lower resolution) spectra. We confirm that the $^{16}$O/$^{18}$O ratio is less than unity for those HdC stars (3 of the known 5) exhibiting CO lines in their spectra. However, $^{16}$O/$^{18}$O=16 was obtained for the cool RCB star SAps; the other RCB stars in our sample were too warm to display CO molecular lines in their spectra. Our result $^{16}$O/$^{18}$O=16 for S Aps contrasts with Clayton et al. 's ratio of 4. Our spectra show the gain in information resulting from the ability at the higher resolution to distinguish clearly different CO isomers. High-resolution spectra are essential to derive reliable and precise oxygen isotopic ratios for these hydrogen-deficient stars. With S Aps showing a nearly 50-fold difference in the $^{16}$O/$^{18}$O ratio from the analysable HdC stars, the question arises - is a higher ratio characteristic of the RCBs? If so, this presumably provides a clue to understand the evolutionary relationship between HdC and RCB stars. Unfortunately, no new information may be drawn about the oxygen isotopic ratio in HdC stars until more HdC stars are discovered, but there are several RCBs with CO bands of a suitable strength not yet observed at high resolution. The example of S Aps shows that RCBs should be reobserved at higher resolution before drawing conclusions about the HdC-RCB connection. Therefore, we have extended high-resolution K-band observations to another five cool RCB stars in order to establish the range of the $^{16}$O/$^{18}$O ratio among RCBs (Clayton et al.'s range from the 6 RCBs was 1 to $\\geq$12). Section 2 describes the high-resolution infrared spectroscopic observations and gives an overview of the observed spectra. Our abundance analysis and the results obtained are presented in Section 3 and 4, respectively. The derived oxygen isotopic ratios are discussed in the context of the HdC-RCB connection in Section 5. Final conclusions are presented in Section 6. ", "conclusions": "Challenges are not a new phenomenon to students of the R Coronae Borealis stars. High-resolution infared spectroscopy promises to provide responses to some of the key challenges regarding the RCBs and their putative relatives the HdC and EHe stars. Here, we suggested an evolutionary link between the HdCs and the RCBs involving the dredge-up of material to the surface. Further exploration of the link is possible. Unfortunately, the list of known HdCs has been exhausted. A few cool RCBs remain to be observed in the K band at high resolution. Of these, the most critical is V CrA. Unfortunately, our observations were obtained when the star was at minimum and infrared emission came from the dust shell and not the stellar photosphere. V CrA is special target because it is one of four known minority RCB stars (Rao \\& Lambert 1994) and the only one likely to show CO lines. Minority RCBs are distinguished by their high S/Fe and Si/Fe ratios (among other abundance anomalies): V CrA has [Si/Fe] $\\simeq$ [S/Fe] $\\simeq$ 2 (Asplund et al. 2000; Rao \\& Lambert 2008). Additionally, V CrA is rich in $^{13}$C with $^{12}$C/$^{13}$C $\\simeq 3$ (Rao \\& Lambert 2008). Among those HdC and RCBs for which the carbon isotopic ratio may be estimated, V CrA is the only known example with a high $^{13}$C abundance. What does the combination of anomalous abundance ratios and high $^{13}$C abundance imply about the origin of V CrA and minority RCBs? Is their origin a variant of the DD scenario or must another (i.e., the FF) scenario be invoked? Measurement of the oxygen isotopic ratios is potentially a valuable clue to the answers to these questions. Infrared spectroscopy offers the possibility of probing the carbon problems posed by the RCBs and HdCs - see Asplund et al. (2000) and Garc\\'{\\i}a-Hern\\'{a}ndez et al. (2009) - in the failure to account for the strengths of the atomic and molecular carbon lines. From the optical to the infrared, the leading contributor to the continuous opacity depending on the star's effective temperature may change from the photoionization of neutral carbon to free-free absorption from neutral helium. Full spectral coverage at high spectral resolution from the optical through at least the K band offers a novel opportunity to probe and hopefully to resolve the carbon problems. With their resolution, the chemical compositions of these fascinating H-deficient stars should be on a firmer footing and insights gained into their origins." }, "1003/1003.3888_arXiv.txt": { "abstract": "Helioseismology has produced unprecedented measurements of the Sun's internal structure and dynamics over the past 25~years. Much of this work has been based on global helioseismology. Now local helioseismology too is showing its great promise. This review summarizes very briefly the principal global results that may be relevant to an understanding of the origins of solar magnetism. Recent results regarding the variation of frequencies over the solar cycle and the temporal variations of subsurface flows are briefly summarized. ", "introduction": "Helioseismology is concerned with the study of the Sun's internal structure and dynamics using the properties of acoustic waves that propagate through the interior and cause observable motion of the photosphere and lower solar atmosphere. The principal properties used for this study are the frequencies of resonant global modes of the Sun set up by these acoustic waves. Since the Sun is to a good approximation spherically symmetric, the horizontal spatial structure of the modes is described by spherical harmonics $Y_l^m(\\theta,\\phi)$, where $\\theta$ is co-latitude and $\\phi$ is longitude. The modes are then labelled by three quantum numbers, the degree $l$ and order $m$ of the spherical harmonic and a radial order $n$ which is essentially the number of nodes in the mode's structure in the radial direction. The frequencies $\\nu_{nlm}$ depend on the conditions in the solar interior that affect wave propagation. In a non-rotating, perfectly spherically star the frequencies would have a degeneracy in that they would not depend on $m$ for given $n$ and $l$: this degeneracy is lifted by rotation, structural asphericities and magnetic fields, and measurements of the resulting frequency splitting can be used to make inferences about these properties. (The frequency splitting within a multiplet of given $n$ and $l$ can be decomposed into parts that are odd and even functions of $m$: the odd component arises from rotation, while the even component arises from magnetic and thermal asphericities and distortions of the shape of the star from spherical symmetry.) Application of inverse techniques provides maps such as of the adiabatic sound speed $c$, density $\\rho$, rotation and wave-speed asphericities in the Sun's otherwise impenetrable interior. Spatially resolved measurements of the Sun's oscillations by the Global Oscillation Network Group (GONG) and the Michelson Doppler Image (MDI) instrument on board the SOHO satellite began in the mid-1990s and thus now provide essentially continuous coverage of one solar cycle. The whole-disk Sun-as-a-star measurements of the Birmingham Solar Oscillation Network (BiSON) extend back even further. Thus helioseismology is able to comment on frequency changes occurring over the solar cycle and possible changes in flows and acoustic asphericities over that time. ", "conclusions": "Helioseismology has produced unprecedented measurements of the Sun's internal structure and dynamics over the past 25~years. These include having mapped the solar rotation over most of the interior, and discovering the solar tachocline. The frequencies of the Sun's global oscillations change over the solar cycle. Observed changes the odd component of the frequency splittings reflect changes in the solar internal rotation in possibly below the convection zone. The behaviour of the banded zonal flows (torsional oscillations) give a length of approximately 12~years for Cycle 23. Most of the changes in mean multiplet frequencies and in the even component of the frequency splittings comes from changes at or very close to the surface, caused by changes in the surface magnetic field (or something that is highly correlated with the surface field). The frequencies of low-degree modes are lower this minimum than during the previous minimum: there may be some subsurface differences in the two minima that account for this. There may also be some small variation in wave speed at the base of the convection zone correlated with surface activity. Local helioseismology clearly detects temporal and spatial variations, and is the fastest developing area of helioseismology. But there is a need for improved forward models in order to make robust inferences about the physical causes of those variations." }, "1003/1003.3841_arXiv.txt": { "abstract": "{The catalogue by Johannes Hevelius with the positions and magnitudes of 1564 entries was published by his wife Elisabeth Koopman in 1690. We provide a machine-readable version of the catalogue, and briefly discuss its accuracy on the basis of comparison with data from the modern Hipparcos Catalogue\\thanks{The full Table \\hevelius\\ (Table\\,\\ref{t:hevelius}) is available in electronic from only at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/} . We compare our results with an earlier analysis by Rybka (1984), finding good overall agreement. The magnitudes given by Hevelius correlate well with modern values. The accuracy of his position measurements is similar to that of Brahe, with $\\sigma=2$\\arcmin\\ for longitudes and latitudes, but with more errors $>5$\\arcmin\\ than expected for a Gaussian distribution. The position accuracy decreases slowly with magnitude. The fraction of stars with position errors larger than a degree is 1.5\\%, rather smaller than the fraction of 5\\%\\ in the star catalogue of Brahe. ", "introduction": "Even though a major improvement on earlier work, the star catalogue produced by Tycho Brahe (1598, 1602), and re-edited by Kepler (1627), contains occasional large errors. Johannes Hevelius decided to produce a better and larger catalogue, which was printed after his death by his wife and collaborator Elisabeth Koopman (Hevelius 1690). The title page has 1687, the year in which the catalogue was printed, but publication followed only in 1690. The extent of the contribution by Koopman to measuring the stars and producing the catalogue is not known; her presence on several images of instruments used by Hevelius suggests that it was significant. A brief but informative and well-illustrated description of the life of Hevelius and of his star catalogue is given by Volkoff et al.\\ (1971) in a book celebrating the acquisition by the Brigham Young University of Hevelius' manuscript for the catalogue. In 1679 Halley visited Hevelius and his observatory, and verified that measurements with the instruments of Hevelius, fitted with naked-eye sights, were more accurate than measurements with contemporary instruments with telescopic sights (Volkoff et al.\\ 1971, p.41-45). Hevelius' star catalogue was studied among others by Baily (1843), and a modern comprehensive analysis was made by Rybka (1984), who confirmed that the measurements by Hevelius were superior to those by his contemporaries. Our study of the star catalogue by Brahe (Verbunt \\&\\ Van Gent 2010, hereafter Paper\\,I), showed the superiority of the modern {\\em Hipparcos Catalogue} (ESA 1997) in the analysis of old star positions, due to its better completeness, accuracy and homogeneity as compared to earlier catalogues. In this paper we present a machine-readable version of the star catalogue of Hevelius, as printed in Hevelius (1690). In addition to the numbers given by Hevelius this version provides a cross-correlation with the catalogue of Brahe; identifications with stars from the {\\em Hipparcos Catalogue} (ESA 1997) and on the basis of these the accuracy of the positions and magnitudes tabulated by Hevelius; and a comparison of our identifications with those of Rybka (1984). \\begin{table*} \\caption{Constellations in the star catalogue of Hevelius. \\label{t:const}} \\begin{tabular}{rc@{}rrr@{ }r@{ }rrl|rc@{}rrr@{ }r@{ }rrl} C & & $N$ & $N_{OT}$ & $N_K$ & $N_{SC}$ & $N_e$ & H & &C & & $N$ & $N_{OT}$ & $N_K$ & $N_{SC}$ & $N_e$ & H\\\\ \\hline 1 & And & 47 & 23 & 0 & 0 & 0& 1 &Andromeda & 31 & Leo & 50 & 40 & 0 & 0 & 1& 852 &Leo \\\\ 2 & Atn & 19 & 7 & 0 & 0 & 0& 48 &Antinous & 32 & LMi & 18 & 0 & 7 & 0 & 0& 902 &{\\bf Leo Minor} \\\\ 3 & Aqr & 48 & 41 & 0 & 0 & 1& 67 &Aquarius & 33 & Lep & 16 & 13 & 0 & 0 & 0& 920 &Lepus \\\\ 4 & Aql & 23 & 12 & 0 & 0 & 0& 115 &Aquila & 34 & Lib & 21 & 15 & 0 & 0 & 1& 936 &Libra \\\\ 5 & Ari & 27 & 21 & 0 & 0 & 0& 138 &Aries & 35 & Lyn & 19 & 0 & 4 & 0 & 0& 957 &{\\bf Lynx sive Tigris} \\\\ 6 & Aur & 40 & 27 & 2 & 0 & 0& 165 &Auriga & 36 & Lyr & 17 & 11 & 2 & 0 & 0& 976 &Lyra \\\\ 7 & Boo & 52 & 26 & 0 & 0 & 0& 205 &Bootes & 37 & Mon & 19 & 12 & 0 & 0 & 0& 993 &{\\bf Monoceros} \\\\ 8 & Cnc & 29 & 17 & 1 & 0 & 0& 257 &Cancer & 38 & Arg & 5 & 4 & 0 & 0 & 1& 1012 &Navis \\\\ 9 & CMa & 22 & 13 & 0 & 4 & 1& 286 &Canis Maior & 39 & Ori & 62 & 52 & 0 & 0 & 0& 1017 &Orion \\\\ 10 & CMi & 13 & 3 & 0 & 0 & 0& 308 &Canis Minor & 40 & Peg & 38 & 23 & 0 & 0 & 0& 1079 &Pegasus \\\\ 11 & CVn & 23 & 0 & 2 & 0 & 0& 321 &{\\bf Canes Venatici} & 41 & Per & 48 & 28 & 0 & 0 & 0& 1117 &Perseus \\\\ 12 & Cam & 32 & 0 & 14 & 0 & 0& 344 &{\\bf Camelopardalis} & 42 & Psc & 39 & 37 & 0 & 0 & 0& 1165 &Pisces \\\\ 13 & Cap & 30 & 28 & 0 & 0 & 1& 376 &Capricornus & 43 & Sge & 5 & 5 & 0 & 0 & 0& 1204 &Sagitta \\\\ 14 & Cas & 38 & 27 & 2 & 0 & 1& 406 &Cassiopeia & 44 & Sgr & 26 & 14 & 0 & 8 & 4& 1209 &Sagittarius \\\\ 15 & Cep & 51 & 11 & 4 & 0 & 0& 444 &Cepheus & 45 & Sco & 20 & 13 & 0 & 0 & 0& 1235 &Scorpius \\\\ 16 & Cer & 4 & 0 & 0 & 0 & 0& 495 &{\\bf Cerberus} & 46 & Sct & 7 & 0 & 0 & 0 & 0& 1255 &{\\bf Scutum Sobiesc.} \\\\ 17 & Cet & 46 & 22 & 1 & 0 & 1& 499 &Cetus & 47 & Oph & 44 & 23 & 1 & 4 & 4& 1262 &Serpentarius \\\\ 18 & Com & 21 & 14 & 0 & 0 & 0& 545 &Coma Berenices & 48 & Ser & 22 & 18 & 0 & 0 & 0& 1306 &Serpens \\\\ 19 & CrB & 8 & 8 & 0 & 0 & 0& 566 &Corona Borealis & 49 & Sex & 12 & 0 & 1 & 0 & 0& 1328 &{\\bf Sextans Uraniae} \\\\ 20 & Crv & 8 & 7 & 0 & 0 & 0& 574 &Corvus & 50 & Tau & 51 & 43 & 0 & 0 & 0& 1340 &Taurus \\\\ 21 & Crt & 10 & 8 & 0 & 0 & 0& 582 &Crater & 51 & Tri & 9 & 4 & 0 & 0 & 0& 1391 &Triangulum \\\\ 22 & Cyg & 47 & 20 & 0 & 0 & 0& 592 &Cygnus & 52 & TrM & 3 & 0 & 0 & 0 & 0& 1400 &{\\bf Triangulum Minus} \\\\ 23 & Del & 14 & 10 & 0 & 0 & 0& 639 &Delphinus & 53 & Vir & 50 & 39 & 0 & 0 & 0& 1403 &Virgo \\\\ 24 & Dra & 40 & 32 & 2 & 0 & 0& 653 &Draco & 54 & UMa & 73 & 34 & 1 & 0 & 0& 1453 &Ursa Maior \\\\ 25 & Equ & 6 & 4 & 0 & 0 & 0& 693 &Equuleus & 55 & UMi & 12 & 9 & 0 & 0 & 0& 1526 &Ursa Minor \\\\ 26 & Eri & 29 & 17 & 0 & 11 & 2& 699 &Eridanus & 56 & Vul & 27 & 0 & 3 & 0 & 0& 1538 &{\\bf Vulpecula} \\\\ 27 & Gem & 38 & 29 & 0 & 0 & 0& 728 &Gemini \\\\ 28 & Her & 45 & 28 & 0 & 0 & 0& 766 &Hercules \\\\ 29 & Hya & \\phantom{15}31 & 23 & 0 & 1 & 0& 811 &Hydra \\\\ 30 & Lac & 10 & 0 & 0 & 0 & 0& 842 &{\\bf Lacerta} & all & &1564 &915 & 47 & 28 & 18 \\end{tabular} \\tablefoot{For each constellation the table gives its abbreviation, the number of stars in it, and the numbers among these of stars indicated to be in \\keplere\\ $N_{OT}$, in \\keplere\\ without indication $N_K$ and in {\\em Secunda Classis} $N_{SC}$, the number of empty entries $N_e$, and the H sequence number of the first star in each constellation. New constallations by Hevelius -- partially based on Plancius -- are indicated with bold face.} \\end{table*} In the following we refer to (our machine-readable version of) the catalogue of Hevelius (1690) as \\hevelius, to Kepler's 1627 edition of Brahe's catalogue as \\kepler, and to our emended version of this edition as \\keplere. As we will see, Hevelius also refers to the {\\em Secunda Classis}, the star list that immediately follows Brahe's catalogue in Kepler (1627), and gives positions and magnitudes of those stars from the catalogue of Hipparchos/Ptolemaios that Brahe omitted from his own catalogue. Individual entries in \\hevelius\\ are numbered according to the order in which they appear, i.e.\\, H\\,350 is the 350th entry. A K-number refers to an entry in \\keplere\\ (K$\\leq$1004) or in {\\em Secunda Classis} (K$\\geq$1005). The sequence number within a constellation is indicated by a number following the abbreviated name of the constellation: Vul\\,3 is the third entry in the constellation Vulpecula in \\hevelius. ", "conclusions": "\\begin{table} \\caption{Frequency of flags B of identifications of corresponding entries in \\keplere\\ as a function of our flags I. \\label{t:kepler}} \\begin{tabular}{l|rrrrr} I\\verb+\\+B & 0 & 1 & 2 & 3 & all \\\\ \\hline 1 & 2 & 873 & 2 & 15 & 892 \\\\ 2 & 0 & 5 & 0 & 0 & 5 \\\\ 3 & 0 & 0 & 0 & 0 & 0 \\\\ 4 & 0 & 3 & 1 & 1 & 5 \\\\ 5 & 2 & 0 & 0 & 0 & 2 \\\\ 6 & 0 & 1 & 0 & 0 & 1 \\\\ all & 4 & 882 & 3 & 16 & 905 \\end{tabular} \\end{table} \\subsection{Comparison with Brahe\\label{s:brahe}} 915 entries in \\hevelius\\ can be matched unambigously with an entry in \\keplere: 911 (including the repeat entry H\\,793) through their {\\em Ordo Tychonis}, and 4 without OT through the position according to Tycho as given by Hevelius. 10 of these 914 independent entries have no position by Hevelius, among them SN\\,1572. Table\\,\\ref{t:kepler} compares our identifications of the entries common to \\hevelius\\ and \\keplere: in most cases we found the same identification for both catalogues. The 16 exceptions arise because Hevelius gives a position very different from the position in \\keplere: striking examples are H\\,14/K\\,463, H\\,1005/K\\,961 and H\\,1099/K\\,441, as detailed in Sect.\\,\\ref{s:notes}. 90 entries in \\keplere, among which 12 which we were unable to identify with an Hipparcos star, have no explicit counterpart in \\hevelius. The 12 unidentified stars probably have a wrong position in \\keplere, which would explain why Hevelius found no star at that position. For the others we checked whether their Hipparcos identifications occur also in \\hevelius. This is the case for 57. It appears likely that Hevelius and Brahe observed the same star in these cases, but we cannot exclude an occasional chance coincidence. Many of the remaining 21 entries in \\keplere\\ that are not matched with an entry in \\hevelius\\ have a very large position error $\\Delta$ in \\keplere. The unmatched entries also include 4 stars in Argo and all 4 stars in Centaurus. Some remarkable features of entries in \\keplere\\ are present in \\hevelius\\ as well. We mention in particular the three stars in Capricornus which are denoted `nebulous' both in Hevelius and Brahe, and H\\,1188/K\\,801 which both in \\kepler\\ and \\hevelius\\ are given a B for northern latitude, whereas the correct latitude is S for southern (see Fig.\\,C.44 in Paper\\,I). \\begin{table} \\caption{Frequency of flags R of identifications by Rybka as a function of our flags I. \\label{t:rybka}} \\begin{tabular}{l|rrrrr|rrrrr} & \\multicolumn{5}{c}{in KeplerE} & \\multicolumn{5}{c}{full catalogue} \\\\ I\\verb+\\+R & 0 & 1 & 2 & 3 & all & 0 & 1 & 2 & 3 & all \\\\ \\hline 1 & 5 & 880 & 3 & 4 & 892 & 20 & 1426 & 4 & 29 & 1479 \\\\ 2 & 0 & 5 & 0 & 0 & 5 & 0 & 10 & 0 & 1 & 11 \\\\ 3 & 0 & 0 & 0 & 0 & 0 & 2 & 9 & 0 & 4 & 15 \\\\ 4 & 0 & 2 & 1 & 2 & 5 & 1 & 5 & 2 & 4 & 12 \\\\ 5 & 0 & 0 & 0 & 2 & 2 & 7 & 0 & 0 & 9 & 16 \\\\ 6 & 0 & 1 & 0 & 0 & 1 & 1 & 11 & 0 & 1 & 13 \\\\ all & 5 & 888 & 4 & 8 & 905 & 31 & 1461 & 6 & 48 & 1546 \\end{tabular} \\tablefoot{Left: for stars in \\keplere, right: for all entries in \\hevelius.} \\end{table} \\subsection{Comparison with Rybka (1984)} In Table\\,\\ref{t:rybka} we compare the identifications as found by us with those given by Rybka (1984), separately for the stars in \\keplere\\ and for all stars in \\hevelius. In most cases the identifications are identical, but there are differences. We have identified 24 stars (among which 1 repeated entry) that Rybka could not identify. In 6 cases where two stars are plausible counterparts we choose the stars that Rybka did not choose. In 48 cases (among which 1 repeated entry) we do not agree with the identification in Rybka (1984); this includes 9 stars which we cannot identify. This number does {\\em not} include the 21 cases where our emendation to Rybka leads to a different Hipparcos identification (see Sect.\\,\\ref{s:rybkae}). In some cases the identification given by Rybka (1984) has such a large positional offset, or is so faint, that we consider our rejection secure; in other cases we choose another closer and/or brighter star as a more plausible counterpart. Details may be found in Sect.\\,\\ref{s:notes}. \\subsection{Accuracy} \\begin{figure} \\includegraphics[angle=270,width=\\columnwidth]{keplerv.pdf} \\includegraphics[angle=270,width=\\columnwidth]{heveliusv.pdf} \\caption{Distribution of the magnitudes for all stars in \\hevelius\\ (below), and for only those stars that have a counterpart in \\keplere\\ (above). In the large frames the histograms indicate the magnitudes according to Hevelius for stars which we have securely identified (red; flags 1-2 ) or not securely identified (blue, flags 3-5), and the magnitudes from the {\\em Hipparcos} catalogue for securely identified stars (black). The numbers of securely and not-securely identified stars are indicated. The small frames give the {\\em Hipparcos} magnitude distributions for securely identified stars for each magnitude according to Hevelius separately. The number of securely identified stars at each (Hevelius) magnitude is indicated. \\label{f:magnitudes}} \\end{figure} Table\\,\\ref{t:rybka} shows that there are 16 stars, i.e.\\ one percent of the total, in \\hevelius\\ that we are not able to identify. We do not count in these the empty entries. As in the case of the Brahe catalogue, we would have to accept fainter counterparts or larger position errors to identify these; in both cases the probability of chance coincidences would increase. Other entries in \\hevelius\\ which do not have an identification in the {\\em Hipparcos Catalogue} are H\\,32 (= M\\,31), H\\,259 (= Praesepe) and H\\,1540 (= Nova Vul 1670; see Sect.\\,\\ref{s:notes}). \\begin{figure} \\includegraphics[angle=270,width=0.48\\columnwidth]{hlamdlam.pdf} \\includegraphics[angle=270,width=0.48\\columnwidth]{hlamdbet.pdf} \\includegraphics[angle=270,width=\\columnwidth]{hdlondlat.pdf} \\caption{Above: Correlations of the differences in longitude $\\Delta\\lambda\\equiv(\\lambda_\\mathrm{HIP}-\\lambda)\\cos\\beta$ and latitude $\\Delta\\beta\\equiv\\beta_\\mathrm{HIP}-\\beta$ of the entries in \\hevelius\\ and their secure {\\em Hipparcos} counterparts (converted to 1661) with longitude. Below left: Distributions of $\\Delta\\lambda$ and $\\Delta\\beta$. The numbers of sources with $\\Delta\\lambda,\\Delta\\beta<-10\\arcmin$, of sources included in the histogram ($-10\\arcmin<\\Delta\\lambda,\\Delta\\beta<10\\arcmin$), and of sources with $\\Delta\\lambda,\\Delta\\beta>10\\arcmin$ are indicated. Below right: correlation between $\\Delta\\lambda$ and $\\Delta\\beta$. \\label{f:dlongdlat}} \\end{figure} Figure\\,\\ref{f:magnitudes} illustrates that the magnitudes assigned by Hevelius correlate well with those of their counterparts in the {\\em Hipparcos Catalogue}. Only the higest magnitudes assigned by Hevelius, 6 and 7, tend to be too high. \\begin{figure} \\includegraphics[angle=270,width=\\columnwidth]{alphahdb.pdf} \\caption{Distribution of the errors in latitude in \\hevelius\\ as a function of right ascension, together with the r.h.s.\\ of Eq.\\,\\ref{e:alphadb}. \\label{f:alphahdb}} \\end{figure} In Figure\\,\\ref{f:dlongdlat} we show the error distributions separately for the longitudes and latitudes, as well as their correlation. The correlation distribution is roughly spherical, i.e.\\ the errors in longitude and latitude are mostly independent. Gaussians that fit the central regions ($-$5\\arcmin,+5\\arcmin) of the distributions of $\\Delta\\lambda$ and $\\Delta\\beta$ both have widths $\\sigma\\simeq2$\\arcmin; both predict fewer points at errors larger than 5\\arcmin\\ than observed. The numbers of errors with absolute values larger than 10\\arcmin\\ correspond to less than 10\\%\\ of the number of identified entries, a similar percentage as in \\keplere. The widths of the peak of the error distributions (near 2\\arcmin) and the fraction of larger errors are thus similar in \\hevelius\\ to those in \\keplere, which is an impressive achievement since the number of stars has increased by more than 50\\%\\, mostly at the fainter magnitudes 5 and 6. Figure\\,\\ref{f:dlongdlat} further shows that the errors in longitude $\\Delta\\lambda$ increase with the distance to the zero point \\aries; and that the errors in latitude $\\Delta\\beta$ have a roughly sinusoidal dependence on longitude. The latter dependence may be explained by an error in the value of the obliquity. Hevelius measured the obliquity in several years, and found values around $\\epsilon_H$=23\\fdg506 (Rybka 1984, p.37); according to modern theory the obliquity in 1661 was $\\epsilon$=23\\fdg483. For small declinations $\\delta$, the resulting error in latitude $\\Delta\\beta\\equiv\\beta-\\beta_H$ due to the error $\\Delta\\epsilon\\equiv\\epsilon-\\epsilon_H$ after converting equatorial to ecliptic coordinates may be written \\begin{equation} \\cos\\beta\\Delta\\beta \\simeq -\\sin\\alpha\\cos\\epsilon\\Delta\\epsilon\\simeq 1\\farcm3\\sin\\alpha \\label{e:alphadb}\\end{equation} The observed relation between $\\Delta\\beta$ and $\\alpha$ is shown in Fig.\\,\\ref{f:alphahdb} together with the curve 1\\farcm3$\\sin\\alpha$ which roughly matches the phase and amplitude of the $\\alpha$-dependence of $\\Delta\\beta$. \\begin{figure} \\includegraphics[angle=270,width=\\columnwidth]{HDelta.pdf} \\caption{Distribution of the position errors $\\Delta$ for all stars in \\hevelius\\ (open histogram) and for only those stars that have a counterpart in \\keplere\\ (solid histograms) for all securely identified stars (left) and for the securely identified stars at each Hevelius magnitude separately (right). The numbers indicate the number of stars included in the plot (i.e.\\ with $\\Delta<10\\arcmin$) and those excluded ($\\Delta\\ge10\\arcmin$) \\label{f:delta}} \\end{figure} \\begin{figure} \\includegraphics[angle=270,width=\\columnwidth]{ksplots.pdf} \\caption{Cumulative distribution of the position errors $\\Delta$ in \\hevelius\\ as a function of Hevelius magnitude, showing a systematic increase in median error with magnitude.\\label{f:ksplots}} \\end{figure} The average offset of longitude is virtually zero; the latitudes have an average offset of -1\\farcm4. This average offset may be due to an underestimate by Hevelius of refraction. The distribution of the total position errors $\\Delta$ in \\hevelius\\ is shown in Fig.\\,\\ref{f:delta}. This distribution peaks roughly at the value of the width 2\\arcmin\\ of the separate distributions in $\\Delta\\lambda$, $\\Delta\\beta$, as expected (see explanation in Paper\\,I). The number of stars with large position errors is markedly smaller in \\hevelius\\ than in \\keplere. In particular, the number of stars with position errors larger than a degree is 21 (on a total of 1517 identified entries) in \\hevelius\\ as compared to 47 (on a total of 977 identified entries) in \\keplere. Similarly, the number of unidentified stars is 16 (of 1533 independent entries) in \\hevelius\\ and 14 (of 992 independent entries) in \\keplere. It may be concluded that the overall accuracy of the star catalogue of Hevelius is better than that of the star catalogue of Brahe/Kepler. \\begin{figure} \\includegraphics[angle=270,width=\\columnwidth]{complet.pdf} \\caption{Completeness of \\hevelius\\ and \\keplere\\ as a function of magnitude and declination, as illustrated by cumulative magnitude distributions. For each range of declination (equinox 1631) the top curve shows the magnitudes for all stars in the {\\em Hipparcos Catalogue}, the middle curve the Hipparcos stars {\\em not} in \\keplere, and the lower curve the Hipparcos stars {\\em not} in \\hevelius. \\label{f:complet}} \\end{figure} In Fig.\\,\\ref{f:ksplots} we show the cumulative error distributions for each Hevelius magnitude separately, taking magnitudes 1 and 2 together, and limiting the distributions to $\\Delta<10$\\arcmin. It is seen that the median error increases slowly but systematically with magnitude. \\subsection{New and old stars: completeness} In Fig.\\,\\ref{f:complet} we investigate the completeness of \\hevelius\\ and \\keplere\\ as function of magnitude, for three declination ranges. For this purpose we select from \\hevelius\\ and \\keplere\\ only those entries which we have identified, and which are not repeat entries, i.e.\\ entries with I=1-4. For selecting the Hipparcos stars in the latitude ranges we convert their positions to an equinox halfway between Brahe and Hevelius, viz.\\ 1631.0. At magnitudes $V$$<$4 there are 348 stars from the {\\em Hipparcos Catalogue} with $\\delta>-30^\\circ$, of which 5 are absent from \\hevelius\\, and 23 from \\keplere\\ (of which 13 with $\\delta<0^\\circ$). At magnitudes $V$$<$5 there are 1138 Hipparcos stars with $\\delta>-30^\\circ$, of which \\hevelius\\ misses 141 stars (89 with $\\delta<0^\\circ$) and \\keplere\\ 389 (156 with $\\delta<0^\\circ$). Finally, about 3500 Hipparcos stars with $V$$<$6 have $\\delta>-30^\\circ$, and of these some 2000 are absent from \\hevelius\\ and 2500 from \\keplere, which is just another way of saying that \\hevelius\\ and \\keplere\\ contain about 1500 and 1000 stars visible to the naked eye, respectively. It may be noted here that the latitude of Gdansk is about 1\\fdg5 further south than that of Hven. Nonetheless, as shown in Fig.\\,\\ref{f:complet} \\keplere\\ is more incomplete already at brighter magnitudes, also in the northern parts of the sky. How many new stars did Hevelius observe? In the manuscript of the catalogue, a note dated 1681 March 31 states that 946 stars of Tycho and 617 new stars were observed (Volkoff et al.\\ 1971, p.72). This gives a total of 1563, very close to the total of 1564 entries given in Table\\,\\ref{t:const}, but spuriously so since Hevelius gives no own measurements for 18 of the 1564 entries. Hevelius indicates, through an OT number or a position from Tycho, for 915 entries that they are from \\kepler\\ (Table\\,\\ref{t:const}). For 905 of these \\hevelius\\ gives his own measurements (see Sect.\\,\\ref{s:brahe}). To obtain the higher number of 946 stars from the note, we have two options. One is to add the 28 stars from the {\\em Secunda Classis}. (These were also measured by Brahe, according to Kepler, albeit with less accuracy.) This option leaves us with too small a number. The other option is to add the 47 entries for which our identification corresponds to an identification in \\keplere. This would imply that Hevelius was aware that more stars from his catalogue corresponded to stars in \\kepler\\ than the 915 entries marked by him as such through OT or Tycho position. If we subtract from the total number of 1564 entries in \\hevelius\\ all 962 that have a counterpart in \\keplere\\ and further subtract the 28 entries that have a counterpart in {\\em Secunda Classis}, we find a number of 574 entries first measured by Hevelius. 573 of these are stars, the other one is M\\,31. Cerberus, Lacerta, Scutum, Sextans and Triangulum Minus, the truly new constellations by Hevelius, contain a total of 36 stars, of which only one possibly corresponds to a star in \\keplere\\ (H\\,1328 / K959: the position error of K\\,959 is 2\\fdg5, so a chance coincidence is possible). Monoceros and Camelopardalis, two constellations retained by Hevelius from Plancius, contain twelve and up to fourteen stars from \\keplere, respectively (Table\\,\\ref{t:const}). The four constellations fashioned by Hevelius from two constellations by Plancius contain up to sixteen stars from \\keplere. We use the qualification `up to' for the stars listed under $N_K$ in Table\\,\\ref{t:const} because some of the correspondences between \\hevelius\\ and \\keplere\\ may be chance coincidences. The six constellations retained of refashioned from Plancius by Hevelius contain 138 stars, so even accepting all 42 correspondences as real, we still find that a large majority of stars in these constellations was first observed by Hevelius." }, "1003/1003.1902_arXiv.txt": { "abstract": "{} {We gathered about 100 high-resolution spectra of three typical HgMn (mercury-manganese) stars, \\object{HD\\,11753}, \\object{HD\\,53244}, and \\object{HD\\,221507}, to search for slowly pulsating B-like pulsations and surface inhomogeneous distribution of various chemical elements. } {Classical frequency analysis methods were used to detect line profile variability and to determine the variation period. Doppler imaging reconstruction was performed to obtain abundance maps of chemical elements on the stellar surface. } {For \\object{HD\\,11753}, which is the star with the most pronounced variability, distinct spectral line profile changes were detected for Ti, Sr, Y, Zr, and Hg, whereas for \\object{HD\\,53244} and \\object{HD\\,221507} the most variable line profiles belong to the elements Hg and Y, respectively. We derived rotation periods for all three stars from the variations of radial velocities and equivalent widths of spectral lines belonging to inhomogeneously distributed elements: P$_{rot}$ (\\object{HD\\,11753})=9.54\\,d, P$_{rot}$ (\\object{HD\\,53244})=6.16\\,d, and P$_{rot}$ (\\object{HD\\,221507})=1.93\\,d. For \\object{HD\\,11753} the Doppler imaging technique was applied to derive the distribution of the most variable elements Ti, Sr, and Y using two datasets separated by $\\sim$65 days. Results of Doppler imaging reconstruction revealed noticeable changes in the surface distributions of \\ion{Ti}{ii}, \\ion{Sr}{ii}, and \\ion{Y}{ii} between the datasets, indicating the hitherto not well understood physical processes in stars with radiative envelopes that cause a rather fast dynamical chemical spot evolution.} {} ", "introduction": "The mercury-manganese (HgMn) stars constitute a well-defined sub-class of chemically peculiar (CP) stars of the B7--B9 spectral types with $T_{\\rm eff}$ between 10\\,000 and 15\\,000\\,K. These stars exhibit marked abundance anomalies of several elements: e.g., overabundances of Hg, Mn, Ga, Y, Cu, Be, P, Bi, Sr, Zr, and deficiencies of He, Al, Zn, Ni, Co. More than two thirds of them belong to spectroscopic binaries (Hubrig \\& Mathys \\cite{hubrig_mathys95}). They are slow rotators ($\\langle v\\,\\sin i\\rangle=29$~km$\\,$s$^{-1}$, Abt \\cite{abt72}). There is no evidence that they would have strong large-scaled organized magnetic fields. Their elemental overabundances/underabundances are believed to be due to radiatively-driven diffusion and gravitational settling. In the H-R diagram many HgMn stars are located in the instability strip of the so-called slowly pulsating B (SPB) stars (see De Cat \\cite{decat03} for a review on the latter), and the sophisticated models predict that pulsations should also be driven in HgMn stars (Turcotte \\& Richard \\cite{turcotte_richard05}). Searches for variability in this group of stars have been made mostly photometrically in the past, but without any success. Very recently, Alecian et al.\\ (\\cite{alecian09}) discovered low amplitude (less than 1.6\\,mmag) periodic variations (4.3 and 2.53~days respectively, with harmonics) in two candidate HgMn stars by means of the high quality light curves provided by the CoRoT satellite. These variations are compatible with theoretically predicted pulsation periods. However, as stated by the authors, only spectroscopic datasets could help to conclusively establish or withdraw this pulsation interpretation. Currently, there is thus no observational proof of pulsation in HgMn stars. The aspect of inhomogeneous distribution of some chemical elements over the surface of HgMn stars was for the first time discussed by Hubrig \\& Mathys (\\cite{hubrig_mathys95}). From the survey of HgMn stars in close spectroscopic binaries (SBs) it was suggested that some chemical elements might be inhomogeneously distributed on the surface, with in particular a preferential concentration of Hg along the equator. In close double-lined systems (SB2s), where the orbital plane has a small inclination to the line of sight, a rather large overabundance of Hg was found. By contrast, in stars with orbits almost perpendicular to the line of sight, mercury is not observed at all. The first indication of variability of the \\ion{Hg}{II} 3984\\,\\AA{} and \\ion{Y}{II} 3983\\,\\AA{} lines was reported for the HgMn SB2 system \\object{AR\\,Aur} by Takeda et al.\\ (\\cite{Takeda79}). Later, Wahlgren et al.\\ (\\cite{wahlgren01}) and Adelman et al.\\ (\\cite{adelman02}) showed that the \\ion{Hg}{II} 3984\\,\\AA{} line of the primary component of \\object{$\\alpha$~And} varies with a 2.8-d period. The spectral line variations were attributed to the surface inhomogeneous mercury distribution along the stellar equator, together with the stellar rotation period. Recently, Kochukhov et al.\\ (\\cite{kochukhov05}) found clear signatures of surface mercury spots in two rapidly rotating HgMn stars by analysing the \\ion{Hg}{II} 3984\\,\\AA{} line profiles. Variability of spectral lines associated to larger number of chemical elements were discovered for the first time by Hubrig et al.\\ (\\cite{hubrig06a}) for the primary component of the eclipsing binary \\object{AR Aur}. The strongest variations were found for the chemical elements Pt, Hg, Sr, Y, Zr, He, and Nd. The first Doppler maps for the elements Mn, Sr, Y, and Hg were recently presented by Savanov et al.\\ (\\cite{sav09}). The study of Hubrig et al.\\ (\\cite{hubrig08}) suggests that spectral variability of various chemical elements is indeed observed in most HgMn stars. We present the first observational study based on a substantial number of spectra, more than one hundred, obtained with the CORALIE \\'echelle spectrograph attached to the 1.2m Leonard Euler telescope in La Silla in Chile. The selected targets, the single-lined (SB1) spectroscopic binaries \\object{HD\\,11753} ($\\phi$\\,Phe, V = 5.1 mag, B\\,9p) and \\object{HD\\,53244} ($\\gamma$\\,CMa, V = 4.1 mag, B8\\,II), and the star \\object{HD\\,221507} ($\\beta$\\,Scl, V = 4.4 mag, B9.5\\,IVmnpe) were chosen as the brightest known southern HgMn stars visible during the periods of observation. The goal of the presented spectroscopic study was twofold: to search for stellar pulsations and/or surface inhomogeneous distribution of various chemical elements. Our observations and data reduction are presented in the appendix A. \\begin{table} \\caption[]{Atmospheric parameters (effective temperature, logarithm of the gravity, microturbulent velocity, and projected rotational velocity) taken from Dolk et al.\\ (\\cite{dolk2003}) for \\object{HD\\,11753} and \\object{HD\\,221507}, and from Woolf \\& Lambert (\\cite{woolf_lambert99}) for \\object{HD\\,53244}.} \\begin{center} \\begin{tabular}{ccccc} \\hline \\hline\\\\[-7pt] & T$_{\\rm eff}$ (K) & $\\log g$ & $\\xi$ (km s$^{-1}$) & $v \\sin i$ (km s$^{-1}$) \\\\[5pt] \\hline\\\\[-7pt] \\object{HD\\,11753} & 10\\ 612$\\pm$200 & 3.79$\\pm$0.10 & 0.5 $\\pm$0.5 & 14$\\pm$0.5\\\\ \\object{HD\\, 53244} & 13\\ 600$\\pm$200 & 3.40$\\pm$0.03 & 2.0 $\\pm$0.5 & 35$\\pm$0.5\\\\ \\object{HD\\,221507} & 12\\ 476$\\pm$200 & 4.13$\\pm$0.10 & 0.0 $\\pm$0.5 & 25$\\pm$0.5\\\\[5pt] \\hline \\end{tabular} \\end{center} \\label{par} \\end{table} \\begin{figure} \\centering \\includegraphics[angle=270,totalheight=0.27\\textwidth]{13775fg1.ps} \\caption{Position of the studied stars in a \\mbox{($\\log T_{\\rm eff}$,$\\log g$)}-diagram. From left to right, the error boxes represent \\object{HD\\,53244}, \\object{HD\\,221507}, and \\object{HD\\,11753}. The full and dashed lines represent the boundary of the theoretical SPB instability strip for a metallicity $Z=0.02$ and $Z=0.01$, respectively (taken from Miglio et al.\\ \\cite{miglio07}).} \\label{HR} \\end{figure} \\begin{figure} \\centering \\includegraphics[angle=270,totalheight=0.35\\textwidth]{13775fg2.eps} \\caption{ Spots on the surface of \\object{HD\\,11753}: time-series spectra phased on the stellar rotation period of 9.54~days. } \\label{HD11753} \\end{figure} \\begin{figure} \\centering \\includegraphics[bb= 50 30 527 720, width=4cm, height=7cm, angle=270]{13775fg3.eps} \\caption{Left: Inhomogeneous distribution of \\ion{Hg}{ii} on the surface of \\object{HD\\,53244} apparent in the time-series spectra around \\ion{Hg}{ii} $\\lambda$3984 phased on the rotation period of 6.16 days. Right: Inhomogeneous distribution of \\ion{Y}{ii} on the surface of \\object{HD\\,221507} apparent in time-series spectra around the spectral line of \\ion{Y}{ii} $\\lambda$4310 phased on the rotation period of 1.93 days.} \\label{HD53244} \\end{figure} ", "conclusions": "All Ti, Sr, and Y abundance maps reveal a structure reminiscent of broken rings of low and high abundance. This elemental distribution is to some extent similar to the maps previously reconstructed for another HgMn star, \\object{AR\\,Aur} (Savanov et al.\\ \\cite{sav09}), where the elements Mn, Y, Sr, and Hg show abundance concentration in equatorial and polar features. Typically, inhomogeneous chemical abundance distributions are observed only on the surface of magnetic chemically peculiar stars with large-scale organised magnetic fields. In these stars, the abundance distribution of certain elements is non-uniform and non-symmetric with respect to the rotation axis. A magnetic field of the order of a few hundred Gauss was detected in hydrogen lines of four HgMn stars by Hubrig et al.\\ (\\cite{hubrig06b}) using low-resolution ($R=2000$) circular polarisation spectra obtained with FORS\\,1 at the VLT. This small sample of HgMn stars also included the spectrum variable HgMn star \\object{$\\alpha$~And} with a magnetic field of the order of a few hundred Gauss. On the other hand, high-resolution spectropolarimetric spectra of some HgMn stars, including \\object{$\\alpha$~And}, were used in studies of Shorlin et al. (\\cite{shorlin02}) and Wade et al. (\\cite{wade06}), where no detection was achieved using all metal lines together in the least-squares deconvolution multi-line profile. Although strong large-scale magnetic fields have not generally been found in HgMn stars, it has never been ruled out that these stars might have tangled magnetic fields of the order of a few thousand Gauss with no net longitudinal component (e.g., Mathys \\& Hubrig \\cite{mat95}; Hubrig et al.\\ \\cite{hubrig99}; Hubrig \\& Castelli \\cite{hubrig01}). It is of interest that magnetohydrodynamical simulations by Arlt et al. \\ (\\cite{arlt03}), which combine a poloidal magnetic field and differential rotation can produce a magnetic field topology that is similar to the broken elemental ring structures seen in \\object{HD\\,11753} and \\object{AR\\,Aur}. These simulations and their implication have been recently discussed by Hubrig et al.\\ (\\cite{hubrig08}). The abundance maps of \\object{HD\\,11753} presented in this work exhibit clear differences between the surface abundance distribution of Ti, Sr, and Y. We also detected clear differences in the spot configurations obtained from the same lines but for different data sets, which indicates a rather fast dynamical evolution of the abundance distribution with time. Kochukhov et al.\\ (\\cite{kochukhov07}) discovered mercury clouds in the atmosphere of a HgMn star \\object{$\\alpha$~And} that showed secular changes with a time period of 2--4 years. In our analysis, using two datasets separated by $\\sim$65 days, we reveal that the changes in the chemical spot configuration of \\object{HD\\,11753} appear much faster and can already be detected at a time scale of months. The results reported in this paper open up new perspectives for our knowledge and understanding of HgMn stars. Different dynamical processes take place in stellar radiation zones. An interaction between the differential rotation, the magnetic field, and the meridional circulation could possibly play a role in the generation of dynamical evolution of chemical spots. From the comparison of maps we find that it is possible that the Y and Sr distributions show indications of an increasing rotation rate towards the rotation pole, so-called differential rotation of anti-solar type. On the other hand, further analyses of the elemental surface distribution in a larger sample of HgMn stars should be carried out before the implication of these new results can be discussed in more detail." }, "1003/1003.1311_arXiv.txt": { "abstract": "Reconstructing the matter density field from galaxy counts is a problem frequently addressed in current literature. Two main sources of error are shot noise from galaxy counts and insufficient knowledge of the correct galaxy position caused by peculiar velocities and redshift measurement uncertainty. Here we address the reconstruction problem of a Poissonian sampled log-normal density field with velocity distortions in a Bayesian way via a maximum a posteriory method. We test our algorithm on a 1D toy case and find significant improvement compared to simple data inversion. In particular, we address the following problems: photometric redshifts, mapping of extended sources in coded mask systems, real space reconstruction from redshift space galaxy distribution and combined analysis of data with different point spread functions. ", "introduction": "Correct and thorough signal analysis is vital in cosmology. This is for one thing due to often low signal-to-noise ratios of many cosmic measurements. An important fact to notice here is that many measurements can \\emph{not} be independently repeated, as nature only grants us with one realisation of the data such as the CMB or the large-scale structure of the universe (LSS). Since the complications in extracting the desired information from data are so fundamental, it is often not possible to draw sensible conclusions from it without some knowledge on the properties of the underlying signals. Although it would be desirable to be completely independent and prejudice-free, one typically has to give up on some freedom in the analysis by restricting to a specific model. In return one gains much better constraints on the measured quantity. Hence a thorough signal analysis must take care of all those aspects, otherwise the wrong conclusions could be drawn. This is most naturally done in the Bayesian framework where all variables are considered to be subject to error and variance. The emphasis in this work is on the reconstruction problem of a log-normal density field that is sampled via a Poissonian process as simple description of galaxy formation and a number of other processes, like a highly structured gamma ray emissivity. We choose the log-normal field, because we think it suited to model the dark matter density distribution of the Universe (see section \\ref{section:log-normal-density field}). In addition, we extend the problem such that the signal is spatially distorted and galaxy counts from one position may show up at other locations. This way a sharp peak in the signal can show up as a broad distribution in the data depending on the underlying distortion process. This allows us to address the real space reconstruction problem of the dark matter density field from redshift distorted galaxy counts and to naturally incorporate photometric redshift errors in our analysis, to name a few. The most generic case of uncertain position is the measurement error from the measurement apparatus itself which comes with \\emph{any} measurement. In many cases, these errors form a Gaussian distribution around a mean value. But there are also other cases like photometric redshift where due to the measurement technique there is considerable chance for `catastrophic outliers' which leads to non-Gaussian probability distributions. In our analysis the chance for such catastrophic errors can be naturally included and dealt with. This permits do deal with cases where such outliers are the rule, for example coded mask detectors in X- and $\\gamma$-ray astronomy where a point source has to be identified via its complex mask shadow on the detector plane. Other areas where spatial distortion should be included in the analysis are $\\gamma$-ray astronomy via \\v Cerenkov telescopes, or Ultra-high-energy cosmic ray detectors due to the extended point spread functions of the measurement devices. In all examples so far the distortion of the data was known a priori and was fixed. However, one can go one step further and allow the distortion to depend on the signal itself that one wants to measure. The paradigm for this problem is the measurement of redshift in galaxy surveys. Here the aim is to measure the real space density distribution of matter in the Universe via galaxy counts, but since the presence of matter has an effect on the peculiar velocities of the observed galaxies, the distortion depends on the details of the signal to be reconstructed. Beyond the linear regime, where a dark matter halo has collapsed to a virialised object like a galaxy cluster, the galaxies have large peculiar velocities. Since only the component along the line of sight adds to the redshift, those collapsed objects appear as dense elongated structures in redshift space pointing towards the observer, therefore this is also called the `\\FOG' effect. Including the feedback of matter on the redshift space distortions is the ultimate goal of LSS reconstruction. The point of our work is to address one very important effect so far often ignored in statistical inference of the LSS: spatial distortions. In order to focus our discussion on this, we work with simplified descriptions of the complex galaxy formation process. The adopted description, however, was previously shown to provide good reconstructions despite its simplicity. Significant progress in the field of large-scale structure reconstruction with a log-normal model for the dark matter over-density has been achieved in a number of recent works. \\cite{2009PhRvD..80j5005E} derive the MAP estimator and loop corrections thereof within the \\emph{information field theory} (IFT) framework. \\citet{2010MNRAS.403..589K} successfully apply the MAP Poissonian log-normal filter on mock data from N-body simulations. \\citet{2009arXiv0911.2498J} even achieved a reconstruction from real world SDSS\\,7 data going beyond the MAP approximation by using a Hamiltonian sampling method. However, all these approaches do not take the spatial uncertainty of redshift measurements -- i.e.~the point spread function (PSF) -- into account. The complications from a non-trivial PSF has been addressed in a number of works with similar data models as ours. \\cite{1992ITSP...40.2290H,Green90bayesianreconstructions,2008PMB....53..593W} consider a similar data model but use a signal clique prior adequate for image reconstruction. \\cite{2009OptCo.282.2489O} have a distorted Poissonian data model but use a smoothing prior based on Fisher information for the signal. \\cite{1990Ap&SS.171..341N} also work with a Poissonian data model but use an ad-hoc image entropy as prior for their signal. \\cite{1978JOSA...68...93F,1978Natur.272..686G} work with maximum entropy prior for the signal distorted by a point spread function and approximate the Poissonian distribution by a Gaussian. Different choices for such entropy priors are discussed in \\cite{1982JApA....3..419N,1985A&A...143...77C}. The outline of this work is as follows. We formulate the Bayesian reconstruction problem in a field-theoretic language in section \\ref{section:theoretic-background}. In section \\ref{chapter:reconstruction-of-spatially-distorted-signals} we address the reconstruction problem from spatially distorted data. A suitable data-model for this purpose is introduced. In particular this includes a discussion of the distortion operator in section \\ref{section:the-distortion-operator}. We address how approximate error bars can be constructed for our reconstructions. In section \\ref{section:photometric-redshift} we apply our method to the LSS reconstruction from photometric redshift measurement in a 1D test-case. We show how data sets with different error characteristics can be naturally combined within our framework in the subsequent section. In section \\ref{section:coded-mask} we apply the same algorithm to a completely different field, namely X- and $\\gamma$-ray astronomy via coded mask telescopes. Finally, in section \\ref{section:s-dependent-distortion} we formulate a distortion problem where the distortion itself depends on the signal to be reconstructed. We therefore propose a model how to approach the forward problem to transform from real into redshift space. We compare our results for this distortion model to a Metropolis-Hastings sampling method in section \\ref{section:metropolis-hastings}. In \\ref{section:conclusions} we summarise our findings. Details about the notation can be found in \\ref{appendix:notation}. ", "conclusions": "\\label{section:conclusions} Many reconstruction problems in cosmology suffer not only from large noise but also from substantial measurement uncertainties. While it is possible that some measurement uncertainties will be ameliorated in the future by more sophisticated techniques, other sources of uncertainty are fundamental such as cosmic variance and galaxy redshift distributions. Areas where this applies are real space LSS reconstructions from galaxy counts in redshift space, but also consistent treatment of photometric redshift. For precision cosmology with galaxies it is therefore of paramount importance to incorporate these uncertainties in the analysis. Here we have presented a novel method how spatially distorted log-normal fields as they occur in density field reconstruction can be reconstructed in a Bayesian way. This method was developed in the framework of information field theory which we outlined in section \\ref{section:theoretic-background}. We showed that the IFT moment calculation ultimately foots on the minimization of the expected $L_2$-weighted error of the reconstruction. Where exact moment calculation from the posterior was not possible, we argued how the correct map -- the posterior mean -- could be approximated by a MAP approach. We developed a data model for a log-normal signal with Poissonian noise where the response can be non-local. We even allowed for the case, in which the distortion of data space could depend on the signal that was to be reconstructed. The resulting problem is so complex that it could only be solved approximately via numerical minimization of its probability Hamiltonian. For a test of our approach we performed simulations where we constructed mock signals, produced mock data thereof and tried to reconstruct the underlying signal by numerical minimization of the probability Hamiltonian. We tested our reconstruction code on three different distortion problems which were \\begin{itemize} \\item data with typical distortions as they appear in photometric redshift measurement, \\item coded mask aperture problems as they appear in X- and $\\gamma$-ray astronomy and \\item real space matter reconstruction from redshift distorted data. \\end{itemize} For the latter we developed a model for the forward problem to construct redshift space data from real space galaxy distributions and where the distortion was dependent on the underlying matter distribution that was to be measured. We were able to tackle this problem with a MAP approach. However, after further complication of the distortion operator we found that the MAP method does not live up to its expectations. Instead, we could show that approximating the posterior via Metropolis-Hastings sampling could give much more accurate reconstructions. Therefore we think, that for such complicated problems the MAP method gives misleading results and should be superseded by more powerful however also computationally more demanding approaches such as sampling the posterior \\PDF. For the coded mask data we were able to identify the largest peaks and showed that it is even possible to reconstruct their substructure if the count rates are high enough. An application of this approach to real X- or $\\gamma$-ray data should be possible but before doing so, some effort must be spent to make the approach robuster to false detection of peaks. At last, the reconstruction of a redshift space signal from photometric redshift data proved to be very fruitful. In many cases we were able to reconstruct the underlying matter distribution remarkably well. Since the colour space distortion is independent of the underlying signal, an application of our approach to large data sets is feasible. We also showed that in the IFT framework it is possible to easily combine data sets with different error characteristics. We considered the problem of combining photometric redshift data with large uncertainties and spectroscopic data that are very accurate in position. Our analysis showed that even with a low abundance of accurate data it is possible to improve the reconstruction from distorted data with large abundance as long as there is room for improvement. In all cases we found that Bayesian analysis of the problem is inevitable for the noise level we were considering. We also showed that the reconstruction becomes significantly worse when the data were distorted, but the data space distortion was neglected during the reconstruction. Therefore we think that including the data space distortions in future precision analysis is inevitable. Since the assumptions of our method are based on a few generic principles we are confident that further areas will be found where our work will be appreciated. \\appendix" }, "1003/1003.4228_arXiv.txt": { "abstract": "% Although there are now some tentative signs that the start of cycle 24 has begun there is still considerable interest in the somewhat unusual behaviour of the current solar minimum and the apparent delay in the true start of the next cycle. While this behaviour is easily tracked by observing the change in surface activity a question can also be asked about what is happening beneath the surface where the magnetic activity ultimately originates? In order to try to answer this question we can look at the behaviour of the frequencies of the Sun's natural seismic modes of oscillation - the p modes. These seismic frequencies also respond to changes in activity and are probes of conditions in the solar interior. The Birmingham Solar Oscillations Network (BiSON) has made measurements of low-degree (low-$\\ell$) p mode frequencies over the last three solar cycles, and so is in a unique position to explore the current unusual and extended solar minimum. We compare the frequency shifts in the low-$\\ell$ p-modes obtained from the BiSON data with the change in surface activity as measured by different proxies and show there are significant differences especially during the declining phase of solar cycle 23 and into the current minimum. We also observe quasi-biennial periodic behaviour in the p mode frequencies over the last 2 cycles that, unlike in the surface measurements, seems to be present at mid- and low-activity levels. Additionally we look at the frequency shifts of individual $\\ell$ modes. ", "introduction": "The magnetic activity of the Sun is known to follow an approximately 11-year cycle ranging from quiet periods to much more active periods. Starting in March 1755 these cycles have been labeled with a number and we are currently in the minimum between cycles 23 and 24. However, this current minimum is proving to be somewhat unusual, both in terms of how long it is lasting and just how quiet the Sun is, at least in terms of the more recently observed cycles. Surface measures of solar actively such as the number of visible sunspots and 10.7cm radio flux, all suggest that we are still in (or possibly just leaving) an extended minimum. Therefore, the question presents itself is there anything different about this minimum that we can learn from looking at the activity of the solar interior. This can be investigated by looking at the acoustic modes of oscillation of the Sun (known as p modes). Modes with low angular degrees (low-$\\ell$) travel deeply into the solar interior and so sample the conditions below the solar surface. Also, low-$\\ell$ modes are truly global in nature and therefore are sensitive to a large fraction of the solar disc. It has been known for some time that the frequencies of these oscillations vary during the solar cycle, with the frequencies being highest when the magnetic activity is at a maximum. Therefore, by looking at the change in frequencies we can gain insight into cycle related processes that are occurring beneath the solar surface. The Birmingham Solar Oscillations Network (BiSON) has been collecting unresolved (Sun-as-a-star) Doppler velocity observations for over three decades. These types of observations are sensitive to the large scale `global' oscillations of the sun, (i.e., those modes with the lowest angular degrees, $\\ell$). Due to the fact the BiSON has observations going back into the 1970's, these data offer a unique opportunity to study the last three solar cycles. However, when the network was first established there was only one station and as such the early observations are sporadic and so the data quality is relatively poor. The quality and duty cycle of the data greatly improves after 1986 once three stations were operating and this allows us to study the last two solar cycles in their entirety. ", "conclusions": "The shifts in oscillation frequencies over time is reasonably well correlated with other activity proxies such as the 10.7cm radio flux and ISN. However, there are some differences. The largest of which, at least over the two cycles we have data for, occur recently, starting on the downward part of cycle 23 and heading into the unusual and extended minimum between cycle 23 and 24. The current minimum in the frequency shifts is considerably deeper than in the proxies when compared with the previous minimum and there is clear structure in the shifts that does not appear in the proxies. Additionally, there is some evidence for pseudo-periodic short term (2-year) variations in activity on top of the 11-year cycle. Frequency shifts show sharper amplitude variations during the quiet sun periods compared with the radio flux. Since the frequency shifts are sensitive to conditions below the surface it is likely that these differences are due to changes in the magnetic flux that are occurring in the solar interior. These changes then may have either yet to manifest on the surface, or are attenuated before ever reaching there. Alternatively the differences could be due to zonal effects, since the oscillations are not all equally sensitive at different latitudes of the Sun. We have looked at the shifts of modes with different $\\ell$ in order to gain a clearer picture of this. The $\\ell$ = 0 modes show the closest match with the radio flux. The higher degree modes all show greater discrepancies. The strongest regions of magnetic flux tend to start at latitudes near the poles and then drift towards the equator as the cycle progresses. Therefore, modes that are only sensitive to ceratin latitudes will see their frequency shift's respond to this drift in magnetic flux. This will ultimately add further structure into the plots of the frequency shifts that may not be apparent in the $\\ell$ = 0 modes or radio flux plots. If we assume that cycle 24 had started at the time of the last few points in the data set used here, then the shape of the plots for $\\ell$ = 0, 1, 2, and 3 might also be explained by this latitudinal sensitivity argument as suggested by \\cite{Salabert2009}. Although the fact that the $\\ell$ = 2 modes showed a similar upturn to the $\\ell$ = 0 modes in the final few observations, may actually work against this theory since the zonal $m$ = 0 component, which does have a greater sensitivity to higher latitudes, only has a small effect on the fitted frequency of the $\\ell$ = 2 multiplet. It is clear that further work on the individual $\\ell$ modes is needed to help clear up this matter. In the future we hope to continue the work on individual $\\ell$ modes by comparing the frequency shifts with decomposed magnetic maps. This will allow us to compare the frequency shifts with the change in magnetic flux over time for similar zonal regions." }, "1003/1003.3252_arXiv.txt": { "abstract": "The population of compact massive galaxies observed at $z > 1$ are hypothesised, both observationally and in simulations, to be merger remnants of gas-rich disc galaxies. To probe such a scenario we analyse a sample of 12 gas-rich and active star forming sub-mm galaxies (SMGs) at $1.810^{13}L_{\\odot}$, \\citealt{RR:00}), containing a similar amount of molecular gas, and an extreme star formation rate \\citep{Farrah:02a, Farrah:02b}. Furthermore, the similarity between the stellar mass surface densities of SMGs and the compact massive galaxies at lower redshifts \\citep{Tacconi:06, Tacconi:08, Cimatti:08} makes sub-mm galaxies natural candidates for being the precursors of these compact galaxies. This paper tests this hypothesis and investigates whether the proposed theoretical scenario for the formation of massive compact galaxies is compatible with present deep imaging and star formation analysis of these sub-mm systems. To explore this, we probe the sizes of the SMGs in various phases, which we construct based on morphology and structure, and test whether the star formation rates of these galaxies are in agreement with theoretical expectations. Second, we probe whether the different observed morphologies of the SMGs can be fitted into the evolutionary sequence proposed by current massive galaxies formation models. The paper is structured as follows. In Sect. \\ref{data} we describe our sample. In Sect. \\ref{anal} we explain the method adopted for our analysis. In Sect. \\ref{res} we present our results and discuss their implication in Sect. \\ref{disc}. Throughout, we assume the following cosmology: $H_0$=70 \\mbox{km s$^{-1}$} \\mbox{Mpc$^{-1}$}, $\\Omega_m=0.3$ and $\\Omega_{\\Lambda}=0.7$ and we use AB magnitudes. ", "conclusions": "\\label{disc} Motivated by high-resolution hydrodynamical simulations \\citep{Dekel:06, Cox:08, Hopk:09, Narayanan:09}, we argue that the three morphological classes outlined above could represent an evolutionary sequence, where the disc-like class represents a pre-merger phase followed by the major merger event, while the compact sources can be interpreted as the end-stages of the merger, caught during or just after the coalescence. Numerical simulations naturally explain how the tidal forces involved in the interactions remove angular momentum from the systems, allowing the gas to fuel towards the centre. At this stage, the gas is compressed into a very small volume, leading to surface densities of the order $\\simeq 10^{5} M_{\\odot} pc^{-2}$ \\citep{Hopk:09b}, close to that of molecular clouds. According to the Kennicutt law, the SFR in such conditions is extremely enhanced. Since the dynamical time-scale that drives the collapse is similar to the star formation rate, the system rapidly exhausts its gas while it contracts \\citep{MH:96, Hopk:08}. Thus, it is not surprising that we observe such compact galaxies with high SFR. The models, indeed, predict a spheroidal-like morphology at the time of coalescence or just after, since the coalescence generally completes at about the same time as the gas first reaches the centre \\citep{Cox:08p}. Moreover, our findings that a fraction of the SMGs have a compact morphology agrees with measurements of the gas distribution from CO maps for many galaxies \\citep{Tacconi:06, Tacconi:08}. Therefore, these compact, highly star-forming systems are likely to be in the final phase of the merger, and are the transition link between starbursts and compact galaxies. A peculiar case is represented by the system SMMJ123635.59+621424.1. We classify it as a merger event, as it appears as a double-nucleus system. However, its effective radius and S\\'ersic index are characteristic of those for compact galaxies. It is likely that we are measuring the compactness of the bulge component, given the faintness of the outer light. This systems also shows faint signs of potential spiral structure over a scale of 5 kpc. The SSFR of this system is the highest in our sample, compatible with being near the peak in star formation during the coalescence phase. It is worth to note that this object has the highest infrared luminosity $L=10^{13.01}L_{\\odot}$ (from \\citealt{Mich:09}) of our sample, hence it can be considered a HLIRG, supporting the picture where HLIRGs are galaxies in their maximal star formation periods triggered by interactions. In order to build an ``illustrative'' evolutionary sequence, we can order galaxies inside a given class according to the values of their SSFR. In Figure \\ref{ssfr} we show as a dashed line, the evolution of the SSFR in a simulated galaxy merger, taken from \\citet{Narayanan:09}. The simulation illustrates the evolution of the SFR in a major merger (with mass ratio 1:1) for a $\\sim 2 \\times 10^{13} M_{\\odot}$ dark matter halo. We have computed the SSFR taking the SFR of their Figure 1 and dividing it by the final stellar mass ($8\\times10^{11}M_{\\odot}$, which is roughly the maximum stellar masses of our sample). The trend in SSFR shows a modest star formation rate in the pre-merger phase, in agreement with the value of our diffuse galaxies, followed by a steep rise during the starburst/merger, a peak in the coalescence phase, and then a rapid decline. The trend depicted by the simulations is matched by our data if we assume that the SSFR increases in the merging phase, and then declines for the compact galaxies. The peak is reached during the coalescence, as it is illustrated by the source SMMJ123635.59+621424.1. In simulations the diffuse/isolated systems are not fuelled by a new gas reservoir, as in the case for merging systems, and they simply exhaust the cold gas available. Hence their SSFR declines with time (see \\citealt{Cox:08}). Therefore, in the same plot, we have indicated by the coloured points the SSFR of our sample of SMGs, with the number in the x-axis indicating the position in the evolutionary sequence (in Figure \\ref{acs} galaxies are ordered according to this sequence). Note that we adopt the SFR derived from the IR luminosity from \\citet{Mich:09}, using the \\citet{Kennicutt:98} relation to computed the star formation rate. The fact that we can put our galaxy sample in a SSFR-sequence that well matches that found in hydrodynamical simulations, strengthens the conclusion that the (morphological) evolutionary sequence described above, where large diffuse systems transform themselves in compact remnants passing through a major merger event, is likely the formation mechanism for these galaxies into compact systems seen at slightly lower redshifts. Another piece of evidence supporting this theoretical merging scenario is the size of the progenitor discs we have found. Theoretically these discs are expected to have effective radius of 2-3 kpc \\citep{Dekel:06, Hopk:09}. This is in fact what our observations show. Given the lack of NICMOS imaging for most of the sources and the insufficient resolution in the near-IR, we were forced to use ACS imaging to constrain our morphological sequence. To better test such picture high-sensitivity imaging in near-IR would be required, and this could be achieved only when WFC3 imaging will be available. Summarising, our data can not reject the evolutionary picture depicted by theoretical models in which the precursors of the superdense galaxies are massive, gas-rich discs at $z \\sim 2-3$, which evolve into compact remnants through dissipative major mergers. Moreover, a cold gas accretion driven scenario for the formation of the compact massive galaxies, as the one proposed by \\citet{Dekel:09} perhaps can not be easily supported by our data, since in this case we would likely not observe such a large diversity of size and structure for these progenitors as we observe here." }, "1003/1003.0317_arXiv.txt": { "abstract": "We present results from the \\suzaku\\ observation of the microquasar \\grs\\ performed during the 2005 October multiwavelength campaign. The data include both stable state (class $\\chi$) and limit-cycle oscillation (class $\\theta$). Correct interstellar absorption as well as effects of dust scattering are fully taken into account in the spectral analysis. The energy spectra in the 2--120 keV band in both states are all dominated by strong Comptonization of disk photons by an optically thick ($\\tau \\approx$7--10) and low temperature ($T_{\\rm e} \\approx$2--3 keV) hybrid plasmas containing non-thermal electrons produced with 10--60\\% of the total power input. Absorption lines of highly ionized Fe ions detected during the oscillation indicate that a strong disk wind is developed. The ionization stage of the wind correlates with the X-ray flux, supporting the photoionization origin. The iron-K emission line shows a strong variability during the oscillation; the reflection is strongest during the dip but disappears during the flare. We interpret this as evidence for ``self-shielding'' that the Comptonizing corona becomes geometrically thick in the flare phase, preventing photons from irradiating the outer disk. The low-temperature and high luminosity disk emission suggests that the disk structure is similar to that in the very high state of canonical black hole binaries. The spectral variability during the oscillation is explained by the change of the disk geometry and of the physical parameters of Comptonizing corona, particularly the fractional power supplied to the acceleration of non-thermal particles. ", "introduction": "\\grs\\ is the brightest microquasar in our Galaxy, providing us with a unique opportunity to study the accretion flow onto a black hole at high fractions of Eddington ratio \\citep[for a review, see][]{fen04}. This source, discovered in 1992 with the WATCH instrument on GRANAT \\citep{cas92}, has been persistently active up to now unlike usual soft X-ray transients. \\grs\\ was recognized to be a superluminal source by VLBA observations of radio outbursts \\citep{mir94}, making it a target of great importance for studying the formation mechanism of relativistic jets. From the kinetics of the radio jets, their intrinsic speed and inclination are determined to be $0.92 c -0.98 c$ ($c$ is the light speed) and $66^\\circ-70^\\circ$, respectively \\citep{mir94,fen99}. The near infrared observations have revealed that the system consists of a black hole with a mass of $14\\pm4$ \\solarmass\\ and a K~III type companion star with an orbital period of 33.5 days \\citep{gre01}. This indicates a huge size of the accretion disk with a continuously high mass transfer rate from the companion, which may account for many unique features of this source compared with normal black hole binary systems. The source has been intensively studied at high energies by many observatories on different occasions. \\rxte /PCA observations revealed that \\grs\\ often exhibits dramatic temporal and spectral variations, so-called limit-cycle oscillations, unique features in accreting stellar mass black holes in our Galaxy \\citep[e.g.,][]{bel97}. \\citet{bel00} found the instantaneous spectral state of \\grs\\ can be divided into states, States~A, B, and C, where the source undergoes frequent transition or stays stable. Based on these transition patterns, they classified them into 12 Classes, including 2 ``stable'' classes (Class~$\\chi$ and $\\phi$) with little variability on a time scale longer than $\\sim 1$ sec. This behavior is quite different from normal black holes observed in canonical states, the low/hard, intermediate (or very high), and high/soft states (\\citet{tan95}; see also \\citet{hom05} and \\citet{rem06} for more recent classification). The correspondence of States~A, B, and C of \\grs\\ to these states is not completely understood, however \\citep[][]{rei03}. Theoretically, the limit-cycle oscillation can be explained by thermal instability in the inner part of the accretion disk under high mass accretion rate, triggering limit-cycle transitions between two stable branches in the surface density vs $\\dot{M}$ plane: the standard disk and optically-thick advection dominated accretion flow, so-called the slim disk \\citep{abr88}. In fact, two-dimensional hydrodynamic simulations are successful to reproduce the limit cycle behaviors of \\grs\\ qualitatively \\citep{ohs06}. However, direct observations that trace the evolution of the disk structure both in the oscillation and stable state are still insufficient to be compared with model predictions, and their interpretations may be subject to large uncertainties since understanding of the origin of the X-ray emission is not fully established. Many early works relied on the spectral modelling by a multi-color disk (MCD; \\citealt{mit84}) model plus a power law, as applied to other black hole binaries. There is no guarantee that this model always holds to \\grs, in particular when the power law component dominates the entire flux. Later observational works \\citep{zdz01,don04,ued09} have suggested that the X-ray spectra of \\grs\\ at least in some states are dominated by an optically-thick Comptonization of the disk photons off low energy electrons, by applying physically more realistic models than the canonical MCD + power law model. \\citet{zdz01,zdz05} analyzed broad band spectra observed with \\rxte /PCA, HEXTE, and \\cgro /OSSE, and found that they can be described by Comptonization from thermal/non-thermal hybrid plasmas. Such low temperature, optically-thick Comptonization may be common characteristics in accretion flows at high Eddington fractions, which may apply to those onto supermassive black holes \\citep{mid09}. Observations of local spectral features such as iron-K emission line and absorption lines give key information to reveal the structure of accretion disks. From high resolution spectra (CCD or the HETGS instrument on \\chandra ) of \\grs, absorption lines from highly ionized ions have been clearly detected \\citep{kot00,lee02,ued09}. This indicates the presence of strong disk wind most probably occurring at $r\\sim 10^5 r_{\\rm g}$ ($r_{\\rm g} \\equiv GM/c^2$ is the gravitational radius where $G$, $M$, and $c$ being the gravitational constant, mass of the black hole, and light velocity, respectively), which carries a huge amount of accreting gas outward the system \\citep{ued09,nei09}. On the other hand, the emission line profile gives constraints on the geometry between the continuum emitting region and reflector, most probably the accretion disk. An iron-K emission line has been detected from \\grs, whose shape and intensity seem to depend on the states \\citep[][]{kot00,nei09}. Since its equivalent width is not large ($<50$ eV) in \\grs, we need both good spectral resolution and large effective area to accurately measure the profile, in particular to trace the change during oscillations. \\suzaku\\ \\citep{mit07}, the 5th Japanese X-ray satellite, observed \\grs\\ on 2005 October 16--18 (UT throughout the paper) as a part of the science working group's observation program. The large effective area with good energy resolution in the 0.2--12 keV band, together with the simultaneous coverage of the 10--600 keV band with unprecedented sensitivities, give us the best opportunities to reveal the origins of the X-ray emission and disk structure of \\grs\\ both in stable and oscillation states. In our work, we refer to the latest results on the metal abundances of interstellar (plus circumstellar) gas toward \\grs\\ as determined by \\chandra /HETGS \\citep{ued09}. In the analysis, we also take into account the effects by dust-scattering, which have been ignored in most of previous works. We show that these are important for accurate modelling of the spectra. \\S~2 describes the observation and data reduction, and \\S~3 the light curves and states of \\grs\\ in our observations. The spectral analysis is presented in \\S~4. We discuss the interpretation of our results in \\S~5. The conclusions are summarized in \\S~6. In our paper, we assume the distance of $D=12.5$ kpc and inclination angle of $i=70^\\circ$ unless otherwise stated. Around the epoch of our \\suzaku\\ observations, a large multiwavelength campaign was conducted involving other space and ground observatories, whose results will be reported in a separate paper (Ueda et al., in preparation). For the preliminary results of the campaign, refer to \\citet{ued06}. The {\\it INTEGRAL} light curves are published in \\citet{rod08a}. ", "conclusions": "\\begin{enumerate} \\item We find that the \\suzaku\\ broad band spectra of \\grs\\ band both in the stable (Class~$\\chi$) and oscillation states (Class~$\\theta$) are commonly represented with a model consisting of a MCD model and its Comptonization with a reflection component, over which the opacity of the disk wind (including iron-K absorption lines) is applied. The Comptonized component dominates the flux above $\\sim$3 keV. \\item The Comptonization is made by optically-thick ($\\tau \\approx$ 7--10), non-thermal / thermal ($T \\approx$ 2--3 keV) hybrid plasmas. The non-thermal electrons are produced with 10--60\\% of the total energy input to the plasma with a power law index of $\\approx$3--4, which account for the hard X-ray tail above $\\sim$50 keV. \\item During the limit-cycle oscillation, the reflection strength, estimated by the iron-K emission line, is the largest during the dip phase but disappears in the flare phase. We interpret this as evidence for self-shielding effects that Comptonized photons are obscured by the surrounding cool region of the expanded disk when viewed at a very high inclination angle from the outer disk. The evolution of the disk geometry is in accordance with theoretical predictions. \\item The disk wind, traced by iron-K absorption lines of \\ion{Fe}{25} and \\ion{Fe}{26}, always exists during the oscillation and its ionization well correlates with the X-ray flux. This supports the photoionization origin. In the stable state, the iron-K absorption lines are not detected probably because it is highly ionized and/or the scale height is small. \\item The disk parameters suggest that the inner disk structure is similar to that in the very high state of black hole binaries. The spectral variability in the oscillation state is explained by the change of the disk geometry and of physical parameters of Comptonizing corona, particularly the fractional power supplied to the acceleration of non-thermal particles. \\end{enumerate}" }, "1003/1003.2638_arXiv.txt": { "abstract": "Observations continue to support the interpretation of the anomalous microwave foreground as electric dipole radiation from spinning dust grains as proposed by Draine \\& Lazarian (1998ab). In this paper we present a refinement of the original model by improving the treatment of a number of physical effects. First, we consider a {\\it disk-like} grain rotating with angular velocity at an arbitrary angle with respect to the grain symmetry axis (i.e., grain wobbling) and derive the rotational damping and excitation coefficients arising from infrared emission, plasma-grain interactions and electric dipole emission. The angular velocity distribution function and the electric dipole emission spectrum for disk-like grains are calculated using the Langevin equation, for cases both with and without fast internal relaxation. Our results show that for fast internal relaxation, the peak emissivity of spinning dust, compared to earlier studies, increases by a factor of $\\sim $ 2 for the Warm Neutral Medium (WNM), the Warm Ionized Medium (WIM), the Cold Neutral Medium (CNM) and the Photodissociation Region (PDR), and by a factor $\\sim$ 4 for Reflection Nebulae (RN). The frequency at the emission peak also increases by factors $\\sim$1.4 to $\\sim$2 for these media. Without internal relaxation, the increase of emissivity is comparable, but the emission spectrum is more extended to higher frequency. The increased emission results from the non-sphericity of grain shape and from the anisotropy in damping and excitation along directions parallel and perpendicular to the grain symmetry axis. Second, we provide a detailed numerical study including transient spin-up of grains by single-ion collisions. The range of grain size in which single-ion collisions are important is identified. The impulses broaden the emission spectrum and increase the peak emissivity for the CNM, WNM and WIM, although the increases are not as large as those due to the grain wobbling. In addition, we present an improved treatment of rotational excitation and damping by infrared emission. ", "introduction": "Introduction} Diffuse Galactic microwave emission in the 10 -- 100 GHz frequency range carries important information on the fundamental properties of the interstellar medium, but it also interferes with Cosmic Microwave Background (CMB) experiments (see Bouchet et al.\\ 1999, Tegmark et al.\\ 2000, Efstathiou 2003). It used to be thought that there were only three major components of the diffuse microwave Galactic foreground: synchrotron emission, free-free radiation from plasma (thermal bremsstrahlung) and thermal emission from dust. However, in the range of frequency from 10 to 100 GHz an anomalous microwave foreground which was difficult to reconcile with the components above was first reported by Kogut et al.\\ (1996a, 1996b). de Oliveira-Costa et al.\\ (2002) gave this emission the nickname ``Foreground X'', reflecting its mysterious nature. This component is spatially correlated with 100 $\\mu$m thermal emission from dust, but its intensity is much higher than one would expect by extrapolating the thermal dust emission spectrum to the microwave range. Draine \\& Lazarian (1998a,b) proposed that this foreground was electric dipole radiation from ultrasmall spinning dust grains. Although such emission from spinning dust had been discussed previously (see Erickson 1957, Ferrara \\& Dettmar 1994), DL98a were the first to include the variety of excitation and damping processes that are relevant for very small grains. As time went on alternative models for the enigmatic foreground have appeared to be inconsistent with observations\\footnote{For instance, the dust-correlated synchrotron emission suggested by Bennett et al.\\ (2003) has now been ruled out (see de Oliveira-Costa et al.\\ 1999, Finkbeiner, Langston, \\& Minter 2004, Boughn \\& Pober 2007, Gold et al.\\ 2009, 2010).}, while the predictions of the spinning dust model have thus far been confirmed. As a result, spinning dust is now the principal explanation for the mysterious ``Foreground X''. Although the model in Draine \\& Lazarian (1998ab) provided quantitative predictions consistent with observational data, the current state of precision measurement of the foregrounds calls for refinement of the model, using a better description of the complex grain dynamics and modifying some of the original assumptions. Recent studies showed that the correspondence of the DL98 model to observations can be improved by adjusting the parameters of the model. For instance, Dobler et al.\\ (2009) used the Wilkinson Microwave Anisotropy Probe (WMAP) 5 year data to show that a broad bump with frequency at $\\sim$ 40GHz correlated with H$\\alpha$, a tracer of the warm ionized medium (WIM). They showed that this bump is consistent with predictions from a DL98 model modified so that grains have a characteristic dipole moment of $3.5$ D at grain size 1 nm, and the gas number density of the WIM is $n_{\\H}=0.15 \\cm^{-3}$ (cf. $n_{\\H}=0.1 \\cm^{-3}$ in the DL98 model). On the theoretical front, Ali-Ha\\\"imound et al.\\ (2009) improved the accuracy of predictions of the model of emission from spinning dust using the Fokker-Planck equation for the angular velocity ${\\bomega}$. The authors quantified the deviations of the grain angular velocity distribution function from the Maxwellian approximation that had been used by DL98b for the sake of simplicity. With the other assumptions being identical to DL98b, their findings are not much different from DL98b's predictions. However, both the DL98 model and the refined model by Ali-Ha\\\"imoud et al. (2009) disregarded the non-sphericity of grains and the anisotropy in the damping and excitation processes. Obviously, this assumption is inexact for non-spherical grains or when there exists any anisotropy in the damping and excitation processes. This present paper is intended to go deeper into studies of grain dynamics for a disk-like grain geometry, relaxing more of the simplifying assumptions in the original DL98b treatment. The main thrust of our present work is to provide a better description of several physical processes which have not been addressed in their complexity either in DL98b or the papers that followed. In particular: (i) the effects on electric dipole emission arising from the wobbling of the axis of major inertia of the {\\it disk-like} grain around the angular momentum due to internal relaxation and the anisotropy of grain rotational damping and excitation (not yet been treated in the literature, so far as we are aware); and (ii) transient spin-up of very small grains due to single-ion collisions. First of all, the former process, i.e., imperfect internal alignment in the non-spherical grain, can alter the frequency at which the electric dipole emits. In the DL98 model, the emission frequency is identical to the angular frequency $\\omega/2\\pi$. However, if the dipole moment is fixed in the grain body, then the complex motion of the grain axes around the angular momentum $\\bJ$ will result in emission at frequencies different from its angular frequency. In addition, imperfect alignment is essential for many astrophysical processes, e.g., for grain alignment (see Lazarian 2007 and Lazarian \\& Hoang 2009 for recent reviews). In our quest to understand the rotational dynamics of grains which do not rotate about their axis of major inertia we capitalize on improved understanding of internal randomization arising from thermal fluctuations within the grain (Lazarian 1994, Lazarian \\& Roberge 1997, Weingartner 2009, Hoang \\& Lazarian 2009). Disalignment of the grain's principal axis from the direction of the angular momentum $\\bJ$ will cause the angular velocity to increase, leading to increased electric dipole emission at high frequency. Our paper provides a quantitative description of the effect on the spinning dust emissivity. The latter process -- collisions with ions -- is important for small grains where the angular momentum of an impinging ion can be larger than the pre-collision grain angular momentum, resulting in a rotational excitation spike. In DL98b it was noted that the effect is expected to increase the spinning dust emissivity, but no quantitative description was given. While the treatment of high impulse ion collisions is easily performed within our Langevin code, the treatment of grain wobbling for non-spherical grains requires a careful and somewhat tedious modification of our treatment of the angular momentum diffusion and damping in the DL98 model. In particular, we have to consider separately parallel and perpendicular contributions to grain damping and excitation. In addition, we provide an improved treatment of infrared emission from spinning dust grains. The structure of the paper is as follows. In \\S 2, we present elements of the DL98 model and our modifications to that model. In \\S 3 we provide detailed calculations for rotational damping and excitations arising from plasma-grain interactions and electric dipole emission for the disk-like grain geometry. Refined calculations for the effects of infrared emission are presented in \\S 4 and 5. In \\S 6, we study the effect of the grain precession on the electric dipole emission spectrum and identify its frequency modes. In \\S 7 and 8 we present our numerical techniques for finding the distribution functions of grain angular velocity and electric dipole emission, and benchmark calculations. In \\S 9, we present our results for emissivities from spinning dust for various idealized environments, and clarify the role of grain shape and differential rotational damping and excitation processes to the increase of peak emissivity and frequency. Discussion and summary are presented in \\S 10 and 11, respectively. ", "conclusions": "\\subsection{Comparison with earlier studies} The DL98 model was the first proposal to explain the $20-40$ GHz anomalous foreground emission in terms of rotational emission from ultra-small dust particles. Numerous observations (Dobler et al.\\ 2009 and references therein) have confirmed the DL98 model predictions. In the DL98 model the calculations were mostly analytical. The present study used a different approach, namely, the Langevin equation. We successfully benchmarked our calculations against the Fokker-Plank calculation in Ali-Ha\\\"imoud et al.\\ (2009) for the perfect alignment case, and then introduced additional effects not included in DL98. In particular, our approach allows us to treat discrete high impact impulses arising from single-ion collisions with the grain. Most importantly, it enables us to study the spinning dust emission for the case of imperfect internal alignment with and without fast internal relaxation. Among grain rotational damping and excitation processes, plasma drag is shown to be important for Molecular Clouds and CNM. Ragot (2002) considered the interactions of the plasma with spinning dust from the point of view of generation of waves. We believe that for the case we consider, the dust-plasma interaction is accurately approximated by interaction of the grain with individual, uncorrelated, passing ions. DL98b noted that the typical rotational quantum number of even a small PAH would be $J/\\hbar\\gtsim 10^2$, allowing the dynamics to be treated classically. Ysard \\& Verstraete (2009) formulated the problem quantum-mechanically; the close correspondence between their results and those of DL98b confirms that a classical treatment is entirely adequate.\\footnote{% In discussing some effects, e.g., dust-plasma interactions, DL98b did include quantum limitations for the angular momentum transfer, i.e. that $\\Delta J$ cannot be less than $\\hbar$.} While our studies confirm the general validity of the DL98 model of spinning dust emission, they also show ways of making the model predictions more accurate. For instance, our refined treatment of grain dynamics provides peak emissivity predictions which differ by factors from $\\sim 2$ to $4$ from the original ones, depending on the environments. These differences may be detectable with high precision future observations. The largest uncertainty within the model is the value of the dipole electric moment. The distribution of dipole moments adopted in DL98b was only an educated guess. Moreover, this distribution could be affected by interstellar processes and change from one media to another. For instance, Dobler et al.\\ (2009) showed the anomalous emission in the 5-year WMAP data exhibits a bump at $\\sim 40$ GHz. This peak frequency is larger than the prediction by the original DL98b model but could be explained if a modified distribution of electric dipole moments is adopted (Dobler et al.\\ 2009), but the new effects considered here -- misaligned rotation and impulsive excitation by ions may also explain the observed emission spectrum. \\subsection{Effect of single-ion collisions} DL98b pointed out that for the CNM, WNM and WIM the angular momentum transferred in a single ion collision can be larger than the angular momentum of the rotating grain, but neither DL98b or the subsequent paper by Ali-Ha\\\"imoud et al.\\ (2009) treated the impulsive nature of the ion collisions. In the present paper, we describe ion-grain collisions as Poisson-distributed discrete events. When an ion hits the grain surface, the grain undergoes transient spin-up and then gets damped, mainly by electric dipole damping. Due to strong electric dipole damping, small grains spend most of their time rotating subthermally between two rotational spikes due to single-ion collisions. Ionic impulses are shown to be important for grains smaller that $\\sim 8.6\\times 10^{-8}$, $1.2 \\times 10^{-7}$ and $3.8\\times 10^{-8}$cm for the WIM, WNM and CNM, respectively. Our quantitative simulations show that the impulses extend the high-frequency tail of the emission spectrum but change the peak frequency only slightly. They can enhance the peak emissivity by $23\\%$ and $11\\%$ for the WIM and WNM, respectively. \\subsection{Effect of Grain Wobbling} The DL98b model assumed perfect alignment of the grain axis of maximal moment of inertia with the instantaneous angular momentum. This assumption is valid only for strong internal relaxation {\\it and} suprathermal (much faster than thermal) rotation of grains (Lazarian 1994). The tiny grains we deal with rotate thermally or subthermally. For such a rotation, the grain axes are {\\it not} coupled with the angular momentum. Due to imperfect internal alignment (i.e., grain wobbling), the anisotropy arising from grain shape, and differential damping and excitation (e.g. due to infrared emission and plasma drag) obviously affect the grain mean rotation rate and therefore the emissivity. Increased grain anisotropy, as measured by the anisotropy ratio $\\eta$ (eq.\\ \\ref{eq:eta}) produces an increase in the peak frequency and emissivity. As we do not know the rates for internal relaxation in microscopic molecular size grains, we considered two cases: very rapid internal relaxation, and very slow internal relaxation. Using the Langevin equation, we showed that for the former case, the peak emissivity of spinning dust increases by a factor of $\\sim 2$, and the peak frequency of the spectrum increases by factors from 1.4 to 1.8, depending on the environment. Grain wobbling has a larger effect in broadening the emission spectrum when the internal relaxation is slow. Especially, for the RN when the anisotropy of damping and excitation from infrared emission is very high, the peak emissivity can be increased by a factor of $4$ and peak frequency is increased by a factor of $2$. For the WIM, the joint effect of grain wobbling and impulses due to single-ion collisions causes the peak frequency to shift from $\\sim23$ to $\\sim 35$ GHz. It is interesting to note that the new predicted peak frequency is close to the bump with the anomalous emission found in the WMAP five year data (Dobler et al.\\ 2009). In their paper, using the DL98b model, they explained the observed spectrum with a modified model of spinning dust in which the grains have broader distribution of electric dipole moments. They also varied the number density of the WIM to produce spectra with higher peak frequency. These modifications do not seem to be necessary with our improved treatment of grain dynamics." }, "1003/1003.5829_arXiv.txt": { "abstract": "We have developed a clumpy stellar wind model for OB supergiants in order to compare predictions of this model with the X-ray behaviour of both classes of persistent and transient High Mass X-ray Binaries (HMXBs). ", "introduction": "The Galactic plane survey performed by \\emph{INTEGRAL} satellite in the last 7 years allowed the discovery of a new class of HMXBs with OB supergiants, called \\emph{Supergiant Fast X-ray Transients} (SFXTs): they are transient sources which sporadically exhibit flares, with duration of a few hours, reaching luminosity of $10^{36}-10^{37}$~erg~s$^{-1}$, and a high dynamic range, spanning 3 to 5 orders of magnitude \\cite{Sguera2005}. The accretion mechanism responsible for the peculiar SFXT behaviour is still not clear (see \\cite{Sidoli2009} and references therein for a recent review). Among the different possible explanations, \\cite{intZand2005} proposed that the SFXTs flares are produced by the accretion of dense blobs of matter from the companion wind. ", "conclusions": "" }, "1003/1003.6127_arXiv.txt": { "abstract": "We consider the effect of lensing magnification on high redshift sources in the case that magnification varies on the sky, as expected in wide fields of view or within observed galaxy clusters. We give expressions for number counts, flux and flux variance as integrals over the probability distribution of the magnification. We obtain these through a simple mapping between averages over the observed sky and over the magnification probability distribution in the source plane. Our results clarify conflicting expressions in the literature and can be used to calculate a variety of magnification effects. We highlight two applications: 1. Lensing of high-$z$ galaxies by galaxy clusters can provide the dominant source of scatter in SZ observations at frequencies larger than the SZ null. 2. The number counts of high-$z$ galaxies with a Schechter-like luminosity function will be changed at high luminosities to a power law, with significant enhancement of the observed counts at $L\\simgt 10\\ L^*$. ", "introduction": "Magnification due to gravitational lensing leads to observable effects, namely changes in the number density of galaxies behind large-scale structure and galaxy clusters (known as magnification bias) and in the moments of the flux distribution due to unresolved sources at high redshift. These and other effects of lensing magnification have been studied extensively in the last few decades, usually assuming simple expressions that apply for constant magnifications. In this brief note, we generalize to the case where magnification varies on the sky -- the variation is taken to be given by a magnification probability in the {\\it source plane}, while quantities of interested are observed as averages in the {\\it image plane}. We apply this calculation to lensing of the intrinsic number count distributions of high redshift galaxies as well as moments of the flux for Poisson distributed high-$z$ galaxies behind galaxy clusters. Our goal is to provide the formulae needed for magnification effects in a variety of physical situations and give estimates of the scale of the main effects. Applications to more detailed models and results for Sunyaev-Zel'dovich surveys have been presented in a separate paper (Lima, Jain \\& Devlin 2009). ", "conclusions": "We have derived expressions for computing average quantities in the observed image plane, given a distribution $P(\\mu)$ of magnifications in the source plane. Our results, summarized in Eqns. 9-12, generalize expressions found in the lensing literature for the case of constant magnifications. We illustrated the effect of lensing on steep number counts of background galaxies % and on boosting the contamination that high redshift galaxies induce in cluster SZ fluxes. The formulae we have presented may be useful for studying the intrinsic properties of high redshift galaxies and for current and upcoming cluster surveys. The quantitative estimates presented here for galaxy clusters are based on analytical models of halos. While these incorporate realistic density profiles as well as halo ellipticity, they miss the full complexity of halo bimodality (due to major mergers) and substructure. These features only enhance the effects of averaging we have considered. Finally we note that the lensing effects discussed here do not impact the predicted (magnification induced) cross-correlation of number counts measured in different redshift bins (Moessner \\& Jain 1998). Such cross-correlations depend on the two-point cross-correlations of magnification with the galaxy density. Thus measurements of magnification bias from galaxy-quasar cross-correlations % or of its contaminating effect on high-$z$ ISW cross-correlations % are unaffected by the spatial averaging issues discussed here. \\smallskip \\noindent {\\it Acknowledgments:} We thank Anna Cabre, Yan-Chuan Cai, Mark Devlin, Mike Jarvis and Ravi Sheth for helpful discussions, and Stefan Hilbert for sharing his simulation results. We especially thank Gary Bernstein and Peter Schneider for sharing their insights and knowledge of the history of the field. This work was supported in part by an NSF-PIRE grant and AST-0607667." }, "1003/1003.6061_arXiv.txt": { "abstract": "A study of two dE/dSph members of the nearby M\\,81 group of galaxies, KDG\\,61 and UGC\\,5442 = KDG\\,64, has been made. Direct Hubble Space Telescope (HST) Advanced Camera for Surveys (ACS) images and integrated-light spectra of 6 m telescope of Special Astrophysical Observatory of Russian Academy of Sciences have been used for quantitative star formation history analysis. The spectroscopic and colour-magnitude diagrams analysis gives consistent results. These galaxies appear to be dominated by an old population (12--14 Gyr) of low metallicity (${\\rm [Fe/H]}\\sim-1.5$). Stars of ages about 1 to 4 Gyr have been detected in both galaxies. The later population shows marginal metal enrichment. We do not detect any significant radial gradients in age or metallicity in these galaxies. Our radial velocity measurement suggests that the \\hii{} knot on the line-of-sight of KDG\\,61 is not gravitationally attached to the galaxy. ", "introduction": "Most of the known galaxies are dwarfs. They span mass ranges from a few $10^9$ \\msol, (or ${\\rm M}_V \\simeq -18$ mag; $10^9$ \\lsol) like NGC~205, to $10^7$ \\msol{} (${\\rm M}_V \\simeq -12$ mag; $10^6$ \\lsol), like Fornax and Sculptor galaxies, satellites of the Milky Way, and probably down to $10^5$ \\msol{} (${\\rm M}_V \\simeq -3$ mag; 1000 \\lsol) for the fainter Milky Way satellites found from the SDSS \\citep{koposov09}. Dwarfs vary considerably in their gas content and morphology, between dwarf irregular galaxies (dIrr) and dwarf elliptical or spheroidal galaxies (dE/dSph). Environmental effects likely drive the evolution of star forming dwarf galaxies to the quiescent ones after gas removal and/or exhaustion. Among the quiescent objects, the fainter ones are traditionally called dSph while the more massive are named dE (diffuse or dwarf ellipticals; e.g. \\citealp{grebel99}). The bound is often set at ${\\rm M}_V \\simeq -14$ mag. However, the physical distinction between these two classes is unclear. Dwarf spheroidals are generally companions of massive galaxies (20 of them were identified around the Milky Way and 15 around Andromeda; \\citealp{irwin08}) while dEs are found in clusters (1141 of them are listed in the Virgo Cluster Catalogue; \\citealp{bst85}), but NGC~205, prototype of dEs, is also a companion of M\\,31, and faint dE/dSph are now found in nearby clusters \\citep{trentham02,mieske07,adami07,derijcke09}. The nearby dSphs are often found to be dominated by dark matter \\citep{mateo98}, while the more massive, dEs are apparently similar to the massive elliptical galaxies: the stellar content of NGC~147, NGC~185 and NGC~205 account for half of their dynamical masses \\citep{derijcke06}. Other properties, like the S\\'ersic index, n, characterizing the shape of the photometric profile, seems to form a continuum over all the mass range. \\cite{derijcke09} proposes that the S\\'ersic index is independent of the luminosity for ${\\rm M}_V \\gtrsim -14$ mag, $n \\approx 0.7$, and increases with the luminosity above this limit. The change in the photometric scaling relation may reflect the different nature of dEs and dSphs, but both the measurement uncertainties and the cosmic dispersion are still large enough to dispute such a dichotomy. Exploring the mass sequence of dEs/dSphs is a significant step toward the understanding of the different processes of their formation and evolution. Is there a sharp transition of the dark matter content at some intermediate mass? Were all these galaxies formed at the same epoch and are there any systematics in the star formation? To address these questions, we have studied two of the nearest dE/dSph galaxies with intermediate masses, located in the M\\,81 group at about 3.6\\,Mpc. They are two of the brightest dE/dSph in this group: UGC\\,5442 = KDG\\,64 and KDG\\,61. The KDG designations (Karachentseva dwarf galaxies) are from the galaxy catalogue by \\citet{karachentseva68}. Both objects are situated in the central region of the group. KDG\\,61 is located 29.5\\,arcmin South of M\\,81, i.e.\\ 31\\,kpc in projected distance. KDG\\,64 is projected 58.2\\,arcmin South of NGC\\,3077 (61.5\\,kpc) and 97.5\\,arcmin South-East of M\\,81 (103\\,kpc). KDG\\,61 is the most luminous dSph in the M\\,81 group and it is one of the closest companions to the M\\,81 galaxy. \\citet{johnson97} has found a \\hii{} knot 34\\,arcsec NE from the centre (0.6 kpc in projection, position J2000 09:57:07.48+68:35:53.9). This sign of recent star formation and the detection of a \\hi{} cloud led \\citet{k00} and \\citet{boyce01} to classify the galaxy as transitional type, between dIrr and dSph. A globular cluster has been discovered in KDG\\,61 \\citep{sharina05} at the projected distance of 0.05\\,kpc (3\\,arcsec) from the galaxy's centre. This cluster has an effective radius of 4.7\\,pc (0.3\\,arcsec), absolute luminosity ${\\rm M}_V\\,=\\,-7.55$ and a de-reddened colour of $V-I\\,=\\,0.92$. This colour is similar to the rest of the galaxy (see Table~\\ref{table:general}). In spite of the fact that KDG\\,64 is fainter than KDG\\,61, it has half magnitude higher surface brightness. The nucleus-like object \\citep{bp94,bremnes98} is a distant galaxy \\citep{k00,sp02}. No globular cluster was detected by \\citet{sharina05}. The internal kinematics of KDG\\,64 was studied using long-slit spectroscopy by \\citet{sp02}. The absolute luminosity, total colour and other general parameters of the both galaxies are shown in the Table~\\ref{table:general}. Coordinates are taken from HyperLEDA\\footnote{\\url{http://leda.univ-lyon1.fr/}} \\citep{patu03}. Total $B$ and $R$ magnitudes are taken from \\citet{bremnes98}, central surface brightness in $V$ and $(V-I)$ colour are from \\citet{sharina08}. The axial ratio was taken from the `A Catalog of Neighboring Galaxies' \\citep{CNG} and the central velocity dispersion from \\citet{ps02}. The rest of values are from this work. All magnitudes and colours are corrected for Galactic extinction using the \\citet{schlegel98} maps. \\begin{table*} \\centering \\caption{General parameters of KDG\\,61 and KDG\\,64}. \\label{table:general} \\begin{tabular}{lll} \\hline & KDG\\,61 & KDG\\,64 \\\\ Position (J2000) & 095703.1$+$683531 & 100701.7$+$674938 \\\\ $E(B-V)$, mag & \\phantom{$-$00}0.072 & \\phantom{$-0$}0.054 \\\\ $B^0_T$, mag & \\phantom{$-$0}14.93 & \\phantom{$-$}15.29 \\\\ Axial ratio $b/a$ & \\phantom{$-$00}0.58 & \\phantom{$-0$}0.47 \\\\ $(B-R)^0_T$, mag & \\phantom{$-$00}1.44 & \\phantom{$-0$}1.35 \\\\ $(V-I)^0_e$, mag & \\phantom{$-$00}0.94 & \\phantom{$-0$}0.99 \\\\ Heliocentric velocity, \\kms{} & \\phantom{$-$}$221\\pm3$ & $-15 \\pm 13$ \\\\ Distance modulus $\\mu(0)$, mag & \\phantom{$-$0}$27.77\\pm0.04$ & \\phantom{$-$}$27.84\\pm0.04$ \\\\ Distance, Mpc & \\phantom{$-$00}$3.58\\pm0.07$ & \\phantom{$-0$}$3.70\\pm0.07$ \\\\ Spatial separation to M~81, kpc & \\phantom{$-$0}40 & 230 \\\\ $B$ absolute magnitude, mag & \\phantom{0}$-12.84$ & $-12.55$ \\\\ Central surface brightness in $V$, mag arcsec$^{-2}$ & \\phantom{$-$0}23.76 & \\phantom{$-$}22.89 \\\\ Central velocity dispersion, \\kms{} & & \\phantom{$-$}$25.0 \\pm 16.0$ \\\\ Fraction of old stars (12--14 Gyr) & \\phantom{$-$0}82--86 \\% & \\phantom{$-$}72--89 \\% \\\\ Metallicity of old stars, [Fe/H], dex & \\phantom{00}$-1.5\\pm0.2$ & \\phantom{0}$-1.6\\pm0.4$ \\\\ \\hline \\end{tabular} \\end{table*} In Sect.\\,2 we analyse colour-magnitude diagrams (CMDs) of KDG\\,61 and KDG\\,64 derived from HST/ACS images. We measure accurately the distance and find the star formation and metal enrichment history. In Sect.\\,3 we present the long-slit spectra that we obtained and their analysis using full spectrum fitting. Section\\,4 discusses the nature of these galaxies and Section\\,5 gives our conclusions. ", "conclusions": "We have derived the star formation histories of two dE/dSph galaxies, KDG\\,61 and UGC\\,5442 = KDG\\,64, belonging to the nearby M\\,81 galaxy group. These galaxies have regular axisymmetric morphologies consistent with an early-type classification. They appear to be dominated by an old stellar population (approximately 12--14\\,Gyr) of low metallicity (${\\rm [Fe/H]}\\simeq-1.5$). We also detected stars formed about 1 to 4\\,Gyr ago in both galaxies with a marginal metal enrichment. KDG\\,64 is slightly less luminous and it has a higher surface brightness than KDG\\,61, but both galaxies are nearly a mid point of the luminosity sequence connecting the bright dEs, like, for example NGC~205, and the smaller local dwarfs, like, for example, Sculptor. The well-known numerical simulations by \\citet{yun99} gives the time since nearest approach between M\\,81--M\\,82--NGC\\,3077 of about 300 Myr ago. Recent star formation events in these and tidal dwarf galaxies agree well with this age. But signs of earlier approaches and, therefore, earlier common star formation events are probably ``erased'' by the recent interaction, and, therefore, we could not relate the recent star formation event in KDG\\,61 and KDG\\,64 (1--4 Gyr ago) to the large scale star formation event in the M\\,81 group. There is no sign of a gaseous component in none of the two galaxies (\\hi{} or ionized gas). On the basis of our radial velocity measurements on the stellar components, we reject previous suggestions for \\hi{} clouds association with the dwarfs. We have also found, that the \\hii{} knot previously suggested to belong to KDG\\,61 is in fact a wind-blown superbubble associated to a \\hi{} tidal stream projected on the line-of-sight. This study was made with two complementary approaches: CMD fitting and full spectrum fitting. The first method uses CMDs usually obtained with HST/ACS, while the last uses long slit spectra that can be acquired on a ground based telescope. The two methods are consistent and it is notable that the full spectrum fitting can give finer constrains on the metallicity. In full spectrum fitting, the metallicity is directly constrained by the metal absorption lines from the stars, while CMDs rely on broad bands colours and on stellar evolution models. It is the first time that a detailed spectroscopic determination of the SFH was successfully applied to low surface brightness galaxy (with low S/N data). We like to notice the agreement between the two approaches. It is promising for the future where both methods will be applicable within a distance of 20\\,Mpc. The spectroscopic analysis of low surface brightness objects may also be used, at moderate cost, in a large volume of the local universe to explore the diversity of SFHs, which can be related to other characteristics of the galaxies and to the environment." }, "1003/1003.1194_arXiv.txt": { "abstract": "The effects of noncommutativity on the phase space of a dilatonic cosmological model is investigated. The existence of such noncommutativity results in a deformed Poisson algebra between the minisuperspace variables and their momenta conjugate. For an exponential dilaton potential, the exact solutions in the commutative and noncommutative cases, are presented and compared. We use these solutions to address the late time acceleration issue of cosmic evolution.\\vspace{5mm}\\newline PACS numbers: 98.80.-k, 04.20.Fy, \\vspace{0.8mm}\\newline Keywords: noncommutative cosmology, deformed phase space, late time acceleration\\vspace{.5cm} ", "introduction": "Noncommutativity between space-time coordinates was first introduced by Snyder \\cite{1}, and in more recent times a great deal of interest has been generated in this area of research \\cite{2, 4}. This interest has been gathering pace in recent years because of strong motivations in the development of string and M-theories. However, noncommutative theories may also be justified in their own right because of the interesting predictions they have made in particle physics, a few examples of which are the IR/UV mixing and non-locality \\cite{7}, Lorentz violation \\cite{8} and new physics at very short distance scales \\cite{8}-\\cite{10}. Noncommutative versions of ordinary quantum \\cite{11} and classical mechanics {\\cite{12, 13}} have also been studied and shown to be equivalent to their commutative versions if an external magnetic field is added to the Hamiltonian. A different approach to noncommutativity is through the introduction of noncommutative fields \\cite{14}, that is, fields or their conjugate momenta are taken as noncommuting. These effective theories can address some of the problems in ordinary quantum field theory, e.g. regularization \\cite{14} and predict new phenomenon, such as Lorentz violation \\cite{15}, considered as one of the general predictions of quantum gravity theories \\cite{16}. In cosmological systems, since the scale factors, matter fields and their conjugate momenta play the role of dynamical variables of the system, introduction of noncommutativity by adopting the approach discussed above is particularly relevant. The resulting noncommutative classical and quantum cosmology of such models have been studied in different works \\cite{17}. These and similar works have opened a new window through which some of problems related to cosmology can be looked at and, hopefully, resolved. For example, an investigation of the cosmological constant problem can be found in \\cite{18}. In \\cite{19} the same problem is carried over to the Kaluza-Klein cosmology. The problem of compactification and stabilization of the extra dimensions in multidimensional cosmology may also be addressed using noncommutative ideas in \\cite{20}. In recent years many efforts have been made in cosmology from string theory point of view \\cite{21}-\\cite{24}. In the pre-big bang scenario, based on the string effective action, the birth of the universe is described by a transition from the string perturbative vacuum with weak coupling, low curvature and cold state to the standard radiation dominated regime, passing through a high curvature and strong coupling phase. This transition is made by the kinetic energy term of the dilaton, a scalar field with which the Einstein- Hilbert action of general relativity is augmented, see \\cite{26} for a more modern review. In this Letter we deal with noncommutativity in a dilaton cosmological model with an exponential dilaton potential and to facilitate solutions for the case under consideration, we choose a suitable metric. Our approach to noncommutativity is through its introduction in phase space constructed by minisuperspace fields and their conjugate momenta. Indeed, in general relativity formulation of gravity in a noncommutative geometry of space-time is highly nonlinear and setting up cosmological models is not an easy task. Here our aim is to study some aspects regarding the application of noncommutativity in cosmology, i.e. in the context of a minisuperspace reduction of the dynamics. Since our model has two degrees of freedom, the scale factor $a$ and the dilaton $\\phi$, with a change of variables, we have a set of dynamical variables $(u,v)$, which are suitable candidates for introducing noncommutativity in the phase space of the problem at hand. We present exact solutions of commutative and noncommutative cosmology and show that commutative model cannot describe the late time acceleration while the noncommutative counterpart of the model clearly points to a possible late time acceleration. ", "conclusions": "In this letter we have studied the effects of noncommutativity in phase space, on classical cosmology of a dilaton model with an exponential dilaton potential. Motivation of such study is that the construction of an effective theory is usually the result of the observation that the corresponding full theory is too complicated to deal with. This, for example, is true in describing the quantum effects on cosmology since the full theory is immensely difficult to handle. Introduction of a deformed phase space can be interpreted as one such effective theory. In the present study, in the case of commutative phase space, the evolution of the classical universe is like the motion of a particle (universe) moving on a plane with a constant acceleration. We saw that in this case the universe decelerates its expansion both in early and late times of cosmic evolution, in contrast to the current observation data. We have investigated the possibility of having a late time accelerated phase of the universe, suggested by recent supernova observation, in the context of a deformed phase space model of string cosmology action. Indeed, we have presented a scenario in which cosmic acceleration occurs late in the history of the universe due to introduction of noncommutativity in phase space, on classical cosmology of a dilaton model with an exponential dilaton potential. We have found that while the usual classical model cannot support this acceleration, a classical model with noncommutative phase space variable can drive this late time acceleration. \\vspace{5mm}\\newline \\noindent {\\bf Acknowledgement}\\vspace{2mm}\\noindent\\newline B. Vakili is grateful to the research council of Azad University of Chalous for financial support." }, "1003/1003.2368_arXiv.txt": { "abstract": "{Phase-referencing is a standard calibration technique in radio interferometry, particularly suited for the detection of weak sources close to the sensitivity limits of the interferometers. However, effects from a changing atmosphere and inaccuracies in the correlator model may affect the phase-referenced images, and lead to wrong estimates of source flux densities and positions. A systematic observational study of signal decoherence in phase-referencing and its effects in the image plane has not been performed yet.} {We systematically studied how the signal coherence in Very-Long-Baseline-Interferometry (VLBI) observations is affected by a phase-reference calibration at different frequencies and for different calibrator-to-target separations. The results obtained should be of interest for a correct interpretation of many phase-referenced observations with VLBI.} {We observed a set of 13 strong sources (the S5 polar cap sample) at 8.4 and 15 GHz in phase-reference mode with 32 different calibrator/target combinations spanning angular separations between 1.5 and 20.5 degrees. We obtained phase-referenced images and studied how the dynamic range and peak flux-density depend on observing frequency and source separation.} {We obtained dynamic ranges and peak flux densities of the phase-referenced images as a function of frequency and separation from the calibrator. We compared our results with models and phenomenological equations previously reported.} {The dynamic range of the phase-referenced images is strongly limited by the atmosphere at all frequencies and for all source separations. The limiting dynamic range is inversely proportional to the sine of the calibrator-to-target separation. Not surpriseingly, we also find that the peak flux densities decrease with source separation, relative to those obtained from the self-calibrated images.} ", "introduction": "Phase-referencing is a standard calibration technique in radio interferometry. It allows the detection of a weak source (target source) by using quasi-simultaneous observations of a closeby strong source (calibrator) (see e.g., Ros \\cite{ros05} and references therein). This technique also allows the user to recover the position of the target source relative to that of the calibrator; a position otherwise lost by the use of closure phases in the imaging. Basically, phase-referencing consists of estimating the antenna-based complex gains with the calibrator fringes, time-interpolating these gains to the observations of the target, and calibrating the visibilities of the target with the interpolated gains. Therefore, it is assumed that for each antenna the optical paths of the signals from both sources are similar. However, atmospheric turbulences and/or inaccuracies in the correlator geometrical model may introduce errors in the estimates of the optical paths of the signals and severely affect the phase-referencing. These errors can be partially corrected by applying the self-calibration algorithm (see, e.g., Readhead \\& Wilkinson \\cite{Readhead1978}) after phase-referencing. However, self-calibration, specially on observations of weak sources, may affect the resulting images in undesirable ways (see, e.g., Mart\\'i-Vidal \\& Marcaide \\cite{mar08b}). The correlator model includes contributions from the dry troposphere, the Earth orientation parameters, gain corrections for the sampling, and feed rotation of the alt-azimuthal mounts of the antennae. An imperfect modelling in any of these contributions and the loss of coherence of the radio waves within the time elapsed between consecutive observations of a given source, have an impact on the quality of phase-referencing. Some authors have stated that the loss of signal coherence in phase-referencing is linearly dependent on the separation between the calibrator source and the target source (Beasley \\& Conway \\cite{bea95}). However, in a recent publication Mart\\'{\\i}-Vidal et al. (\\cite{mar10}) suggest a phenomenological model different from that of Beasley \\& Conway (\\cite{bea95}), based on Monte Carlo simulations of atmospheric turbulences. To empirically establish this dependence, we need to compare phase-referenced images to those obtained from the self-calibrated visibilities (i.e., the images obtained by applying the {\\em optimum} phase gains; those that maximize the signal coherence of the target sources). For this purpose, both calibrator and target must be strong enough to generate fringes with a signal-to-noise ratio (SNR) high enough to allow for an accurate estimate of the phase, delay, and rate of the fringe peaks, thus avoiding the bias effects related to self-calibration of weak signals (e.g., Mart\\'i-Vidal \\& Marcaide \\cite{mar08b}). The S5 polar cap sample (a subset of 13 sources from the S5 survey, see K\\\"uhr et al. \\cite{Kuhr1981}, Eckart et al. \\cite{Eckart1986}) is an ideal set of sources to perform such a study. The spectra of these sources are relatively flat at radio wavelenghts and their flux densities range from hundreds of mJy to several Jy. Therefore it is possible to study the loss of phase coherence as a function of observing frequency and source separation. We observed the S5 polar cap sample at 8.4\\,GHz and 15\\,GHz. These observations are part of a large astrometry programme (Ros et al. \\cite{ros01}, P\\'erez-Torres et al. \\cite{per04}, Mart\\'{\\i}-Vidal et al. \\cite{mar08}, and Jim\\'enez-Monferrer et al. in preparation) and were performed in phase-reference mode with many different calibrator/target combinations. From these observations, we studied the loss of signal coherence in phase-referenced observations by comparing phase-referenced and self-calibrated images for all possible calibrator/target combinations allowed by the observations. In the next section we describe our observations and the process of data calibration and reduction. In Sect. \\ref{sec:results} we report on the results obtained. In Sect. \\ref{sec:conclusions} we summarize our conclusions. ", "conclusions": "\\label{sec:conclusions} We report how phase referencing affects the signal coherence (and the fidelity of the images) in VLBI observations at different frequencies (8.4\\,GHz and 15\\,GHz) and for different calibrator-to-target separations (from 1.5 to 20.5 degrees). We determined the loss of dynamic range and peak flux-density of the phase-referenced images and compared the results with the model predictions given in Mart\\'i-Vidal et al. (\\cite{mar10}). The dynamic range of the phase-referenced images is strongly limited by the atmosphere at all frequencies and for all calibrator-to-target separations. The maximum achievable dynamic range using the VLBA is given by Eqs. \\ref{DL1} and \\ref{FAT}, with $K \\sim 12.4$ (if $\\nu$ is given in GHz and $\\Delta t$ in hours). If the target source is not strong (as is usually the case), the thermal noise of the receiving system cannot be ignored and the dynamic range should be estimated with Eqs. \\ref{DPR} and \\ref{FAT}. The flux-density (computed as the peak flux density of the phase-referenced image in units of the real peak flux-density of the source) decreases as the separation to the calibrator increases and is given by Eqs. \\ref{PEAKRATIO} and \\ref{FAT}, with $k_1 \\sim 63$ and $k_2 \\sim 1.87$. This model roughly predicts the peak ratios observed at 8.4\\,GHz for calibrator-to-target separations below $\\sim 15$ degrees (the results at 15\\,GHz are too noisy to reach a robust conclusion). For separations larger than 15 degrees, the observed peak ratios are higher than the model predictions, possibly due to a saturation in the power spectrum of the tropospheric turbulences at large scales, which was not considered in Mart\\'i-Vidal et al. (\\cite{mar10}). It is remarkable that for relatively small separations (below 5 degrees), which are typical in many phase-referencing observations, the flux-density loss can be as large as 20\\% at 8.4\\,GHz and $30-40$\\% at 15\\,GHz (and even larger at higher frequencies, according to Eq. \\ref{PEAKRATIO}). It must be also taken into account that the phase-referenced observations here reported were performed under good weather conditions and when the sources were close to their transits at nearly all stations (except Mauna Kea and St. Croix, see Mart\\'i-Vidal et al. \\cite{mar08} for details on the observing schedule). Therefore even larger flux-density losses and lower dynamic ranges may be obtained when the observing conditions are far from optimal. However, we must also notice that the typical calibrator-to-calibrator cycle times in our observations were about $120-180$ seconds. For observations at 8.4\\,GHz, these cycle times are of the order of (and slightly shorter than) the coherence time of the signal (which takes typical values of $180-240$ seconds in VLBI observations). Therefore it is not expected that a changing atmosphere may have introduced strong effects in the coherence of the phase-referenced visibilities at 8.4\\,GHz. At 15\\,GHz, the shorter coherence time (which takes typical values of $100-140$ seconds in VLBI observations) might be short enough to imply an additional phase degradation in the phase-referenced visibilities due to a changing atmosphere. Nevertheless, this issue is unlikely to significantly affect the results of our analysis, because the time evolution of the self-solutions of the antenna-based phase gains at 15\\,GHz are smooth, which indicates that the weather conditions were good enough to allow for a well-behaved phase connection between contiguous scans of each source." }, "1003/1003.5548_arXiv.txt": { "abstract": "Several anomalies appear to be present in the large-angle cosmic microwave background (CMB) anisotropy maps of WMAP, including the alignment of large-scale multipoles. Models in which isotropy is spontaneously broken (e.g., by a scalar field) have been proposed as explanations for these anomalies, as have models in which a preferred direction is imposed during inflation. We examine models inspired by these, in which isotropy is broken by a multiplicative factor with dipole and/or quadrupole terms. We evaluate the evidence provided by the multipole alignment using a Bayesian framework, finding that the evidence in favor of the model is generally weak. We also compute approximate changes in estimated cosmological parameters in the broken-isotropy models. Only the overall normalization of the power spectrum is modified significantly. ", "introduction": "Our understanding of cosmology has advanced extremely rapidly in the past decade. These advances are due in large part to observations of cosmic microwave background (CMB) anisotropy, particularly the data from the Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{wmap1yr1,wmap1yr2,wmap1,wmap7yrbasic}. As a result of these and other observations, a ``standard model'' of cosmology has emerged, consisting of a Universe dominated by dark energy and cold dark matter, with a nearly scale-invariant spectrum of Gaussian adiabatic perturbations \\cite{wmap7yrparams,wmap7yrinterp} of the sort that would naturally be produced in an inflationary epoch. The overall consistency of the CMB data with this model is quite remarkable. In particular, the CMB observations are very nearly Gaussian, and the angular power spectrum matches theoretical models very well from scales of tens of degrees down to arcminutes. However, several anomalies have been noted on the largest angular scales, including a lack of large-scale power \\cite{wmap1yr2,copi2,dOCTZH}, alignment of low-order multipoles \\cite{copi1,schwarz,dOCTZH,hajian}, and hemispheric asymmetries \\cite{eriksen2004,freeman,hansen}. Some anomalies seem to be associated with the ecliptic plane, suggesting the possibility of a systematic error associated with the WMAP scan pattern, perhaps related to coupling of the scan pattern with the asymmetric beam \\cite{hanson}. If the anomalies have cosmological significance, then naturally the correlation with the ecliptic plane must be a coincidence. The significance of and explanations for these puzzles are hotly debated. In particular, it is difficult to know how to interpret a posteriori statistical significances: when a statistic is invented to quantify an anomaly that has already been noticed, the low $p$-values for that statistic cannot be taken at face value. One can (and from a formal statistical point of view, arguably one must) dismiss this entire subject on the ground that all such anomalies are characterized only by invalid a posteriori statistics \\cite{wmap7yranomalies}. Nonetheless, the number and nature of the anomalies (in particular, the fact that several seem to pick out the same directions on the sky) seem to suggest that there may be something to explain in the data. Given the potential importance of new discoveries about the Universe's largest observable scales, and the difficulty in obtaining a new data set that would allow for a priori statistical analysis, we believe that the potential anomalies are worth further examination. In this paper, we will tentatively assume that there is a need for an explanation and consider what that explanation might be. One of the most robust of the large-scale anomalies found in WMAP is a lack of large-scale power, as quantified either by the low quadrupole or the vanishing of the two-point correlation function at large angles \\cite{wmap1yr2,copi2,dOCTZH}. If this anomaly is real, then it provides strong evidence {\\it against} a broad class of nonstandard models. To be specific, all models in which a statistically independent contaminant (whether due to a foreground, systematic error, or exotic cosmology) is added to the data will necessarily fare worse than the standard model in explaining this anomaly \\cite{gordon,bunnbourdon}. There is a simple reason for this: a statistically independent additive contaminant always increases the root-mean-square power in any given mode, reducing the probability of finding low power. It is natural, therefore, to seek an explanation of the anomalies among models that do not involve a mere additive contaminant. One simple phenomenological model is a multiplicative contaminant, in which the original statistically isotropic CMB signal $T^{(0)}(\\theta,\\phi)$ is modulated by a multiplicative factor, leading to an observed signal \\beq T(\\theta,\\phi)=f(\\theta,\\phi)T^{(0)}(\\theta,\\phi). \\label{eq:modulate} \\eeq This model arises naturally in the framework of spontaneous isotropy breaking by a scalar field \\cite{gordon}. Moreover, models based on the existence of a vector field specifying a preferred direction during inflation \\cite{ackerman,bohmer} produce similar modulation, but with $f$ having specifically a quadrupolar form. To be precise, the modulation in these models takes place in the primordial power spectrum $P(\\vec k)$, which acquires a quadrupolar dependence on the direction of $\\vec k$. The full effect on the CMB anisotropy is more complicated than the above model, but the dominant effect on large scales is, at least approximately, a quadrupolar moduloation of the above form.\\footnote{The stability of the specific model of ref.\\ \\cite{ackerman} has been questioned \\cite{acwunstable}; nonetheless, we believe it is worthwhile to consider models of this general class.} Since our goal is to explain the observed large-scale anomalies while maintaining the success of the standard model on smaller scales, it is natural to consider models in which $f$ has power only on large scales. We will consider three classes of model: one in which $f$ has only monopole and dipole terms, one in which it has monopole and quadrupole, and one in which it has all three. We will refer to these as the dipole-only, quadrupole-only, and dipole-quadrupole models. The quadrupole-only model is inspired by the theory of a preferred direction during inflation, while the others are inspired by the general isotropy-breaking framework. This paper addresses the following central question. Do the broken-isotropy models provide an explanation for one of the main observed anomalies, namely the surprising alignment between the quadrupole and octupole (multipoles $l=2,3$)? To examine this question, we choose statistics to quantify the anomaly and use these statistics to assess goodness of fit of the data to the different models. Several different statistics are chosen in order to assess the robustness of the results. Because the statistics are most naturally computed in spherical harmonic space, we use the all-sky internal linear combination (ILC) maps from the five-year WMAP data release \\cite{wmap5yrbasic}. There is bound to be residual foreground contamination in the ILC maps \\cite{eriksenilc04,eriksenilc05}. Section \\ref{sec:foregrounds} contains a brief discussion of the effects of this contamination. Naturally, because the anisotropic models have more free parameters than the standard model (and indeed include the standard model as a special case), there will generically be parameter choices that make the anisotropic model fit the data better. We adopt the Bayesian evidence criterion to assess whether this improved fit is sufficient to justify the additional complexity of the anisotropic model. Bayesian evidence has been used in addressing this sort of question in the past \\cite{landmag,eriksen2007,hoftuft}. Although some controversy has arisen over its use in cosmology (e.g., \\cite{marshall,lidmukpar,liddle,efstathioubayes}), in this context it is both a simple and a natural criterion to adopt. In some cases, the Bayesian evidence ratios are greater than one, meaning that one's assessment of the probability that the broken-isotropy models are true should rise as a result of the CMB anomalies. However, in all cases, the improvement is modest, providing at most weak support for the adoption of the anisotropic models. We also consider the changes in parameter estimates that would arise if the anisotropic models are correct. To be specific, because we assume that the modulation is a perturbation to the standard model, we assume that the unmodulated temperature map $T^{(0)}$ is derived from the cosmological parameters in the usual way -- i.e., its power spectrum is given by CMBFAST \\cite{cmbfast}. If there is a nonconstant modulation function $f$, then parameter estimates from based on the observed data will naturally differ from the true values. We estimate the resulting parameter shifts, finding them to be minor. The remainder of this paper is structured as follows. Section \\ref{sec:simulate} specifies precisely the anisotropic models under consideration and describes how we simulate these models. In Section \\ref{sec:bayes}, we review the method for computing Bayesian evidence ratios. Section \\ref{sec:align} contains our main results, indicating the degree to which the multipole alignment, quantified in several different ways, favor the broken-isotropy models. In Section \\ref{sec:params} we quantify the degree to which best-fit cosmological parameters are modified by changing from the standard model to the broken-isotropy models. Section \\ref{sec:foregrounds} discusses some aspects of the issue of foreground contamination. Finally, we provide a brief discussion of our results in Section \\ref{sec:discussion}. ", "conclusions": "\\label{sec:discussion} The various anomalies that have been noted in the large-angle CMB may provide hints of departures from the standard cosmological model, possibly including violations of statistical isotropy. Although the statistical significance of these anomalies is difficult or even impossible to quantify a posteriori, these possibilities are exciting enough to warrant closer examination. We have considered several classes of physically-motivated models that might explain the anomalies. We have calculated Bayesian evidence ratios to assess the degree to which the purported anomalies in the multipoles $l=2,3$ favor the anisotropic models over the standard model. According to the pioneering work of Jeffreys \\cite{jeffreys}, a Bayesian evidence ratio constitutes ``substantial'' evidence if $\\ln\\Lambda>1$ and ``strong'' evidence if $\\ln\\Lambda>2.5$. As the results in the Section \\ref{sec:align} make clear (note that what is plotted in each case is $\\Lambda$, not $\\ln\\Lambda$), only for the most judicious choice of prior do the tests performed here reach the ``substantial'' level, and they never come close to being ``strong.'' Of course, Jeffreys's criteria are somewhat arbitrary, but in this case they seem to describe the situation fairly well. Recall that the evidence ratio $\\Lambda$ is simply the factor by which the ratio of prior probabilities must be adjusted, in the light of the observations, in order to get the posterior probability ratio. Presumably, the prior probability distribution assigns very low weight to the less natural anisotropic models, so even after applying an evidence ratio $\\Lambda\\sim 3$, the anisotropic models are still considered unlikely. One would require an exponentially large evidence ratio before assigning significant probability to the anisotropic models. We used several different statistical approaches to to characterize the observed multipole alignment. Some ($\\lambda,S$) are adopted from previous work, while others $(A,p)$ are of our own devising. In the latter case, we attempted to minimize (although not eliminate) the problem of a posteriori statistics by choosing a method blindly that seemed to us to naturally encapsulate the observed phenomena with minimal arbitrary choices. In any case, the general consistency of the results based on the different statistics indicates that the approach we have followed is robust. We have estimated the changes in cosmological parameter estimates that would arise if the anisotropic models were shown to be correct. The chief effect of the modulation is on the estimate of the overall power spectrum normalization, which would of course have consequences for studies of large-scale structure. Our calculations are valid only if the modulation is applied to the CMB at all $l$-values measured by WMAP. If a more complicated model is correct (e.g. \\cite{hou}), in which only some scales are modulated, then the parameter changes would presumably be smaller. We have used simulations of the ILC mapmaking process to evaluate the degree to which foreground contamination might affect our results. The statistic $S$ appears least affected by this problem: ILC maps with high values of $S$ are very likely to correspond to high values of $S$ in the intrinsic CMB. A thorough treatment of foregrounds in our analysis would generically reduce the (already modest) enhancements in the evidence ratio, so due to the effects of foregrounds our results can be regarded as upper limits. In this paper, we have tentatively adopted the point of view that there are anomalies to be explained. Of course, one would greatly prefer to settle this question in a way that was not plagued by the problem of a posteriori statistics. To do this, we would require a new data set that probes similar scales to the large-angle CMB. All-sky polarization maps may provide some insight into these issues \\cite{frommert,dvorkin}. Another possibility is to survey the ``remote quadrupole'' signal found in the polarization of CMB photons scattered in distant clusters \\cite{kamloeb}, which can be used to reconstruct information on gigaparsec-scale perturbations \\cite{bunnrq,abramorq}. Although gathering data on these scales is a difficult task, the potential for learning about the structure of the Universe on the largest observable scales makes it worth pursuing." }, "1003/1003.5915_arXiv.txt": { "abstract": "In cosmic inflation driven by a scalar gauge singlet field with a tree level Higgs potential, the scalar to tensor ratio $r$ is estimated to be larger than $0.036$, provided the scalar spectral index $n_s \\geq 0.96$. We discuss quantum smearing of these predictions arising from the inflaton couplings to other particles such as GUT scalars, and show that these corrections can significantly decrease $r$. However, for $n_s \\geq 0.96$, we obtain $r \\geq 0.02$ which can be tested by the Planck satellite. ", "introduction": " ", "conclusions": "" }, "1003/1003.2829_arXiv.txt": { "abstract": "We investigate the impact of the existence of a primordial magnetic field on the filter mass, characterizing the minimum baryonic mass that can form in dark matter (DM) haloes. For masses below the filter mass, the baryon content of DM haloes are severely depressed. The filter mass is the mass when the baryon to DM mass ratio in a halo is equal to half the baryon to DM ratio of the Universe. The filter mass has previously been used in semianalytic calculations of galaxy formation, without taking into account the possible existence of a primordial magnetic field. We examine here its effect on the filter mass. For homogeneous comoving primordial magnetic fields of $B_0 \\sim 1$ or 2 nG and a reionization epoch that starts at a redshift $z_s=11$ and is completed at $z_r=8$, the filter mass is increased at redshift 8, for example, by factors 4.1 and 19.8, respectively. The dependence of the filter mass on the parameters describing the reionization epoch is investigated. Our results are particularly important for the formation of low mass galaxies in the presence of a homogeneous primordial magnetic field. For example, for $B_0\\sim 1\\nG$ and a reionization epoch of $z_s\\sim 11$ and $z_r\\sim7$, our results indicate that galaxies of total mass $M\\sim5 \\times 10^8\\msun$ need to form at redshifts $z_F\\gtrsim 2.0$, and galaxies of total mass $M\\sim10^8\\msun$ at redshifts $z_F\\gtrsim 7.7$. ", "introduction": "The formation of galaxies is one of the most important areas in cosmology. Within the simplified hierarchical scenario of a Universe governed by a cosmological constant and cold dark matter (DM), the density profiles of forming DM haloes have been well characterized. Numerical calculations show that the first generation of galaxies formed at very high redshifts inside collapsing haloes starting at $z \\sim 65$, corresponding to high peaks of the primordial DM density field \\citep{NNB}. CMB observations suggest that reionization began at high redshifts. This means that a high abundance of luminous objects must have existed at that time, since these first luminous objects are expected to have heated and reionized their surroundings \\citep{rev,WL03,HH03,Cen}. The formation of a luminous object inside a halo necessarily requires the existence of baryonic gas inside the halo. Even in haloes that are too small for cooling via atomic hydrogen, the gas content can have substantial, and observable, astrophysical effects. In addition to the possibility of hosting astrophysical sources, such as stars, small haloes may produce a 21-cm signature \\citep{Kuhlen,Shapiro06,NB08, Furlanetto06}, and can block ionizing radiation and produce an overall delay in the global progress of reionization \\citep{bl02, iliev2, iss05, mcquinn07}. As shown by \\citet{Gnedin2000a} and \\citet{Gnedin1997}, both in the linear and non-linear regimes, the accretion of gas into DM haloes is suppressed below a characteristic mass scale called the filter mass, $M_F$. This mass scale coincides with the Jeans mass, $M_J$, if the latter does not vary in time. Otherwise, $M_F$ is a time average of $M_J$. Thus, an increase in the ambient pressure in the past, causes an increase in $M_J$ and suppresses the accretion of baryons into DM haloes in a cumulative fashion, producing an increase in $M_F$. Until now, studies focused on the UV heating of the neutral interstellar gas as the main source of pressure, for determining the filter mass. These results are widely used in many semianalytic models (e.g. \\citealt{Maccio2009}), particularly those designed to study the properties of small galaxies (due to the high redshift character of the UV heating). In this paper we add the effect of a possible homogeneous primordial magnetic field as another important source of ambient pressure. The magnetic field contributes to pressure support, which changes the Jeans mass and, consequently, the filter mass and the quantity of gas that is accreted by DM haloes. The paper is organized as follows. In section \\ref{sec:magnetic} we briefly review our knowledge about primordial magnetic fields and in section \\ref{sec:filter} we present our results, calculating the filter mass with a Jeans mass corrected for the presence of a primordial magnetic field. Finally, in section \\ref{sec:conclusions}, we present our conclusions. ", "conclusions": "\\label{sec:conclusions} The Jeans mass was modified in order to account for the presence of a homogeneous primordial magnetic field. We then calculated the implications on the modification of the filter mass. The modification leads to a significant increase in the filter mass. We found a change in the behaviour of the filter mass as a function of redshift, which becomes flatter with more intense primordial magnetic fields, as opposed to the tendency of a decline of the filter mass with redshift, which is observed in the absence of magnetic fields. Our results are shown in figures \\ref{fig:filter}--\\ref{fig:filterar}. To see the effect of a primordial magnetic field, let us examine its effect at a specific redshift, for example $z=z_s-1$. In figure \\ref{fig:filter}, for $z_s=11$ and $z_r=8$ we find that the filter mass increases from $M_F=5.0\\times 10^6\\msun$ for $B_0=0\\nG$ to $M_F=6.0\\times 10^7\\msun$ for $B_0=1\\nG\\,$ and $M_F=3.8\\times 10^8\\msun$ for $B_0=2\\nG$. In figure \\ref{fig:filteras}, for $z_r=8$ we find that for $z_s=12$ that the filter mass increases from $M_F=4.1\\times 10^6\\msun$ for $B_0=0\\nG$ to $M_F=5.8\\times 10^7\\msun$ for $B_0=1\\nG\\,$ and $M_F=3.8\\times 10^8\\msun$ for $B_0=2\\nG$. In figure \\ref{fig:filterar}, for $z_s=11$ we find that for $z_r=6$ that the filter mass increases from $M_F=5.0\\times 10^6\\msun$ for $B_0=0\\nG$ to $M_F=6.0\\times 10^7\\msun$ for $B_0=1\\nG\\,$ and $M_F=3.8\t\\times 10^8\\msun$ for $B_0=2\\nG$. The results of figures \\ref{fig:filter}--\\ref{fig:filterar} are particularly important for the formation of low mass galaxies. For example, for a reionization epoch $z_s\\sim 11$ and $z_r\\sim 7$ and a homogeneous primordial magnetic field $B_0=1\\nG$, our results indicate that galaxies of total mass $M\\sim5 \\times10^8\\msun$ need to form at redshifts $z_F\\gtrsim 2.0 $ , and galaxies of total mass $M\\sim 10^8\\msun$ need to form at $z_F\\gtrsim 7.7 $." }, "1003/1003.4515_arXiv.txt": { "abstract": "Primordial Black Holes (PBHs), which may have been created in the early Universe, are predicted to be detectable by their Hawking radiation. The Fermi Gamma-ray Space Telescope observatory offers increased sensitivity to the gamma-ray bursts produced by PBHs with an initial mass of $\\sim 5\\times 10^{14}$ g expiring today. PBHs are candidate progenitors of unidentified Gamma-Ray Bursts (GRBs) that lack X-ray afterglow. We propose spectral lag, which is the temporal delay between the high and low energy pulses, as an efficient method to identify PBH evaporation events with the Fermi Large Area Telescope (LAT). ", "introduction": "\\label{aba:sec1} The formation of Primordial Black Holes (PBHs) has been postulated in many theories of the early Universe (for a recent review see Ref.~\\refcite{C05}). Black holes of mass $M_{\\rm bh}$ continually emit Hawking radiation~\\cite{H} with a temperature of $T_{\\rm bh}= 1.06\\ {\\rm GeV}/\\left( M_{\\rm bh}/10^{13}\\ {\\rm g} \\right)$ in the form of all available fundamental particle species. The emitted particles decay quickly on astrophysical timescales into $\\gamma$, $\\nu_{e,\\mu,\\tau}$, $\\bar{\\nu}_{e,\\mu,\\tau}$, $p$, $\\bar{p}$, $e^+$ and $e^-$. PBHs with an initial mass\\cite{MCP} of $M_*\\sim 5\\times 10^{14}$ g should be expiring today with a burst of high energy particles including gamma-rays. The current upper limit on the number expiring today per volume per unit time is\\cite{MC} \\begin{equation} R\\lesssim 10^{-7}\\eta_{\\rm \\, local}\\, \\rm{pc^{-3} \\, yr^{-1}} \\label{aba:eq1} \\end{equation} where $\\eta_{\\rm \\, local}$ is the density enhancement of PBHs in the local region. Typically $\\eta_{\\rm \\, local}$ is $\\sim 10^6$ (for clustering in the Galactic halo) or larger. Such PBH bursts may be detectable by the Fermi Gamma-ray Space Telescope observatory's Large Area Telescope (LAT). Conversely, non-detection by the LAT may lead to tighter bounds on the PBH distribution. ", "conclusions": "The Fermi LAT offers greater sensitivity to local PBH bursts than ground-based detectors. We have proposed\\cite{T} spectral lag measurements (the temporal delay between high and low energy pulses) of the incoming light curve in two different energy bands as a method to identify PBH bursts. A PBH burst arriving at the detector will exhibit positive to negative evolution with increasing energy because the black hole temperature and $\\overline{E}_{\\gamma}$ increase over time as the black hole loses mass. Because spectral lag measurements require counts in only two energy bands, and not the full spectrum, spectral lag can be measured even for weak events that last for very short time scales. Work is in progress to calculate quantitative values for the PBH spectral lags for the characteristics of the Fermi LAT." }, "1003/1003.4665_arXiv.txt": { "abstract": "{We present observations of the recently discovered supernova 2008iz in M82 with the VLBI High Sensitivity Array at 22 GHz, the Very Large Array at frequencies of 1.4, 4.8, 8.4, 22 and 43 GHz, and the Chandra X-ray observatory. The supernova was clearly detected on two VLBI images, separated by 11 months. The source shows a ring-like morphology and expands with a velocity of $\\sim$ 23000 \\kms. The most likely explosion date is in mid February 2008. The measured expansion speed is a factor of $\\sim$2 higher than expected under the assumption that synchrotron self-absorption dominates the light curve at the peak, indicating that this absorption mechanism may not be important for the radio emission. We find no evidence for an asymmetric explosion. The VLA spectrum shows a broken power law, indicating that the source was still optically thick at 1.4 GHz in April 2009. Finally, we report upper limits on the X-ray emission from SN\\,2008iz and a second radio transient recently discovered by MERLIN observations. } ", "introduction": "Radio supernovae (RSNe) are rare events and difficult to study. So far only about two dozen have been detected \\citep[e.g][]{Weiler2002} and most of them are quite distant and rather weak. Only few radio supernovae have been imaged with Very Long Baseline Interferometry (VLBI) techniques and for only four has it been possible to study the evolution of the expanding shell (SN\\,1979C, SN\\,1986J, SN\\,1987A, SN\\,1993J). The best studied radio supernova so far is SN\\,1993J in M81. Following the expansion of the supernova allowed many different phenomena to be studied \\citep{Marcaide1997, Marcaide2009, Bietenholz2001, Bietenholz2003, Perez2001, Perez2002, Bartel2002, Bartel2007}, including a measurement of the expansion speed, the deceleration of the shock front, and the proper motion of the supernova shell, for which a limit was obtained. The recent discovery of a new bright radio supernova in M82, SN\\,2008iz, \\citep{BrunthalerMentenReid2009,BrunthalerMentenReid2009b} at a similar distance as SN\\,1993J offers the rare opportunity to study the evolution of another supernova in great detail and to make a comparison to SN\\,1993J. So far, SN\\,2008iz was only detected in the radio band, with the VLA at 22 GHz \\citep{BrunthalerMentenReid2009}, MERLIN at 5 GHz \\citep{BeswickMuxlowPedlar2009}, and the Urumqi telescope at 5 GHz \\citep{MarchiliMartiVidalBrunthaler2009}. There are no reported detections in visible light and \\cite{FraserSmarttCrockett2009} report only a non detection in the near infrared on 2009 June 11. The non-detections at other wavebands indicate that the supernova exploded behind a large gas or dust cloud in the central part of M82. Thus, it has not been possible to classify this supernova. However, since type Ia supernovae are not known to show strong radio emission, SN\\,2008iz is most likely a core collapse supernova, i.e. type Ib/c or II. Here we present the first VLBI images of SN\\,2008iz taken $\\sim$2.5 and $\\sim$13.5 month after the explosion, a radio spectrum (at an age of $\\sim$14.5 month) from 1.4 to 43 GHz and Chandra X-ray observations. Throughout the paper, we adopt a distance of 3.6 Mpc \\citep[based on a Cepheid distance to M81 determined by][] {FreedmanHughesMadore1994}. ", "conclusions": "\\subsection{Column density and extinction} \\label{coldens} The non-detection of SN\\,2008iz in the optical, infrared, and X-rays indicates that it exploded inside or behind a very dense cloud. Indeed, the $^{12}$CO (J=2$\\rightarrow$1) line intensity map in \\cite{WeissNeiningerHuettenmeister2001} shows a prominent cloud exactly at the position of the supernova (Fig.~\\ref{fig:co}). The line intensity at the position of SN\\,2008iz is $\\sim$ 1800 K kms$^{-1}$. Using the Galactic conversion factor X$_\\mathrm{CO}=1.6\\times 10^{20}$ cm$^{-2}$ (K kms$^{-1}$)$^{-1}$, this corresponds to a H$_2$ column density $N(\\rm{H}_2)$ of $\\sim 29\\times 10^{22}$ cm$^{-2}$. However, \\cite{WeissNeiningerHuettenmeister2001} find much smaller and spatially variable conversion factors from radiative transfer calculations, that lead to smaller H$_2$ column densities. At the position of SN\\,2008iz, the conversion factor is $\\sim 0.3\\times10^{20}$ cm$^{-2}$ (K kms$^{-1}$)$^{-1}$ (their Fig. 10). This leads to a H$_2$ column density of $\\sim 5.4\\times 10^{22}$ cm$^{-2}$. However, the CO observations were performed with a linear resolution of $\\sim$25 pc. Thus it is possible, that the supernova is located behind a smaller cloud with much higher column density. The $^{12}$CO line intensity at the position of the MERLIN transient is almost 4 times smaller ($\\sim$500 K kms$^{-1}$). \\begin{figure} \\centering \\includegraphics[width=9cm]{f8.eps} \\caption{$^{12}$CO (J=2$\\rightarrow$1) intensity map of the central region of M82 from \\cite{WeissNeiningerHuettenmeister2001} with the positions of SN\\,2008iz (star) and the MERLIN transient (cross). The contours correspond to 200, 400, 600, 800, 1200 1600, 2000, and 2400 K kms$^{-1}$. The resolution of the observation is $1.5''\\times 1.4''$ ($\\sim$25 pc).} \\label{fig:co} \\end{figure} Taking the latter value ($5.4\\times 10^{22}$ cm$^{-2}$) for N$({\\rm H}_2)$ and assuming that all of the column is between SN 2008iz and us, we derive a visual extinction, $A_V$, of 24.4 mag. Here we have used the relation between optical extinction and hydrogen nucleus column density, N(${\\rm H}$), derived by \\citet{GueverOezel2009} from X-ray absorption data of a sample of Galactic SNRs and assumed that all the hydrogen is in molecular form, i.e., N(${\\rm H}_2$) = 2~N(${\\rm H}$). Such a high value of the extinction would explain the lack of an optical counterpart. The derived extinction is much higher than toward SN\\,1993J, for which \\cite{Richmond1994} discuss values of $A_V = 0.7$ and 1.0 mag. With the extinction law given by \\citet{Cardelli1989}, $A_K = 0.114~A_V$, we calculate 2.8 mag for the $K-$band ($2.2~\\mu$m) extinction $A_K$. \\cite{FraserSmarttCrockett2009} report an even higher $K$-band extinction of up to 11 mag based on their non-detection (3$\\sigma$ upper limit on the \\ absolute $K$-band magnitude of -5 mag) and the assumption that the infrared light curve of SN\\,2008iz behaves similar to the one of SN\\,1993J. Thus, SN\\,2008iz is either very weak in the infrared compared to SN\\,1993J, or behind a smaller but even denser cloud than estimated here. \\subsection{Expansion velocity and synchrotron self-absorption} \\cite{Chevalier1998} proposed a way to estimate the mean expansion velocity of a supernova based on the radio-light curve and assuming that synchrotron self-absorption (SSA) dominates the light curve at the peak. If SSA is not the dominant absorption process at the peak of the light curve, then the estimate of the expansion velocity from the \\cite{Chevalier1998} model is a lower bound to the real expansion velocity of the radio shell. Following \\cite{Chevalier1998}, the estimated mean expansion velocity at the peak of the radio-light curve, assuming dominant SSA, is: \\begin{eqnarray} V_\\mathrm{SSA}(\\mathrm{km}\\,\\mathrm{s}^{-1}) = 5.3786\\times10^{6} ({\\beta\\phi})^{-0.053} \\left (\\frac{F_{\\mathrm{p}}}{1 \\mathrm{Jy}}\\right )^{0.47}\\times \\nonumber \\\\ \\left (\\frac{D}{1 \\mathrm{ Mpc}}\\right )^{0.95} \\left (\\frac{\\nu}{1.0~\\mathrm{ GHz}}\\right )^{-1} \\left (\\frac{t_{p}}{1 \\mathrm{ day}}\\right)^{-1} \\label{VelSSA} \\end{eqnarray} \\noindent where $\\beta$ is the ratio of the relativistic particle energy density to the magnetic field energy density \\cite[if we assume energy equipartition, $\\beta = \\frac{4}{3(1+k)}$, where $k$ ranges from 1 to 2000, see ][chapter 7]{Pacholczyk1970}, $\\phi$ is the filling factor of the emitting region to a sphere (we assume a shell of width equal to 30\\% of the outer radius, which yields $\\phi = 0.66$), $F_{\\mathrm{p}}$ is the flux density at the peak, $D$ is the distance, $\\nu$ is the observing frequency, $t_{p}$ is the supernova age at the peak, and we use a spectral index $\\alpha = -1$ \\cite[see Eqs. 11 and 13 of ][]{Chevalier1998}. For SN2008iz, Eq. \\ref{VelSSA} yields a mean expansion velocity (depending on k) in the range $8100 - 11800$\\,km\\,s$^{-1}$ at the peak of the 5\\,GHz radio light curve. Using an expansion index of 0.89, these velocities translate into a mean expansion velocity in the range $7000 - 10300$\\,km\\,s$^{-1}$ at the epoch of 27 April 2009. This range of velocities is a factor $\\sim 2$ smaller than the velocities estimated from our VLBI observations, thus indicating that, in contrast to the case of SN1993J, SSA may not be an important absorption mechanism in the SN2008iz radio emission. This is also consistent with the results in \\cite{MarchiliMartiVidalBrunthaler2009} who were able to model the radio light curve of SN2008iz assuming that SSA effects are much smaller than free-free absorption (FFA) during the whole supernova expansion. \\subsection{Comparison with other type II radio supernovae} \\begin{table} \\caption{Comparison between SN\\,1993J and SN\\,2008iz.} \\label{tab:comp} \\begin{tabular}{lccc} \\hline\\hline Property&&SN\\,1993J&SN\\,2008iz \\\\ \\hline spectral index $\\alpha$& &-0.99$^a$&-1.08$^b$\\\\ L$_{5 \\mathrm{GHz}}$& $[10^{27}$erg~s$^{-1}$~Hz$^{-1}]$&1.5$^c$&2.5$^d$\\\\ t$_\\mathrm{peak}$-t$_0$& [days]&180$^c$&$\\sim$120$^d$\\\\ v$_\\mathrm{VLBI}$&[\\kms]&14900$^e$ & 21200$^b$\\\\ L$_\\mathrm{X-ray\\,at\\,t\\sim220\\,days}$& $[10^{38}$ erg~s$^{-1}]$&$\\sim8^f$ &$<15^b$\\\\ \\hline % \\end{tabular} References: a) \\cite{vanDykWeilerSramek1994}; b) this work; c) \\cite{WeilervanDykMontes1998}; d) \\cite{MarchiliMartiVidalBrunthaler2009}; e) \\cite{MarcaideAlberdiRos1995}; f) \\cite{ZimmermannAschenbach2003}. \\end{table} Estimates of the expansion velocities of other type II radio supernovae have been estimated from VLBI observations, and very different results have been obtained. For instance, the mean expansion velocity of SN\\,1979C during the first year after explosion is estimated to be $\\sim 10000 - 11000$ \\kms\\, \\citep{BartelBietenholz2003, MarcaideMartiVidalPerezTorres2009}; for SN\\,1986J, a velocity of $\\sim$ 14700 \\kms\\, was obtained by \\cite{PerezTorresAlberdiMarcaide2002}, while \\cite{BietenholzBartelRupen2002} find 20000 \\kms\\, 3 month after the explosion; \\cite{Staveley-SmithBriggsRowe1993} report a mean expansion speed of $\\sim$ 35000 \\kms\\ during the first years for SN\\,1987A before it slowed down to $\\sim$ 4800 \\kms; for SN\\,2004et, the expansion velocity was $>$ 15700 \\kms\\, \\citep{MartiVidalMarcaideAlberdi2007}; and for SN\\,2008ax, an expansion velocity as large as 52000 \\kms\\, was obtained \\citep{MartiVidalMarcaideAlberdi2009}. Estimates of the expansion velocities of other supernova remnants in M82 (the host galaxy of SN\\,2008iz), have been also reported, which range between $\\sim$ 1500 and 11000 \\kms\\, \\citep{BeswickRileyMartiVidal2006}. These later velocities are much higher (a factor of $3-22$) than the predicted velocities from the model of \\cite{ChevalierFransson2001}, based on the high pressure expected in the interstellar medium (ISM) of M\\,82. The expansion velocity reported in this paper for SN2008iz is indeed a factor $\\sim 40$ larger than the predicted velocities in \\cite{ChevalierFransson2001}, although of the same order of magnitude than the velocities reported in \\cite{BeswickRileyMartiVidal2006} for the other remnants in M82, and the typical velocities of the other type II supernovae observed to date. \\cite{WeilervanDykMontes1998} find a correlation between peak radio luminosity at 5 GHz and the time between the explosion and the peak in the 5 GHz light curve for type II supernovae. The 5 GHz light curve of SN\\,2008iz from \\cite{MarchiliMartiVidalBrunthaler2009} gives a peak luminosity of $\\sim 2.5\\times10^{27}$erg~s$^{-1}$~Hz$^{-1}$ at $\\sim$120 days after the explosion. These values are well within the scatter of the correlation. Thus it seems plausible that SN\\,2008iz is also a type II supernova. Since SN\\,2008iz and SN\\,1993J are located at very similar distances, this allows a detailed comparison between these two supernovae. Several properties of both supernovae are summarized in Table~\\ref{tab:comp}. The radio spectral indices, the peak radio luminosities, rise times, and early VLBI expansion velocities are similar (considering that the rise times and peak radio luminosities can vary by several orders of magnitudes for type II radio supernovae). The non-detection in X-rays of SN\\,2008iz can be attributed to absorption by the dense molecular cloud seen in the CO data." }, "1003/1003.4386_arXiv.txt": { "abstract": "{ We perform kinetic simulations of diffusive shock acceleration (DSA) in Type Ia supernova remnants (SNRs) expanding into a uniform interstellar medium (ISM). Bohm-like diffusion due to self-excited Alfv\\'en waves is assumed, and simple models for Alfv\\'enic drift and dissipation are adopted. Phenomenological models for thermal leakage injection are considered as well. We find that the preshock gas temperature is the primary parameter that governs the cosmic ray (CR) acceleration efficiency and energy spectrum, while the CR injection rate is a secondary parameter. For SNRs in the warm ISM of $T_0 \\lsim 10^5$K, if the injection fraction is $\\xi \\gsim 10^{-4}$, the DSA is efficient enough to convert more than 20 \\% of the SN explosion energy into CRs and the accelerated CR spectrum exhibits a concave curvature flattening to $E^{-1.6}$, which is characteristic of CR modified shocks. Such a flat source spectrum near the knee energy, however, may not be reconciled with the CR spectrum observed at Earth. On the other hand, SNRs in the hot ISM of $T_0\\approx 10^6$K with a small injection fraction, $\\xi < 10^{-4}$, are inefficient accelerators with less than 10 \\% of the explosion energy getting converted to CRs. Also the shock structure is almost test-particle like and the ensuing CR spectrum can be steeper than $E^{-2}$. With amplified magnetic field strength of order of 30$\\mu$G, Alfv\\'en waves generated by the streaming instability may drift upstream fast enough to make the modified test-particle power-law as steep as $E^{-2.3}$, which is more consistent with the observed CR spectrum. } ", "introduction": "It is believed that most of the Galactic cosmic rays (CRs) are accelerated in the blast waves driven by supernova (SN) explosions (\\eg Blandford \\& Eichler 1987, Reynolds 2008 and references therein). If about 10 \\% of Galactic SN luminosity, $L_{SN}\\approx 10^{42} {\\rm erg~s^{-1}}$, is transfered to the CR component, the diffusive shock acceleration (DSA) at supernova remnants (SNRs) can provide the CR luminosity, $L_{CR} \\approx 10^{41} {\\rm erg~s^{-1}}$ that escapes from the Galaxy. Several time-dependent, kinetic simulations of the CR acceleration at SNRs have shown that an order of 10 \\% of the SN explosion energy can be converted to CRs, when a fraction $\\sim 10^{-4}$ of incoming thermal particles are injected into the CR population at the subshock (\\eg Berezhko, \\& V\\\"olk 1997; Berezhko et al 2003; Kang 2006). X-ray observations of young SNRs such as SN1006 and RCW86 indicate the presence of 10-100 TeV electrons emitting nonthermal synchrotron emission immediately inside the outer SNR shock (Koyama \\etal 1995; Bamba \\etal 2006, Helder \\etal 2009). They provide clear evidence for the efficient acceleration of the CR electrons at SNR shocks. Moreover, HESS gamma-ray telescope detected TeV emission from several SNRs such as RXJ1713.7-3946, Cas A, Vela Junior, and RCW86, which may indicate possible detection of $\\pi^0$ $\\gamma$- rays produced by nuclear collisions of hadronic CRs with the surrounding gas (Aharonian \\etal 2004, 2009; Berezhko \\& V\\\"olk 2006; Berezhko \\etal 2009; Morlino \\etal 2009, Abdo \\etal 2010). It is still challenging to discern whether such emission could provide direct evidence for the acceleration of hadronic CRs, since $\\gamma$-ray emission could be produced by inverse Compton scattering of the background radiation by X-ray emitting relativistic electrons. More recently, however, Fermi LAT has observed in GeV range several SNRs interacting with molecular clouds, providing some very convincing evidence of $\\pi^0$ decay $\\gamma$-rays (Abdo \\etal 2009, 2010). In DSA theory, a small fraction of incoming thermal particles can be injected into the CR population, and accelerated to very high energies through their interactions with resonantly scattering Alfv\\'en waves in the converging flows across the SN shock (\\eg Drury \\etal 2001). Hence the strength of the turbulent magnetic field is one of the most important ingredients, which govern the acceleration rate and in turn the maximum energy of the accelerated particles. If the magnetic field strength upstream of SNRs is similar to the mean interstellar medium (ISM) field of $B_{\\rm ISM} \\sim 5 \\mu$G, the maximum energy of CR ions of charge $Z$ is estimated to be $E_{\\rm max} \\sim 10^{14}Z$ eV (Lagage \\& Cesarsky 1983). However, high-resolution X-ray observations of several young SNRs exhibit very thin rims, indicating the presence of magnetic fields as strong as a few $100 \\mu$G downstream of the shock (\\eg Bamba \\etal 2003, Parizot \\etal 2006). Moreover, theoretical studies have shown that efficient magnetic field amplification via resonant and non-resonant wave-particle interactions is an integral part of DSA (Lucek \\& Bell 2000, Bell 2004). If there exist such amplified magnetic fields in the upstream region of SNRs, CR ions might gain energies up to $E_{\\rm max} \\sim 10^{15.5}Z$ eV, which may explain the all-particle CR spectrum up to the second knee at $\\sim10^{17}$ eV with rigidity-dependent energy cutoffs. A self-consistent treatment of the magnetic field amplification has been implemented in several previous studies of nonlinear DSA (\\eg Amato \\& and Blasi 2006, Vladimirov et al. 2008). In Kang 2006 (Paper I, hereafter), we calculated the CR acceleration at typical remnants from Type Ia supernovae expanding into a uniform interstellar medium (ISM). With the upstream magnetic fields of $B_0=30\\mu$G amplified by the CR streaming instability, it was shown that the particle energy can reach up to $10^{16}Z$ eV at young SNRs of several thousand years old, which is much higher than what Lagage \\& Cesarsky predicted. But the CR injection and acceleration efficiencies are reduced somewhat due to faster Alfv\\'en wave speed. With the particle injection fraction $\\sim 10^{-4}-10^{-3}$, the DSA at SNRs is very efficient, so that up to 40-50 \\% of the explosion energy can be transferred to the CR component. We also found that, for the SNRs in the warm ISM ($T_0=10^4$K), the accelerated CR energy spectrum should exhibit a concave curvature with the power-law slope, $\\alpha$ (where $N(E)\\propto E^{-\\alpha}$) flattening from 2 to 1.6 at $E > 0.1$ TeV. In fact, the concavity in the CR energy spectrum is characteristic of strong ($M>10$) CR modified shocks when the injection fraction is greater than $10^{-4}$. (\\eg Malkov \\& Drury 2001, Berezhko \\& V\\\"olk 1997, Blasi \\etal 2005) Recently, Ave \\etal (2009) have analyzed the spectrum of CR nuclei up to $\\sim 10^{14}$ eV measured by TRACER instrument and found that the CR spectra at Earth can be fitted by a single power law of $J(E) \\propto E^{-2.67}$. Assuming an energy-dependent propagation path length ($\\Lambda \\propto E^{-0.6}$), they suggested that a soft source spectrum, $N(E)$ with $\\alpha \\sim 2.3-2.4$ is preferred by the observed data. However, the DSA predicts that $\\alpha=2.0$ for strong shocks in the test-particle limit and even smaller values for CR modified shocks in the efficient acceleration regime as shown in Paper I. Thus in order to reconcile the DSA prediction with the TRACER data the CR acceleration efficiency at typical SNRs should be minimal and perhaps no more than 10 \\% of the explosion energy transferred to CRs (\\ie test-particle limit). Moreover, recent Fermi-LAT observations of Cas A, which is only 330 years old and has just entered the Sedov phase, indicate that only about 2\\% of the explosion energy has been transfered to CR electrons and protons, and that the soft proton spectrum with $E^{-2.3}$ is preferred to fit the observed gamma-ray spectrum (Abdo \\etal 2010). According to Paper I, such inefficient acceleration is possible only for SNRs in the hot phase of the ISM and for the injected particle fraction smaller than $10^{-4}$. One way to soften the CR spectrum beyond the canonical test-particle slope ($\\alpha > 2$) is to include the Alfv\\'enic drift in the precursor, which reduces the velocity jump across the shock. Zirakashvili \\& Ptuskin (2008) showed that the Alfv\\'enic drift in the amplified magnetic fields both upstream and downstream can drastically soften the accelerated particle spectrum. We will explore this issue using our numerical simulations below. Caprioli \\etal (2009) took a different approach to reconcile the concave CR spectrum predicted by nonlinear DSA theory with the softer spectrum inferred from observed $J(E)$. They suggested that the CR spectrum at Earth is the sum of the time integrated flux of the particles that escape from upstream during the ST stage and the flux of particles confined in the remnant and escaping at later times. They considered several cases and found the injected spectrum could be softer than the concave instantaneous spectrum at the shock. The main uncertainties in their calculations are related with specific recipes for the particle escape. It is not well understood at the present time how the particles escape through a free escape boundary ($x_{\\rm esc}$) located at a certain distance upstream of the shock or through a maximum momentum boundary due to lack of (self-generated) resonant scatterings above an escape momentum. The escape or release of CRs accelerated in SNRs to the ISM remains largely unknown and needs to be investigated further. One of the key aspects of the DSA model is the injection process through which suprathermal particles in the Maxwellian tail get accelerated and injected into the Fermi process. However, the CR injection and consequently the acceleration efficiency still remain uncertain, because complex interplay among CRs, waves, and the underlying gas flow (\\ie self-excitation of waves, resonant scatterings of particles by waves, and non-linear feedback to the gas flow) is all model-dependent and not understood completely. In this paper, we adopted two different injection recipes based on thermal leakage process, which were considered previously by us and others. Then we have explored the CR acceleration at SNR shocks in the different temperature phases (\\ie different shock Mach numbers) and with different injection rates. Details of the numerical simulations and model parameters are described in \\S II. The simulation results are presented and discussed in \\S III, followed by a summary in \\S IV. ", "conclusions": "The evolution of cosmic ray modified shocks depends on complex interactions between the particles, waves in the magnetic field, and underlying plasma flow. We have developed numerical tools that can emulate some of those interactions and incorporated them into a kinetic numerical scheme for DSA, CRASH code (Kang \\etal 2002, Kang \\& Jones 2006). Specifically, we assume that a Bohm-like diffusion arises due to resonant scattering by Alfv\\'en waves self-excited by the CR streaming instability, and adopt simple models for the drift and dissipation of Alfv\\'en waves in the precursor (Jones 1993; Kang \\& Jones 2006). In the present paper, using the spherical CRASH code, we have calculated the CR spectrum accelerated at SNRs from Type Ia supernova expanding into a uniform interstellar medium. We considered different temperature phases of the ISM, since the shock Mach number is the primary parameter that determines the acceleration efficiency of DSA. One of the secondary parameters is the fraction of particles injected into the CR population, $ \\xi$, at the gas subshock. Since detailed physical processes that governs the injection are not known well, we considered two injection recipes that are often adopted by previous authors. The main difference between the two recipes is whether the ratio of injection momentum to thermal peak momentum, \\ie $p_{\\rm inj}/p_{th}$, is constant or depends on the subshock Mach number. It turns out the CR acceleration and the evolution of SNRs are insensitive to such difference as long as the injection fraction is similar. For example, the models with injection recipe A with $\\epsilon_B=0.23$ and the models with injection recipe B with $R_{\\rm inj}=3.6$ show almost the same results with similar injection fractions, $\\xi \\approx 10^{-3.5} -10^{-3}$. In general the DSA is very efficient for strong SNR shocks, if the injection fraction, $\\xi \\gsim 10^{-3.5}$. The CR spectrum at the subshock shows a strong concavity, not only because the shock structure is modified nonlinearly by the dominant CR pressure, but also because the SNR shock slows down in time during the ST stage. Thus the concavity of the CR spectrum in SNRs is more pronounced than that in plane-parallel shocks. However, the volume integrated spectrum, $G(p)$, (\\ie the spectrum of CRs confined by the shock including the particles in the upstream region) is much less concave, which is consistent with previous studies (\\eg Berezhko \\& V\\\"olk 1997). We have shown also that $G(p)$ approaches roughly to time-asymptotic states, since the CR pressure decreases as $t^{-6/5}$ while the volume increases as $R_{\\rm ST}^3 \\propto t^{6/5}$. This in turn makes the total CR energy converted ($E_{\\rm CR}$) asymptotes to a constant value. If we assume that CRs are released at the break-up of SNRs, then the source spectrum can be modeled as $N(E)dE=G(p)p^2dp$. However, it is a complex unknown problem how to relate $G(p)$ to the source spectrum $N(E)$ and further to the observed spectrum $J(E)$. In the warm ISM models ($T_0=3\\times10^4$K, $n_H=0.3{\\rm cm^{-3}}$), the CR acceleration at SNRs may be too efficient. More than 40\\% of the explosion energy ($E_o$) is tranferred to CRs and the source CR spectrum, $N(E)\\propto E^{-\\alpha}$ with $\\alpha \\approx 1.5$, is too flat to be consistent with the observed CR spectrum at Earth (Ave \\etal 2009). In these models with efficient injection and acceleration, the flow structure is significantly modified with $\\rho_2/\\rho_0 \\approx$ 7.2-7.5 for WA/WB models. In the intermediate temperature ISM models ($T_0=10^5$K, $n_H=0.03{\\rm cm^{-3}}$), the flow structure is still significantly modified with $\\rho_2/\\rho_0 \\approx$ 5.7-6.0 and the fraction of energy conversion, $E_{\\rm CR}/E_0 \\approx 0.2-0.4$ for MA/MB models. Only in the hot ISM model ($T_0=10^6$K, $n_H=0.003{\\rm cm^{-3}}$) with inefficient injection ($\\epsilon_B=0.2$ or $R_{\\rm inj}>3.8$), the shock structure is almost test-particle like with $\\rho_2/\\rho_0 \\approx$ 4.2-4.4 and the fraction of energy conversion, $E_{\\rm CR}/E_0 \\approx 0.1-0.2$ for HA/HB models. The predicted source spectrum $G(p)$ has a slope $q=4.1-4.3$ for $10^{11}< E< 10^{15}$ eV. Here drift of Alfv\\'en waves relative to the bulk flow upstream of the subshock plays an important role, since the modified test-particle slope, $q_{\\rm tp}=3(u_0-v_A)/(u_0-v_A-u_2)$, can be steeper than the canonical value of $q=4$ for strong unmodified shocks. With magnetic fields of $B_0=30\\mu$G, the Alfv\\'en speed is $v_A\\approx 1000 {\\rm km s^{-1}}$, and so the modified test-particle slope is $\\alpha \\approx 2.3$. This may imply that SN exploding into the hot ISM are the dominant sources of Galactic CRs below $10^{15}$eV. One might ask if the magnetic field amplification would take place in the case of such inefficient acceleration, since the magnetic field energy density is expected to be proportional to the CR pressure. An alternative way to enhance the downstream magnetic field was suggested by Giacalone \\& Jokipii (2007). They showed that the density fluctuations pre-existing upstream can warp the shock front and vortices are generated behind the curved shock surface. Then vortices are cascade into turbulence which amplifies magnetic fields via turbulence dynamo. Finally, in all models considered in this study, for Bohm-like diffusion with the amplified magnetic field in the precursor, indicated by X-ray observations of young SNRs, the particles could be accelerated to $E_{\\rm max} \\approx 10^{15.5}Z$eV. The drift and dissipation of {\\it faster} Alfv\\'en waves in the precursor, on the other hand, soften the CR spectrum and reduce the CR acceleration efficiency." }, "1003/1003.3273_arXiv.txt": { "abstract": "Galaxies are missing most of their baryons, and many models predict these baryons lie in a hot halo around galaxies. We establish observationally motivated constraints on the mass and radii of these haloes using a variety of independent arguments. First, the observed dispersion measure of pulsars in the Large Magellanic Cloud allows us to constrain the hot halo around the Milky Way: if it obeys the standard NFW profile, it must contain less than 4-5\\% of the missing baryons from the Galaxy. This is similar to other upper limits on the Galactic hot halo, such as the soft X-ray background and the pressure around high velocity clouds. Second, we note that the X-ray surface brightness of hot haloes with NFW profiles around large isolated galaxies is high enough that such emission should be observed, unless their haloes contain less than 10-25\\% of their missing baryons. Third, we place constraints on the column density of hot haloes using nondetections of OVII absorption along AGN sightlines: in general they must contain less than 70\\% of the missing baryons or extend to no more than 40 kpc. Flattening the density profile of galactic hot haloes weakens the surface brightness constraint so that a typical L$_*$ galaxy may hold half its missing baryons in its halo, but the OVII constraint remains unchanged, and around the Milky Way a flattened profile may only hold $6-13\\%$ of the missing baryons from the Galaxy ($2-4 \\times 10^{10} M_{\\odot}$). We also show that AGN and supernovae at low to moderate redshift - the theoretical sources of winds responsible for driving out the missing baryons - do not produce the expected correlations with the baryonic Tully-Fisher relationship and so are insufficient to explain the missing baryons from galaxies. We conclude that most of missing baryons from galaxies do not lie in hot haloes around the galaxies, and that the missing baryons never fell into the potential wells of protogalaxies in the first place. They may have been expelled from the galaxies as part of the process of galaxy formation. ", "introduction": "The so-called ``missing baryon problem'' actually refers to two separate, but related, issues. First is the realization that the mean density of all detected baryons in the local universe accounts for less than half of the cosmological baryon density. The universal baryon fraction is known precisely from, among other sources, the Wilkinson Microwave Anisotropy Probe 5-year data (Dunkley et al. 2009): $f _b \\equiv \\Omega_b / \\Omega_m = 0.171 \\pm 0.006$ and $\\Omega_b h^2 = 0.0227 \\pm 0.0006$. But studies that attempt to count baryons in the low-redshift universe (Fukugita, Hogan, and Peebles 1998, Fukugita and Peebles 2004, Nicastro et al. 2005a) find that only about a tenth of this figure is observed in stars and cool gas in galaxies, while another third of $\\Omega_b$ is divided between the hot plasma in the intracluster medium of galaxy clusters, and the cool intercluster gas detected in Ly$\\alpha$ absorption lines (for more on the latter, see e.g. Penton, Stocke, and Shull 2004). This result introduces the second missing baryon problem, which is the realization that most galaxies are severely baryon-depleted relative to the cosmological fraction (e.g. Bell et al. 2003). Based on cosmological simulations, the rest of the baryon budget is thought to reside in the ``warm-hot intergalactic medium'' (WHIM), at temperatures between $10^5$ and $10^7$ K (Cen and Ostriker 1999, Dav\\'{e} et al. 2001). This is thought to include most of the baryons associated with galaxies, since the optically-luminous portion of galaxies is severely baryon-depleted relative to the universal fraction (Dai et al. 2010). Observational evidence for the existence of the WHIM has subsequently begun to accrue, primarily through detections with ultraviolet and X-ray telescopes of highly-ionized Oxygen absorption lines (for a recent review, see Bregman 2007). These observations do not yet constrain the nature and distribution of the WHIM. Essentially all models agree that the WHIM exists as a cosmic filamentary web with some material in hot haloes around galaxies, but the mass fraction as a function of WHIM density is still unknown. Some numerical models (e.g Dav\\'{e} et al. 2009) find a diffuse WHIM at $z=0$ with most of its baryonic content lying outside of galactic haloes. However, many simulations either assume (e.g. Bower et al. 2006, Croton et al. 2006) or self-consistently derive (e.g. Cen and Ostriker 2006, Tang et al. 2009, Kim, Wise, and Abel 2009) the result that cosmologically significant reservoirs of baryons are embedded in hot haloes around massive galaxies ($kT \\sim $ a few hundred eV). These hot haloes have also been modeled in separate theoretical calculations (Fukugita and Peebles 2006, Dekel and Birnboim 2006, Kaufmann et al. 2009) and have important consequences in the galaxy formation process such as differentiating the ``blue cloud'' from the ``red sequence'' (Dekel and Birnboim 2006, Bouch\\'{e} et al. 2009). Observationally, the evidence for hot haloes is still unclear: gas at these temperatures has been observed in X-rays around disk galaxies out to a few kpc beyond the disk (Strickland et al. 2004, Li et al. 2006, Li, Wang, and Hameed 2007), but this gas seems to result from starburst activity in the host galaxy instead of emerging as a byproduct of galaxy formation (Rasmussen 2009), and its inferred mass falls significantly short of theoretical predictions (T\\\"{u}llmann et al. 2006). In this paper, we examine the implications of this paradigm in more detail. If most of the baryons associated with galaxies reside in hot haloes around the optically luminous part, there are several predictions that can be tested with existing observations. We seek to answer two questions. First, how much baryonic mass resides in these haloes? And, what can these haloes and the second missing baryon problem tell us about the process of galaxy formation? To answer these questions, we present several independent lines of argument. First, we consider observational constraints on the density of hot gas around the Milky Way using the dispersion measure of pulsars in the Large Magellanic Cloud, as well as the ambient pressure around high velocity clouds, and the galactic soft X-ray background. We then extend this argument to other galaxies. We constrain the characteristic density and radius of hot haloes using existing X-ray observations of absorption along quasar sightlines and emission from galactic haloes. We also discuss the energetics of driving a galactic-scale wind, arguing that existing mechanisms for expelling most of the missing baryons seem incomplete. Finally, we discuss flattened density profiles and other prospects for resolving the second missing baryon problem. Before engaging these arguments, however, we will begin by describing the assumptions we use to model hot haloes, which are very similar to those of Fukugita and Peebles (2006). ", "conclusions": "The above arguments place separate and independent constraints on the baryons around galaxies. These arguments lead us to two general conclusions: First, the majority of the missing baryons from spiral galaxies do not lie in hot haloes around these galaxies. For the Milky Way, the combination of observed constraints suggests that the common assumption of a hot halo of missing baryons obeying an NFW density profile with a concentration of 12 can only hold 2-3\\% of the baryons missing from the Galaxy. The emission measure of theoretical haloes places an upper limit of 11-24\\% on the fraction of the missing baryons that can reside in similar hot haloes around other large galaxies. The analysis of O VII sightlines constrains hot haloes around a wider mass range of galaxies. The free electrons in these haloes should be detected through O VII absorption in the spectra of several strong background X-ray sources. This absorption has not been observed, so the column must be several times smaller than the prediction. This would occur if the halo has a radius of $\\lapprox 40$ kpc or if its mass is reduced by factor of 3-4. The former case leads to further difficulties. The density would be 0.01-0.1 cm$^{-3}$, so the halo would cool quickly and should be far too bright in X-rays. The latter case is possible, but such a rarefied halo would not account for the missing baryons from galaxies. The proposal of hot haloes obeying flattened density profiles instead of NFW profiles was also considered. Around the Milky Way, a typical flattened profile could hold $6-13\\%$ of the missing baryons without violating any of our constraints. The emission measure of hot haloes with a flattened profile around other large disk galaxies sets an upper limit on the mass of about 59\\% of the missing baryons from these galaxies. The OVII constraint is approximately unchanged, which also limits the mass and extent of flattened density profiles. Collectively, these arguments rule out the existence of a large reservoir of hot gas around large galaxies that could account for the missing baryons. The most likely scenario is the presence of a small gaseous halo around $L_*$ galaxies containing $10^9 - 10^{10} M_{\\odot}$ and extending to a radius of 50-100 kpc. This picture also satisfies pulsar DM observations, measurements of the pressure around high-velocity clouds, and X-ray surface brightness constraints. If such a limited halo exists, it is not the primary reservoir of the missing baryons from galaxies. Our second general conclusion is that galaxies do not expel their baryons primarily by galactic winds driven by supernovae or AGN activity. We infer this from the lack of correlation between missing baryon fraction in galaxies and the stellar fraction or bulge mass in the galaxy. Since the latter observables are taken to be proxies for the energy available in a galaxy from supernovae or from a supermassive black hole, the lack of correlation implies these energy sources do not explain the observed missing baryon fractions. Additionally, there are galaxies where most of the baryons are still gaseous yet they have the same missing baryon fraction as their counterparts where the baryons are nearly all in stars. To explain the baryon depletion of some of these star-poor, gas-rich galaxies, the energy required is several times that expected from galactic winds powered by supernovae. It is beyond the scope of this paper to model the remaining mechanisms for resolving the second missing baryon problem, but we can comment on the possible solutions. Based on our conclusions, a straightforward way to account for the missing baryons from galaxies is if the baryons never fell into galaxies in the first place. This would require pre-heating before or just as the galaxy is forming. The most obvious source of non-gravitational energy at these times is early supernovae, whose contribution to heating would be much more effective in the epoch of protogalaxies (when potential wells are shallow) than it is today. These early supernovae could even come from the hypothetical ``Population III'' stars, which have the advantage of a top-heavy IMF, so this process would not convert much of the missing baryonic mass into low-mass stars. If the missing baryons never accreted onto protogalaxies because of early supernova heating, there are several observable consequences in addition to resolution of the second missing baryon problem. These early supernovae pre-enrich gas in addition to pre-heating it, and this helps to explain phenomena like the G-dwarf problem and the \"floor\" in iron abundance noted in section 4. For more discussion of this possibility, see Benson and Madau (2003). Due to the top-heavy IMF, they would have short lifetimes, so few of these early unbound stars would survive today, and the strict limits on intracluster light (e.g Krick and Bernstein 2007) would not be violated. A top-heavy IMF would also increase the efficiency per unit mass of metal production, which has further implications for the metallicity in galaxy clusters, as discussed in Bregman, Anderson, and Dai (2010). There are also observable constraints on this process. For example, if supernova heating is to prevent the baryons from accreting onto galaxies, the supernovae must occur before the galaxies form, i.e. $z \\gapprox 6$. There is also an important period at $z \\sim 8$ when Compton cooling can efficiently cool the baryons (Cen 2003). Depending on the details of the supernovae and the cosmology, the bulk of the early star formation and baryon expulsion might occur either before or after this period. More detailed observation and theoretical work can further constrain these possibilities. The remarkably tight correlation for galaxies between dark matter halo mass and baryon fraction seems to offer a powerful clue for understanding galaxy formation. The correlation persists over variations in star/gas mass ratio The mechanism and details of this connection remain unclear, however, and still need to be understood. Further observations could also help to constrain more tightly the mass and extent of hot haloes around galaxies. Direct imaging around large galaxies in soft X-rays could strengthen the surface brightness constraints. Measurements of extragalactic pulsars or additional OVII lines could also provide useful constraints on the electron column density around galaxies. Detailed examination of simulations could also be helpful, predicting, for example, the fraction of dark matter in the Universe that falls into galactic haloes for comparison to the behavior of the baryons. Finally, more modeling of galaxy formation could greatly assist in producing testable predictions that would differentiate between various solutions to this second missing baryon problem." }, "1003/1003.3790_arXiv.txt": { "abstract": "{The AKARI FU-HYU mission program carried out mid-infrared imaging of several well studied Spitzer fields preferentially selecting fields already rich in multi-wavelength data from radio to X-ray wavelengths filling in the wavelength desert between the Spitzer IRAC and MIPS bands. We present the initial results for the FU-HYU survey in the GOODS-N field. We utilize the supreme multiwavelength coverage in the GOODS-N field to produce a multiwavelength catalogue from infrared to ultraviolet wavelengths, containing more than 4393 sources, including photometric redshifts. Using the FU-HYU catalogue we present colour-colour diagrams that map the passage of PAH features through our observation bands. We find that the longer mid-infrared bands from AKARI (IRC-L18W 18 micron band) and Spitzer (MIPS24 24 micron band) provide an accurate measure of the total MIR emission of the sources and therefore their probable total mid-infrared luminosity. We also find that colours incorporating the AKARI IRC-S11 11 micron band produce a bimodal distribution where an excess at 11 microns preferentially selects moderate redshift star-forming galaxies. These powerful colour-colour diagnostics are further used as tools to extract anomalous colour populations, in particular a population of Silicate Break galaxies from the GOODS-N field showing that dusty starbursts can be selected of specific redshift ranges (z=1.2 -- 1.6) by mid-infrared drop-out techniques. The FU-HYU catalogue will be made publically available to the astronomical community. } ", "introduction": "\\label{sec:introduction} Studies with the Infrared Space Observatory ({\\it ISO}) of the Hubble deep fields, North and South (HDF-N, HDF-S) have revealed star formation rates at least comparable to or higher than those of the optical/UV studies (\\cite{mann02}) . At submillimetre wavelengths, surveys in the same HDF-N field with SCUBA on the JCMT has revealed a large ($>$3000 deg$^{-2}$ at $S_{850} > $2 mJy) population of strongly evolving sources with bolometric luminosities $>10^{12}L_{\\sun}$ and star formation rates of $\\sim$300--1000M$_{\\sun}$yr$^{-1}$ with a median redshift of 2.4 (\\cite{chapman05}). The overwhelming conclusion is that the star formation rate at z$\\sim$1--2 requires significant evolution in the IR galaxy population from the current epoch. Deep observations with the {\\it Spitzer} Space Telescope have confirmed this strong evolution in the galaxy population out to the intermediate redshift range probed by {\\it ISO} (z$\\sim$0.3--1) and furthermore provided insight into the higher redshift Universe in the so called redshift desert z$\\sim$1-3. To connect the local and intermediate redshift Universe to the higher and high z Universe observed by {\\it Spitzer} and SCUBA, comprehensive multiwavelength imaging and spectroscopy is required throughout the extragalactic population. The Fields surveyed by the {\\it Spitzer} space telescope are some of the most richest fields in multiwavelength data in the entire sky. The SWIRE wide area survey (50 square degrees over 6 fields, \\cite{lonsdale04}), the GOODS survey (300 square arcmins over 2 fields, \\cite{giavalisco04}), the other GTO fields and the Spitzer First Look Survey (FLS, e.g. \\cite{frayer06}) as well as having coverage over the 3.5-160$\\mu$m wavelength range also include a wealth of ancillary data at other wavelengths from radio to X-rays. However, Spitzer has limited imaging capability between 8$\\mu$m$<\\lambda<$24$\\mu$m and no spectroscopic capability shortward of 5$\\mu$m leaving a conspicuous gap in these ranges in both wavelength and redshift. The {\\it AKARI} satellite, (formerly known as {\\it ASTRO-F}, \\cite{murakami07}) is the first Japanese space mission dedicated to infrared astrophysics and was launched on board JAXA's M-V Launch Vehicle No. 8 (M-V-8) at 6:28 a.m. on February 22, 2006 Japan Standard Time, JST (February 21st, 9.28 p.m. UT) from the Uchinoura Space Center (USC) on the southern tip of Japan. The satellite is in a Sun-synchronous polar orbit at an altitude of 700 km and a period of 100 minutes. {\\it AKARI} has a 68.5 cm cooled telescope with two focal plane instruments, namely the Far-Infrared Surveyor (FIS) and the Infrared Camera (IRC). The FIS has two 2-dimensional detector arrays and observes in four far-infrared bands between 50 and 180 $\\mu$m (\\cite{kawada07}). The IRC consists of three cameras covering 1.7-- 26 $\\mu$m in 9 bands (see Table \\ref{table:irc}) with fields of view of approximately 10$\\arcmin$ $\\times$ 10$\\arcmin$ (\\cite{onaka07}). Both instruments have low- to moderate-resolution spectroscopic capability. \\begin{table} \\caption{Photometric Filter band definitions for the three cameras of the IRC.} \\begin{center} \\begin{tabular}{cccccc} \\hline\\hline Camera & Array Area$^*$ & Pixel Scale & Band & $\\lambda_{\\rm ref}$ & $\\Delta \\lambda$ \\\\ & (pixels) & (arcsec) & & ($\\mu$m) & ($\\mu$m) \\\\ \\hline NIR & 319$\\times$412 & 1.46$\\times$1.46 & N2 & 2.4 & 1.9--2.8 \\\\ & & & N3 & 3.2 & 2.7--3.8 \\\\ & & & N4 & 4.1 & 3.6--5.3 \\\\ MIR-S & 233$\\times$256& 2.34$\\times$2.34 & S7 & 7.0 & 5.9--8.4 \\\\ & & & S9W & 9.0 & 6.7--11.6 \\\\ & & & S11 & 11.0 & 8.5--13.1 \\\\ MIR-L & 246$\\times239$ & 2.51$\\times$2.39 & L15 & 15.0 & 12.6--19.4 \\\\ & & & L18W& 18.0 & 13.9--25.6 \\\\ & & & L24 & 24.0 & 20.3--26.5 \\\\ \\hline \\multicolumn{6}{l}{* This is the effective imaging area of the array}\\\\ \\end{tabular} \\end{center} \\label{table:irc} \\end{table} The {\\it AKARI} mission is primarily a survey mission with an All-Sky Survey in the far-infrared and large mid-infrared legacy survey programs being carried out in the region of the North Ecliptic Pole (\\cite{matsuhara06}) and the Large Magellanic Cloud. However in addition to these Large Surveys (LS), a campaign of guaranteed time observations, referred to as Mission Programs (MP) was also undertaken. We report the initial results from the {\\it AKARI}-FU-HYU Mission Program (Follow-Up Hayai-Yasui-Umai). The FU-HYU MP strategically targeted well-studied fields to maximise the legacy value of the {\\it AKARI} data and in particular to target {\\it Spitzer} fields in the above mentioned conspicuous wavelength \"gap\", at wavelengths of 11, 15 \\& 18$\\mu$m. The FU-HYU observations fill in vital gaps in the wavelength coverage (see Figure \\ref{fig:bandcoverage}) and provide invaluable insight into the connection between {\\it ISO} \\& {\\it Spitzer} populations and linking the far-infrared Universe to the high redshift sub-mm Universe as observed by the SCUBA instrument on the James Clerk Maxell Telescope. The FU-HYU program imaged three major fields (GOODS-N, Lockman Hole, ELAIS-N1, see Section \\ref{sec:FU-HYU}) and in this work we report on the initial data reduction and results from the GOODS-N observations. Observation in the Lockman Hole will be reported by \\cite{serjeant09}. Unfortunately, the FU-HYU observations in the GOODS-N field were not dithered during the observation operation resulting in many hot and bad pixels leaving {\\it holes} in the final co-added images. Moreover the distortion correction in the standard IRC pipeline toolkit blurs the hot pixels compounding the problem. In Section ~\\ref{sec:Reduction} we describe in detail our extensive additional processing utilizing intrinsic {\\it jitter} in the spacecraft observations as an effective {\\it dithering} and interrupting the standard IRC pipeline toolkit before the distortion correction stage to carry out independent re-gridding of the pixels onto a finer mesh, co-adding and re-binning, resulting in final dithered images without the hot pixels and line artifacts seen in the original processed data. The cross associations with other data sets in the GOOD-N field, most notably the {\\it Spitzer} data and the construction of the FU-HYU catalogue is described in Section \\ref{sec:Association}. The results and conclusions are given in Sections \\ref{sec:results} \\& \\ref{sec:summary} respectively. \\begin{figure} \\centering \\centerline{ \\psfig{ figure=13383FG01.eps,width=7.5cm} } \\caption{The comprehensive coverage of the {\\it AKARI} and {\\it Spitzer} bands in the FU-HYU fields. The combination of the {\\it AKARI} and {\\it Spitzer} bands completely cover the features in the mid-infrared spectra of star-forming galaxies (STFG, shown in the figure at three different redshifts) locally, out to redshifts of $\\sim$2. \\label{fig:bandcoverage}} \\end{figure} ", "conclusions": "\\label{sec:summary} We have presented the initial results from the {\\it AKARI} FU-HYU Mission Program in the GOODS-N field describing the data reduction process including the additional processing steps required to analyze the FU-HYU data due to the omission of dithering cycles when the observations were originally taken. Combining all the available data in the GOODS-N field, a final FU-HYU catalogue has been produced containing more than 4393 sources with almost 200 sources detected in the {\\it AKARI} bands. Using the combination of the {\\it AKARI} and {\\it Spitzer} multi-wavelength coverage a total mid-infrared flux has been defined to be representative of the total mid-infrared emission of the sources and a measure of the total luminosity of the galaxies (c.f. \\cite{elbaz02}). This total MIR flux is tightly correlated with the longer wavelength mid-infrared (IRC L18W, MIPS 24) bands but less correlated with the shorter IRAC and IRC bands. The implication is that the {\\it Spitzer} MIPS24 or {\\it AKARI} L18W band are representative of the total mid-infrared luminosity of these galaxies in agreement with the results of \\cite{bavouzet08} who have indicated that the 24$\\mu$m flux is a good indicator of the mid and bolometric infrared luminosity of star forming galaxies. The mid-infrared colours have been used to track the passage of the PAH emission features through the observation bands and to segregate the star-forming population at z$\\sim$1 which have higher total MIR / single band colours than the more quiescent population. In particular we have shown that an excess in emission in the {\\it AKARI} S11 band is indicative of a moderate redshift population. Using the {\\it AKARI} IRC L18W to {\\it Spitzer} MIPS 24 band colour as an example diagnostic, we have shown that it is possible to segregate specific populations in the colour colour plane, thus we have used the \"{\\it Silicate-Break}\" technique to extract extinct, dusty galaxies from our FU-HYU sample. This population is sensitive to the passage of the 9.7$\\mu$m absorption feature through the MIPS 24 band in the redshift range $\\sim$1.2 -- 1.6 and this technique has been successful in identifying 8 possible candidates in the GOODS-N field. Models fits of spectra using a photometric redshift code have indeed confirmed that the sources lie in the redshift range expected for Silicate-Break galaxies confirming that the silicate break selection method can provide a powerful means to detect dusty ULIRGs at moderate redshift." }, "1003/1003.4453_arXiv.txt": { "abstract": "The polar condensation/sublimation of $\\mathrm{CO_{2}}$, that involve about one fourth of the atmosphere mass, is the major Martian climatic cycle. Early observations in visible and thermal infrared have shown that the sublimation of the Seasonal South Polar Cap (SSPC) is not symmetric around the geographic South Pole. Here we use observations by OMEGA/Mars Express in the near-infrared to detect unambiguously the presence of $\\mathrm{CO_{2}}$ at the surface, and to estimate albedo. Second, we estimate the sublimation of $\\mathrm{CO_{2}}$ released in the atmosphere and show that there is a two-step process. From Ls=180\\textdegree{} to 220\\textdegree{}, the sublimation is nearly symmetric with a slight advantage for the cryptic region. After Ls=220\\textdegree{} the anti-cryptic region sublimation is stronger. Those two phases are not balanced such that there is $22\\%\\pm9$ more mass the anti-cryptic region, arguing for more snow precipitation. We compare those results with the MOLA height measurements. Finally we discuss implications for the Martian atmosphere about general circulation and gas tracers, e.g. Ar. ", "introduction": "The \\textquotedblleft{}cryptic region\\textquotedblright{} is a dark region covered by ice in the South Polar Region of Mars (\\citet{Kieffer_SouthRecessionTes_JGR2000,Titus_MarsPolarProcess_Book2008}). It appears to be a region where most of the {}``spiders'' features are located (\\citet{Piqueux_SublimationSpiderCryptic_JGR2003,Kieffer_ClodJets_JGR2007}). In order to simplify, we define in this article the cryptic region from longitude 50\\textdegree{}E to 230\\textdegree{}E (through longitude 90\\textdegree{}E), and the anti-cryptic region from longitude 130\\textdegree{}W to 50\\textdegree{}E (the complementary sector, passing through longitude 0\\textdegree{}E) (see fig. \\ref{fig:DefCrypticAnti-cryptic}). During the phase of $\\mathrm{CO_{2}}$ frost accumulation, the Seasonal South Polar Cap (SSPC) is mainly formed by direct condensation, but some snow events can occur (\\citealp{Forget_CO2snow_Icarus1998}). GCM studies show that a topographic forcing by the Hellas basin creates an asymmetry in the mode of deposition. Precipitation events are more frequent for the anti-cryptic sector than for the cryptic sector (\\citealp{Colaprete_AlbedoSouthPole_Nature_2005,Giuranna_PfsCondensingSSPC_Icarus2008}). This suggests that the texture should be more granular (smaller grain sizes) for the former compared to the latter. This texture difference produces a relatively higher albedo in the anti-cryptic region and also possibly a higher accumulated mass. The direct measurement of the sublimating $\\mathrm{CO_{2}}$ mass on Mars has been done only recently by three different techniques: gravity (\\citet{Smith_SaisonalSnowDepth_Science_2001,Karatekin_globalCO2_JGR2006}), neutron flux (\\citet{Litvak_PolarRegionsHEND-Odyssey_Icarus_2006}) and gamma ray flux (\\citet{Kelly_CO2-GRS_JGR2006}). But the time and space resolutions of these methods prohibit any conclusion at the regional scale. At the present time, only direct measurements of the seasonal variations in altitude by MOLA are able to estimate the local $\\mathrm{CO_{2}}$ mass sublimation (\\citet{Smith_SaisonalSnowDepth_Science_2001,Aharonson_Co2sheetMOLA_jgr2004,Jian_MOLAcryptic_ASR2009}). But this impressive technique is limited by low signal to noise ratio because seasonal topographic variations are typically less that one meter! To enhance the climatic signal, \\citet{Aharonson_Co2sheetMOLA_jgr2004} use a Fourier transform and apply a filter on the MOLA time variation keeping only the annual period. Results show that the amplitude of this annual Fourier cycle is asymmetric, with a maximum in the anti-cryptic region (see fig. 4 (a) in \\citet{Aharonson_Co2sheetMOLA_jgr2004}).", "conclusions": "We estimate the $\\mathrm{CO_{2}}$ mass balance of the SSPC using the D-frost model and estimate the total sublimated mass to be around $(5.6\\pm1,3)\\times10^{15}$ kg which is compatible with gravity measurements (\\citet{Karatekin_globalCO2_JGR2006}), with measurements by gamma ray and neutron spectroscopy (\\citet{Litvak_PolarRegionsHEND-Odyssey_Icarus_2006,Kelly_CO2-GRS_JGR2006}) and with previous GCM studies (\\citet{Forget_MarsGCM_JGR1999,Kelly_CO2-GRS_JGR2006}). This agreement validates our current approach. The effect of the Ar content, in the case of a well mixed atmosphere is below 1\\% on the total sublimated mass. A more realistic case of denser Ar enriched air in the bottom part of the atmosphere should be studied in the future. We show that the SSPC is not symmetric in mass around the geographic pole, i.e.: the cryptic and anti-cryptic regions have an accumulated mass with a difference in order of $22\\pm9$ \\%. in favour of the anti-cryptic region. Previous studies suggested that in addition to direct condensation, during the southern fall and winter seasons, snow precipitation is more intense in the anti-cryptic region than in the cryptic one (\\citet{Colaprete_CO2DustStorm_JGR2002,Giuranna_PfsCondensingSSPC_Icarus2008}). Our results show that, the actual mass deposited (possibly as snow) should be in order of 20 \\%. This conclusion is still valid if the $\\mathrm{CO_{2}}$ emissivity has regional and/or temporal variations within the range 0.9 to 1 during the spring and summer, as observed by TES (\\citet{Eluszkiewicz_MicrophysicsRadiativeIces_Icarus2003}) ans supported by fast $\\mathrm{CO_{2}}$ metamorphism (\\citet{Cornwall_ColdSpot_JGR2010inpress}). The previous estimation on small regions of interests (ROI) by \\citet{Kieffer_SouthRecessionTes_JGR2000} was 852 $\\mathrm{kg.m}^{-1}$ in the {}``cryptic'' ROI and 841$\\mathrm{kg.m}^{-1}$ in the {}``bright cap'' ROI. The difference is probably due to the crocus date measurement that is delayed for OMEGA in comparison to TES. In case of subpixel mixing, due to the thermal emission dependence on the temperature at power four, bare soil is dominating in the thermal infrared. This effect is not present in near infrared domain and the mixing is really linear. Further studies should be done to precisely compare visible, near infrared and thermal infrared datasets. From Ls=180\\textdegree{} to 220\\textdegree{}, we point out that the cryptic region is sublimating slightly stronger than the anti-cryptic region. In the second phase, anti-cryptic region sublimation is dominating with a maximum difference in sublimation of $1.0\\times10^{8}$ kg.$\\mathrm{s}^{-1}$. This results is compatible with the MOLA height measurements difference that reach a maximum for Ls=210\\textdegree{}-240\\textdegree{}. We use a simple box model to fit the GRS measurements of Ar and estimate that the atmosphere advection flux entering the south polar vortex is similar to the flux from the sublimating SSPC. More precise studies on atmosphere dynamics and minor species concentration should be carried out with a GCM or a local scale climate model. At present time, SSPC albedo is a constant parameter in any GCM. Future studies should be done to include the SSPC asymmetry of albedo leading to an asymmetry of sublimation in order to simulate realistic dynamics and dilution factors. Those results will have to be compared to the atmospheric observations of pressure (for instance using the method proposed by \\citet{Forget_RetrievalPressureOMEGA_JGR2007,Spiga_PressureOMEGA_JGR2007}) and minor species concentration (using gamma rays, i.e.: \\citet{Sprague_ArSouthPolar_Science2004}, optical measurements on the ground i.e: \\citet{Encrenaz_isotopH2O-CO2_Icarus2005}, or orbiters i.e.: \\citet{Melchiorri_WaterVaporOMEGA_PSS2007}). In particular, future local climate studies should explain the actually asymmetric Ar enrichment (\\citet{Sprague_ArSouthPolar_Science2004}) that could play an indirect role in the stability of the SSPC itself. If the atmosphere is locally enriched in Ar, the partial pressure of $\\mathrm{CO_{2}}$ is lower and then $\\mathrm{CO_{2}}$ ice stability is lower (\\citet{Forget_Ar_MOMA2009})." }, "1003/1003.3741.txt": { "abstract": "{}{}{}{}{} % 5 {} token are mandatory \\abstract{The various IBIS/ISGRI catalogues contain a large population of hard X-ray sources whose nature is still unknown. Even if the $>20$ keV positional uncertainty provided by ISGRI is unprecedented, it is still too large to pinpoint the counterpart at other wavelengths, which is the only secure way of obtaining a source identification. We continue the work of trying to reveal the nature of these hard X-ray sources, starting with analysis of X-ray data collected via focusing X-ray telescopes, in order to obtain arcsec accurate X-ray positions. We can then identify counterparts at infrared and optical wavelengths and try to unveil the nature of the sources. We analysed data from observations of 13 \\integral\\ sources made with the \\swift\\ satellite. The X-ray images obtained by the X-Ray Telescope instrument allowed us to find possible counterparts to the IGR sources with a positional accuracy of a few arcsec. We then browsed the online catalogues (e.g., NED, SIMBAD, 2MASS, 2MASX, USNO B1.0) to search for counterparts at other wavelengths. We also made use of the X-ray spectral parameters in trying to identify the nature of those objects. For the 13 objects, we found possible counterparts at X-ray energies and identified the IR/optical and/or UV counterparts as seen with \\swift/UVOT. We also discuss the likelihood of association of the X-ray and \\integral\\ source in each case. We confirm the previously proposed classification of IGR~J02524$-$0829 (Sey 2 AGN), J08023$-$6954 (RS CVn star), and J11457$-$1827 (Sey 1 AGN). For 7 of these sources we give the first identification of their nature: IGR J02086$-$1742, J12060+3818, J12070+2535, J13042$-$1020, and J13412+3022 are AGN, and J14488$-$5942 is a probable X-ray binary. For J03184$-$0014, although we question the association of the IGR and \\swift\\ sources, we classify the latter as an AGN. We suggest that IGR J15283$-$4443 is a Galactic source, but we cannot classify the source further. Finally, we question the association of IGR J11457$-$1827 and J23130+8608 with the X-ray sources we found, and go on to question the genuineness of the former IGR source. } ", "introduction": "The most recent version of the IBIS catalogue contains more than 700 hard X-ray sources \\citep{bird10}. While a certain number were known as (hard) X-ray emitters prior to the launch of \\integral, about half of them have been detected for the first time above 20 keV with IBIS/ISGRI \\citep{lebrun03}. In this paper we refer to these sources as `IGRs'\\footnote{An up-to-date online catalogue of all IGRs can be found at http://irfu.cea.fr/Sap/IGR-Sources/ note the new address for the site}. \\citet{arash07} has collected known parameters (e.g., the absorption column density, \\nh, the pulse period for Galactic sources with X-ray pulsations, the redshift for AGN, etc.) of all sources detected by \\integral\\ during the first four years of activity. With this they could study the parameter spaces occupied by different families of sources and therefore deduce important aspects of the physics of high-energy sources. However, many of these IGRs have still not been identified, and therefore any attempt to study, understand, and model populations of high-energy sources will be incomplete. The determination of the nature of these object is therefore extremely important if one wants to have the most complete view of the content of our Galaxy and our Universe. \\\\ \\indent In this paper, we continue our work of identifying the unknown IGRs that we started soon after the discovery of the first IGRs. A first step is to provide an $\\sim$arcsec position with soft X-ray telescopes such as {\\it {XMM-Newton}} \\citep{rodrigue03_1632, rodrigue06_1632, bodaghee06} {\\it{Chandra}} \\citep{john06_4igr, john08_igr,john09_igr}, and also \\swift\\ \\citep{rodrigue08_igr,rodrigue09_igr, rodrigue09_19294}. We then search for counterparts at a position consistent with the refined X-ray position of a given source. As in \\citet{rodrigue08_igr} and \\citet{rodrigue09_igr} (Papers 1 and 2 in the remainder of this article), we report here the analysis of \\swift\\ observations (XRT imaging and spectral analysis and UVOT imaging) of 13 IGRs that still lacked precise arcsec X-ray positions at the time of the writing of the paper. We also present the identification of IR and optical counterparts obtained from online catalogues such as SIMBAD, the United States Naval Observatory (USNO), the 2 Micron All Sky Survey point source and extended source catalogues\\footnote{http://www.ipac.caltech.edu/2mass/} \\citep[2MASS and 2MASX][]{skrutskie06}, and the NASA/IPAC Extragalactic Database (NED\\footnote{http://nedwww.ipac.caltech.edu/index.html}). Although the presence of a bright \\swift\\ source within a given \\integral\\ error circle usually renders the association between the two sources likely, there is a slight probability that the two sources are not associated, especially in the case of dim X-ray sources. This is, also, exemplified by the few cases where several \\swift\\ sources are found within the \\integral\\ error circle. Given the wide range of association probabilities from possible associations to nearly certain associations, no general statement can be given for the probability of associations. A low Galactic latitude source will have a higher chance of spurious association than a high latitude one. For all sources, we discuss the likelihood of association between the \\integral, \\swift, and counterparts at other wavelengths. Dubious cases (such as multiple possible counterparts) are discussed in more detail. \\\\ \\indent We start by introducing the \\swift\\ observations, and we briefly present the data reduction techniques in Sect.~2. We then give the results of X-ray (Sect.~3) and IR/Optical/UV (Sect.~4) candidate counterparts identification. In Sect.~5 we describe the results for each source, including the results of the X-ray spectral analysis, and discuss their possible nature. We conclude the paper by summarising the results in Sect.~6. \\begin{table} \\caption{Journal of the \\swift\\ observations analysed in this paper.} \\begin{tabular}{lllll} \\hline \\hline Source Id & Id & Date Obs & Tstart & Exposure\\\\ (IGR) & & & (UTC) & (s) \\\\ \\hline J02086$-$1742 & 00038021001 & 2009-09-10 & 14:08:56 & 4938 \\\\ J02524$-$0829 & 00036970001 & 2008-01-27 & 00:41:02 & 11117 \\\\ & 00036970002 & 2008-06-11 & 00:51:48 & 4310 \\\\ J03184$-$0014 & 00030995001 & 2007-11-07 & 00:12:58 & 9192 \\\\ & 00036969001 & 2008-02-28 & 01:48:52 & 8378 \\\\ & 00036969002 & 2008-02-29 & 19:32:52 & 4711 \\\\ & 00036969003 & 2009-06-24 & 19:25:35 & 1030 \\\\ J08023$-$6954 & 00036095001 & 2006-12-27 & 07:45:23 & 701 \\\\ & 00036095002 & 2006-12-30 & 08:04:22 & 5736 \\\\ J11457$-$1827 & 00035645001 & 2006-07-27 & 00:05:11 & 10000 \\\\ & 00035645002 & 2006-07-29 & 00:19:43 & 3384 \\\\ J12060+3818 & 00037838001 & 2008-10-13 & 00:31:19 & 5363 \\\\ J12070+2535 & 00037837001 & 2008-10-27 & 18:03:14 & 2032 \\\\ J12482-5828 & 00038349001 & 2008-12-19 & 00:43:32 & 3718 \\\\ & 00038349002 & 2009-05-08 & 01:06:36 & 1847 \\\\ J13042-1020 & 00031153001 & 2008-03-03 & 15:34:26 & 2948 \\\\ & 00031153002 & 2008-03-05 & 14:11:09 & 816 \\\\ & 00031153003 & 2008-03-05 & 14:14:18 & 4889 \\\\ & 00031153004 & 2008-03-07 & 00:21:55 & 185 \\\\ & 00031153005 & 2008-03-07 & 00:22:18 & 6621 \\\\ & 00031153006 & 2008-03-09 & 11:45:39 & 2676 \\\\ & 00031153007 & 2008-03-13 & 13:44:13 & 3327 \\\\ & 00031153008 & 2008-03-16 & 11:00:24 & 3279 \\\\ & 00031153009 & 2008-04-23 & 05:01:41 & 3246 \\\\ J13412+3022 & 00037380001 & 2008-08-24 & 00:42:39 & 2694 \\\\ & 00037835001 & 2008-08-25 & 15:17:06 & 3562 \\\\ J14488-5942 & 00039094001 & 2009-09-25 & 19:25:22 & 16413 \\\\ J15283-4443 & 00036114001 & 2007-01-06 & 07:36:11 & 5534 \\\\ J23130+8608 & 00037078002 & 2007-07-09 & 01:33:08 & 7908 \\\\ & 00037078003 & 2007-07-11 & 00:03:00 & 11191 \\\\ \\hline \\hline \\end{tabular} \\label{tab:log} \\end{table} ", "conclusions": "In this paper, we reported the \\swift\\ X-ray analysis of the field of thirteen IGRs that still lacked an arcsec accurate position. The refined X-ray positions provided by the \\swift\\ observations (Table~\\ref{tab:xray}) allowed us to pinpoint the possible IR, optical, and UV counterparts in most of the cases. We also analysed the X-ray spectra of the sources and used these results as additional arguments to confirm or refute the association of the \\swift\\ source with the \\integral\\ one. This also helped us to tentatively give a possible classification for the X-ray source. Table~\\ref{tab:results} reports the conclusions of our analysis \\begin{table}[htbp] \\caption{Summary of the possible type for each counterpart of the thirteen sources, obtained through our analysis.} \\label{tab:results} \\begin{tabular}{ll} \\hline\\hline Name & Type \\& Comment\\\\ (IGR) & \\\\ \\hline J02086$-$1742 & AGN, possibly Sey 1 \\\\ J02524$-$0829 & Sey 2 AGN at z=0.016721 \\\\ J03184$-$0014 & possible AGN, \\swift\\ and IGR associated?\\\\ J08023$-$6954 & RS CVn \\\\ J11427+0854 & ?, possible spurious IGR\\\\ J11457$-$1827 & Sey 1 AGN at z=0.0329 \\\\ J12060+3818 & QSO, possibly Sey 1, \\swift\\ and IGR associated?\\\\ J12070+2535 & AGN, possibly Sey 1, \\swift\\ and IGR associated?\\\\ J13042$-$1020\\#1 & AGN, possibly Sey 2\u00a0\\\\ J13412+3022 & Sey 2 AGN at z=0.03986 \\\\ J14488$-$5942\\#1 & probable XRB\\\\ J15283$-$4443 & Galactic source, \\correc{\\swift\\ and IGR associated?}\\\\ J23130+8608 & ? , \\swift\\ and IGR associated?\\\\ \\hline \\hline \\end{tabular} \\end{table} We can summarise our results as follows.\\\\ \\begin{itemize} \\item We identify IGR J02086$-$1742, IGR J12060+3818, IGR J12070+2535, IGR J13042$-$1020, and IGR J13412+3022 as AGN. We, however, do question the associations of IGR J12060+3818 and IGR J12070+2535 with the X-ray counteparts we found. We suggest that IGR J02086$-$1742 is a possible Sey 1 and that IGR J13042$-$1020 is a possible Sey 2. Our analysis permits us to clearly identify IGR J13412+3022 as a Sey 2. \\item We confirm the previously proposed associations of IGR J02524$-$0829, IGR J08023$-$6954, and IGR J11457$-$1827. These objects are respectively classified as a Sey 2 AGN, an RS CVn star, and a Sey 1 AGN. \\item We classify IGR J14488$-$5942 as a probable XRB, \\item Apart from classifying IGR J15283$-$4443 as a Galacric source, we cannot conclude much more about the nature of this source . \\item We provide new data for IGR J03184$-$0014. We confirm the presence of a source found in Paper 2, but we discuss its possible association with the IGR source. We provide a new SDSS identification for the counterpart, which is classified as a galaxy. The X-ray source is therefore an AGN, which leads us to tentatively associate it with the IBIS source, although we notice that the latter could be spurious. \\item We are not able to give a classification for IGR J11427+0854 and IGR J23130+8608. In both cases, the association of the X-ray and the IGR can be questioned. We further question the genuineness of the former IGR source. \\end{itemize} Caution is needed with these proposed identifications, as definitive conclusions will only come from optical/IR spectroscopy. We are, however, confident wbout the objects proposed as AGN, because they come from identifying extended counterparts within the XRT error box. Note that IGR J13412+3022 is also a known Sey 2 object." }, "1003/1003.0270_arXiv.txt": { "abstract": "{AKARI is the first Japanese astronomical satellite dedicated to infrared astronomy. One of the main purposes of AKARI is the all-sky survey performed with six infrared bands between 9 and 200\\,$\\mu$m during the period from 2006 May 6 to 2007 August 28. In this paper, we present the mid-infrared part (9\\,$\\mu$m and 18\\,$\\mu$m bands) of the survey carried out with one of the on-board instruments, the Infrared Camera (IRC). } {We present unprecedented observational results of the 9 and 18\\,$\\mu$m AKARI all-sky survey and detail the operation and data processing leading to the point source detection and measurements.} { The raw data are processed to produce small images for every scan and point sources candidates, above the 5$\\sigma$noise level per single scan, are derived. The celestial coordinates and fluxes of the events are determined statistically and the reliability of their detections is secured through multiple detections of the same source within milli-seconds, hours, and months from each other. } { The sky coverage is more than 90\\% for both bands. A total of 877,091 sources (851,189 for 9\\,$\\mu$m, 195,893 for 18\\,$\\mu$m) are confirmed and included in the current release of the point source catalogue. The detection limit for point sources is 50\\,mJy and 90\\,mJy for the 9\\,$\\mu$m and 18\\,$\\mu$m bands, respectively. The position accuracy is estimated to be better than 2$''$. Uncertainties in the in-flight absolute flux calibration are estimated to be 3\\% for the 9\\,$\\mu$m band and 4\\% for the 18\\,$\\mu$m band. The coordinates and fluxes of detected sources in this survey are also compared with those of the IRAS survey and found to be statistically consistent.} {} ", "introduction": "Unbiased and sensitive all-sky surveys at infrared wavelengths are important for the various fields of astronomy. The first extensive survey of the mid- to far- infrared sky was made by the Infrared Astronomy Satellite (IRAS) mission launched in 1983 \\citep{IRAS}. IRAS surveyed 87\\% of the sky in four photometric bands at 12, 25, 60 and 100\\,$\\mu$m and substantially pioneered the various new fields of astronomy, such as circumstellar debris disks around Vega-like stars \\citep{Vega}, and a new class of galaxies that radiate most of their energy in the infrared \\citep{LIRG}. A decade later than IRAS, the Midcourse Space Experiment \\citep[MSX;][]{MSX} surveyed the Galactic plane as well as the regions not observed by or confused in the IRAS mission with higher sensitivity and higher spatial resolution (18.3$''$) in four mid-infrared broad bands centered at 8.28, 12.13, 14.65 and 21.23\\,$\\mu$m and two narrow bands at 4.29 and 4.35\\,$\\mu$m. The MSX catalogue (version 1.2) of the Galactic plane survey contains 323,052 sources, 3 times as many as IRAS listed for the same region. In 2006, {\\it AKARI}, the first Japanese space mission dedicated to infrared astronomical observations \\citep{Murakami}, was launched and brought into a sun-synchronous polar orbit at an altitude of 700\\,km. It has two scientific instruments, the Infrared Camera \\citep[IRC;][]{IRC} for 2--26\\,$\\mu$m and the Far-Infrared Surveyor \\citep[FIS; ][]{FIS} for 50--200\\,$\\mu$m. {\\it AKARI} has a Ritchey-Chretien type cooled telescope with a primary-mirror aperture size of 685\\,mm \\citep{Tel}, which is operated at 6\\,K by liquid helium and mechanical coolers. One of the major observational objectives of {\\it AKARI} is an all-sky survey observation. The survey was executed during the life time of the cooling medium between 2006 May 8 and 2007 August 28. The 9 and 18\\,$\\mu$m bands of the IRC and the 60, 90, 140, and 160\\,$\\mu$m bands of the FIS were used for the all-sky survey. In this paper, we present the mid-infrared part of the all-sky survey performed with the IRC. The IRC was originally designed for imaging and spectroscopic observations in the pointing mode, however, the all-sky observation mode was added as an operation mode following ground tests, in which the acceptable performance of continuous survey-type observations was confirmed \\citep{ScanOpe}. The data of the IRC all-sky survey observation have been processed by a dedicated program and a point source catalog has been prepared. The content of this paper is based on the $\\beta$-1 version of the {\\it AKARI}/IRC All-Sky Survey Point Source Catalogue. The outline of the observation is presented in \\S 2. The data reduction is described in \\S 3. The quality of the catalogue is statistically evaluated in \\S 4, and a summary is given in \\S 5. ", "conclusions": "The {\\it AKARI} Mid-Infrared all-sky survey was performed with two mid-infrared broad bands centered at 9 and 18\\,$\\mu$m. More than 90\\% of the entire sky was observed in both bands. A total of 877,091 sources (851,189 for 9\\,$\\mu$m, 195,893 for 18\\,$\\mu$m) are detected and included in the present release of the point source catalogue. This {\\it AKARI} Mid-Infrared point source catalogue is scheduled for the public release in 2010 after the prioritized period for the team members \\citep{relnote}. We present the spatial distribution, flux distribution, flux accuracy, position accuracy, and completeness of the sources in the {\\it AKARI} MIR All-Sky Survey Catalogue version $\\beta$-1. The {\\it AKARI} mid-infrared survey provides a unique data-set relevant to interstellar, circumstellar-, and extra galactic astronomy. This new deep, large survey is well suited to research in the various fields of astronomy, such as search for warm debris disks, of asteroid analogue (Fujiwara et al. 2009a, 2009b), and provides valuable information for the study of planet formation and other fields. The chemical compositions of an unprecedented number of dust forming asymptotic giant branch (AGB) stars can also be investigated \\citep{Ita} making use of the characteristics of the filter bands of this survey. Such studies enable new discussions on the structure, the star formation history and the cycle of matter in our Galaxy. The mid-infrared all-sky survey also provides a chance to detect highly obscured active galactic nuclei (AGNs), which are difficult to observe in previous optical or X-ray surveys. The population of obscured AGNs is expected to contribute to the hard X-ray background, which is not competely resolved into individual sources yet \\cite{Ueda}." }, "1003/1003.2333_arXiv.txt": { "abstract": "{} {Our aim is to study the photospheric flux distribution of a twisted flux tube that emerges from the solar interior. We also report on the eruption of a new flux rope when the emerging tube rises into a pre-existing magnetic field in the corona.} {To study the evolution, we use 3D numerical simulations by solving the time-dependent and resistive MHD equations. We qualitatively compare our numerical results with MDI magnetograms of emerging flux at the solar surface.} {We find that the photospheric magnetic flux distribution consists of two regions of opposite polarities and elongated magnetic tails on the two sides of the polarity inversion line (PIL), depending on the azimuthal nature of the emerging field lines and the initial field strength of the rising tube. Their shape is progressively deformed due to plasma motions towards the PIL. Our results are in qualitative agreement with observational studies of magnetic flux emergence in active regions (ARs). Moreover, if the initial twist of the emerging tube is small, the photospheric magnetic field develops an undulating shape and does not possess tails. In all cases, we find that a new flux rope is formed above the original axis of the emerging tube that may erupt into the corona, depending on the strength of the ambient field.} {} ", "introduction": " ", "conclusions": "" }, "1003/1003.0336_arXiv.txt": { "abstract": "We use results from a constrained, cosmological MHD simulation of the Local Universe to predict radio haloes and their evolution for a volume limited set of galaxy clusters and compare to current observations. The simulated magnetic field inside the clusters is a result of turbulent amplification within them, with the magnetic seed originating from star-burst driven, galactic outflows. We evaluate three models, where we choose different normalisations for the cosmic ray proton population within clusters. Similar to our previous analysis of the Coma cluster \\citep{2010MNRAS.401...47D}, the radial profile and the morphological properties of observed radio halos can not be reproduced, even with a radially increasing energy fraction within the cosmic ray proton population. Scaling relations between X-ray luminosity and radio power can be reproduced by all models, however all models fail in the prediction of clusters with no radio emission. Also the evolutionary tracks of our largest clusters in all models fail to reproduce the observed bi-modality in radio luminosity. This provides additional evidence that the framework of hadronic, secondary models is disfavored to reproduce the large scale diffuse radio emission. of galaxy clusters. We also provide predictions for the unavoidable emission of $\\gamma$-rays from the hadronic models for the full cluster set. None of such secondary models is yet excluded by the observed limits in $\\gamma$-ray emission, emphasizing that large scale diffuse radio emission is a powerful tool to constrain the amount of cosmic ray protons in galaxy clusters. ", "introduction": "The thermal gas, that is the dominant component in the Inter-Galactic-Medium (IGM), is mixed with magnetic fields and relativistic particles, as proven by radio observations which detected Mpc-sized diffuse radio emission from the IGM, in the form of radio halos and relics \\citep[e.g.][]{2003astro.ph..1576F,2008SSRv..134...93F}. These Mpc-scale radio sources are found in a fraction of massive clusters with complex dynamics, which suggests a connection between non-thermal emission and cluster mergers \\citep[e.g.][]{2001ApJ...553L..15B,2008A&A...484..327V,2009A&A...507..661B} Cluster mergers are the most energetic events in the universe and a fraction of the energy dissipated during these mergers may be channelled into the amplification of the magnetic fields \\citep[e.g.][]{2002A&A...387..383D,2006MNRAS.366.1437S,2008Sci...320..909R} and into the acceleration of relativistic, primary, electrons and protons via shocks and turbulence \\citep[e.g.][]{1998A&A...332..395E,1999astro.ph.11439S,2001MNRAS.320..365B, 2004MNRAS.350.1174B,2001ApJ...557..560P,2003ApJ...583..695G,2003ApJ...593..599R, 2005MNRAS.357.1313C,2006MNRAS.367..113P,2007MNRAS.378..245B,2009A&A...504...33V} Relativistic protons in the IGM have long life-times and remain confined within galaxy clusters for a Hubble time \\citep[e.g.][]{1996SSRv...75..279V,1997ApJ...487..529B}. As a consequence they are expected to be the dominant non-thermal particle component. Collisions between these relativistic protons and the thermal protons in the IGM generate secondary particles that combined with the primary relativistic particles are expected to produce a complex emission spectrum from radio to $\\gamma$-rays \\citep[e.g.][]{2001APh....15..223B,2009A&A...507..661B}. Only upper limits to the $\\gamma$-ray emission from galaxy clusters have been obtained so far \\citep[][]{2003ApJ...588..155R,2006ApJ...644..148P, 2009A&A...495...27A,2009A&A...502..437A,2009arXiv0909.3267T} however the FERMI Gamma-ray telescope will shortly allow a step forward, having a chance to obtain first detections of galaxy clusters or to put stringent constraints on the energy density of the relativistic protons. Most importantly, in a few years the Low Frequency Array (LOFAR) and the Long Wavelength Array (LWA) will observe galaxy clusters at low radio frequencies with the potential to discover the bulk of the cluster-scale synchrotron emission in the Universe \\citep[e.g.][]{2002A&A...396...83E,2009arXiv0910.2025C,2006MNRAS.369.1577C}. The emerging theoretical picture is very complex and modern numerical simulations provide an efficient way to obtain detailed models of non thermal emission from galaxy clusters to compare with present and future observations. Advances in this respect have been recently obtained by including aspects of cosmic-ray physics into cosmological Lagrangian simulations mostly focussing on the acceleration of relativistic particles at shocks and on the relative production of secondary electrons \\citep[e.g.][]{2008MNRAS.385.1211P}. In this work we investigate the non-thermal emission from secondary particles in galaxy clusters extracted from Lagrangian cosmological simulations and, for the first time, we report on an adequate comparison between our expectations and observations. ", "conclusions": "We use a constrained, cosmological MHD SPH simulation with a semianalytic model for galactic magnetic outflows, to obtain a sample of 16 galaxy clusters with thermal properties similar to clusters in the Local Universe. Further we assume 3 different models for secondary cosmic rays motivated by simulations \\citep{2007MNRAS.378..385P} using the proper high energy approximation for the pion cross-section. In the first model we assume a constant CRp normalisation relative to the thermal density and the simulated magnetic field. In the second model we keep the magnetic field and introduce a radius dependent CRp normalisation infered from non-radiative simulations by \\citet{2007MNRAS.378..385P}. The third model uses the same CRs and flatten the simulated magnetic field to be $B \\propto \\sqrt{\\rho}$.\\\\ \\par Although our simulations do not include a (internally) self-consistent treatment of CRp as done in other simulations \\citep[e.g.][]{2007MNRAS.378..385P}, contrary to previous work they allow to properly simulate the properties of the magnetic field in the ICM which is important for modeling the cluster-scale synchrotron emission. \\par For the first time, we carry out a detailed comparison between the observed properties of giant radio haloes and those of simulated halos according to secondary models and under different assumptions for the spatial distribution of CRp. \\par In an earlier paper we presented a detailed comparison of the simulated Coma cluster \\citep{2010MNRAS.401...47D}. In this work we focus on global sample properties and compare with recent observations. \\\\ In particular, as a first step, we show that : \\begin{itemize} \\item The radial profiles of Faraday rotation of the median of our cluster sample is in line with that obtained from a number of observations of different clusters. This confirms that the properties of the clusters's magnetic fields in our simulations are rather similar to the observed ones. \\item The normalised radial profiles of the radio emission at 1.4 GHz of our simulated hadronic-halos show a deficit at radii $\\geq 0.1 r_{\\mathrm{vir}}$ with respect to the synchrotron profiles observed for a sample of well studied radio halos. This is in line with previous claims based on semi-analytic calculations in the context of the hadronic model \\citep{2004MNRAS.350.1174B}. In addition our results show that, even by assuming a flat profile for both the magnetic field and CRp spatial distributions (Model 3), secondary electrons may account for less than about 10-15 \\% at radii 0.15 - 0.3 $r_{\\mathrm{vir}}$. \\item A Point to Point comparison of radio vs. X-ray emission, obtained for the 4 largest clusters of our sample, confirms the results obtained for the profiles showing that the radial distribution of radio emission is too steep. All three models do not fit the observations over the whole range. Furthermore an excess in radio emission of the innermost patches suggests that haloes from the simulation are to centrally peaked. \\end{itemize} As a second step we compare scaling relations obtained for the hadronic-halos in our simulated cluster sample with those given in \\citep{2007MNRAS.378.1565C} that are obtained from a sample of observed radio halos. We find that : \\begin{itemize} \\item The geometrical correlation between the radius of radio halos and the cluster-mass contained within this radius is well reproduced by all three models. Due to the expected self-similarity of cluster in thermal properties this result implies that -at least- the simulated and observed clusters share similar physical properties. \\item A quasi self-similar behaviour is found for the non-thermal properties of our simulated clusters. In particular the radius of our hadronic halos is found to scale (approximately) with the virial radius of the simulated clusters. This is contrary to observations that found a steeper correlation between halo-radius and virial radius of the hosting clusters and implies that our simulated halos are systematically smaller than the observed ones. \\item A correlation between the monochromatic luminosity of our hadronic halos and the X-ray luminosity of the simulated hosting-clusters is found. As soon as the population of our hadronic halos is normalised, by scaling the radio luminosity of the simulated Coma halo with that of the observed one, the correlation is similar to that observed for radio halos. However, since at this point all the simulated clusters show radio emission at the level of the observed halos, we find that the cluster radio {\\it bi-modality}, observed for X-ray selected clusters, cannot be reproduced. A radio {\\it bi-modality} in the radio -- X-ray diagram would require a fast ($<$ Gyr) evolution of the radio luminosity in connection with cluster mergers, on the other hand, we find that the time evolution of our simulated massive-clusters in this diagram happens on cosmological, long, time-scales. Finally, we show that once the radio -- X-ray correlation is approached at low X-ray luminosities, clusters follow the correlation closely due to the saturation of the magnetic field. \\end{itemize} As a final point we calculate the $\\gamma$-ray emission from our simulated clusters once the radio luminosity of the simulated Coma halo is anchored to that of the observed one (essentially by scaling the number density of CRp in simulated clusters). We find that : \\begin{itemize} \\item The $\\gamma$-ray emission expected from our simulated clusters is well below the sensitivity of present Cherenkov Arrays, e.g. the VERITAS experiment, for all the adopted models for hadronic halos. The $\\gamma$-ray fluxes at $> 100$ MeV expected from our simulated clusters would allow for a marginal detection by the FERMI telescope in next years, at least by assuming Models 2 and 3. \\item The integrated $\\gamma$-ray flux from our simulated clusters is expected to scale with their radio emission, although with rather large scatter. On the other hand, a tight correlation is found between $\\gamma$-ray and X-ray fluxes, in which case (due to the scalings between thermal and non-thermal CRp in Models 1--3) the correlation is essentially driven by the density of the thermal gas in our simulated clusters. \\end{itemize} Considering all these results, as well as the outcome from our earlier work on the Coma cluster \\citep{2010MNRAS.401...47D}, we conclude that hadronic models alone are not able to explain the observed properties of giant radio haloes, in terms of their radial extension, observed clusters's radio {\\it bi-modality}, scaling relations and spectral properties. Therefore we conclude that radio emission observed in galaxy clusters in form of giant radio halos are a powerful tool to infer the amount of cosmic ray protons within galaxy cluster, confirming previous results that the energy content of cosmic ray protons in clusters can not exceed percent level \\citep[e.g.][]{2007ApJ...670L...5B,2008MNRAS.388.1062C}. The problem of halo's extension could be alleviated in {\\it extended} hadronic models, where the contribution from primary (shock accelerated) electrons at the cluster outskirts is combined in the simulations with that from secondary electrons (\\citet{2008MNRAS.385.1211P}). However, as shown in Sect. \\ref{rradprof}, the problem of the halo profiles arises at distances 0.1 -- 0.3 $R_\\mathrm{vir}$ from cluster centers where, based on the same simulations, the contribution to the diffuse synchrotron emission from these shock accelerated electrons is not yet dominant. Further we would like to note that the detailed morphology of the radio emission caused by electrons injected at shocks \\citep[e.g.][]{1999ApJ...518..603R} \\footnote{These are base on Eularian simulation, which are most appropriate to describe the morphology of shocks outside the cluster core.} will be significantly different from the morphology of giant radio halos and therefore is in general unlikely able to mimic large scale, cluster centric radio emission. Thus additional physical processes are required to explain observations. The problem of the radio {\\it bi-modality} in massive clusters can be alleviated by assuming that MHD turbulence (and the rms magnetic field) decays as soon as clusters approach a relaxed state after a major merger \\citep{2009A&A...507..661B,2009JCAP...09..024K}. In this case it might also be thought that CRp in these relaxed clusters would undergo less scattering on magnetic field irregularities escaping from the cluster volume and reducing the source of secondary electrons. All these processes are not properly included in our MHD cosmological simulations. However, \\citet{2009A&A...507..661B} have shown that the fast decay of the rms field in galaxy clusters, that is necessary to explain the observed {\\it bi-modality}, would imply a {\\it rather unphysical} situation where the magnetic field power spectrum peaks at the smaller scales, and leads to the consequence that an {\\it extremely} large flux of energy in clusters goes into magnetic field amplification/dissipation. In addition, because turbulent cascade should start to decay on largest scales, which are resolved by Faraday rotation measurements, this scenario would predict a {\\it bi-modality} in rotation measures and depolarisation in largest clusters, which is not observed (Govoni et al. 2010).\\\\ The discovery of giant radio halos with {\\it very} steep spectrum ($\\alpha \\sim 1.8-2$) in merging clusters proves that {\\it inefficient} particle acceleration mechanisms are responsible for the origin of these sources and, based on simple energy arguments, disfavours hadronic models \\citep{2008Natur.455..944B}. Observations of the best studied halo, in the Coma cluster, have shown that its spectrum steepens at higher frequencies \\citep{2003A&A...397...53T}, due to the competition between energy losses and acceleration of the emitting particles, in which case the particle acceleration time-scale would be of the order of 0.1 Gyr. More recently we have shown that the steepening cannot be a result of the inverse Compton (SZ) decrement \\citep{2010MNRAS.401...47D}, confirming these previous finding and supporting a scenario of {\\it inefficient} particle acceleration mechanisms at the origin of the halo. \\noindent Particle acceleration due to micro-turbulence in merging galaxy clusters has been proposed for the origin of radio halos \\citep[e.g.][]{2001MNRAS.320..365B,2001ApJ...557..560P}. In this case {\\it gentle} particle acceleration mechanisms would generate radio halos in connection with (massive) cluster mergers, while the radio emission would decay as soon as clusters approach a relaxed state, due to dissipation of (at least a fraction of) this turbulence and the fast electron radiative cooling. It will therefore be necessary to consider these processes in future simulations used to study radio halos. That includes an estimation of the locally merger injected turbulence as well as a more detailed description of CR electron spectra." }, "1003/1003.4103_arXiv.txt": { "abstract": "Turbulent magnetic fields fill most of the volume of the solar atmosphere. However, their spatial and temporal variations are still unknown. Since 2007, during the current solar minimum, we are periodically monitoring several wavelength regions in the solar spectrum to search for variations of the turbulent magnetic field in the quiet Sun. These fields, which are below the resolution limit, can be detected via the Hanle effect which influences the scattering polarization signatures ($Q/I$) in the presence of magnetic fields. We present a description of our program and first results showing that such a synoptic program is complementary to the daily SOHO magnetograms for monitoring small-scale magnetic fields. ", "introduction": "Spectropolarimetry is a powerful tool for diagnosing magnetic properties of the Sun. Circular polarization, characterized by Stokes $V/I$, is produced by line-of-sight magnetic fields and is a basis of solar magnetograms. Linear polarization, characterized by Stokes $Q/I$ and $U/I$, is produced by transverse magnetic fields or through coherent scattering. A known disadvantage of magnetograms is their dependence on spatial resolution. For instance, let us imagine a resolution element that contains a line-of-sight magnetic field pointing towards us and another of equal strength pointing away from us. Their respective Stokes $V/I$ signals will be of opposite sign and, if they are at the same Doppler shift, their sum is zero. Hence, we observe a grey part in the magnetogram even though there may be strong fields below our resolution limit. This cancellation effect limits the use of magnetograms for small-scale, turbulent fields, and another technique has to be used for their detection. \\begin{figure*}[!ht] \\begin{center} \\plotone{kleint_fig1.eps} \\end{center} \\caption{ {\\it Left}: Scattering geometry showing an equal amount of photons being scattered at disk center from every direction and thus no resulting polarization. At the limb, however, a radiation anisotropy results in a net polarization parallel to the nearest solar limb. {\\it Right}: For our synoptic program, the spectrograph slit is positioned at five different heliographic angles, always at an equal distance to the solar limb in order to search for spatial variations of the turbulent fields (drawing not to scale).} \\label{scatt} \\end{figure*} Coherent scattering processes on the Sun produce a linearly polarized spectrum, like scattering in the Earth's atmosphere produces polarization of the blue sky. The principle can be seen in Fig.~\\ref{scatt}. Scattering polarization only arises if the spatial symmetry is broken by an anisotropic incident radiation field (the arrows). At the disk center, an equal amount of photons are scattered from every direction relative to the line-of-sight, and with no preferred direction there is no resulting polarization. However, at the solar limb the situation is different. The anisotropy introduces a preferred direction and a net linear polarization perpendicular to this direction, i.e., parallel to the next solar limb. On the Sun, the anisotropy is caused by the temperature gradient, which manifests itself in the limb darkening. The net polarization is produced in the continuum as well as in atomic and molecular lines which can depolarize the continuum or add on top of it. The linearly polarized $Q/I$ spectrum, with $Q/I$ being defined to be parallel to the next solar limb, is called the second solar spectrum \\citep{ivanov91, stenflokeller97}. It does not resemble Stokes $I$ and cannot be easily deduced. Magnetic fields modify scattering polarization via the Hanle effect. Depending on the properties of atomic or molecular transitions, magnetic fields in the range from a fraction of Gauss to kilo Gauss can produce visible depolarization in $Q/I$ or rotate the plane of polarization and induce a signal in $U/I$. Independently of their spatial orientation, magnetic fields modify the $Q/I$ spectrum in the same way. Thus, in contrast to the Zeeman effect, the absence of cancellation enables us to detect fields below the resolution limit. Our synoptic program monitors turbulent magnetic fields, their spatial distribution and variation with the solar cycle since 2007, starting before the appearance of the first sunspot of the cycle 24. This is the first systematic study of turbulent magnetic fields which can lead to a better understanding of small-scale fields governing most of the solar photosphere. ", "conclusions": "The turbulent magnetic field as measured with the C$_2$ lines has not shown any variation during this current solar minimum. However, the average level of the polarization amplitudes is different from the single measurement obtained at the solar maximum. The chromospheric Cr~{\\small I} lines at 5206 \\AA \\ show more variations, independent of position angle. More measurements are needed for their statistical analysis, and a model has to be developed for a quantitative interpretation. Our synoptic program represents a complementary way of measuring small-scale fields in addition to, e.g., polar field measurements. It is important to know how the turbulent field varies with time, also for different activity minima and maxima." }, "1003/1003.3179_arXiv.txt": { "abstract": "{We study the structure of neutron stars in perturbative $f(R)$ gravity models with realistic equations of state. We obtain mass--radius relations in a gravity model of the form $f(R)=R+\\alpha R^2$. We find that deviations from the results of general relativity, comparable to the variations due to using different equations of state (EoS'), are induced for $|\\alpha| \\sim 10^{9}$ cm$^{2}$. Some of the soft EoS' that are excluded within the framework of general relativity can be reconciled with the 2 solar mass neutron star recently observed for certain values of $\\alpha$ within this range. For some of the EoS' we find that a new solution branch, which allows highly massive neutron stars, exists for values of $\\alpha$ greater than a few $10^9$ cm$^2$. We find constraints on $\\alpha$ for a variety of EoS' using the recent observational constraints on the mass--radius relation. These are all 5 orders of magnitude smaller than the recent constraint obtained via Gravity Probe B for this gravity model. The associated length scale $\\sqrt{\\alpha}\\sim 10^5$ cm is only an order of magnitude smaller than the typical radius of a neutron star, the probe used in this test. This implies that real deviations from general relativity can be even smaller.} ", "introduction": "The current accelerated expansion of the universe has been confirmed by many independent observations. The supporting evidence comes from the supernovae Ia data \\cite{Perlmutter, Riess1, Riess2}, cosmic microwave background radiation \\cite{Spergel1, Spergel2, Komatsu}, and the large scale structure of the universe \\cite{Tegmark1, Tegmark2}. Although the cosmological constant is arguably the simplest explanation and the best fit to all observational data, its theoretical value predicted by quantum field theory is many orders of magnitude greater than the value to explain the current acceleration of the universe. This problematic nature of cosmological constant has motivated an intense research for alternative explanations and the reasonable approaches in this direction can be divided into two main categories, both of them introducing new degrees of freedom \\cite{Uzan}: The first approach is to add some unknown energy-momentum component to the right hand side of Einstein's equations with an equation of state $p/\\rho \\approx -1$, dubbed \\textit{dark energy}. In the more radical second approach, the idea is to modify the left hand side of Einstein's equations, so-called \\textit{modified gravity}. Trying to explain such perplexing observations by modifying gravity rather than postulating an unknown dark energy has been an active research area in the last few years and in this paper we adopt this path. A modified theory of gravity has to explain the late time cosmology, and also be compatible with the constraints obtained from solar system and laboratory tests. However, it is not easy to construct theories of gravity with these requirements. Nevertheless, a class of theories, called $f(R)$ models \\cite{Odintsov-rev, Sotiriou-rev, deFelice-rev}, has attracted serious attention possibly because of its (deceptive) simplicity. Today there exist viable $f(R)$ models which are constructed carefully to be free of instabilities, and to pass the current solar system and laboratory tests \\cite{Nojiri1, Nojiri2, Nojiri3, Cognola, Hu-Sawicki, Appleby-Battye, Starobinsky, Miranda}. The strong gravity regime \\cite{Dimitri-rev} of these theories is another way of checking their viability. In this regime, divergences stemming from the functional form of f(R) may prevent the existence of relativistic stars in such theories \\cite{Briscese, Abdalla, Nojiri4, Bamba, Kobayashi-Maeda, Frolov, Nojiri5}, but thanks to the chameleon mechanism the possible problems jeopardizing the existence of these objects may be avoided \\cite{Tsujikawa, Upadhye-Hu}. Furthermore, there are also numerical solutions corresponding to static star configurations with a strong gravitational field \\cite{Babichev1, Babichev2} where the choice of the equation of state for the star is crucial for the existence of solutions, and therefore a polytropic equation of state is used in these works to overcome the possible problems related to the equation of state. Another approach to probe the viability of f(R) theories in the strong gravity regime is to use a method called perturbative constraints, or order reduction \\cite{Eliezer, Jaen}. The motivation behind using this technique in the present context is the thought that the reason of all the problems encountered in modifying gravity may be the outcome of considering these modifications as exact ones. The main issue with the exact modifications are the problems arising in curvature scales which are not originally aimed by these modifications. In the perturbative constraints approach, the modifications are viewed as next to leading order terms to the terms coming from Einstein's General Relativity. Treating f(R) gravity via perturbative constraints at cosmological scales is considered in \\cite{DeDeo, Cooney1} and for compact objects in \\cite{Cooney2}. In this manuscript we further examine the existence and properties of relativistic stars in the context of f(R) models via perturbative constraints. Specifically, we study a $f(R)$ model of the form $f(R)=R+\\alpha R^2$ and constrain the value of $\\alpha$ with the recent constraints on the mass-radius relation \\cite{OBG10}. Such a gravity model in the weak field limit is known to reduce to Yukawa-like potentials and has been recently constrained by binary pulsar data \\cite{naf} as $\\alpha \\lesssim 5 \\times 10^{15}$ cm$^2$. Here we find that in the strong gravity regime the constraint on perturbative parameter is $\\alpha \\lesssim 10^{10}$ cm$^2$. This value does not contrast with the value obtained in \\cite{naf} as they argue that the value of $\\alpha$ could be different at different length scales. The plan of the paper is as follows: In section II, we assume a perturbative form of $f(R)$ modified gravity model and obtain the field equations. Assuming also perturbative forms of metric functions and the hydrodynamical quantities we obtain the modified Tolman--Oppenheimer--Volkoff (TOV) equations. In section III, the modified TOV equations are solved numerically for various forms of the equation of state, the functional form of $f(R)$, and various values of the perturbation parameter $\\alpha$. Finally, in the discussion section, we comment on the results of numerical study and on the significance of the scale of the perturbation parameter $\\alpha$. ", "conclusions": "In this work we analyze the neutron star solutions with realistic EoS' in perturbative $f(R)$ gravity. Among the modified gravity theories the $f(R)$ theories are relatively simple to handle. However, even for these theories, the field equations are complicated and obtaining modified TOV equations in a standard fashion is difficult. This difficulty is mainly due to field equations being fourth order unlike in the case of general relativity, which has second order field equations. In order to resolve this situation, we adapt a perturbative approach \\cite{Cooney2} in which the extra terms in the gravity action are multiplied by a `dimensionful' parameter $\\alpha$. The extra terms with some appropriate value of $\\alpha$ are supposed to act perturbatively and modify the results obtained in the case of general relativity. We present how the perturbative $f(R)$ modifications affect the TOV equations. A drawback in using neutron stars for testing alternative theories of gravity has been the weakly constrained M-R relation. After the tight constraints obtained in \\cite{OBG10}, it seems this is no longer quite true. The result of \\cite{OBG10} and the measured mass of PSR J1614-2230 \\cite{dem10} excludes many EoS' in the framework of GR. In the $f(R)=R+\\alpha R^2$ gravity model, the value of $\\alpha$ provides a new degree of freedom and we show in this paper that some of the EoS', which are excluded within the framework of GR, can now be reconciled with the observations for certain values of $\\alpha$. This then brings the question of degeneracy between using different EoS' and modifying gravity. In the gravity model studied here, variations in M-R relation comparable to that of using different EoS are induced for $\\alpha$ being in the order of $10^9$ cm$^2$, which we specify for each EoS. An order of magnitude larger values of $\\alpha$, which is still 5 orders of magnitude smaller than the constraint obtained via Gravity Probe B \\cite{naf}, does not produce neutron stars with observed properties. Thus, we argue that $|\\alpha| \\lesssim 10^{10}$ cm$^2$ is a reasonable constraint independent of the EoS. We conclude that the presence of uncertainties in the EoS does not cloak the effect of the free parameter $\\alpha$ on the results. For some EoS' (AP4 and GS1) a new stable solution branch ($dM/d\\rho_c>0$), which does not exist in general relativity, is found. This solution branch, for larger values of $|\\alpha|$ has a larger domain and joins the conventional stable branch beyond some $\\alpha$. In this case there is no critical maximum mass to the neutron stars. One might be curious whether the perturbative approach followed in this paper holds for the range of $\\alpha$ considered. For a neutron star the typical value of the Ricci curvature is calculated to be roughly on the order of $\\sim 10^{-12}$ cm$^{-2}$. Therefore in the case of $f(R)=R+\\alpha R^2$ model one easily sees that the perturbative term is $10^{-2}$ orders of magnitude smaller than the Einstein--Hilbert term $R$, which justifies our approach. Although we consider 6 representative EoS' here, the above analysis repeated with other realistic EoS' will not alter the order of magnitude of the constraint on $\\alpha$ for the following reason: the constraints we obtained implies a length scale of $R_0^{-1/2}\\sim 10^5$ cm which is only an order of magnitude smaller than the typical radius of a neutron star, the probe used in this test. This is actually the length scale below which the gravity models used here \\emph{should} induce modifications on the M-R relation of an object of size 10 km. This implies that real deviations from general relativity should be hidden at even much smaller values of $|\\alpha|$." }, "1003/1003.3981_arXiv.txt": { "abstract": "% Thirty-four years of WSO (Wilcox Solar Observatory) and thirteen years of SOHO/MDI (Michelson Doppler Imager on the Solar and Heliospheric Observatory) magnetograms have been studied to measure the east-west inclination angle, indicating the toroidal component of the photospheric magnetic field. This analysis reveals that the large-scale toroidal component of the global magnetic field is antisymmetric around the equator and reverses direction in regions associated with flux from one solar cycle compared to the next. The toroidal field revealed the first early signs of cycle 24 at high latitudes, especially in the northern hemisphere, appearing as far back as 2003 in the WSO data and 2004 in MDI. As in previous cycles, the feature moves gradually equatorward. Cycles overlap and the pattern associated with each cycle lasts about 17 years. Even though the polar field at the current solar minimum is significantly lower than the three previous minima, the toroidal field pattern is similar. ", "introduction": "Solar dynamo models predict that the toroidal and poloidal components of the global magnetic field are regenerated from one other. The poloidal field is transformed into the toroidal field from differential rotation (the $\\Omega$-effect), and the toroidal field is twisted into the poloidal field (the $\\alpha$-effect). These alternate and repeat in a 22 year cycle [e.g. \\citeauthor{dg2008}, 2008]. This toroidal field component, previously measured by \\citeauthor{ss1994} [1994] using WSO data over the period of 1977-1992, provided evidence for an extended activity cycle of 16-18 years. \\citeauthor{ub2005} [2005] also measured the toroidal field using Mount Wilson data from 1986-2004 and verified the dynamo model for the creation and reversal of the toroidal field. Since we are facing a peculiar solar minimum, we want to investigate whether there is also some peculiar activity in the toroidal fields leading up to the minimum. Three solar cycles of WSO magnetic field data are available to us now, as well as a complete cycle observed with SOHO/MDI [\\citeauthor{sea1995}, 1995]. ", "conclusions": "From the inclination difference summary averages, we detected the first emergence of Cycle 24 some time during 2003-2004. This emergence pattern is very much like the pattern observed in the previous two cycles. The new emerging inclination angles, which are related to the toroidal fields, are consistent with those from the previous cycles. This suggests that the toroidal field is not a clear indicator of the duration of the following solar minimum or the generation of the poloidal fields at the concurrent minimum. However, if the toroidal field is generated from the preceding poloidal field, we might expect to see that the next cycle's toroidal field will be weak. Also, the inclination angle that indicates the toroidal field at the Sun's surface is only a distant reflection of the toroidal field that is present wherever the true solar dynamo is acting. Particularly in the higher resolution MDI data, the weak field regions that contribute to the analysis may also reflect the spatially averaged characteristics of a surface dynamo process whose contribution to the solar cycle dynamo is uncertain. Further work needs to be done to solve the calibration issues of MDI." }, "1003/1003.6089_arXiv.txt": { "abstract": "{We present the ground-based activities within the different working groups of the {\\it Kepler} Asteroseismic Science Consortium (KASC). The activities aim at the systematic characterization of the 5000+ KASC targets and at the collection of ground-based follow-up time-series data of selected promising {\\it Kepler} pulsators. So far, 36 different instruments at 31 telescopes on 23 different observatories in 12 countries are in use and a total of more than 530 observing nights has been awarded.} ", "introduction": "The {\\it Kepler} Asteroseismic Science Consortium, KASC\\footnote{http://astro.phys.au.dk/KASC}, u\\-ni\\-tes hundreds of asteroseismologists from institutes all over the world in different topical Working Groups, with the aim of performing seismic studies of all types of pulsating stars across the Hertzsprung-Russell diagram, based on {\\it Kepler} time-series space photometry. The ground-based observational Working Groups (GBOsWG) take care of the organisation of ground-based observations in support of the {\\it Kepler} space data. Additional ground-based multi-colour and spectral information are indispensable for a successful seismic modelling (see, e.g., Uytterhoeven et al. 2008a, 2009; Uytterhoeven 2009). The need for ground-based support data is motivated by two objectives: 1) the characterization of all {\\it Kepler} targets in terms of fundamental stellar parameters, 2) the identification of mode parameters from multi-colour and spectral time-series observations for selected pulsators. The KASC GBOsWG is making great efforts in organising and planning telescope time on various instruments around the world to meet these objectives and to ensure an optimal seismic exploitation of the {\\it Kepler} data. So far, 36 different instruments at 31 telescopes on 23 different observatories in 12 countries are involved and a total of more than 530 observing nights has been awarded. ", "conclusions": "" }, "1003/1003.2380_arXiv.txt": { "abstract": "These notes cover some of the topics associated with direct detection of dark matter at an introductory level. The general principles of dark matter search are summarized. The current status of some experiments is described, with an emphasis on bolometric and noble liquid techniques. Plots and illustrations associated to these notes may be found on transparencies presented during the lecture, on the web site of Gif school 2009~\\footnote{Ecole de Gif, Batz sur mer, 21-25 September 2009} (in French) : \\url{http://www-subatech.in2p3.fr/gif2009.html} ", "introduction": "\\label{sec:introduction} Current cosmological observations have lead to a concordance model of our Universe, in which a major role is played by a still mysterious dark matter. Dark matter drives the dynamics of galaxies and clusters, and generates the growth of large-scale structures, from initial density perturbations of order $\\sim 10^{-5}$ that are measured in the CMB anisotropies, to inhomogeneities of amplitude $\\sim$ 1 observed in the large scale structures at low redshift. In the concordance model, dark matter cannot be of baryonic nature since both primordial abundance measurements and CMB anisotropies point towards a small global baryonic density $\\Omega_b$. Dark matter, as well as dark energy, might be a manifestation of the break-down of general relativity on large scales or in the regime of low accelerations. A large litterature is now devoted to possible modifications of the laws of gravitation : in general, these modifications are equivalent to the introduction of new fields in addition to the metric tensor $g_{\\mu\\nu}$. Another approch is to assume that dark matter is constituted of a gaz of collisionless, stable and massive particles whose nature is still unknown. Indeed several arguments strongly suggest that the Standard Model of particle physics is not complete, at least for energies larger than the electroweak scale. A large number of extensions of the Standard Model provide quite naturally a possible dark matter candidate in their spectrum~\\cite{Steffen:2008qp}. In these notes we present the experimental efforts for direct detection of possible dark matter (DM) particles. All experiments for DM direct detection aim at testing the hypothesis that our own galactic halo is filled with particle dark matter. For several plausible DM hypothesis, like the gravitino in mSUGRA~\\cite{Ellis:2003dn}, the interaction rate of DM with ordinary matter is unfortunately far too small to be detected with current technologies. Two main categories of DM particles can currently be probed by direct detection experiments: \\begin{enumerate} \\item Weakly Interacting Massive Particles (WIMPs) which have cross-sections with ordinary matter driven by physics at the electroweak scale. Currently, the search for WIMPs represents by far the strongest effort in direct detection experiments. Several models of new physics at the electroweak scale provide such candidates, like the neutralino in the minimal supersymetric extension of the standard model (MSSM) or the lightest Kaluza-Klein particle (LKP) in universal extra dimension (UED) models. The case for WIMPs is particularly strong due to the so-called ``WIMP miracle\": if dark matter particles are thermal relics, then from thermodynamical considerations their current density $\\Omega_M \\sim 0.3$ implies a typical annihilation cross-section during their freeze-out $\\langle \\sigma_{a} v \\rangle \\sim 3\\times 10^{-26}$ cm$^3$/s which is characteristic of weak interactions. \\item Axions or axion-like particles are hypothetical pseudo-scalar particles. Axions were initially introduced to solve the CP violation problem in the strong interaction sector through the so-called Peccei-Quinn mechanism, but axion-like particles constitute a generic prediction of some string models. They can constitute (non-thermal) relics of the Big-Bang. The direct search for axion dark matter is possible thanks to their coupling to photons, and will be briefly mentionned in Section~\\ref{sec:axions}. \\end{enumerate} \\noindent We now concentrate mostly on the search for WIMPs, which started during the 80s after the seminal works of, eg.~\\cite{Goodman:1984dc,Primack:1988zm}. In the most studied models, typical WIMP masses may range from $\\sim 10$ GeV to 10 TeV. For kinematic reasons, direct detection experiments aim in general to observe nuclear recoils due to an elastic scattering of WIMPs on the nuclei of a target. The WIMP-nucleon cross section is poorly constrained by current measurements in frameworks like MSSM or UED . It may range typically from $10^{-6}$ to $10^{-12}$ pb~\\footnote{1 pb = 1 picobarn = $10^{-36}$ cm$^2$}. The last value corresponds to less than one interaction per ton of detector per year. All the efforts of WIMP experiments are therefore devoted to the ability to detect such a low rate of interactions. ", "conclusions": "" }, "1003/1003.2349_arXiv.txt": { "abstract": "{} {We study the emergence of a toroidal flux tube into the solar atmosphere and its interaction with a pre-existing field of an active region. We investigate the emission of jets as a result of repeated reconnection events between colliding magnetic fields.} {We perform 3D simulations by solving the time-dependent, resistive MHD equations in a highly stratified atmosphere.} {A small active region field is constructed by the emergence of a toroidal magnetic flux tube. A current structure is build up and reconnection sets in when new emerging flux comes into contact with the ambient field of the active region. The topology of the magnetic field around the current structure is drastically modified during reconnection. The modification results in a formation of new magnetic systems that eventually collide and reconnect. We find that reconnection jets are taking place in successive recurrent phases in directions perpendicular to each other, while in each phase they release magnetic energy and hot plasma into the solar atmosphere. After a series of recurrent appearance of jets, the system approaches an equilibrium where the efficiency of the reconnection is substantially reduced. We deduce that the emergence of new magnetic flux introduces a perturbation to the active region field, which in turn causes reconnection between neighboring magnetic fields and the release of the trapped energy in the form of jet-like emissions. This is the first time that self-consistent recurrency of jets in active regions is shown in a three-dimensional experiment of magnetic flux emergence.} {} ", "introduction": " ", "conclusions": "" }, "1003/1003.5796_arXiv.txt": { "abstract": "Discovery of the first planetary system by direct imaging around HR\\,8799 has made the age determination of the host star a very important task. This determination is the key to derive accurate masses of the planets and to study the dynamical stability of the system. The age of this star has been estimated using different procedures. In this work we show that some of these procedures have problems and large uncertainties, and the real age of this star is still unknown, needing more observational constraints. Therefore, we have developed a comprehensive modeling of HR\\,8799, and taking advantage of its $\\gamma$ Doradus-type pulsations, we have estimated the age of the star using asteroseismology. The accuracy in the age determination depends on the rotation velocity of the star, and therefore an accurate value of the inclination angle is required to solve the problem. Nevertheless, we find that the age estimate for this star previously published in the literature ([30,160] Myr) is unlikely, and a more accurate value might be closer to the Gyr. This determination has deep implications on the value of the mass of the objects orbiting HR\\,8799. An age around $\\approx$1 Gyr implies that these objects are brown dwarfs. ", "introduction": "The discovery of the first planetary system by direct imaging around HR\\,8799 \\citep{Marois} was an important milestone in the field of exoplanet research. Up to now, eleven planets have been discovered with this procedure, and only one possible planetary system. One of the advantages of this technique is the direct measurement of the luminosity and projected orbits of the planets, making the physical characterization of the system and the individual planets possible. \\cite{Marois} used a procedure for estimating the mass of the objects around HR\\,8799 in order to discriminate whether they were real planets or brown dwarfs (BD). This procedure can be applicable to any direct imaging detection, and it is based on the comparison of theoretical evolutionary tracks of BD and giant planets with observations in an Age - Luminosity diagram. Since luminosity is a direct observable, this technique is limited by the accuracy of the age determination and the theoretical models used for this comparison \\citep{reider}. The A5V spectral type star HR\\,8799 (V342 Peg, HD\\,218396, HIP\\,114189) has been extensively studied. \\cite{schuster} firstly reported this star as a possible SX Phoenicis object (these stars are pulsating subdwarf stars with periods larger than one day). \\cite{zerbi} classified it as one of the 12 first $\\gamma$ Doradus pulsators known. This pulsating stellar type is composed of Main Sequence (MS) stars in the lower part of the classical instability strip \\citep{tdcma}, with periods around one day. That means that their pulsating modes are asymptotic $g$-modes and they are suitable to study the deep interior of the star. \\cite{gray} obtained an optical spectrum of HR\\,8799, and assigned an spectral type of kA5 hF0 mA5 V $\\lambda$ Bootis (see that paper for the meaning of this specific nomenclature), reporting a metallicity of [M/H]$=-0.47$. The $\\lambda$ Bootis nature of the star means that it has solar surface abundances of light elements, and subsolar abundances of heavy elements, the internal metallicity of the star being unknown. They also noted that HR\\,8799 may be also a Vega-type star, characterized by a far IR excess due to a debris disk. Up to now, only three $\\lambda$ Bootis stars have been reported to be $\\gamma$ Doradus pulsators: HD\\,218427 \\citep{rodriguez06a}, HD\\,239276 \\citep{rodriguez06b}, and HR\\,8799. \\cite{Marois} estimated the age of the star using four different methods (see Section 2 for details). The conclusion of that work was that HR\\,8799 is a young MS star with an age in the range $[30, 160]$ Myr. \\cite{reider} added another element for estimating the age of this star: the infrared excess ratio, studied by \\cite{su2006} and also used by \\cite{chen}. None of these determinations are conclusive, as we explain in Section 2. In this work, a comprehensive modelling of HR\\,8799 has been done in order to estimate its age and mass. The $\\lambda$ Bootis nature of this star, and its $\\gamma$ Doradus pulsations are the bases of the determinations presented here. All the technical details are presented in an accompanying paper \\citep[Paper I]{hr8799lambda}, where, a complementary study on the $\\lambda$ Bootis nature of this star is developed. ", "conclusions": "In the present work, an analysis of the age determination of the planetary system HR\\,8799 has been done. The results found in the literature are not conclusive, and the only valid argument to estimate the age of the star is that using its radial velocity and proper motion, but it is an estatistical argument needed of additional estimations. The only valid argument used to estimate the age of the star is its radial velocity and proper motion, but it is a statistical argument needed of additional estimations. The main complementary argument is the position of the star in the HR diagram. In this work we have demostrated that this procedure does not provide accurate age estimations due to the $\\lambda$ Bootis nature of HR\\,8799. This nature hides the real internal metallicity of the star. The main consequence of this result is that the models fulfilling observations are in a range of ages [10,2337] Myr, a much broader range that one estimated by other authors of [30,160] Myr \\citep{Marois}. Only a small amount (18.1$\\%$) of models in our representative grid have ages in the range claimed in the literature. Therefore we need aditional constraints for an accurate estimation of the age and mass of the star (these two quantities are the physical characteristics of the star with larger impact in the understanding of the planetary system). We have taken advantage of the $\\gamma$ Doradus pulsations of the star to better estimate these values with the help of asteroseismology. A comprehensive asteroseismological study of this star has been developed. This study is described in detail in \\cite[Paper I]{hr8799lambda}. The main source of uncertainty of the procedure is the unknown rotation velocity of the star. We have analysed the possible results depending on the inclination angle $i$. There is a range of angles where this study is not accurate ($i=[18^\\circ,36^\\circ]$). For angles around $i=36^\\circ$, the models fulfilling all the observational constraints have masses in two separate ranges of M=[1.32, 1.33], [1.44, 1.45] $M_{\\odot}$. The age of the system is constrained in two separate ranges: [1123, 1625] Myr and [26, 430] Myr respectively. A percentage of 16.7$\\%$ of the models are in the range given in the literature, i.e., young ages. A consequence of this result is that, in the case of the youngest age range, the predicted masses of the observed planets are [5,14] $M_{\\rm{Jup}}$ for the most luminous planets, and [3,13] $M_{\\rm{Jup}}$ for the less luminous one. The oldest age range predicts masses for the three objects in the brown dwarfs domain \\citep[see Fig. 4 of][]{Marois}. In the most favourable case for the procedure used in this work, i.e. inclination angles of around $i\\approx 50^\\circ$, asteroseismology is very accurate, and the star would have an age in the range [1126, 1486] Myr, and a mass M=1.32 $M_{\\odot}$. This age range implies that the observed objects orbiting HR\\,8799 would be brown dwarfs \\citep[following Fig. 4 in][]{Marois}. The lack of an accurate determination of the inclination angle is the main source of uncertainty of the present study. This angle has not been unambiguously obtained up to now, and its value would say whether the results of this study are actually applicable, and then, it can provide a very accurate determination of the age and mass of the star. In any case, one of the main conclusions of our study is that the range of ages assigned to this star in the literature is unlikely to be the correct one. Only a stellar luminosity larger than that reported would allow young models with solar metallicity to fulfill all the observational constraints." }, "1003/1003.2808_arXiv.txt": { "abstract": "\\noindent We investigate the extent to which the uncertainties associated with the propagation of Galactic cosmic rays impact upon estimates for the \\gray flux from the mid-latitude region. We consider contributions from both standard astrophysical background (SAB) processes as well as resolved point sources. We have found that the uncertainties in the total \\gray flux from the mid-latitude region relating to propagation parameter values consistent with local B/C and \\bebe~data dominate by 1-2 orders of magnitude. These uncertainties are reduced to less than an order of magnitude when the normalisations of the SAB spectral components are fitted to the corresponding Fermi LAT data. We have found that for many propagation parameter configurations (PPCs) our fits improve when an extragalactic background (EGB) component is simultaneously fitted to the data. We also investigate the improvement in our fits when a flux contribution from neutralino dark matter (DM), described by the Minimal Supersymmetric Standard Model, was simultaneously fitted to the data. We consider three representative cases of neutralino DM for both Burkert and Einasto DM density profiles, in each case simultaneously fitting a boost factor of the DM contribution together with the SAB and EGB components. We have found that for several PPCs there are significant improvements in our fits, yielding both substantial EGB and DM components, where for a few of these PPCs the best-fit EGB component is consistent with recent estimates by the Fermi Collaboration. ", "introduction": "\\label{sec:intro} \\noindent A plethora of astrophysical data ranging from Galactic to cosmological scales indicate the (gravitational) influence of otherwise non-interacting particles in our Universe. This highly abundant, yet elusive, dark matter (DM), undoubtedly occupies one of the hot seats in current astroparticle physics research. This is especially true in light of the recent analyses of data from direct detection experiments, which have yielded new improved bounds but also intriguing hints of a possible signal (see, e.g., \\cite{Ahmed:2009zw, Aalseth:2010vx}). In addition, the Large Hadron Collider (LHC) \\cite{lhc}, set to probe new areas of physics, in particular, supersymmetry (SUSY), which has spawned some of the most compelling DM candidates, has also been recently activated. (For a review of SUSY DM see, e.g.,~\\cite{Jungman:1995df,Bertone:2004pz}.) An alternative strategy to identifying DM is to reconcile discrepancies between current astrophysical predictions for various radiation (e.g., positrons, antiprotons and $\\gamma$-rays) and current observations, then subsequently constrain the properties of DM particles whose annihilation (or decay) products may result in such discrepancies. To do this, one crucially has to determine how the particle and electromagnetic radiation emerging from DM annihilations, {\\it as well as} the radiation emerging from the interactions of cosmic rays (CRs) with background nuclei and interstellar radiation fields, propagate through the interstellar medium (ISM), whilst specifying the distribution of contributing DM and astrophysical sources. Unfortunately, the number of Galactic properties relevant to CR propagation are vast, their effect on these {\\it indirect signals} are largely ambiguous, and the data constraining them (e.g., local nuclear abundance ratios) is limited. However, the Large Area Telescope (LAT) of the recently launched Fermi Gamma-ray Space Telescope \\cite{Fermi} (Fermi LAT henceforth) is designed to observe \\grays with energies 20\\,MeV$\\lesssim E\\lesssim\\,$300\\,GeV, which are typical of the energies expected from the annihilation products of the popular group of DM candidates known as weakly interacting massive particles (WIMPs) \\cite{Baltz:2006sv}. Fermi LAT has acquired data from both the mid-latitude region (MLR), $10^{\\circ}<|b|<20^{\\circ}, 0^{\\circ}\\gamma\\gtrsim-1$, which also includes the EGB estimate determined by the Fermi Collaboration. We then investigated how these fits may be further improved by simultaneously fitting a \\gray flux component from dark matter described by three different candidate points within the MSSM, together with the SAB and the EGB to the Fermi LAT data. We found that the un-boosted flux, generated from motivated versions of the Burkert and Einasto profiles, was insufficient to impact the fits to the data. However, when we artificially re-normalised the DM flux using boost factors of approximately 40, 1500 and 30 with our best-fit, gaugino-dominated and mixed candidate points together with our Einasto profile (and BF's approximately 3.3 times larger when using our Burkert profile), we found that there exist many PPCs that provide further improvement in their fits to the Fermi data. Moreover, there exist PPCs that not only provide good fits to the Fermi data when including the SAB, EGB and DM components, that are significantly better than the corresponding fits when omitting the DM flux, but also require both substantial EGB and DM fluxes which are crucial in determining the quality of the fit with the parameters of the EGB component similar to those estimated previously by the Fermi Collaboration. A possible source of these enhancements could arise from local DM substructures within the Galactic halo, thought to be able to give rise to boost factors as large as 20.\\\\ \\noindent {\\bf Acknowledgments:} We would like to thank Gudlaugur Johannesson, Tim Linden, Igor Moskalenko, Stefano Profumo and Andrew Strong for their helpful comments and technical assistance. DTC is funded by the Science and Technology Facilities Council." }, "1003/1003.0672_arXiv.txt": { "abstract": "The Universe may harbor relics of the post-inflationary epoch in the form of a network of self-ordered scalar fields. Such fossils, while consistent with current cosmological data at trace levels, may leave too weak an imprint on the cosmic microwave background and the large-scale distribution of matter to allow for direct detection. The non-Gaussian statistics of the density perturbations induced by these fields, however, permit a direct means to probe for these relics. Here we calculate the bispectrum that arises in models of self-ordered scalar fields. We find a compact analytic expression for the bispectrum, evaluate it numerically, and provide a simple approximation that may be useful for data analysis. The bispectrum is largest for triangles that are aligned (have edges $k_1\\simeq 2 k_2 \\simeq 2 k_3$) as opposed to the local-model bispectrum, which peaks for squeezed triangles ($k_1\\simeq k_2 \\gg k_3$), and the equilateral bispectrum, which peaks at $k_1\\simeq k_2 \\simeq k_3$. We estimate that this non-Gaussianity should be detectable by the Planck satellite if the contribution from self-ordering scalar fields to primordial perturbations is near the current upper limit. ", "introduction": "A wealth of precise cosmological data are in good agreement with the predictions of the simplest single-field slow-roll (SFSR) inflationary models \\cite{inflation}. Still, no theorist considers these as anything more than toy models. Realistic models must surely be more complicated, and they generically predict that there should arise, at some point, observable phenomena that depart from the predictions of SFSR inflation. Some possible directions for physics beyond the SFSR approximation include multi-field models \\cite{larger,curvaton} and inflaton models with non-standard kinetic terms \\cite{Dvali:1998pa}. There has also been investigation of the consequences of topological defects \\cite{topdefects} produced toward the end of or after inflation \\cite{endinflation}. If inflation was followed by a transition associated with the breaking of a global $O(N)$ symmetry, then self-ordering scalar fields (SOSFs) are another possibly observable early-Universe relic, even if there are no topological defects (i.e., if $N > 4$). Here, the alignment of the scalar field as the Universe expands gives rise to a scale-invariant spectrum of isocurvature perturbations, without topological defects \\cite{Turok:1991qq}. Sample variance on the current data limit these perturbations to contribute no more than $\\sim10\\%$ of large-angle cosmic-microwave-background (CMB) anisotropy power \\cite{Bevis:2004wk,Crotty:2003rz}. SOSF models are parametrized simply by the number $N$ of scalar fields and the vacuum expectation value $v$. The CMB constraint implies $(v/N^{1/4}) \\lesssim 5\\times10^{15}$~GeV, as we explain below. At this low amplitude, it is unlikely that any surviving relics leave a distinct imprint on the CMB power spectrum \\cite{CMBtexture}. In recent years, non-Gaussianity has been developed as a novel tool to investigate beyond-SFSR physics \\cite{NGReview,Bernardeau:2001qr}. SFSR models do not predict that primordial perturbations should be Gaussian, but the departures from Gaussianity that they predict are unobservably small \\cite{localmodel,Wang:2000,Maldacena:2002vr}. Multi-field models \\cite{larger}, such as curvaton models \\cite{curvaton}, string-inspired DBI \\cite{Dvali:1998pa,Equil} models, and models with features in the inflaton potential \\cite{Wang:2000,Hannestad:2009yx} can all produce larger, and possibly observable, deviations from non-Gaussianity. For example, the detailed shape (triangle dependence) of the bispectrum may also help distinguish these different scenarios. The ``local-model'' bispectrum, like that which arises in curvaton and multi-field models, has a very different shape dependence than ``equilateral-model'' bispectra, like those in DBI models. Non-Gaussianity can be sought in the CMB \\cite{Luo:1993xx}, large-scale structure (LSS) \\cite{LSS}, and the abundances and properties of gravitationally-bound objects \\cite{abundances} or voids \\cite{Kamionkowski:2008sr}. Biasing may significantly amplify the effects of non-Gaussianity \\cite{Dalal:2007cu} in the galaxy distribution. The energy-density perturbations in self-ordering scalar fields are quadratic in the scalar-field perturbation, which may itself be approximated as a Gaussian field. The density perturbations induced by SOSFs are thus expected to be highly non-Gaussian \\cite{Turok:1991qq,Jaffe:1993tt,DefectBITRISPECTRA}, even in the absence of topological defects. It is thus plausible that the non-Gaussianity induced by SOSFs might be detectable, even if they provide only a secondary contribution to primordial perturbations. In this paper, we perform the first calculation of the full shape (triangle) dependence of the bispectrum from SOSFs. We follow the formalism for non-Gaussianity developed in Ref.~\\cite{Jaffe:1993tt}. We find considerably simplified formulas for the bispectrum, evaluate them numerically, and find a simple approximation to aid in data-analysis efforts. We estimate the current non-Gaussianity constraint to the model parameter space and find it to be comparable to that from the upper limit to isocurvature perturbations from CMB fluctuations. The plan of this paper is as follows: In Section \\ref{sec:scalarfielddynamics}, we define the model, write the scalar-field equations of motion, show that the dynamics are those of a nonlinear-sigma model, and introduce the large-$N$ scaling limit for the nonlinear sigma model. In Section \\ref{sec:matterperturbations}, we write the relation between the matter-density perturbation and the scalar-field perturbation. In Section \\ref{sec:powerspectrum}, we derive the power spectrum for density and curvature perturbations, discuss the normalization, and derive current constraints to the $v$-$N$ parameter space from upper limits to the SOSF contribution to CMB fluctuations. In Section \\ref{sec:bispectrum}, we discuss the calculation of the bispectrum, the central focus of this paper. We present a simplified version, our Eq.~(\\ref{eqn:biresult}), of the matter-bispectrum expression in Ref.~\\cite{Jaffe:1993tt}, evaluate it numerically, and provide a simple analytic approximation for the results. We write the bispectra for matter and curvature perturbations, define a non-Gaussianity parameter $\\fnl^\\sigma$ for the model, and estimate the current constraint to $\\fnl^\\sigma$ from the CMB. Section \\ref{sec:galaxysurveys} presents the matter bispectrum for modes that entered the horizon during radiation domination, those relevant for galaxy surveys. The central results of the paper are Eq.~(\\ref{eqn:centralresult}) for the curvature bispectrum; Eq.~(\\ref{eqn:fnlsigma}) which defines $\\fnl^\\sigma$ in terms of the SOSF model parameters $v$ and $N$; Eq.~(\\ref{eqn:g3approx}) which approximates the bispectrum function $g_3(k_1,k_2,k_3)$; and Eqs.~(\\ref{eqn:galaxybispectrum}) and (\\ref{eqn:galaxybispectrum2}) which present the matter bispectrum in a form useful for galaxy surveys. We make concluding remarks in Section \\ref{sec:discussion}. An Appendix contains some calculational details and useful approximations. ", "conclusions": "\\label{sec:discussion} If some post-inflationary physics involves the spontaneous breaking of an exact $O(N)$ symmetry with $N>4$, then the ordering of these scalar fields may provide a secondary contribution to primordial perturbations. Current constraints allow up to $\\sim10\\%$ of the power in primordial perturbations to be due to SOSFs. SOSF models are appealing from the theoretical perspective because they are simple, well-defined, and parametrized only by the symmetry-breaking scale $v$ and number $N$ of fields. In this paper we have calculated the matter and curvature bispectra induced by the ordering of such scalar fields. Given that the density perturbation is quadratic in the scalar-field perturbation, SOSF density perturbations are expected to be highly non-Gaussian, and if so, measurements of non-Gaussianity may provide the means to test these models. Here we have calculated analytically the bispectrum due to SOSFs and presented results in a way that should be easily accessible to those doing measurements with the CMB and large-scale structure. We find that the triangle-shape dependence of the bispectrum peaks for aligned triangles, unlike the local-model bispectrum, which is largest for squeezed triangles, and the equilateral bispectrum, which is largest for equilateral triangles. We have estimated a current upper limit to the non-Gaussianity parameter $\\fnl^\\sigma$ for the model and find that the implied constraints to the $v$-$N$ SOSF parameter space are competitive with those from the upper limit to CMB temperature fluctuations. Finally, we have already argued above, in Section \\ref{sec:currentconstraints}, that the correlation of modes will be similar for the large-scale modes as they enter the horizon, those relevant for large-angle CMB fluctuations. We therefore believe that rough constraints to the model can be derived from CMB measurements by assuming that the curvature bispectrum we calculate is the primordial one. Clearly, there is room for further numerical work to test our assumptions and to make our predictions more precise. In the meantime, though, we believe that our analytic approximation captures the essential physics and that our bispectrum can be used in the meantime as a ``working-horse'' model to derive constraints, from non-Gaussianity measurements, to this interesting class of models for secondary contributions to primordial perturbations. Finally, we note that the model makes a number of other predictions. Given that density perturbations are actively generated as new modes come within the horizon, vector and tensor modes will be excited, and these may give rise to interesting polarization signals \\cite{Seljak:1997ii} in the CMB and perhaps excite B modes \\cite{Kamionkowski:1996ks} in the CMB that might be distinguished from those due to inflation \\cite{Baumann:2009mq}. There will also be a scale-invariant spectrum of primordial gravitational waves produced \\cite{Krauss:1991qu} that can be sought in gravitational-wave observatories." }, "1003/1003.0444_arXiv.txt": { "abstract": "We have detected narrow \\hi~21cm and \\ci\\ absorption at $z \\sim 1.4 - 1.6$ towards Q0458$-$020 and Q2337$-$011, and use these lines to test for possible changes in the fine structure constant $\\alpha$, the proton-electron mass ratio $\\mu$, and the proton gyromagnetic ratio $g_p$. A comparison between the \\hi~21cm and \\ci\\ line redshifts yields $\\Delta X/X = [+6.8 \\pm 1.0] \\times 10^{-6}$ over $0 < \\langle z \\rangle < 1.46$, where $X = g_p \\alpha^2/\\mu$, and the errors are purely statistical, from the gaussian fits. The simple line profiles and the high sensitivity of the spectra imply that statistical errors in this comparison are an order of magnitude lower than in previous studies. Further, the \\ci\\ lines arise in cold neutral gas that also gives rise to \\hi~21cm absorption, and both background quasars are core-dominated, reducing the likelihood of systematic errors due to local velocity offsets between the hyperfine and resonance lines. The dominant source of systematic error lies in the absolute wavelength calibration of the optical spectra, which appears uncertain to $\\sim 2$~km/s, yielding a maximum error in $\\Delta X/X$ of $\\sim 6.7 \\times 10^{-6}$. Including this, we obtain $\\Delta X/X = [+6.8 \\pm 1.0 (statistical) \\pm 6.7 (max. systematic)] \\times 10^{-6}$ over $0 < \\langle z \\rangle < 1.46$. Using literature constraints on $\\Delta \\mu/\\mu$, this is inconsistent with claims of a smaller value of $\\alpha$ from the many-multiplet method, unless fractional changes in $g_p$ are larger than those in $\\alpha$ and $\\mu$. ", "introduction": "\\label{sec:intro} A critical assumption in the standard model of particle physics is that low-energy coupling constants and particle masses do not vary with space or time. This assumption breaks down in most theories that attempt to unify the standard model and general relativity (e.g. \\citealp{marciano84}). The detection of such spatio-temporal variation in coupling constants like the fine structure constant $\\alpha$, or the ratios of particle masses (e.g. the proton-electron mass ratio $\\mu \\equiv m_p/m_e$), would imply new physics beyond the standard model and is hence of great interest (e.g. \\citealp{uzan03}). Comparisons between the redshifts of spectral lines detected in distant galaxies provide an important tool to probe changes in $\\alpha$, $\\mu$ and the proton gyromagnetic ratio $g_p$ over cosmological times (e.g. \\citealp{wolfe76,dzuba99,chengalur03,flambaum07b}). Most such studies have yielded constraints on changes in $\\alpha$, $\\mu$ and $g_p$, with different systematic effects [see \\citet{kanekar08b} for a recent review]. At present, the only technique that has found statistically-significant evidence for a change in one of the fundamental constants is the ``many-multiplet'' method \\citep{dzuba99}. \\citet{murphy04} obtained $\\dal = (-5.7 \\pm 1.1) \\times 10^{-6}$ from Keck High Resolution Ultraviolet Echelle Spectrometer (HIRES) optical spectra of 143~absorbers with a mean redshift $\\langle z \\rangle = 1.75$, suggesting that $\\alpha$ was smaller at earlier times. Other studies, applying a similar technique to Very Large Telescope (VLT) data on smaller samples, have not confirmed this result (e.g. \\citealp{levshakov06,srianand07b}). However, \\citet{murphy08b} argue that the errors in these studies have been under-estimated, and that results from small samples are more prone to systematic effects related to the fitting of spectral components to complex absorption profiles. In the case of $\\mu$, \\citet{king08} used VLT Ultraviolet Echelle Spectrograph (UVES) spectra of ro-vibrational H$_2$ lines in three damped Ly$\\alpha$ systems at $z \\sim 2.6 - 3$ to find $\\dmu = (-2.6 \\pm 3.0) \\times 10^{-6}$. Note that none of the above error estimates include recently-detected systematic effects due to distortions in the wavelength scales of the HIRES and UVES spectrographs (e.g. \\citealt{griest10}). Finally, constraints on changes in $\\alpha$, $\\mu$ and $g_p$ have also been obtained from radio techniques (e.g. \\citealp{carilli00,kanekar04b,kanekar05,murphy08}), although at lower redshifts. Comparisons between the redshifts of \\hi~21cm (hyperfine) and ultraviolet resonance dipole transitions are sensitive to changes in $X \\equiv g_p \\alpha^2/\\mu$ \\citep{wolfe76}. The best resonance lines for this method are those arising from {\\it neutral} atomic species (e.g. \\ci, Fe{\\sc i}, etc), as these species are most likely to be physically associated with the \\hi. The \\ci\\ multiplets are likely to be the best among the neutral resonance transitions as they typically arise in cold gas which also gives rise to the \\hi~21cm absorption (e.g. \\citealt{jenkins01,srianand05}). The ionization potentials of \\ci\\ and \\hi\\ are also similar, 11.3~eV for \\ci\\ and 13.6~eV for \\hi\\ (cf. Mg{\\sc i}, which is more easily detectable than \\ci\\ in high-$z$ absorbers, has an ionization potential of 7.6~eV, as well as a high dielectronic recombination rate that can give significant Mg{\\sc i} absorption in warm, ionized gas; \\citealp{pettini77}). Finally, absorbers with a single (or dominant) spectral component in both \\hi~21cm and \\ci\\ transitions are best-suited for such studies. \\hi~21cm and \\ci\\ absorption have hitherto both been detected in only one high-$z$ absorber, the $z \\sim 1.776$ system towards 1331+170, yielding a weak constraint on changes in $X \\equiv g_p \\alpha^2/\\mu$ \\citep{cowie95}. However, the \\ci\\ line in this absorber has two clear spectral components \\citep{dessauges04}, implying ambiguities (and thus, large systematic errors) in the comparison with the \\hi~21cm line. In this {\\it Letter}, we report the detection of narrow, single-component \\hi~21cm and \\ci\\ absorption in two absorbers at $z \\sim 1.4-1.6$, that allow a high-sensitivity study of changes in the fundamental constants. \\begin{figure*}[t!] \\centering \\epsfig{file=fig1a.eps,width=3.4in,height=3.4in} \\epsfig{file=fig1b.eps,width=3.4in,height=3.4in} \\caption{\\hi~21cm and \\ci\\ spectra towards [A]~Q2337$-$011 and [B]~Q0458$-$020. In each panel, the solid line shows the 1-gaussian fit to the spectrum. The small-dashed and large-dashed vertical lines indicate the \\ci\\ and \\hi~21cm redshifts, respectively; the \\ci\\ redshift is seen to be higher than the \\hi~21cm redshift in both panels.} \\label{fig:spectra} \\end{figure*} ", "conclusions": "\\label{sec:discuss} Prior to this work, the most sensitive result constraining changes in $X \\equiv g_p \\alpha^2/\\mu$ using the hyperfine/resonance comparison was that of \\citet{tzanavaris07}. These authors compared redshifts of the deepest absorption in \\hi~21cm and low-ionization metal lines to obtain $\\Delta X/X = (6.3 \\pm 9.9) \\times 10^{-6}$ from a sample of nine absorbers at $0.23 < z < 2.35$. Note that these error estimate do not include systematic effects. Considering only statistical errors, our result, $\\Delta X/X = [+6.8 \\pm 1.0 (statistical)] \\times 10^{-6}$ over $0 < \\langle z \\rangle < 1.46$, is an order of magnitude more sensitive than that of \\citet{tzanavaris07}. Further, the low-ionization metal lines used by \\citet{tzanavaris07} could also arise in warm \\hi\\ or ionized gas, and are not necessarily associated with the cold \\hi\\ that gives rise to the \\hi~21cm absorption. Most of the absorbers used by \\citet{tzanavaris07} also have complex \\hi~21cm and metal line profiles, and it is not necessary that the deepest absorption in the two types of transitions arises in the same spectral component \\citep{kanekar06}. This could give systematic errors of $\\gtrsim 10$~km/s, far larger than the statistical errors of \\citet{tzanavaris07}, or the systematic errors in the present result. The comparison between hyperfine and resonance transitions directly probes changes in $X \\equiv g_p \\alpha^2/\\mu$, and one cannot obtain {\\it independent} constraints on the individual constants without additional assumptions. Our result gives $2\\times\\dal + [\\Delta g_p/g_p] - \\dmu = [+6.8 \\pm 1.0 (statistical) \\pm 6.7 (systematic)] \\times 10^{-6}$. Note that $6.7 \\times 10^{-6}$ is the {\\it maximum} estimated error due to systematics in the absolute wavelength calibration of the optical spectra, (and not a $1 \\sigma$ estimate, as in the case of the statistical error). This implies that $2\\times\\dal + [\\Delta g_p/g_p] - \\dmu \\ge [+0.1 \\pm 1.0] \\times 10^{-6}$. The Keck/HIRES result of \\citet{murphy04} is $\\dal = (-5.7 \\pm 1.1) \\times 10^{-6}$, with $\\langle z \\rangle = 1.75$. Consistency between the two results would require either (1)~an additional wavelength calibration error of $\\sim 1.7$~km/s in the \\ci\\ spectra, which appears unlikely, (2)~a ``local'' velocity offset of $\\sim 1.7$~km/s (and with the appropriate sign) between the \\hi\\ and \\ci\\ lines in {\\it both} absorbers, (3)~under-estimated errors in the many-multiplet result [e.g. the distortions in the wavelength scale found by \\citet{griest10}], or (4)~$[\\Delta g_p/g_p] - \\dmu \\ge (+1.14 \\pm 0.24) \\times 10^{-5}$ at $\\langle z \\rangle = 1.46$. In other words, assuming that the errors in one (or both) of the results have not been under-estimated, consistency between the results requires that fractional changes in $\\mu$ and/or $g_p$ are comparable to those in $\\alpha$. Strong constraints are available on fractional changes in $\\mu$ at both higher and lower redshifts, $\\dmu < 1.6 \\times 10^{-6}$ at $z \\sim 0.685$ \\citep{murphy08} and $\\dmu < 6.0 \\times 10^{-6}$ at $z \\sim 2.8$ \\citep{king08}, making it unlikely that $\\dmu \\sim 10^{-5}$ at $z \\sim 1.46$. We thus conclude that the present result appears inconsistent with the smaller value of $\\alpha$ at $\\langle z \\rangle \\approx 1.75$ found by \\citet{murphy04}, unless fractional changes in the proton gyromagnetic ratio $g_p$ are larger than those in $\\alpha$ and $\\mu$. In summary, we have detected narrow \\hi~21cm and \\ci\\ absorption in two absorbers at $z \\sim 1.4-1.6$ towards Q0458$-$020 and Q2337$-$011, using the Keck telescope, the GMRT and the GBT. The \\ci\\ and \\hi~21cm line frequencies have different dependences on the fundamental constants $\\alpha$, $\\mu \\equiv m_p/m_e$ and $g_p$, allowing us to use the measured \\ci\\ and \\hi~21cm line redshifts to test for putative changes in these constants. Comparing the \\hi~21cm and \\ci\\ redshifts in the two absorbers yields the result $\\Delta X/X = [+6.8 \\pm 1.0 (statistical) \\pm 6.7 (systematic)] \\times 10^{-6}$ over $0 < \\langle z \\rangle 1.46$, where $X \\equiv g_p \\alpha^2/\\mu$. This is inconsistent with evidence for a smaller value of $\\alpha$ at similar redshifts from the many-multiplet method, unless fractional changes in $g_p$ are larger than those in $\\alpha$ and $\\mu$. Systematic errors in the hyperfine/resonance comparison are currently dominated by errors ($\\sim 2$~km/s) in the absolute wavelength calibration of the optical spectra. However, the comparison between \\hi~21cm and \\ci\\ lines has a high sensitivity, and it should be possible to significantly reduce systematic effects by the use of new calibration techniques (e.g. laser frequency combs; \\citealt{steinmetz08}). Increasing the number of detections of redshifted, narrow \\hi~21cm and \\ci\\ absorption is thus of much importance." }, "1003/1003.5057_arXiv.txt": { "abstract": "We discuss generation of non-Gaussianity in density perturbation through the super-horizon evolution during inflation by using the so-called $\\delta N$ formalism. {\\bfs We first provide a general formula for the non-linearity parameter generated during inflation. We find that it is proportional to the slow-roll parameters, multiplied by the model dependent factors that may enhance the non-Gaussianity to the observable ranges.} Then we discuss three typical examples to illustrate how difficult to generate sizable non-Gaussianity through the super-horizon evolution during inflation. {\\bfs First example is the double inflation model, which shows that temporal violation of slow roll conditions is not enough for the generation of non-Gaussianity. Second example is the ordinary hybrid inflation model, which illustrates the importance of taking into account perturbations on small scales. Finally, we discuss Kadota-Stewart model. This model gives an example in which we have to choose rather unnatural initial conditions even if large non-Gaussianity can be generated. } ", "introduction": "Inflation solved the various cosmological problems like horizon, homogeneity, flatness problem, just by assuming that the potential energy of scalar fields dominates the expansion of the universe. Moreover, inflation scenario naturally explains the origin of the almost scale invariant density perturbation. The observations in the past decades verified many of the predictions of inflation. Almost scale invariant spectrum is now confirmed and now we even know that the violation of the exact scale invariance is 4\\% level~\\cite{Komatsu:2010fb}, which is also consistent with the simplest slow roll inflation scenario. \\subsection{Density perturbation} The amplitude of quantum fluctuation of inflaton $\\phi$ during slow-roll inflation is determined by the unique relevant mass scale at that time, e.g. the Hubble rate $H$: \\begin{equation} \\delta\\phi={H\\over 2\\pi}. \\end{equation} However, on scales as large as the Horizon radius the meaning of the amplitude of field perturbation becomes subtle because it depends on the choice of gauge. For the flat slicing gauge, in which the trace of the spatial curvature is kept unperturbed, the perturbation equation for a scalar field becomes very simple and looks very similar to the one without gravitational perturbation~\\cite{text}. Therefore the amplitude mentioned above can be in fact understood as that in the flat slicing gauge. This amplitude of perturbation can be interpreted as the dimensionless curvature perturbation on a uniform energy density surface $\\zeta$. Transformation law is given by \\begin{equation} \\zeta= H\\delta t=H{\\delta\\phi\\over\\dot\\phi}={H^2\\over 2\\pi\\dot \\phi}, \\end{equation} where $\\delta t$ is the shift in time coordinate for this transformation, which is given by $\\delta\\phi/\\dot\\phi$ using the time derivative of the background scalar field. Writing the perturbation amplitude in terms of $\\zeta$ is useful because the curvature perturbation does not evolve on super-horizon scales when the evolution path of the universe is unique. This constancy of $\\zeta$ simplifies the analysis of density perturbation during inflation for a single inflaton model. An extension of this argument to the case of multi-component inflaton is the $\\delta N$ formalism~\\cite{Starobinsky:1986fx,Salopek:1990jq,Sasaki:1995aw,Sasaki:1998ug,Lyth:2004gb,Lyth:2005fi}, as we explain below. \\subsection{Further steps from observations} Further detailed comparison between the theoretical predictions and future observations is awaited. One important observation will be tensor-type perturbation in cosmic microwave background. Tensor-type perturbation is the transverse-traceless part of the spatial metric perturbation, which is generated by the same mechanism as the perturbation of the inflaton field. Gravitational action has a overall prefactor $m_{pl}^2$, and hence we define the canonically normalised metric perturbation $\\psi_{\\mu\\nu}=m_{pl}\\delta g_{\\mu\\nu}$ so as to absorb this factor from the quadratic action. Then, the quadratic action for the transverse traceless part of $\\psi_{\\mu\\nu}$ becomes identical to the one for a massless scalar field. Therefore the amplitude of perturbation $\\psi_{\\mu\\nu}$ generated through almost de Sitter inflation is also $O(H)$. Using this relation, we have \\begin{equation} \\left\\vert{\\delta T\\over T}\\right\\vert_{tensor}=O(\\delta g_{\\mu\\nu}) =O\\left({\\psi_{\\mu\\nu}\\over m_{pl}}\\right)= O\\left({H\\over m_{pl}}\\right). \\end{equation} Thus we find that the CMB temperature fluctuation caused by the tensor perturbation directly probes the value of $H$ during inflation, i.e. the energy scale of inflation. This can be discriminated from the scalar-type perturbation by looking at B-mode polarisation~\\cite{Mandolesi:2010yw,Baumann:2008aq,quiet}. Another important observation is the non-Gaussianity of the temperature fluctuations~\\cite{Komatsu:2001rj,Bartolo:2004if}. The non-Gaussianity is caused by the non-linear dynamics of cosmological perturbation. Once we have completely understood the evolution of density perturbation at late time, the remaining non-Gaussianity which is not accounted for should have their origin in the earlier universe. The results of WMAP 7 years indicate that the local-type non-Gaussianity parameter is given as $f_{NL}=32\\pm 21$ at 68\\% confidence level~\\cite{Komatsu:2010fb}. For the Planck satellite, it is expected that the window of $f_{{\\rm NL}}$ is expected to be reduced to $O(10)$~\\cite{Mandolesi:2010yw}. Recently, non-Gaussianity of the primordial perturbation also has been studied by many authors ~\\cite{Lyth:2004gb,Lyth:2005fi,Komatsu:2001rj, Bartolo:2004if,Maldacena:2002vr,Seery:2005wm,Seery:2006vu,Seery:2005gb,Malik:2003mv,Malik:2005cy, Yokoyama:2007uu,Yokoyama:2007dw,Byrnes:2006vq,Byrnes:2007tm,Yokoyama:2008by, Creminelli:2003iq,Alishahiha:2004eh,Koyama:2010xj,Alabidi:2006hg,Malik:2006pm,Sasaki:2006kq,Dvali:2003em,Bernardeau:2002jy,Bernardeau:2002jf,Lyth:2005qk,Salem:2005nd,Seery:2006js,Alabidi:2006wa,ByrnesTasinato,Byrnes:2008wi,Byrnes:2010em,Cogollo,Rodriguez,Enqvist:2004ey,Jokinen:2005by, Chambers:2007se,Chambers:2008gu,Barnaby:2006cq,Mulryne:2009ci}. The main reason for non-Gaussianity to attract much attention is the expectation to future observations mentioned above. These observations may bring us valuable information about the dynamics of inflation. Here in this short paper, we would like to clarify the difficulties in generating large non-Gaussianity from the non-linear dynamics during inflation. ", "conclusions": "Multi-field inflation generates entropy perturbation as well as the adiabatic one. This entropy perturbation in general affects the evolution of the curvature perturbation even after the horizon crossing. Possible mechanisms of production of non-Gaussianity in the super-horizon regime can be classified into two cases. One is the production of non-Gaussianity during inflation and the other is that at the end of or after inflation. Although we could not give a general proof, it seems very difficult to produce large non-Gaussianity during inflation as far as we consider potential without fine-tuning, although there are several claims that suggest it possible \\footnote{ In the examples given in Refs.~\\cite{Byrnes:2010em,Cogollo}, it is not clear if we should say that the non-Gaussianity is generated during super-horizon evolution during inflation because distribution in field space stays Gaussian. In fact, the origin of non-Gaussianity is concentrated in the term neglected in Eq.~(\\ref{formula}).}. {\\bfs Here are several subtle issues. In this paper we claimed it difficult to construct a natural model in which large non-Gaussianity originating from the non-linear evolution is produced before inflation terminates. By contrast, generation of large non-Gaussianity at the end of or after inflation is rather easy. However, it may not be so trivial in general to identify when non-Guassianity is generated. We gave an example of double inflation in \\S~\\ref{doubleinf}. Although we did not mention it in the text, in this model the non-linear parameter $f_{NL}$ as a function of $t_F$ becomes very large temporally when the background trajectory changes its direction, but this large $f_{NL}$ does not persist long. This means that the time dependence of $f_{NL}$ is not always appropriate to read when non-Gaussianity is mainly generated. Another important point to note is the limitation of neglecting small scale perturbations. The small scale perturbations often become important for the field components that become massive during inflation. A typical example is the hybrid inflation model discussed in \\S~\\ref{hybrid}. To take into account the effect of small scale perturbations, $\\delta N$ formalism based on the background evolution for spatially homogeneous spacetime is not sufficient. We need to extend $\\delta N$ formalism incorporating collective variables that characterise statistical state such as the averaged magnitude of small scale fluctuation. Finally we should mention the naturalness of the initial conditions. It might be possible to obtain large non-Gaussianity by tuning the initial conditions, but it will not be realised in reality if they are extremely fine-tuned. As an example, we discussed Kadota-Stewart model in \\S~\\ref{KS}. However, once we start to discuss the likeliness of the chosen initial conditions, the long-standing measure problem arises. At the moment we are very far from a conclusive answer on this issue. } \\ack This work is supported by the JSPS through Grants Nos. 19540285, 21244033. We also acknowledge the support of the Grant-in-Aid for the Global COE Program ``The Next Generation of Physics, Spun from Universality and Emergence'' and the Grant-in-Aid for Scientific Research on Innovative Area N.21111006 from the Ministry of Education, Culture, Sports, Science and Technology (MEXT) of Japan." }, "1003/1003.4860_arXiv.txt": { "abstract": "Within the hierarchical framework for galaxy formation, minor merging and tidal interactions are expected to shape all large galaxies to the present day. As a consequence, most seemingly normal disk galaxies should be surrounded by spatially extended stellar 'tidal features' of low surface brightness. As part of a pilot survey for such interaction signatures, we have carried out ultra deep, wide field imaging of 8 isolated spiral galaxies in the Local Volume, with data taken at small ($D=$0.1-0.5m) robotic telescopes that provide exquisite surface brightness sensitivity ($\\mu_{lim}(V)\\sim 28.5$ mag/arcsec$^2$). This initial observational effort has led to the discovery of six previously undetected extensive (to $\\sim 30$kpc ) stellar structures in the halos surrounding these galaxies, likely debris from tidally disrupted satellites. In addition, we confirm and clarify several enormous stellar over-densities previously reported in the literature, but never before interpreted as tidal streams. Even this pilot sample of galaxies exhibits strikingly diverse morphological characteristics of these extended stellar features: {\\it great circle}-like features that resemble the Sagittarius stream surrounding the Milky Way, remote shells and giant clouds of presumed tidal debris far beyond the main stelar body, as well as jet-like features emerging from galactic disks. Together with presumed remains of already disrupted companions, our observations also capture surviving satellites caught in the act of tidal disruption. A qualitative comparison with available simulations set in a $\\Lambda$Cold Dark Matter cosmology (that model the stellar halo as the result of satellite disruption evolution) shows that the extraordinary variety of stellar morphologies detected in this pilot survey matches that seen in those simulations. The common existence of these tidal features around 'normal' disk galaxies and the morphological match to the simulations constitutes new evidence that these theoretical models also apply to a large number of other Milky Way-mass disk galaxies in the Local Volume. ", "introduction": "\\label{int} Galactic mergers have long been recognized as crucial agent in shaping and evolving galaxies (Toomre \\& Toomre 1972). Within the hierarchical galaxy formation framework (e.g., White \\& Frenk, 1991), dark matter halo mergers are a dominant evolutionary driver on the scale of galaxies. For all mass scales lower than entire galaxy clusters, the merger of two DM halos is followed quickly by the merger of the (stellar) galaxies that had been sitting at the halo's centers (e.g., Kauffmann et al. 1993). The most spectacular manifestations of this process may be {\\it major mergers} (i.e.,the coalescence of galaxies with comparable mass), that usually entails the destruction of any pre-existing stellar disk and may lead to star-bursts. Such events have been relatively rare at least since $z\\sim 1$, with only a few percent of luminous galaxies being involved in an ongoing major merger at any point in time (e.g., Robaina et al 2009). However, {\\it minor mergers} (i.e., the coalescence of a satellite galaxy and its halo with a much more luminous and massive companion) are expected to be significantly more common (e.g., Cole et al. 2000). Indeed, such minor mergers should remain frequent to the present epoch in a $\\Lambda$CDM cosmogony. As minor mergers do not destroy pre-existing stellar disks (e.g., Robertson et al. 2006), signs of recent or on-going minor mergers should be apparent around spirals, the most common type of large galaxy. If the satellite galaxies become tidally disrupted while still in an orbit that extends beyond the stellar body of the larger galaxy companion, then they should form stellar tidal 'features', which extend into the halo of the central galaxy. The observational consequences of this scenario, where the stellar halo of spiral galaxies is essentially comprised of tidal stellar debris from merged satellite galaxies, has been explored by Bullock \\& Johnston (2005, BJ05) and others (e.g., Tumlinson et al. 2009; Cooper et al. 2010). Satellites that merged on compact orbits or a long time ago have phase-mixed into a seemingly smooth component by now. In contrast, merger remnants that are only a few dynamical periods old, either because they occurred recently or on orbits with $t_{orbit}\\gtrsim 1$~ Gyr, should leave stellar streams, rings or plumes as the 'fossil record' of their interactions. BJ05 showed that such tidal debris can exhibit a wide range of morphologies and that such distinctive structural features should be common, perhaps ubiquitous around normal disk galaxies. They also showed that most of the features occur at very low surface brightness ($\\mu_{V}\\gtrsim 28.5$ mag/arcsec$^2$) and would therefore not be recognizable in traditional images of nearby galaxies. The Milky Way and the Andromeda galaxy, both resolvable into individual stars so that low surface brightness streams can more readily be seen, show a wealth of (sub-)structure in the stellar distribution of their outskirts. The most spectacular cases are the Sagittarius tidal stream surrounding the Milky Way (e.g., Majewski et al. 2003) and the Great Southern stream around the Andromeda galaxy (e.g., Ibata et al. 2001), which have become archetype fossil records of satellite galaxy mergers. But overall, the stellar halo structure of both galaxies is complex (e.g., Majewski, Munn \\& Hawley 1994; Belokurov et al 2006; Bell et al 2008; McConnachie et al 2009). Both simulations and empirical evidence suggest that there is a great deal of galaxy-to-galaxy variation in the level and the epoch of merging and hence variation in the amount and morphology of tidal debris. Therefore, a more than qualitative comparison between the predicted and observed prevalence of stellar debris around disk galaxies requires a much larger sample that necessarily must include galaxies well beyond the Local Group. The current models predict that a survey of between 50 and 100 parent galaxies reaching to a surface brightness of $\\sim 30 mag/arcsec^2$ should reveal many tens of tidal features, perhaps nearly one detectable stream per galaxy (Bullock \\& Johnston 2005; Johnston et al. 2008; Cooper et al. 2010). However, at present the evidence for tidal streams beyond the Local Group is mostly anecdotal, rather than systematic. The first cases of candidate extragalactic tidal stream candidates were reported a decade ago by Malin \\& Hardlin (1997). Using special contrast enhancement techniques on deep photographic plates, these authors were able to highlight two possible tidal streams surrounding the galaxies M83 and M104. Subsequently, deep CCD images of the nearby, edge-on galaxy NGC 5907 by Shang et al. (1998) revealed an elliptically-shaped loop in the halo of this galaxy. This was the most compelling example of an external tidal stream up to now. More recently, very deep images have clearly revealed large scale, complex structures of arcing loops in the halos of several nearby, NGC galaxies (NGC 5907: Martinez-Delgado et al. 2008; NGC 4013: Martinez-Delgado et al. 2009; NGC 891: Mouhcine, Ibata \\& Rejkuba 2010) as well as more distant, anonymous galaxies (Forbes et al. 2003). These results suggest that a more systematic survey for tidal streams in the nearby universe is not only practical but required as a new way to constrain models of galaxy formation. During the past few years, we have initiated a pilot survey of stellar tidal streams in a select number of spiral systems using modest telescopes operating at very dark sites. Ultimately, the most basic question we will seek to answer concerns the frequency of stellar streams in the Local Volume. Our aim is to test theoretical predictions by comparing substructure counts from our galaxy sample to cosmological simulations. But the models also make predictions about a number of direct observational characteristics (such as the colors, morphologies, spatial coherence and extent of halo substructures) that can be tested with the results of our survey. This paper describes the initial results of our pilot study on eight nearby spiral galaxies. These systems were selected for study because they were already suspected of being surrounded by diffuse-light over-densities based on data collected from available surveys (e.g., POSS-II; SDSS-I) and previously published deep images posted on the internet by amateur astronomers. While based on a biased sample of systems pre-selected for substructures, our pilot study serves as a proof of concept for the intended, more systematic survey of halo substructure around spiral galaxies. It also enabled us to resolve the required observing strategies and data reduction methodologies. The results presented here come from a productive collaboration between amateur and professional astronomers, dedicated to exploiting the scientific potential of modest aperture telescopes. ", "conclusions": "\\label{discussion} Our pilot survey of tidal streams associated with nearby galaxies has revealed that many spiral galaxies in the Local Universe contain significant numbers of gigantic stellar structures that resemble the features expected from hierarchical formation. Although we have only explored a handful of galaxies, our collection already presents a wide spectrum of morphologies for these stellar features. Some of them maybe have analogs in the Milky Way --- e.g., (i) great arc-like features (labeled {\\it A} in Fig.~2) that resemble the Milky Way's Sagittarius, Orphan and Anticenter streams (e.g., Majewski et al. 2003, Belokurov et al. 2006, 2007b, Grillmair 2006) and (ii) enormous clouds of debris that resemble our current understandings of the expansive Tri-And and Virgo overdensities and the Hercules-Aquila cloud in the Galactic halo (Rocha-Pinto et al. 2004, Belokurov et al. 2007a, Martinez-Delgado et al. 2007, Juric et al. 2008). Our observations also uncover enormous features resembling giant ``umbrellas'' (labeled {\\it U} in Fig.~2), isolated shells, giant plumes of debris (labeled {\\it GP} in Fig.~2), spike-like patterns (labeled {\\it S} in Fig.~2) emerging from galactic disks, long, tighly coherent streams with a central remnant core (labeled {\\it PD} in Fig.~2) and large-scale diffuse forms that are possibly related to the remnants of ancient, fully disrupted satellites. Remarkably, the diverse morphologies of stellar tidal features detected in our pilot data nearly span the range of morphologies seen in cosmologically motivated simulations. Therefore, they already represent the most comprehensive evidence matching and supporting the detailed hierarchical formation scenario predictions for galaxies similar to the Milky Way. We illustrate this through comparison with the set of eleven available snapshots featuring stellar halo models from BJ05. Each model was constructed with different merger histories in a $\\Lambda$CDM Universe and provides an external, panoramic view of surviving tidal debris (see Fig.2) from about one hundred satellites. These low-mass systems were ``injected'' into a central halo potential along orbits whose distributions is consistent with current cosmological models. Even when limiting the output of the simulations to a surface brightness comparable to our observational limit, the different stream morphologies seen fossilized in these nearby spiral halos (see Fig.~1) can be easily identified in snapshots of the model halos as well. This is illustrated by Figure 2, which compares the most conspicuous types of tidal debris detected in our survey with those visible in the model snapshots for three different assembly histories.\\footnote{This tight correspondence of observation and theory also holds for the comparison with the more recent simulations by Cooper et al. (2010).\"} From an analysis of their models, Johnston et al. (2008: see their Fig. 1) concluded that the observed stream morphology is principally dependent upon the progenitor satellite's orbit and accretion epoch. For example, {\\it great circle} features (like those seen around NGC 5907 and M63) apparently arise from near circular orbit accretion events that occurred 6-10 Gigayears ago. Straight narrow features with associated shells (e.g., the spikes in NGC 5866 and the umbrella shaped structure in NGC 4651, labeled ``Sp\" and ``U\" respectively in Fig.~2) were formed in a similar epoch from low-mass satellites in almost radial orbits. Finally, the large-scale diffuse structures observed around NGC 1055 and NGC 5866 could correspond to the mixed-type category pointed out by these authors, in which case they represent the debris of one accreted satellite that occurred longer than 10 Gigayears ago that have had time to fully mix along its orbit. Figure 2 also illustrates how the stochastic nature of halo formation (Cooper et al. 2010) leads to a large variety of substructures in the outer regions of the galaxies: each halo displays a unique and very complex pattern of stellar debris caused by different defunct companions. This morphological variance among different galaxies should be largest for brighter streams, typically formed by massive and quite recent mergers. This would explain why our modest sample, which was constructed to contain comparatively prominent streams in the Local Universe, reveals such a wide variety of detected streams. Despite the biased sample selection, the results presented here may constitute the first comprehensive observational evidence to support that the predicted great diversity of stellar halos/stream morphologies is actually present in Nature. Encouraged by this pilot survey, we have embarked on the first systematic search for stellar tidal streams in a complete, volume limited sample of spiral galaxies up to 15 Mpc (i.e., the Local Volume). This will result in the first comprehensive census of stellar stream structures within the Local Volume and it will enable meaningful statistical comparisons with cosmological simulations. The frequency of streams, their stellar populations and their morphologies will help reveal the nature of the progenitors and lend insights into the underlying structure and gravitational potential of the massive dark matter halos in which they reside. This will thereby offer a unique opportunity to study the apparently still dramatic last stages of galactic assembly in the Local Universe. In this regard, the survey will be complementary to (and directly inform the interpretation of) local galactic `archaeological' data from resolved galaxies like M31 and the Milky Way. We thank K.V. Johnston for providing the models used in this paper and for useful discussion. We also thank G. van de Venn, David Valls-Gabaud, Andrew Cooper and M. A. Gomez-Flechoso for useful discussions. D. M-D acknowledges funding from the Spanish Ministry of Education (Ramon y Cajal fellowship and research project AYA 2007-65090) and the Instituto de Astrofisica de Canarias (proyect 3I. SRM appreciates support from NSF grant AST-0807945." }, "1003/1003.1674_arXiv.txt": { "abstract": "We discuss a method to constrain the distance of blazars with unknown redshift using combined observations in the GeV and TeV regimes. We assume that the VHE spectrum corrected for the absorption through the interaction with the Extragalactic Background Light can not be harder than the spectrum in the {\\it Fermi}/LAT band. Starting from the observed VHE spectral data we derive the EBL-corrected spectra as a function of the redshift $z$ and fit them with power laws to be compared with power law fits to the LAT data. We apply the method to all TeV blazars detected by LAT with known distance and derive an empirical law describing the relation between the upper limits and the true redshifts that can be used to estimate the distance of unknown redshift blazars. Using different EBL models leads to systematic changes in the derived upper limits. Finally, we use this relation to infer the distance of the unknown redshift blazar PKS~1424+240. ", "introduction": "The extragalactic TeV sky catalogue ($E>100$ GeV), counts nowadays 35 objects\\footnote{for an updated list see: http://www.mppmu.mpg.de/$\\sim$rwagner/sources/}. Many of these sources have recently been detected also at GeV energies by the {\\it Fermi} satellite~(Abdo~et~al. 2009), allowing for the first time a quasi-continuous coverage of the spectral shape of extragalactic VHE emitters over more than 4 decades of energy. Except for two starburst galaxies and two radiogalaxies, all the others are blazars, radio-loud active galactic nuclei with a relativistic jet closely oriented toward the Earth, as described in Urry $\\&$ Padovani~(1995). The apparent luminosity of the non-thermal radiation emitted by the jet is then largely enhanced by relativistic beaming and dominates the observed high energy emission. Typically, the spectral energy distribution (SEDs) emitted from these objects, extending from radio waves to gamma-ray frequencies, is composed of two broad humps. In the case of TeV detected blazars, the first component usually peaks in the UV-X-ray band, and the second peak is located at GeV-TeV energies. The first component is identified as electron synchrotron radiation, whilst the second component is widely attributed to inverse Compton scattering of ambient photons by the same synchrotron emitting electrons. Relativistic electrons are accelerated within a region in bulk relativistic motion along the jet~(e.g. Tavecchio~et~al.~1998). VHE photons emitted by cosmological sources are effectively absorbed, through the pair production process, $\\gamma \\gamma \\rightarrow e^{+-}$, by the interaction with the so-called Extragalactic Background Light (EBL) (Stecker, de Jager $\\&$ Salamon 1992). EBL is composed of stellar light emitted and partially reprocessed by dust throughout the entire history of cosmic evolution. The expected EBL spectrum is composed by two bumps at near-infrared and far-infrared wavelengths~(Hauser~$\\&$~Dwek~2001). Direct measurement of the EBL has proved to be a difficult task, primarily due to the zodiacal light that forms a bright foreground which is difficult to suppress. Due to the lack of direct EBL knowledge, many models have been elaborated in the last years~(Stecker, Malkan $\\&$ Scully~2006; Franceschini, Rodighiero $\\&$ Vaccari~2008; Gilmore~et~al.~2009; Kneiske $\\&$ Dole~2010). Moreover, for some blazars the derivation of the intrinsic spectrum is also difficult due to the uncertainty or lack of a redshift measurement. In particular a direct spectroscopic measure of the redshift is often difficult in BL Lac objects, which are characterized by extremely weak emission lines (equivalent width $<$ 5 $\\AA$). \\begin{center} \\begin{table*} {\\small \\centering \\begin{tabular}{llcccccc} \\hline Source Name & $z[real]$ &{\\it Fermi}/LAT &VHE & $z^*$ & $z^*$ & $z^*$ & $z[rec]$ \\\\ & & slope &slope &low EBL model & mean EBL model &high EBL model&mean EBL model \\\\ \\hline \\hline Mkn 421 & 0.030 & 1.78$\\pm$0.03 & 2.3$\\pm$ 0.1$^{(1)}$ & 0.101$^{+0.021}_{-0.022}$ & 0.078$^{+0.016}_{-0.018}$ & 0.054$^{+0.012}_{-0.012}$ & 0.009$^{+0.012}_{-0.014}$ \\\\ \\hline Mkn 501 & 0.034 & 1.73$\\pm$0.06 & 2.3$\\pm$ 0.1$^{(2)}$ & 0.122$^{+0.025}_{-0.024}$ & 0.096$^{+0.018}_{-0.018}$ & 0.067$^{+0.013}_{-0.014}$ & 0.029$^{+0.013}_{-0.013}$ \\\\ \\hline 1ES 2344$+$514 & 0.044 & 1.76$\\pm$0.27 &2.9$\\pm$0.1$^{(3)}$ & 0.248$^{+0.080}_{-0.077}$ & 0.196$^{+0.056}_{-0.060}$ & 0.139$^{+0.043}_{-0.045}$ & 0.105$^{+0.041}_{-0.043}$ \\\\ \\hline Mkn 180 & 0.045 & 1.91$\\pm$0.18 &3.3$\\pm$0.7$^{(4)}$ & 0.248$^{+0.146}_{-0.150}$ & 0.196$^{+0.116}_{-0.118}$ & 0.147$^{+0.091}_{-0.089}$ & 0.104$^{+0.085}_{-0.086}$ \\\\ \\hline 1ES 1959$+$650 & 0.047 & 1.99$\\pm$0.09 &2.6$\\pm$ 0.2$^{(5)}$ & 0.111$^{+0.055}_{-0.049}$ & 0.086$^{+0.040}_{-0.040}$ & 0.058$^{+0.028}_{-0.017}$ & 0.022$^{+0.029}_{-0.029}$ \\\\ \\hline BL Lacertae & 0.069 & 2.43$\\pm$0.10 &3.6$\\pm$0.5$^{(6)}$ & 0.299$^{+0.152}_{-0.162}$ & 0.234$^{+0.116}_{-0.116}$ & 0.172$^{+0.085}_{-0.086}$ & 0.132$^{+0.085}_{-0.085}$ \\\\ \\hline PKS 2005$-$489 & 0.071 & 1.91$\\pm$0.09 & 3.2$\\pm$0.2$^{(7)}$ & 0.240$^{+0.049}_{-0.047}$ & 0.186$^{+0.036}_{-0.036}$ & 0.129$^{+0.028}_{-0.027}$ & 0.098$^{+0.026}_{-0.026}$ \\\\ \\hline W Comae & 0.102 & 2.02$\\pm$0.06 & 3.7$\\pm$0.2$^{(8)}$& 0.298$^{+0.065}_{-0.055}$ & 0.234$^{+0.046}_{-0.046}$ & 0.179$^{+0.036}_{-0.033}$ & 0.133$^{+0.034}_{-0.034}$ \\\\ \\hline PKS 2155$-$304 & 0.116 & 1.87$\\pm$0.03 & 3.4$\\pm$0.1$^{(9)}$ & 0.281$^{+0.018}_{-0.017}$ & 0.220$^{+0.014}_{-0.014}$ & 0.162$^{+0.011}_{-0.011}$ & 0.126$^{+0.010}_{-0.010}$ \\\\ \\hline 1ES 0806$+$524 & 0.138 & 2.04$\\pm$0.14 & 3.6$\\pm$1.0$^{(10)}$& 0.281$^{+0.180}_{-0.186}$ & 0.226$^{+0.138}_{-0.152}$ & 0.181$^{+0.107}_{-0.123}$ & 0.126$^{+0.101}_{-0.111}$\\\\ \\hline 1ES 1218$+$304 & 0.182 & 1.63$\\pm$0.12 &3.1$\\pm$0.3$^{(11)}$& 0.264$^{+0.107}_{-0.103}$ & 0.212$^{+0.080}_{-0.082}$ & 0.169$^{+0.061}_{-0.067}$ & 0.114$^{+0.058}_{-0.059}$ \\\\ \\hline 1ES 1011$+$496 & 0.212 & 1.82$\\pm$0.05 &4.0$\\pm$0.5$^{(12)}$& 0.667$^{+0.188}_{-0.193}$ & 0.490$^{+0.118}_{-0.124}$ & 0.348$^{+0.112}_{-0.090}$ & 0.323$^{+0.087}_{-0.092}$ \\\\ \\hline S5 0716$+$714 & 0.310$^{*,a}$ & 2.16$\\pm$0.04 &3.4$\\pm$0.5$^{(13)}$& 0.264$^{+0.107}_{-0.117}$ & 0.210$^{+0.086}_{-0.090}$ & 0.157$^{+0.064}_{-0.068}$ & 0.114$^{+0.063}_{-0.066}$\\\\ \\hline PG 1553+113 & 0.400$^b$ & 1.69$\\pm$0.04 & 4.1$\\pm$0.2$^{(14)}$& 0.779$^{+0.075}_{-0.064}$ & 0.568$^{+0.046}_{-0.046}$ & 0.395$^{+0.031}_{-0.030}$ & 0.338$^{+0.029}_{-0.029}$ \\\\ \\hline 3C66A & 0.444$^*$ & 1.93$\\pm$0.04 & 4.1$\\pm$0.4$^{(15)}$& 0.446$^{+0.076}_{-0.069}$ & 0.344$^{+0.050}_{-0.048}$ & 0.265$^{+0.038}_{-0.039}$ & 0.213$^{+0.037}_{-0.035}$\\\\ \\hline 3C279 & 0.536 & 2.34$\\pm$0.03 & 4.1$\\pm$0.7$^{(16)}$& 1.095$^{+1.066}_{-1.066}$& 0.746$^{+0.716}_{-0.716}$ & 0.440$^{+0.422}_{-0.422}$ & 0.507$^{+0.524}_{-0.524}$ \\\\ \\hline \\end{tabular} \\caption {TeV blazars used in this study. In the first column the list of sources, their redshift (second column), their {\\it Fermi}/LAT slope (third column), the VHE slope of the observed spectrum fit (fourth column). The next 3 columns show the redshift values obtained by de-absorbing the VHE spectra until the slope is the one observed by LAT, using three different EBL models, while the last column lists the corresponding reconstructed redshift of each source obtained by using $z^*$ and the fits parameters, as described in the text. $^*$: uncertain. $^a$: from Nilsson et al. (2008). $^b$:private communication with C.~W.~Danforth. 1: Acciari et al. (2009d); 2: Albert et al. (2007d); 3: Albert et al. (2007a); 4: Albert et al. (2006); 5: Tagliaferri et al. (2008); 6: Albert et al. (2007b); 7: Acero et al. (2010); 8: Acciari et al. (2009e); 9: Aharonian et al. (2005); 10: Acciari et al (2009a); 11 Acciari et al. (2009c); 12: Albert et al. (2007c); 13: Anderhub et al. (2009); 14: Prandini et al. (2009); 15: Acciari et al. (2009b); 16: Albert et al. (2008)} \\label{table_values} } \\end{table*} \\end{center} In this paper we discuss a method to derive upper limits on the redshift of a source based on the comparison between the spectral index at GeV energies as measured by LAT (unaffected by the cosmological absorption up to redshifts far beyond those of interest here) and the deabsorbed TeV spectrum. Basically, for larger distances the deabsorbed spectrum becomes harder. A solid upper limit to the redshift can be inferred deriving the redshift at which the slope of the deabsorbed spectrum coincides with that measured by LAT. Our approach can be considered complementary to those used by Stecker $\\&$ Scully (2010) and Georganopoulos,~Finke $\\&$ Reyes (2010) (see also Abdo~et~al.~2009), where the comparison of the spectral slopes at GeV and TeV energies of blazars at known distances is used to derive limits on the EBL. Starting from the derived limits, we find a simple law relating these values to real redshift, that can be used to guess the distance of unknown redshift blazars. We assume a cosmology with $h=0.72$, $\\Omega_M=0.3$ and $\\Omega_\\Lambda=0.7$. ", "conclusions": "We presented a method that allows the estimation of the quantity $z^*$, upper limit on the redshift of a TeV emitting blazar with a GeV counterpart observed by {\\it Fermi}/LAT, obtained by deabsorbing the observed TeV spectrum. In order to use the largest sample of spectra for this study, we made several assumptions: first of all we combined GeV and TeV data even if the observations in the different energy bands were not simultaneous. The impact of this choice, however, is probably moderated by fact that we do not use the flux but only the values of the slopes, less variable than the flux (unless in extreme states). For example the spectral slope of the HBL 1ES~1218+304 recently measured by the Veritas Collaboration (Acciari~et~al.~2010b) during a high flux level, matches within the errors the slope determined during its quiescient state. Secondly, in this work we use TeV spectra observed with various Cherenkov experiments, characterized by different sensitivities. This difference, especially at high energies, could affect the result, leading to systematic effects in the distance limit determination. Another possible cause of systematics could be the use of all the blazar sample, independently from the nature of the source: we didn't apply any distinction between HBL, LBL and FSRQ, characterized by a different position of the IC peak. Despite all these approximations, the method presented in this paper applied to a sample of test sources gave satisfactory results. The $z^*$ values obtained by correcting the spectra from the EBL absorption, are, in fact, all above the real redshift values if we use a mean background photon level. This suggest the use of this method for constraining the distance of unknown redshift sources. We applied the $z^*$ estimate also to two sources with uncertain distance: in both cases the limit lies below the quoted values. This result could be due to some intrinsic properties of the sources (specifically, a more moderate intrinsic spectral break between the GeV and TeV bands than that of the other sources), or to a wrong estimate of their distances. In the latter case, our method would constrain the redshift of S5~0716+714 below $0.21\\pm0.09$ and that of 3C~66A below $0.34\\pm0.05$. It can be pointed out that in the case of S5~0716+714, the redshift of 0.31 used in this work, recently reported by Nilsson et al. (2008), is estimated by assuming the luminosity of its host galaxy. Another estimate on the blazar distance, based on the spectrography of the three galaxies close to this source, gives the value of $\\sim 0.26$ for its redshift, more in agreement with our derived limit. The same procedure was applied to our sample using two extreme EBL models. The low density one gives even safer upper limits on the sources distances, while the $z^*$ obtained with the high density model are closer to the real redshift values. Even with this model, all the estimates are above or on the bisector, confirming our assumption that the deabsorbed TeV slope cannot be harder than the GeV slope reported by {\\it Fermi}/LAT. Following previous works, we tested the possibility of a linear relation between our $z^*$ estimates and the real distances of the sources. We found that the linear fit describes quite well our results, independently on the EBL model considered, although the slope and intercept of the fits are different in the three cases (Table~\\ref{params}). The relation found suggests to use the $z^*$ estimate not only to set an upper limit on unknown distances of blazars, but also, via the inverse-formula, to try an evaluation of this distance. In order to investigate this opportunity, we tested it on our sample of sources using the mean EBL model, paying a special attention to avoid biases in the calculation. The distribution of the difference $\\Delta z$ between the reconstructed and the real redshift is well described by a Gaussian peaked in zero with a $\\sigma$ of $0.05$. Once again, the uncertain redshift sources are outside the expected interval. The value of the redshift of S5~0716+714 obtained with this method is $0.11\\pm0.05$, where the error quoted is the $\\sigma$ of the $\\Delta z$ distribution. For 3C~66A, the same procedure leads to a redshift estimate of $0.213\\pm0.05$. As a final example of application, we use our procedure to PKS~1424+240, a blazar of unknown redshift recently observed in the VHE regime by Veritas~(Acciari~et~al.~2010a). The {\\it Fermi}/LAT spectrum slope measured between $0.2$ to $300$ GeV is $1.85\\pm0.05$ and the corresponding $z^*$ redshift at which the deabsorbed TeV spectrum slope equals it is $0.382\\pm0.105$, using Franceschini et al. EBL model. This result is in agreement with the value of $0.5\\pm0.1$, reported by Acciari~et~al.~(2010a), calculated applying the same procedure but only simultaneous {\\it Fermi} data. Our estimate on the most probable distance for PKS~1424+240, obtained by inverting the $z^*$ formula, is $z[rec]=0.24\\pm0.05$, where the error, as before, is assumed as the $\\sigma$ of the Gaussian fitting the $\\Delta z$ of Fig.~\\ref{dispersion_plot}." }, "1003/1003.0608_arXiv.txt": { "abstract": "% We present UBVRI photometry of three symbiotic stars ZZ CMi, TX CVn and AG Peg carried out from 1997 to 2008 in Piwnice Observatory near Toru\\'n. To search orbital periods of these stars Fourier analysis was used. For two of them, TX CVn and AG Peg, we have confirmed the earlier known periods. For ZZ CMi we found a relatively short period 218.59~days. Assuming, that the orbital period is twice longer (P=437.18~days), the double sine wave in the light curve can be interpreted by ellipsoidal effect. ", "introduction": " ", "conclusions": "" }, "1003/1003.0322_arXiv.txt": { "abstract": "We examine the alignment between Brightest Cluster Galaxies (BCGs) and their host clusters in a sample of $7031$ clusters with $0.080.65$ mag brighter than the mean of the second and third ranked galaxies) show stronger alignment than do clusters with less dominant BCGs at the $4.4\\sigma$ level. Rich clusters show a stronger alignment than do poor clusters at the $2.3\\sigma$ level. Low redshift clusters ($z<0.26$) show more alignment than do high redshift ($z>0.26$) clusters, with a difference significant at the $3.0\\sigma$ level. Our results do not depend on the algorithm used to select the cluster sample, suggesting that they are not biased by systematics of either algorithm. The correlation between BCG dominance and cluster alignment may be a consequence of the hierarchical merging process which forms the cluster. The observed redshift evolution may follow from secondary infall at late redshifts. ", "introduction": "\\indent The present cosmological paradigm \\citep{Komatsu:2009p1215} predicts the existence of coherent structures in the universe as well as alignments between structures on various scales. In a cold dark-matter dominated universe with structures forming hierarchically from the bottom up, filamentary structure formation and tidal forces may cause alignments between different elements of the hierarchy \\citep[e.g.][]{Torlina:2007p58,Faltenbacher:2008p67}. Simulations using a $\\Lambda$CDM cosmology have predicted alignments between clusters and super-clusters \\citep[e.g.][]{Basilakos:2006p69}, alignments between clusters \\citep{Hopkins:2005p348}, and alignments between galaxies and clusters \\citep{Dubinski:1998p366}. Observational studies have found alignments between galaxies and large-scale structures \\citep{Faltenbacher2009RAA.....9...41F}; clusters and clusters \\citep[e.g..][]{Binggeli:1982p376}, galaxies and clusters \\citep[e.g..][]{Sastry:1968p242,Binggeli:1982p376,Struble:1990p48,Kim:2002ASPC..268..395K,Hashimoto:2008MNRAS.390.1562H}, and galaxies within groups \\citep[e.g..][]{Yang:2006p82,Faltenbacher:2007p60,Wang:2008p65}. Most of the earlier observational probes of galaxy-cluster alignment were based on small samples and only have shallow redshift coverage. With the availability of large imaging surveys such as the Sloan Digital Sky Survey \\citep[SDSS][]{York:2000p367} we can now extend these earlier observations. In this spirit we study the alignment of brightest cluster galaxies (BCGs) with their host clusters from SDSS data in this paper. \\\\ \\indent BCGs are not simply the brightest galaxy picked from the luminosity function of a cluster. \\citet{Humason:1956p416,Tremaine:1977p403,Lauer:1988p402,Postman:1995p237,Garijo:1997p54,Dubinski:1998p366,Loh:2006p425} and others point out that BCGs have properties that suggest a formation scenario distinct from that of other galaxies, including photometric and colour homogeneity, a metric luminosity independent of cluster richness and significantly higher than that of lower-ranked galaxies, disturbed morphologies, extended low surface brightness stellar haloes and central location. \\\\ \\indent The shapes of clusters, the special properties of brightest cluster galaxies, and their alignment are clues to the process by which clusters formed. Understanding the alignment offers an observational test of the current cosmological paradigm. The alignment of BCGs with their hosts was first noted by \\citet{Sastry:1968p242} and studied by \\citet{Carter:1980p229}. Detailed observations were carried out by \\citet{Binggeli:1982p376} (hence the alignment is also known as the Binggeli effect); \\citet{Struble:1985p234} (who found no alignment with a small cluster sample); \\citet{Rhee:1987p231,Djorgovski1987,Lambas:1988p43,Struble:1990p48,Trevese:1992p39,Fuller:1999p61,Kim:2002ASPC..268..395K,Donoso:2006p66} and \\citet{Siverd:2009p516}. The effect is observed when the cluster shape is defined by the distribution of member galaxies or by the observed distribution of x-ray emitting hot gas \\citep{Hashimoto:2008MNRAS.390.1562H}. The later studies were able to confirm the alignment for clusters with redshifts up to $z=0.5$. The effect is strongest for cluster shapes measured for red, centrally concentrated galaxies, and it weakens with increasing separation between the central galaxy and the satellites. \\\\ \\indent The two leading candidate mechanisms to explain the alignment of the BCG with their parent cluster are that both are formed from infall along preferred directions in filaments \\citep[e.g.][]{Dubinski:1998p366}, and that BCGs are aligned by tidal interactions \\citep[e.g.][]{Faltenbacher:2008p67}. The points in cosmic history at which these mechanisms act are very different: whereas filamentary infall is likely to affect alignments during and immediately after cluster virialization tidal affects can act during the cluster's entire lifetime. \\citet{Ciotti:1994p478} have shown that the time scale on which a prolate galaxy's orientation is affected by a cluster's tidal field is much shorter than a Hubble time. Present-day alignments may therefore either be the result of primordial alignments stemming from the period of cluster formation or can have grown during the cluster's lifetime. By studying the redshift evolution of the alignment effect we can hope to distinguish these two cases. \\\\ \\indent In this paper we use Sloan Digital Sky Survey Data Release 6 data \\citep{York:2000p367,AdelmanMcCarthy:2008p362} to study the alignment effect in $12,755$ clusters extending out to $z=0.44$. In \\S 2 we describe our cluster selection and in \\S 3 we analyse the dependence of alignment on BCG dominance, cluster richness, and redshift. We discuss the implications of our findings in \\S 4 and summarise in \\S 5. Throughout this paper we assume an $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$, and $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$ cosmology. ", "conclusions": "\\end{table*} \\begin{figure} \\centering \\includegraphics[width=3in]{bcgdistcomp3.pdf} \\caption{Histogram of projected distances of the BCG from the cluster centre. The Dong et al. sample is shown in the solid line and the Koester et al. sample in the dashed line. In the Dong et al. sample we find that $80\\%$ of galaxies are within $350$ kpc and $50\\%$ are within $200$ kpc. In the Koester et al. sample we find $90\\%$ within $350$ kpc and $63\\%$ within $200$ kpc.} \\label{fig:disthisto} \\end{figure} \\indent To compare the two catalogues we have matched cluster centres, requiring that their projected distance be less than $2$ Mpc and their redshifts agree to within $0.02$. The sample of clusters we compare is restricted to the redshift range of the Koester catalogue ($0.14L_{\\star}$) red galaxies today are bound up in a single galaxy by $z=0.7$ and that the stellar mass in red galaxies evolves only by a factor of $~2$ since $z=1$. In the case of massive red galaxies this is mostly due to dry mergers \\citep{Faber:2007p1433, Brown:2007ApJ...654..858B}. The observation that BCG dominance evolves little with redshift (at least in the range considered here) suggests that the mechanism that causes the BCG to become dominant acts early on in the cluster's history. However, as pointed out by \\citet{Faltenbacher:2007p60} any primordial alignment on small scales is likely reduced by non-linear processes during cluster virialization. A possibility is that the BCG is somehow resistant to realignment after it has settled at the bottom of the potential well. Secondary infall episodes may reinforce the primordial alignments, which may be the cause of the slight redshift evolution of the alignment that we observe. If galaxy mergers early in a cluster's history are the cause of both BCG dominance (with more mergers creating more dominant BCGs) and its alignment with the host cluster, the uniformity of the BCG population as a whole remains puzzling, since BCGs in poorer clusters will have undergone fewer mergers than those in rich clusters. Perhaps the colour and luminosity of BCGs are not affected by mergers after a certain threshold of merging activity. Additional mergers then only affect the BCG orientation within the cluster removing other galaxies from the cluster's population (preferentially brighter more massive galaxies) letting the BCG grow in dominance. Further investigations of clusters in the redshift range $0.51$ and decline in frequency towards the present epoch as energy and metal rich gas are expelled from the galaxies into the surrounding medium. For a representative galaxy of final stellar mass $\\simeq 3\\,10^{11}\\Msun$, roughly $3\\,10^{10}\\Msun$ of recycled gas has been added to the ISM since $z\\simeq 2$ and, of that, roughly 63\\% has been expelled from the galaxy, 19\\% has been converted into new metal rich stars in the central few hundred parsecs, and 2\\% has been added to the central supermassive black hole, with the remaining 16\\% in the form hot X-ray emitting ISM. The bursts occupy a total time of $\\simeq 170$ Myr, which is roughly 1.4\\% of the available time. Of this time, the central SMBH would be seen as an UV or optical source for $\\simeq 45$\\% and $\\simeq 71$\\% of the time, respectively. Restricting to the last 8.5 Gyr, the burst occupy $\\simeq 44$ Myr, corresponding to a fiducial duty-cycle of $\\simeq 5\\,10^{-3}$. ", "introduction": "In a previous paper (Ciotti, Ostriker \\& Proga 2009, hereafter Paper I), we described in detail the physical processes that are included in our current hydrodynamical modelling of the co-evolution of a massive elliptical galaxy that contains a central supermassive black hole (hereafter SMBH). Important elements include gas shed by evolving stars, cooling flow driven infall to the central regions of this gas and the associated star bursts, accompanied by accretion onto the central SMBH and followed by nuclear and galactic winds driven from the galaxy. Feedback in both radiative and mechanical forms is taken into account, with the sources being the central SMBH, SNII and UV from the newly formed stars, SNIa from older populations, and thermalization of stellar mass losses. To our surprise in Paper I and in the companion Paper II (Shin, Ostriker \\& Ciotti 2010a) we found that despite the richness of the modeling we could not adequately represent the co-evolution of elliptical galaxies and their central SMBHs, using the conventional physics and {\\it either} purely radiative or purely mechanical feedback from the central SMBHs. In retrospective, this is not surprising because both processes operate in Nature, so that presumably both are required to produce outcomes in agreement with observations. We can summarize the main results of our previous work as follows: 1) SNIa are energetically quite important and will drive winds from elliptical galaxies (e.g, Ciotti et al.~1991; Ciotti \\& Ostriker 2007, hereafter CO07) but are only effective on the kpc scale, where the gas densities are low. They cannot prevent cooling flows and massive accumulations of gas into the inner regions of medium to massive ellipticals when their present-day rate is in accordance with the most recent observational estimates and the time evolution follows the current theoretical indications (e.g., see Pellegrini \\& Ciotti 1998). 2) Radiative feedback from central SMBHs (primarily the X-ray component) and the young star generated feedback consequent to central star bursts (e.g. Thompson, Quataert \\& Murray 2005) can balance and consume the cooling flow gas at the $10^2-10^3$ pc scale, but they will not sufficiently limit the growth of the central SMBHs. These processes - radiation feedback and energy input from stellar evolution - regulate the starburst phenomenon (Ciotti \\& Ostriker 2001, CO07). 3) Mechanical feedback from the central SMBH on the $10^1-10^2$ pc scale, mediated by a nuclear jet and the Broad Line Region winds (e.g., Binney \\& Tabor 1995; Begelman \\& Nath 2005; Begelman \\& Ruszkowski 2005; Di Matteo, Springel \\& Hernquist 2005), is efficient in limiting the growth of the SMBH, but, absent the processes noted in Point 2 above, would leave elliptical galaxies with more central star formation (fed either by cooling flows or mergers) than is observed (Papers I and II). Thus, we concluded that all three sets of processes 1, 2, and 3 acting on three different radial scales are in fact required (and, of course, as noted all three -- SNIa, AGN radiation, Broad Line Region winds -- are observed phenomena) to match what we know of the properties of elliptical galaxies. Note that characteristic time-scales, relevant from the observational point of view, are associated in a natural way to the different processes mentioned above. In practice, while SNIa, stellar mass losses and gas cooling in the central regions of the galaxy drive the global evolution of the galaxy gas mass budget on temporal scales of several Gyrs (determined by the stellar evolution clock), AGN feedback acts on shorter time scales, as each major feedback event usually spans $10^7-10^8$ yrs. The simulations also revealed that each major central outburst actually consist of several feedback events, on time scales of $\\approx 1$ Myr or less, i.e., the sound crossing time of the central kpc-scale region of the galaxy. Finally, the last characteristic time scale ($\\approx 0.5-1$ Gyrs) is that of dissipation of feedback effects on galactic scales, setted by the sound crossing time over the galaxy body. The purpose of this paper is to refine those conclusions, to show which {\\it combined models} (i.e., in which both radiative and mechanical feedback effects are allowed) best fit observations, and finally to propose observational tests of the overall picture. We recall that all the presented simulations represent the evolution of an intermediate luminosity, isolated elliptical galaxy (i.e., no external pressure is imposed at the numerical grid outer boundary), while the case of a galaxy in a cluster will be the focus of future works. We assume, following long standing observational data and recent simulations, that the elliptical galaxy in question was made earlier and is close to its current state when the calculation begins at $z\\simeq 2$ (e.g. Renzini 2006, Naab et al.~2007). The overall situation is inherently complex, and so, before diving into the details of our new computations, it may be useful to present a very rough, cartoon-level picture of the results. In Figure~\\ref{fig:cartoon} the arrows show the direction of time as the galaxy/SMBH passes through four (of the many) phases of evolution, and, since we are describing a cyclic phenomenon, we can start at any point. A very rough estimate of the fraction of a cycle spent in any given phase is shown in the sub-boxes at the upper right of each box as a ``duty cycle'', $f_{\\rm duty}$. 1) We arbitrarily begin at the upper left hand corner of the figure in the more or less quiescent phase that occupies most of the time for normal elliptical galaxies. Planetary nebulae and other sources of secondary gas, processed through stellar evolution, are added to the ambient gas everywhere in the galaxy at a rate proportional to the stellar density and with an energy due to the stellar motions which gurantees that, when the gas is thermalized, it will be approximately at the local ``virial'' temperature without need for extra energy input or output. Supernovae, primarily of type Ia are also distributed like the stars and will tend to drive a mild wind from the outer parts of the galaxy, with the inner parts being quite luminous in thermal X-rays. This is a ``normal'' giant elliptical galaxy. 2) But, the gas in the dense inner part of the galaxy is radiating far more energy than can be replaced by SNIa, stellar outbursts, cosmic rays, conduction or any other energy source and thus a ``cooling catastrophe'' occurs with a collapsing cold shell forming at $\\approx 1$ kpc from the center. As this falls towards the center, a starburst occurs of the type described by Thompson et al. (2005), and the galaxy seen as an ULIRG. A radio jet may be emitted, but the AGN flare up is at first heavily obscured and the central source will only be seen in hard X-rays. 3) Gradually, the gas is consumed, as it is transformed to new stars, and some of it is driven out in a strong wind by the combined effects of feedback from the starburst and the central SMBH, which is now exposed as an optical and then UV ``quasar'', complete with Broad Line Region (hereafter BLR) wind, optically thick disc of gas, and young stars. 4) As gas is used up or blown away, a hot cavity is formed at the center of the system and, since a shock has propagated through that volume, it is essentially like a giant supernova remnant and one expects there to be particle acceleration and non-thermal radiation from the central region. This phase has been studied in detail in Jiang et al. (2010). Then, gradually this hot bubble cools and collapses and one returns to the normal elliptical phase at stage 1. One is reminded of the Shakespearean seven ages of man in ``As You Like It''. To paraphrase: ``And one galaxy, in its time, plays many parts'' - the thermal X-ray source, the starburst, the quasar, the A+E system, etc. \\begin{figure} \\includegraphics[angle=0,scale=0.74]{f1.eps} \\caption{This diagram shows the four main phases of the feedback cycle in the life of a galaxy. Secondary gas from stellar evolution leads to a cooling flow thermal instability that feeds a central SMBH, the outbursts from which leads to an expanding hot bubble which terminates the inflow. The cycle may be repeated several times, and in each box we give the characteristic duty-cycle $f_{\\rm duty}$ associated with each phase in a standard simulation.} \\label{fig:cartoon} \\end{figure} The paper is organized as follows. In Section 2, we briefly summarize the main properties of the models adopted for the simulations, referring to Paper I for all details of galaxy model construction and input physics. In Section 3, we present the results obtained when adopting a fixed mechanical efficiency of the nuclear radiatively driven wind. In Section 4, we describe in detail how the different problems encountered in the models with fixed mechanical efficiency are solved in the preferred class of simulations, in which the mechanical feedback depends on the instantaneous accretion luminosity; some relevant observational properties of these last types of model are also described. In Section 5, we also present preliminary results obtained in a more advanced modelling of mechanical feedback. Finally, in Section 6, we discuss the main results obtained, observational tests and future developments. ", "conclusions": "In this paper we have addressed, with the aid of 1D hydrodynamical simulations, the combined effects of radiative and mechanical feedback from the central SMBH on the gas flows in elliptical galaxies. The investigation is in the line of previous papers, where the input physics and the galaxy models have been substantially improved over time. In Papers I and II we focused on purely radiative and purely mechanical feedback models, and for both the cases we found difficulties in reproducing some of the most important observational features of observed galaxies. In the present paper we explored the behavior of {\\it combined} models, i.e. when both a physically motivated implementation of radiative and mechanical feedback effects is active in the code. We briefly recall here the main secure points on which our framework is based. First of all, it is known from stellar evolution theory, and firmly supported by observations, that the recycled gas from dying stars is an important source of fuel for the central SMBH, both in its amount (summing up to 20-30\\% of the total mass in stars) and in its availability over cosmic epochs. It is also obvious that the recycled gas, arising from stars in the inner several kpc of the galaxy (assumed a giant elliptical), is necessarily a subject of a classical radiative cooling instability, leading to a collapse towards the center of metal rich gas. As a consequence, star-bursts occur and also the central SMBH is fed. The details of how much is accreted on the central SMBH vs. consumed in stars vs. ejected from the center by energy input from the starburst and AGN are uncertain. But order of magnitude estimates would have the bulk going into stars or blown out as a galactic wind, with a small amount going into the central SMBH. In addition, since at the end of a major outburst, an hot bubble remains at the center, both processes shut themselves off, and it will take a cooling time for the cycle to repeat. In other words, relaxation oscillations are to be expected, but their detailed character is uncertain. Finally, order of magnitude estimates would indicate that during the bursting phase the center would be optically thick to dust, so one would observe a largely obscured starburst and largely obscured AGN with most radiation in the far IR; as gas and dust are consumed, the central sources become visible. Much of the AGN output occurs during obscured phases: then there is a brief interval when one sees a ``normal'' quasar, and finally one would see a low X-ray luminosity and E+A spectrum galaxy, with A dominating in the central several hundred pc for $10^{7-8}$ yrs. Such figures are consistent with observed statistical occurrence of E+A galaxies in homogeneous samples that are nowadays available (e.g., Goto 2007). In sum, we find that inclusion of a modest amount of (momentum and energy driven mechanical) feedback significantly improves the correspondance with reality, with the optimal level being $<\\epsw>\\simeq 10^{-4.5}$. In the models where we took feedback to be $10^2$ more efficient, as is commonly assumed, and assuming conservation of both mass and momentum, we found that the results were unsatisfactory with too low a SMBH mass growth and too low thermal X-ray luminosity from the gas. The logic is simple. Very efficient mechanical feedback keeps gas from infalling to the central SMBH and strips the galaxy of thermal gas. In order to better study the problem, we presented three different versions of the same models, namely combined models (type A) in which the mechanical energy associated with the nuclear wind is computed under the assumption of a fixed opening angle for the wind cone; combined models (type B) in which the opening angle of the nuclear wind and the mechanical feedback efficiency are function of the Eddington normalized istantaneous bolometric luminosity of the SMBH. Finally, in the third family we briefly consider B models with the full differential equation of mechanical feedback, in which also the propagation velocity of the nuclea wind is considered. In each family of models we assumed different (but physically plausible) combinations of parameters. Overall, we have confirmed the results in CO07, and completed the model investigation in Papers I, II and V. More in detail, the main results can be summarized as follows: 1) Radiative heating (primarily due to X-rays) without any mechanical energy input greatly reduces the ``cooling flow catastrophe'' problem, but leads to a result that is still defective as compared to detailed observations of local elliptical galaxies in that the central SMBH would be too bright and too massive and the galaxy would be too blue at $z=0$ (CO07, Paper I). 2) Utilizing mechanical energy alone from an AGN wind with fixed efficiency can address the problems but does not give a solution that in detail satisfies the observations (Papers I and II). If the chosen efficiency is large, then we obtain (consistent with Di Matteo et al.~2005 and others) a giant burst, and a not too large SMBH but do not get any late time AGN and the overall duty cycle is too small. If the fixed efficiency is made low enough to avoid these problems, then one simply reverts to case (1) above, the radiative case. Purely mechanical models with luminosity-dependent mechanical efficiency (which seems to be what is indicated both by observations and detailed multi-dimensional hydrodynamical simulations of radiatively driven winds by Proga 2007, Proga et al.~2008, Kurosawa \\& Proga 2009, Kurosawa et al.~2009) perform better. 3) The combined models explored in this paper, in which both radiative and mechanical feedback are allowed are clearly better than the two limit cases described above. This family of models, with mechanical energy efficiency proportional to the luminosity, when combined with a realistic treatment of the radiative effects, does seems it to be consistent with all observations for a range of efficiencies $\\epsw$ which includes the values thought to obtain either from an analysis of the observations or from theoretical modeling of the central engines. And, of course, actual observations indicate that both feedback processes do occur: both the radiative output and the broad-line and narrow-line winds are observed from AGNs (e.g., Alexander et al.~2009, 2010). 4) It is found that radiative and mechanical feedback affect different regions of the galaxy at different evolutionary stages. During the ``quiescent'' phases, when the ISM is optically thin, radiative heating is distributed over all the galaxy body, while the mechanical feedback is deposited in a region of a kpc scale radius. This produces a characteristic feature which is not present in purely radiative models, i.e. the density profile is flatter within $\\sim 1$ kpc than in models without mechanical feedback. This in turn produce a more flat core in the observed X-ray surface brightness profile (Pellegrini, Ciotti \\& Ostriker 2010, Paper IV, in preparation). However, during the burst, the collapsing cold shells are optically thick, and most of the radiation is intercepted and re-radiated in the IR. It is found that during these short phases actually is the mechanical feedback that is playng the major role. This means that a proper description of SMBH feedback requires both the physical components, as discussed in detail in Ostriker et al.~2010. We summarize the overall situation with a cartoon in Figure \\ref{fig:scheme} in which we indicate the spatial regions within which each physical mechanism is dominant. Since the detailed hydro is highly time-dependent, the cartoon greatly oversimplifies the complex situation but shows the regions where a process is most important during the time interval that it is important. \\begin{figure} \\hskip -3truecm \\includegraphics[angle=-90,scale=0.75]{f11.eps} \\caption{This diagram shows the radial intervals where the different physical mechanisms affect the evolution of gas flows driven by stellar evolution, in a representative (isolated) elliptical galaxy.} \\label{fig:scheme} \\end{figure} Tests of the overall picture presented in these papers are numerous and obvious. i) The E+A spectrum should be found to come from $\\lsim 100$ pc regions, be younger and be significantly more metal rich than the bulk of the stars in a given galaxy. ii) The duration of the AGN bursts should be quite short with regard to cosmic time, in the range of burst duration say $0.3\\lsim \\Delta t\\lsim 30$ Myrs. iii) Since satellite galaxies in clusters have more difficulty in retaining the recycled gas than do central galaxies and isolated galaxies, the E+A phenomenon should be rarer in these systems, the central stellar cusps should be weaker and the incidence of the AGN phenomenon rarer. iv) The fraction of the time that normal elliptical galaxies spend in the bursting state should be small and fairly steeply declining with increasing cosmic time. The final SMBH luminosities will be typically $\\sim 10^{-5}\\ledd$. v) If cooling instabilities in recycled gas dominate (at late times) the fueling of AGN bursts, then evidence for merging activity should be relatively rare and the gas seen during the outbursts should be relatively rich in metals (including S-process elements). We conclude by recalling that we restricted our study for simplicity to the case of an isolated elliptical galaxy, therefore excluding the confining or stripping effect of the ICM on the galactic X-ray emitting corona. The two physical phenomena have opposite effects. On one hand, as shown in Shin et al. (2010b), the ram-pressure stripping of the ICM on the galactic atmosphere has the effect of reducing the global X-ray luminosity of the galaxy, and so also to retard or even suppress the cooling catastrophe and the associated AGN bursts. In turn, this will also reduce the starburst activity that we (invariably) find associated with strong bursts. On the other hand, galaxies residing in the central regions of a cluster are presumably more affected by external pressure effects than by ram pressure stripping. Such galaxies would be on average more X-ray luminous, and show not only larger duty-cycle values, but also younger and bluer nulear cusps near the SMBH, on the 100 pc scale, than similar galaxies orbiting in gthe cluster outskirts. Therefore, we expect that the coronal X-ray luminosity of the hot gas, and the number of central bursts of the models presented in this paper, are lower limits for the case of galaxies immersed in a realistic ICM in the central regions of a cluster. Also, the effect of mechanical feedback due to a jet, which is relevant at low accretion luminosities (i.e., during the hot phase accretion common at late times, e.g. see Allen et al.~2006, Merloni \\& Heinz 2007), is not included. The associated reduction of the accretion luminosity in the low-luminosity states will bring our models nearer to the observed Eddington ratios of low-luminosity AGNs. However, we stress that the maintenance of the SMBH masses to the observed level, in presence of the important amounts of recycled metal-rich gas produced over an Hubble time by the aging stellar populations, is mainly due to the combined effect of the feedback terms included in the simulation, i.e. SNIa and SNII heating, radiative heating, and nuclear wind feedback. The (metal rich) recycled gas from stellar evolution is present {\\it even in absence of external phenomena such as galaxy merging or input from cold gas flows}, that are often considered as the natural way to induce QSO activity. Therefore, one of the main results of our simulations (also considering all the simplifications in the treatment of physics, and of the geometry of the code), is that {\\it the evolution of an isolated galaxy, subject to internal evolution only, can be quite complicated} (e.g., see Pierce et al.~2007), and that AGN feedback will lead to central AGN and starburst activity in many widely spaced brief intervals (e.g., Shi et al.~2009). We note that the possibility of QSO activity even in absence of merging, has also been recently proposed also by others (e.g., Li et al.~2008, Kauffmann \\& Heckman 2009, Tal et al.~2009). Clearly, the main limitation of the models explored in our series of papers is the adopted spherical symmetry. This choice has been necessary as the main focus of the study has been the understanding of the physics behind AGN feedback. The necessity to explore the parameter space, and the time-expensive numerical integration of heating/cooling, star formation, and radiative trasport equations from the pc scale near the SMBH up to the hundred-kpc scale of the galaxy outer regions, in presence of multiple mutually interacting shocks, forced us to use a 1-D code. However, we are now working on 2-D simulations, that allows for a more realistic description of mechanical feedback, of the cold shell stability and fragmentation, and of the effect of angular momentum of the accreting gas. Work performed to date indicates that several aspects of the gas evolution are quite similar in the 2-D and 1-D simulations." }, "1003/1003.0052_arXiv.txt": { "abstract": "{ We have analyzed data from a multi-site campaign to observe oscillations in the F5 star Procyon. The data consist of high-precision velocities that we obtained over more than three weeks with eleven telescopes. A new method for adjusting the data weights allows us to suppress the sidelobes in the power spectrum. Stacking the power spectrum in a so-called \\'echelle diagram reveals two clear ridges that we identify with even and odd values of the angular degree ($l=0$ and 2, and $l=1$ and 3, respectively). We interpret a strong, narrow peak at 446\\,\\muHz\\ that lies close to the $l=1$ ridge as a mode with mixed character. We show that the frequencies of the ridge centroids and their separations are useful {diagnostics} for asteroseismology. In particular, variations in the large separation appear to indicate a glitch in the sound-speed profile at an acoustic depth of $\\sim$1000\\,s. We list frequencies for 55 modes extracted from the data spanning 20 radial orders, {a range comparable to the best solar data,} which will provide valuable constraints for theoretical models. A preliminary comparison with published models shows that the {offset between observed and calculated frequencies} for the radial modes is very different for Procyon than for the Sun and other cool stars. We find the mean lifetime of the modes in Procyon to be $1.29^{+0.55}_{-0.49}$\\,days, which is significantly shorter than the 2--4\\,days seen in the Sun. } ", "introduction": "{The success of helioseismology and the promise of asteroseismology have motivated numerous efforts to measure oscillations in solar-type stars. These began with ground-based observations \\citep[for recent reviews see][]{B+K2007c,AChDC2008} and now extend to space-based photometry, particularly with the {\\em CoRoT} and {\\em Kepler Missions} \\citep[e.g.,][]{MBA2008,GBChD2010}.} We have carried out a multi-site spectroscopic campaign to measure oscillations in the F5 star Procyon~A (HR 2943; HD 61421; HIP 37279). We obtained high-precision velocity observations over more than three weeks with eleven telescopes, with almost continuous coverage for the central ten days. In Paper~I \\citep{AKB2008PaperI} we described the details of the observations and data reduction, measured the mean oscillation amplitudes, gave a crude estimate for the mode lifetime and discussed slow variations in the velocity curve that we attributed to rotational modulation of active regions. In this paper we describe the procedure used to extract the mode parameters, provide a list of oscillation frequencies, and give an improved estimate of the mode lifetimes. ", "conclusions": "We have analyzed results from a multi-site campaign on Procyon that obtained high-precision velocity observations over more than three weeks \\citep[][Paper~I]{AKB2008PaperI}. We developed a new method for adjusting the weights in the time series that allowed us to minimize the sidelobes in the power spectrum that arise from diurnal gaps and so to construct an \\'echelle diagram that shows two clear ridges of power. To identify the odd and even ridges, we summed the power across several orders. We found structures characteristic of $l=0$ and 2 in one ridge and $l=1$ and 3 in the other. This identification was confirmed by comparing our Procyon data in a scaled \\'echelle diagram \\citep{B+K2010} with other stars for which the ridge identification is known. We showed that the frequencies of the ridge centroids and their large and small separations are easily measured and are useful {diagnostics} for asteroseismology. In particular, an oscillation in the large separation appears to indicate a glitch in the sound-speed profile at an acoustic depth of $\\sim$1000\\,s. We identify a strong narrow peak at 446\\,\\muHz, which falls slightly away from the $l=1$ ridge, as a mixed mode. In Table~\\ref{tab.freq.matrix} we give frequencies, extracted using iterative sine-wave fitting, for 55 modes with angular degrees $l$ of 0--3. These cover 20 radial orders and a factor of more than 4 in frequency, which reflects the broad range of excited modes in Procyon and the high S/N of our data, especially at low frequencies. Intensity measurements will suffer from a much higher stellar background at low frequencies, making it unlikely that even the best data from the {\\em Kepler Mission} will yield the wide range of frequencies found here. This is a strong argument in favor of {continuing efforts towards ground-based Doppler studies, such as} the SONG network (Stellar Observations Network Group; \\citealt{GChDA2008}), {which is currently under construction, and the proposed Antarctic instrument SIAMOIS (Seismic Interferometer to Measure Oscillations in the Interior of Stars; \\citealt{MAC2008}). } We estimated the mean lifetime of the modes by comparing the ``peakiness'' of the power spectrum with simulations and found a value of $1.29^{+0.55}_{-0.49}$\\,days, significantly below that of the Sun. A global fit to the power spectrum using Bayesian methods confirmed this result and provided evidence that the lifetime increases towards lower frequencies. {It also casts some doubt on the mode identifications. We still favor the identification discussed above, but leave open the possibility that this may need to be reversed.} Finally, comparing the observed frequencies of radial modes in Procyon with published theoretical models showed an offset that {appears to be constant with frequency, making it very different from that seen in} the Sun and other cool stars. Detailed comparisons of our results with theoretical models will be carried out in future papers. {We would be happy to make the data presented in this paper available on request.}" }, "1003/1003.1865_arXiv.txt": { "abstract": "A rigorous QED theory of the multiphoton decay of excited states in hydrogen atom is presented. The \"two-photon\" approximation is formulated which is limited by the one-photon and two-photon transitions including cascades transitions with two-photon links. This may be helpful for the strict description of the recombination process in hydrogen atom and, in principle, for the history of the hydrogen recombination in the early Universe. ", "introduction": "The recent accurate astrophysical observations and measurements of the cosmic microwave background (CMB) tempreture and polarization anisotropy \\cite{Hin}, \\cite{Page} triggered a new interest to the theory of the two-photon processes in hydrogen in view of the important role of these processes in the cosmological hydrogen recombination. The history of the hydrogen recombination in the early Universe is described in many reviews, for example \\cite{Seager}. The bound-bound one-photon transitions from the upper levels to the lower ones did not permit the atoms to reach their ground states: each photon released in such a transition in one atom was immidiately absorbed by another one. In particular, the Lyman-alpha radiation 2p-1s, being reabsorbed, reemitted and again reabsorbed, did not allow the radiation to escape the interaction with the matter. As it was first established in \\cite{Zeld}, \\cite{Peebles} the two-photon 2s-1s radiative transition presents one of the main channels for the radiation escape from the interaction with matter. Hence, the recent properties of the CMB are essentially defined by the two-photon decay processes during the cosmological recombination epoch. In \\cite{Dubrovich}, \\cite{Wong} it was argued that the $ns\\rightarrow 1s$ ($n>2$) and $nd\\rightarrow 1s$ two-photon transitions can also give a sizeable contribution to the process of the radiation escape from the interaction with the matter. Recently this problem was investigated thoroughly in the theoretical astrophysical studies in \\cite{J.Chluba}, \\cite{Hirata}. There is a crucial difference between the decay of the $ns$ ($n>2$) or $nd$ levels and the $2s$ decay level. This difference is due to the presence of the cascade transitions as the dominant decay channels in case of $ns$ ($n>2$) and $nd$ levels. For the $2s$ level the cascade transitions are absent. Since the cascade photons can be effectively reabsorbed, the problem of separation of the \"pure\" two-photon contribution from the cascade contribution arises. An interference between the two decay channels also should be taken into account. This problem appears to be not at all trivial and requires an application of the methods of the Quantum Electrodynamocs (QED) for the bound electrons. Quantum Mechanical theory for the two-photon transitions was first developed by G\\\"{o}ppert-Mayer \\cite{Goepp} and the first evaluation of the two-photon $2s\\rightarrow 1s+2\\gamma(E1)$ decay rate in hydrogen was performed by Breit and Teller \\cite{Breit}. The accurate nonrelativistic calculation for this transition rate was given in \\cite{Klarsfeld}; fully relativistic calculations, valid also for the H-like Highly Charged Ions (HCI) with arbitrary $Z$ (nuclear charge) values were performed in \\cite{GD82}-\\cite{Santos}. The most accurate recent calculation for this transition rate with the QED radiative corrections taken into account belongs to Jentschura \\cite{Jent}. As well as for the neutral hydrogen, the cascade problem does not arise for the transition $2s\\rightarrow 1s+2\\gamma(E1)$ in the HCI with arbitrary $Z$ values. The two-photon transitions were investigated theoretically and experimentally also in the few-electron and many-electron atoms and ions. In particular, the two-photon transition $1s2s\\,^1S_0\\rightarrow (1s)^2\\,^1S_0+2\\gamma(E1)$ transition rate for the neutral He atom was first evaluated in \\cite{Dalg}. This decay channel also does not contain cascade contribution. The cascade problem first did arise in connection with the decay of the metastable $2^3P_0$ level in He-like Uranium: $2^3P_0\\rightarrow1^1S_0+\\gamma(E1)+\\gamma(M1)$. In this case there are two possible cascade transitions: $2^3P_0\\rightarrow 2^3S_1+\\gamma(E1)\\rightarrow 1^1S_0+\\gamma(E1)+\\gamma(M1)$ and $2^3P_0\\rightarrow 2^3P_1+\\gamma(M1)\\rightarrow 1^1S_0+\\gamma(M1)+\\gamma(E1)$. The corresponding decay rate was first evaluated by Drake \\cite{Drake}. Later Savukov and Johnson \\cite{Savukov} performed similar calculation for a variety of He-like ions ($50\\leq Z\\leq 92$). In \\cite{Drake}, \\cite{Savukov} the \"pure\" two-photon contribution was obtained by subtraction of a Lorentzian fir for the cascade contribution from the total two-photon decay frequency distribution. A rigorous QED approach for the evaluation of the two-photon decay probability in presence of cascades was developed in \\cite{LabShon} on the basis of the Line Profile Approach (LPA) in QED, i.e. the QED theory of the spectral line profile (see \\cite{AndrLab}). The LPA consists of a standard evaluation of the decay probability as a transition probability to the lower levels. In the presence of cascades the integral over emitted photon frequency distribution becomes divergent due to the singular terms, corresponding to the cascade resonances. To avoid such a singularity, the resummation of an infinite series of the electron self-energy insertions into the electron propagator was performed in \\cite{LabShon}. This resummation converts into a geometric progression and in this way the electron self-energy matrix element (and the level width as its imaginary part) enters the energy denominator and shifts the pole from the real axis into the complex energy plane, thus making the transition probability integral finite. With this approach F. Low \\cite{Low} first derived the Lorentz profile from QED. In \\cite{Drake}, \\cite{Savukov} the level widths in the energy denominators were also introduced, though as the empirical parameters. In \\cite{LabShon} the ambiguity of the separation of the \"pure\" two-photon decay and cascades was first revealed for HCI; it was shown also that the interference terms can essentially contribute to the total decay probability. Nearly at the same time when the paper \\cite{Drake} did arrive, the cascade problem was discussed also for the $ns$ ($n>2$), $nd$ transitions in the hydrogen atom \\cite{cea86}, \\cite{fsm88}. In these works the \"pure\" two-photon contribution was obtained simply by omitting the resonant (singular) terms, responsible for the cascades. This approach was criticized later in \\cite{J.Chluba}. Another (\"alternative\") method which formally allows for the separate determination of the \"pure\" two-photon contribution in case of the two-photon transitions with cascades was developed in a series of works by U. Jentschura \\cite{Jent1}-\\cite{Jent3}. This approach contradicts to the LPA results. The LPA was applied to the $3s-1s$ transition (including cascade) in hydrogen in \\cite{LSP}, where the ambiguity of the separation of the \"pure\" two-photon and the cascade contributions was again demonstrated numerically, as in the case of the HCI \\cite{LabShon}. Very recently a paper \\cite{Amaro} did arrive where an attempt was made to find a compromise between LPA and \"alternative\" approach. A reasonable agreement between the numerical results obtained by both methods was found. However, to our mind, the disagreement between the LPA and \"alternative\" approach is of conceptual character and cannot be eliminated. Thus from the QED point of view only the total two-photon frequency distribution has a direct physical sense in case of the two-photon decays with cascades. This quantity should be a basic tool for the description of the two-photon processes in astrophysics. The employment of the \"1+1\" approximation for the description of cascades should be avoided. Along this way the most recent astrophysical theories \\cite{J.Chluba}, \\cite{Hirata} are built. Still the \"1+1\" approximation is not fully excluded from the considerations in \\cite{J.Chluba}, \\cite{Hirata}. In view of the recent very accurate (with relative accuracy $\\sim 1\\%$) measurements of the properties of CMB \\cite{Hin}, \\cite{Page} and with expectation of the even more accurate ($\\sim 0.1\\%$) measurements in the near future, the theory of the cosmological recombination free of any uncertainties connected with the separation of the \"pure\" two-photon and cascade contributions should be formulated. In the present paper we will formulate such a theory for the two-photon and the multiphoton decays in hydrogen. In this theory only two types of the level decays should be present: the direct one-photon decays when they are allowed and the total two-photon decays without separation of the \"pure\" two-photon decays and cascades. The total solution of the problem formulated above consists of two steps. First, the pure QED problem of the description of the multiphoton transitions in hydrogen in the \"two-photon\" approximation should be resolved. That is, all the decays of the excited levels should be classified and described either as the direct one-photon transitions to the ground state or as the two-photon transitions with cascades. In the \"two-photon\" approximation transitions with more than two nonresonant photons should be neglected. The formulation of the \"two-photon\" approximation should finalize the first step of the studies. The present paper will concern only this first step. An important feature of the rigorous QED treatment of the process of recombination is that we have to trace the decay of every particular level up to the ground state. This is of course not the full picture of the recombination process. To be more close to the cosmological recombination one has to consider the transitions from the continuous states (plasma electrons) down to the ground state, taking into account the rescattering processes. This would correspond to the second step mentioned above. However, the accurate treatment of the recombination process from the particular excited level, as presented in this paper, also may be of interest. In particular, we demonstrate that the consequent QED treatment of the $3p$ level decay should include the two-photon contribution comparable with the widely discussed two-photon decay of $3s$ level \\cite{cea86}-\\cite{Amaro}. In this paper we limit ourselves only with electric dipole transitions (both in one- and two-photon decays) and ignore $n'd\\rightarrow ns$ transitions which also are of importance (\\cite{Dubrovich}-\\cite{Hirata}). At the second step one should modify the basic astrophysical equations describing the level population in hydrogen in such a way that the imput data for them should be, apart from the direct one-photon transition probabilities, only the total two-photon decay rates, including cascades, without separating out the \"pure\" two-photon decay rates. The use of the \"1+1\" approximation should be fully avoided. This task is beyond the scope of our paper. Our paper is organized as follows. In Section II we formulate the basic concepts for the LPA-based theory for the two-photon decay with cascades. The two-photon approximation for the description of the multiphoton transitions in hydrogen is introduced. In Section III the standard derivation of the transition rate for the Lyman-$\\alpha$ $2p-1s$ transition is presented and in Section IV the standard QED derivation of the Lorentz profile for this emission process is given. The same is done in Section V for the two-photon decay $2s-1s$. The decay of the $3s$ level is considered in Section VI and the ambiguity of separation of the \"pure\" two-photon and cascade contributions is demonstrated. The decay of $3p$ level in the two-photon approximation is described in Section VII, where it is shown that this decay also contains the two-photon contribution comparable with the two-photon contribution to the decay of $3s$-level. In Section VIII an investigation of the decay of $4s$ level in the \"two-photon\" approximation is performed which gives the clue to the general formulation of the two-photon approximation in the theory of the multiphoton transitions. Section IX contains discussion of the results and conclusions. ", "conclusions": "In our paper we have considered the processes of multiphoton transitions for hydrogenic atom. Recent astrophysical investigations necessitate the detailed analysis of the multiphoton emission processes and, namely, of the separation of the \"pure\" multiphoton radiation and cascade emission. The \"pure\" two-photon emission leads to the photon escape from the matter and, thus, presents the formation mechanism for the background radiation. We began with the standard QED derivation of the Lorentz profile for emission process. After this we have investigated the decay of the $3s$ level and showed the ambiguity of separation of the \"pure\" two-photon and cascade contributions. We demonstrated that the strict separation of the \"pure\" two-photon and cascade contributions for $3s$-level decay in hydrogen is impossible. Moreover, we show that even the approximate separation of these two decay channels cannot be achieved with an accuracy, required in modern astrophysical investigations (i.e. at 1\\% level) of the recombination history of hydrogen in the early Universe. We formulated the \"two-photon\" approximation which make it possible to separate out cascade emission with two-photon links and show that this type of cascades gives the contribution to the two-photon transitions, i.e. to the radiation escape, comparable with the contribution of the \"direct\" two-photon transitions. On a basis of this approximation the decay of $3p$ level was described. It was shown that the cascade two-photon decay rates are comparable with the \"pure\" two-photon contribution of the $3s\\rightarrow 1s+2\\gamma(E1)$ channel. Unlike the $3s\\rightarrow 1s+2\\gamma(E1)$ case, where the contribution of the \"pure\" two-photon decay is nonseparable from the cascade contribution, for $3p\\rightarrow 1s+3\\gamma(E1)$ decay only the cascade contributions with the two-photon links should be taken into account. The reason is that the \"pure\" 3-photon contribution is beyond the accuracy of the \"two-photon\" approximation, adopted in this paper. The four-photon emission process of the $4s$ level was considered also. The result is: the decay rates for the cascades with the two-photon links are comparable with the contribution of the \"direct\" two-photon transitions. The main goal of our paper is the formulation of the \"two-photon\" approximation which allows for the rigarous incorporation of all types of the two-photon processes. This may be important for the more accurate astrophysical investigations of the cosmical radiation background. \\begin{center} Acknowledgments \\end{center} The authors acknowledge financial support from RFBR (grant Nr. 08-02-00026). The work of D.S. was supported by the Non-profit Foundation ``Dynasty'' (Moscow). The authors acknowledge also the support by the Program of development of scientific potential of High School, Ministry of Education and Science of Russian Federation (grant Nr. 2.1.1/1136 and goskontrakt $\\Pi$1334)." }, "1003/1003.4392_arXiv.txt": { "abstract": "We consider the spin angular momentum evolution of the accreting components of Algol-type binary stars. In wider Algols the accretion is through a disc so that the accreted material can transfer enough angular momentum to the gainer that material at its equator should be spinning at break-up. We demonstrate that even a small amount of mass transfer, much less than required to produce today's mass ratios, transfers enough angular momentum to spin the gainer up to this critical rotation velocity. However the accretors in these systems have spins typically between $10$ and~$40\\,$per cent of the critical rate. So some mechanism for angular momentum loss from the gainers is required. Unlike solar type chromospherically active stars, with enhanced magnetic activity which leads to angular momentum and mass loss, the gainers in classical Algols have radiative envelopes. We further find that normal radiative tides are far too weak to account for the necessary angular momentum loss. Thus enhanced mass loss in a stellar wind seems to be required to spin down the gainers in classical Algol systems. We consider generation of magnetic fields in the radiative atmospheres in a differentially rotating star and the possibility of angular momentum loss driven by strong stellar winds in the intermediate mass stars, such as the primaries of the Algols. Differential rotation, induced by the accretion itself, may produce such winds which carry away enough angular momentum to reduce their rotational velocities to the today's observed values. We apply this model to two systems with initial periods of 5\\,d, one with initial masses $5$ and $3\\,\\rm{M}_{\\odot}$ and the other with $3.2$ and $2\\,\\rm{M}_{\\odot}$. Our calculations show that, if the mass outflow rate in the stellar wind is about $10\\,$per cent of the accretion rate and the dipole magnetic field is stronger than about $1\\,$kG, the spin rate of the gainer is reduced to below break-up velocity even in the fast phase of mass transfer. Larger mass loss is needed for smaller magnetic fields. The slow rotation of the gainers in the classical Algol systems is explained by a balance between the spin-up by mass accretion and spin-down by a stellar wind linked to a magnetic field. ", "introduction": "Evolution of single stars is now well modelled \\citep[see for example][]{pols1995}. There remain concerns with mass loss, rotation and convection but appropriate and successful empirical treatments exist. Evolution of a binary star has several additional complications associated with interaction between the components. Since solving the mystery of Algol systems \\citep{hoyle1955,crawford1955}, the prototype of semi-detached Algol-type binary stars with one evolved and one main-sequence component, we realize that there are some stages of evolution when interaction between the components is unavoidable. We must therefore take into account, in our calculations, the mass transfer and mass loss together with any angular momentum and magnetic interaction between the components, at least in some critical phases, to fully understand evolution of a binary system. Over the last few decades, the evolution of Algols has been modelled with well defined approximations such as conservation of total mass and total orbital angular momentum. The effect of mass transfer on the structure of both stars can be modelled reasonably well. The angular momentum transfer during mass exchange, however, is not well understood. As we shall see in section~\\ref{accdiscs} there are some episodes of mass transfer in Algols when accretion discs or disc-like structures form around the mass gainer. Current approximations of binary star evolution do not adequately explain the spin angular momentum of the mass-gaining components because the high specific angular momentum of the disc material should easily spin these stars up to their critical break-up rotational velocities in less than the time needed to reverse the mass ratio of system and enter the Algol phase. Here we discuss formation of discs in classical Algol systems and consider the spin angular momentum evolution of mass accreting components, taking into account discs, tides and magnetic stellar winds. We demonstrate that tidal effects play a minor role in the removal of excess angular momentum from the gainers and rely on a magnetically locked stellar wind to do this. \\begin{table*} \\caption{The absolute parameters of Algol primary components \\citep{ibanoglu2006}. For each star, columns $2-8$ are the orbital period, mass ratio, masses, radii and inclination. Then $v_{\\rm syn}$ would be the equatorial velocity of the mass gainer (star~1) if it were synchronous while $v_{\\rm eq}\\sin i$ is the measured projected velocity and $F = v_{\\rm eq}/v_{\\rm syn}$. \\label{table}} \\begin{tabular}{lccccccccccr} \\hline Name\t&$P/\\rm d$\t&$q$\t&$M_{1}/\\rm{M}_{\\odot}$\t&$M_{2}/\\rm{M}_{\\odot}$\t&$R_{1}/\\rm{R}_{\\odot}$\t&$R_{2}/\\rm{R}_{\\odot}$\t&$i$ \t&$v_{\\rm syn} \\sin i$ \t&$v_{\\rm eq} \\sin i$\t&$F$\t&Ref.\\\\ \\hline TW And\t&4.12\t&0.210\t&1.68\t&0.32\t&2.19\t&3.37\t&87\t&27\t&32\t&1.19\t&5\\\\ KO Aql\t&2.86\t&0.217\t&2.53\t&0.55\t&1.74\t&3.34\t&78\t&30\t&41\t&1.37\t&4\\\\ IM Aur\t&1.25\t&0.311\t&2.24\t&0.76\t&2.57\t&1.74\t&75\t&101\t&135\t&1.34\t&1\\\\ R CMa\t&1.14\t&0.170\t&1.07\t&0.17\t&1.50\t&1.15\t&80\t&76\t&98\t&1.29\t&1\\\\ S Cnc\t&9.48\t&0.090\t&2.51\t&0.23\t&2.15\t&5.25\t&83\t&12\t&174\t&14.5\t&5\\\\ RZ Cas\t&1.20\t&0.351\t&2.10\t&0.74\t&1.67\t&1.94\t&83\t&62\t&87\t&1.40\t&3\\\\ TV Cas\t&1.81\t&0.470\t&3.78\t&1.53\t&3.15\t&3.29\t&79\t&90\t&79\t&0.88\t&3\\\\ U Cep\t&2.49\t&0.550\t&3.57\t&1.86\t&2.41\t&4.40\t&88\t&56\t&437\t&7.80\t&3\\\\ RS Cep\t&12.4\t&0.145\t&2.83\t&0.41\t&2.65\t&7.63\t&87\t&11\t&170\t&15.4\t&1\\\\ XX Cep\t&2.34\t&0.150\t&2.03\t&0.33\t&2.12\t&2.25\t&85\t&46\t&47\t&1.02\t&2\\\\ U CrB\t&3.45\t&0.289\t&4.74\t&1.46\t&2.79\t&4.83\t&82\t&48\t&60\t&1.25\t&1\\\\ SW Cyg\t&4.57\t&0.190\t&2.50\t&0.50\t&2.60\t&4.30\t&83\t&22\t&196\t&8.91\t&3\\\\ WW Cyg\t&3.32\t&0.310\t&2.10\t&0.60\t&2.00\t&7.00\t&89\t&31\t&41\t&1.32\t&2\\\\ TW Dra\t&2.81\t&0.470\t&1.70\t&0.80\t&2.40\t&3.40\t&86\t&43\t&37\t&0.86\t&2\\\\ AI Dra\t&1.20\t&0.429\t&2.86\t&1.34\t&2.17\t&2.42\t&78\t&83\t&85\t&1.02\t&1\\\\ S Equ\t&3.44\t&0.130\t&3.24\t&0.42\t&2.74\t&3.24\t&87\t&40\t&52\t&1.30\t&4\\\\ AS Eri\t&2.66\t&0.110\t&1.92\t&0.21\t&1.57\t&2.19\t&80\t&29\t&36\t&1.02\t&5\\\\ RX Gem\t&12.2\t&0.254\t&4.40\t&0.80\t&4.80\t&7.00\t&85\t&20\t&157\t&7.85\t&2\\\\ RY Gem\t&9.30\t&0.193\t&2.04\t&0.39\t&2.38\t&6.19\t&83\t&13\t&70\t&5.38\t&5\\\\ AD Her\t&9.77\t&0.350\t&2.90\t&0.90\t&2.60\t&7.70\t&84\t&13\t&143\t&11.0\t&2\\\\ TT Hya\t&6.95\t&0.224\t&2.63\t&0.59\t&1.95\t&5.87\t&84\t&15\t&164\t&10.9\t&2\\\\ $\\delta$ Lib\t&2.33\t&0.345\t&4.70\t&1.70\t&4.12\t&3.88\t&81\t&89\t&68\t&0.76\t&1\\\\ AU Mon\t&11.1\t&0.199\t&5.93\t&1.18\t&5.28\t&10.04\t&79\t& 24\t&124\t&5.17\t&5\\\\ TU Mon\t&5.09\t&0.210\t&12.60\t&2.70\t&5.60\t&7.10\t&89\t&56\t&153\t&2.73\t&2\\\\ AT Peg\t&1.15\t&0.484\t&2.50\t&1.21\t&1.91\t&2.11\t&76\t&80\t&82\t&1.02\t&1\\\\ $\\beta$ Per\t&2.87\t&0.217\t&3.70\t&0.81\t&2.74\t&3.60\t&82\t&51\t&52\t&1.02\t&3\\\\ RW Per\t&14.2\t&0.150\t&2.56\t&0.38\t&2.80\t&7.30\t&81\t&10\t&161\t&16.1\t&3\\\\ RY Per\t&6.86\t&0.271\t&6.24\t&1.69\t&4.06\t&8.10\t&83\t&30\t&213\t&7.10\t&1\\\\ Y Psc\t&3.77\t&0.250\t&2.80\t&0.70\t&3.06\t&3.98\t&87\t&37\t&38\t&1.03\t&3\\\\ RZ Sct\t&15.2\t&0.216\t&5.50\t&1.50\t&11.00\t&14.00\t&83\t&36\t&222\t&6.17\t&3\\\\ V356 Sgr\t&8.89\t&0.380\t&12.20\t&4.70\t&8.50\t&15.40\t&85\t&48\t&212\t&4.42\t&1\\\\ V505 Sgr\t&1.18\t&0.520\t&2.68\t&1.23\t&2.24\t&2.17\t&80\t&85\t&101\t&1.19\t&1\\\\ U Sge\t&3.38\t&0.370\t&4.45\t&1.65\t&4.00\t&5.00\t&90\t&60\t&76\t&1.27\t&3\\\\ $\\lambda$ Tau\t&3.95\t&0.263\t&7.18\t&1.89\t&6.40\t&5.30\t&76\t&80\t&88\t&1.10\t&1\\\\ TX Uma\t&3.06\t&0.248\t&4.76\t&1.18\t&2.83\t&4.24\t&82\t&46\t&63\t&1.37\t&2\\\\ Z Vul\t&2.45\t&0.430\t&5.40\t&2.30\t&4.30\t&4.50\t&89\t&89\t&135\t&1.52\t&1\\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[References:]{\\small (1) \\cite{vanhamme90} , (2) \\cite{etzel93},(3) \\cite{mukher96}, (4) \\cite{soydugan2007} and (5) \\cite{glazunova2008}.} \\end{list} \\end{table*} ", "conclusions": "It has been known for some time that many long-period Algol systems have accretion discs. Accreting material from such a disc should increase the spin rate of the more massive component up to its break-up speed as soon as even a small fraction of the mass has been transferred \\citep{demink2007}. All the classical Algols have a less massive evolved and a more massive main-sequence component. Therefore a substantial amount of mass from the initially more massive star must be either lost or transferred to its companion. The angular velocities of the more massive components in many Algols are measured with great accuracy. These measurements show that the gainers rotate somewhat more slowly than their break-up rates. Thus there must be a mechanism which removes spin angular momentum from the rapidly rotating hot star. We have demonstrated that tidal interaction between the components and the orbit is too weak to do this fast enough. We show, however, that magnetic braking, driven by a magnetic dynamo that is maintained by the accretion itself, can remove sufficient angular momentum. \\citet{mm2004} hypothesized that a dynamo operating in a sheared radiative region can generate flux tubes which are able to rise the surface of the star. A shear dynamo in the extended radiative envelope of a massive star can therefore serve as an efficient supplier of magnetic flux to the surface of the star. We propose a similar self-consistent model in which a dynamo operates in shear-unstable material throughout the radiative envelopes of the gainers in Algols with orbital periods longer than $4-5\\,$d. Such stars can lose angular momentum in magnetized stellar winds, leading to non-conservative angular momentum evolution. Under this non-conservative evolution, an Algol system continues to evolve with relatively little mass loss at a rate $\\dot{M}_{\\rm w}\\approx 0.1 \\dot{M}_{\\rm acc}$ but corotating to a relatively large Alfv\\'en radius $R_{\\rm A}$. When it reaches the classical Algol phase the less massive component is observed to be more evolved than the massive one. Spruit-Tayler instabilities seem most likely to be responsible for the operation of a magnetic dynamo in massive stars with radiative envelopes \\citep{mm2005}. However we need to construct two dimensional models with an accretion disc to test how such a dynamo really operates in Algols. Such numerical models could also supply more realistic Alfv\\'{e}n radii based on the field geometry. Using such detailed models we should be able to find an equilibrium spin-ratio which can explain the observations in Fig~\\ref{figprot}. We might also take into account the effect of wind induced hydrodynamic instabilities suggested by \\citet{lignieres1996} which may also effect the differential rotation parameter. Using a very simple model we computed the spin angular momentum evolution of gainers with discs in two systems with an initial orbital period of $5\\,$d but different masses. The results show that even a small amount of mass, about 10\\,per cent of the transferred material, lost by the gainer with a magnetic field of $1\\,$kG is sufficient to slow down the star to below its break-up velocity. As the magnetic field strength increases the rotational velocity of the star decreases for the same amount of mass loss." }, "1003/1003.4671_arXiv.txt": { "abstract": "The Cartwheel galaxy harbours more Ultra--Luminous X--ray sources (ULXs) than any other galaxy observed so far, and as such it is a particularly interesting target to study them. In this paper we analyse the three {\\sl Chandra} observations of the brightest ULX (N10) in the Cartwheel galaxy, in light of current theoretical models suggested to explain such still elusive objects. For each model we derive the relevant spectral parameters. Based on self--consistency arguments we can interpret N10 as an accreting binary system powered by a~$\\sim100\\msun$ black hole. A young supernova strongly interacting with its surroundings is a likely alternative, that can be discarded only with the evidence of a flux increase from future observations. ", "introduction": "\\label{s-intro} The Ultra--Luminous X--ray sources (ULXs) are a class of very bright ($L_X\\simeq 10^{39}-10^{41}\\ergs$), point--like X--ray sources detected off the nuclei in several galaxies \\citep{Fab06}. Although they were discovered more than 20 years ago (by the {\\sl Einstein} X--ray satellite: \\citealp{Lon83,Fab89}), their nature remains unclear. If the ULXs are powered by accretion, their luminosity exceeds the Eddington limit for a stellar--sized object, sometimes by a factor~$\\sim1000$. This has led some authors \\citep{Col99} to postulate that the ULXs are black holes with masses $10^2-10^4\\msun$, intermediate between the black holes expected from the final evolution of massive stars and the super--massive black holes of $10^6-10^9\\msun$ powering the Active Galactic Nuclei. Possible formation scenarios of these intermediate mass black holes (IMBH) are discussed in \\citet{Mil02} (formation in globular clusters), and \\citet{Por04} (in young super--massive star clusters). An alternative view places the ULXs in the more familiar realm of stellar--sized black holes. According to \\citet{Kin01}, the ULXs are black holes of $M_\\bullet \\simeq 10\\msun$, accreting mass from a disc, at a rate above their Eddington limit. In this regime, the inner accretion disc is thick, and may collimate the outgoing radiation. The anisotropy leads the observer to overestimate the luminosity (and the mass of the hole) by a large factor, up to $10-100$. Alternatively, some models suggest that the ULXs may actually radiate above the Eddington limit. They are slim disc models \\citep{Ebi03}, photon bubble dominated discs \\citep{Beg02, Beg06} and two--phase radiatively efficient discs \\citep{Soc06}. A combination of beaming and super--Eddington accretion is also possible \\citep{Pou07, Kin08, Kin09}. More recently, a third possibility has been emerging, namely that the ULXs are black holes of $30-90\\msun$ produced by the final evolution of massive stars in a metal poor environment \\citep{Zam09}. In such an environment the black hole's progenitor does not loose much of its original mass by the action of line--driven winds, and might collapse directly into a black hole without exploding as a supernova, producing more massive black holes than in a metal rich environment. The accretion--powered models do not exhaust the possible scenarios able to explain the nature of the ULXs. In a sample of $154$ ULXs observed with {\\sl Chandra}, \\citet{Swa04} argue that about~$20\\%$ of them may be well described by a thermal model, consistent with young supernova remnants strongly interacting with the surrounding environment. In this interpretation, the Eddington limit is clearly not an issue, and indeed the ULXs have X--ray luminosities comparable to those of the brightest supernovae ($L_X\\sim 10^{41}\\ergs$, \\citealp{Imm03}). The Cartwheel galaxy is a peculiar nursery of ULXs. This galaxy underwent a collision with a smaller galaxy about $10^8\\yr$ ago \\citep{Hig96, Map08}. This episode not only conferred the galaxy its peculiar shape, but also triggered a massive star formation in its ring. The Cartwheel harbours the largest number of ULXs than any other galaxy: 15 sources are more luminous than $10^{39}\\ergs$ \\citep{Wol04}, and at least some of them are known to be variable \\citep{Cri09}. In this paper we address the properties of the ULX labelled N10 by \\cite{Wol04}. In the {\\sl Chandra} observation of~2001 it was the brightest ULX in the Cartwheel (and also among the brightest known ULXs), but a couple of subsequent {\\sl XMM--Newton} observations taken in 2004 and 2005 showed that its luminosity is fading \\citep{Wol06}. Two additional {\\sl Chandra} observations were performed in~2008 to follow the variability pattern of this ULX. We present here the {\\sl Chandra} observations taken in 2008 and compare them with that of 2001. The {\\sl ACIS} angular resolution is instrumental to reduce to a minimum the possible contamination of N10 from the surrounding diffuse gas and the neighbour sources crowding the Cartwheel ring. In our analysis we do not use the {\\sl XMM--Newton} observations, since the analysis of the variability with {\\sl XMM--Newton} data was already published by \\citet{Wol06}. In addition, the use of {\\sl XMM--Newton} data would require specific modelling of the contamination of the neighbouring sources and the Cartwheel's ring due to the relatively large PSF of {\\sl XMM--Newton}, incrementing the uncertainties in the derived spectral parameters. Our main goal is to compare the spectral parameters of N10 to the theoretical models put forward to explain the ULXs. In particular, we shall discuss several accretion models and the supernova model. Although N10 is a bright source, its large distance ($D=122\\Mpc$, \\citealp{Wol04} and references therein) limits the number of collected photons, preventing us from rejecting or confirming a spectral model on purely statistical grounds. For this reason, we shall assess the likelihood of any model more in the light of its own self--consistency than of its statistical evidence. The outline of the paper is the following. In section~\\ref{s-prep} we describe the {\\sl Chandra} data and their preparation. In Section~\\ref{s-plaw} we present a simple spectral model of N10, which will be a thread for a more detailed analysis. The accretion and supernova models are presented in section~\\ref{s-accretion} and~\\ref{s-supernova}, respectively. Finally, we discuss and summarise our results in section~\\ref{s-sum}. ", "conclusions": "\\label{s-sum} In this paper we considered the spectral modelling of the ULX N10, located in the Cartwheel galaxy in order to assess the nature of this source. Although the source is intrinsically very bright, the spectra have few counts on account of the large distance ($122\\Mpc$). For this reason, we are unable to reject (or confirm) a model on purely statistical grounds. We first discussed several accretion models (multicolour disc around a Schwarzschild or Kerr black hole, slim disc, hyperaccretion). All models indicate a black hole of~$\\sim 100\\msun$, at the high end of the mass distribution of a black hole generated by stellar evolution, possibly in a metal poor environment. This result is consistent with the conclusions of~\\citet{Map09} and~\\citet{Zam09} that stellar evolution in a metal depleted environment may produce black holes of $30-90\\msun$. The theoretical studies of the evolution of very massive stars have been rekindled by the discovery of the ULXs, with the result that black holes of $\\sim 100\\msun$ are more common than earlier works suggested. Isolated solar--abundance stars may have masses up to $150\\msun$, but in dense environments they may coalesce to form stars as massive as $1000\\msun$ \\citep{Yun08}. Recent studies on the mass loss from hot, massive stars (\\citealp{Pul08}, and references therein) have shown that the occurrence of weak and/or clumpy winds may reduce the mass loss up by a factor $10$, thus allowing the star to retain a higher fraction of its original mass at the end of its life. Also lower metallicities entail weaker wind mass losses \\citep{Heg03}. Zero metallicity (Population~III) stars may be significantly more massive than solar--abundance stars \\citep{Ohk06}. All these factors determine the mass of the final black hole, that may range from $\\sim 70\\msun$ for a solar--abundance progenitor \\citep{Yun08}, up to $500\\msun$ for low metallicity stars \\citep{Ohk06}. Observations indicate that the low metallicity scenario is more likely for N10: the Cartwheel galaxy is metal--poor, and although the available measurements \\citep{Fos77} are not at the exact location of N10, they should nevertheless be representative of the average abundance in the ring. We therefore conclude that the interpretation of N10 as a black hole of $\\sim 100\\msun$ suggested by the X--ray data is consistent with the theoretical predictions of the evolution of massive stars in a dense environment. In summary, if ordinary stellar evolution is the correct scenario to form the black hole, N10 could be an extreme High Mass X--ray Binary (HMXB). The mass accretion rate $\\dot M$ on the black hole may be inferred from the luminosity $L$ via Equation~\\eqref{e-eff}, and is of the order of ${\\dot M}\\simeq10^{-6}~\\msun\\yr^{-1}$. This value is comparable with the loss rate of massive (i.e., few tens solar masses, \\citealp{Fra02}) donor stars on a thermal time scale. The accretion flow from the donor star to the BH is most likely to occur through a Roche lobe overflow; the BH capture of a wind blown by the companion is less favoured, since the small capture radius would require an unlikely strong mass loss from the companion. Can the observed decay of the luminosity of the source be explained in the framework of the accretion model? The longest characteristic time of an accretion disc is the so--called ``viscous time'' $t_{\\rm visc}$, i.e. the characteristic time taken by the disc to adapt to new conditions. For a standard $\\alpha$-disc orbiting a black hole of $100\\msun$ with an accretion rate ${\\dot M}\\simeq10^{-6}~\\msun\\yr^{-1}$, $t_{\\rm visc}\\simeq 10^2-10^3\\s$ \\citep{Fra02}, much shorter than the time scale of the variability of N10 (few years). Therefore, the observed variability cannot be due to a sort of disc instability, which would affect the disc and the luminosity of the source over a time scale $t_{\\rm visc}$. The simplest alternative is that the variability is due to a change of the mass transfer rate from the donor star. This may be due to an intrinsic decay of the mass loss from the donor, but also to other effects occurring at the inner Lagrangian point, since the instantaneous mass flow here is very sensitive to the relative sizes of the donor star and the Roche lobe \\citep{Fra02}. New observations of this interesting source would help to settle the question. \\bigskip We also explored the possibility that N10 is a young supernova strongly interacting with the circumstellar medium. Both the interpretation of this model and the best--fitting spectral parameters are consistent with this view. A new flux brightening of the source in the future would rule out this hypothesis." }, "1003/1003.1923_arXiv.txt": { "abstract": "{ In July 2010 the ESA spacecraft Rosetta will fly by the main belt asteroid 21 Lutetia. Several observations of this asteroid have been performed so far, but its surface composition and nature are still a matter of debate. For a long time Lutetia was supposed to have a metallic nature due to its high IRAS albedo. Later on it has been suggested that the asteroid has a surface composition similar to primitive carbonaceous chondrite meteorites, while further observations proposed a possible genetic link with more evolved enstatite chondrite meteorites. } {We performed visible spectroscopic observations of 21 Lutetia in November 2008 at the Telescopio Nazionale Galileo (TNG, La Palma, Spain) to make a decisive contribution to solving the conundrum of its nature. } {Thirteen visible spectra were acquired at different rotational phases and subsequently analyzed. } {We confirm a narrow spectral feature at about 0.47-0.48 $\\mu$m which was already found by Lazzarin et al. (2009) in the spectra of Lutetia. We also confirm an earlier find of Lazzarin et al. (2004), who detected a spectral feature at about 0.6 $\\mu$m in one of their Lutetia's spectra. More remarkable is the difference of our spectra though, which exhibit different spectral slopes between 0.6 and 0.75 $\\mu$m and, in particular, we found that up to 20\\% of the Lutetia surface could have flatter spectra. } {We detected a variation of the spectral slopes at different rotational phases that could be interpreted as possibly due to differences in the chemical/mineralogical composition as well as to inhomogeneities of the structure of the Lutetia's surface (e.g., to craters or albedo spots) in the southern hemisphere.} ", "introduction": "Rosetta is the ESA cornerstone mission devoted to the study of minor bodies of the solar system. The main target is comet 67P/Churyumov-Gerasimenko that will be reached in 2014, after three Earth and one Mars gravity assisted swing-bys. During its journey the mission investigates also two main belt asteroids, 2867 Steins (fly-by in September 2008) and 21 Lutetia (fly-by in July 2010). The asteroid 21 Lutetia was discovered in 1852 by H.~Goldschmidt at the Paris Observatory. Zappal\\`a et al. (1984) measured a rotational period of P=8.17$\\pm$0.01 hours, a value later refined by Torppa et al. (2003), who found P=8.165455~hours and also computed the pole coordinates obtaining a model with axis ratios a/b=1.2 and b/c=1.2. Although several observations are available since a couple of decades, the nature of this asteroid is still controversial. Radiometric measurements gave albedo values included in the range 0.19--0.22 (IRAS, M\\\"{u}ller et al. 2006, Lamy et al. 2008). More recently, Carvano et al. (2008) obtained a geometric albedo of 0.129, significantly lower than all the previous estimations, and explained the wide range of computed albedo values as the evidence of inhomogeneities on the surface of Lutetia (e.g., one or more large craters on the northern hemisphere). Due to the first estimation of its albedo by IRAS, Lutetia was supposed to have a metallic nature (Barucci et al. 1987; Tholen 1989). Later on it has been suggested that the asteroid has a surface composition similar to primitive carbonaceous chondrite meteorites (Howell et al. 1994; Burbine \\& Binzel 2002; Lazzarin et al. 2004, 2009; Birlan et al. 2004; Barucci et al. 2005, 2008), while Vernazza et al. (2009) proposed a possible genetic link with more evolved enstatite chondrite meteorites. On the basis of the observational evidence, Lutetia appears to be an atypical puzzling asteroid, whose nature is still far from fully understood. In order to enhance our knowledge of this unusual object, in November 2008 we performed visible spectroscopic observations of Lutetia to investigate its surface and check if some inhomogeneities are present. ", "conclusions": "An observational campaign of Lutetia was carried out on November 2008. Visible spectroscopy was performed for this main belt asteroid and thirteen spectra were acquired at different rotational phases. All of them exhibit absorption features centered at about 0.47--0.48 $\\mu$m and around 0.6 $\\mu$m. The spectral slope shows a variation through the rotational phase, suggesting that some inhomogeneities must be present in a portion of up to 20\\% of the observed surface of Lutetia in its southern hemisphere. These differences in the acquired spectra can be due to inhomogeneities in the chemical/mineralogical composition, or to albedo spots or craters exposing regions with different albedo/age. These data are useful in the assessment of the physical nature of this object, and will constitute a fundamental basis for the calibration, analysis and interpretation of the data acquired by the instruments onboard the Rosetta spacecraft." }, "1003/1003.4501_arXiv.txt": { "abstract": "We report on submillimetre bolometer observations of the isolated neutron star RX\\,J1856.5--3754 using the LABOCA bolometer array on the Atacama Pathfinder Experiment (APEX) Telescope. No cold dust continuum emission peak at the position of RX\\,J1856.5--3754 was detected. The $3 \\sigma$ flux density upper limit of 5\\,mJy translates into a cold dust mass limit of a few earth masses. We use the new submillimetre limit, together with a previously obtained H-band limit, to constrain the presence of a gaseous, circumpulsar disc. Adopting a simple irradiated-disc model, we obtain a mass accretion limit of $ \\dot{M} \\lsim 10^{14}$ g s$^{-1}$, and a maximum outer disc radius of $\\sim 10^{14}$\\,cm. By examining the projected proper motion of RX\\,J1856.5--3754, we speculate about a possible encounter of the neutron star with a dense fragment of the CrA molecular cloud a few thousand years ago. ", "introduction": "Since the discovery of planets around the pulsar PSR~1257+12 by \\citet{WF92}, dusty discs around pulsars have become interesting to observers, who have been trying to detect them in the infrared (IR) or at (sub-)millimeter wavelengths. Most comprehensive are the searches by \\citet{loehmer2004} and \\citet{Greaves2000}. Previous searches had concentrated mainly on recycled, thus formerly accreting, radio pulsars as the most likely objects to be surrounded by dusty discs. Neither of the above mentioned surveys detected dust emission around the planet-hosting PSR~1257+12, nor did the recent search with $Spitzer$ at 24\\,$\\mu$m and at 70\\,$\\mu$m by \\citet{Bryden2006}. However, dusty discs could also be present around non-recycled neutron stars (NSs). Following the supernova explosion, which creates the NS, some of the explosion ejecta may fail to escape and remain bound - forming a fallback disc. Such fallback discs are a general prediction of current supernova models \\citep{Michel1981,Chevalier1989}, and have been invoked by a number of authors to explain a variety of phenomena related to NSs (e.g., \\citealt{Chatterjee2000,Alpar2001,Menou2001,Blackman2004}). Recently, \\citet{Wang2006} reported the discovery of mid-infrared emission from a cool disc around the Anomalous X-ray pulsar AXP 4U~0142+61. \\citet{Wang2006} interpreted their detection as a passive disc, while \\citet{Ertan2007} argued that it could also originate from an actively accreting disc. To date, AXP 4U~0142+61 is the only isolated neutron star for which a fallback disc is believed to have been detected. Such discs appear to be rare. According to \\citet{Eksi2005} fallback disc are rare because they are likely to be disrupted when the newly born NS spins rapidly through the propeller stage, at which in-flowing matter, instead of being accreted, would be expelled. The fallback discs can survive if the initial NS spin is slow enough ($\\geq 40$~ms at a magnetic moment of $\\mu=10^{30}$~G cm$^{3}$). \\citet{Jones2007} studied the effect of pulsar wind induced ablation of fallback discs. He concluded that long-lived discs could be present in many pulsars without exceeding published limits on IR luminosity.\\\\ In the following, we report on submillimetre observations of RX\\,J1856.5--3754, which is the brightest and closest member of a class of NSs neglected so far in the search for circumstellar discs, the X-ray thermal isolated neutron stars. They are peculiar because they show pure thermal soft X-ray spectra without any (confirmed) non-thermal emission, especially no confirmed radio emission; for reviews, see \\citet{Kaplan2008A} or \\citet{Haberl2007}. These objects have periods in the range of 4 to 12~s. Thus, they are much slower than the bulk of radio pulsars. Their X-ray pulse periods and period derivatives are similar to those of the AXPs and soft gamma-ray repeaters, SGRs, (see, e.g., \\citealt{Kaplan2009}). This has led to discussions about whether they may be related to those objects. For example, \\citet{Alpar2001, Alpar2007} suggested that the X-ray thermal isolated NSs may simply have accretion discs with smaller masses than those of the AXPs. We note that it is currently not clear whether AXPs host accretion discs. The prevailing model for AXPs and SGRs is the magnetar model -- isolated, young neutron stars with exceptionally high ($\\approx 10^{14}$\\,G) surface magnetic fields \\citep{Duncan1992}.\\\\ RX\\,J1856.5--3754 has been observed at radio wavelengths (to date without detection), near infrared (without detection), optical and X-ray wavelengths. The striking absence of deviations from a pure blackbody in its X-ray spectrum has led to lively discussions on the nature of the object and of its X-ray emission, including highly magnetized atmospheres (e.g., \\citealt{Ho2007a}), condensed surfaces (e.g., \\citealt{Burwitz2003}) and quark stars (e.g., \\citealt{Drake2002}). The most recent, preliminary, parallax by \\citet{Kaplan2007a} is $167^{+18}_{-15}$\\,pc and marks RX\\,J1856.5--3754 as one of the closest NSs currently known. \\citet{Tiengo2007} discovered the $\\approx 7$\\,s pulsations of RX\\,J1856.5--3754, which has an extremely small pulsed fraction in the 0.15-1.2 keV range of only 1.2\\,\\%. A timing study by \\citet{Kerkwijk2008} inferred a magnetic field of $1.5 \\times 10^{13}$\\,G assuming spin-down by dipole radiation. Interestingly, the spin-down luminosity is not high enough to explain the H{$\\alpha$} nebula around RX\\,J1856.5--3754, discovered by \\citet{Kerkwijk2001}. The nature of the H{$\\alpha$} nebula remains unclear to date (see, e.g., \\citealt{Kerkwijk2008}). These authors also derived a characteristic age of $~4$\\,Myrs, which is much older than the kinematic age, $~0.5$\\,Myrs, obtained assuming an origin in the Upper Scorpius OB association (e.g., \\citealt{Walter2002}). At such ages, any fallback disc around the source would be expected to be rather cool.\\\\ Within distance error bars, $117 \\pm 12$\\,pc by \\citet{Walter2002} to $167^{+18}_{-15}$\\,pc by \\citet{Kaplan2007a}, RX\\,J1856.5--3754 appears to be in the outskirts of the Corona Australis star-forming cloud, whose distance is relatively well known: $129 \\pm 11$\\,pc from the orbit solution of the double-lined spectroscopic binary TY CrA by \\citet{Casey1998}; for a detailed discussion we refer to Sect.\\,1.3 in \\citet{Neuh2008}. Bondi-Hoyle accretion \\citep{Bondi1944,Bondi1952} from the interstellar medium (ISM) is unlikely to be a major contributor to the observed X-ray luminosity today, given the high velocity of this neutron star \\citep{Kerkwijk2001}. Furthermore, it is likely that isolated NSs accrete at sub-Bondi rates \\citep{Perna2003}. The accretion rate scales inversely with the magnetic moment as $~{\\mu}^{-2.1}$ according to \\citet{Toropina2006, Romanova2003}. RX\\,J1856.5--3754 has a high magnetic field, $1.5 \\times 10^{13}$\\,G \\citep{Kerkwijk2008}, reducing the possibility of accretion even further. However, as \\citet{Drake2003} noted, RX\\,J1856.5--3754 might have passed more dense regions than it now resides in. Indeed the IRAS 100\\,$\\mu$m image shows some cloud fragments projected along the past trajectory of RX\\,J1856.5--3754 (see Fig.~\\ref{fig:IRAS}), opening the possibility that the neutron star collected material in the past.\\\\ Here, we investigate the surroundings of RX\\,J1856.5--3754 for cold dust as one might expect either in an old fallback disc or in a dense gas disc. The gas content of a fallback disc is not well known. The formation from supernova ejecta would be expected to produce a disc mainly consisting of heavy elements. \\citet{Currie2007} showed that gas from viscous, circumpulsar discs typically depletes on very short, $\\approx 10^5$ yr, timescales due to high temperatures during the early evolutionary stages, subsequent rapid spread of the disc and cooling allowing for the heavy elements to condense onto grains. In addition, \\citet{Phillips1994} have shown that for a pulsar moving with $\\approx 100$\\,km\\,s$^{-1}$ through the ISM, small dust grains with radii of $\\leq 0.1$\\,$\\mu$m are spiraled in from a circumpulsar disc to the neutron star on a time scale of order $~ 10^6$ years. Thus, a circumpulsar disc at the age of RX\\,J1856.5--3754 is likely to be similar to debris discs, which are dust-rich and dominated by larger grains (see, e.g., \\citealt{Krivov2009}) as opposed to protostellar discs, which are gas-rich. If, on the other hand, RX\\,J1856.5--3754 has collected new, gas-rich material from the outskirts of the Corona Australis star-forming cloud, the disc composition could more closely resemble a protostellar disc. Given the uncertainties in the gas content and composition of a possible disc around RX\\,J1856.5--3754, we will consider here both types of discs -- dense, gas-rich protostellar-type discs, and gas-poor, dust-rich debris-type discs. \\begin{figure} \\includegraphics[width=8.5cm]{RXJ1856_IRIS100_191009a.eps} \\caption{The IRAS/IRIS $100\\mu$m band of the neighbourhood of RX\\,J1856.5--3754. Strong infrared emission, signaling the presence of dust, is shown by darker areas. The peak dust emission outlines the CrA star forming region. The position of RX\\,J1856.5--3754, as well the direction of its proper motion, is indicated. The dashed white line shows the rough projection of the proper motion into the past. About 3000 years ago the neutron star might have crossed a region of the denser dust filament -- depending on the actual distances of the filament and RX\\,J1856.5--3754. The boxes 1 and 2 cover the same area, see section\\,\\ref{past} for discussion. } \\label{fig:IRAS} \\end{figure} ", "conclusions": "We investigated RX\\,J1856.5--3754 for $870$\\,$\\mu$m continuum emission which would be indicative of a cold dusty disc around the neutron star. The derived deep flux density limit translates into a dust mass limit of few earth masses. Applying the irradiated (gas-rich) accretion disc model by \\citet{Perna2000}, together with further observational constraints, we obtained a mass accretion limit of $ \\dot{M} \\lsim 10^{14}$ g s$^{-1}$, and a constraint on the outer disc radius, $R_{out}$, to be smaller than $10^{14}$\\,cm or 7\\,AU. Looking at the projected proper motion of RX\\,J1856.5--3754, we note that the neutron star might have passed a dense fragment of the CrA molecular cloud a few thousand years ago which could have affected a potential circumstellar disc, as well as have enabled a brief history of accretion form the ISM." }, "1003/1003.5754_arXiv.txt": { "abstract": "We discuss the properties of homogeneous and isotropic flat cosmologies in which the present accelerating stage is powered only by the gravitationally induced creation of cold dark matter (CCDM) particles ($\\Omega_{m}=1$). For some matter creation rates proposed in the literature, we show that the main cosmological functions such as the scale factor of the universe, the Hubble expansion rate, the growth factor and the cluster formation rate are analytically defined. The best CCDM scenario has only one free parameter and our joint analysis involving BAO + CMB + SNe Ia data yields ${\\tilde{\\Omega}}_{m}= 0.28\\pm 0.01$ ($1\\sigma$) where $\\tilde{{\\Omega}}_{m}$ is the observed matter density parameter. In particular, this implies that the model has no dark energy but the part of the matter that is effectively clustering is in good agreement with the latest determinations from large scale structure. The growth of perturbation and the formation of galaxy clusters in such scenarios are also investigated. Despite the fact that both scenarios may share the same Hubble expansion, we find that matter creation cosmologies predict stronger small scale dynamics which implies a faster growth rate of perturbations with respect to the usual $\\Lambda$CDM cosmology. Such results point to the possibility of a crucial observational test confronting CCDM with $\\Lambda$CDM scenarios trough a more detailed analysis involving CMB, weak lensing, as well as the large scale structure. ", "introduction": "The analysis of high quality cosmological data (e.g. supernovae type Ia, CMB, galaxy clustering, etc) have converged towards a cosmic expansion history that involves a spatially flat geometry and some sort of dark energy in order to explain the recent accelerating expansion of the universe \\cite{Spergel07,essence,komatsu08,Teg04,Eis05,Kowal08,Hic09}. The simplest dark energy candidate corresponds to a cosmological constant, $\\Lambda$ (see \\cite{Weinberg89,Peebles03,Pad03} for reviews). In the standard concordance $\\Lambda$CDM model, the overall cosmic fluid contains baryons, cold dark matter plus a vacuum energy that fits accurately the current observational data and thus it provides an excellent scenario to describe the observed universe. On the other hand, the concordance model suffers from, among others \\cite{Peri08}, two fundamental problems: (a) {\\it The fine tuning problem} i.e., the fact that the observed value of the vacuum energy density ($\\rho_{\\Lambda}=\\Lambda c^{2}/8\\pi G\\simeq 10^{-47}\\,GeV^4$) is more than 120 orders of magnitude below the natural value estimated using quantum field theory \\cite{Weinberg89}, and (b) {\\it the coincidence problem} \\cite{coincidence} i.e., the fact that the matter energy density and the vacuum energy density are of the same order just prior to the present epoch. Such problems have inspired many authors to propose alternative candidates to dark energy such as $\\Lambda(t)$ cosmologies, quintessence, $k-$essence, vector fields, phantom, tachyons, Chaplygin gas and the list goes on (see \\cite{Reviews,Ratra88,Oze87,Free187,Wetterich:1994bg,Caldwell98,Brax:1999gp, KAM,fein02,Caldwell,chime04,Brookfield:2005td,Bauer05,Grande06, Boehmer:2007qa,Bas09b} and references therein). Nowadays, the nature of the dark energy is considered one of the most fundamental and difficult problems in the interface uniting Astronomy, Particle Physics and Cosmology. However, there are other possibilities. For instance, one may consider non-standard gravity theories where the present accelerating stage of the universe is driven only by cold dark matter, that is, with no appealing to the existence of dark energy. Such a reduction of the so-called dark sector is naturally obtained in the so-called $f(R)$ gravity theories \\cite{FR} (see, however, \\cite{Nat10}). Even in the framework of the standard general relativity theory, is also possible to reduce the dark sector by considering the presence of inhomogeneities \\cite{Inhom}, quartessence models \\cite{Quart}, as well as the gravitationally induced particle creation mechanism \\cite{Parker,BirrellD,Prigogine,LCW,LG92,ZP2,LGA96,LA99,Susmann94,ZSBP01,SLC02,freaza02,Makler07,LSS08,SSL09,LJO09}. In what follows we focus our attention to the last approach by considering the gravitational creation of cold dark matter in the expanding Universe. The basic microscopic description for gravitational particle production in an expanding universe has been investigated in the literature by many authors starting with Schr\\\"odinger \\cite{SCHO39}. In the late 1960s, independently of Schr\\\"odinger's work, this issue was investigated by L. Parker and others \\cite{Parker,BirrellD} by considering the Bogoliubov mode-mixing technique in the context of quantum field theory in curved spacetime. The basic physical effect is that a classical non-stationary background influences bosonic or fermionic quantum fields in such a way that their masses become time-dependent (see, e.g., \\cite{Parker1} for more recent works). For a real scalar field in a flat Friedmann-Lema{\\^i}tre-Robertson-Walker (FLRW) geometry described in conformal coordinates, for example, the key result is that the effective mass reads \\cite{MW07}, $m^2_{eff}(\\eta)\\equiv m^2a^2-{a''/a}$, where m is the ``Minkowskian\" constant mass of the scalar field, $a(\\eta)$ is the scale factor, and the derivatives are computed with respect to the conformal time. This time dependent mass accounts for the interaction between the scalar field and the geometry of the Universe. When the field is quantized, this leads to particle creation, with the energy for newly created particles being supplied by the classical, time-varying gravitational background. Macroscopically, the first self-consistent formulation for matter creation was put forward by Prigogine and coworkers \\cite{Prigogine} and somewhat clarified by Calv\\~{a}o, Lima and Waga \\cite{LCW} (see also \\cite{LG92}). It was shown that matter creation, at the expenses of the gravitational field is also macroscopically described by a negative pressure, and, potentially, can accelerate the Universe. Within this framework, various interesting features of cosmologies with particle creation have been discussed in \\cite{ZP2,LGA96,LA99,Susmann94,ZSBP01,SLC02,freaza02} (see also \\cite{Makler07} for recent studies on this subject). More recently, the corresponding effects on the global dynamics of the particle creation regime has been investigated extensively by a number of authors (see \\cite{LSS08,SSL09,LJO09} and references therein). In particular, it was also found that a subclass of such models depends only of one free parameter and that the evolution of the scale factor is exactly the same of $\\Lambda$CDM models \\cite{LJO09}. Naturally, in order to know if creation of cold dark matter (CCDM) models provide a realistic description of the observed Universe, its viability need to be tested by discussing all the constraints imposed from current observations both for background and perturbative levels (structure formation). In this context, the basic aim of the present work is twofold. First, we place constraints on the main parameters of CCDM cosmologies by performing a joint likelihood analysis involving the shift parameter of the Cosmic Microwave Background (CMB) \\cite{komatsu08} and the observed Baryonic Acoustic Oscillations (BAO) \\cite{Eis05}, and the latest SN Ia data (Constitution) \\cite{Hic09}. Secondly, for a wide class of matter creation cosmologies, we also develop the linear approach for the density perturbations. Analytical solutions for the differential equation governing the evolution of the growth factor and some properties of the large scale structure (cluster formation) are also discussed and compared to the predictions of the $\\Lambda$CDM model. The paper is planned as follows. The basic theoretical elements of the problem are presented in section 2, where we introduce the basic FLRW cosmological equations for CCDM cosmologies. In section 3 and 4 we present the specific CCDM scenarios and derive the constraints on their parameters based on a statistical joint analysis involving the shift parameter of the Cosmic Microwave Background (CMB)\\,\\cite{komatsu08}, the Baryonic Acoustic Oscillations (BAO) \\cite{Eis05}, and the latest SN Ia data (Constitution) \\cite{Hic09}. In section 5 we study the evolution of linear perturbations, while in section 6 we present the corresponding theoretical predictions regarding the formation of the galaxy clusters with basis on the Press-Schecther formalism. Finally, in section 7 we draw our conclusions. Throughout the paper we consider $H_{0}=70.5$ km/sec/Mpc as given by the WMAP 5-years data \\cite{komatsu08}. ", "conclusions": "In this work, we have investigated (analytically and numerically) the overall dynamics of two cosmological models in which the dark matter creation process provides the late time accelerating phase of the cosmic expansion without the need of dark energy. Such scenarios termed here by LSS \\cite{LSS08} and LJO \\cite{LJO09} are phenomenologically characterized by two distinct creation rates, $\\Gamma$ (see section 3). In the first scenario (LSS), the creation rate is $\\Gamma = 3\\gamma H_{0} + 3\\beta H$ while in the second one it is proportional to the Hubble parameter, namely: $\\Gamma=3\\tilde{\\Omega}_{\\Lambda}(\\rho_{co}/{\\rho_{dm}})H$, where $\\rho_{co}=3H^{2}_{0}/8\\pi G$ is the present day value of the critical density. It should also be stressed that the phenomenological approach adopted here cannot determine the mass of the particles, as well as whether their nature is fermionic or bosonic. In order to access the mass of the particles, and, therefore, the nature of the CDM particles, it is necessary to determine the creation rate, $\\Gamma$, from quantum field theory in FLRW space time. In principle, such a treatment must somewhat incorporate a source of entropy since the matter creation process is truly an irreversible process. In particular, in the case of adiabatic matter creation considered here, both the entropy (S) and the number of particles (N) in a comoving volume increase but the specific entropy (S/N) remains constant \\cite{LCW,LG92,SLC02}. In this context, by using current observational data (BAOs, CMB shift parameter and SNIa) we have first performed a joint likelihood analysis in order to put tight constraints on the main cosmological parameters. Subsequently, trough a linear analysis we have also studied the growth factor of density perturbations for both classes of creation cold dark matter models, and, finally, by using the Press-Schecther formalism we have discussed how the cluster formation rates evolve in such scenarios. The result of our joint statistical analysis shown that the fit provided by the LSS model with $\\gamma=0.66$ and $\\beta=0$ turns out to be of poor quality because it is unable to adjust simultaneously the observational data at low and high redshift. This result confirm the conclusion by Steigman et al \\cite{SSL09} using only SNe Ia and the determination of $z_{eq}$, the redshift of the epoch of matter radiation equality. For the LSS scenario, we also find that the amplitude and the shape of the linear density contrast are significantly different with respect to those predicted by the LJO and $\\Lambda$ models (see Fig. 1). In particular, for $z\\le 0.8$ the matter fluctuation field of the LSS model practically decays thereby implying that the corresponding cosmic structures cannot be formed via gravitational instability (see also Figs. 2 and 3). On the other hand, the creation cold dark matter scenario (termed here as LJO) provides good quality fits of the cosmological parameters at all redshifts and it resembles the global dynamics of the concordance $\\Lambda$CDM cosmology by including only one free parameter. Despite, the latter the corresponding growth factor has the following evolution with respect to that of the usual $\\Lambda$ cosmology: (i) prior to the present epoch ($z<0.4$) the evolution of the LJO growth factor reaches almost a plateau, which means that the matter fluctuations are effectively frozen, (ii) between $0.4 \\le z < 1.6$, the growth factor in the LJO model is greater than that of the concordance $\\Lambda$ cosmology, while for $1.6 \\le z \\le 2.8$ the two growth factors have converged and (iii) for $z>2.8$ we find that $D(z) Re_M^C$, dynamo action results. Two different types of dynamo can be found, depending on the presence or absence of net flow helicity \\citep{MeFrPo1981}. With net helicity, magnetic energy grows at scales larger than the energy-containing scale of the fluid motions: large-scale dynamo (LSD) or mean-field dynamo. LSDs are often studied with mean-field theory; the production of large-scale magnetic energy is approximately the alpha-effect (see, e.g., \\citealt{Br2003}). Without net helicity, dynamo action is harder to achieve and magnetic every grows at scales smaller than the forcing scale. This latter defines small-scale dynamo (SSD) or fluctuation dynamo action. Near the solar surface, the convective time scale is much shorter than the rotation period, the effects of rotation can be neglected, and a flow with no net helicity results. Any surface dynamo will thus be a SSD. This is also suggested by observation of the small-scale magnetic field in the quiet Sun. In high resolution magnetograms we see mixed polarity fields on small scales, which is variously called the magnetic carpet or the salt-and-pepper pattern \\citep{TiSc1998,HaScTi2003}, consistent with the idea of SSD action. As turbulent convection can drive small-scale dynamo action (in numerical simulations of Boussinesq convection without rotation \\citealt{C99,CaEmWe2003}), observations and simulations together provide evidence of a SSD driven by turbulence at the solar surface (likely deeper as well). Several arguments can be put forward against a small-scale solar dynamo. Firstly, there exists the possibility that the small-scale field is solely produced by the shredding up of large-scale field by turbulence. However, observationally the amount of small-scale flux is not dependent on the solar cycle \\citep{HaScTi2003,TrBuShAsRa2004}. This might not be the case if the small-scale flux is the result of the shredding up of the field from the global dynamo. It might also then show some latitudinal dependence (among low latitudes). Assuming the existence of both SSD and shredding, small-scale dynamo is predicted to create small-scale magnetic field at the turbulent rate of stretching ($\\propto Re_M^{1/2}$, see, e.g., \\citealt{IsScCo+2007}) which is much faster either than large-scale field can be produced (at time scales associated with the kinetic-energy-containing length scales) or be shredded up by the turbulence \\citep{ScHaBr+2005}. Secondly, small-scale dynamo action may not be possible at the magnetic Prandtl number, $P_M$, of the solar plasma, $P_M \\equiv Re_M/Re \\approx 10^{-5}$: $Re_M^C$ sharply increases with decreasing $P_M$ (increasing $Re$) since eddies smaller than the characteristic scale of the magnetic field diffuse the field and inhibit dynamo action. The two asymptotic possibilities are $Re_M^C\\rightarrow\\,$const as $Re\\rightarrow\\infty$ (SSD at low $P_M$; \\citealt{RoKl1997,BoCa2004}) or $Re_M^C/Re=P_M^C\\rightarrow\\,$const as $Re\\rightarrow\\infty$ (no SSD at low $P_M$; \\citealt{ScHaBr+2005}). Numerical simulations of low $P_M$ SSDs have focused on the existence (or not) of a time-averaged mean flow. Such a flow exists for the Sun and many other astrophysical cases. Both with \\citep{PMM+05} and without a mean flow \\citep{IsScCo+2007}, a plateau in $Re_M^C$ was found. This suggests such dynamo action should be possible on the Sun. However, the $P_M$ of the sun is 3 orders of magnitude smaller that that accessible to present numerical computations; the observed plateaus may not represent the asymptotic behavior. Fortunately, a laboratory dynamo resulting from unconstrained turbulence in liquid sodium ($P_M\\approx10^{-5}$) has been demonstrated \\citep{MoBeBo+2007} establishing that a turbulent dynamo is possible at values of $P_M$ corresponding to the solar plasma. ", "conclusions": "A small-scale solar dynamo near the surface is likely. MURaM simulations show dynamo action and its properties are found to be in agreement with {\\sl Hinode} observations. The arguments against a solar small-scale dynamo, thus far, have failed. Whatever its source, the small-scale magnetic field is turbulent and fractal and this should be taken into consideration when interpreting observations: 1) the PDF of Stokes $V$ field estimates do not accurately represent the PDF of the $B_z$ and are consistent with a much greater prevalence of weak $B_z$; 2) the multi-fractal self-similar pattern of the quiet-Sun photospheric magnetic field extends down to the resolution limit, $200\\,$km. This constitutes observational evidence that the smallest scale of magnetic structuring in the photosphere is $20\\,$km or smaller. The power law also constrains the quiet-Sun true mean unsigned vertical flux density: the lower bound, $\\approx43\\,$G, is consistent with estimates based solely on numerical simulations ($\\sim55\\,$G). The order of magnitude disparity between Hanle and Zeeman-based estimates may be resolved by a proper consideration of the cancellation properties of the full vector field." }, "1003/1003.2174_arXiv.txt": { "abstract": "We report the detection of $158$ $\\mu$m [CII] fine-structure line emission from MIPS J142824.0+352619, a hyperluminous ($L_\\mathrm{IR}\\sim10^{13}$ $L_\\sun$) starburst galaxy at $z=1.3$. The line is bright, and corresponds to a fraction $L_\\mathrm{[CII]}/L_\\mathrm{FIR}\\approx2\\times10^{-3}$ of the far-IR (FIR) continuum. The [CII], CO, and FIR continuum emission may be modeled as arising from photodissociation regions (PDRs) that have a characteristic gas density of $n\\sim10^{4.2}$ cm$^{-3}$, and that are illuminated by a far-UV radiation field $\\sim$$10^{3.2}$ times more intense than the local interstellar radiation field. The mass in these PDRs accounts for approximately half of the molecular gas mass in this galaxy. The $L_\\mathrm{[CII]}/L_\\mathrm{FIR}$ ratio is higher than observed in local ULIRGs or in the few high-redshift QSOs detected in [CII], but the $L_\\mathrm{[CII]}/L_\\mathrm{FIR}$ and $L_\\mathrm{CO}/L_\\mathrm{FIR}$ ratios are similar to the values seen in nearby starburst galaxies. This suggests that MIPS J142824.0+352619 is a scaled-up version of a starburst nucleus, with the burst extended over several kiloparsecs. ", "introduction": "The $^2$P$_{3/2}$ $\\rightarrow$ $^2$P$_{1/2}$ transition of C$^+$ ($\\lambda=157.74$ $\\mu$m) is one of the brightest emission lines in star-forming galaxies, typically accounting for $0.1-1\\%$ of the far-IR (FIR) continuum~\\citep{Stacey1991,Malhotra2001}. Much of this emission arises from the warm and dense photodissociation regions (PDRs) that form on the UV-illuminated surfaces of molecular clouds. The [CII] transition is a primary PDR coolant, and is a sensitive probe of both the physical conditions of the photodissociated gas, and the intensity of the ambient stellar radiation field~\\citep{Hollenbach1999}. In ultraluminous infrared galaxies (ULIRGs), the $L_\\mathrm{[CII]}/L_\\mathrm{FIR}$ ratio is a factor of $\\sim$$7$ times lower than in less luminous systems~\\citep{Luhman2003}, for reasons that are not well understood. Large aperture (sub)millimeter telescopes have been used to search for [CII] emission from sources at $z>3$, and 3 FIR-luminous QSOs at $z=4.4-6.4$ have been detected thus far. SDSS J1148~\\citep{Maiolino2005} and BR 1202N~\\citep{Iono2006quasar} have low $L_\\mathrm{[CII]}/L_\\mathrm{FIR}$ ratios similar to or smaller than the mean local ULIRG value, while BRI 0952~\\citep{Maiolino2009} has a somewhat larger ratio falling at the lower end of the range seen in normal galaxies. We have initiated a search for [CII] emission from galaxies at $z=1-2$, concentrating on star-formation--dominated systems selected by their FIR or submillimeter continuum brightness. Here we report our first detection, from the $z=1.3$ hyperluminous starburst galaxy MIPS J142824.0+352619 (hereafter MIPS J1428). MIPS J1428 was identified in a Spitzer/MIPS blank field survey as a bright 160 $\\mu$m source with red optical/near-IR colors, and subsequent spectroscopic and photometric measurements established it as a hyperluminous ($L_\\mathrm{IR}[8-1000$ $\\mu$m$]=3.2\\times10^{13}$ $L_\\sun$) galaxy at $z=1.325$~\\citep{Borys2006}. There are several indications that this IR emission is powered by star formation, including the brightness of the PAH features and the nondetection in hard X-ray emission, and with a starburst origin the estimated $L_\\mathrm{IR}$ corresponds to a star formation rate of $\\sim$$5500$ $M_\\sun$ yr$^{-1}$~\\citep{Borys2006,Desai2006}. MIPS J1428 was detected in CO($2$$\\rightarrow$$1$) and CO($3$$\\rightarrow$$2$), and the line luminosities correspond to a large gas mass of $M_\\mathrm{H_2}\\sim10^{11}$ $M_{\\sun}$~\\citep{Iono2006CO}. However, a gravitational lens may amplify the flux by as much as a factor of 10~\\citep{Borys2006,Iono2006SMA,Iono2006CO}, making the actual star formation rate and gas mass correspondingly lower. Independent of the unknown lensing amplification, the large value of $L_\\mathrm{IR}/L\\arcmin_\\mathrm{CO}\\approx320$ $L_\\sun$ (K km s$^{-1}$)$^{-1}$~\\citep{Iono2006CO} indicates that MIPS J1428 is forming stars with high efficiency, similar to that seen in local ULIRGs and high-redshift submillimeter galaxies (SMGs)~\\citep{Solomon2005}. This is the first detection of [CII] in the $z=1-3$ epoch of peak star formation activity in the Universe, and the first detection from a high-redshift galaxy not associated with a QSO. ", "conclusions": "\\label{discussion} \\subsection{[CII], CO, and the FIR Continuum: A Comparison to Other Galaxies} \\label{starburst} To place our [CII] measurement of MIPS J1428 in context, we compare the [CII], CO, and FIR continuum emission from this source with that from other galaxies and Galactic sources. In Figure~\\ref{fig3} we plot the $L_\\mathrm{[CII]}/L_\\mathrm{FIR}$ and $L_{\\mathrm{CO}(1\\rightarrow0)}/L_\\mathrm{FIR}$ ratios for samples of Galactic star-forming regions and nearby gas-rich galaxies~\\citep{Stacey1991}, normal galaxies~\\citep{Malhotra2001}, local ULIRGs~\\citep{Luhman2003}, and MIPS J1428 and the 3 other high-redshift sources detected in [CII]~\\citep{Maiolino2005,Iono2006quasar,Maiolino2009}. We estimate the CO($1$$\\rightarrow$$0$) luminosity of MIPS J1428 assuming an $L\\arcmin_{\\mathrm{CO}(2\\rightarrow1)}/L\\arcmin_{\\mathrm{CO}(1\\rightarrow0)}=0.9$ ratio intermediate between the ratios observed in local ULIRGs~\\citep[$\\approx$$0.7$;][]{Downes1998} and in starburst nuclei~\\citep[$\\approx$$1.1$;][]{Harrison1999,Weiss2005M82}. For reference, we overplot the PDR model curves from~\\citet{Kaufman1999}. A vector attached to the MIPS J1428 data point indicates how the measured ratios were adjusted for the PDR analysis in section~\\ref{pdranalysis} (to account for non-PDR contributions to the [CII] emission, and optical depth effects in the CO transititions), and we further note that MIPS J1428 is placed in Figure~\\ref{fig3} assuming an $L\\arcmin_{\\mathrm{CO}(2\\rightarrow1)}/L\\arcmin_{\\mathrm{CO}(1\\rightarrow0)}$ ratio $\\approx$$20\\%$ lower than our best fit PDR solution. We compare MIPS J1428 with each of the populations shown in Figure~\\ref{fig3} in turn. \\begin{figure*} \\epsscale{1.1} \\plotone{f3.eps} \\caption{$L_\\mathrm{[CII]}/L_\\mathrm{FIR}$ vs. $L_{\\mathrm{CO}(1\\rightarrow0)}/L_\\mathrm{FIR}$ for Galactic star-forming regions (\\textit{crosses}), starburst nuclei (\\textit{filled squares}), non-starburst nuclei (\\textit{open squares}), normal galaxies (\\textit{triangles}), local ULIRGs (\\textit{circles}), and high-redshift sources (\\textit{asterisks}). CO($1$$\\rightarrow$$0$) luminosities of the normal galaxies and local ULIRGs are taken from the literature and from J. Graci{\\'a}-Carpio et al. (2010, in preparation). CO($1$$\\rightarrow$$0$) luminosities for MIPS J1428 and 2 of the other high-redshift sources are estimated from the measurements of higher-$J$ lines as described in the text. Starburst galaxies and Galactic star-forming regions are characterized by $L_\\mathrm{[CII]}/L_{\\mathrm{CO}(1\\rightarrow0)}\\approx4100$ (\\textit{dashed}). Overplotted are the PDR model calculations for gas density ($n$) and FUV field strength ($G_0$) from~\\citet{Kaufman1999}, and a vector indicates how the MIPS J1428 data point was shifted for the PDR analysis in section~\\ref{pdranalysis}.\\label{fig3}} \\end{figure*} \\citet{Stacey1991} study the [CII] emission from the central regions of a sample of nearby galaxies, and find that the $L_\\mathrm{[CII]}/L_{\\mathrm{CO}(1\\rightarrow0)}$ and $L_\\mathrm{[CII]}/L_\\mathrm{FIR}$ ratios are sensitive tracers of the star formation activity in these sources. Starburst galaxies are well characterized by the same $L_\\mathrm{[CII]}/L_{\\mathrm{CO}(1\\rightarrow0)}=4100$ ratio seen in Galactic star-forming regions, which is a factor of $\\sim$$3$ times larger than in non-starburst galaxies and Galactic molecular clouds (Fig.~\\ref{fig3}). % However, the $L_\\mathrm{[CII]}/L_\\mathrm{FIR}$ ratios in starburst galaxies are higher than those observed in the most intense Galactic OB star formation regions. These results imply that while much of the molecular gas in a starburst nucleus is photodissociated by the FUV radiation from young stars, the very intense fields found in the most extreme Galactic star-forming regions are not providing the bulk of the luminosity. This picture is supported by PDR modeling, which indicates that the FUV fields producing the emission in the centers of starburst galaxies are intermediate in intensity between those producing the emission in non-starburst galaxies, and those found in Galactic OB star formation regions (Fig.~\\ref{fig3}). As the data point for MIPS J1428 falls near the high-$G_0$ end of the distribution of starburst points in Figure~\\ref{fig3}, we conclude that in MIPS J1428, like in nearby starburst nuclei, the bulk of the molecular gas is exposed to moderately-intense FUV radiation produced by young stars. \\citet{Malhotra2001} observe [CII] and other FIR fine-structure line emission from a sample of 60 normal galaxies with varying levels of star formation activity, which are all sufficiently distant that the FIR emission is contained within the $\\approx$$70\\arcsec$ \\textit{ISO}-LWS beam. These authors find that the FIR continuum and much of the [CII] emission arises from PDRs, which in many of their galaxies are characterized by similarly elevated values of $G_0$ as deduced here for MIPS J1428. However, the~\\citet{Malhotra2001} sources have lower $L_\\mathrm{[CII]}/L_\\mathrm{CO}$ ratios than MIPS J1428 or its starburst analogs, indicating that most of the CO emission is produced by molecular gas residing in less active star-forming regions than those that produce the fine-structure line emission. This difference helps reinforce our conclusion that in MIPS J1428 it is the entire galaxy, rather than just an active central region, that is host to vigorous star formation. The median $L_\\mathrm{[CII]}/L_\\mathrm{FIR}$ ratio in local ULIRGs is a factor of $\\sim$$7$ times lower than in normal star-forming galaxies~\\citep{Luhman2003}, and a factor of $\\sim$$4$ times lower than in MIPS J1428. One possible explanation for this low ratio is that the [CII] emission is produced in dense PDRs illuminated by intense FUV radiation~\\citep[e.g.,][]{Papadopoulos2007}, as indicated by the PDR model overlays in Figure~\\ref{fig3}. Alternatively, the [CII] emission may originate in PDRs with more modest values of $n$ and $G_0$, but an additional source of FIR continuum emission may lower the global $L_\\mathrm{[CII]}/L_\\mathrm{FIR}$ ratio~\\citep{Luhman2003,Abel2009}. Regardless of the origins of the weak [CII] emission in local ULIRGs, it is clear that these galaxies provide poor templates for MIPS J1428. We also include the 3 high-redshift QSOs detected in [CII] in Figure~\\ref{fig3}. SDSS J1148 has been detected in CO($3$$\\rightarrow$$2$) and BRI 0952 in CO($5$$\\rightarrow$$4$), and we use the measured $L\\arcmin_{\\mathrm{CO}(3\\rightarrow2)}/L\\arcmin_{\\mathrm{CO}(1\\rightarrow0)}=0.72$ and modeled $L\\arcmin_{\\mathrm{CO}(5\\rightarrow4)}/L\\arcmin_{\\mathrm{CO}(1\\rightarrow0)}=0.58$ ratios for Mrk 231~\\citep{Papadopoulos2007} to estimate the CO($1$$\\rightarrow$$0$) luminosities for these sources. To account for the uncertainties in the CO excitation of high-redshift systems~\\citep[e.g.,][]{Hainline2006}, we show these data points with errors bars corresponding to 0.2 dex and 0.3 dex, respectively. The low $L_\\mathrm{[CII]}/L_\\mathrm{FIR}$ and $L_\\mathrm{CO}/L_\\mathrm{FIR}$ ratios in SDSS J1148 and BR 1202N indicate similar physical conditions as in local ULIRGs. Indeed,~\\citet{Maiolino2005} found that the [CII], CO, and FIR continuum emission of SDSS J1148 is consistent with a high-density, high-$G_0$ PDR, and~\\citet{Walter2009} showed that the star formation in this system is confined to a compact region with a high surface brightness comparable to that at the center of Arp 220. However, BRI 0952 has a larger $L_\\mathrm{[CII]}/L_\\mathrm{FIR}$ ratio and a $L_\\mathrm{[CII]}/L_\\mathrm{CO}$ ratio similar to that in MIPS J1428. We suggest that while SDSS J1148 and BR 1202N are likely dissimilar to MIPS J1428, the connection drawn here between MIPS J1428 and local starburst nuclei may also provide insight into the nature of the host galaxy of BRI 0952. Finally, it is important to note that while we assume that the $L_\\mathrm{[CII]}/L_\\mathrm{FIR}$ and $L_\\mathrm{CO}/L_\\mathrm{FIR}$ ratios in MIPS J1428 are unaffected by the potential gravitational lensing, we cannot rule out the possibility that these ratios are altered by differential magnification. For example, one could construct a model in which the intrinsic emission ratios of MIPS J1428 are similar to those in local ULIRGs, but where the $L_\\mathrm{[CII]}/L_\\mathrm{FIR}$ ratio is enhanced by the relatively large magnification of a localized HII region that falls near a caustic. However, in the absence of other evidence supporting such a scenario, we assume that the measured flux ratios are equal to the intrinsic luminosity ratios. \\subsection{MIPS J1428 as an Extended Starburst} The molecular gas in MIPS J1428 is exposed to similar UV radiation fields, and has similar excitation, as the gas in the central region of a typical starburst galaxy, and we suggest that MIPS J1428 may be modeled as a scaled-up version of a starburst nucleus. As an example, we consider scaling up the central region of the prototypical starburst galaxy M82 in such a manner as to match the much larger luminosity of MIPS J1428, while at the same time conserving the strength of the characteristic FUV radiation field that controls the PDR emission. We assume that in a starburst environment this characteristic field is determined by the global star formation density, rather than by the purely local properties of the interactions between individual star clusters and their natal molecular clouds. This interpretation was adopted for M82 and NGC 253 based on multiple-line PDR analysis~\\citep{Wolfire1990,Carral1994,Lord1996}, and is also motivated by the non-linear relation between $I_\\mathrm{[CII]}$ and $I_\\mathrm{FIR}$ observed in larger samples of galactic nuclei~\\citep{Crawford1985,Stacey1991}. For a starburst region in which young stellar clusters and molecular clouds are randomly distributed,~\\citet{Wolfire1990} modeled the relationship between the average FUV flux incident on the molecular gas ($G_0$), and the size ($D$) and total luminosity ($L_\\mathrm{IR}$) of the region. This relationship depends on the mean free path of a FUV photon ($\\lambda$), but in the limits of $\\lambda\\ll D$ and $\\lambda\\gg D$ these models give $G_0\\propto L_\\mathrm{IR}/D^3$ and $G_0\\propto L_\\mathrm{IR}/D^2$, respectively. Adopting $L_{IR}\\sim3\\times10^{10}$ $L_\\sun$ for the central $D\\sim300$ pc starburst region of M82~\\citep{Telesco1980,Joy1987}, and assuming that MIPS J1428 and M82 have the same average $G_0$ (Fig.~\\ref{fig3}), the factor of $\\sim$$1000$$\\mu^{-1}$ greater luminosity of MIPS J1428 then translates into a size scale of $D\\approx3$$\\mu^{-1/3}-10$$\\mu^{-1/2}$ kpc. In short, if the large luminosity of MIPS J1428 is produced by only moderate-intensity radiation, the star formation generating this radiation must be extended over a large area. While MIPS J1428 remains unresolved in dust continuum and molecular gas emission, integral field spectroscopy shows H$\\alpha$ emission extended over $\\approx$$0.7\\arcsec$ ($\\approx$$6$ kpc)~\\citep{Swinbank2006}. This result is consistent with our interpretation of MIPS J1428 as a source powered by extended star formation, although the degree to which the H$\\alpha$ image is distorted by the foreground lens, or by the presence of large and spatially-varying extinction, is unknown. Further interpretation of the connection between the bright [CII] emission in MIPS J1428 and the potentially large extent of the star formation will benefit from a size measurement in an extinction-free tracer, and better constraints on the lensing magnification. If MIPS J1428 is indeed powered by a starburst extending over several kiloparsecs this would be in sharp contrast to local ULIRGs, in which most of the emission traced by CO interferometry is contained within the central $\\sim$$1$ kpc~\\citep{Downes1998}. However, there is ample evidence that the most luminous star-forming galaxies at higher redshifts are more extended. The best studied sample of such sources are the SMGs, which like MIPS J1428 are submillimeter-bright galaxies with luminosities of $L_\\mathrm{IR}\\sim10^{13}$ $L_\\sun$ that are generated predominantly by star formation~\\citep{Alexander2005ApJ}. Interferometric imaging of the radio continuum of $z=1-3$ SMGs has shown that many of these sources have detectable emission extended over $\\sim$$10$ kpc~\\citep{Chapman2004extended}, and that the population has a median FWHM size of $\\sim$$5$ kpc~\\citep{Biggs2008}. Similarly, many of the SMGs studied by~\\citet{Tacconi2006,Tacconi2008} have CO FWHM sizes of $\\sim$$4$ kpc. The model of extended star formation we suggest here for MIPS J1428 is therefore consistent with the large sizes of SMGs, which indicate that the most luminous star-forming galaxies at $z=1-3$ are powered by starbursts extending over much larger regions than in local systems." }, "1003/1003.2497_arXiv.txt": { "abstract": "The radio astronomy community is currently building a number of phased array telescopes. The calibration of these telescopes is hampered by the fact that covariances of signals from closely spaced antennas are sensitive to noise coupling and to variations in sky brightness on large spatial scales. These effects are difficult and computationally expensive to model. We propose to model them phenomenologically using a non-diagonal noise covariance matrix. The parameters can be estimated using a weighted alternating least squares (WALS) algorithm iterating between the calibration parameters and the additive nuisance parameters. We demonstrate the effectiveness of our method using data from the low frequency array (LOFAR) prototype station. ", "introduction": "\\label{sec:intro} The radio astronomical community is currently constructing a number of large scale phased array telescopes such as the low-frequency array (LOFAR) \\cite{Vos2009-1} and the Murchison wide field array (MWA) \\cite{Lonsdale2008-1}. These instruments need to be calibrated regularly to track variations in the electronics of the antennas and receivers, as well as direction dependent variations of the ionosphere \\cite{Tol2007-1}. E.g., LOFAR will consist of order 50 stations (distributed over an area of a hundred kilometers or more), where each station consists of 96 ``low band'' dual-polarized dipole antennas (10--90 MHz) and 96 ``high band'' antennas (110--240 MHz). The latter antennas are in turn composed of 16 beamformed dual-polarized droopy dipole antennas. Each station provides a number of beamformed outputs, which in turn are correlated at a central location to form images and other astronomy products. In this paper we focus on the calibration of the station antennas. The general aim is to estimate the direction independent gains and phases of each sensor, as well as the direction dependent gains corresponding to each source. This is done for the 2--10 brightest sources in the sky, assuming a point source model. The problem is complicated by the fact that the covariances of signals from closely spaced antennas within a station are sensitive to noise coupling (for the lowest frequencies, the antennas are spaced closer than half a wavelength). Also, the point source model does not entirely hold because of bright emission from the plane of the galaxy, extending over the entire sky. Fortunately, this emission is spatially smooth, which implies that it is dominant on the short spatial scales in the array aperture, i.e. on the short baselines. In this paper we propose to model both short baseline effects by an additive noise covariance matrix, which in this case is not diagonal and has unknown entries for each short baseline. If we can estimate this matrix, a simple point source model will be sufficient to calibrate the array, which reduces the problem to a problem for which solutions are readily available \\cite{Fuhrmann1994-1, Wijnholds2006-1, Wijnholds2008-1}. Direction finding problems for calibrated arrays in the presence of unknown correlated noise have been extensively studied in the 1990s. It was proven that the general problem is not tractable without imposing some appropriate constraints on the noise covariance matrix or exploiting differences in temporal characteristics between source and noise signals \\cite{Stoica1992-1}. Radio astronomical signals generally behave like noise, thus temporal techniques (instrumental variables) are not applicable. Instead, we should rely on an appropriately constrained parameterized model of the noise covariance matrix. Starting with \\cite{bohme88icassp}, a series of papers were published; see \\cite{Goransson1999-1} for an overview. ML estimators for the source and instrument parameters under a generalized noise covariance parameterization is provided in \\cite{friedlander95tsp, Goransson1999-1}, whereas nonlinear least squares estimators were studied in \\cite{friedlander95tsp, Wax1996-1, Ottersten1998-1}. In either case, an analytic source and instrument parameter dependent solution is derived for the noise model parameters which is substituted back into the cost function. This cost function then has to be minimized using a generalized solving technique, such as Newton iterations. This approach works well if the number of instrument and source parameters is small. For larger problems (we consider 100 antenna/source parameters and over 750 noise covariance parameters), it is convenient to exploit suboptimal but closed-form analytic solutions, at least for initialization. We therefore propose a weighted alternating least squares (WALS) approach which iterates over noise, source and instrument parameters. The proposed method can thus be regarded as an extension to the methods proposed in \\cite{Wijnholds2008-1} \\emph{Notation}: The transpose operator is denoted by $^T$, the complex conjugate (Hermitian) transpose by $^H$, complex conjugation by $\\overline{(\\cdot)}$ and the pseudo-inverse by $^\\dagger$. An estimated value is denoted by $\\widehat{(\\cdot)}$. $\\otimes$ denotes the Kronecker product and $\\circ$ is used to denote the Khatri-Rao or column-wise Kronecker product of two matrices. $\\vec(\\cdot)$ converts a matrix to a vector by stacking the columns of the matrix. ", "conclusions": "We have demonstrated using data from a LOFAR prototype station that the effects of noise coupling, receiver noise powers and extended emission on a radio astronomical phased array can be phenomenologically described by a non-diagonal noise covariance matrix with non-zero entries on short baselines. These entries can be computationally efficient and accurately estimated by a WALS algorithm alternating between estimation of the correlated noise parameters and calibration parameters." }, "1003/1003.5562_arXiv.txt": { "abstract": "{ AU~Mic is a young, nearby X-ray active M-dwarf with an edge-on debris disk. Debris disk are the successors of the gaseous disks usually surrounding pre-main sequence stars which form after the first few Myrs of their host stars' lifetime, when -- presumably -- also the planet formation takes place. Since X-ray transmission spectroscopy is sensitive to the chemical composition of the absorber, features in the stellar spectrum of AU~Mic caused by its debris disk can in principle be detected. The upper limits we derive from our high resolution {\\it Chandra}~LETGS X-ray spectroscopy are on the same order as those from UV absorption measurements, consistent with the idea that AU~Mic's debris disk possesses an inner hole with only a very low density of sub-micron sized grains or gas. } ", "introduction": "The disks around young stars undergo dramatic changes during the first $\\sim10$~Myrs after their host stars' birth, when the gas content of the disk largely disappears \\citep{Alexander_2008, Meyer_2007, Hernandez_2006}, leaving behind a so-called debris disk. The Kuiper belt and the asteroid belt are the solar system's analogs of stellar debris disks. The main components of an optically thin debris disk are small grains with about sub-micrometer sizes, larger bodies in the cm range and, possibly, planets, which are thought to form in the same time-span. The initial composition of the material in the debris disks after the transition phase is not well known. Collisions of already formed smaller bodies replenish the dust in the ``older'' debris disks, while it is not clear, whether this is also the source of the initial dust in the debris disk or whether it is remnant proto-planetary dust. \\subsection{AU Mic and its activity} \\au is a $12^{+8}_{-4}$~Myr old M1 dwarf at a distance of about 10~pc, which belongs to the $\\beta$~Pic moving group \\citep[e.g.][]{Navascues_1999, Zuckerman_2001}. \\au is one of the brightest nearby X-ray emitters ($\\log L_X \\approx29.3$) and shows strong flaring activity, making \\au a valuable target for flare studies as shown by e.g. UV observations \\citep{Robinson_2001}. At X-ray wavelengths \\au has been observed many times. The first \\chan~ observation provided the highest resolution spectrum, but was limited to the wavelength range below $\\sim 25$~\\AA~ \\citep{Linsky_2002}. The \\xmm~ observation of \\au in 2000 was simultaneous with UV and VLA observations revealing several flares \\citep[][for line fluxes]{Smith_2005,Mitra-Kraev_2005,Ness_2003}. Furthermore, \\au was the target of FUSE and Hubble Space Telescope STIS observations \\citep[e.g.][]{Pagano_2000, Robinson_2001, DelZanna_2002}, aiming at the determination of the temperature structure of its chromosphere and corona. \\subsection{AU~Mic and its debris disk} The first indications for cold material around \\au go back to IRAS data, which exhibit excess emission at 60~$\\mu$m \\citep{Mathioudakis_1991,Song_2002}. A clear infrared excess at 850~$\\mu$m in the spectral energy distribution (SED) of \\au was detected by \\citet{Liu_2004}, clearly pointing to the existence of a debris disk. By assuming an optical thin disk at a single temperature, \\citet{Liu_2004} derived a mass of 0.01~$M_\\oplus$ and a temperature of 40~K for the disk. These values have been confirmed by \\citet{Rebull_2008} from Spitzer data \\citep[see also][]{Chen_2005}. \\citet{Metchev_2005} also confirmed the dust mass of about 0.01~$M_\\oplus$ composed of grains in the submicron regime by modelling the optical, near-IR and SED data. Optical observations by \\citet{Kalas_2004}, initiated shortly after \\citet{Liu_2004}, clearly showed the presence of a debris disk around \\au with a radial extent of at least 210~AU and almost perfectly edge-on. The disk has then been subsequently studied at optical wavelengths with, e.g., the Hubble Space Telescope (HST) and adaptive optics, making it one of the most well studied debris disks. The models derived by \\citet{Krist_2005} from their HST observation restrict the disk inclination to $i \\gtrsim 89^\\circ$; they also confirmed the small-scale brightness variations detected by \\citet{Kalas_2004}. These disk inhomogeneities can be readily explained by the existence of orbiting planets, however, no clear signatures of a planet have been found to date \\citep[][]{Hebb_2007, Metchev_2005}. Two studies aimed at the detection of circum-stellar gas in the \\au disk. By their non-detection of far-UV H$_2$ absorption, \\citet{Roberge_2005} derived an upper limit on the gas column density along the line of sight of $N_{H_2} < 10^{19}$~cm$^{-2}$. The detection of $H_2$ in fluorescence enabled \\citet{France_2007} to derive a column density of $3\\times10^{15}\\, \\text{cm}^{-2} < N_{H_2} < 2\\times10^{17}\\, \\text{cm}^{-2}$ ($T_{H_2}=800$~K and 2000~K, respectively). Comparing their $H_2$ value with the the CO results of \\citet{Liu_2004}, they conclude that \\mbox{$H_2$} contributes less than about 1/30$^{th}$ to the total disk mass. The $H$ absorption mainly traces interstellar rather than circum-stellar material and has a column of $N_H = 2.3\\times10^{18}$~cm$^{-2}$ \\citep{Wood_2005}, thus corresponding to within a factor of five to the Mg~{\\sc ii} absorption measurements of \\citet[][, $N_{Mg} = 1.6\\times10^{13}$~cm$^{-2}$]{Redfield_2002}, assuming solar abundances and Mg~{\\sc II} to be the dominant Mg species \\citep{Slavin_2002}. \\subsection{Disk models} \\citet{Krist_2005} used three-dimensional models of the scattering cross-section densities throughout the disk to interpret their HST images. They find that in the inner disk region \\mbox{(12~AU $ 15~\\rm mag$) are less spatially extended on the sky with LEDA~166099 having an angular size of $D_{\\rm ext}<1$ arcminute. \\section[]{Observations} \\label{chap3:DataAcqu} Deep NIR imaging of the six dwarf galaxies was obtained with the 1.4m Infrared Survey Facility\\footnote{A description of the telescope can be found at http://www.z.phys.nagoya-u.ac.jp/$\\sim$telescop/index\\_e.html~.} (IRSF) telescope in Sutherland, South Africa \\citep{Gla2000}. The \\textit{Sirius} detector of the IRSF telescope consists of three 1024$\\times$1024 HgCdTe arrays. The array system allows for simultaneous 3-channel $J$- ($\\rm 1.25\\mu m$), $H$- ($\\rm 1.65\\mu m$) and $K_s$-band ($\\rm 2.15\\mu m$) imaging. The individual CCDs give a total FoV of 7\\farcm8$\\times$7\\farcm8 together with a 0\\farcs45 pixel scale. The NIR data were obtained during three different observing runs over the period of 2006--2007. A log of all the galaxy observations is presented in Table~\\ref{NIRobs1}. The observations performed in 2006 February and June were carried out using shared telescope time. Single galaxy observations of NGC~3115~DW01 and NGC~59 were obtained during these respective observing runs. A further week in 2007 March was dedicated to the imaging of dwarf and low SB galaxies. Simultaneous $J$-, $H$- and $K_s$-band imaging of the other four galaxies was obtained during this observing run. The total integration times for the galaxy observations vary from 60--132 minutes. All observations were performed during grey time under photometric conditions. The aim of the NIR observations is to conduct a detailed photometric analysis of the dwarf galaxies out to at least the $\\mu_{K_s}\\sim23~\\rm~mag~arcsec^{-2}$ isophote. R. Metcalfe \\& M. McCall (private communication) derived an exposure time for detecting nearby dwarf irregular (dIrr) galaxies down to this isophote using the IRSF telescope. The exposure time was calculated based on NIR observations of these galaxies with the 3.6m Canada-France-Hawaii telescope (CFHT) and the 2.1m OAN-SPM telescope \\citep{Vad05}. They found that a limiting magnitude of $\\mu_{K_s}\\sim23~\\rm~mag~arcsec^{-2}$ can be reached in $\\sim$70 minutes. The goal was therefore to observe each galaxy in principle for at least 70 minutes. An exception was made for the luminous galaxy NGC~59 which was only observed for 60 minutes on the shared night of June 11 2006. The low SB dwarf galaxies, LEDA~166099 and UGCA~200, were both observed for 96~minutes to ensure their detection (see Table~\\ref{NIRobs1}). The quality of deep NIR images is greatly affected by temporal and spatial variations in the sky background. A good estimate of the sky-level is particularly important when observing low SB dwarf galaxies such as LEDA~166099 and UGCA~200. The NIR observing technique of \\citet{Vad04} was employed to ensure an optimal extraction of the sky-level in the images. They propose an observing sequence of the form: \\begin{equation} \\rm{sky-galaxy-sky-~....~-sky-galaxy-sky}~, \\label{eq:ObsSeq} \\end{equation} when observing faint extended sources in the NIR. Temporal variations in the sky-level were accounted for by allowing equal exposures for the galaxy and sky frames. A dithering step of 10\\arcsec\\ was applied to each new galaxy and sky exposure in the observing sequence. This was necessary for the removal of bad pixels and contaminants from the images. \\begin{table} \\centering \\caption{Observing Log \\label{NIRobs1}} \\begin{tabular}{@{}lrcccc@{}} \\hline & & & Exp & Total Exp \\\\ Galaxy & {Date (UT)} & Filters & (min) & (min) \\\\ \\hline NGC~3115~DW01 & 2006 Feb 14 & $JHK_s$ & 48 & 132\\\\ & Feb 16 & $JHK_s$ & 24 & \\\\ & Feb 17 & $JHK_s$ & 60 & \\\\ NGC~59 & 2006 Jun 11 & $JHK_s$ & 60 & 60 \\\\ LEDA~166099 & 2007 Mar 08 & $JHK_s$ & 96 & 96 \\\\ ESO~384-016 & 2007 Mar 08 & $JHK_s$ & 24 & 84 \\\\ & Mar 09 & $JHK_s$ & 60 & \\\\ UGCA~200 & 2007 Mar 10 & $JHK_s$ & 24 & 96 \\\\ & Mar 11 & $JHK_s$ & 24 & \\\\ & Mar 13 & $JHK_s$ & 48 & \\\\ NGC~5206 & 2007 Mar 10 & $JHK_s$ & 48 & 60 \\\\ & Mar 13 & $JHK_s$ & 12 & \\\\ \\hline \\end{tabular} \\end{table} The sequential sky and galaxy frames in (\\ref{eq:ObsSeq}) were exposed for 60 seconds each. Individual exposures were sub-divided into 3$\\times$20s non-dithered frames to avoid saturation of the pixel arrays. The background was sampled by choosing a sky region in close proximity to the galaxy that shows the least amount of stellar contamination. The sky-level was sampled either 10\\arcmin\\ North or South of the galaxy center. Twilight sky images were obtained in the evening and morning. A sequence of equal-duration exposures was taken as the twilight brightened or faded in each filter. These sky flats were used to remove the spatial variations in the images. A series of dark exposures were taken every morning in the $J$-, $H$-, and $K_s$-bands for removal of the detector signature. \\section[]{Data Reduction and Calibration} \\label{chap3:Reduction} The data reduction was carried out using standard tasks in \\iraf\\footnote{\\iraf\\ is distributed by the National Optical Astronomy Observatory (NOAO), which is operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation.}. The first step in the reduction procedure involves the removal of the bias-level from the target images (\\ie sky flats, galaxy and sky exposures). Dark frames having the same exposure time as the target image were used to create a master dark frame which was then subtracted from the individual target exposures. A master flat was created in each filter from a sequence of 30--40 equal-duration twilight sky exposures. The master flat was used to correct for the pixel-to-pixel sensitivity in the sky and galaxy exposures. The non-dithered 20sec galaxy and sky exposures were combined giving an observational sequence of the form: \\begin{equation} S_1 - T_1 - S_2 - ... - S_i - T_i - S_{i+1} - ... - S_n - T_n - S_{n+1}, \\label{sky3} \\end{equation} where a dithered region of the sky $S_i$ is sampled before and after each galaxy exposure $T_i$. This technique of straddling the sky frames allows for the interpolation of the background at the time of the galaxy observation. The sky subtraction procedure of \\citet{Vad04} was followed in subtracting the background from the galaxy images. Basically, this procedure creates a smooth background by first subtracting adjacent sky frames using two different combinations \\ie $S_i - S_{i+1}$ and $S_{i+1} - S_i$. The background level at the time of the galaxy observation is obtained by averaging the resulting sky frames. The background level and corresponding uncertainty were measured at various locations in the galaxy image. These measurements were carried out at a radius of $r\\ga4\\arcmin$ from the galaxy center to ensure that the results are not influenced by light from the galaxy itself. We find that the reduction method described above removes the sky to an accuracy of $\\sim$0.02\\% in the $K_s$-band relative to the original signal. \\begin{table} \\centering \\caption{Properties of the reduced NIR images \\label{NIRobs2}} \\begin{tabular}{@{}llcc@{}} \\hline & & {Exp} & {Ave seeing} \\\\ {Galaxy} & {Filter} & {(min)} & {(arcsec)} \\\\ \\hline NGC 3115 DW01 & $J$ & 57 & 1.3 \\\\ & $H$ & 57 & 1.2 \\\\ & $K_s$ & 57 & 1.2 \\\\ NGC 59 & $J$ & 56 & 1.6 \\\\ & $H$ & 57 & 1.5 \\\\ & $K_s$ & 58 & 1.4 \\\\ LEDA 166099 & $J$ & 89 & 1.4 \\\\ & $H$ & 93 & 1.3 \\\\ & $K_s$ & 94 & 1.3 \\\\ ESO 384-016 & $J$ & 76 & 1.5 \\\\ & $H$ & 83 & 1.5 \\\\ & $K_s$ & 83 & 1.4 \\\\ UGCA 200 & $J$ & 74 & 1.5 \\\\ & $H$ & 73 & 1.4 \\\\ & $K_s$ & 81 & 1.4 \\\\ NGC 5206 & $J$ & 44 & 1.3 \\\\ & $H$ & 44 & 1.3 \\\\ & $K_s$ & 47 & 1.2 \\\\ \\hline \\end{tabular} \\end{table} The total on-source exposure time of the reduced galaxy images is listed in Table \\ref{NIRobs2}. Those galaxy exposures largely affected by atmospheric and telescopic defects were not used in creating the final images. This results in the different exposure times in the $J$-, $H$- and $K_s$-bands for a single galaxy observation. Defects in the images include extreme blurring due to changes in the observing conditions as well as problems with the CCD readout. The seeing in the final galaxy images varies from 1\\farcs2 in the $K_s$-band to 1\\farcs6 in the $J$-band. The reduced $K_s$-band galaxy images are shown in Fig.~\\ref{NIR_im1}. A distinct nucleus is seen in the galaxies NGC~3115 DW01, NGC~59, LEDA~166099 and NGC~5206. The brighter galaxies NGC~3115 DW01, NGC~59 and NCG~5206 show extended light profiles at least 1\\arcmin\\ beyond the nuclear component. The white features seen for example in the $K_s$-band image of LEDA~166099 are negative residuals left behind by the sky subtraction. These features are observed in the crowded stellar fields where it is more challenging to obtain a smooth sky background. The central region of the galaxy NGC~59 is shown in Fig.~\\ref{n59zoom}. Two intensity peaks aligned in the northeast-southwest direction can be distinguished in the center of this galaxy. The two peaks are separated by $\\sim$2\\farcs3 with the northern component being more luminous in all three wavelength bands. Recent star formation activity has been detected in the center of NGC~59 by \\citet{Ski03}. In this paper, we assume that the northern component is the ``true\" nucleus of the galaxy while the second component is a star-forming region. The photometry of the two nuclear components together with a detailed kinematic study of NGC~59, will be presented in a companion paper (de Swardt et al., in preparation). \\subsection[]{Photometric Calibration} The photometric calibrations of the IRSF data involve a direct comparison of the instrumental magnitudes of point sources in the field to their corresponding apparent magnitudes given by 2MASS. A colour correction was applied to the instrumental magnitudes to account for differences in the IRSF and 2MASS filter systems. The point sources used to calibrate the galaxy images were selected from the 2MASS Point Source Catalogue \\citep{Skr06} and were chosen to satisfy the following criteria: \\begin{enumerate} \\item{Point sources should be brighter than the 2MASS completeness limit: $J\\la15.8\\rm~mag$, $H\\la15.1\\rm~mag$ and $K_s\\la14.3\\rm~mag$.} \\item{`AAA' quality photometry is available for all stellar sources in the 2MASS catalogue. Point sources are rated as having `AAA' quality photometry if they have a magnitude uncertainty of less than 10\\% in all three wavelength bands.} \\item{Saturated stellar sources having magnitudes brighter than $10\\rm~mag$ in the IRSF images were not used in the photometric calibrations.} \\end{enumerate} The instrumental magnitudes of the point sources were transformed to apparent magnitudes using the equations: \\begin{equation} j = J + j_1 + j_2(J - K_s)~, \\label{jcalib} \\end{equation} \\begin{equation} h = H + h_1 + h_2(J - H)~, \\label{hcalib} \\end{equation} \\begin{equation} k_s = K_s + k_1 + k_2(J - K_s)~, \\label{kcalib} \\end{equation} where $J$, $H$ and $K_s$ are the apparent magnitudes of the stars. The instrumental magnitudes are given by $j$, $h$ and $k_s$ in the $J$-, $H$- and $K_s$-bands, respectively. The nightly zero-points in the different wavelength bands are $j_1$, $h_1$ and $k_1$. These were determined for the individual galaxy images of each night. The offset between the instrumental and apparent magnitude (given by the 2MASS Point Source catalogue) was calculated for each of the point sources. The uncertainty in the magnitude offset is given by the square of the internal errors associated with the 2MASS and instrumental magnitudes, respectively. The magnitude offsets of the point sources show little scatter ($< 0.05\\rm~mag$) so that the nightly zero-point was taken as the mean magnitude offset between the IRSF and 2MASS point sources. \\begin{figure*} \\includegraphics[width=6.5cm]{n3115image_Kv1.ps} \\includegraphics[width=7.5cm]{n3115_sb_paper.ps} \\includegraphics[width=6.5cm]{n59image_Kv1.ps} \\includegraphics[width=7.5cm]{n59_sb_paper.ps} \\includegraphics[width=6.5cm]{l166099image_Kv1.ps} \\includegraphics[width=7.5cm]{l166099_sb_paper.ps} \\caption{The reduced $K_s$-band image of the dwarf galaxies is shown on the left. The horizontal bar at the bottom of each image indicates a scale of 2\\arcmin. North is up and East is left. The $J$-, $H$- and $K_s$-band SB profiles of the galaxies are shown on the right. The corresponding $J$-$K_s$ colour profile is displayed below the SB profiles of the galaxy.} \\label{NIR_im1} \\end{figure*} \\begin{figure*} \\includegraphics[width=6.5cm]{e384016image_Kv1.ps} \\includegraphics[width=7.5cm]{e384016_sb_paper.ps} \\includegraphics[width=6.5cm]{n5206image_Kv1.ps} \\includegraphics[width=7.5cm]{n5206_sb_paper.ps} \\includegraphics[width=6.5cm]{u200image_Kv1.ps} \\includegraphics[width=7.5cm]{u200_sb_paper.ps} \\contcaption{} \\end{figure*} \\begin{figure} \\begin{center} \\includegraphics[width=6cm]{h_zoomn59.ps} \\caption[Double nuclear component of the Scl group dwarf galaxy NGC~59 in the NIR.]{Double nuclear component of the Scl group dwarf galaxy NGC~59 as seen in the $H$-band. The two peaks are separated by $\\sim$2\\farcs3. The image size shown is 68\\arcsec$\\times$58\\arcsec. North is up and East is left.} \\label{n59zoom} \\end{center} \\end{figure} The final term in equations (\\ref{jcalib})--(\\ref{kcalib}) represents the colour correction to the magnitudes. The colour coefficients for the IRSF observations have been found to be $j_2=-0.018\\rm~mag$, $h_2=0.050\\rm~mag$ and $k_2=0.079\\rm~mag$ in the $J$-, $H$- and $K_s$-bands \\citep{Kot07}, respectively. A correction for the airmass was not directly applied to the instrumental magnitudes. This calibration term is accounted for in the nightly zero-point correction derived from the 2MASS apparent magnitudes. Overall, the IRSF apparent magnitudes derived from equations (\\ref{jcalib})--(\\ref{kcalib}) agree within $0.05\\rm~mag$ with those given by 2MASS. \\subsection[]{Star Subtraction} The $K_s$-band galaxy images in Fig.~\\ref{NIR_im1} show that the foreground contamination varies from one image to the next. A careful removal of the foreground stars is essential to obtain photometric results not influenced by resolved stellar sources. The star subtraction routine of T. Nagayama (Kyoto University, private communication) was used in removing the foreground stars from the galaxy images. This routine builds the point spread function (psf) of the image by carrying out three iterative runs of the psf-fitting to the stars. This ``automated\" building of the psf distinguishes the star subtraction routine of Nagayama from the conventional \\textsc{killall}\\ routine of \\citet{But1999}. The effectiveness of the star-subtraction routine is illustrated in Fig.~\\ref{starsub} which displays the original and star-subtracted image for the galaxy NGC~5206. Most of the stellar sources are cleanly removed by the routine. The residuals from background galaxies, bright and saturated stars were interactively removed using the IMEDIT task in \\iraf. An extended source can vaguely be seen south-west of the centre of NGC~3115~DW01 (see Fig.~\\ref{NIR_im1}). The NIR colour of this source indicates that it is a background galaxy which is seen through the fairly bright central region of NGC~3115~DW01. The $B$-band image of NGC~3115~DW01 from \\citet{Par2002} reveals that this source is indeed a background spiral galaxy. The background galaxy was retained in the NIR images of NGC~3115~DW01 to minimize the effect of smoothing on the light distribution of the galaxy. \\begin{figure} \\includegraphics[width=4.1cm]{n5206_Hstars.ps} \\includegraphics[width=4.1cm]{n5206_Hstarsub.ps} \\caption{The original (\\textit{left}) and star-subtracted image (\\textit{right}) of the galaxy NGC~5206. The surface photometry was performed on the cleaned galaxy image.} \\label{starsub} \\end{figure} \\section[]{Surface Photometry} \\label{sec5:SBprofiles} The ELLIPSE\\ task in \\iraf\\ was used to measure the intensity distribution along the semi-major axis of the galaxy. The fitting process requires a measure of the position angle of the galaxy on the sky, as well as its ellipticity. These parameters were derived from the outer isophotes where the older stellar population is expected to set the overall geometry of the galaxy. The position angle (PA) and ellipticity ($\\epsilon$) of the galaxies are listed in Table~\\ref{2mass_comp} which represent the mean value obtained from the three NIR bands. The $B$-band value of the position angle and ellipticity given by \\citet{Par2002} is adopted for UGCA~200 whose low signal-to-noise (S/N) levels do not allow for the measurement of these parameters. By keeping the geometrical parameters fixed it is possible to measure the intensity distribution of the galaxy out to even fainter levels. The radial surface brightness of the dwarf galaxies is measured down to $\\mu_{\\rm lim}\\simeq24\\rm~mag~arcsec^{-2}$ in the $J$- and $H$-bands, and $\\mu_{\\rm lim}\\simeq23\\rm~mag~arcsec^{-2}$ in the $K_s$-band. The near-infrared SB profiles of the six dwarf galaxies are displayed in Fig.~\\ref{NIR_im1}. The error associated with each point in the SB profile was computed as the RMS scatter in intensity given by the isophotal fitting. The NIR light profiles of the galaxies NGC~3115~DW01, NGC~59, NGC~5206 and LEDA~166099 can be divided into two components: the nuclear and low SB component. The SB profiles for these galaxies show a distinct nucleus whereas pure exponential profiles are observed for ESO~384-016 and UGCA~200. The background galaxy can be identified in the $J$- and $H$-band SB profiles of NGC~3115~DW01 which show a slight increase at the location of the galaxy. NGC~5206 is the most spatially extended galaxy in the sample reaching the $\\mu_J=23\\rm~mag~arcsec^{-2}$ isophote at $r$$\\sim$3\\arcmin. The central SB component of UGCA~200 lies up to $\\sim3\\rm~mag~arcsec^{-2}$ above the NIR detection limit. The SB profile for this faint galaxy is measured out to a radius of $r$=40--50\\arcsec\\ despite its extremely low intensity levels. \\begin{table*} \\centering \\begin{minipage}{170mm} \\caption{Measured parameters for six dwarf galaxies from deep near-infared imaging: position angle (PA) measured from North to East, ellipticity ($\\epsilon$), total magnitude ($m_t$) and corresponding aperture radius ($r_t$). The photometric parameters from the 2MASS All-Sky Extended Source Catalogue are listed for the brightest dwarf galaxies ($B\\la13.4\\rm~mag$) in the sample. The total magnitudes together with the distances from Table~\\ref{dwarf_properties1} were used to compute the absolute magnitudes $M_{\\rm abs}$ in the different wavelength bands. The extinction coefficients $A_\\lambda$ from \\citet{Sch1998} are listed. \\label{2mass_comp}} \\begin{tabular}{@{}lcccrccccccrc@{}} \\hline & \\multicolumn{7}{c}{IRSF} & \\multicolumn{5}{c}{2MASS} \\\\ & {PA} & & & {$m_t~~~~~$} & {$r_t$} & $M_{\\rm abs}$ & $A_\\lambda$ & {PA} & & & {$m_t~~$} & {$r_t$} \\\\ {Galaxy} & {(deg)} & {$\\epsilon$} & {Filter} & {(mag)$~~~~$} & {(arcsec)} & (mag) & (mag) & {(deg)} & {$\\epsilon$} & {Filter} & {(mag)} & {(arcsec)} \\\\ \\hline NGC~3115~DW01 &\t7 & 0.14 & $J$ & $10.70\\pm0.06$ & 90 & $-19.23\\pm0.20$ & 0.05 & 45 & 0.32 & $J$ & 10.99 & 70 \\\\ &\t& & $H$ & $10.15\\pm0.07$ & & $-19.78\\pm0.20$ & 0.03 & & & $H$ & $10.15$ & \\\\ &\t& & $K_s$ & $9.94\\pm0.07$ & & $-19.99\\pm0.20$ & 0.02 & & & $K_s$ & $9.94$ & \\\\ NGC~59 & 121 & 0.41 & $J$ & $10.84\\pm0.03$ & 99 & $-17.38\\pm0.16$ & 0.02 & 115 & 0.50 & $J$ & $10.89$ & 89 \\\\ &\t& & $H$ & $10.24\\pm0.03$ & & $-17.98\\pm0.16$ & 0.01 & & & $H$ & $10.26$ & \\\\ &\t& & $K_s$ & $10.07\\pm0.03$ & & $-18.15\\pm0.16$ & 0.01 & & & $K_s$ & $10.10$ & \\\\ LEDA~166099 & 125 & 0.34 & $J$ & $13.72\\pm0.03$ & 45& $-16.24\\pm0.20$ & 0.18 & -- & -- & & & \\\\ &\t& & $H$ & $13.18\\pm0.04$ & & $-16.78\\pm0.20$ & 0.11 & & & & \\\\ &\t& & $K_s$ & $12.95\\pm0.04$ & & $-17.01\\pm0.20$ & 0.07 & & & & \\\\ ESO~384-016 & 82 & 0.31 & $J$ & $13.12\\pm0.05$ & 54& $-15.00\\pm0.13$ & 0.07 & -- & -- & & & \\\\ &\t& & $H$ & $12.49\\pm0.07$ & & $-15.63\\pm0.13$ & 0.04 & & & & \\\\ &\t& & $K_s$ & $12.35\\pm0.07$ & & $-15.77\\pm0.13$ & 0.03 & & & & \\\\ NGC~5206 & 23 & 0.35 & $J$ & $8.91\\pm0.04$ & 180 & $-18.87\\pm0.20$ & 0.11 & 45 & 0.16 & $J$ & 9.39 & 114 \\\\ &\t& & $H$ & $8.35\\pm0.05$ & & $-19.43\\pm0.20$ & 0.07 & & & $H$ & 8.55 & \\\\ &\t& & $K_s$ & $8.05\\pm0.05$ & & $-19.73\\pm0.20$ & 0.04 & & & $K_s$ & 8.49 & \\\\ UGCA~200 & -31 & 0.30 & $J$ & $13.84\\pm0.06$ & 45& $-16.09\\pm0.20$ & 0.04 & -- & -- & & & \\\\ &\t& & $H$ & $13.32\\pm0.07$ & & $-16.61\\pm0.20$ & 0.03 & & & \\\\ &\t& & $K_s$ & $13.26\\pm0.09$ & &$-16.67\\pm0.20$ & 0.02 & & & \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*} The NIR colour profiles (\\ie $J\\mbox{-}K_s$, $H\\mbox{-}K_s$ and $J\\mbox{-}H$ profiles) were derived by subtracting the SB profiles as function of radius in the respective wavelength bands. We only display the $J\\mbox{-}K_s$ colour profile below each of the galaxy SB profiles in Fig.~\\ref{NIR_im1}. Generally, the galaxy colours remain almost constant for surface brightnesses of $\\mu_{K_s}\\la 22\\rm~mag~arcsec^{-2}$ after which the noise levels dominate the $K_s$-band photometry. The NIR colour of the extremely low SB galaxy, UGCA~200, shows the largest scatter from the mean colour of up to $0.25\\rm~mag$ within the $\\mu_{K_s}\\la 22\\rm~mag~arcsec^{-2}$ limit. The mean colour of each galaxy was computed by averaging the data points above the $\\mu_{K_s}\\sim 22\\rm~mag~arcsec^{-2}$ detection limit. The mean colours for each galaxy are listed in Table~\\ref{dwarf_colors}. The six dwarfs exhibit typical NIR colours observed for dwarf galaxies in the 2MASS Extended Source Catalogue \\citep{Jar2000}. The galaxies NGC~3115~DW01, LEDA~166099 and NGC~5206 are found to have redder colours of $J\\mbox{-}K_s > 0.8\\rm~mag$, which is in agreement with their optical $B\\mbox{-}R$ colours shown in Table~\\ref{dwarf_properties1}. A mean colour of $0.7 < J\\mbox{-}K_s < 0.8\\rm~mag$ is measured for the remaining galaxies with NGC~59 and ESO~384-016 showing similar NIR colours. \\subsection[]{Total Magnitudes} \\label{Total_mag} The total apparent magnitude is given by the integrated flux within the detection limit of the galaxy. A growth curve was constructed for each galaxy by integrating the intensity in circular apertures. The apertures were defined at a radius step of $r=\\sqrt{ab}$ where $a$ and $b$ are the respective major and minor axis of the galaxy. The total apparent magnitude, $m_t$, corresponds to the asymptotic intensity of the growth curve which is measured down to the background level of the image. A good measure of the background level was found by systematically varying the sky brightness in the image. The growth curve converges asymptotically to a flat background when the correct sky level is achieved. The total magnitudes ($m_t$) of the six dwarf galaxies are listed in Table~\\ref{2mass_comp}. The accuracy of the total magnitudes depends on the data reduction and calibration procedures. The sky subtraction technique gives an uncertainty of $\\la 0.02\\rm~mag$ in the background level for the three wavelength bands. A larger source of uncertainty is introduced by the photometric calibrations of the images with the accuracy in the measured zero points varying from $0.03\\rm~mag$ in the $J$-band to $0.09\\rm~mag$ in the $K_s$-band. The total magnitudes are thus derived with an accuracy of $<0.1\\rm~mag$ given the errors introduced through the reduction and calibration procedures. \\begin{figure} \\includegraphics[width=4.1cm]{n3115image_2massH.ps} \\includegraphics[width=4.1cm]{n3115image_H.ps} \\caption{A comparison of the $H$-band image of the galaxy NGC~3115~DW01 obtained from 2MASS (\\textit{left}) and the IRSF (\\textit{right}). The image size is 3\\farcm5$\\times$3\\farcm5. North is up and East is left.} \\label{IRSFvs2MASS} \\end{figure} NIR imaging and photometry are available for the three brightest dwarfs ($B\\la13.4\\rm~mag$) in our sample from the 2MASS All-Sky Extended Source Catalogue\\footnote{The 2MASS Extended Source Catalogue can be accessed online at http://irsa.ipac.caltech.edu/applications/Gator/.}. A comparison of the IRSF and 2MASS $H$-band image of the galaxy NGC~3115~DW01 is shown in Fig.~\\ref{IRSFvs2MASS}. The four times higher spatial resolution of the IRSF images gives a clear distinction between the nucleus and the extended low SB component of the galaxy. The 2MASS total magnitudes for the galaxies NGC~3115~DW01, NGC~59 and NGC~5206 are listed in Table~\\ref{2mass_comp}. These magnitudes are the extrapolated total magnitudes $r_{\\rm ext}$ in the 2MASS Extended Source catalogue and correspond to the total integrated flux of the galaxy measured down to the background level. The deep observations of the dwarf galaxies allow for their detection out to larger radii compared to 2MASS. In particular, the IRSF observations of NGC~5206 show that this galaxy has an extended low SB component which was not completely detected by 2MASS. The difference between the IRSF and 2MASS total magnitudes for the three dwarfs is plotted in Fig.~\\ref{2masscomp}. It is seen that 2MASS underestimates the flux of the three galaxies by up to $\\la0.5\\rm~mag$. The largest deviation of $J=0.48\\rm~mag$ from the 2MASS total magnitude is observed for the galaxy NGC~5206 which shows an extended low SB component in the deep NIR images. This underestimation of the galaxy fluxes obtained by 2MASS has been observed by \\citet{Kir2008} and \\citet{And2002}. Overall, the 2MASS survey fails to detect low SB galaxies so that the NIR photometric measurements of the galaxies ESO~384-016, LEDA~166099 and UGCA~200 are presented here for the first time. The deep NIR observations emphasize the serious selection biases of the 2MASS galaxy survey. \\subsection[]{Structural Parameters from One-dimensional Analytical Profile Fits} \\label{sec6:SBfits} The $J$-, $H$- and $K_s$-band SB profiles (shown in Fig.~\\ref{NIR_im1}) were fit with an analytical function to characterize the underlying structure of the six dwarf galaxies. Those galaxies having a nuclear component (NGC~3115~DW01, NGC~59, LEDA~166099 and NGC~5206) were fit with a combination of an exponential and S\\'ersic law to model the stellar disk component and nucleus, respectively. The S\\'ersic law can take the form of \\citep{Cao1993} : \\begin{equation} I(r) = I_e{\\rm exp}\\left[-b_n\\left( \\left(\\frac{r}{r_e}\\right)^{1/n} - 1\\right)\\right]~, \\label{sersic} \\end{equation} where $I_e$ is the intensity at the effective radius $r_e$ which contains half of the total integrated light from the model. These parameters were derived for the stellar disk of the galaxy by setting the shape parameter $n = 1$ giving an exponential profile fit. The parameter $b_n$ is directly related to the shape parameter $n$ and can be approximated by $b_n\\simeq1.9992n - 0.3271$ for $0.5 \\la n \\la 10$. The light profiles of the remaining dwarf galaxies in the sample were appropriately modeled using a single S\\'ersic function as given by equation~(\\ref{sersic}). The fits to the radial SB profiles were performed using the NFIT1D task in \\iraf\\ which employs a $\\chi^2$-minimization technique in finding the best-fit to the observed light profiles. An exponential and S\\'ersic function were fit simultaneously to the SB profiles of galaxies hosting a nucleus. For these galaxies, the stellar disk can be modeled by an exponential law out to the detection limit. The central region of the light profiles was fit with a S\\'ersic function to model the steep rise in the SB profile towards the center of the galaxy. The structural parameters (\\ie the effective radius $r_e$ and corresponding SB $\\mu_e$) of the extended stellar disk are given in Table~\\ref{1Dfits} for the nucleated dwarf galaxies. A less than $\\sim1\\%$ variation is observed in these parameters if instead a single exponential law is fit to the light profile. This implies that the derived structural parameters are associated purely with the stellar disk as they are not sensitive to the S\\'ersic fit in the central parts of the galaxy. The structural parameters for the non-nucleated dwarfs ESO~384-016 and UGCA~200 are given by a S\\'ersic fit with shape parameter in the range of $0.7-15.5\\rm~mag$). It is not clear if a selection effect results in the lack of low luminosity, red dIrr galaxies at the faint end of this relation. \\begin{figure} \\begin{center} \\includegraphics[width=8.5cm]{JK_paper.ps} \\caption[$J$-$K_s$ colour-magnitude diagram for early-type galaxies and dIrr's.]{$J$-$K_s$ colour-magnitude diagram for the Virgo early-type (dE, dS0 and dE/dS0) galaxies (\\textit{top}) and dIrr's (\\textit{bottom}). The six IRSF dwarfs are shown in both panels where the filled symbols indicate those dwarfs which have been detected in \\ion{H}{i}. The linear least-squares fit and $1\\sigma$ standard deviation for each galaxy sample are represented by the dashed and dotted lines, respectively. The slope $m$ of the linear fit is indicated. For comparison, the linear fit for the dIrr galaxies is re-plotted (solid line) in the top panel.} \\label{JK_dwarfs} \\end{center} \\end{figure} \\subsection{Dwarf Galaxy Evolutionary Sequence and Morphology} The NIR photometric results for the six dwarf galaxies were combined with existing optical measurements to assess their evolutionary state and morphology. The $B$-$K_s$ colour is known to be a good indicator of the galaxy morphological type \\citep[\\eg][]{Jar2003}. In Fig.~\\ref{BK}, the $B$-$K_s$ colours of the six IRSF dwarfs are plotted against their corresponding $K_s$-band luminosity. The total apparent $B$-band magnitudes of the dwarfs (see Table~\\ref{dwarf_properties1}) are taken from \\citet{Kara2004}. For the faintest galaxies in the sample ($B\\ga15.2\\rm~mag$), the total $B$-band magnitudes have been determined with an accuracy of $\\sim0.5\\rm~mag$. The $B$-$K_s$ colours were calculated as the difference between the total apparent $B$ and $K_s$-band (from Table~\\ref{2mass_comp}) magnitudes of the galaxies. To gain more perspective of where the IRSF dwarfs are located relative to other galaxy morphologies, we have added four different galaxy samples to the $B$-$K_s$ colour-magnitude diagram: early-type dwarfs (dE, dS0, dE/dS0) and elliptical galaxies from the Virgo Cluster (taken from the the Goldmine database), dIrr galaxies from \\citet{Vad05}, blue compact dwarf (BCD) galaxies from \\citet{Cai03} and \\cite{Noe2003}. The general trend shows redder $B$-$K_s$ colours for the elliptical galaxies, while bluer colours are measured for both early and late-type dwarfs. An average $B$-$K_s$ colour of $\\sim4\\rm~mag$ was measured for early-type (E, S0) galaxies in the 2MASS Extended Source Catalog \\citep{Jar2003}. The early-type dwarfs form a continuous sequence between the more luminous elliptical galaxies and late-type dwarfs. The emission from the BCD galaxies is dominated by the younger, starburst component indicated by the blue colours. Figure~\\ref{BK} shows that the six IRSF dwarfs fit very well with the sequence of early-type dwarf galaxies. \\begin{figure} \\begin{center} \\includegraphics[width=8.5cm]{BK_paper.ps} \\caption{$B$-$K_s$ colour-magnitude diagram. The six IRSF dwarfs are shown by the black points where the filled symbols indicate those dwarfs which have been detected in \\ion{H}{i}. The elliptical and early-type dwarf galaxies from the Virgo Cluster are also shown. The dIrr's \\citep{Vad05} are indicated by the open circles while the BCD galaxies from \\citet{Cai03} and \\citet{Noe2003} are represented by the filled triangular points.} \\label{BK} \\end{center} \\end{figure} A detailed study of the relationship between the $B$- and $H$-band luminosity of galaxies was made by \\citet{Kir2008}. They found a tight correlation between the absolute $B$- and $H$-band magnitudes for their sample of nearby galaxies ($D\\la10\\rm Mpc$). This correlation was extended to include more luminous galaxies such as the sample of spirals from \\citet{Kas2006} and galaxies from the Virgo Cluster \\citep{Gav2003}. The six IRSF dwarfs are found to closely follow the linear relation of \\citet{Kir2008}. The correlation between the $B$- and $H$-band luminosity implies that the early-type dwarf galaxies are minimally affected by dust. This result is supported by the optical and NIR images of the IRSF dwarfs which reveal similar galaxy morphologies in the different wavelength bands. \\subsection{Stellar Masses of the Dwarf Galaxies} \\label{NIRmass} We have adopted the galaxy evolutionary models of \\citet{Bel2001} in determining the stellar $M/L$ ratio of the six dwarf galaxies. The linear coefficients given in Table~1 of \\citet{Bel2001} were used to derive the $M/L$ ratios. These coefficients were measured by adopting a formation epoch model (with bursts) and a Salpeter initial mass function (IMF). It is expected that the choice of the IMF will not introduce uncertainties greater than 10\\% in the NIR stellar $M/L$ ratio \\citep[see][]{Gal2002}. \\begin{table*} \\centering \\begin{minipage}{106mm} \\caption{Total stellar masses ($M_{stars}$) of six IRSF dwarfs. The columns represent, (2): Extinction-corrected, absolute $H$-band magnitudes; (5): \\ion{H}{i} masses of galaxies taken from the literature: NGC~59 and ESO~384-016 from \\citet{Bea2006}; NGC~5206 from \\citet{Cot97}. The upper \\ion{H}{i} mass limit is quoted for the galaxy NGC~5206. \\label{ML_ratio1}} \\begin{tabular}{@{}lcrlrr@{}} \\hline & {$M_H^0$} & {$L_H$} & {$M_{stars}$} & {$M_{\\ion{H}{i}}$} & \\\\ {Galaxy} & {(mag)} & {($\\times10^9L_\\odot$)} & {($\\times10^9M_\\odot$)} & {($M_\\odot$)} & {$M_{\\ion{H}{i}}/M_{stars}$} \\\\ {(1)} & {(2)} & {(3)} & {(4)} & {(5)} & {(6)} \\\\ \\hline NGC~3115 DW01 & -19.81 & $1.84\\pm0.76$ & $2.58\\pm1.17$ & -- &\\\\ NGC~59 & -17.99 & $0.34\\pm0.13$ & $0.48\\pm0.21$ & $1.5\\times10^7$ & 0.031 \\\\ LEDA~166099 & -16.89 & $0.12\\pm0.05$ & $0.17\\pm0.08$ & -- & \\\\ ESO~384-016 & -15.67 & $0.04\\pm0.01$ & $0.06\\pm0.03$ & $6.0\\times10^6$ & 0.100 \\\\ NGC~5206 & -19.50 & $1.38\\pm0.57$ & $1.93\\pm0.87$ & $5.5\\times10^5$ & $<0.001$ \\\\ UGCA~200 & -16.64 & $0.10\\pm0.04$ & $0.14\\pm0.06$ & -- & \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*} The $H$-band stellar $M/L$ ratio of the galaxies was computed using the linear relation: \\begin{equation} \\Upsilon_H = a_H + b_H(B-H)~. \\label{ML_eqn} \\end{equation} The stellar $M/L$ ratio was calculated for each of the IRSF dwarfs by substituting their individual $B$-$H$ colour into equation (\\ref{ML_eqn}). To increase the statistics by using a larger galaxy sample, galaxies having absolute magnitudes in the range of $-15.6\\la M_H \\la -19.8~\\rm mag$ in the \\citet{Kir2008} sample were also used in calculating the $M/L$ ratio. These magnitudes span the range of $H$-band absolute magnitudes observed for the IRSF dwarf galaxies (see Table~\\ref{2mass_comp}). A total of 33 dwarf galaxies were identified in this magnitude range. It should be noted that the dIrr galaxy AM~0521-343 in the \\citet{Kir2008} sample was excluded as the photometry of this galaxy is compromised by a bright foreground star. The $M/L$ ratio of the dwarf galaxies was computed by taking the mean of the individual ratios obtained for both the IRSF and \\citet{Kir2008} galaxy samples. This gives an $H$-band stellar $M/L$ ratio of $\\Upsilon_H = 1.4\\pm0.8$ where the error in the $M/L$ ratio represents the standard deviation of the mean. This value overlaps with the $M/L$ ratio of $\\Upsilon_H = 0.9\\pm0.6$ obtained by \\citet{Kir2008} for their full sample of 57 galaxies (consisting mostly of irregulars). The $M/L$ ratio of $\\Upsilon_H = 1.4\\pm0.8$ was used to calculate the stellar mass of each of the IRSF dwarfs. The extinction-corrected, absolute $H$-band magnitudes $M_H^0$ of the galaxies were converted into luminosity using the standard relation \\begin{equation} L_H = 10^{0.4(M_{H,\\odot} - M_H^0)} \\end{equation} where $M_{H,\\odot} = 3.35\\rm~mag$ is the $H$-band luminosity of the sun \\citep{Col1996}. The total stellar masses of the galaxies are listed in Table~\\ref{ML_ratio1}. The two brightest dwarfs (NGC~3115~DW01 and NGC~5206) in the sample ($M_H^0<-19\\rm~mag$), have the largest stellar masses which are of the order of $10^{9} M_\\odot$. \\citet{Puz2000} measure $(4.8\\pm2.3)\\times10^{10}\\Msun$ as the lower mass estimate of NGC~3115~DW01. This mass estimate was obtained from the kinematics of seven GCs in the galaxy. We have derived an $H$-band stellar mass of $(2.6\\pm1.2)\\times10^{9}\\Msun$ for NGC~3115~DW01 which is $\\sim25$ times lower than the mass estimate imposed by the kinematics of the seven GCs. The dynamical mass of NGC~3115~DW01 is measured out to a projected radius of $r\\simeq2.7\\arcmin$, making it almost twice as large as the radius in which the stellar mass of the galaxy was determined. This result suggests that NGC~3115~DW01 is a very DM dominated galaxy with no direct evidence of an extended DM halo \\citep[\\eg][]{Kle2002}. On the other hand, it is also possible that the GC system of \\citet{Puz2000} is not virialised around NGC~3115~DW01 giving the higher dynamical mass estimate for the galaxy. Given the \\ion{H}{i} masses for the galaxies NGC~59 and ESO~384-016, we are able to derive the \\ion{H}{i} gas-to-star mass fractions for these systems (see Table~\\ref{ML_ratio1}). The \\ion{H}{i} mass of NGC~59 corresponds to $\\sim$3\\% of its total stellar mass. This galaxy contains the largest amount of \\ion{H}{i} gas compared to the other early-type dwarfs in the Scl group \\citep{Bou2005,Bea2006}. The relatively large \\ion{H}{i} content of NGC~59 together with the ionized gas in the galaxy center \\citep{Ski03} suggests that this galaxy cannot be classified as a genuine dS0 galaxy. Instead, this detection supports the claim of \\citet{Bou2005} that this dwarf should rather be classified as a mixed-type dS0/Im galaxy as it exhibits the characteristics of both early and late-type galaxies. The galaxy ESO~384-016 shows a larger gas fraction to that detected for the star-forming galaxy NGC~59. This correlates with the \\ion{H}{i} mass to $B$-band luminosity ratios measured by \\citet{Bea2006} where the lower \\ion{H}{i} gas fraction was measured for NGC~59. These were found to be $M_{\\ion{H}{i}}/L_B = 0.21$ for ESO~384-016 while NGC~59 has $M_{\\ion{H}{i}}/L_B = 0.07$. ESO~384-016 represents one of four mixed-type dwarfs in the Cen~A group \\citep{Bou2007}. Only two of the mixed-type dwarfs (including ESO~384-016) have been detected in \\ion{H}{i}. These galaxies show $M_{\\ion{H}{i}}/L_B$ ratios similar to that found for mixed-morphology dwarfs in the Local Group \\citep{Ger1999,Bou2006}. In addition, the \\ion{H}{i} distribution of ESO~384-016 shows an eastern extension of the gas \\citep[see][]{Bea2006}. \\citet{Bou2007} suggest that the \\ion{H}{i} distribution of this galaxy is a result of mild ram pressure exerted by an intergalactic medium of density $\\rho_{\\rm IGM}\\sim 10^{-3}\\rm cm^{-3}$. The upper \\ion{H}{i} mass limit for NGC~5206 was used to estimate the gas fraction of this dwarf. The stellar mass is the main contributor to the baryonic mass of this galaxy. The \\ion{H}{i} gas content of NGC~5206 is found to be less than 0.1\\% of the total stellar mass. ", "conclusions": "Deep NIR $J$-, $H$- and $K_s$-band imaging were obtained for a sample of six early-type dwarf galaxies in the LV ($D\\la10\\rm~Mpc$). The galaxies are detected down to a SB limit of $\\mu\\simeq24\\rm~mag~arcsec^{-2}$ in the $J$- and $H$-bands, and $\\mu\\simeq23\\rm~mag~arcsec^{-2}$ in the $K_s$-band. The low SB galaxies LEDA~166099 and UGCA~200 are detected in the NIR for the first time. A detailed NIR photometric study was conducted for each of the dwarf galaxies to explore the properties and various characteristics of the old stellar component of the galaxy. The deep NIR observations allow for an accurate measure of the total magnitudes of even the faintest galaxies in our sample. The total magnitudes of the three brightest ($M_B\\la-15.5\\rm~mag$) galaxies NGC~3115~DW01, NGC~59 and NGC~5206 were compared to those obtained by 2MASS. For these galaxies, we find that 2MASS can underestimate the magnitudes by up to $\\la0.5\\rm~mag$. The remaining galaxies in our sample were not detected by 2MASS. These findings highlight the selection biases faced when using photometric data from the 2MASS galaxy survey. The structure of the underlying stellar component was determined by fitting an analytical function to the one-dimensional light profile of the galaxy. An exponential law was fit to the stellar disk of those galaxies hosting a nucleus. The effective radius $r_e$ and corresponding SB $\\mu_e$ associated with the disk component were derived for these galaxies. The light profiles of the non-nucleated galaxies were best modeled by a S\\'ersic law. The six dwarf galaxies are found to have similar NIR structure to early-type dwarf systems in the Virgo Cluster. A first indication of the extension of the relations between the structural parameters ($\\mu_e$ and $r_e$) and $H$-band luminosity to low luminosities ($M_H\\ga-18\\rm~mag$) is given by the faintest galaxies in our sample. This parameter space remains largely unexplored due to the lack of NIR data for low SB dwarf galaxies. The $J$-$K_s$ colour was derived for the individual dwarf galaxies to explore its variation with the $K_s$-band luminosity of the galaxy. We found that the six dwarfs exhibit almost constant $J$-$K_s$ which is independent of their luminosity. In addition, the nearly constant $J$-$K_s$ colours with luminosity are observed for early-type dwarf galaxies in the Virgo Cluster. These results suggest that the $J$-$K_s$ colour is not a strong tracer of the galaxy metallicity in early-type dwarf systems. The six dwarfs were also found to have typical $B$-$K_s$ colours to those seen in early-type dwarf galaxies which are independent of environment. The correlation between the $B$- and $H$-band luminosities implies that the dwarf galaxies are not strongly affected by dust attenuation so that similar galaxy morphologies are revealed at optical and NIR wavelengths. Finally, the stellar masses of the six galaxies were determined from the $H$-band observations. The dwarf galaxies are found to have stellar masses in the range of $10^8-10^{10}$\\MSUN. For the case of NGC~3115~DW01, the stellar mass was compared to its dynamical mass estimate indicating that this galaxy is DM dominated and a possible candidate for hosting a DM halo. The \\ion{H}{i} gas-to-star mass fractions were determined for those galaxies in which neutral hydrogen gas has been detected. The galaxies NGC~59 and ESO~384-016 show \\ion{H}{i} gas-to-star mass fractions of $\\ga3$\\% providing further support that these are mixed-type systems rather than pure dwarf elliptical (dE) or lenticular (dS0) galaxies." }, "1003/1003.3034_arXiv.txt": { "abstract": "We present results from the first three-dimensional radiation hydrodynamical calculations to follow the collapse of a molecular cloud core beyond the formation of the stellar core. We find the energy released by the formation of the stellar core, within the optically-thick first hydrostatic core, is comparable to the binding energy of the disc-like first core. This heats the inner regions of the disc, drives a shock wave through the disc, dramatically decreases the accretion rate on to the stellar core, and launches a temporary bipolar outflow perpendicular to the rotation axis that travels in excess of 50 AU into the infalling envelope. This outburst may assist the young protostar in launching a conventional magnetic jet. Furthermore, if these events are cyclic, they may provide a mechanism for intense bursts of accretion separated by long periods of relatively quiescent accretion which can potentially solve both the protostellar luminosity problem and the apparent age spread of stars in young clusters. Such outbursts may also provide a formation mechanism for the chondrules found in meteorites, with the outflow transporting them to large distances in the circumstellar disc. ", "introduction": "\\label{introduction} More than four decades ago, \\cite{Larson1969} performed the first numerical calculations of the collapse of a molecular cloud core to stellar core formation and beyond. These one-dimensional radiation hydrodynamical calculations revealed the main stages of protostar formation: an almost isothermal collapse until the inner regions become optically thick, the almost adiabatic formation of the first hydrostatic core (typical radius $\\approx 5$~AU and initial mass $\\approx 5$~M$_{\\rm J}$), the growth of this core as it accreted from the infalling envelope, the second collapse within this core triggered by the dissociation of molecular hydrogen, the formation of the stellar core (initial radius $\\approx 2$~R$_\\odot$ and mass $\\approx 1.5$~M$_{\\rm J}$), and, lastly, the long accretion phase of the stellar core to its final mass. Subsequent one-dimensional \\citep[e.g.][]{MasInu2000} and two-dimensional \\citep{Tscharnuter1987, Tscharnuteretal2009} calculations have not changed this qualitative picture substantially, although the latter have allowed the disc-like structure of rotating first cores to be studied. The first three-dimensional hydrodynamical calculations to follow the collapse to stellar core formation were performed more than a decade ago by \\cite{Bate1998}. These calculations of rotating molecular cloud cores examined the non-axisymmetric evolution of the first core and the second collapse phase. If the first core was rotating rapidly enough it was found to be dynamically unstable to a bar-mode leading to the formation of trailing spiral arms. Gravitational torques removed angular momentum and rotational support from the inner regions of the first core, quickening the onset of the second collapse and preventing fragmentation during the second collapse phase to form close binaries. Several subsequent studies have investigated this phenomenon in more detail \\citep*{SaiTom2006, SaiTomMat2008, MacInuMat2010}, some also including magnetic fields and finding outflows \\citep{Machidaetal2005, MacInuMat2006}. However, all these calculations used barotropic equations of state rather than solve the radiation hydrodynamical problem. \\begin{figure*} \\centering \\vspace{-2.7cm} \\includegraphics[width=15.5cm]{Fig1.pdf}\\vspace{-13.3cm} \\caption{The time evolution of the maximum (central) density (left two panels) and temperature (right panel). For the left and right panels, the lines are for cloud cores with $\\beta=0,5\\times 10^{-4},0.001,0.005,0.01$ from left to right. For the central panel, for each calculation the time in years has been zeroed when the density first exceeded $10^{-3}$~g~cm$^{-3}$ and the $\\beta=0$ calculation is plotted with a solid line, $\\beta=5\\times 10^{-4}$ with a dotted line, and the other calculations are indistinguishable. The free-fall time of the initial cloud core, $t_{\\rm ff}=1.8\\times 10^{12}$~s (56,500 yrs). Each calculation was performed with $10^6$ SPH particles. } \\label{evolutions} \\end{figure*} The first three-dimensional calculations including radiative transfer that followed collapse to the point of stellar core formation (but not beyond) were \\citet{WhiBat2006}, using the flux-limited diffusion approximation, and \\cite{Stamatellosetal2007}, using a radiative cooling approximation. Most recently, radiation magnetohydrodynamical calculations of cloud collapse have been performed \\citep{Tomidaetal2010}, but were stopped before the onset of the second collapse. In this paper, we report results from the first three-dimensional radiation hydrodynamical calculations to follow the collapse of rotating molecular cloud cores {\\it beyond} the formation of the stellar core. We find that the use of radiation hydrodynamics rather than a barotropic equation of state has little effect up until the formation of the stellar core. However, with radiative transfer, the energy released by the formation of the stellar core has a dramatic effect on the surrounding disc and envelope and launches a temporary outflow {\\it even in the absence of a magnetic field}. ", "conclusions": "\\label{conclusions} We have presented results from the first three-dimensional radiation hydrodynamical calculations to follow the collapse of a molecular cloud core beyond the formation of the stellar core. We find the evolution before the formation of the stellar core is very similar to that found in the past using barotropic equations of state. In particular, the evolution of the first hydrostatic core which, depending on its rotation rate, may be dynamically unstable to the growth of non-axisymmetric perturbations and the generation of spiral structure, is very similar to that found in barotropic calculations. As found in earlier calculations, a rapidly rotating first core actually evolves into a disc so that the disc actually forms {\\it before} the stellar core. However, we find that the evolution following the formation of the stellar core is {\\it qualitatively} different in radiation hydrodynamical calculations. In barotropic calculations, the formation of the stellar core deep inside the first core (or disc) has no effect on the surrounding disc because the temperature of the gas is simply set by the density of the gas. However, with radiation hydrodynamics, the energy released by the formation of the stellar core within the optically thick disc is similar to the binding energy of the disc. This heats the inner regions of the disc, drives a shock wave outwards through the disc, dramatically decreases the accretion rate on to the stellar core, and launches a bipolar outflow perpendicular to the rotation axis that can travel in excess of 50 AU out into the infalling envelope in less than 50 years. We speculate that such outflows may assist the young protostar in launching a conventional magnetic jet by clearing a path perpendicular to the disc and adding substantial thermal pressure to the force provided by the magnetic field. It may also be that such events are cyclic, occurring every time the accretion rate onto the protostar exceeds a certain level rather than simply being a one-off event associated with the formation of the stellar core. If so, such events may provide the mechanism for intense bursts of accretion separated by long periods of relatively quiescent accretion which may be necessary to solve the protostellar luminosity problem and the apparent age spread of young stars. Finally, such outbursts may provide another mechanism for the formation of chondrules in meteorites and the associated outflows may be able to transport them to large distances in the circumstellar disc." }, "1003/1003.1207_arXiv.txt": { "abstract": "{A star will become brighter and brighter with stellar evolution, and the distance of its habitable zone will become farther and farther. Some planets outside the habitable zone of a host star during the main sequence phase may enter the habitable zone of the host star during other evolutionary phases. A terrestrial planet within the habitable zone of its host star is generally thought to be suited to life existence. Furthermore, a rocky moon around a giant planet may be also suited to life survive, provided that the planet-moon system is within the habitable zone of its host star. Using Eggleton's code and the boundary flux of habitable zone, we calculate the habitable zone of our Solar after the main sequence phase. It is found that Mars' orbit and Jupiter's orbit will enter the habitable zone of Solar during the subgiant branch phase and the red giant branch phase, respectively. And the orbit of Saturn will enter the habitable zone of Solar during the He-burning phase for about 137 million years. Life is unlikely at any time on Saturn, as it is a giant gaseous planet. However, Titan, the rocky moon of Saturn, may be suitable for biological evolution and become another Earth during that time. For low-mass stars, there are similar habitable zones during the He-burning phase as our Solar, because there are similar core masses and luminosities for these stars during that phase. ", "introduction": "Typically, stellar habitable zone (HZ) is defined as a region near the host star where water at the surface of a terrestrial planet is in the liquid phase (e.g., Kasting et al. 1993; Franck et al. 2000; Noble et al. 2002; Jones et al. 2006). The inner HZ boundary is determined by the loss of water via photolysis and hydrogen escape. And the outer HZ boundary is determined by the condensation of carbon dioxide crystals out of the atmosphere (von Bloh et al. 2007). Previously, it was generally paid attention to the HZs of host stars during the main sequence (MS) phase, because the evolution from biochemical compounds to primary life needs long enough time. And a terrestrial planet around a host star covers a presumed heavy bombardment phase as on Earth (Jones 2004; Lal 2008), during the initial evolutionary stage of the host star. And the heavy bombardments make the temperature on the surface of a terrestrial planet become very high, which is not suitable for biological evolution. Only when the temperature on the surface of a terrestrial planet goes down, the evolution from biochemical compounds to primary life may start. These mean that the evolutionary age of a host star when life comes into being on a terrestrial planet around the host star is longer than the evolutionary timescale from biochemical compounds to primary life on the terrestrial planet. The evolutionary timescale of a host star during the MS phase is generally much longer than the timescale of heavy bombardments on a planet around the host star. Therefore, all the planets around a host star had become stable, when the MS phase of the host star is terminated. Without heavy bombardments, the temperature on the surface of a planet is mainly decided by the luminosity of its host star and the distance from the planet to the host star. Generally speaking, a host star becomes brighter and brighter during the post main sequence (post-MS) phase, mainly including subgiant branch phase, red giant branch (RGB) phase, horizontal branch (HB) phase and asymptotic giant branch (AGB) phase. And some terrestrial planets outside the HZ of a host star during the MS phase will enter the HZ of the host star during the post-MS phase. Once these terrestrial planets enter the HZ of the host star, the evolution from biochemical compounds to primary life may start immediately. The timescale of the evolution from biochemical compounds to prokaryote life in habitable environment may be as short as about 100 million years (Watson 2008). And all the evolutionary timescales of low-mass stars (defined in subsection 3.2) during the subgiant branch phase, the RGB phase and the HB phase are longer than 100 million years. Furthermore, positional changes of the HZ with stellar age are considered in future studies of long-term changes of the planetary biospheric conditions (Franck et al. 2000; Noble et al. 2002). Hence, it is meaningful to study the HZs of host stars during the post-MS phase. Life-seeds may also migrate from one terrestrial planet to another (Buccino et al. 2007). A 1.9 kg meteorite (ALH84001) discovered in 1984 in Allan Hills, Antarctica, is believed to have been blasted off the Martian surface due to an asteroid or comet impact about 15 million years ago (McKay et al. 1996; Hoyle and Wickramasinghe 1999a). There are also some meteorites found on Earth, whose source is the Moon (Warren 1994). These proves that a major meteorite impact on a planet or a moon would eject fragments into space and the fragments may finally fall down another planet. If there are life-seeds (such as spore, fungi, bacterial and plant seed) in the fragments, they can be ejected into space. Some microbes may also survive after a long journey through space from one planet to another planet (Joseph 2009; Wickramasinghe et al. 2009). And some plant seeds resist deleterious conditions found in space, e.g., ultra low vacuum, extreme temperatures and intense ultraviolet light. In a receptive environment, life-seeds could liberate a viable embryo, viable higher cells or a viable free-living organism (Tepfer and Leach 2006). Therefore, life can transfer from one planet to another through ejected fragment. In addition, life seeds may also be saved on comets and can be sent to remote places by the comets, and the life seeds may take root, provided that the conditions they visit on become habitable (Hoyle and Wickramasinghe 1999a, 1999b, 1999c; Wickramasinghe et al. 2009). Generally speaking, a terrestrial planet in the HZ of a host star is suited to life existence. Moreover, a rocky moon orbiting a giant planet or a brown dwarf could also be habitable, provided that the planet-moon system or the brown dwarf-moon system is within the HZ of the host star (Williams et al. 1997; Lal 2010). There are several rocky moons in solar system, such as Europa, Ganymede and Titan. These rocky moons will be habitable when Solar becomes brighter enough during the RGB phase or the HB phase. And there may be lives, which can generate from biochemical evolution or be transferred from other planets or moons, living on the rocky moons during that time. Furthermore, it is even thought that there may be life existence on these moons now, considering the potential of extremophiles to survive in highly inhospitable environments on Earth (Lal 2008). Using the boundary distances of HZ (Jones et al. 2006), we calculate the HZs of host stars during the post-MS phase. It is found that Mars' orbit, Jupiter's orbit and Saturn's orbit will enter the HZ of Solar during the subgiant branch phase for about 1.48 Gyr, the RGB phase for about 175 million years and the HB phase for about 137 million years, respectively. For low-mass stars, there are similar HZs as Solar during the HB phase. This means that a plant may also enter the HZ of a low-mass star during the HB phase, provided that the distance from the planet to the host star is just equal to the distance from Saturn to Solar. For intermediate-mass stars (defined in subsection 3.3), the HZs of host stars during the He-burning phase become farther and farther, with the stellar masses increasing. The outline of the paper is as follows: we describe our methods in Section 2, show our results in Section 3, present some discussions in Section 4, and then finally in Section 5 we give our conclusions. ", "conclusions": "Firstly, we calculate the HZ of Solar during the post-MS phase. It is found that Mars' orbit, Jupiter's orbit and Saturn's orbit will enter the HZ of Solar during the subgiant branch phase, the RGB phase and the HB phase, respectively. These mean that Titan will be suited to biological evolution, when Solar evolves to the HB phase. Secondly, we calculate the HZs of low-mass stars during the He-burning phase. As there are similar core masses, $T_{\\rm eff}$ and luminosities for low-mass stars during the HB phase, there are also similar HZs for them. Therefore, Saturn's orbit is also within the HZs for all low-mass stars during the HB phase. Thirdly, we calculate the HZs of intermediate-mass stars during the He-burning phase. It is found that the HZs of host stars become farther and father with the stellar masses increasing. Finally, we present discussions about the outward movement of planet's orbit, the UV radiation and the terrestrial planets within the HZs of low-mass stars during the HB phase. One may also send any special request to \\it{guojianpo1982@hotmail.com} \\normalfont{or} \\it{guojianpo16@163.com}\\normalfont{.}" }, "1003/1003.4278_arXiv.txt": { "abstract": "% Cygnus X-3 is a unique microquasar. Its X-ray emission shows a very strong 4.8-hour orbital modulation. But its mass-donating companion is a Wolf-Rayet star. Also unlike most other X-ray binaries Cygnus X-3 is relatively bright in the radio virtually all of the time (the exceptions being the quenched states). Cygnus X-3 also undergoes giant radio outbursts (up to 20 Jy). In this presentation we discuss and review the flaring behavior of Cygnus X-3 and its various radio/X-ray states. We present a revised set of radio/X-ray states based on Cygnus X-3's hardness-intensity diagram (HID). We also examine the connection of a certain type of activity to the reported AGILE/Fermi gamma-ray detections of Cygnus X-3. ", "introduction": "Cygnus X-3 represents one of the most unusual X-ray binaries to have been observed (see Bonnet-Bidaud \\& Chardin 1988 for a review). It is at a distance of $\\sim$ 9 kpc (Predehl et al. 2000) in the galactic plane and is heavily obscured at optical wavelengths. Cygnus X-3 does not fit well into any of the established classes of X-ray binaries. It has a 4.8-hour orbital period, observed both in the X-ray (Parsignault et al. 1972) and the infrared (Mason et al. 1986), which is typical of a low mass X-ray binary. But infrared spectroscopic observations (van Kerkwijk et al. 1992, Fender et al. 1999) indicate that the mass-donating star is a Wolf-Rayet star making the system a high mass X-ray binary. In addition, Cygnus X-3 undergoes giant radio outbursts and there is strong evidence of jet-like structures moving away from Cygnus X-3 at 0.3--0.9c (Molnar et al. 1988, Schalinski et al. 1995, Mioduszewski et al. 2001; see Fig. 1). The nature of the compact object is uncertain, but recent spectral studies (Szostek et al. 2008; hereafter SZM08 and Hjalmarsdotter et al. 2009) indicate that it may be a black hole. ", "conclusions": "We have discussed the various correlations and radio/X-ray states of Cyg X-3. From the construction of a unique HID for Cygnus X-3 with the radio flux as a third dimension we have gained additional insight into the nature and behavior of Cygnus X-3. We have identified three broad states and several sub-states which are delineated by intensity, X-ray hardness, and radio flux. We have also identified a ``new\" very soft state that we call the {\\it hypersoft} state. This state directly relates to the $\\gamma$-ray emission that has been observed from Cyg X-3 (Tavani et al. 2009, Abdo et al. 2009). This analysis lays the groundwork for a much better understanding of the workings and nature of Cygnus X-3." }, "1003/1003.0798_arXiv.txt": { "abstract": "The log-normal shape of the mass function for metal-poor halo globular clusters is proposed to result from an initial $M^{-2}$ power law modified rapidly by evaporation, collisions with clouds, and mutual cluster interactions in the dense environment of a redshift $z\\sim5-15$ disk galaxy. Galaxy interactions subsequently spray these clusters into the galaxy group environment, where they fall into other growing galaxies and populate their halos. Clusters forming later in $z\\sim2-5$ galaxies, and those formed during major mergers, produce metal-rich globulars. Monte Carlo models of evolving cluster populations demonstrate the early formation of a log-normal mass function for typical conditions in high-redshift galaxies. ", "introduction": "The globular cluster mass function (GCMF) in present-day galaxy halos is approximately log-normal with a peak $M_{\\rm p}\\sim10^{5.3}\\;M_\\odot$ \\citep[see reviews in][]{mc03,bs06}. The origin of this peaked distribution is not understood. Wherever massive dense clusters like globular clusters (GCs) are formed today, they have a power-law mass function like $dN/dM\\propto M^{-2}$, or a Schechter function with a similar power law at low mass and a cutoff at high mass \\citep[see review in][]{gieles09}. Consequently, one theory for halo GCs is that they begin with a power-law mass function and then lose their low-mass members through dispersal over a Hubble time \\citep{fr77, ot95,ee97,fz01}. The dispersal rate works out about right if the disruption process is thermal evaporation \\citep{mf08}. The problem with this model is that the GCMF does not vary with radius in several nearby galaxies \\citep{tam06,jor07}, and evaporation is expected to occur faster in the inner regions where tidal forces are larger, thereby shifting the peak toward higher masses there. A model in which GCs have a small range of pericenters can fix this problem \\citep{fz01}, but the resulting GC velocities disagree with observations in M87 \\citep{ves03b}. What is expected to drive the radial gradient in evaporation rate is a gradient in the tidal density \\citep{gb08}, which is the average GC density inside the GC tidal radius. This tidal density is difficult to observe directly and is not necessarily proportional to the average density inside the half-light radius, which is observed directly. The half-light density does not correlate well with radial distance. Noting this, \\citet{chan07} and \\citet{mf08} fit the GCMFs in M104 and the Milky Way to evolved Schechter functions for three bins of half-light density using GC evaporation rates proportional to the square roots of these densities. The results are consistent with faster evaporation, i.e., higher peak mass, at greater half-light density, regardless of galactocentric radius. \\citet{mf08} also fit the Milky Way GCMF for three bins of tidal density using evaporation rates given by that density. An alternative model is that GCs are born with a peaked mass function and then keep it over a Hubble time \\citep{ves03b,pg05}. Log-normal mass functions evolve somewhat self-similarly during evaporation, and their peak mass and width may even converge to the observed values \\citep{ves98}. The peak mass would also be uniform with galactocentric radius after a while. Initially peaked GCMFs could result from a lack of low-mass clouds in the early GC environment \\citep{pg07}. Another model considers variable star formation efficiencies with a greater probability for lower cluster masses to disperse when the gas leaves \\citep{parm08,baum08}. A third model is that low mass clusters were born with lower central concentrations and so evaporated more quickly than high mass clusters \\citep{vz03}. There is no direct evidence for any of these models because all known clusters today are born with power-law mass functions. For one of these models to be viable, the qualitative nature of cluster formation would have to be different in the early Universe. We consider here a model where halo GCs form with power-law mass functions that are quickly converted into peaked functions in very dense cloudy environments. The GCMF stays peaked and insensitive to environment thereafter \\citep{ves98}. This model involves the same physical process of cluster formation that is present today, i.e., gravitational collapse in giant gas complexes, and the cluster-forming cloud-cores are probably similar as well, considering that GC densities, masses, and IMFs are not unusual. What differed in young galaxies was a much denser and more turbulent interstellar medium (ISM) than we have in main galaxy disks today \\citep{fs06,g06,law07}. GC formation in $z>10$ dwarf galaxies was discussed by \\citet{br04}; bulge-GC formation in $z\\sim2$ galaxy clumps was discussed by \\citet{shap10}, and simulations of GC formation in young galaxy disks were made by \\citet{krav05}. The main point here is that when the density is high and the motions are fast, clusters should frequently collide with cloud clumps and other clusters during their first several hundred Myrs. These collisions destroy the lowest mass clusters and produce a log-normal GCMF. ", "conclusions": "We propose that old halo GCs formed by normal processes much like clusters form today, but that all of these processes occurred in young galaxies at very high densities, higher than today's ISM densities by factors of 10 to 100. Then collisional disruption and evaporation was rapid enough to produce a peaked mass function in the young galaxy. Subsequent evolution by evaporation in lower density environments preserved this function, although slight changes in peak mass and width probably occurred. Given this basic model, there are many possibilities for the delivery of these clusters into modern galaxies. The dense galaxies they formed in are most likely not the same as the galaxies or inner parts of the galaxies that currently host them. Their presence in galaxy halos today implies that they were delivered to the host in a non-dissipative way, perhaps by infall along with most of the host's other baryons. For example, the oldest GCs could have formed in small dense galaxies that formed as condensations in larger-scale cold gas flows. The cold flows make today's spiral disks \\citep{dekel09,atm09,keres09}, and the GCs along with other condensations in the flow populate modern halos. The small dense galaxies could also have collided before they entered the modern galaxy's potential well, freeing up the GCs which would then enter the well as free-floaters along with the other material. Some small galaxies are still adding GCs to modern halos \\citep[e.g.,][]{carraro07,gao07,casetti09,smith09}. The oldest GC populations in elliptical galaxies presumably formed in small dense galaxies and fell into spiral halos too in the same way, but then ended up in ellipticals after major mergers of the spirals. These oldest GCs would be the blue, metal-poor populations in spiral and elliptical halos. GCs formed during major mergers would be redder and more metal-rich \\citep{bs06}. The uniform properties of old halo globular clusters (color, density, IMF, peak GCMF mass, metallicity) follow from this model if we postulate that GCs are the first examples of star formation at high enough metallicity to make a normal IMF. Then stellar evolution will not remove excessive amounts of gas and cause the cluster to come unbound. Top-heavy IMFs, such as those thought to produce the carbon-enhanced metal-poor stars \\citep{tumlinson07,komiya09}, lose too high a fraction of their mass during stellar evolution to keep a cluster bound. Boundedness would require them to occupy the nuclei of small galaxies so they can retain their stellar wind material in the galactic potential well. Then contamination from subsequent generations of stars would also enrich them \\citep[e.g.,][]{marc07,bailin09}. For clusters born in small dense disks, the uniformity of both the clusters and the GCMFs suggest that the first epoch of normal IMFs occurred when galaxies were still small and dense, and this was long before today's spirals were assembled. The bulge GCs of modern spirals could have a similar origin as the halo GCs, but there is no similar constraint on the constancy of their GCMFs over a range of galactocentric radii. Also, the bulge GCs formed in metal-rich environments and they are now located deep in the potential wells of their current hosts. Thus they need not have formed in separate small galaxies. They could have formed in their hosts and settled to the center after energy dissipation. This is the model by \\citet{shap10}. They proposed that bulge GCs formed in the dense clumps of young massive $z\\sim2$ galaxies, much like we propose that metal-poor GCs formed in the dense clumps of young and small $z\\sim5-15$ galaxies. The local environment is about the same for each, namely high-density clumpy gas, so collisional disruption would have occurred quickly for the Shapiro et al. model too. The densities in the massive clumps of $z\\sim2$ spirals are not expected to have been as high as the densities in the disk clumps of smaller galaxies at higher redshift, however. Thus collisional disruption of low mass GCs may not have been as rapid for the Shapiro et al. model. Our model accounts for the number of halo GCs today if a high fraction of early star formation made bound clusters. Such a high fraction is expected when the ISM pressure and velocity dispersion are large \\citep{e08}. The space density of GCs today is $\\sim8$ Mpc$^{-3}$ \\citep{pm00}. There could have been a factor of $\\sim2$ more when they formed, considering evaporation \\citep{ves98}, but blue halo GCs are only half of the total. The co-moving space density of clumpy galaxies at $z\\sim1-2$ is $n_{\\rm g}\\sim2\\times10^{-3}$ Mpc$^{-3}$ \\citep{e07}. Each galaxy has $\\sim5$ clumps of stellar mass $\\sim10^8\\;M_\\odot$, which would be $N_{\\rm c}\\sim5\\times10^3$ clusters of mass $\\sim10^5\\;M_\\odot$ each. Thus the space density of massive clusters would have been $n_{\\rm g}N_{\\rm c}\\sim10$ Mpc$^{-3}$, with a factor-of-3 uncertainty either way. \\citet*{shap10} got the same result a different way. If metal-poor GCs formed in younger, smaller galaxies, then the $M_{\\rm gal}^{-2}$ cosmological galaxy mass function means that each log interval of galaxy mass contains the same total mass. Thus we would get these same $\\sim10$ GC per Mpc$^3$ for GC formation in galaxies 10 or 100 times less massive at higher redshifts. Modern examples of this model may occur in dense galactic nuclear regions where we predict a rapid destruction of low mass clusters because of heightened collision and evaporation rates. Acknowledgements: I am grateful to Enrico Vesperini and Tanuka Chattopadhyay for discussions about the GCMF, and to Dean McLaughlin for comments on the manuscript." }, "1003/1003.5497_arXiv.txt": { "abstract": "Fifty-nine quasars in the background of the Magellanic Clouds had brightness records monitored by the MACHO project during the years 1992 - 99. Because the circumpolar fields of these quasars had no seasonal sampling defects, their observation produced data sets well suited to further careful analysis. Following a preliminary report wherein we showed the existence of reverberation in the data for one of the radio-quiet quasars in this group, we now show that similar reverberations have been seen in all of the 55 radio-quiet quasars with adequate data, making possible the determination of the quasar inclination to the observer's line of sight. The reverberation signatures indicate the presence of large-scale elliptical outflow structures similar to that predicted by the Elvis (2000) and ``dusty torus'' models of quasars, whose characteristic sizes vary within a surprisingly narrow range of scales. More importantly the observed opening angle relative to the polar axis of the universal elliptical outflow structure present was consistently found to be on the order of 78 degrees. ", "introduction": "The wide-field imaging MACHO project produced brightness curves for approximately 6.5 million stars (Alcock et al, 1999), among which are brightness curves of 59 background quasars. An initial analysis of these (Schild, Lovegrove and Protopapas, 2009; PaperI) showed that these are well suited for systematic analysis because the Magellanic Clouds are circumpolar and the data do not suffer seasonal dropouts from when the quasar was too close to the sun. The initial analysis was inspired by the fact that the 25-year brightness record of the gravitationally lensed Q0957 quasar had shown evidence of reverberation structure (Thomson and Schild, 1997), that was easily detected in autocorrelation. The brightness amplitudes were typically 0.2 mag and the reverberations were on time scales of several hundred days. Because the pulse trains were frequently overlapping, it was difficult to recognize the pulse train from simple inspection. In our Paper I initial analysis we chose a single quasar, MACHO 13.5962.237, and demonstrated that the structure found in autocorrelation was realized by averaging together multiple wave trains, and the wave trains showed multiple peaks as had been already inferred from Q0957, and interpreted successfully by Schild (2005) as reverberations from the multiple surfaces that occur in the region of the dusty torus. Such structure had already been inferred by Elvis (2000) from a comprehensive analysis of a great wealth of published spectroscopic quasar data. A surprising discovery was that in addition to positive reverberations, in which an initial narrow brightening pulse is followed after approximately 80 days by several broader pulses, an almost equal number of $fading$ initial pulses with comparable widths and lags were found. A key finding was that the UV-optical reverberations originate at the same radial distance as the dusty torus already inferred to exist from emission line and infrared continuum reverberation. The existence of luminous outer structure was also inferred from microlensing studies of quasars, where failure of the standard luminous accretion disc model of Shakura and Sunyanaev (1973; S-S) to reproduce the observed microlensing in Q2237 has been a persistent problem. The first attempt to simulate the Q2237 microlensing with a S-S disc by Wyithe, Webster and Turner (2001) produced a prediction for a large microlensing brightening event that was to occur in 2001 but never happened. In a next attempt, Kochanek (2004) applied the model, but found that it could work only by seizing upon highly improbable intrinsic quasar brightening in coincidence with the microlensing. It was also noted that for expected cosmological velocities the inferred masses of the microlenses were sub-stellar, and like the Wyithe simulation, it predicted (Fig 10) the existence of 2.5 magnitude microlensing events, which have never been observed. A final attempt to apply the model by Eigenbrod et al (2008) encountered the same problems, and introduced a bias factor as a prior to force the calculation to fit the improving data set with stellar mass microlenses, but again predicted 2.5 magnitude microlensing events that have never been observed. A successful simulation by Vakulik et al (2007) abandoned the S-S disc model and simply allowed the calculation to find the size parameters describing the diameter of the microlensed inner structure, the masses of the microlenses, and the fraction of the total luminosity responsible for the observed brightness curves. These are the physical parameters to which the Q2237 microlensing would be most sensitive. This objective simulation produced a successful fit to the data with no bias factors, and with events that were highly probable from a quasar having only 1/3 of its UV-optical continuum radiation originating in the central region, within 6 gravitational radii, and the remaining 2/3 in a larger outer structure. The model predicts smaller amplitude microlensing events as observed, and determines that for expected cosmological velocities the microlenses have planetary mass. It had previously been demonstrated by Schild and Vakulik (2003) that such a population of microlenses was necessary to explain the rapid microlensing seen in another lensed quasar, Q0957. Additional microlensing results have hinted at extended UV-optical emission. Analysis of quadruple-image lenses by Pooley et al (2007) showed that the optical emission comes from a region approximately $ 3 \\times 10^{16}$ cm in radius, although this is ambiguous if the emitting region is an extended ring with inner radius and thickness dimensions, especially since this size was also inferred by SLR06 from estimates of the ring thickness (Fig. 1; $ \\delta r = 2 \\times 10^{16} cm$). Microlensing of the emission line region of SDSS 1004+4112 by Richards et al (2004) demonstrated that the characteristic size of the emitting region must be of order $10^{16} cm$, based upon the duration of an emission line microlensing event. We propose that the dusty torus, the broad-emission line region, the UV-optical continuum reverberation region, and the infrared reverberation region are all the same outer quasar structure. Henceforth we simply refer to it as the dusty torus. The central radius to this structure is of order $10^{17.3} cm$, its thickness parameter is $10^{16.3} cm$, and the height of of the luminous ring above and below the accretion disc plane is $10^{16.7} cm$. A cross-sectional view of the quasar with this picture is given as Fig. 1 of SLR06. The microlensing signature of such a structure has been simulated by Schild \\& Vakulik (2003) and by Abajas et al (2007). The preponderance of microlensing and reverberation evidence now supporting the dusty torus model of quasar structure now challenges the theoretical community. The Dusty Torus model was comprehensively discussed in 1993, (Antonucci 1993) but as yet has no explanation in a standard model of quasar structure. However a model proposed to explain the existence of the different observed quasar spectral states (radio loud - radio quiet) offers a possibility. While the standard black hole structure model has a magnetic field originating in the accretion disc, the relationship of the accretion disk to the size and location of the dusty torus is not clear. On the other hand Robertson and Leiter (2006; SLR06) have shown that a strong magnetic field anchored to a compact, highly red shifted, rotating central object called a magnetospheric eternally collapsing object (MECO) is also a viable solution to the Einstein-Maxwell field equations. Such an object would exhibit the effects of a co-rotation radius in the accretion disk to explain the observed quasar spectral states (Schild, Leiter, and Robertson, 2008; SLR08) and the location of the dusty torus would be associated with magnetic reconnection effects generated by the twisting of the dipole field lines to into toroidal near the light cylinder. In addition such a model would be radiatively inefficient near the central object\u2019s surface as a relativistic effect of the large intrinsic redshift there. Because our discovery that reverberation of the UV-optical continuum evidences important outer structure seriously challenges the standard black hole model, our report will focus on demonstrating that $all$ $quasars$ show evidence for such outer structure. We adopt the standard definition of quasars as luminous cosmological objects with a stellar appearance at 1 arcsec resolution, and broad blue-shifted emission lines. Previous work (SLR 06, 08) has made the connection that the reverberation radius is approximately the same as the central distance of the outer region originating the broad blue-shifted lines as an outflow wind. Our conclusion that the Elvis flow occurs where all quasars have reverberating luminous outer structure would seem to prove that such quasar structure is universal. \\section {Preliminary Processing of MACHO Quasar Brightness Data} Our preliminary processing has followed the procedures of Schild, Lovegrove, and Protopapas (2009; Paper 1). Raw data from standard V and R filters was corrected for CCD defects by simple removal of any 5 sigma data points. The data records were rebinned into uniformly spaced super-bins, with the number of super-bins equal to the number of observations, and all data within such a super-bin averaged. Super-bins containing no data were linearly interpolated over. Since the original brightness records had typically 600 data points spread over 2600 nights, a typical super-bin has time resolution of 4 observers' days. The timing of these observations was then rescaled for cosmological redshift. For the highest-redshift object in the survey, MACHO 208.15799.1085 at z = 2.77 there is insufficient observing time to produce more than one full reverberation pattern in quasar proper time. Therefore, we have excluded the source from further processing or plotting. For this reason our original sample of 59 quasars less 3 radio loud objects reduces to a sample of 55 radio quiet objects. In Paper I a quasi-periodicity with amplitude of 30 \\% was found in some but not all quasars. This variability has been called ``red-noise'' in some contexts, but it is well observed and real and in need of explanation, but we defer its further discussion. We remove this signature by forming a 300-day running boxcar smoothing algorithm over the binned and interpolated brightness record. An example of this procedure for MACHO quasar 13.5962.237 is given as Fig. 1 of Paper I. Following this preprocessing we have computed the autocorrelation function separately for the V and R filter data. The autocorrelation function always shows important structure, with a strong central peak having a brightness amplitude of order $ 30\\%$ and lags up to approximately 50 days, followed by a 100-day string of lags with negative autocorrelation, and then several positive autocorrelation peaks with only 10 \\% autocorrelation amplitude. In Paper 1 we have shown from a noise simulation that the peaks are real, since a realisitc noise simulation shows noise that occurs on time scales of our super-pixel resolution, approximately 4 observer's days, but which we do not see in our real data. We also show in Paper I that the autocorrelation peaks are confirmed to be real brightness enhancements with a reverberation pattern because we can co-add data segments containing the pattern and directly find a repeating wave form in brightness that confirms the autocorrelation pattern. Throughout this series of papers we presume that any structure found in brightness records for lag $t$ reflects quasar structure on size scale $ct$. In other words, we presume that the brightness features are caused by structure excited in the central region and propagating outward at the speed of light. \\section {The Mean Autocorrelation Due To Central Structure} In Fig. 1 we show a plot of the lag of the first minimum of the autocorrelation function separately for the V-filter and R-filter data. As a function of this we plot the (anti-)correlation amplitude. Based upon the theoretical error bars associated with the autocorrelation function, we see that the formal errors are comparable to the symbol size. We find in Fig. 1 that all the points lie in the lag range 30 - 260 days. The anti-correlation amplitude is between 0 and -0.5. The points scatter around this delimited area, with no trend evident. For a typical quasar with an anticorrelation peak lag of 100 days and an anticorrelation peak of -.25, we infer that significant anticorrelation exists and that an upper limit to the size of the excited central structure is 100 light days, or $2 \\times 10^{17} cm$. \\begin{figure*} \\begin{center} \\includegraphics{fig1.eps} \\caption{ Autocorrelation minimum values of 55 MACHO quasars as a function of the lag time of the autocorrelation minimum (the anti-correlation maximum) for the V data (top) and for the R data (bottom). The anti-correlation peak lag seems unrelated to the amplitude of the peak. The anti-correlation lag is surprisingly limited to only a small range, less than a factor ten, particularly in the bottom plot. One might interpret this anti-correlation peak as an indicator of a dark region around the central luminous UV-optical source, implying that all the MACHO quasars are approximately the same size.} \\label{fig. 1} \\end{center} \\end{figure*} In this context the siginificant result of our research was the finding that the radio quiet MACHO quasars are quite homogenous in their observed physical properties. It appears that the inner structures of these radio quiet are very similar. ", "conclusions": "In this report we have shown, from the analysis of 55 radio-quiet MACHO quasars, that reverberation patterns in the brightness of their UV-optical continuum indicate the presence of a universal large-scale elliptical outflow structure, similar to that predicted by the Elvis (2000) \"dusty torus\" model for quasars illuminated by their central regions. Of course there might be other models that could cause the pattern of brightness fluctuations found. While it may yet be argued that the autocorrelation features observed may be reproduced by some artificially tuned red-noise process, the fact that our simulations predict a 17\\% failure rate of red-noise to produce significant autocorrelation structure leads us to conclude that noise is a highly unlikely source of these patterns and that they are in fact real reverberation processes. The combined conclusions from Paper 1 that reverberation peaks originating at the region of the Elvis outflow structures can be identified from the autocorrelation function, and also can be seen by co-adding brightness record segments, suggests that they are real and easily recognized and studied. Similar structure has been reported in the radio loud gravitationally lensed Q0957 quasar by SLR06. Since the broad blue-shifted high excitation spectrum is a characteristic of all quasars, and since it is now understood that the dusty torus with Elvis outflow wind creates the emission lines, we expect the reverberation signature in brightness records to be present in radio loud quasars, and therefore to be a fundamental physical element of quasar structure that is easily studied with reverberation. The 55 radio quiet quasars in our sample showed a surprisingly small dispersion in the various structural properties of the radius of their broad line regions and of the polar opening angle of the outflow wind with respect to their polar axis. In particular we found that if a luminosity selected field quasar is observed to be in the radio quiet state, an average broad-line region of 545 light days size may be adopted for it, accurate to within a 50 \\% error. More importantly the average of the polar opening angle of the outflow wind (the complementary angle to the internal structure variable $\\epsilon$ of Schild (2005)) was found to be 78 degrees within a 3 degree variation. The very large 78 degree polar opening angles observed for the outflow winds in these quasars, which are similar to that seen for the centrally driven magnetic outflows seen in cataclysmic variables and young stellar objects, are difficult to explain in the context of standard black hole accretion disk models for quasars. Although many readers may find it surprising that our observations of the 55 MACHO quasars show such small differences in inferred reverberation region size, related to the broad-line region structure, this result is not in contradiction currently known observations of the difference between quasars and AGN. Quasars have long been known to have very similar emission line spectra, and no elaborate system of quasar spectral classification has emerged from 50 years of quasar studies. This is to be contrasted to the case of AGN which includes Seyfert galaxies that probably evidence galaxy interactions which bring a great deal more complexity to study. For the Seyferts, an elaborate classification system has in fact emerged. On the other hand quasars have long been recognized to have large differences in radio luminosity and high-energy X-ray emission, which have as yet not been closely associated with significant difference in their broad-line spectra." }, "1003/1003.5204_arXiv.txt": { "abstract": "{After leaving the main sequence, massive stars undergo complex evolution, which is still poorly understood. With a population of hundreds of OB stars, the starburst cluster Westerlund~1 offers an unparallelled environment to study their evolutionary tracks. } {We characterise a large sample of evolved OB stars in the cluster, with the aim of determining cluster parameters and place stars in an evolutionary sequence.} {We used the FORS2 instrument on the VLT to obtain intermediate-resolution spectroscopy over the range 5800--9000\\AA\\ of about a hundred stars selected as likely members of the cluster based on their photometry. We developed criteria for their spectral classification using only spectral features in the range observed. We discuss these criteria, useful for spectral classification of early-type stars in the GAIA spectral region, in the appendix. Using these criteria, we obtain spectral classifications, probably accurate to one subtype, for 57 objects, most of which had no previous classification or a generic classification.} {We identify more than 50 objects as OB supergiants. We find almost 30 luminous early-B supergiants and a number of less luminous late-O supergiants. In addition, we find a few mid B supergiants with very high luminosity, some of them displaying signs of heavy mass loss. All these stars form a sequence compatible with theoretical evolutionary tracks. In addition, two early B supergiants also show indication of heavy mass loss and may represent the evolutionary phase immediately prior to the Wolf-Rayet stage. We investigate cluster properties using the spectral types and existing photometry. We find that the reddening law to the cluster does not deviate strongly from standard, even though extinction is quite variable, with an average value $A_{V}=10.8$. Though evolutionary tracks for high-mass stars are subject to large uncertainties, our data support an age of $\\ga5$~Myr and a distance $d\\approx5$~kpc for Westerlund~1.} {The spectral types observed are compatible with a single burst of star formation (the age range is very unlikely to be $>1$~Myr). Westerlund~1 shows its potentiality as a laboratory for massive star evolution, which can be fulfilled by detailed study of the population presented here.} ", "introduction": "Massive stars play a crucial role in the dynamical and chemical evolution of galaxies, providing a major source of ionising UV radiation, mechanical energy, and chemical enrichment \\citep[e.g.,][]{mas03,lanza}. As they evolve from the main sequence (MS) towards the Wolf-Rayet (WR) phase, massive stars must shed most of their outer layers. Models predict that massive stars will evolve redwards at approximately constant $L_{{\\rm bol}}$, but there is an apparent dearth of yellow supergiants with $L_{{\\rm bol}}$ comparable to the brightest hot stars. The lack of yellow hypergiants above a certain luminosity defines the observational Humphreys-Davidson (HD) limit, which seems to imply that stars hit some kind of instability when they reach this area of the HR diagram, and then lose mass at a formidable rate. This phase of enhanced mass loss is generally identified with the luminous blue variable (LBV) stage, but there is no clear understanding of the actual evolution through the HR diagram of massive stars during the H-shell burning phase \\citep{hd94,vGen01}. \\defcitealias{main}{C05} There is general agreement that the WR stage corresponds to the He-core burning phase, and the fact that WR stars are very hot implies that massive stars do actually loop bluewards. But there is a complex zoo of transitional objects, comprising blue supergiants (BSGs), red supergiants (RSGs), yellow hypergiants (YHGs), LBVs and OBfpe/WNVL stars, whose identification with any particular evolutionary phase is a matter of guesswork. Understanding this evolution is, however, crucial because the mass loss during this phase completely determines the contribution that the star will make to the chemistry of the ISM and even the sort of post-supernova remnant it will leave. Unfortunately, massive stars are scarce and, as this phase is very short on evolutionary terms, examples of massive stars in transition are rare. For most of them, distances are unknown and so luminosities can be determined at best to order-of-magnitude accuracy, with the uncertainty feeding through to other parameters ($M_{*}$, $R_{*}$). As a consequence, the difficulty in placing these objects within an evolutionary sequence is obvious. The starburst cluster Westerlund~1 (Wd~1) may represent a unique laboratory for addressing these issues. With a mass $\\sim10^{5}\\:M_{\\sun}$ \\citep[][]{brandner,mt07}, it is young enough to contain a large number of OB stars, but old enough to contain stars in all the evolutionary stages. Large populations of WR stars and transitional massive stars \\citep[][henceforth C05]{wrs,main} have been identified. In this paper, we set their evolutionary context by characterising the population of moderately evolved massive stars in the cluster, i.e., objects that have started to move away from the MS, but have not yet reached the region of instability. With this information, we are able to constrain the cluster parameters and check the agreement with evolutionary tracks, which supports the idea of an (approximately) single-age population. ", "conclusions": "Using new classification criteria, developed in Appendix~A, we provide spectral types for almost 60 OB supergiants in Wd~1. Only a small fraction of them had previous accurate classifications. The stars leaving the cluster main sequence form a well populated clump at spectral types O9--B0, displaying a range of luminosities, which likely reflect a difference in initial rotational velocities and, perhaps, formation over a timespan. Sixteen luminous supergiants, with spectral types in the range B1--B4 range span the temperature range between this clump and the three blue hypergiants that bridge the gap with the A/F hypergiants. After a careful analysis of colour excesses, we find no strong reason to suspect that the extinction law towards Wd~1 deviates strongly from the standard, and we attribute the slightly divergent $E(B-V)$ values to the difficulty in obtaining precise $B$ photometry. The measured values of $E(V-I)$ and $E(V-R)$ show that the extinction varies strongly across the face of the cluster, with average values implying $A_{V}=10.8$. This value is fully consistent with the $A_{K_{{\\rm S}}}=1.13\\pm0.03$ derived by \\citet{brandner} and a standard law. The value $A_{K_{{\\rm S}}}=0.96$ found by \\citet{crowther}, though subject to much higher uncertainty, is also compatible. The dereddened magnitudes of the OB supergiants support a distance to Wd~1 similar to that found by \\citet{crowther} from the infrared magnitudes of WR stars, namely $DM=13.5$ ($d=5.0$~kpc). With this distance modulus, the intrinsic magnitudes of OB supergiants correspond closely to their spectral types according to typical calibrations. The shorter $DM=13.0$ ($d=4.2$~kpc) found by \\citet{kothes} is disfavoured (though their error bars reach $d=4.9$~kpc). The $DM=12.75$ ($d=3.5$~kpc) found by \\citet{brandner} is strongly disfavoured, as it would mean that all our Ia supergiants would have intrinsic magnitudes typical of Iab supergiants, implying masses $\\la25\\,M_{\\sun}$, very difficult to reconcile with their age derivation of 3.6~Myr. We suspect this to be due to the systematically younger ages that pre-MS isochrones indicate \\citep[cf.][]{naylor}. The picture drawn from the distribution in spectral types seems compatible with a single burst of star formation and the predictions of evolutionary models until the stars start to move quickly towards lower temperatures. Further studies, including the use of the population of unevolved stars to nail down the cluster age and detailed analysis of the evolved population (which will provide accurate stellar parameters and element abundances), will undoubtedly result in very firm constraints on the evolutionary tracks followed by massive stars after they leave the main sequence. The high binary fraction found amongst massive stars in Wd~1 \\citep{clark08,ritchie09a} will likely allow a better understanding of the role that binarity plays in this evolution. Furthermore, characterisation of binary properties (frequency, mass ratio, period distribution, etc.) in the homogeneous massive star population of Wd~1 can provide strong constraints on the formation mechanism of massive stars in such dense environments\\fnmsep\\footnote{Note, however, that at an age $\\sim5$~Myr, Wd~1 must have already undergone dynamical evolution. Indeed, \\citet{brandner} find some evidence for mass segregation, which could partially erase the signature of the formation mechanism.}, as the different formation mechanisms proposed will predict different outcomes \\citep[for a review]{zy07}. Long-term spectroscopic monitoring of binaries in the cluster is underway \\citep{ritchie09a}, and will deliver some of the observational characteristics." }, "1003/1003.0674.txt": { "abstract": "\\noindent We analyze Aharonov-Bohm radiation of charged fermions from oscillating solenoids and cosmic strings. We find that the angular pattern of the radiation has features that differ significantly from that for bosons. For example, fermionic radiation in the lowest harmonic is approximately isotropically distributed around an oscillating solenoid, whereas for bosons the radiation is dipolar. We also investigate the spin polarization of the emitted fermion-antifermion pair. Fermionic radiation from kinks and cusps on cosmic strings is shown to depend linearly on the ultraviolet cut-off, suggesting strong emission at an energy scale comparable to the string energy scale. ", "introduction": "The Aharonov-Bohm (AB) interaction \\cite{Aharonov:1959fk,AlfordWilczek} between charged particles and thin magnetic fluxes is of much interest as it provides a physical consequence of a pure gauge field with vanishing field strength but non-trivial topology. Further, the physical effects emerge only in quantum theory and hence the AB interaction provides an example of a quantum, topological interaction. The classic Aharonov-Bohm setup involves scattering an electron off a solenoid, with non-trivial scattering obtained even for an arbitrarily thin solenoid, whereby the electron is exclusively localized in a region of vanishing magnetic field. A novel feature of the scattering is a periodic dependence of the scattering cross-section on the magnetic flux through the solenoid. If $\\Phi$ denotes the magnetic flux in the solenoid and $e$ the electron charge, the cross-section is proportional to $\\sin^2(\\pi\\epsilon)$ where $\\epsilon \\equiv e\\Phi/2\\pi$. The classic AB setup was recently extended in another direction \\cite{AB_Bosons}, where it was shown that an oscillating solenoid in vacuum can produce charged particle-antiparticle bosons from the vacuum due to the AB interaction. The AB radiation rate also has the characteristic $\\sin^2(\\pi \\epsilon)$ dependence on the magnetic flux. AB radiation is relevant to the evolution of cosmic strings, which are similar to solenoids, except the magnetic flux within them is massive, unlike electromagnetic fluxes in laboratory solenoids. Moreover, the gravitational analog of the AB effect can cause cosmic strings to emit light, even if the fields composing the cosmic string are unrelated to electromagnetic fields \\cite{Garriga:1989bx,AB_Bosons}. In this paper we will investigate {\\it fermionic} AB radiation. One motivation is that the electron is a fermion. Hence fermionic AB radiation is what is relevant to oscillating solenoids. The investigation is also relevant, for instance, to neutrino emission from cosmic strings by the AB process. A second motivation is that the spin of the fermion adds another degree of freedom to the emission and the polarization properties of the radiation are of interest. Our results show a significant difference between AB radiation of bosons and fermions. For example, if a solenoid aligned with the $z$-axis oscillates along the $x$-direction, the angular distribution of bosonic AB radiation is peaked in the $y$-direction. Fermionic AB radiation, however, is (approximately) isotropically distributed. The outline of the paper is as follows. In Secs.~\\ref{abfermionpairproduction} through \\ref{cosmicstringloops}, we use conventional interaction picture perturbation theory to calculate the fermion-antifermion pair production rate in the small AB phase ($\\epsilon$) limit. We consider AB radiation from an infinite straight solenoid oscillating perpendicular to its length in Sec.~\\ref{infinitestraightstring}, a cosmic string loop with kinks in Sec.~\\ref{kinkyloops}, and a cosmic string loop with cusps in Sec.~\\ref{cuspyloops}. In Sec.~\\ref{movingframesperturbationtheory}, we solve the problem using a different technique that does not assume that $\\epsilon$ is small. We call this the ``moving frames'' scheme and use it to obtain the $\\sin(\\pi\\epsilon)$ dependence of the radiation on the AB phase, provided the motion of the solenoid is slow. We conclude in Sec.~\\ref{conclusions} and describe our conventions in Appendix~\\ref{conventions}. ", "conclusions": "\\label{conclusions} We have solved for fermionic radiation first from oscillating electromagnetic solenoids and then from cosmic string loops. For the solenoid we have done the calculation in two different ways, first using a small AB phase approximation, and second by considering slowly moving solenoids. We have evaluated the angular distribution of the fermionic radiation from the solenoid, and the total power emitted from cosmic string loops and cusps. Our results can be compared to the results of Ref.~\\cite{AB_Bosons}. The total power emitted in bosons and fermions is very comparable. For example, both are proportional to $\\epsilon^2 v_0^2 \\Omega^2$ for the lowest harmonic of the oscillating solenoid (see Eq.~\\eqref{pprodrateleq1}). However, the angular distributions of the radiation in the two cases are quite distinct. To highlight the difference, we show the angular distribution in both cases for the lowest harmonic emission from an oscillating solenoid in Fig.~\\ref{angdistn}. We also find that the fermion and antifermion are preferably emitted in opposite helicity states and discuss the spin distribution. Just like in the bosonic case, fermionic AB radiation from kinks and cusps on cosmic strings is ultra-violet divergent for massless fermions, with a linear dependence on the cut-off. This may translate into a significant amount of radiation of neutrinos from strings with which neutrinos have an AB interaction. Our results ought to apply to the low energy end of the emission spectrum of electrons from idealized solenoids, where the wavelengths of the particle pairs are much longer than the diameter of the solenoid. A more realistic theoretical investigation would have to take into account the finite width of the solenoid itself. In Ref.~\\cite{AB_Bosons}, the gravitational analog of AB radiation was also discussed. Via the same analogy, we also expect cosmic strings to radiate fermions. We leave that calculation for future work." }, "1003/1003.3108_arXiv.txt": { "abstract": "{This is the second paper of a series in which we attempt to put constraints on the flattening of dark halos in disk galaxies. For this purpose, we observe the \\HI\\ in edge-on galaxies, where it is in principle possible to measure the force field in the halo vertically and radially from gas layer flaring and rotation curve decomposition respectively. To calculate the force fields, we need to analyse the observed XV diagrams to accurately measure all three functions that describe the planar kinematics and distribution of a galaxy: the radial \\HI\\ surface density, the rotation curve and the \\HI\\ velocity dispersion. In this paper, we discuss the improvements and limitations of the methods previously used to measure these \\HI\\ properties. We extend the constant velocity dispersion method to include determination of the \\HI\\ velocity dispersion as a function of galactocentric radius and perform extensive tests on the quality of the fits. We will apply this `radial decomposition XV modelling method' to our \\HI\\ observations of 8 \\HI\\ rich, late-type, edge-on galaxies in the third paper of this series. } {} ", "introduction": "In paper I \\citep{ofk2008} in this series we presented \\HI\\ observations of a sample of 8 edge-on, \\HI\\ rich, late-type galaxies. The aim of the project has been described there in detail. Briefly, we attempt to put constraints on the flattening of dark halos around disk galaxies by measuring the force field of the halo vertically from the flaring of the \\HI\\ layer and radially from rotation curve decomposition. For the vertical force field we need to determine in these galaxies both the velocity dispersion of the \\HI\\ gas (preferably as a function of height from the central plane of the disk) and the thickness of the \\HI\\ layer, all of this as a function of galactocentric radius. In addition we also need to extract information on the rotation curve of the galaxy and the deprojected \\HI\\ surface density, as a function of galactocentric radius. In this second paper in the series we will review earlier determinations of these properties and then describe the methods we have developed to analyse our sample. \\begin{table*}[t] % \\caption[Summary of previous measurements of gas velocity dispersion in face-on galaxies] {Summary of previous observations of gas velocity dispersion in face-on galaxies.} \\label{tab:ch4-vel-disp-meas-lowi} \\begin{footnotesize} \\begin{tabular}{lcccccccccl} \\hline Galaxies & i & Type & H{\\scriptsize I} Velocity & \\multicolumn{2}{c}{Resolution} & \\multicolumn{3}{c}{H{\\scriptsize I} radius} & Dist. & Reference \\\\ \\cline{5-9} & & & dispersion & Spectral & Spatial & & & & & \\\\ \\cline{5-6} & \\dg & & & \\kms & kpc & kpc & $R_{25}$ & $\\theta_{maj}$ & Mpc & \\\\ \\hline\\hline NGC3938 & $10\\pm5$ & ScI & constant$^{\\rm h}$ & 8.2$^{\\rm a}$ & $1.2 \\times 1.7$ & 13.1 & 1.7 & 7.5 & 10 & \\citet{vdks1982c} \\\\ NGC628 & $5\\pm5$$^{\\rm e}$ & ScI & falling$^{\\rm h}$ & 8.2$^{\\rm a}$ & $2.9 \\times 2.9$ & 40.7 & 2.8 & 14.0 & 10 & \\citet{svdk1984} \\\\ NGC1058 & $8\\pm2$$^{\\rm f}$ & Sc{\\sc II-III} & constant$^{\\rm h}$ & 8.2$^{\\rm a}$ & $2.2 \\times 2.2$$^{\\rm c}$ & 18.9 & 4.3 & 8.7 & 10$^{\\rm c}$ & \\citet{vdks1984} \\\\ NGC1058 & $8\\pm2$$^{\\rm f}$ & Sc{\\sc II-III} & falling$^{\\rm h}$ & 2.58$^{\\rm a}$ & $2.1 \\times 2.0$$^{\\rm d}$ & 14.1 & 3.2 & 6.7 & 10$^{\\rm d}$ & \\citet{dhh1990} \\\\ NGC6946 & 34 & Scd/Sc{\\sc I} & falling$^{\\rm h}$ & 8.25$^{\\rm a}$ & $0.7 \\times 0.8$ & 17.2 & 2.0 & 21.2 & $5.9$$^{\\rm b}$ & \\citet{bv1992} \\\\ NGC5474$^{\\rm g}$ & $21^{+4}_{-6}$ & ScdIV & falling$^{\\rm h}$ & 5.2$^{\\rm a}$ & $1.2 \\times 1.1$ & 10.5 & 4.6 & 8.6 & 7 & \\citet{rdh1994} \\\\ NGC1058 & $8\\pm2$$^{\\mathrm f}$ & Sc{\\sc II-III} & falling$^{\\rm h}$ & 2.58$^{\\rm a}$ & $1.5 \\times 1.5$ & 21.8 & 5.0 & 15.0 & 10 & \\citet{pr2006} \\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] Hanning smoothed. \\item[$^{\\mathrm{b}}$] Distance shown is mean distance calculated from brightest stars in NGC6946 group. \\item[$^{\\mathrm{c}}$] To facilitate comparison with the other observations for NGC1058, we have used the adopted distance of $10$ Mpc used by \\citet{pr2006}, rather than the \\citet{vdks1984} preferred value of $9$ Mpc. \\item[$^{\\mathrm{d}}$] To facilitate comparison with the other observations for NGC1058, we have used the adopted distance of $10$ Mpc used by \\citet{pr2006}, rather than the \\citet{dhh1990} preferred value of $10.2-14.5$ Mpc. \\item[$^{\\mathrm{e}}$] Warps from $\\approx5^{\\circ}$ within the optical disk to $i\\approx0^{\\circ}$, then back to $i\\approx10^{\\circ}$ at the H{\\scriptsize I} edge \\citep{svdk1984}. \\item[$^{\\mathrm{f}}$] Warps from $\\approx8\\pm2^{\\circ}$ within the optical disk to exactly face-on at the H{\\scriptsize I} edge \\citep{vdks1984} \\item[$^{\\mathrm{g}}$] Classified as peculiar, and known to be interacting with its much larger neighbour M101 which is only $90$ kpc distance assuming a distance of $7$ Mpc \\citep{rdh1994}. \\item[$^{\\mathrm{h}}$] Direct measurement of the velocity dispersion of the broadening of the vertical velocity dispersion from a near face-on disk. Measurements corrected for instrumental broadening and inclination, but not beam-smearing. \\end{list} \\end{footnotesize} \\end{table*} To meet our goals we need to measure the vertical structure of galaxies out to the lowest surface densities possible and use galaxies with relatively large fractions of their mass in their halos, i.e. high total mass-to-light ratios galaxies. For that purpose we defined a sample of nearby, \\HI\\ rich, late-type edge-on galaxies. From paper I we recall that for a vertically isothermal gas sheet with a vertically Gaussian density distribution, the gradient of the total force in the vertical direction can be written as \\begin{equation} \\label{eq:ch1-hydro} \\frac{\\partial K_z}{\\partial z} = - \\frac{\\sigma_{v,g}^2}{({\\rm FWHM}_{z,g}/2.35)^2}, \\end{equation} where $\\sigma_{v,g}$ is the vertical velocity dispersion of the gas and FWHM$_{z,g}$ the gas layer thickness. The vertical force field comes from the gas itself, plus the dark halo and the stars of the disk. In regions where the halo is a significant contributor to $K_z$, the \\HI\\ flaring will be smaller for the more flattened halos with a smaller axis ratio $q=c/a$ (the ratio of the halo polar axis $c$ to its major axis in the galactic plane $a$). In the present paper we will first present a new method to accurately determine the rotation curve, deprojected \\HI\\ surface density and \\HI\\ velocity dispersion at all radii in a edge-on gas disk. The motivation for this measurement method came from the strong dependence in Eqn.~(2) in paper I of the derived dark halo flattening $q$ on the measured \\HI\\ velocity dispersion and vertical \\HI\\ gas disk flaring. The superposition of velocity profiles from many radii in each sightline through an edge-on \\HI\\ disk tends to cause an overestimate of the velocity dispersion with most measurement methods, as explained below. Measuring the radial flaring profile requires a model of the galaxy rotation and face-on surface density; this also necessitates high accuracy rotation curve measurement and \\HI\\ surface density deprojection. The \\HI\\ velocity dispersion has not previously been measured systematically in edge-on galaxies. Most measurements of the \\HI\\ velocity dispersion have been conducted on face-on galaxies, and these measurements were largely conducted 15-20 years ago with relatively low spatial resolution FWHM$_{\\theta}$= 1-3 kpc. In Table~\\ref{tab:ch4-vel-disp-meas-lowi} and \\ref{tab:ch4-vel-disp-meas-highi}, we show the resolution of each observation used to measure the velocity dispersion of near face-on and near-edge-on galaxies, respectively. The radial structure of the \\HI\\ velocity dispersion has been measured in only 6 galaxies -- 5 from near face-on galaxies and one from an edge-on galaxy. The 5 face-on galaxies were observed at $\\sim$1 kpc or larger resolution. The uncertain inclination of low resolution face-on \\HI\\ observations can cause additional broadening, due to projection of the gas rotation along the line-of-sight. Galaxies inclined slightly away from face-on will suffer beam smearing of the gradients in the projected rotation as well as beam smearing of the gradients of the intrinsic velocity dispersion. \\begin{table*}[t] % \\caption[Summary of previous measurements of gas velocity dispersion in highly-inclined galaxies]{Summary of previous observations of gas velocity dispersion of highly inclined galaxies.} \\label{tab:ch4-vel-disp-meas-highi} \\begin{footnotesize} \\begin{tabular}{lcccccccccc} \\hline Galaxies & i & Type & H{\\scriptsize I} Velocity & \\multicolumn{2}{c}{Resolution} & \\multicolumn{3}{c}{H{\\scriptsize I} radius} & Distance & Ref. \\\\ \\cline{5-9} & & & dispersion & Spectral & Spatial & & & & & \\\\ \\cline{5-6} & \\dg & & & \\kms & kpc & $R_{25}$ & kpc & $\\theta_{maj}$ & Mpc & \\\\ \\hline\\hline NGC4244 & $84.5$$^{\\rm c}$ & Scd & constant$^{\\rm d}$ & 5.2$^{\\rm a}$ & 0.17 & 630 & 11.0 & 2.0 & 3.6 & \\citet{olling1996a} \\\\ ESO142-G24 & 90$^{\\rm b}$ & Scd & constant$^{\\rm e}$ & 6.6$^{\\rm a}$ & 3.8 & 19.0 & 1.6 & 5.0 & 26.6 & \\citet{kvdkdb2004} \\\\ ESO157-G18 & 90$^{\\rm b}$ & Scd & constant$^{\\rm e}$ & 6.6$^{\\rm a}$ & 2.6 & 11.1 & 1.3 & 4.3 & 18.1 & \\citet{kvdkdb2004} \\\\ ESO201-G22 & 90$^{\\rm b}$ & Sc & constant$^{\\rm e}$ & 6.6$^{\\rm a}$ & 9.0 & 34.1 & 1.5 & 3.8 & 55.4 & \\citet{kvdkdb2004} \\\\ ESO240-G11 & 90$^{\\rm b}$ & Sc & constant$^{\\rm e}$ & 6.6$^{\\rm a}$ & 1.7 & 41.5 & 1.4 & 24.2 & 38.1 & \\citet{kvdkdb2004} \\\\ ESO263-G15 & 90$^{\\rm b}$ & Sc & constant$^{\\rm e}$ & 6.6$^{\\rm a}$ & 5.1 & 24.8 & 1.5 & 4.8 & 35.2 & \\citet{kvdkdb2004} \\\\ ESO269-G15 & 90$^{\\rm b}$ & Sc & constant$^{\\rm e}$ & 6.6$^{\\rm a}$ & 7.5 & 19.9 & 1.0 & 2.7 & 46.7 & \\citet{kvdkdb2004} \\\\ ESO321-G10 & 90$^{\\rm b}$ & Sa & constant$^{\\rm e}$ & 6.6$^{\\rm a}$ & 7.3 & 14.5 & 1.2 & 2.0 & 41.0 & \\citet{kvdkdb2004} \\\\ ESO416-G25 & 90$^{\\rm b}$ & Sb & constant$^{\\rm e}$ & 6.6$^{\\rm a}$ & 17.7 & 36.7 & 1.6 & 2.1 & 67.0 & \\citet{kvdkdb2004} \\\\ ESO435-G14 & 90$^{\\rm b}$ & Sc & constant$^{\\rm e}$ & 6.6$^{\\rm a}$ & 7.7 & 18.9 & 1.5 & 2.5 & 35.2 & \\citet{kvdkdb2004} \\\\ ESO435-G25 & 90$^{\\rm b}$ & Sc & constant$^{\\rm e}$ & 6.6$^{\\rm a}$ & 2.2 & 36.1 & 1.2 & 16.3 & 32.4 & \\citet{kvdkdb2004} \\\\ ESO435-G50 & 90$^{\\rm b}$ & Sc & constant$^{\\rm e}$ & 6.6$^{\\rm a}$ & 7.1 & 13.1 & 1.4 & 1.9 & 35.2 & \\citet{kvdkdb2004} \\\\ ESO446-G18 & 90$^{\\rm b}$ & Sb & constant$^{\\rm e}$ & 6.66$^{\\rm a}$ & 12.2 & 31.1 & 1.3 & 2.5 & 64.1 & \\citet{kvdkdb2004} \\\\ ESO564-G27 & 90$^{\\rm b}$ & Sc & constant$^{\\rm e}$ & 6.66$^{\\rm a}$ & 10.2 & 28.7 & 1.6 & 2.8 & 29.5 & \\citet{kvdkdb2004} \\\\ NGC5170 & 90$^{\\rm b}$ & Sc & constant$^{\\rm e}$ & 6.66$^{\\rm a}$ & 7.4 & 31.9 & 1.3 & 4.3 & 20.9 & \\citet{kvdkdb2004} \\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] Hanning smoothed. \\item[$^{\\mathrm{b}}$] From LEDA database. \\item[$^{\\mathrm{c}}$] The unwarped inner disk has an inclination of $i=84.5^{\\circ}$, outside of which there is a small warp to $i=82.5\\pm1^{\\circ}$, before warping back to the plane of the inner disk at large radii \\citep{olling1996a}. \\item[$^{\\mathrm{d}}$] Direct measurement of the velocity dispersion of the broadening of the extreme velocity envelope of the XV map of an edge-on galaxy. Measurements corrected for instrumental broadening, projection and beam-smearing, via Olling's method (see Sect.~\\ref{sec:olling-RC-disp-meth}). \\item[$^{\\mathrm{e}}$] Iterative envelope fitting of the XV map of an high-inclination galaxy integrated over the galaxy minor axis \\citep{kvdkdb2004}. This method was unable to measure radial structure of the velocity dispersion, as the method did not include beam correction and the galaxies in the sample were observed with very large beams. \\end{list} \\end{footnotesize} \\end{table*} \\begin{table*}[t] % \\caption[Summary of previous measurements of rotation curves in edge-on galaxies]{Summary of previous measurements of rotation curves in edge-on galaxies ($i > 80^{\\circ}$).} \\label{tab:ch4-RC-meas} \\begin{footnotesize} \\begin{tabular}{lp{4.0cm}ccccl} \\hline Galaxies & Method & Line & \\multicolumn{2}{c}{Resolution} & Distance & Reference \\\\ \\cline{4-5} & & & Spectral & Spatial & & \\\\ & & & \\kms & kpc & Mpc & \\\\ \\hline\\hline UGC7321 & Peak Flux$^{\\rm a}$ & H{\\scriptsize I} & 5.2 & 0.8 & 10 & \\protect\\citet{um2003} \\\\ \\hline NGC891 & Envelope Tracing$^{\\rm b}$ & H{\\scriptsize I} & 27 & 1.7 & 9.5$^{\\rm f}$ & \\protect\\citet{sa1979} \\\\ NGC891$^{\\ddagger}$ & Envelope Tracing$^{\\rm b}$ & H{\\scriptsize I} & 33 & 0.7 & 9.5$^{\\rm f}$ & \\protect\\citet{ssvdh1997} \\\\ 9 dwarfs & Envelope Tracing$^{\\rm b,e}$ & H{\\scriptsize I} & 4-8 & 0.5-2.8 & 3.4-19 & \\protect\\citet{swaters1999} \\\\ 4 spirals & Envelope Tracing$^{\\rm b}$ & HI & 20-41 & 0.9-2.3 & 9.5-15.5 & \\protect\\citet{sofue1996} \\\\ 4 LSBs$^{\\#}$ & Envelope Tracing$^{\\rm b}$ & H{\\scriptsize I} & 4-8 & 0.7-2.7 & 5-19 & \\protect\\citet{dbb2002} \\\\ \\hline NGC891$^{\\dagger}$ & Fixed $\\sigma_v$ \\& Parametric $v(R)$ XV Modelling$^{\\rm b,c}$ &H{\\scriptsize I} & 27 & 1.7 & 9.5 & \\protect\\citet{vdkruit1981} \\\\ NGC5023 & Fixed $\\sigma_v$ \\& Parametric $v(R)$ XV Modelling$^{\\rm b,c}$ & H{\\scriptsize I} & 17 & 0.8 & 7.9 & \\protect\\citet{bsvdk1986} \\\\ \\hline NGC3079 & Iterative Envelope Tracing$^{\\rm b}$ & CO & 5.2 & 0.12 & 15.5 & \\protect\\citet{ts2002} \\\\ \\hline 14 spirals & Fixed $\\sigma_v$ 1D Gaussian Envelope Fitting$^{\\rm b}$ & H{\\scriptsize I} & 6.6 & 1.7-17.7 & 18-67 & \\protect\\citet{kvdkdb2004} \\\\ \\hline 24 spirals & Fixed $\\sigma_v$ 2D Gaussian Envelope Fitting$^{\\rm b}$ & H{\\scriptsize I} & 5.0-20.1 & 0.8-4.8 & 5-33 & \\protect\\citet{grsk2002} \\\\ \\hline 8 spirals$^{\\star}$ & Fixed $\\sigma_v$ XV Modelling & H{\\scriptsize I} & 6.6-33.0 & 0.9-9.8 & 9.5-54$^{\\rm f}$ & \\protect\\citet{kvdk2004} \\\\ \\hline NGC4244 & Corrected Gaussian Envelope Fitting$^{\\rm d}$ & H{\\scriptsize I} & 5.2 & 0.17 & 3.6 & \\protect\\citet{olling1996a} \\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] Corrected for instrumental broadening \\& gas velocity dispersion using the ``equivalent rectangular measure'' \\citep{sbs1966}. \\item[$^{\\mathrm{b}}$] Corrected for instrumental broadening \\& gas velocity dispersion by quadratic subtraction, see Eqn.~(\\ref{eq:ch4-instr-broadening-corr}). \\item[$^{\\mathrm{c}}$] Uses the standard \\citet{fe1980} parametric rotation curve form, constant $\\sigma_z$, and a 2D exponential surface density to model the XV diagram. \\item[$^{\\mathrm{d}}$] Corrected for broadening \\& beam-smearing (see Olling's method in Sect.~\\ref{sec:olling-RC-disp-meth}). \\item[$^{\\mathrm{e}}$] Asymmetric drift correction applied. \\item[$^{\\mathrm{f}}$] Using distance measured by \\citet{vdkruit1981} for NGC891. \\item[$^{\\mathrm{\\dagger}}$] Re-analysis of NGC891 data observed by \\citet{sa1979}. \\item[$^{\\mathrm{\\ddagger}}$] New deeper WSRT NGC891 data -- $144$ hr. \\item[$^{\\mathrm{\\#}}$] Re-analysed subset of the \\citet{swaters1999} dataset. \\item[$^{\\mathrm{\\star}}$] Includes re-analysed NGC891 data observed by \\citet{ssvdh1997}. \\end{list} \\end{footnotesize} \\end{table*} From Table~\\ref{tab:ch4-vel-disp-meas-lowi}, it is seen that only 3 of the low-inclination velocity dispersion studies were undertaken with good spectral resolution (channel FWHM$_v\\lesim 5$ \\kms). Of these three high spectral resolution studies, only the recent measurements of NGC1058 by \\citet{pr2006} had sufficient spatial resolution to sample the \\HI\\ radius by at least 10 synthesised beam FWHMs ($\\theta _{maj}$). The main difference between the observations of \\citet{pr2006} and the other low inclination \\HI\\ velocity dispersion measurements are that the observations by \\citet{pr2006} were deeper ($23$ hr total integration time on 3 VLA array configurations) and with good spectral resolution. The two low \\HI\\ inclination velocity dispersion studies that found a constant velocity dispersion with radius were only marginally spectrally resolved (channel FWHM$_v$= 8.2 \\kms), and likewise also marginally spatially resolved ($5 < R_{HI}/{\\rm FWHM}_{\\theta} < 10$). Of the \\HI\\ velocity dispersion measurements of high inclination galaxies (see Table~\\ref{tab:ch4-vel-disp-meas-highi}), only the study of NGC4244 \\citep{olling1996a} has the spatial resolution ($\\sim$100 pc) to investigate the radial velocity dispersion on scales relevant to common galactic structure (star formation, spiral density waves, bars). However, Olling's method (\\citeyear{olling1996a}) used only part of the \\HI\\ emission distribution, and may not be accurate (see Sect.~\\ref{sec:olling-RC-disp-meth}). The other high inclination velocity dispersion measurements obtained by \\citet{kvdkdb2004} used an \\HI\\ velocity dispersion fit that was held fixed with radius, consequently no information was obtained about the radial variability of the velocity dispersion (see the discussion of the method in Sect.~\\ref{sec:fixed-disp-iter-envtrac-RC}). As a result, very little is known about the velocity dispersion of the \\HI\\ in the ISM. The dependence of the \\HI\\ velocity dispersion on galaxy mass, surface brightness, star formation rate, and density variations such as spiral arms or bars is unknown. Proper modelling of the full \\HI\\ emission cube gives the radial distribution of surface density, rotation and velocity dispersion. This approach could be expanded to also model the \\HI\\ transparency of the \\HI\\ emission distribution of a galaxy, whether edge-on or inclined. In the past, low telescope resolution and deprojection issues made such studies nonviable. In this paper, we show that it is possible to accurately perform the deprojection required to measure the velocity dispersion, rotation and surface density as functions of radius in edge-on galaxies. The method is also applicable to galaxies of moderate inclination. The edge-on orientation of galaxies in our study is only required to minimise the uncertainty of the measured vertical gas flaring, which is required to determine the dark halo flattening. If fitting galaxies less inclined than edge-on, the inclination must be accurately fitted at all radii, and the galaxy must not be too close to face-on or else the line-of-sight velocity will be not be sufficiently resolved to model the galaxy rotation. Application of our new \\HI\\ emission modelling method to nearby non-edge-on spiral galaxies, would allow investigation of the velocity dispersion, transparency, temperature and density structure of neutral hydrogen in the ISM of a large sample of galaxies. This could explain how this small-scale ISM structure varies with the mass and star formation rate of spiral galaxies. Many methods have been devised to measure the rotation curves of edge-on galaxies. In Table~\\ref{tab:ch4-RC-meas}, we display the previous rotation curve measurements of edge-on galaxies, showing the method used and the resolution of the observations. In the next section, we will briefly review these methods, and the few previous methods to measure the gas velocity dispersion and surface density, before presenting our new method in Sect.~\\ref{sec:iter-raddecomp-all3}. In the paper III of this series we will present the results of applying this new method to the \\HI\\ observations of the edge-on galaxy sample presented in paper I. In that paper, we will paper also derive the radial flaring of the \\HI\\ layer of these galaxies. A determination of the flattening of the dark halo of one galaxy in our sample, UGC7321, will be presented in paper IV. ", "conclusions": "Using our new iterative radial decomposition method, we show that it is possible to accurately measure all three functions that describe the planar kinematics and distribution of a galaxy: the radial \\HI\\ surface density, the rotation curve and the \\HI\\ velocity dispersion. This method is a significant improvement on previous methods to model edge-on galaxies which approximated one or more of these kinematic functions. Measurements of simulated galaxies show that a peak signal-to-noise of $\\gesim 30$ is needed to accurately measure all three kinematic functions on a galaxy with a intrinsically steep (exponential over all R) \\HI\\ surface density. However, reliable measurements would probably also be obtained with observations at a peak signal-to-noise ($S/N \\gesim 20$) Furthermore, the method should also be successful on galaxies with an intrinsically shallower surface density --- as is common in most disk galaxies due to depletion of \\HI\\ in the stellar disk. In the next paper in this series, we present the derived \\HI\\ kinematics and radial surface density of our galaxies as measured with this method. Paper III also contains the vertical \\HI\\ flaring results, which is measured as a function of galactocentric radius using the deprojected \\HI\\ distribution and kinematics." }, "1003/1003.6031_arXiv.txt": { "abstract": "{% Global magnetohydrodynamic simulations show the growth of Kelvin-Helmholtz instabilities at the contact surface of two merging neutron stars. That region has been identified as the site of efficient amplification of magnetic fields. However, these global simulations, due to numerical limitations, were unable to determine the saturation level of the field strength, and thus the possible back-reaction of the magnetic field onto the flow. } {% We investigate the amplification of initially weak magnetic fields in Kelvin-Helmholtz unstable shear flows, and the back-reaction of the field onto the flow. } {% We use a high-resolution finite-volume ideal MHD code to perform 2D and 3D local simulations of hydromagnetic shear flows, both for idealized systems and simplified models of merger flows. } {% In 2D, the magnetic field is amplified on time scales of less than $0.01\\,$ms until it reaches locally equipartition with the kinetic energy. Subsequently, it saturates due to resistive instabilities that disrupt the Kelvin-Helmholtz unstable vortex and decelerate the shear flow on a secular time scale. We determine scaling laws of the field amplification with the initial field strength and the grid resolution. In 3D, the hydromagnetic mechanism seen in 2D may be dominated by purely hydrodynamic instabilities leading to less filed amplification. We find maximum magnetic fields $\\sim 10^{16}\\,$G locally, and r.m.s.\\ maxima within the box $\\sim 10^{15}\\,$G. However, due to the fast decay of the shear flow such strong fields exist only for a short period ($< 0.1\\,$ms). In the saturated state of most models, the magnetic field is mainly oriented parallel to the shear flow for rather strong initial fields, while weaker initial fields tend to lead to a more balanced distribution of the field energy among the components. In all models the flow shows small-scale features. The magnetic field is at most in energetic equipartition with the decaying shear flow. } {% The magnetic field may be amplified efficiently to very high field strengths, the maximum field energy reaching values of the order of the kinetic energy associated with the velocity components transverse to the interface between the two neutron stars. However, the dynamic impact of the field onto the flow is limited to the shear layer, and it may not be adequate to produce outflows, because the time during which the magnetic field stays close to its maximum value is short compared to the time scale for launching an outflow (i.e., a few milliseconds).} ", "introduction": "\\label{Sec:Intro} The merger of two neutron stars is considered the most promising scenario for the generation of short gamma-ray bursts (GRBs). After a phase of inspiral due to the loss of angular momentum and orbital energy by gravitational radiation, the merging neutron stars are distorted by their mutual tidal forces. Finally, they touch each other at a contact surface. Due to a combination of the orbital motion and the rotation of the neutron stars, the gas streams along that surface, the flow directions on either side of the surface being anti-parallel with respect to each other. As a consequence of this jump in the tangential velocity, the contact surface is Kelvin-Helmholtz (KH) unstable. Growing within a few milliseconds, the KH instability leads to the formation of typical KH vortices between the neutron stars. These vortices can modify the merger dynamics via the dissipation of kinetic into thermal energy. The generation of KH vortices is observed in actual merger numerical simulations \\citep[e.g.,][]{Oechslin_etal__2007__AA__NS-NS-merger--EOS-1}. The exponential amplification of seed perturbations can lead to very strong magnetic fields as shown by \\cite{Price_Rosswog__2006__Sci__NS-NS-merger-B-amplif}, and \\cite{Rosswog__2007__RevMex__NS-merger}. These fields, in turn, can modify the dynamics of the instability described above, either already during its linear growth phase or, for weak fields, in the saturated state. Exerting stresses and performing work on the fluid, the magnetic field does lose part of its energy. Thus, the maximum attainable field strength is limited by the non-linear dynamics. In their merger simulations, \\cite{Price_Rosswog__2006__Sci__NS-NS-merger-B-amplif}, and \\cite{Rosswog__2007__RevMex__NS-merger} observed fields exceeding by far $10^{15}\\,$G. Their numerical resolution, however, did not allow them to follow the detailed evolution of the KH instability in the non-linear phase. Thus, they could not draw any definite conclusions on the maximum strength of the field nor its back-reaction onto the fluid. They observed that the maximum field strength is a function of the numerical resolution: the better the resolution, the stronger becomes the field. Performing numerical convergence tests, these authors did not find an upper bound for the field strength attainable in the magnetized KH instability. Thus, \\citet{Price_Rosswog__2006__Sci__NS-NS-merger-B-amplif} discussed, based on energetic arguments, but not supported by simulation results, two different saturation levels: the field growth saturates when the magnetic energy density equals either the kinetic (\\emph{kinetic equipartition}) or the internal energy of the gas (\\emph{thermal equipartition}), corresponding to fields of the order of $10^{16}\\,$G and $\\sim 10^{18}\\,$G, respectively. From their simulations they were not able to identify the saturation mechanism applying to the KH instability in neutron-star mergers. Thus, we address this question here again using highly resolved simulations and independent numerical methods. Most simulations of neutron-star mergers, including the ones by \\cite{Price_Rosswog__2006__Sci__NS-NS-merger-B-amplif}, and \\cite{Rosswog__2007__RevMex__NS-merger}, are performed using smoothed-particle hydrodynamics (SPH) \\citep{Monaghan__1992__ARAA__SPH}. This free Lagrangian method is highly adaptive in space, and allows on to follow large density contrasts without ``wasting'' computational resources in areas of very low density. This property of SPH makes it highly advantageous for the problem of mergers. On the other hand, its relatively high numerical viscosity renders SPH inferior compared to Eulerian grid-based schemes for the treatment of (magneto-)hydrodynamic instabilities and turbulence \\citep{Agertz_etal__2007__MNRAS__SPH-grid-comparison}. Moreover, the spatial resolution of most merger simulations is rather low, i.e., the reliability of their results concerning the details of the KH instability is limited. A grid-based code such as ours is well suited for a study of flow instabilities and turbulence. Using it to simulate the entire merger event, however, is cumbersome due to the large computational costs required to cover the entire system with an appropriate computational grid. In spite of this fact, \\cite{Giacomazzo_Rezzolla_Baiotti__2009__PRL__NS_mergers_MHD} (see also \\cite{Liu_etal__2008__prd__GRMHD_NS_mergers, Anderson_etal__2008__PRL__NS_mergers_MHD_GW}) have performed full general-relativistic MHD simulations using vertex-centered mesh refinement to assess the influence of magnetic fields on the merger dynamics and the resulting gravitational waveform. But, as we shall show below, even their (presently world-best) grid resolution ($h \\sim 350\\,$m) is still too crude to properly capture the disruptive dynamics after the KH amplification of the field. For comparison, we note here that our merger models employ a grid resolution of $h \\sim 0.1\\,$m in 2D (\\secref{sSek:mrgr-2d}) and $h \\sim 0.8\\,$m in 3D (\\secref{sSek:mrgr-3d}), respectively. We performed a set of numerical simulations of the KH instability to understand the dynamics of magnetized shear flows and to draw conclusions on the evolution of merging neutron stars. The main issues we address in our study are motivated by two different, albeit related, intentions: \\begin{itemize} \\item We strive for a better understanding of the magneto-hydrodynamic (MHD) KH instability. This includes the influence of numerical parameters such as the grid resolution on the dynamics, and generic properties of the saturation of the instability. We address these questions by a series of \\emph{dimensionless} models that use scale-free parameters as most previous studies focusing on the generic properties of the KH instability instead of a particular astrophysical application. \\item We further want to verify the results of \\cite{Price_Rosswog__2006__Sci__NS-NS-merger-B-amplif} and reassess their estimates of the saturation field strength. Hence, we consider the growth time of the instability that has to compete with the dynamical time scale of the merger event (a few milliseconds), the saturation mechanism, the saturation field strength, and generic dynamical features of supersonic shear flows. Our results should also allow us to reassess the findings of global simulations extending the ones performed by \\cite{Price_Rosswog__2006__Sci__NS-NS-merger-B-amplif}, e.g., the simulations by \\cite{Anderson_etal__2008__PRL__NS_mergers_MHD_GW} and \\cite{Liu_etal__2008__prd__GRMHD_NS_mergers}. \\end{itemize} To this end we utilize a newly developed multidimensional MHD code \\citep{Obergaulinger_etal__2009__AA__Semi-global_MRI_CCSN} that employs various explicit finite-volume algorithms, and that is particularly well suited for simulating instabilities and turbulent systems. As the code is based on Eulerian high-resolution methods instead of SPH as in \\cite{Price_Rosswog__2006__Sci__NS-NS-merger-B-amplif}, our results are complementary to theirs, serving as an independent check. Since we are unable to simulate the entire merger event using fine resolution, we focus on the evolution of a small, representative volume around the contact surface. This \\emph{local} simulation allows us to concentrate on the dynamics of the magnetohydrodynamic KH instability. However, as our simulations lack the feedback from the dynamics occurring on scales larger than the simulated volume, its influence has to be mimicked by suitably chosen boundary conditions. We neglect neutrino radiation, and the gas obeys either an ideal-gas or a hybrid (barotropic and ideal-gas) equation of state (EOS), the latter serving as a rough model for nuclear matter. This paper is organized as follows. We describe the physics of the magnetohydrodynamic KH instability in \\secref{Sek:KH-Phys}, and our numerical code in \\secref{Sek:PhysNum}. We discuss the simulations addressing generic properties of the KH instability in two and three spatial dimensions in \\secref{Sek:KH2} and \\secref{Sek:3d-mod}, respectively. The results applying to neutron-star mergers are given in \\secref{Sek:mrgr}. Finally, we present a summary and conclusions of our work in \\secref{Sek:Summary}. ", "conclusions": "\\label{Sek:Summary} Global simulations indicate that the contact layer between two merging neutron stars is a site of very efficient field amplification. The layer is prone to the Kelvin-Helmholtz instability, and thus, exponential growth of any weak seed field is possible, as observed by \\cite{Price_Rosswog__2006__Sci__NS-NS-merger-B-amplif} (see also \\cite{Giacomazzo_Rezzolla_Baiotti__2009__PRL__NS_mergers_MHD, Anderson_etal__2008__PRL__NS_mergers_MHD_GW, Liu_etal__2008__prd__GRMHD_NS_mergers}). The limitations of their simulations, mainly concerning grid resolution, did not allow these authors to determine the saturation level of the instability firmly and accurately. Thus, the implications of magnetic fields for the merger dynamics remains unclear. On the basis of energetic arguments, the instability might lead to a field in equipartition with the kinetic or the internal energy of the shear flow, corresponding to field strengths of the order of $10^{16}\\,$G or $10^{18}\\,$G, respectively. We reassessed these arguments by means of local high-resolution simulations of magnetized shear layers in two and three spatial dimensions. To this end we performed more than 220 simulations focusing on properties of the hydromagnetic KH instability in general as well as on the contact surfaces of merging neutron stars. We refer to these two classes of simulations as \\emph{dimensionless} and \\emph{merger-motivated} models, respectively. We employed a recently developed multi-dimensional Eulerian finite-volume ideal MHD code based on high-order spatial reconstruction techniques and Riemann solvers of the MUSTA-type \\citep{Obergaulinger__2008__PhD__RMHD, Obergaulinger_etal__2009__AA__Semi-global_MRI_CCSN}. We set up a KH-unstable shear flow in Cartesian coordinates in a quadratic (2D) and cubic (3D) computational domain imposing periodic boundary conditions in the direction of the shear flow and reflecting or open ones in the transverse directions. Focussing on the effects of a magnetic field on the instability, we used a simplified equation of state (ideal gas EOS and a hybrid barotropic/ideal gas EOS for the dimensionless and the merger-motivated models, respectively) and neglected additional physics, e.g., such as neutrino transport. Under these simplifications, the shear flows are characterized by two parameters, the initial Mach number $\\Mach$ and the initial Alfv\\'en number $\\Alfv$, measuring the magnitude of the jump in shear velocity in units of the sound speed and the Alfv\\'en velocity, respectively. Analytic considerations and previous simulations of non-magnetized shear flows show that the growth rate of the KH instability as well as its saturation level (i.e., the kinetic energy of the circular KH vortex formed by the instability) increase with increasing $\\Mach$ for subsonic shear flows . A magnetic field is known to reduce the growth rate and potentially, i.e., for $\\Alfv < 2$, even to suppress the instability \\citep{Chandrasekhar__1961__Buch__HD-HM-stab, Miura_Pritchett__1982__JGR__MHD-KH-stability, Keppens_etal__1999__PP__MHD-KH}. However, less is known about the saturation level and the dynamic back-reaction, in particular for weak initial fields. \\cite{Frank_etal__1996__ApJ__MHD-KH-2d-1,Jones_etal__1997__ApJ__MHD-KH-2d-2, Jeong_etal__2000__Apj__MHD-KH-2d-3, Ryu_etal__2000__ApJ__MHD-KHI-3d} studied the evolution of the hydromagnetic KH instability in two and three dimensions. In 2D, models undergo a transition from \\emph{non-linear stabilization} of the KH vortex to its violent \\emph{disruption} or more gradual \\emph{dissipation} when the initial field strength is reduced, while in 3D purely hydrodynamic \\emph{elliptic} instabilities of the vortex tube may dominate over MHD effects. The study of these non-linear effects is hampered by high requirements on the grid resolution that is necessary to follow the development of increasingly thin magnetic flux sheets and tubes. This limits the range of Alfv\\'en numbers for which numerical convergence can be achieved to rather modest values. It also reduces the predictive power for merger systems, where rather weak initial fields are expected. This limitation can be overcome only when using large grids in combination with a highly accurate code. We evolved subsonic, transsonic, and supersonic shear flows with $\\Mach \\in [0.5; 1; 4]$, while using the maximum Alfv\\'en numbers for which convergence is achievable. The resulting broad range of Alfv\\'en numbers covered by our simulations allows us to establish scaling laws governing the field amplification as a function of the initial field strength. The main results of our simulations are: \\begin{enumerate} \\item In 2D, we confirm the results of analytic work (in the linear regime) and previous simulations concerning the growth rate and the saturation of the transverse kinetic energy densities for strong initial fields due to \\cite{Chandrasekhar__1961__Buch__HD-HM-stab,Miura_Pritchett__1982__JGR__MHD-KH-stability,Keppens_etal__1999__PP__MHD-KH}. This agreement supports the viability of our numerical approach for the problem at hand. \\item For subsonic shear flows ($\\Mach = 0.5, 1$) we explored a wide range of initial field strengths covering Alfv\\'en numbers up to $\\Alfv = 5000$ in 2D. \\begin{enumerate} \\item For intermediate and weak fields, we distinguish two phases: the KH growth phase during which the field grows at the KH growth rate, and after formation of a KH vortex, a phase of kinematic field amplification by the overturning vortex. The growth rate during the latter phase depends on the velocity of the vortex. The field is highly intermittent and concentrated in flux sheets, which are stretched by the flow leading to an exponential growth of the field strength while the sheet width decreases. \\item The termination of the kinematic amplification phase occurs either numerically, when the flux sheets get too thin to be resolved on a given computational grid, or dynamically by back-reaction of the field onto the flow. The most important mode of back-reaction is the growth of secondary resistive instabilities feeding off the magnetic energy of the flux sheets. These instabilities terminate the kinematic field growth and initiate the non-linear saturation phase during which the KH vortex is destroyed by the ensuing MHD turbulence and the shear flow is gradually decelerated. This scenario is equivalent to that of the \\emph{disruption} models of \\cite{Frank_etal__1996__ApJ__MHD-KH-2d-1}. \\item We quantified the amount of field amplification during the kinematic amplification phase by computing the ratio of the volume-averaged Maxwell stress component $M_{xy}$ at the beginning and at the end of that phase. The amplification factor scales with the initial \\Alfven\\,number as $\\Alfv^{3/4}$, corresponding to a scaling of the maximum Maxwell stress with the initial field strength as $b_0^{5/4}$. If the simulation is under-resolved, the amplification factor is reduced by a factor $\\propto m^{7/8}$ ($m$ being the number of zones per dimension). The maximum local field strength corresponds to a \\emph{local} equipartition between the magnetic energy density of a flux sheet and the kinetic energy density of the shear flow; it depends only weakly on the initial field. \\item The secondary resistive instabilities observed in our simulations are triggered by numerical resistivity instead of a physical one. The numerical resistivity, which is a function of the grid resolution $\\Delta$, is important only for small thin structures having a spatial size of the order of $\\Delta$ or less. In our simulations, it causes current sheets to become unstable when their width approaches the grid spacing $\\Delta$. Although only simulations with arbitrarily high resolution can sustain arbitrarily thin and intense current sheets, we observe nevertheless convergence: the field amplification becomes independent of the grid resolution, if $\\Delta$ is smaller than some threshold which depends on the initial field strength. The reason for this independence is the fact that the most unstable current sheets do not consist of individual flux sheets but of pairs or triples of coalescing flux sheets. Thus, decreasing the distance between flux sheets does not lead to a stronger field (which would be the case, if a single flux sheet is compressed in transverse direction). \\item The disruption of the vortex and the efficient dissipation set these \\emph{resisto-dynamic} models apart from the class of \\emph{dissipation} models with even weaker initial fields where the KH vortex remains intact, and only very slow dissipation is provided by turbulence. In the simulations of \\cite{Frank_etal__1996__ApJ__MHD-KH-2d-1}, secondary instabilities do not modify the flow field qualitatively. Our simulations indicate that this is, partially at least, a resolution effect. If a simulation is under-resolved and the field growth is not limited by dynamic back-reaction but by the resolvable width of flux sheets, no disruption will occur, and the deceleration time of the shear flow is very long. Converged simulations show, on the other hand, the disruption of the KH vortex by secondary magnetic instabilities when the magnetic field strength approaches a local maximum close to equipartition with the kinetic energy density of the shear flow. This happens in all converged models, but for weak initial fields, the deceleration time can be very long. \\item Models with initially anti-parallel and parallel magnetic fields, but otherwise identical, give qualitatively similar results, the above discussed effects being somewhat less pronounced in case of the former field configuration. \\end{enumerate} \\item The contact layer of merging neutron stars resembles supersonic shear flows. In principle, these are stable. We find, however, that an exponentially growing instability may occur when closed boundary conditions are imposed in the direction transverse to the shear flow. The instability is mediated by shock waves traveling through the computationl domain. The corresponding growth rates are much smaller than for subsonic shear flows. The effects of a magnetic field on a supersonic shear flow are qualitatively similar to those on subsonic shear flows. \\item In 3D the disruption of the KH vortex tube can be induced by a purely hydrodynamic secondary so-called \\emph{elliptic} instability as discussed, e.g., by \\cite{Ryu_etal__2000__ApJ__MHD-KHI-3d}. It leads to a very rapid growth of the kinetic energy densities corresponding to all components of the flow velocity once the KH vortex tube forms, and decelerates the shear flow more efficiently than the MHD mechanisms outlined above. Which of the two possible disruption mechanisms, elliptic or hydromagnetic, operates depends on the initial field strength $b_0$ and the value of volume-averaged kinetic energy density $e^z_\\mathrm{kin}$. The magnetic mechanism will dominate only if $e_\\mathrm{mag} > e^z_\\mathrm{kin}$, i.e., as long as the magnetic energy density exceeds the transverse kinetic energy in $z$-direction. Due to the very fast growth of the elliptic instability, this may be the case only for a short time, if at all. A rather strong initial field and small perturbations in $z$-direction are required for a hydromagnetic disruption. \\item 2D and 3D simulations of shear flows with merger-motivated initial conditions performed in a cubic computational domain of constant density and pressure having an edge size of $200\\,$m show the same overall dynamics as corresponding dimensionless models. The initial Mach number of the shear flow was chosen to be $\\Mach = 1$ and $\\Mach = 4$ corresponding to a density of $10^{13}\\,$g cm$^{-3}$, and shear velocities of $1.83 \\times 10^{9}\\,$cm/s, and $7.2 \\times 10^{9}\\,$cm/s, respectively. The initial magnetic field strength was varied between $5\\times 10^{13}\\,$G and $4 \\times 10^{14}\\,$G. \\begin{enumerate} \\item The instability grows rapidly: saturation occurs within $\\lesssim 0.1\\,$msec, and the disruption and deceleration times are much less than $1\\,$msec. \\item The dynamics is the same as that of the dimensionless models. Field amplification leads to a maximum field strength $\\lesssim 10^{16}\\,$G, and a r.m.s.\\ value of $\\lesssim 1.6 \\times 10^{15}\\,$G. These values are the same for 3D models suffering hydrodynamic and hydromagnetic disruption. \\end{enumerate} \\end{enumerate} From our results, we may draw a few conclusions concerning the growth and the influence of magnetic fields in neutron-star mergers. The foremost implication is that the maximum field strength, independent whether it refers to a single point or a spatial average, is not amplified to equipartition with the thermal energy density. We can, hence, exclude saturation fields of the order of $10^{18}\\,$G in the contact layers of neutron star mergers. Instead, local equipartition with the kinetic energy density is reached with corresponding maximum fields $\\sim 10^{16}\\,$G, as speculated by \\cite{Price_Rosswog__2006__Sci__NS-NS-merger-B-amplif}. Due to the high degree of intermittency in the case of weak initial fields, the (r.m.s.) average of the field strength is smaller, i.e, its direct dynamic impact (e.g., disruption of the KH vortex tube or deceleration of the shear flow) on the flow is probably rather limited. This is even more the case if the geometry of the system and the perturbations resulting from the merger dynamics enhance the importance of purely hydrodynamic instabilities. More indirect effects can, however, not be excluded, e.g., whether magnetic flux tubes created at the shear layer are transported rapidly far away by large-scale flows. The short period of time during which the magnetic field stays close to its maximum value and its fast decay impose severe constraints on the impact that the amplified fields may have on any hydromagnetic or electromagnetic jet-launching mechanism in a merger of two neutron stars. We note that magnetically driven relativistic outflows may need much longer time scales ($\\sim$ a few msec) to tap the rotational energy of either the black hole or the accretion disk resulting after the merger. Though these results limit the prospect for magnetic effects to play a dynamic role in neutron star mergers, their proper inclusion in current and forthcoming simulations may be necessary, because magnetic fields influence the dissipation rates in the shear layer, i.e., their neglect may lead to an underestimation of the temperature in the shear layer, and hence in the accretion disk. Given the resolution requirements imposed by weak initial fields, a more sophisticated treatment of the problem probably also has to abandon the assumption of ideal MHD and to consider the formulation of a turbulence model for unresolved magnetic field structures." }, "1003/1003.4202_arXiv.txt": { "abstract": "The paper investigates the overall and detailed features of cosmic ray (CR) spectra in the knee region using the scenario of nuclei-photon interactions around the acceleration sources. Young supernova remnants can be the physical realities of such kind of CR acceleration sites. The results show that the model can well explain the following problems simultaneously with one set of source parameters: the knee of CR spectra and the sharpness of the knee, the detailed irregular structures of CR spectra, the so-called ``component B'' of Galactic CRs, and the electron/positron excesses reported by recent observations. The coherent explanation serves as evidence that at least a portion of CRs might be accelerated at the sources similar to young supernova remnants, and one set of source parameters indicates that this portion mainly comes from standard sources or from a single source. ", "introduction": "Since the discovery by \\cite{KK1958}, the ``knee'' of cosmic ray spectra has been one fundamental problem of CR physics for half a century. Many theoretical works try to explain this interesting and important phenomenon. The most popular explanation attributes the knee to the inefficient acceleration of the Galactic CRs by the accelerators above PeV energies \\citep[e.g.,][]{2002PhRvD..66h3004K,2003A&A...409..799S,2001JPhG...27.1005E}. Alternative possibilities include the leakage of CRs when propagating in the Galaxy \\citep{1993A&A...268..726P,2004IJMPA..19.1133R, 2001NuPhS..97..267L}, interactions between CRs and the background light \\citep{1993APh.....1..229K,2002APh....17...23C} or neutrinos \\citep{2003APh....19..379W} before arriving at the earth, or exotic interaction of CRs in the atmosphere where undetectable particles are produced and missing the detection \\citep{2000PAN....63.1799N, 2001ICRC....5.1760K}. From the experimental aspects, the measurements of the CR spectra around the knee region become increacingly precise, which can even reveal some fine structures of the knee. After a long term operation, the Tibet Air Shower array reported a very good measurement of the knee spectra with unprecedented high statistics and low systematics \\citep{2008ApJ...678.1165A}. Especially interesting signature of the Tibet result is that the CR spectra show a very sharp break of the spectrum index around 4 PeV. At almost the same time several experiments have reported their new measurements with the similar behavior, such as KASKADE\\citep{2009APh....31...86A}, ARAGATS-GAMMA\\citep{2008JPhG...35k5201G}, Yakutsk\\citep{2009NJPh...11f5008I}, and MAKET-ANI\\citep{2007APh....28...58C}. Such a sharp knee challenges the traditional interpretations of the knee \\citep{Shibata2009,2009arXiv0906.3949E,2009arXiv0911.3034H}. It is shown that if adopting an exponential-like cutoff of each component of CR species with low He flux, it will be very difficult to reproduce the sharp knee data \\citep{Shibata2009}. It is suggested by \\cite{2009arXiv0911.3034H} that a double power-law may well fit the observational data with high He flux, which indicates that He may be the main component around the knee. Furthermore Hillas suggested that there should be another Galactic component, ``component B'', to explain the CR spectra above $10$ PeV \\citep{2005JPhG...31R..95H}. Besides the sharp transition of the knee, \\cite{2009arXiv0906.3949E} carefully analyzed the CR spectra of individual experiment. By renormalizing the energy with respect to the break point(measured knee energy)of individual experiment, the problem related to the uncertainty of the absolute energy scale can be avoided, so the deviations of the observed spectra from the fitted spectra can be combined for all experiments. The result clearly shows the peculiarities at the positions expected for CNO group and Fe group if the knee is corresponding to the position of He. Interestingly the energies of these fine bumps are proportional to the mass number of the several major nuclei species: proton, He, CNO and Fe. The sharp knee and the irregularities of CR spectra are regarded as evidence for the single source origin of CRs \\citep{2009arXiv0906.3949E}. Another important development in CR physics is the new discovery of electron/positron excesses by several experiments \\citep{2009Natur.458..607A,2008Natur.456..362C,2008PhRvL.101z1104A, 2009A&A...508..561A,2009PhRvL.102r1101A}. To explain the positron fraction and electron spectrum excesses simultaneously one may need to introduce some exotic sources of e$^+$e$^-$ pairs \\citep{2009PhRvD..79b1302S}. \\cite{2009ApJ...700L.170H} (hereafter Paper I) proposed a model resorting to e$^+$e$^-$ pair through interactions of CR nuclei and ambient photons around the acceleration sources, which can explain the knee of the CR spectra and the electron/positron excesses at the same time. Based on that model we further study the detailed structures of the CR spectra in this work, intending to reproduce the sharp knee and fine structures mentioned above. In our interaction model the threshold energy of different chemical compositions is $A$-dependent, which will result in an $A$-dependent knee of each composition. This feature is consistent with the property of the fine structures found in experimental data \\citep{2009arXiv0906.3949E}. In addition, the interaction will cause a pile-up of the particles below the threshold. We expect this effect can contribute to the sharp knee and irregular bumps of the CR spectra. This paper is organized as follows. In Sec. 2 we will first go over the model describing CR-photon interactions. In Sec. 3 we present the calculated results and comparisons with observational data of the sharp knee and irregular structures of CR spectra. Finally Sec. 4 is the conclusion. ", "conclusions": "In summary we use the pair production interaction model between CR nuclei and ambient radiation field proposed in Paper I to explain the features of the CR spectra, including the sharp knee and fine structures. Results show that the spectra of CRs agree well with the observations. In our model, the He composition dominates around the knee at $\\sim 4$ PeV. The sharp knee observed by Tibet air shower array and confirmed by more and more experiments, can be well reproduced through the pile-up of He particles. In addition, this model can explain the fine structures of CR spectrum through the pile-up effects of CNO group and Fe nuclei. As an additional result, our model can provide a natural explanation of the ``component B'' problem of individual composition in tens of PeV energy range as suggested in \\cite{2005JPhG...31R..95H}. As in Paper I, the electron/positron excesses can also be explained. Furthermore we would like to discuss some implications of the model parameters. We note that only one single set of parameters is enough to explain the data. As discussed in Paper I and \\cite{2009arXiv0911.3034H}, it may indicate that the sources are ``standard'' which have similar parameters such as the temperature evolution of the radiation field, the relative abundances and spectral indices for individual elements, or the observed fine structures of CRs spectra are mainly due to one single source, either nearby or not so nearby but intensive (e.g. possibly the Galactic center). This work can be regarded as one progress approaching the origin of CRs." }, "1003/1003.3746_arXiv.txt": { "abstract": "{The QSO HE0450-2958 and the companion galaxy with which it is interacting, both ultra luminous in the infrared, have been the subject of much attention in recent years, as the quasar host galaxy remained undetected. This led to various interpretations on QSO and galaxy formation and co-evolution, such as black hole ejection, jet induced star formation, dust obscured galaxy, or normal host below the detection limit. } { We carried out deep observations in the near-IR in order to solve the puzzle concerning the existence of any host.} { The object was observed with the ESO VLT and HAWK-I in the near-IR J-band for 8 hours. The images have been processed with the MCS deconvolution method (Magain, Courbin \\& Sohy, 1998), permitting accurate subtraction of the QSO light from the observations.} {The compact emission region situated close to the QSO, called the blob, which previously showed only gas emission lines in the optical spectra, is now detected in our near-IR images. Its high brightness implies that stars likely contribute to the near-IR emission. The blob might thus be interpreted as an off-centre, bright and very compact host galaxy, involved in a violent collision with its companion. }% {} ", "introduction": "While there is convergence in the scientific community on the idea that the super massive black holes (SMBHs) lying at the centre of galaxies are growing in parallel with the galactic bulge they are hosted in, a few outstanding objects seem to lie outside this relation. This is probably the case of HE0450-2958. Located at z=0.286, 04h 52m 30.10s $-$29d 53m 35.3s, this system was first detected as an IRAS source (\\cite{degrijp}), then identified as a QSO and studied as an interacting active galaxy (\\cite{low}, \\cite{canalizo}). The interest in this system grew with the claim by Magain et al.\\ (2005) that no host galaxy was detected on optical HST images (ACS F606W), implying that it was either abnormally compact or more than six times fainter than the QSO hosts described in \\cite{floyd04}. The only significant extended emission arose from two objects in the nearby field: first, from a strongly disturbed companion galaxy 7\\,kpc distant in projection from HE0450-2958, and also from a cloud of highly ionized gas showing no trace of stellar light, centered at 0.9\\,kpc NW from the QSO. Several studies followed, proposing various scenarios for explaining this system. \\cite{haenelt} and \\cite{hoffman07} proposed that it was the first observation of a BH ejected during a galactic merger. \\cite{merrit} observed that the QSO spectrum was similar to those of Narrow Line Seyfert 1 galaxies, suggesting accretion above the Eddington limit and thus argued that the BH mass had been overestimated and that the host could then be just below the detection limit. Radio observations of the CO molecule by \\cite{papad} implied that the bulk of star formation was located in the companion galaxy and not around the QSO, and that both the QSO and galaxy have the properties of Ultra Luminous Infrared Galaxies (ULIRGs). \\cite{letg09} and \\cite{knud09} showed that the center of the companion galaxy is highly obscured by dust. \\cite{letg09} also presented evidence that a second AGN lies in the companion galaxy, hidden by the thick dust cloud. Lately \\cite{elbaz} proposed an alternative scenario of galactic building, by jet induced star formation. None of these studies has however been able to directly detect a host galaxy. This is the purpose of the present deep near-IR observations with the ESO VLT, aimed at detecting a co-centered host galaxy, if any, around the QSO. \\begin{figure} \\centering \\includegraphics[width=9cm]{illudec.eps} \\caption{HAWK-I observation of HE0450-2958 in the J-band. Top left: one of the 12 exposures; Top right: deconvolved image, where the point sources are modeled by 2D gaussians, here appearing as black dots; Bottom left: diffuse background only (companion galaxy + other extended emissions) as obtained by simultaneous deconvolution of all exposures (see text); Bottom right: average residuals (model minus observations, in units of the standard deviation; the intensity scale goes from $-3 \\sigma$ to $3 \\sigma$) . North is up and East to the left.} \\label{imgdec} \\end{figure} ", "conclusions": "The detection of a plausible host rules out the scenario of the ejection of a BH from the companion galaxy, that was previously introduced by \\cite{haenelt} or \\cite{hoffman07} and discussed in several following papers concerning HE0450-2958, such as \\cite{merrit} and \\cite{letg09}. If the BH had been ejected with some gas to feed it, it would not be accompanied by a galaxy with stars. The well-detected stellar emission from the blob sheds new light on recent attempts to explain this unusual system. Indeed, \\cite{elbaz} proposed that the activity of the quasar and associated radio jets have been creating the companion galaxy and parts of the host in construction, still partially disjoint from the QSO as seen in a primordial step of its evolution. The observation of stars in the direct vicinity of the QSO shows that there is already a host candidate, independently of the impact of the radio jet on the surrounding objects. Moreover, from the CO observations of \\cite{papad}, no star formation is found at the blob location, excluding the possibility that the QSO could have created the blob. These findings support the hypothesis already presented by \\cite{papad}, involving a collision between a gas rich galaxy (the companion), and a smaller one (the blob). Two galaxies have collided, each of them probably harbouring an active BH, the interaction enhancing the star formation in at least one of the galaxies and disrupting the pre-existing structures. The disrupted parts will probably merge in the future into a single galaxy. The radio jet would play a secondary role in the enhancement of the star formation in some parts of the system (e.g.\\ in the ``tail\" in between the main QSO and the companion galaxy, and maybe in the N-E part of the companion galaxy itself). Moreover, the QSO radiation would ionize the gas expelled all around the system (Letawe et al., 2008b) In fact, this scenario explains all the observations: disturbed morphology in both galaxies (including off-centre nuclei), presence of the second AGN, enhanced star formation, N-E stellar emission and extended emission line regions. The previous non detection of the stellar content in the blob was due to its faintness in the optical, its compactness and its proximity to the bright QSO, bringing the optical continuum just below the detection limits. The secure detection of stars near the QSO, even if not centred on the nucleus, makes the HE0450-2958 system consistent with a highly perturbed merger, probably not requiring more exotic scenarios. As such, it fits reasonably well with the hierarchical models of galactic building by successive mergers." }, "1003/1003.1575_arXiv.txt": { "abstract": "We highlight the role of the light elements (Li, Be, B) in the evolution of massive single and binary stars, which is largely restricted to a diagnostic value, and foremost so for the element boron. However, we show that the boron surface abundance in massive early type stars contains key information about their foregoing evolution which is not obtainable otherwise. In particular, it allows to constrain internal mixing processes and potential previous mass transfer event for binary stars (even if the companion has disappeared). It may also help solving the mystery of the slowly rotating nitrogen-rich massive main sequence stars. ", "introduction": "A large effort has been undertaken in the last decades to measure and understand the surface chemical composition of massive main sequence stars. In particular, the detection of nitrogen enhancements in quite a number of such stars (e.g., Gies \\& Lambert 1992) has triggered the idea that internal mixing processes can bring material from the stellar core to the surface in rapid rotators (Meynet \\& Maeder 2000, Heger and Langer 2000). The picture has become more complicated by the recent analysis of a large sample of early B~type main sequence stars of Hunter et al. (2008), who showed that the nitrogen-rich stars found by Gies \\& Lambert (1992), who restricted their analysis to objects with low projected rotational velocities, are likely part of a population of {\\em intrinsically} slowly rotation main sequence stars. This view is supported by the work of Morel et al. (2006, 2008), who indeed identifies such a population in our Galaxy (see also Morel, 2009). The origin of the nitrogen enrichment in these stars is not understood, but as they are slow rotators it appears difficult to reconcile them with the idea of rotational mixing. On the other hand, Hunter et al. (2008) also identified a nitrogen-rich population of rapidly rotating early B stars, which appears to be well in line with the predictions of theoretical models including rotational mixing (cf., Maeder et al. 2008). The caveat here is that evolutionary models of massive close binaries --- whether or not they include rotationally induced chemical mixing (Langer et al. 1998) --- appear to predict essentially the same trend of nitrogen enrichment with rotational velocity as the single star models (Meynet \\& Maeder 2000, Heger and Langer 2000). The light elements lithium, beryllium and boron may play a key role to resolve this issue. They are so rare in the interstellar medium that they can not influence the course of stellar evolution, and thus are often neglected in massive star models. Furthermore, they are fragile nuclei which are generally not synthesised in stars, but rather destroyed, at least certainly on the main sequence. In the cool low mass stars, lithium is most interesting, as it can be observed rather easily, and it is indeed used extensively to constrain internal mixing processes as can be seen in many contributions to this book. In massive stars this role can be played by boron as will be outlined below. ", "conclusions": "It remains a major challenge to the theory of massive star evolution to explain the observed surface abundances of massive main sequence stars. While until recently, the incorporation of rotational mixing was thought to lead to a much better agreement, the discovery that a significant fraction of early B dwarfs are nitrogen-rich and intrinsically slowly rotating has cast some doubts on the previous ideas. We argue that boron observations of early type main sequence stars, as performed by Venn et al. (1996, 2002) and Morel et al. (2006, 2008), have the potential to move towards a solution. Clearly, binary evolution needs to be considered at the same time. Finally, we may be facing the situation that still not all mixing processes which can operate in massive main sequence stars have been described. Amongst possible candidates is mixing due to magnetic processes, and mixing induced by subsurface convection zones in hot massive stars." }, "1003/1003.1605_arXiv.txt": { "abstract": "We have calculated the chameleon pressure between two parallel plates in the presence of an intervening medium that affects the mass of the chameleon field. As intuitively expected, the gas in the gap weakens the chameleon interaction mechanism with a screening effect that increases with the plate separation and with the density of the intervening medium. This phenomenon might open up new directions in the search of chameleon particles with future long range Casimir force experiments. ", "introduction": " ", "conclusions": "" }, "1003/1003.1119_arXiv.txt": { "abstract": "We analyse the Fundamental Plane (FP) relation of $39,993$ early-type galaxies (ETGs) in the optical (griz) and $5,080$ ETGs in the Near-Infrared (YJHK) wavebands, forming an optical$+$NIR sample of $4,589$ galaxies. We focus on the analysis of the FP as a function of the environment where galaxies reside. We characterise the environment using the largest group catalogue, based on 3D data, generated from SDSS at low redshift ($z < 0.1$). We find that the intercept $``c''$ of the FP decreases smoothly from high to low density regions, implying that galaxies at low density have on average lower mass-to-light ratios than their high-density counterparts. The $``c''$ also decreases as a function of the mean characteristic mass of the parent galaxy group. However, this trend is weak and completely accounted for by the variation of $``c''$ with local density. The variation of the FP offset is the same in all wavebands, implying that ETGs at low density have younger luminosity-weighted ages than cluster galaxies, consistent with the expectations of semi-analytical models of galaxy formation. We measure an age variation of $\\sim 0.048$~dex ($\\sim 11\\%$) per decade of local galaxy density. This implies an age difference of about $32 \\%$ ($\\sim 3 \\, Gyr$) between galaxies in the regions of highest density and the field. We find the metallicity decreasing, at $\\sim 2$~$\\sigma$, from low to high density. We also find $2.5 \\, \\sigma$ evidence that the variation in age per decade of local density augments, up to a factor of two, for galaxies residing in massive relative to poor groups. The velocity dispersion slope of the FP, $``a''$, tends to decrease with local galaxy density, with galaxies in groups having smaller $``a''$ than those in the field, independent of the waveband used to measure the structural parameters. Environmental effects (such as tidal stripping) may elucidate this result, producing a steeper variation of dark-matter fraction and/or non-homology along the ETG's sequence at higher density. In the optical, the surface brightness slope, $``b''$, of the FP increases with local galaxy density, being larger for group relative to field galaxies. The difference vanishes in the NIR, as field galaxies show a small ($\\sim 2.5\\%$), but significant increase of $``b''$ from $g$ through $K$, while group galaxies (particularly those in rich clusters) do not. The trend of $``b''$ with the environment results from galaxies residing in more massive clusters, since for groups no variation of $``b''$ with local density is detected. A possible explanation for these findings is that the variation of stellar population properties with mass in ETGs is shallower for galaxies at high density, resulting from tidal stripping and quenching of star formation in galaxies falling into the group's potential well. We do not detect any dependence of the FP coefficients on the presence of substructures in parent galaxy groups. ", "introduction": "\\label{sec:intro} A robust prediction of hierarchical theories of galaxy formation is that the environment plays an important role in shaping the galaxian properties. From the observational viewpoint, this is successfully confirmed by the existence, for instance, of the well established morphology-density relation~\\citep{Dressler:80}, for which denser environments are preferably populated by galaxies with early-type morphology in contrast to the late-type dominated field regions. The role of environment in shaping the properties of ETGs -- the most massive galaxies in the local Universe -- is still an open question in our understanding of galaxy formation and evolution. Studies of the colour--magnitude (hereafter CM) relation in the low-redshift-Universe have found no significant difference in the average colour of ETGs among high- and low-density environments, implying either a tiny difference or a strong anti-correlated variation of stellar population properties, i.e. age and metallicity, with environment~\\citep{Bern:03b, Hogg:04, Haines:06}. Even small differences in the star formation history of galaxies residing in different environments would be magnified as one approaches the galaxy formation epoch. Several studies have found no remarkable difference in the properties of cluster and field galaxies at high redshift. For instance, \\citet{Koo:05a, Koo:05b} found that both field and cluster ETGs have very similar CM relations at redshift $z \\sim 0.8$. \\citet{Cool:06} also reported that the offset and slope of the CM relation of field and cluster galaxies are fully consistent up to redshift $z \\sim 0.4$. \\citet{Pannella:09} found that ETGs have similar characteristic ages, independent of the environment they belong to. An opposite picture comes out of spectroscopic studies of ETGs at redshift $z \\sim 0$. Many authors have reported younger luminosity-weighted ages for field relative to cluster galaxies, with the age difference amounting to $\\sim 1$--$2$~Gyr (e.g.~\\citealt{GLC:92, Trager:00, Kunt:02, TF:02, Thomas:05, BERN:06, Clemens:09}). Such difference seems to exist also when considering galaxies residing in high-density, low-velocity dispersion systems, such as the Hickson Compact Groups~\\citep{deLaRosa:07}. A even more puzzling issue is the difference in metal content of ETGs among different environments, as different studies have found that field ETGs are more metal-rich~\\citep{Thomas:05, deLaRosa:07, Kunt:02, Clemens:09}, as metal-rich as~\\citep{BERN:06, Annibali:07}, or even more metal-poor than their cluster counterparts~\\citep{GALL:06}. A main feature of the spectroscopic investigations is that the galaxy spectra, and hence the inferred stellar population properties, refer to the galaxy inner region, typically inside one effective radius. On the other hand, ETGs are know to possess internal colour gradients, with their central part being redder than the outskirts { (see~\\citealt{Wu:05, Roche:09}; and references therein)}. These gradients are mainly due to metallicity~\\citep{Pel:90}, with a small, but significant contribution from age~\\citep{LdC09, Clemens:09}. Moreover, as shown by~\\citet{LaB:05}, colour gradients seem to depend on the environment where galaxies reside, hence complicating the environmental comparison of spectroscopic properties. Due to its small intrinsic dispersion (see e.g.~\\citealt{Gargiulo:09, HB:09}; and references therein), the Fundamental Plane (FP) relation of ETGs~\\citep{Dressler87, George87}, i.e. the scaling law involving radius, velocity dispersion, and surface brightness, is a powerful tool to measure their mass-to-light ratio (see e.g.~\\citealt{vDF:96}). { This results from interpreting the FP as a consequence of mass-to-light ratio systematically varying with mass and the non-homology of ETGs (see e.g.~\\citealt{HjM95, CdC95, CL97, GrC97, BCC97, Tru04, BBT07, TNR:09})}. Several studies have used the FP to constrain the environmental variation, at a given mass, of galaxy luminosity. Even in this case, a contradictory picture emerges. ~\\citet{BERN:06} found the FP relation at low- and high-density to have consistent slopes but a significant offset, interpreting it as a difference of $\\sim 1$~Gyr in the formation epoch of field and cluster ETGs. On the contrary, \\citet{Donofrio:08}, analysing galaxies in massive nearby clusters, found also significant variations of the slopes with local density. High redshift studies of the FP have reported a consistent evolution of the mass-to-light ratio of ETGs among different environments, implying the same formation epoch of field and cluster galaxies~{ ~\\citep{JCF:06, vDvM:06}}. \\citet{vW:05} also found that the difference in mass-to-light ratios between field and cluster galaxies depends on galaxy mass, with low mass systems exhibiting a strong differential evolution (see also~\\citealt{PDdC01, dSA:06}). This work is the third paper of a series aimed to investigate the properties of bright ($M_r< -20$) ETGs as a function of the environment where they reside, in the low-redshift-Universe ($0.05$ = 800 km/s). Moreover, as discussed in Sec.~\\ref{sec:offset_fp}, the trend with global environment can be entirely explained by the dependence of FP coefficients on local density. D08 found a strong variation of FP coefficients with the local environment. Both $``a''$ and $``b''$ were found to increase, and $``c''$ was found to decrease with respect to the (projected) local galaxy density. The trends of $``b''$ and $``c''$ from D08 are qualitatively consistent with those we find here (see Secs.~\\ref{sec:offset_rband} and~\\ref{sec:slopes_rband}). However, after correcting for systematic effects (see red circles in Fig.~\\ref{fig:FP_DENS}), we detect a $2.5\\sigma$ indication that the $``a''$ {\\it decreases} from the low to the higher density regions, in contrast with D08. Considering ETGs in the most massive groups (see Fig.~\\ref{fig:FP_MASS_DENS_a}), we still do not find any strong variation with local environment, as in D08. A possible explanation for this discrepancy can be the fact that we have corrected the FP slopes of different samples of ETGs for biases due to the environmental dependence of the average mass-to-light ratio of galaxies and the distribution of ETGs in the space of effective parameters, while D08 did not account for either effects. We also notice that the variation of $``b''$ with local environment from D08 is much larger than that detected here. From Fig.~\\ref{fig:FP_DENS}, we see that the end-to-end variation of $``b''$ with respect to local density amounts to less than $0.02$, while D08 find a variation of $\\sim 0.04$ (see their fig.~15). This discrepancy can be explained by the different mass regime covered by our cluster sample and that of D08, and by the fact that we find the variation of $``b''$ with local environment to depend on the mass of parent galaxy groups. For groups as massive as $\\sim 5.8 \\times 10^{14} \\, M_\\odot$, we find an end-to-end variation of $``b''$ of $\\sim 0.04$ when none of the above corrections is applied (in agreement with D08). This variation reduces to $\\sim 0.02$ after the corrections are performed. \\subsection{Environmental dependence of the FP from g through K} \\label{sec:env_dep_FP_gtoK} For what concerns the waveband variation of FP coefficients, we find that the $``a''$ increases, by $\\sim 16\\%$, from $g$ through $K$, independent of the environment, while the variation of $``b''$ with passband is smaller than $\\sim 3\\%$. For field galaxies, (see Fig.~\\ref{fig:FP_slopes_grizYJHK_field_group}), there is a $3~\\sigma$ evidence that the $``b''$ increases, by $\\sim 2.5\\%$, from the optical to NIR, with the difference in $``b''$ between field and group galaxies vanishing in the NIR. For group galaxies, in particular those residing in most massive systems (Fig.~\\ref{fig:FP_slopes_grizYJHK_mass}), no variation of $``b''$ with waveband is detected. As discussed in paper II, the waveband dependence of FP slopes informs on the variation of stellar population properties with mass. For the entire SPIDER sample, the relative variation of $``a''$ and $``b''$ from $g$ through $K$ implies that ETGs, as a whole, have essentially coeval stellar populations, with metallicity being larger in more massive galaxies. As shown from eq.~7 of paper II, a smaller value of $``b''$ (as that we find here for field galaxies) implies a steeper variation of the mass-to-light ratio of ETGs along the galaxy (mass) sequence. As the $``b''$ increases from the optical to NIR for field galaxies, the corresponding mass-to-light ratio variation is related to the way stellar population properties change with mass. In other terms, the environmental variation of $``b''$ might imply that group galaxies exhibit a shallower relation between stellar population parameters (i.e. age and metallicity) and mass. Interestingly, this is consistent with the recent findings of~\\citet{Pasquali:10} (hereafter PGF10), who found that galaxies in denser environments have shallower age-- and metallicity--mass relations. This is also a natural expectation of hierarchical models of galaxy formation, as galaxies in denser environments are those accreted in the group environment at earlier times, having their star formation quenched earlier through environmental-driven effects (e.g. strangulation), implying a shallower relation between (luminosity-weighted) age and stellar mass. A flatter metallicity--mass relation in denser environment might also result from the effect of tidal stripping on the more bound stellar material in galaxies (e.g.~\\citealt{Kli09}) as discussed by PGF10.\\\\ On the other hand, the difference in $``a''$ between field and group galaxies is the same at both optical and NIR wavebands. Group galaxies have a larger FP tilt, as measured by the coefficient $``a''$, likely because of a waveband-independent origin. A possible explanation is a different variation of dark-matter fraction and/or non-homology along the mass sequence of ETGs between galaxies in different environments. For instance, group galaxies might have a larger tilt because of a steeper relation between dark-matter fraction and mass, due to (i) tidal interactions between galaxy dark-matter halos and parent group halos, i.e. lower, relative to higher, mass galaxies are stripped more of their dark-matter halo as they fall into the cluster potential well; (ii) central galaxies have higher dark-matter fractions, as they acquire part of the dark-matter halo of the in-falling satellites. The first scenario might result from the fact that the total time span, during which galaxies are stripped, is larger for low mass galaxies (see e.g. eq.~9 of~\\citealt{Cattaneo10}). However, tidal forces are stronger for more massive galaxies, as a short orbital decay time for $\\rm L > L^*$ galaxies is expected~\\citep{Barnes89}. The scenario (ii) might be more plausible, provided that stripped galaxies have mostly late-type morphology. This would not change significantly the peak value of dynamical mass of ETGs as a function of environment, consistent with Fig.~\\ref{fig:mdyn_r_dens}, and increase only the dark-matter fraction for higher mass systems, implying a smaller tilt of the FP relation. { Notice that a steeper relation between stellar to dynamical mass ratio and mass in ETGs for group (relative to field) galaxies has recently been suggested also from~\\citet{SB:09}, as a result of the wider range of luminosity-weighted ages characterizing cluster relative to field galaxies.} These points will be further analysed in a forthcoming contribution in this series, using estimates of dynamical and stellar mass for ETGs and analysing their dependence on environment. It is interesting to notice that the NIR FP, and its environmental dependence, can be connected to the stellar mass FP recently derived by~\\citet{HB:09}. The stellar mass FP is obtained by replacing luminosity (i.e. surface brightness) in the FP equation with stellar mass. Since the K-band light closely follows the stellar mass distribution in galaxies and the contribution of stellar populations to the tilt of the FP vanishes in the NIR (see paper II), the K-band FP should essentially coincide with the stellar mass FP. In fact, \\citet{HB:09} report a value of $a \\sim 1.54\\pm0.02$ for the stellar mass FP, while we measure $a \\sim 1.55\\pm0.02$ for the K-band FP. This reinforces the fact that the environmental difference of FP coefficients we detect in the NIR bands, among different environments (see above), implies an environmental difference in the relation between the ratio of dynamical-to-stellar mass and galaxy mass. Given the excellent agreement of the K-band and stellar mass FP's, one might actually use either relations to derive ETG's stellar masses, and test how realistic are the estimates from stellar population models. \\subsection{Stellar content of ETGs in different environments} \\label{sec:env_sp} The environmental variation of the FP intercept, `$``c''$, implies that the average mass-to-light ratio of ETGs increases from low to high density regions. { As seen in Eq.~\\ref{eq:age_met_dens}, this variation can be interpreted as a difference in the stellar content of ETGs, at fixed mass, provided that the ratio of dynamical to stellar mass, $M/M_{\\star}$, does not increase, on average, with local density. There are several arguments against such an increase. First, Fig.~\\ref{fig:mdyn_r_dens} shows that the peak value of the quantity $\\sigma_0^2 \\times R_e$ -- a proxy for dynamical mass -- does not change with the environment. So, the $M/M_{\\star}$ can increase with local density, only if stellar mass, at fixed $M$, decreases with $\\Sigma_N$. However, as galaxies are accreted into groups and clusters, they are expected to be stripped off of their dark-matter halos, resulting, eventually, into a larger $M_{\\star}$, for fixed $M$, at high density. These arguments make plausible our assumption of constant $M/M_{\\star}$, that allows us to use the offset of the FP as a tool to analyse the variation of stellar population properties as a function of environment. It is interesting to note that, as shown in fig.~8 of~\\citet{BERN:06}, at fixed luminosity, $L$, dynamical mass does not depend on environment, while the $g-r$ colour (which is a good proxy for $M_{\\star}/L$) does. This implies that, at fixed $L$, the $M/M_{\\star}$ should change with environment. We notice that this is not inconsistent with the above assumption of constant $M/M_{\\star}$, as we are assuming that the $M/M_{\\star}$ is constant at fixed $M$ (not $L$)}. We find a significant difference of the FP offset, in r band, between field and group galaxies, fully consistent with that measured from B06 (see previous section). Using spectral abundance indicators, B06 concluded that the difference in $``c''$ is due to a difference in age between field and cluster ETGs, with field galaxies being younger by $\\sim 1~Gyr$. Using the information encoded in the FP offset at different wavebands, we show that the variation of the FP intercept with local galaxy density implies a variation in logarithmic (luminosity-weighted) age per decade of local density of $\\delta \\log t/\\delta \\log \\Sigma_N =0.048 \\pm 0.006$. From Eq.~\\ref{eq:c_dens}, taking into account the difference of $``c''$ between field and group galaxies (see previous section), the value of $\\delta \\log t/\\delta \\log \\Sigma_N$ translates to a relative difference in age of $\\sim 18 \\pm 2\\%$ between the field and group environments. For a galaxy formation epoch of $10~Gyr$, this implies an age absolute difference of $\\delta t \\sim 1.8\\pm 0.2~Gyr$, larger than what found by B06. The finding that field ETGs are younger than those in high density environments is qualitatively consistent with some previous studies (see Sec.~\\ref{sec:intro}).~\\citet{Clemens:09} found that field ETGs are $\\sim 2~Gyr$ younger than their cluster counterparts, in agreement with the value of $\\delta t$ we measure here. The same amount of difference in age was also found by~\\citet{Thomas:05}. Similar results, but based on smaller samples of ETGs, were also obtained by (e.g.) \\citet{Kunt:02, TF:02, Sanch:06, deLaRosa:07}. { Our results are also consistent with those of~\\citet{Cooper:10}, who found that, at given colour and luminosity (and hence stellar mass), galaxies with older stellar populations favour regions of higher overdensity.} We notice that all these studies have been based on the analysis of spectral features (i.e. line indices) of ETGs. Hence, the inferred differences in age refer to the inner galaxy region, typically inside one effective radius. As already noticed in paper II, the information provided by the waveband dependence of the FP is more related to the global properties of the light distribution in galaxies (as measured by the structural parameters). In this regard, it is not affected by radial population gradients in ETGs and their possible dependence on environment~\\citep{LaB:05}. Recently, \\citet{Rogers:10} have performed a Principal Component Analysis of the SDSS spectra of ETGs residing in groups and clusters, spanning a wide range in mass (similar to that of our FoF catalogue). They find galaxies populating the lowest mass halos ($M_{group} \\sim 10^{12} M_\\odot$) to be younger, by $\\sim 1$~Gyr, than those in most massive clusters. Using the linear fit in Fig.~\\ref{fig:offset_r_mass}, we estimate a variation of the FP offset, over the entire range of parent group mass, of $\\sim -0.015$. From the measured values of $b_{_{\\Sigma_N}}$ and $\\delta \\log t / \\delta \\Sigma_N$ (see Sec.~\\ref{sec:offset_sp}), this difference translates to a difference in age of $\\sim 12 \\%$, i.e. $\\sim 1.2$ Gyr (for a formation epoch of $10$~Gyr), consistent with~\\citet{Rogers:10}. Recently, \\citet{Thomas:09} have analysed a large sample of morphologically selected ETGs from the SDSS, deriving stellar population properties, i.e. age, metallicity, and $\\alpha$-enhancement, for galaxies in low- and high-density environments. In contrast to the works listed above, they concluded that the luminosity-weighted age of the bulk of ETGs is independent of the environment. Low- and high-density environments actually differ because of the fraction of galaxies showing signs of ongoing star formation. The percentage of $``rejuvenated''$ ETGs is found to increase significantly in the lower density environments. It is not clear if this scenario is also able to explain the environmental dependence of the FP. As noticed in Sec.~\\ref{sec:offset_rband}, the variation of FP intercept in r band remains essentially the same when restricting the analysis to the subsample of ETGs with lower contamination from objects with non-genuine ETG morphology. Moreover, the FP offset is computed from the median values of the distributions of effective parameters and velocity dispersions, hence reflecting more closely the bulk properties of the ETG's population. Another important finding of the present work is the lack of any evidence for ETGs in the field to be less metal-rich than those in the cluster environment, as found, for instance, by~\\citet{GALL:06}. The optical+NIR variation of the FP offset implies a difference in metallicity of $-0.043 \\pm 0.023$ per decade in local galaxy density, corresponding to a difference of $\\sim 16 \\pm 8\\%$ between ETGs in low- and high-density environments. Hence, our FP analysis is consistent with field galaxies being more metal-rich than ETGs in groups. Rose et al.~(1994) found a similar result and interpreted this difference as reflecting the fact that star formation and chemical enrichment in ETGs in clusters were truncated at an early epoch, so ETGs were influenced by the environment at an early phase of galaxy evolution. For comparison, ~\\citet{BERN:06} found no detectable difference in the metallicity of field and cluster ETGs, while other studies, e.g.~\\citet{Thomas:05, deLaRosa:07, Clemens:09} found evidence for galaxies in dense environments to be less metal rich than those in the field. Different from most of the works listed above, we do not only analyse the average difference between the properties of group and field galaxies, but also follow the variation of FP coefficients as a function of local galaxy density. We find that the offset of the FP smoothly changes from the highest density regions, in the cores of parent galaxy groups, through the outskirt group regions and the field. Such variation is caused by a variation in age, i.e. ETGs have progressively younger stellar populations from the cluster cores through the field, in agreement with the recent results of~\\citet{Bernardi:09} and~PGF10, who found that, for a given parent halo mass, central galaxies are systematically older than satellite galaxies. In contrast to~PGF10, we find a positive (or eventually no) gradient in galaxy metallicity as a function of local density, that would imply centrals to be either less metal-rich or as metal-rich as satellites (see above). Our findings can be compared to the expectations of semi-analytic models (SAMs) of galaxy formation~\\citep{deLucia:06} (hereafter deL06). The model of deL06 predicts an age variation, from cluster cores to the field, of $\\sim 2$~Gyr (see their fig.~8). The trend is mostly due to the fact that, in a hierarchical scenario, halos in a region of the Universe doomed to become a cluster are those collapsing earlier and merging faster. This would also imply that the age variation as a function of local density might be larger for more massive halos, in agreement with our finding that the variation of the FP offset with the local environment is larger for galaxies residing in clusters, relative to groups. Over the large environmental range spanned by the SPIDER sample (three decades in local galaxy density), we find an age variation of $\\delta \\log t = 0.14 \\pm 0.02$, i.e. $\\delta t / t \\sim 32 \\pm 5\\%$ ($\\sim 3.2 \\pm 0.5$~Gyr, for a galaxy formation epoch of $10$~Gyr). For a pure age model of the variation of FP offset with environment (see Sec.~\\ref{sec:offset_sp}), this difference reduces to $\\delta t / t \\sim 28 \\pm 3 \\%$ ($\\sim 2.8 \\pm 0.3$~Gyr). Both values of $\\delta t /t $ are larger than that predicted from the SAM, but not much different considering all assumptions and different definitions of environments in both cases. We can also notice that SAMs predict field galaxies to be less-metal rich than their counterparts in groups, in contrast to our FP analysis. As discussed by PGF10, current SAMs still do not provide an accurate description of galaxy's metallicities: several recipes, such as the treatment of SN feedback (see e.g.~\\citealt{Bertone:07}), tidal interactions, and recycling of SN ejecta, might be able to improve the comparison with the the environmental trends seen in the data." }, "1003/1003.4811_arXiv.txt": { "abstract": "The two apparently distinct phenomena of dark energy (or late-time cosmic acceleration) and quantum gravity dominate physics on extremely low, and extremely high energies, but do not seem to have any apparent empirical connection. Nevertheless, the two have a theoretical connection, through the {\\it cosmological constant problem}. I argue that the finite temperature quantum gravitational corrections to black hole entropy yields a pressure for the gravitational vacuum (or gravitational aether). Assuming that the relative corrections are linear in horizon temperature (i.e. are suppressed by one power of Planck energy), the pressure is comparable to that of dark energy for astrophysical black holes. This implies that the observation of late-time cosmic acceleration may have provided us with the first precision measurement of quantum gravity, i.e. that of black hole entropy. ", "introduction": " ", "conclusions": "" }, "1003/1003.2430_arXiv.txt": { "abstract": "{Direct imaging of brown dwarfs as companions to solar-type stars can provide a wealth of well-constrained data to ``benchmark'' the physics of such objects, since quantities like metallicity and age can be determined from their well-studied primaries.} {We present results from an adaptive optics imaging program on stars drawn from the Anglo-Australian and Keck Planet Search projects, with the aim of directly imaging known cool companions. } {Simulations have modeled the expected contrast ratios and separations of known companions using estimates of orbital parameters available from current radial-velocity data and then a selection of the best case objects were followed-up with high contrast imaging to attempt to directly image these companions. } {These simulations suggest that only a very small number of radial-velocity detected exoplanets with consistent velocity fits and age estimates could potentially be directly imaged using the VLT's Simultaneous Differential Imaging system and only under favorable conditions. We also present detectability confidence limits from the radial-velocity data sets and show how these can be used to gain a better understanding of these systems when combined with the imaging data. For HD32778 and HD91204 the detectabilities help little in constraining the companion and hence almost all our knowledge is drawn from the SDI images. Therefore, we can say that these stars do not host cool methane objects, out to on-sky separations of $\\sim$2$''$, with contrasts less than 10--11~magnitudes. However, for HD25874, HD120780 and HD145825, the contrasts and detectabilities can rule out a number of possible solutions, particularly at low angular separations, and for the best case, down to strong methane masses of 40M$_{\\rm{J}}$ at 1$''$ separation. The contrast curves constructed for these five stars show 5$\\sigma$ contrasts ($\\Delta$F1) of $\\sim$9.2--11.5 magnitudes at separations of $\\ge$0.6$''$, which correspond to contrasts of $\\sim$9.7--12.0 magnitudes for companions of mid-T spectral type. Such limits allow us to reach down to 40M$_{\\rm{J}}$ around fairly old field dwarfs that typically constitute high precision radial-velocity programs. Finally, the analysis performed here can serve as a template for future projects that will employ extreme-AO systems to directly image planets already indirectly discovered by the radial-velocity method. } {} ", "introduction": "The detection of over 400 planets orbiting Sun-like stars has revolutionised our knowledge of our local neighbourhood and our position therein. Yet planets are not the sole close companions to solar-type stars. For instance, \\citet{duquennoy,duquennoy91} and \\citet{duquennoy92} have examined stellar multiplicity in a series of papers. Radial-velocity surveys have revealed few brown dwarfs orbiting solar-type stars (e.g. \\citealp{wittenmyer09}; \\citealp{jenkins09a}) leading to the phrase `brown dwarf desert' being coined to describe this paucity (\\citealp{marcy}). However, beyond $\\sim$4AU one would expect few radial-velocity planetary or brown dwarf companions to be known due to the limited temporal coverage at the required precision levels necessary to fully sample such companions. In addition, radial-velocity surveys also have strong biases against the detection of long-period companions, as the radial-velocity amplitude is a strong function of orbital period and also since this technique requires the observation of at least half an orbit (e.g. \\citealp{wright07}) to constrain companion properties. Only now are we sensitive enough to detect solar system-like gas giant planets in solar system-like orbits (e.g. \\citealp{jones10}). Conversely, direct and coronographic imaging techniques can probe much wider separations than current radial-velocity programs can reach. For example, \\citet{kalas08} and \\citet{marois08} have directly imaged planetary mass companions to the stars Fomalhaut and HR~8799, located at angular separations of 14.9$''$ and 1.73$''$, or 115AU and 68AU, respectively. \\citet{mccarthy04} found another deficit of brown dwarf companions between 75-1200 AU. \\citet{liu02} used the Gemini-North and Keck Adaptive Optics (AO) systems to obtain three epochs of images of the brown dwarf companion to HR~7672, which had initially been detected by its radial-velocity signature. The flux ratio at 2.16$\\mu$m was found to be 8.6 magnitudes at a separation of 0.79$''$. This level of contrast pushed the instrumentation used in this detection to its very limits. However the introduction of Simultaneous Differential Imaging (SDI) on the VLT's NACO facility permits the achievement of higher contrasts, at smaller separations, for the coolest stellar companions. For example, contrasts on the order of $\\Delta$H$\\sim$13 have been demonstrated at $\\sim$0.5$''$ by \\citet{mugrauer} and \\citet{biller07}. ", "conclusions": "We have performed a targeted direct imaging program to detect cool companions orbiting within 2$''$ of their parent stars. The stars were drawn from the AAPS and Keck planet search projects and consist of objects that exhibit large radial-velocity variation over several years. These radial-velocities indicate the presence of a massive companion on a long period orbit. Five stars were examined with the NACO-SDI system on the VLT in Paranal, Chile. From the five, two possible detections were found around the stars HD25874 and HD120780. However, further analysis of these detections indicate they are probable residual artifacts since they are found to be located at the same distance from the central pixel and the same position angle in each of the images for each star. In addition we also present detectabilities for each system by analysing the radial-velocity information we have acquired. Each of these detections lie within sensitivity boundaries for these stars of between 50-70\\%, meaning they could not be ruled out with any high degree of certainty from the radial-velocity data. We have summarised the results of this work in Table~\\ref{tab:results}, which shows both the broad-band AO and narrow-band SDI reduced contrasts and mass limits at separations of 0.5$''$ and 1.0$''$. Also the sensitivity limits from the radial-velocity data have been summarised at confidence limits of 70,~90~and~95\\% for each system, highlighting the minimum mass possibly detectable at each confidence level. The table shows the contrasts in magnitudes and mass limits in Jupiter-masses at 0.5$''$ and 1.0$''$ angular separations for AO and SDI images and shows the AO performs almost as well as the SDI at separations of $\\ge$1.0$''$. Also since the stars are fairly old the AO reduction tends to reach only a few Jupiter-masses above the SDI reduction for most of the sample, however when the star is fairly young (HD145825) the SDI reaches far deeper than AO i.e. 13M$_{\\rm{J}}$ in this case. Also from this table we see that the radial-velocity confidence limits are all generally the same, apart from HD32778 due its limited number of data points and temporal coverage. As the stars chosen are fairly old, a consequence of the radial-velocity selection method, the companions must be sufficiently massive to lie within the detection threshold of the instrument. Also since three of the five stars were found to have 'liner' trends, and we expect these to be sufficiently massive, then we suspect either the uncertainty in age means we are underestimating our mass thresholds, the models are over estimating the magnitudes of the companions, the companions are aligned such that they are found behind, or very close to the central star ($\\le$0.1$''$) or the companions are sufficiently far out in the system that they are off the 2$''$ contrast limit. A combination of the these mechanisms are probably at work. We find 5$\\sigma$ ($\\Delta$F1) contrasts of 11.5~magnitudes are possible using this method around bright F-K type stars (and 5$\\sigma$ $H$-band contrasts of 12~magnitudes for mid T-like objects). Such contrasts allow access to long period, massive ($\\sim\\ge$40M$_{\\rm{J}}$) methane objects for stars that typically constitute the bulk of radial-velocity programmes. In the future similar analyses as those employed here will lead to a greater understanding of the properties of exoplanets and brown dwarfs when extreme-AO systems can gain the contrasts necessary to directly image known planetary-mass companions detected by ongoing Doppler programs." }, "1003/1003.5951_arXiv.txt": { "abstract": "The Sweet-Parker layer in a system that exceeds a critical value of the Lundquist number ($S$) is unstable to the plasmoid instability. In this paper, a numerical scaling study has been done with an island coalescing system driven by a low level of random noise. In the early stage, a primary Sweet-Parker layer forms between the two coalescing islands. The primary Sweet-Parker layer breaks into multiple plasmoids and even thinner current sheets through multiple levels of cascading if the Lundquist number is greater than a critical value $S_{c}\\simeq4\\times10^{4}$. As a result of the plasmoid instability, the system realizes a fast nonlinear reconnection rate that is nearly independent of $S$, and is only weakly dependent on the level of noise. The number of plasmoids in the linear regime is found to scales as $S^{3/8}$, as predicted by an earlier asymptotic analysis (Loureiro \\emph{et al.}, Phys. Plasmas \\textbf{14}, 100703 (2007)). In the nonlinear regime, the number of plasmoids follows a steeper scaling, and is proportional to $S$. The thickness and length of current sheets are found to scale as $S^{-1}$, and the local current densities of current sheets scale as $S^{-1}$. Heuristic arguments are given in support of theses scaling relations. ", "introduction": "Recent studies of nonlinear reconnection in large high-Lundquist-number ($S$) plasmas, based on resistive magnetohydrodynamics (MHD) \\cite{BhattacharjeeHYR2009} as well as fully kinetic simulations that include a collision operator \\cite{DaughtonRAKYB2009} have produced a surprise. It is seen in these studies that above a critical value of the Lundquist number, the system deviates qualitatively from the predictions of Sweet-Parker theory \\cite{Sweet1958,Parker1963} which has been the standard model for reconnection in the high-$S$ regime. In the Sweet-Parker model, the reconnection layer has the structure of Y-points, with a length of the order of the system size, and a width given by $\\delta_{SP}=L/S^{1/2}$, where $S=LV_{A}/\\eta$ is the Lundquist number based on the system size $L$, the Alfv\\'en speed $V_{A}$, and the magnetic diffusivity $\\eta$. The Sweet-Parker model is usually considered to be a model of \\textquotedblleft{}slow\\textquotedblright{} reconnection as it predicts the reconnection rate to scale as $S^{-1/2}$. (The Petschek model \\cite{Petschek1964} predicts a much weaker dependence on $S$, with the maximum reconnection rate $\\sim1/\\log S$. However, it has become clear over the years that Petschek model is realizable only when the resistivity is locally enhanced around the reconnection site. \\cite{Biskamp2000}) For weakly collisional systems such as the solar corona, the Lundquist number is typically very large ($\\sim10^{12}-10^{14}$) and the Sweet-Parker reconnection time scale is of the order of years, much too slow to account for fast events such as solar flares. The Sweet-Parker model is based on the assumption of the existence of a long thin current layer. Although it has been known for some time that such a thin current layer may be unstable to a secondary tearing instability (referred to hereafter as the plasmoid instability) which generates plasmoids,\\cite{BulanovSS1979,LeeF1986,Biskamp1986,YanLP1992,ShibataT2001,Lapenta2008} it has been realized only fairly recently that the Sweet-Parker layer actually becomes more unstable as the Lundquist number increases, with a linear growth rate $\\gamma\\sim S^{1/4}$ and the number of plasmoids $\\sim S^{3/8}$.\\cite{LoureiroSC2007,BhattacharjeeHYR2009} In a recent paper (Ref. \\cite{BhattacharjeeHYR2009}, hereafter referred to as Paper I), Bhattacharjee\\emph{ et al}. have presented numerical results that suggest strongly that as a consequence of the plasmoid instability, the system evolves into a nonlinear regime in which the reconnection rate becomes weakly dependent on $S$. The primary goal of this paper is to strengthen the results obtained in Paper I in two significant ways: first, to present new simulation results with a modified initial condition that enables us to obtain stronger scaling results on the nonlinear reconnection rate and the number of plasmoids generated in the nonlinear regime, and second, a simple heuristic model that is consistent with the results of the simulation and fortifies the claim in Paper I that the reconnection rate in the nonlinear regime of the plasmoid instability is fast and independent of $S$. ", "conclusions": "In summary, we have shown through a series of simulations that resistive MHD can achieve a fast reconnection rate in the high-Lundquist-number regime. Fast reconnection is facilitated by the plasmoid instability. The resultant reconnection rate is independent of $S$ and is weakly dependent on the noise level. We have verified the $S^{3/8}$ scaling of the number of plasmoids in the linear regime, as predicted in Refs. \\cite{LoureiroSC2007,BhattacharjeeHYR2009}. In the nonlinear regime, the number of plasmoids follows a steeper scaling and is proportional to $S$. We also have done statistical studies of the local current sheets, and found that the current sheet thickness and length both scale as $S^{-1}$, while the current density scales as $S$. These findings are consistent with our heuristic argument and the claim that the reconnection rate is independent of $S$ in the high-$S$ regime. The fast reconnection rate we have obtained is approximately $0.01V_{A}B$, which is similar to the values from other recent resistive MHD studies,\\cite{BhattacharjeeHYR2009,CassakSD2009} but is smaller than the typical reconnection rate from collisionless two-fluid or particle-in-cell simulations by an order of magnitude. Which rate will be realized depends on how collisional the system is. If the Sweet-Parker thickness $\\delta_{SP}$ is greater than the ion skin depth $d_{i}$ (or the ion Larmor radius at the sound speed, $\\rho_{s}$, if there is a guide field) in a system, a Sweet-Parker layer is likely to form first. On the other hand, if $\\delta_{SP}d_{i}$ (or $\\rho_{s}$) but $S>S_{c}$. Then we expect the plasmoid instability to set in and the primary Sweet-Parker layer will break into segments. This brings the thickness further down to $\\delta\\sim\\delta_{SP}(S_{c}/S)^{1/2}$. If $\\delta>d_{i}$ (or $\\rho_{s}$) then the system is still dominated by collisional effects and we may end up getting a reconnection rate of $0.01V_{A}B$. However, if $\\delta2$ galaxy bright enough to be observed spectroscopically is driving out material at velocities of at least several hundred \\kms, but we do not know how far this material travels, or even where it is with respect to the galaxy as we observe it. The mass flux associated with such flows has been measured in only one case at high redshift, MS1512-cB58 (\\citealt{pettini00}), and even in this case the result depends very sensitively on the assumed physical location of the absorbing material responsible for the bulk of the observed absorption. The best hope for constraining the location of outflowing gas is by observing objects lying in the background, but at small angular separation, relative to the galaxy of interest. In this case the challenge is to find background objects bright enough in the rest-frame far-UV but close enough to the foreground galaxy to provide interesting constraints. A great deal of effort, over a large range of redshifts, has been invested using QSOs as background sources, where absorption by metallic ions or H{\\sc I} in the spectrum of the QSO is compared with galaxies with known redshifts and projected separations (e.g., \\citealt{bergeron91, steidel94,chen2001,steidel02,lanzetta95,danforth08,bowen95,bouche07,kaczprak10}); most of this work has focused on redshifts $z < 1$ because of the increasing difficulty obtaining spectra of the foreground galaxies at higher redshifts. Even if galaxies are identified and have redshifts that correspond closely with observed absorption, the association of particular absorption systems with identified galaxies is almost always ambiguous, since the dynamic range for identifying faint galaxies is limited, and often the bright background QSOs make it challenging to observe galaxies within a few arcsec of the QSO sightline. Finally, there is the controversial issue of whether metals seen near to, but outside of, galaxies are a direct result of recent star formation or AGN activity, or are simply tracing out regions of the universe in which {\\it some} galaxies, perhaps in the distant past, polluted the gas with metals (e.g., \\citealt{madau01,scannapieco02,mori02, ferrara03} ; cf. \\citealt{assp03,ass+05}). In this paper, our goal is to try to understand the spatial distribution of cool gas seen in absorption against the stellar continuum of every galaxy observed at high redshift. The objective is to empirically track the kinematics and structure in the CGM from the central parts of galaxies all the way to large galactocentric radii. In this work, we use only galaxy spectra, primarily in the rest-frame far-UV, but we calibrate the velocity zero-point using a set of nearly 100 H$\\alpha$ measurements in the observed-frame near-IR for galaxies in the redshift range $1.9 \\simlt z \\simlt 2.6$. The near-IR measurements are drawn primarily from the sample of \\cite{erb+06b}, after which we use a much more extensive set of rest-UV spectra from a nearly completed UV-selected galaxy survey targeting the same range of redshifts. The paper is organized as follows: in \\S\\ref{sec:bulkvel} we examine the statistics of the kinematics of the outflows using a sub-sample of galaxies for which both near-IR nebular \\Ha\\ spectroscopy and reasonably high-quality optical (rest-UV) spectra are available. We also present new empirical formulae for estimating galaxy systemic redshifts for the typical case in which only low S/N rest-UV spectra are available. In \\S\\ref{sec:correlations}, we seek correlations between the interstellar absorption line kinematics, in particular the bulk velocities measured from strong low-ionization transitions, and other measured galaxy properties. \\S\\ref{sec:composites} describes further inferences on the structure and kinematics of outflowing material from high S/N composite far-UV spectra, while \\S\\ref{sec:lya} examines the observed behavior of \\lya\\ emission and its relationship to the IS absorption features, and attempts to understand the observations with simple models. We introduce in \\S\\ref{sec:galgal} the use of close angular pairs of galaxies at different redshifts for mapping the spatial distribution of the circumgalactic gas around the foreground galaxies, while \\S\\ref{sec:model} develops a simple geometrical and kinematic model for outflows consistent with both the line profiles in galaxy spectra and the larger-scale distribution of gas in the CGM. \\S\\ref{sec:model_implications} discusses the observational results and their interpretation in the context of the models, and \\S\\ref{sec:discussion} summarizes the conclusions and discusses the prospects for improvement in the future. We assume a $\\Lambda$-CDM cosmology with $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$, and $h=0.7$ throughout, unless specified otherwise. ", "conclusions": "\\label{sec:discussion} We have used a relatively large sample of $1.9 \\simlt z \\simlt 2.6$ galaxies with accurate measurements of systemic redshift $z_{sys}$ and reasonably high quality rest-frame far-UV spectra to examine the relationship between galaxy properties and UV spectral morphology. Our main focus has been on the kinematics and strength of interstellar absorption and \\lya\\ emission and their implications for the galaxy-scale outflows observed in all rapidly star-forming galaxies at high redshifts. Using this well-observed subset as a calibration, we then combined the rest far-UV galaxy spectra drawn from a much-larger parent sample with additional spatial information for the same ensemble of galaxies provided by the spectra of background galaxies with small angular separations. Using these joint constraints, we constructed simple models of the kinematics and geometry of the ``circumgalactic medium''-- the interface between star forming galaxies and the IGM. Our principal conclusions are as follows: 1. We have used the \\Ha\\ sample of $z\\simeq 2-2.6$ galaxies together with their far-UV spectra to produce a revised calibration that allows deriving systemic redshifts for star-forming galaxies from their far-UV spectral features (strong IS lines and \\lya\\ emission). In the absence of stellar photospheric absorption lines or nebular emission lines in the rest-frame optical, the most accurate estimates of the systemic redshift are derived from the centroids of strong IS absorption lines. Applying a shift of $\\Delta v = +165$ \\kms ($\\Delta z = +0.0018$) at $z=2.3$ to a measured absorption redshift provides an estimate of a galaxy's systemic redshift accurate to $\\pm 125$ \\kms; redshifts measured from the \\Ha\\ emission line (when available) have a precision of $\\pm 60$ \\kms. Both methods for estimating redshifts have no significant systematic offset relative to the redshift defined by stellar (photospheric) absorption features. 2. Using only the \\Ha\\ sample of BX galaxies, we found mean velocity offsets for the centroids of the strong IS lines and \\lya\\ emission line of $\\dvis = -164\\pm16$ \\kms and $\\dvla = +445\\pm27$ \\kms, respectively. We searched for significant correlations between the kinematics defined by the centroid velocities of IS and \\lya\\ lines and other measured or inferred galaxy properties. Within the \\Ha\\ sample, the only significant correlations found were between $\\dvis$ and galaxy mass estimated from the sum of inferred stellar and gas mass ($M_{bar}$) as well as the independently estimated $M_{dyn}$. The sense of the correlation is that $\\dvis$ is smaller (i.e., less blue-shifted) in galaxies with larger masses. 3. Despite the trend described in point 2 above, the velocity $|v_{max}|$ of the maximum blue-shift observed in the absorption profiles is essentially identical for the sub-samples of galaxies with ${\\rm M_{bar}} > 3.7\\times 10^{10}$ M$_{\\sun}$ and ${\\rm M_{bar}} < 3.7\\times10^{10}$ M$_{\\sun}$, with $|v_{max}| \\simeq 800$ \\kms. However, the higher-mass subset has both $\\dvis$ and $\\dvla$ shifted toward positive velocities (i.e., more redshifted) by $\\sim 200$ \\kms. The differences in the composite IS and \\lya\\ profiles can be explained by an additional component of absorption at $v \\simgt 0$ that appears to be absent in the lower-mass sub-sample. The extra absorption, which has a peak near $v = 0$ and a centroid velocity of $v\\simeq +150$ \\kms, could conceivably be a signature of infalling gas or stalled winds falling back on the galaxy, but is most likely to be gas at small galactocentric radii based on the line profiles. In any case, the velocity centroids of the IS lines, which are commonly used as a proxy for the ``wind velocity'' for high redshift galaxy samples, are actually modulated almost entirely by gas that is not outflowing. 4. The \\lya\\ emission profile, like $\\langle \\dvis \\rangle$, is modulated by the covering fraction (or apparent optical depth) of material near $v \\simeq 0$. We show that, by using the information on the covering fraction near $v \\simeq 0$, the kinematics of blue-shifted gas, and assuming spherical symmetry, the behavior of \\lya\\ emission with respect to IS absorption can be reproduced using simple models. In general, the red wing of \\lya\\ emission is produced by \\lya\\ photons scattered from outflowing gas on the opposite side of the galaxy. The apparent redshift of the \\lya\\ centroid is a manifestation of the fact that only photons scattering from material having a (redshifted) velocity large enough to take the photons off the \\lya\\ resonance for any material between the last scattering and the observer can escape in the observer's direction. The spectral morphologies of \\lya\\ emission and IS absorption are most easily understood if the highest velocity material (either redshifted or blue-shifted) is located at the largest distances from the galaxy, i.e., that $v(r)$ is monotonically increasing with $r$ in a (roughly) spherically symmetric radial flow. The same geometric picture works well to reproduce both the line-of-sight ($b=0$) absorption profiles and the absorption line strength as a function of impact parameter (see point 6 below). 5. We have demonstrated that the use of the far-UV spectra of galaxies within projected angular pairs (with discrepant redshifts) provides an opportunity to constrain the physical location of circumgalactic gas around typical galaxies. A set of 512 such galaxy pairs, on angular scales ranging from 1\\arcs\\ to 15\\arcs\\ ($\\simeq 8-125$ kpc), have been combined with the $b=0$ spectra of the foreground galaxies in each pair. These are used together to measure each foreground galaxy's CGM along two independent lines of sight. Composite spectra stacked according to galactocentric impact parameter allow us to measure the dependence of the line strength of several ionic species on impact parameter $b$ from 0-125 kpc (physical). These lines are strongly saturated, so their strength is determined by a combination of the covering fraction and the velocity spread in the absorbing gas. The fall-off of $W_0$ with $b$ implies $f_c(r) \\propto r^{-\\gamma}$, with $0.2 \\le \\gamma \\le 0.6$, with \\lya\\ and C{\\sc IV} having smaller values of $\\gamma$ compared to (e.g.) Si{\\sc II} and Si{\\sc IV}. For each observed species, the strength of absorption declines much more rapidly beginning at $b \\simeq 70-90$ kpc, with the exception of \\lya\\, which remains strong to $b \\simeq 250$ kpc. Combining the known space density of the galaxies in our sample (i.e., including only those with apparent magnitude ${\\cal R} \\le 25.5$) with the average cross-section for absorption accounts for $\\simeq 45$\\% of all intergalactic \\ion{C}{4} absorption with $W_0(1548) > 0.15$ \\AA\\, and $\\simeq 70$\\% of LLSs at $z \\sim 2$ observed along QSO sightlines. 6. We have proposed a simple model constrained by a combination of $W_0(b)$ from the galaxy pairs and the kinematics of IS species observed toward each galaxy's own stars ($b=0$). The IS absorption line shapes depend on a combination of $v_{out} (r)$, the radial dependence of outflow velocity which in turn depends on the cloud acceleration $a(r)$ and $f_c(r,v)$, the radial dependence of the covering fraction, which depends on cloud geometry and is inferred from $W_0(b)$. We show that a self-consistent model using parametrized forms of $a(r)$ and $f_c(r)$ can simultaneously reproduce the observed line shapes and strengths (apparent optical depth vs. velocity) as well as the observed $W_0(b)$ relation from the galaxy-galaxy pair measurements. In the model, higher velocity gas is located at larger galactocentric radii, and gas clouds are accelerated to $\\simeq 800$ \\kms before the absorption strength (i.e., covering fraction) is geometrically diluted. There is little evidence in the line profiles for stalling or infalling wind material, and the transverse sightlines show metal-enriched gas to at least $\\sim 125$ kpc\\footnote{\\cite{assp03}, and \\cite{asp+05} have argued that there is evidence that the metals affect regions of $\\simeq 300$ kpc (proper) based on \\ion{C}{4}--galaxy cross-correlation, and that outflows would be expected to stall at approximately this distance due to a number of separate considerations.} Taken together, these suggest that high velocity outflows from $\\simeq L*$ galaxies at $z \\sim 2-3$ are able to deposit enriched material to very large radii. They also suggest that the observed metals are causally related to the observed galaxies, rather than remnants of winds generated by other (earlier) galaxies deposited in the nearby volume. 7. The observed $v_{max} \\simeq 800$ \\kms\\ that appears to be a general property of the galaxies in our spectroscopic sample exceeds $v_{esc}(r_{vir})$ estimated for galaxy halo masses of $\\simeq 9\\times10^{11}$ M$_{\\sun}$, the average halo mass expected given the observed space density and clustering of the $z \\simeq 2$ galaxies. Similarly, the galactocentric radius at which the strength of low-ionization absorption line species begins to decrease rapidly is similar to the estimated virial radius of $\\simeq 80$ kpc for the same halo mass. 8. We consider the observations in the context of ``cold accretion'' or ``cold flows'', which are believed to be active over the same range of galactocentric radii and redshifts, for galaxies in the same range of total mass as those in our sample ($4\\times10^{11} \\simlt M_{tot} \\simlt 2\\times 10^{12}$ M$_{\\sun}$). There is possible evidence for the presence of infalling gas in the far-UV spectra of galaxies with greater than the median baryonic mass $M_{bar} = 4\\times10^{10}$ M$_{\\sun}$, while the lower-mass sub-sample shows no evidence for the expected redshifted components of IS absorption lines. The kinematics of \\lya\\ emission also seem inconsistent with the expectations for accreting material, while they are as expected in the context of the picture of outflows we have presented. We expect that the same gas seen in absorption to $b\\simeq 50-100$ kpc also acts as a scattering medium for escaping \\lya\\ photons initially produced in the galaxy \\ion{H}{2} regions. These \\lya\\ ``halos'' are probably a generic property of all high redshift LBG-like galaxies, the more extreme of which have probably already been observed as ``\\lya\\ Blobs'' (\\citep{steidel00,matsuda04}). More typical LBGs likely exhibit this diffuse \\lya\\ emission, which is of low surface brightness and therefore difficult to observe without extremely deep \\lya\\ images. 9. The inferred mass of cool gas in the CGM (within $\\sim 100$ kpc) is comparable to the sum of cold gas and stars in the inner few kpc; together they account for $\\simgt 50$\\% of the total baryonic mass ($\\sim 1.5\\times10^{11}$ M$_{\\sun}$) associated with the average $M\\sim 9\\times 10^{11}$ M$_{\\sun}$ dark matter halo. If the qualitative picture of the CGM we have proposed is correct, the wind material would still be traveling at $v_{out} > v_{esc}(r_{vir})$ when it crosses the virial radius at $r \\simeq 80$ kpc. Even if $v_{out} \\simeq v_{esc}$, the low-ionization wind material would reach $r \\sim 100$ kpc within $\\sim 150$ Myr of the onset of the star formation episode if it is slowed only by gravity and does not accumulate a large amount of swept-up ISM on its way. Since we have shown that the CGM is outflowing and itself comprises a large baryonic reservoir, it is possible that in most directions there is little to impede outflows from propagating ballistically into the IGM. The use of resolved background sources for studying the CGM in absorption is qualitatively different from using QSO sightlines. It seems clear that the multi-phase ISM observable in a number of ionic species indicates differential changes in covering fraction of the outflowing gas as a function of ionization level. This means that the radial behavior of the strength of saturated transitions is providing information on the ``phase-space'' density of gas having physical conditions amenable to the presence of each ion. Because absorption seen in the spectra of background {\\it galaxies} averages over a $\\simeq 2-3$ kpc region, as compared to background QSOs, whose apparent angular sizes are $\\simeq 5$ orders of magnitude smaller, statistics such as IS line strength, covering fraction, and velocity range vs. impact parameter, are much less subject to large variance along a single sightline. All of the galaxy spectra used in this paper were obtained as part of large surveys, and therefore have low to moderate spectral resolution and S/N. It is possible with current generation telescopes and instruments to obtain spectra of relatively high S/N and moderate resolution (e.g., \\citealt{shapley06}) in 10-20 hours integration time (i.e., $\\sim 10$ times longer than most in the current samples). However, spectra similar to those currently possible only for strongly lensed galaxies ($R \\sim 5000$, S/N$\\simgt 30$) will be routine at ${\\cal R} \\simeq 24-24.5$ using future 30m-class telescopes equipped with state of the art multi-object spectrometers (\\citealt{steidel+astro2010}). When such capability is realized, full observational access to the three-dimensional distribution of gas near to and between galaxies will revolutionize the study of baryonic processes in the context of galaxy formation. The results described above are intended as an illustration of the application of rest-far-UV galaxy spectra toward a better understanding of feedback processes and the CGM/IGM interface with the sites of galaxy formation. The feasibility of sharpening our understanding and interpretation with higher quality data in the future is very encouraging. \\bigskip \\bigskip This work has been supported by the US National Science Foundation through grants AST-0606912 and AST-0908805 (CCS), and by the David and Lucile Packard Foundation (AES, CCS). CCS acknowledges additional support from the John D. and Catherine T. MacArthur Foundation. DKE is supported by the National Aeronautics and Space Administration under Award No. NAS7-03001 and the California Institute of Technology. We would like to thank Juna Kollmeier and Joop Schaye for interesting and useful conversations, and Patrick Hall, Martin Haehnelt, and the referee for comments that significantly improved the final version of the paper. Kurt Adelberger played a major role in the early days of the survey used in this paper; his intellectual contributions have remained crucial though he has moved on to new challenges in the ``real'' world. Marc Kassis and the rest of the W.M. Keck Observatory staff keep the instruments and telescopes running effectively, for which we are extremely grateful. We wish to extend thanks to those of Hawaiian ancestry on whose sacred mountain we are privileged to be guests." }, "1003/1003.1762_arXiv.txt": { "abstract": "We present here observations of the transit of WASP-10b on 14 October 2009 UT taken from the University of Arizona's 1.55 meter Kuiper telescope on Mt. Bigelow. Conditions were photometric and accuracies of 2.0 mmag RMS were obtained throughout the transit. We have found that the ratio of the planet to host star radii is in agreement with the measurements of Christian et al. (2008) instead of the refinements of Johnson et al. (2009), suggesting that WASP-10b is indeed inflated beyond what is expected from theoretical modeling. We find no evidence for large ($> 20$ s) transit timing variations in WASP-10b's orbit from the ephemeris of Christian et al. (2008) and Johnson et al. (2009). ", "introduction": "Transiting Extrasolar planets are relatively rare among all known planet discoveries (69 transiting planets out of 429 known planetary systems)\\footnote{http://exoplanet.eu}. However, transits are unique in that they allow for the direct measurement of the radius of the transiting planet relative to the host star. Combined with radial velocity data, it is possible to determine the density of the transiting planet as well. Knowledge of the heating of the star can allow models of the bulk properties of these planets to be formulated and then compared to observation (for example, Baraffe et al. 2008; Fortney et al. 2007; Burrows et al 2007). Since the depth of the transit is directly related to its radius, the size of a transiting planet is one of the easiest parameters to measure. This makes transiting planets especially interesting to discover and study in detail, however this method becomes increasingly difficult when trying to probe to lower mass planets, although the Kepler Space Telescope (Borucki et al. 2008) is expected to detect many planets as small, or smaller, than the Earth. While most transiting planets are in nearly circular orbits, there are a small number of planets that maintain eccentric orbits despite being older than the circularization time scale for the system. The most notable of these systems are the transiting Neptunes GJ 436b (Butler et al. 2004) and HAT-P-11b (Bakos et al. 2010). GJ 436b has a significantly nonzero eccentricity of $0.15 \\pm 0.012$ (Deming et al. 2007), despite being older than its circularization time scale (Maness et al. 2007), which has led to speculation about possible additional (possibly resonant) planets in the system gravitationally ``pumping\" the eccentricity. Gravitational interaction with a third body would also lead to transit timing variations as the line of nodes precesses, but searches for this effect have found no evidence for transit timing variations in this system (Pont et al. 2009; Ballard et al. 2010). A search in the similar HAT-P-11 exo-Neptune system has yielded no positive indication for large transit timing variations (Dittmann et al. 2009a). WASP-10b is a recently discovered 2.96 M$_J$, 1.28 R$_J$ hot Jupiter in orbit around a K5 dwarf star (Christian et al. 2008). Due to the small size of the star ($0.775^{+0.043}_{-0.040} R_\\sun$), the transit depth is relatively deep, at 29 mmag (Christian et al. 2008). WASP-10b is in a $3.0927636^{+0.0000094}_{-0.000021}$ day, $0.059^{+0.014}_{-0.004}$ eccentricity orbit (Christian et al. 2008). Christian et al. (2008) note that the existence of a significantly nonzero eccentricity is surprising given the system's age (600 Myr - 1Gyr), and that studies into tidal dissipation by Jackson et al. (2008) suggest that WASP-10b's orbit should currently be circularized. Johnson et al. (2009) recently performed high quality follow up observations of a transit of WASP-10b using the Orthogonal Parallel Transfer Imaging Camera (OPTIC) on the University of Hawaii 2.2 m telescope on Mauna Kea. Using OPTIC's unique charge transfer\\footnote{To our knowledge, this was the first time this technique was used for an extrasolar planet transit.}, they were able to significantly improve the precision in the measurement of various parameters of the system, including the transit time to within 7 seconds uncertainty and the planetary radius to within $1.8\\%$ uncertainty. Johnson et al. (2009) note that their error in the planetary radius is dominated by the error in the radius of the star. Furthermore, their value of $1.080 \\pm 0.020$ R$_J$ represents a surprising $16\\%$ downward change ($5\\sigma$ lower than the value of Christian et al. (2008)) in the radius of the planet, corresponding to the value expected by models from Fortney et al. (2007). An attempt to confirm a smaller radius for WASP-10 and investigate the possiblity of transit timing variations in the system due to a second planet acting as an eccentricity pump were the primary motivations for this paper. ", "conclusions": "We have investigated one follow-up transit of WASP-10b in order to investigate the ($5\\sigma$) discrepancy between the planet radii measurements of Christian et al. (2008) and Johnson et al. (2009). We have found that, while Johnson et al. (2009) apparently collected more precise data, our radii measurement agrees with that of Christian et al. (2008) and that WASP-10b does appear to be inflated beyond the level expected from theoretical models. Furthermore, we have no evidence that WASP-10b's eccentricity is due to perturbation by a third body. If there was a third body in the system, then we would expect successive transits of WASP-10b to exhibit signs of transit timing variations. However, we have found that our ephemeris is consistent to within $2 \\sigma$ of Christian et al. (2008) and Johnson et al. (2009)'s ephemeris, and therefore large TTVs are unlikely in this system." }, "1003/1003.2627_arXiv.txt": { "abstract": "{\\em Spitzer Space Telescope} IRAC $3-8\\,\\mu$m and {\\em AKARI} IRC $2-4\\,\\mu$m photometry are reported for ten white dwarfs with photospheric heavy elements; nine relatively cool stars with photospheric calcium, and one hotter star with a peculiar high carbon abundance. A substantial infrared excess is detected at HE\\,2221$-$1630, while modest excess emissions are identified at HE\\,0106$-$3253 and HE\\,0307$+$0746, implying these latter two stars have relatively narrow ($\\Delta r < 0.1\\,R_{\\odot}$) rings of circumstellar dust. A likely 7.9\\,$\\mu$m excess is found at PG\\,1225$-$079 and may represent, together with G166-58, a sub-class of dust ring with a large inner hole. The existence of attenuated disks at white dwarfs substantiates the connection between their photospheric heavy elements and the accretion of disrupted minor planets, indicating many polluted white dwarfs may harbor orbiting dust, even those lacking an obvious infrared excess. ", "introduction": "The large-scale composition of extrasolar minor planets can be indirectly measured using white dwarfs as astrophysical detectors \\citep{jur09b,zuc07}. Rings of warm dust (and in some cases gaseous debris) revealed at over one dozen white dwarfs \\citep{far09a} are situated within the Roche limits \\citep{gan06} of their respective stellar hosts, and likely originated via the tidal disruption of minor (or possibly major) planets \\citep{jur03}. These closely orbiting disks rain metals onto the stellar photosphere, and contaminate an otherwise pristine hydrogen or helium atmospheric composition. For any particular cool white dwarf, the timescales for individual heavy elements to sink below the high gravity photosphere differ by less than a factor of three, yet are always orders of magnitude shorter than the cooling age \\citep{koe09a}. With an appropriate treatment of the accretion history, photospheric heavy element abundances can be used as a measure of the composition of the tidally destroyed, polluting parent body or bodies \\citep{kle10,jur08}. The spectacularly debris-polluted white dwarf GD\\,362 is an excellent example of this potential. Its circumstellar disk re-emits at least 3\\% of the incident stellar radiation; one-third of this is carried by a strong $10\\,\\mu$m silicate emission feature \\citep{jur07b}. The star itself is polluted by at least 15 elements heavier than hydrogen and helium, in an arrangement that is rich in refractory elements and deficient in volatiles; a pattern that broadly mimics the inner Solar system, and comparable with the outer composition of the Earth and Moon \\citep{zuc07}. A recent analysis of the unusual atmospheric mix of hydrogen and helium, together with X-ray constraints on the current hydrogen accretion rate at GD\\,362, suggests that if a single parent body gave rise to both the currently observed disk and the panoply of heavy elements in its atmosphere, then its mass would be larger than Callisto, and possibly larger than Mars \\citep{jur09b}. Observations of a large number and variety of metal-polluted white dwarfs are necessary to better understand and constrain the connection between rocky extrasolar parent bodies, the circumstellar dust and gas, and the subsequent photospheric pollutions they create in white dwarfs \\citep{zuc03}. The lifetimes of the disks at white dwarfs are poorly constrained at present, though there is some indication their lifetimes do not significantly exceed $10^5$\\,yr \\citep{far09a,kil08}. Another outstanding issue is the nature of the large metal abundances in stars lacking excess infrared emission; tenuous, gaseous, or previously-accreted disks are strong possibilities. While the evolution of optically thick dust at white dwarfs should be dominated by viscous dissipation-spreading, to date only G166-58 stands out as a possible indicator of how competing mechanisms, such as additional impacts or gas drag at the inner disk edge, may play a role \\citep{jur08,far08}. {\\em Spitzer} IRAC and {\\em AKARI} IRC photometric imaging observations are presented for nine metal-polluted white dwarfs chosen for their high calcium abundances and relatively short cooling ages. The observations are presented in \\S2, the results in \\S3, and the implications for the minor planet accretion hypothesis and future dust searches are discussed in \\S4. \\begin{deluxetable*}{ccccccc} \\tabletypesize{\\footnotesize} \\tablecaption{{\\em Spitzer} and {\\em AKARI} \\ White Dwarf Targets\\label{tbl1}} \\tablewidth{0pt} \\tablehead{ \\colhead{WD}\t\t\t\t\t& \\colhead{Name}\t\t\t\t& \\colhead{SpT}\t\t\t\t\t& \\colhead{$V$}\t\t\t\t\t& \\colhead{[Ca/H(e)]}\t\t\t\t& \\colhead{Telescope}\t\t\t\t& \\colhead{Refs}\t\t\t\t\t\\\\ & & &(mag) & & &} \\startdata 0047$+$190\\tablenotemark{a}\t&HS\\,0047$+$1903\t&DAZ\t&16.1\t&$-6.1$\t\t&{\\em Spitzer}\t\t\t&1\\\\ 0106$-$328\\tablenotemark{a}\t&HE\\,0106$-$3253\t&DAZ\t&15.5\t&$-5.8$\t\t&{\\em Spitzer}\t\t\t&1\\\\ 0307$+$077\\tablenotemark{a}\t&HS\\,0307$+$0746\t&DAZ\t&16.4\t&$-7.1$\t\t&{\\em Spitzer}\t\t\t&1\\\\ 0842$+$231\\tablenotemark{a}\t&Ton\\,345 \t\t&DBZ\t&15.9\t&$-6.9$\t\t&{\\em AKARI}\t\t\t&2\\\\ 1011$+$570\t\t\t\t&GD\\,303\t\t\t&DBZ\t&14.6\t&$-7.8$\t\t&{\\em Spitzer / AKARI}\t&3\\\\ 1225$-$079\t\t\t\t&PG\t\t\t\t&DZAB\t&14.8\t&$-7.9$\t\t&{\\em Spitzer}\t\t\t&3,4\\\\ 1542$+$182\t\t\t\t&GD\\,190\t\t\t&DBQ\t&14.7\t&\\nodata\t\t&{\\em Spitzer}\t\t\t&5\\\\ 1709$+$230\t\t\t\t&GD\\,205\t\t\t&DBAZ\t&14.9\t&$-8.0$\t\t&{\\em Spitzer / AKARI}\t&4,6\\\\ 2221$-$165\\tablenotemark{a}\t&HE\t2221$-$1630\t&DAZ\t&16.1\t&$-7.2$\t\t&{\\em Spitzer}\t\t\t&1\\\\ 2229$+$235\\tablenotemark{a}\t&HS\t2229$+$2335\t&DAZ\t&15.9\t&$-5.9$\t\t&{\\em Spitzer}\t\t\t&1 \\enddata \\tablerefs{ 1) \\citealt{koe06} 2) \\citealt{gan08} 3) \\citealt{wol02} 4) \\citealt{koe05} 5) \\citealt{pet05} 6) \\citealt{vos07}} \\tablenotetext{a}{The WD numbers for these stars are unofficial designations, but correctly reflect the conventional use of epoch B1950 coordinates. Abundances are expressed as [X/Y] = $\\log\\,[ n({\\rm X})/n({\\rm Y})]$.\\\\} \\end{deluxetable*} ", "conclusions": "Three (possibly four) newly identified dust disks orbiting metal-polluted white dwarfs are discovered by {\\em Spitzer} IRAC as mid-infrared excess emission over $3-8\\,\\mu$m. The fractional infrared (dust continuum) luminosity from HE\\,2221$-$1630 is relatively large and around 0.008. In contrast, HE\\,0106$-$3253 and HE\\,0307$+$0746 with their much smaller infrared luminosities are modeled by narrow rings of orbiting dust. Even narrower disks at white dwarfs are possible, and perhaps likely given the necessity for ongoing heavy element pollution at the bulk of DAZ stars. An attenuated, flat disk as narrow as half an earth radius can harbor up to $10^{22}$\\,g of material and be sufficient to reproduce the inferred 1) metal accretion rates of DAZ stars for at least 10$^4$\\,yr, and 2) mass of heavy elements in the convective envelopes of DBZ stars. Yet such a tenuous ring of dust would not exhibit an infrared excess at the 10\\% level for any inclination. The only evidence of accretion may be the photospheric pollution; other signatures such as X-ray emission seen in accreting binaries require much higher mass infall rates than inferred for metal-enriched white dwarfs. Therefore, elemental abundances for polluted stars without an infrared excess are still likely to represent the bulk composition of accreted minor planets. PG\\,1225$-$079 is a tantalizing potential disk that emits only at 7.9\\,$\\mu$m, but which may confirm a previously suspected class of dust ring with a large inner hole. While the excess emission from G166-58 is somewhat similar in temperature, a flat ring model predicts the inner hole at PG\\,1225$-$079 would be significantly larger and may represent a disk that is near to being fully accreted; forthcoming data may test this scenario." }, "1003/1003.0308_arXiv.txt": { "abstract": "{The existence of time-energy correlations in flare occurrence is still an open and much debated problem.} {This study addresses the question whether statistically significant correlations are present between energies of successive flares as well as energies and waiting times.} { We analyze the GOES catalog with a statistical approach based on the comparison of the real catalog with a reshuffled one where energies are decorrelated. This analysis reduces the effect of background activity and is able to reveal the role of obscuration.} { We show the existence of non-trivial correlations between waiting times and energies, as well as between energies of subsequent flares. More precisely, we find that flares close in time tend to have the second event with large energy. Moreover, after large flares the flaring rate significantly increases, together with the probability of other large flares.} {Results suggest that correlations between energies and waiting times are a physical property and not an effect of obscuration. These findings could give important information on the mechanisms for energy storage and release in the solar corona. } ", "introduction": "Solar flares are violent explosions of magnetic energy in the solar corona. The theory of magnetic reconnection (see Priest \\& Forbes \\cite{Pri} for a review) represents the most plausible and widely accepted explanation for flare occurrence, although a definite and clear explanation of the mechanisms at the basis of flare triggering is still lacking. A well-established property of flare occurrence is the power law decay of the peak-flux energy distribution (Lee et al \\cite{Lee}; Aschwanden et al \\cite{Asc}; Crosby et al \\cite{Cro}), a property shared by other stochastic physical phenomena like earthquakes (de Arcangelis et al \\cite{deA}). Several mechanisms have been proposed to reproduce the above experimental findings. In the Rosner and Vaiana (RV) model (Rosner \\& Vaiana \\cite{Ros}) it is assumed that flare occurrence is an uncorrelated Poisson process where the energy grows at a rate proportional to the internal energy of the system. Given this assumption, the energy storage follows an exponential temporal growth, interrupted at random when the system releases all the stored energy in a flare. This mechanism accounts for the power law in the size distribution and predicts correlations between the storage time and the released energy. More precisely, the later a flare will occur the higher the energy will be. Also the avalanche model (Lu \\& Hamilton \\cite{Lu}; Hamon et al \\cite{Ham}), which describes solar flares as energy relaxation events in a system driven at a constant rate, correctly predicts scale-free behavior for the peak-flux energy distribution. Within this approach, flare occurrence is again a Poisson process, as in the RV model, but occurrence times and energies are uncorrelated like in standard self-organized models (Jensen \\cite{jen}). Finally, power law behavior for the size distribution is also consistent with models assuming correlations between the waiting times $\\Delta t$ of successive bursts, as in a shell model of magneto-hydrodynamic turbulence (Boffetta et al \\cite{Bof}). This model reproduces the experimental power law decay of the waiting time distribution $P(\\Delta t)$ which can also be obtained however by means of an uncorrelated piecewise Poisson process (Wheatland et al \\cite{Whe3}; Norman et al \\cite{Nor}). The investigation of correlations between flare energies and waiting times is a useful tool to distinguish among different triggering mechanisms. This problem has been addressed in a series of papers (Crosby et al \\cite{Cro}; Wheatland et al \\cite{Whe1}; Wheatland \\cite{Whe2}), providing no clear evidence for time-energy correlations. In particular, Wheatland et al. (\\cite{Whe1}) detected small correlations at short time and interpreted them as a spurious effect of obscuration, i.e. long-lasting events may hide subsequent small events. By properly taking into account obscuration effects, Wheatland concluded that no significant correlations are present. In this paper we present a novel analysis of experimental catalogs providing evidence for time-energy correlations in flare occurrence. ", "conclusions": "In conclusion, we presented a statistical analysis of the GOES catalog indicating the existence of time-energy correlations between successive events not to be attributed to obscuration effects. More precisely, we observed that for couples of events close in time ($T<1$h), the second event tends to have a high energy. Moreover couples of events distant in time between 1h and 10h have the first event with a high energy. The analysis of the rate decay after large flares shows evidence that the largest number of events is detected about 4h after the occurrence of the main event. Finally the distribution of flare energies confirms that the higher the flare energy, the larger the number of subsequent events with high energy. The existence of time-energy correlations suggests the possibility of scaling laws relating time with the energy released in a flare. This is a still open question, which could provide interesting insights in the energy storage and release mechanisms at the origin of solar flare occurrence." }, "1003/1003.0414_arXiv.txt": { "abstract": "River streamflows are excellent climatic indicators since they integrate \\textbf{precipitation} over large areas. Here we follow up on our previous study of the influence of solar activity on the flow of the Paran\\'a River, in South America. We find that the unusual minimum of solar activity in recent years have a correlation on very low levels in the Paran\\'a's flow, and we report historical evidence of low water levels during the Little Ice Age. We also study data for the streamflow of three other rivers (Colorado, San Juan and Atuel), and snow levels in the Andes. We obtained that, after eliminating the secular trends and smoothing out the solar cycle, there is a strong positive correlation between the residuals of both the Sunspot Number and the streamflows, as we obtained for the Paran\\'a. Both results put together imply that higher solar activity corresponds to larger \\textbf{precipitation}, both in summer and in wintertime, not only in the large basin of the Paran\\'a, but also in the Andean region north of the limit with Patagonia. ", "introduction": "Usually, studies focusing on the influence of solar activity on climate have concentrated on Northern Hemisphere temperature or sea surface temperature. However, climate is a very complex system, involving many other important variables. Recently, several studies have focused in a different aspect of climate: atmospheric moisture and related quantities like, for example, \\textbf{precipitation}. Perhaps the most studied example is the Asian monsoon, where correlations between solar activity and \\textbf{precipitation} have been found in several time scales. For example, \\cite{2001Natur.411..290N} found strong coherence between solar variability and the monsoon in Oman between 9 and 6 kyr ago. \\cite{2002E&PSL.198..521A} found that Indian monsoon intensity followed the solar irradiance variability on centennial time scales during the last millennium. \\cite{2003Sci...300.1737F} studied Holocene forcing of the Indian monsoon, and found that intervals of weak (strong) solar activity correlates with periods of low (high) monsoon \\textbf{precipitation}. On shorter time scales, \\cite{1997GeoRL..24..159M}, found that, at decadal-multi\\-de\\-ca\\-dal time scales, the correlation between the El Ni\\~no 3 index and the monsoon rainfall is stronger when solar irradiance is above normal and {\\bf{viceversa}}. Correlations between solar activity and Indian monsoon in decadal time scales were also found by \\cite{2005GeoRL..3205813B} and \\cite{2004GeoRL..3124209K}, among others. \\cite{2005Sci...308..854W} studied the monsoon in southern China over the past 9000 years, and found that higher solar irradiance corresponds to stronger monsoon. They proposed that the monsoon responds almost immediately to solar changes by rapid atmospheric responses to solar forcing. All these studies reported a positive correlation, where periods of higher solar activity correspond to periods of larger \\textbf{precipitation}. In contrast, \\cite{2001E&PSL.185..111H} studied a 6000-year record of drought and \\textbf{precipitation} in northeastern China, and found that most of the dry periods agree well with stronger solar activity and {\\bf{viceversa}}. In the American continent, droughts in the Yucatan Peninsula have been associated with periods of high solar activity and have even been proposed to explain the Mayan decline \\citep{2001E&PSL.192..109H}. In the same sense, studies based on the water level of Lakes Naivasha \\citep{2000Natur.403..410V} and Victoria \\citep{stager05} in East Africa, report severe droughts during phases of high solar activity and increased \\textbf{precipitation} during periods of low solar irradiation. To explain these differences it has been proposed that increased solar irradiation causes more evaporation in equatorial regions, enhancing the net transport of moisture flux to the Indian sub-continent via monsoon winds \\citep{2002E&PSL.198..521A}. However, these relationships seem to have reversed sign around 200 years ago, as severe droughts developed over much of tropical Africa during the Dalton sunspot minimum, ca. AD 1800-1820 \\citep{stager05}. Furthermore, \\cite{2007JGRD..11215106S} studied recent water levels in Lake Victoria, and found that peaks in the $\\sim$11-year sunspot cycle were accompanied by water level maxima throughout the 20th century, due to the occurrence of positive rainfall anomalies $\\sim$1 year before solar maxima. Similar patterns also occurred in at least five other East African lakes, indicating that these sunspot-rainfall relationships were broadly regional in scale. A different approach was taken by \\cite{2005MmSAI..76.1002M} who proposed to study the streamflow of a large river, the Paran\\'a in southern South America, as an indicator of \\textbf{precipitation}. In fact, flows of continental-scale rivers are excellent climatic indicators since they integrate \\textbf{precipitation}, infiltrations and evapotranspiration over large areas and smooth out local variations. Signals of solar activity have recently been found with spectral analysis techniques in the river Nile by \\cite{Ruzmaikin2006}, who found a low-frequency 88-year variation present in solar variability and in the Nile records. Similarly, \\cite{2008JGRD..11312102Z} found that the \\textbf{discharge} of the Po river appear to be correlated with variations in solar activity, on decadal time scales. In \\cite{2008PhRvL.101p8501M} (hereinafter Paper I) we presented the results of our study of the Paran\\'a. We found that the streamflow variability of the Paran\\'a river has three temporal components: on the secular scale, it is probably part of the global climatic change, which at least in this region of the world is related with more humid conditions; on the multidecadal time scale, we found a strong correlation with solar activity, as expressed by the Sunspot Number, and therefore probably with solar irradiance, with higher activity coincident with larger discharges; on the yearly time-scale, the dominant correlation is with El Ni\\~no. In the present paper we follow up on the study of the influence of solar activity on the flow of South American rivers. In Section 2 we expand in time the study of the multi-decadal component of the Paran\\'a's streamflow, to include the most recent years, which have shown particularly low levels of solar activity. In Section 3 we study other South American rivers, to see whether the influence extends to other areas of the continent. Finally, in Section 4 we discuss the implications of our findings. ", "conclusions": "In this paper we analyzed the influence of solar activity in the streamflow of South American rivers of different regimes. First, we extended in time the study of the correlation between Sunspot Number and the Paran\\'a's streamflow we reported in Paper I. On one hand, we found that the unusual minimum of solar activity in recent years have a correlation on very low levels in the Paran\\'a's flow. On the other, we reported historical evidence of low water levels during the Little Ice Age. We also found that the correlation is stronger with sunspot number than with neutron count, which confirms that what is affecting climate is most probably solar irradiance, and not GCRs. The fact that the river's behaviour follows $S_N$ through one more minimum strongly enhances the significance of the correlation and its predictive value. In particular, the low levels of activity expected for Solar Cycle 24 anticipate that the dry period in the Paran\\'a will continue well into the next decade. To study whether the solar influence extends to other areas of the continent, we analyzed the streamflow of three South American rivers: the Colorado and two of its tributaries, the San Juan and Atuel rivers. We also used snow level from a station at the origin of the Colorado. We obtained that, after eliminating the secular trends and smoothing out the solar cycle, there is a strong correlation between the residuals of both the Sunspot Number and the streamflows. In all cases, the correlation we found on multi-decadal time scales is positive, i.e., higher solar activity corresponds to larger snow accumulation and, therefore, to larger discharges of all these rivers, as we obtained for the Paran\\'a river. Therefore, both results put together imply that higher solar activity corresponds to larger \\textbf{precipitation}, not only in the large basin of the Paran\\'a, but also in the Andean region north of the limit with Patagonia. Furthermore, since streamflow variability of rivers on central Chile are controlled by the same mechanisms that regulate the rivers studied in this paper, one might expect the same correlation to be found west of the Andes. Solar activity can affect \\textbf{precipitation} through the position of the Inter Tropical Convergence Zone (ITCZ), which has been shown to correlate with variations in solar insolation \\citep{2004GeoRL..3112214P,2001Sci...293.1304H}. In fact, it has been proposed that a displacement southwards of the ITCZ would increase \\textbf{precipitation} in southern tropical South America \\citep{2006GeoRL..3319710N}. We point out that increased \\textbf{precipitation} occur both in the Southern Hemisphere's summer when the ITCZ is over the equator, close to the origin of the Paran\\'a, and in wintertime, when the ITCZ displaces north, and \\textbf{precipitation} increase further South. \\ack{We thank the Subsecretar\\'\\i a de Recursos H\\'\\i dricos de la Naci\\'on (Argentina), and particularly Lic. Daniel Cielak, for facilitating the data used in this study.}" }, "1003/1003.2557_arXiv.txt": { "abstract": "We discuss a new one-dimensional non-LTE time-dependent radiative-transfer technique for the simulation of supernova (SN) spectra and light curves. Starting from a hydrodynamical input characterizing the homologously-expanding ejecta at a chosen post-explosion time, we model the evolution of the {\\it entire} ejecta, including gas and radiation. The boundary constraints for this time-, frequency-, space-, and angle-dependent problem are the adopted initial ejecta, a zero-flux inner boundary and a free-streaming outer boundary. This relaxes the often unsuitable assumption of a diffusive inner boundary, but will also allow for a smooth transition from photospheric to nebular conditions. Non-LTE, which holds in all regions at and above the photosphere, is accounted for. The effects of line blanketing on the radiation field are explicitly included, using complex model atoms and solving for all ion level populations appearing in the statistical-equilibrium equations. Here, we present results for SN1987A, evolving the model ``lm18a7Ad'' of Woosley from 0.27 to 20.8\\,d. The fastest evolution occurs prior to day 1, with a spectral energy distribution peaking in the range $\\sim$300-2000\\AA, subject to line blanketing from highly ionized metal and CNO species. After day 1, our synthetic multi-band light curve and spectra reproduce the observations to within 10-20\\% in flux in the optical, with a greater mismatch for the faint UV flux. We do not encounter any of the former discrepancies associated with the He\\,{\\sc i} and H\\,{\\sc i} lines in the optical, which can be fitted well with a standard Blue-supergiant-star surface composition and no contribution from radioactive decay. The effects of time dependence on the ionization structure, discussed in Dessart \\& Hillier, are recovered, and thus nicely integrated in this new scheme. Despite the 1D nature of our approach, its high physical consistency and accuracy will allow reliable inferences to be made on explosion properties and pre-SN star evolution. ", "introduction": "Core-collapse supernovae (SNe) are extraordinary events situated at the crossroads of many fields of astrophysics. They mark the birth of compact objects, either neutron stars or stellar-mass black holes which, owing to their high compactness, are often the site of tremendous magnetic fields, as in magnetars, or the site of tremendous rotation rates, as in millisecond-period pulsars. Their ejecta make a significant contribution to the chemical enrichment of galaxies, while affecting their dynamics and energetics. The connection to $\\gamma$-ray bursts, for a subset of these SNe, may turn them into excellent probes of the early Universe. Our focus, however, is to characterize the SN ejecta itself and extract information that can help us understand the properties of the explosion and of the progenitor star. Such inferences are based on the analysis of the SN light, using photometric, spectroscopic, or spectropolarimetric data. Hence, developing accurate radiative-transfer tools capturing the key physics controlling the interaction of light and matter is of prime importance. Of all SNe, only core-collapse, and in particular Type II SNe, have well-identified progenitors. Numerous Type II-Plateau (II-P) SNe have now been associated with the explosion of low-mass Red-Supergiant (RSG) stars \\citep{smartt_09}, while the progenitor of the Type II-peculiar SN1987A was a Blue-Supergiant (BSG) star, named Sk -69 202 (for a review, see, e.g., \\citealt{ABK89_rev}). The excellent quality of the observational data for SN1987A makes it ideal for detailed modeling. Numerous radiative-transfer studies of the early-time spectra of SN1987A have been done, with some success \\citep{EK89_87A,hoeflich_87,lucy_87,hoeflich_88, lucy_87,SAR90_87A,mazzali_etal_92}. However, these studies are now twenty years old. They generally assumed Local Thermodynamic Equilibrium (LTE; at best treating only a fraction of the species/levels in non-LTE) and/or steady-state for the radiative transfer, and employed small model atoms. \\citet{lucy_87} and \\citet{SAR90_87A} have emphasized the effects of line blanketing in the UV. One explicit question raised by these studies was the problematic observation of He\\,{\\sc i} lines in optical spectra, which required unacceptable helium enrichments. More recently, \\citet{MBB01_mixing,MBB02_87A} argued for significant mixing of $^{56}$Ni beyond 5000\\,\\kms\\ in order to reproduce Balmer line profiles during the first weeks after explosion. The Balmer line strength problem has now been associated with the erroneous neglect of important time-dependent terms that appear in the energy and statistical-equilibrium equations (\\citealt{UC05_time_dep}; \\citealt{DH08_time}), while reproducing the characteristics of He\\,{\\sc i} lines seems to require a non-LTE treatment. Hence, some of these discrepancies may simply reflect the shortcomings of the radiative-transfer tools employed rather than a genuine peculiarity of the progenitor. Taking a new look at this dataset therefore seems warranted. Furthermore, studying SN1987A is a good exercise to gauge the level of accuracy of radiative-transfer codes and assumptions, facilitated by the considerable advances in computer technology over the last twenty years. Because atomic data represent a fundamental and essential ingredient of radiative-transfer computations, the considerable improvements in that domain over that period make such calculations more accurate than calculations done when SN1987A went off. In this paper, we present a new approach for non-LTE time-dependent radiative transfer modeling of SN ejecta using the code {\\sc cmfgen} \\citep{HM98_lb,DH05_qs_SN, DH08_time}. We discuss the conceptual aspects of the method and emphasize the key equations that are solved, delaying a more comprehensive presentation of the technical details to a forthcoming paper (Hillier \\& Dessart 2010, in preparation). We illustrate this new capability with results obtained for SN1987A, starting our time evolution from an hydrodynamical input of an exploded BSG star (Woosley, priv. comm.; model ``lm18a7Ad\"). We do not present an in-depth study of SN1987A, but merely use this well-observed SN to check and confront our model results. A preliminary set of results were presented in \\citet{DH09_boulder}. In the following section, we discuss the various approaches we have used in the recent years for SN spectroscopic modeling, emphasizing their merits and limitations. We then describe in \\S\\ref{sect_setup} the hydrodynamical model we employ as a basis for this calculation, summarizing its various properties, before presenting our setup for the radiative-transfer calculation. In \\S\\ref{sect_results}, we describe the ejecta evolution, such as temperature and ionization structure, and present checks on our numerical technique. We then discuss in \\S\\ref{sect_synthetic_rad} the radiative properties of our non-LTE time-dependent models, covering in turn synthetic spectra and the bolometric light curve, and detailing in particular the sources of line blanketing at various epochs. In \\S\\ref{sect_comp_to_obs}, we compare our theoretical predictions to UV and optical observations of SN1987A. Finally we present our conclusions in \\S\\ref{sect_ccl}, and lay out the various projects ahead. ", "conclusions": "\\label{sect_ccl} We have presented a new approach for radiative-transfer modeling of SN ejecta, retaining the key assets of our former approaches. By treating non-LTE explicitly and incorporating time-dependent terms in the statistical-equilibrium, energy, and moments of the radiative transfer equations, we improve the physical consistency of our computations. In particular, the simulations are performed on the entire ejecta, starting from a given hydrodynamical input at a given post-explosion time, making this an initial-value problem, while spatial boundaries are zero-flux at the inner edge and free-streaming at the outer edge. As diffusion at the inner boundary does not need to hold we can evolve the ejecta from its photospheric to its nebular phase. Importantly, integrating our synthetic spectra over filter-transmission functions, we compute non-LTE light-curves that account for the explicit role of line-blanketing, without recourse to the standard expansion opacities and the assumption of LTE for the gas, as is typically used (see, e.g., \\citealt{blinnikov_etal_2000}). Delaying a presentation of all the technical details of this new approach (Hillier \\& Dessart 2010, in preparation), we illustrate in this paper these modeling improvements with results for SN1987A, using as our starting conditions the ejecta composition and structure of the hydrodynamical input model ``lm18a7Ad'' of Woosley (priv. comm.). We evolved the ejecta gas and radiation from 0.27 to 20.8\\,d after explosion, corresponding to a spatial expansion of a factor of 77, and considerable cooling. At the photosphere, the ejecta conditions evolve from fully-ionized to once-ionized and neutral. During the 21 day time sequence, the photosphere recedes in mass, and thus in velocity, while its radius increases steadily. After a few days, the photospheric location is primarily set by the location of the hydrogen ionization front. Our computed temperature and electron-density evolution reproduces predictions for a radiation-dominated homologously-expanding gas, modulo the effect of cooling at the photosphere and heating from radioactive decay in the inner ejecta. At the last time in the simulation, the photosphere is at 4500\\,\\kms, and is not influenced, either directly or indirectly, by energy deposition from radioactive decay occurring below 2000\\,\\kms\\ in our model. The surface conditions computed here and prior to 20\\,d are thus not affected by unstable nuclei. Since the SN1987A model ejecta we employ are homogeneous above 4500\\,\\kms, the changes we observe in the spectrum over this 21-day time span are conditioned by modulations in ionization rather than composition. Adopting a distance of 50\\,kpc and a reddening $E(B-V)=0.15$, we obtain very good agreement between synthetic and observed light curves in the optical. Our $U$-band synthetic magnitudes are, however, too bright by a fraction of a magnitude at early times, and this discrepancy slowly grows as the flux in the blue ebbs. Our synthetic spectra computed for the epoch 0.3--1\\,d cannot be compared to observations, which started one day after explosion. At such early times, we find an SED with a peak flux in the far-UV, where strong line blanketing occurs due to 3-4-times ionized metal species (O, Fe, Ni), while the nearly-featureless optical range shows weak lines associated with H\\,{\\sc i}, He\\,{\\sc ii} and He\\,{\\sc i}. After day 1 and until the 14th of March 1987, we compare our predictions to observations. In general, the flux agreement in the optical is at the 10\\% level (in an absolute sense, i.e., without any normalization), while in the UV range it is within 30-50\\%. In contrast with \\citet{EK89_87A}, we reproduce the strength and morphology of He\\,{\\sc i} lines at early times, under a photospheric composition that is compatible with a BSG progenitor star. Throughout the time span of our sequence, we reproduce well the multi-band light curve, the overall SED, and in particular the strength and shape of H$\\alpha$. Since radioactive decay does not influence the computed photospheric conditions in any way at such times, and given the relatively good fit to observations (in particular if we focus on the bulk of the emergent luminosity, which falls at optical wavelengths), it does not play a pivotal role for understanding the radiative properties of SN1987A up to 20\\,d, in contradiction with the proposition of \\citet{MBB01_mixing,MBB02_87A}. In contrast, our good reproduction of the Balmer line profiles at all times, and in particular when there are no longer any photons in the Lyman and Balmer continua, gives strong support to our proper modeling of the time-dependent ionization structure of the SN1987A ejecta, an important issue emphasized by \\citet{UC05_time_dep} and \\citet{DH08_time}. \\citet{de_etal_09} suggest that such time-dependent effects are much weaker, but their presentation for the recombination phase of SN1987A is unconvincing since they use models that are fully-ionized and UV bright (in contradiction with observations). Furthermore, despite the one-dimensionality of our approach, we achieve very competitive fits to observations. One may argue the unresolved problems we encounter are caused by those neglected multi-dimensional effects, but we surmise this departure from sphericity would also alter line profile shapes etc., which are fairly well reproduced (see, e.g., H$\\alpha$). We predict the fastest spectral evolution in SN1987A occurred prior to day 1. From 0.3\\,d to 1\\,d, most of the flux was emitted in the range 300-3000\\AA\\ which could have been captured by the IUE satellite. Prior to 0.27\\,d, the photosphere would have shone at shorter wavelengths, but the UV range would have revealed the long-wavelength tail of that SED, testifying for the hotter conditions at earlier post-breakout times. This suggests that obtaining multiple observations prior to day 1, across as broad a wavelength range as possible, would have captured the phase of fastest evolution of SN1987A. While the UV shows the largest changes, multiple spectra, rather than a unique spectrum, of SN1987A during both the first and the second night of its evolution would have been valuable.\\footnote{Covering the full optical range, from 3500\\AA\\ to 1$\\mu$m allows us to gather information on Ca and CNO lines in the red part of the spectrum, as well as constrain the reddening with the blue part.} This applies to some extent to Type II-P SN as well, although their evolution is slower. In Type I SNe, the SED likely evolves too fast after breakout to allow us to capture it, but then, the first few days after explosion would be very useful for constraining the surface composition of the progenitor. Overall, observing frequently at early times would provide important constraints that are generally missing in current observations. The work presented here is an essential benchmarking of our code since SN1987A is one of the best understood and best observed supernova. As it places numerous and tight observational constraints, it offers {\\it a valuable alternative, although not a replacement,} to benchmarking against other codes. The combined approach of using non-LTE, time dependence, and line blanketing seems very promising. The future from here is to systematize such investigations by evolving a wide range of ejecta for all supernova types and compare results with multi-epoch multi-wavelength observations. The aim is then to use such models and observations to make quantitative inferences on the progenitor properties, pre-SN star evolution, and the explosion mechanism. In the context of SN1987A, we now need to repeat the present exploration with a variety of progenitor stars and explosion properties, in order to delineate the systematics associated with core-collapse SN explosions of BSG stars. In parallel, it would be valuable to gather additional and high quality observational data for similar Type II-peculiar events like SN1987A." }, "1003/1003.4883_arXiv.txt": { "abstract": "{}{A model of jet precession driven by a neutrino-cooled disc around a spinning black hole is present in order to explain the temporal structure and spectral evolution of gamma-ray bursts (GRBs).}{The differential rotation of the outer part of a neutrino dominated accretion disc may result in precession of the inner part of the disc and the central black hole, hence drives a precessed jet via neutrino annihilation around the inner part of the disc.}{Both analytic and numeric results for our model are present. Our calculations show that a black hole-accretion disk system with black hole mass $M \\simeq 3.66 M_\\odot$, accretion rate $\\dot{M} \\simeq 0.54 M_\\odot \\rm s^{-1}$, spin parameter $a=0.9$ and viscosity parameter $\\alpha=0.01$ may drive a precessed jet with period $P=1$ s and luminosity $L=10^{51}$ erg s$^{-1}$, corresponding to the scenario for long GRBs. A precessed jet with $P=0.1$s and $L=10^{50}$ erg s$^{-1}$ may be powered by a system with $M \\simeq 5.59 M_\\odot$, $\\dot{M} \\simeq 0.74 M_\\odot \\rm s^{-1}$, $a=0.1$, and $\\alpha=0.01$, possibly being responsible for the short GRBs. Both the temporal and spectral evolution in GRB pulse may explained with our model.}{GRB central engines likely power a precessed jet driven by a neutrino-cooled disc. The global GRB lightcurves thus could be modulated by the jet precession during the accretion timescale of the GRB central engine. Both the temporal and spectral evolution in GRB pulse may be due to an viewing effect due to the jet precession.} ", "introduction": "Internal shock models are extensively discussed for gamma-ray bursts (GRBs) (Rees \\& M\\'{e}sz\\'{a}ros 1992; M\\'{e}sz\\'{a}ros \\& Rees 1993; Zhang \\& M\\'{e}sz\\'{a}ros 2004), in which an individual shock episode of two collision shells gives rise to a pulse, and random superposition of pulses results in the observed complexity of GRB light curves (e.g., Daigne \\& Mochkovitch 1998; Kobayashi et al. 1999). The observed flux rapidly increases in the dynamic timescale of two shell collision, then decays due to the delayed photons from high latitudes with respect to the line of sight upon the abrupt cessation of emission after the crossing timescale, shaping the observed fast-rise-exponential-decay (FRED) pulses. However, some well-separated GRB pulses show symmetric structure, and their peak energy of the $\\nu F_\\nu$ spectrum ($E_p$) traces the lightcurve behavior (Liang \\& Kargatis 1996; Liang \\& Nishimura 2004; Lu \\& Liang 2009; Peng et al. 2009). Both the temporal and spectral properties of these symmetric pulse are difficult to be explained with internal shocks. In addition, the observed $E_{\\rm iso}-E_p$ relation (Amati et al. 2002) or $L_{\\rm iso}-E_p$ relation (Wei \\& Gao 2003; Liang et al. 2004; Yonetoku et al. 2004) also challenge the internal shock models (e.g., Zhang \\& M{\\'e}sz{\\'a}ros 2002). Quasi-periodic feature observed in some GRB light curves motivated ideas that the GRB jet may be precessed (Blackman et al. 1996; Portegies Zwart et al. 1999; Portegies Zwart \\& Totani 2001; Reynoso et al. 2006; Lei et al. 2007). It is generally believed that the progenitors of short and long GRBs are the mergers of two compact objects (Eichler et al. 1989; Paczy\\'{n}ski 1991; Narayan et al. 1992; see recent review by Nakar 2007) and core collapsars of massive stars (Woosley 1993; Paczy\\'{n}ski 1998; see reviews by Woosley \\& Bloom 2006), respectively. Although the progenitors of the two types of GRBs are different, the models for their central engines are similar, and essentially all can be simply classed as a rotating black hole with a rapidly hyper-accreting process of a debris torus surrounding the central black hole. Such a black hole-disk system drives an ultra-relativistic outflow to produce both the prompt gamma-rays and afterglows in lower energy bands. The most popular one is neutrino dominated accretion flows (NDAFs), involving a black hole of $2\\sim10 M_\\odot$ and a hyper-critical rate in the range of $0.01 \\sim 10 M_\\odot {\\rm s}^{-1}$ (Popham et al. 1999; Narayan et al. 2001; Kohri \\& Mineshige 2002; Di Matteo et al. 2002; Kohri et al. 2005, 2007; Lee et al. 2005; Gu et al. 2006; Chen \\& Beloborodov 2007; Liu et al. 2007, 2008, 2010; Kawanaka \\& Mineshige 2007; Janiuk et al. 2007). The different direction of angular momentum of two compact objects and the anisotropic fall-back mass in collapsar may conduct precession between black hole and disc. In this scenario, the inner part of the disc is driven by the black hole during the accretion process. The differential rotation between the inner and outer parts may result in precession of the inner part of the disc and the central black hole, hence drive a precessed jet produced by neutrino annihilation around the inner part of the disc, forming an S- or Z-shaped jet as observed in many extragalactic radio sources (see, e.g. Florido et al. 1990). A tilted accretion disc surrounding a black hole would also make the precession of the black hole and result in an S-shaped jet as observed in SS 433 (Sarazin et al. 1980; Lu 1990; Lu \\& Zhou 2005), although the angle between angular momentum of black hole and disc is small due to that evolution of a two compact object system may decrease the angle between them in mergers or the anisotropic fall-back mass cannot produce large angle between black hole and fall-back mass in collapsars. In this paper, we propose a model of jet precession driven by a neutrino-cooled disc around a spinning black hole in order to explain the temporal structure and spectral evolution of GRBs. In our model, the global profile of a GRB lightcurve may be modulated by the jet procession. The temporal structure and spectral evolution may signal an on-axis/off-axis cycle of the light of sight (LOS) to a precessed jet axis, as proposed by some authors to explain the nature of low luminosity GRBs 980425 and 031203 (Nakamura 1998; Eichler \\& Levinson 1999; Waxman 2004; Ramirez-Ruiz et al. 2005) or to present a unified model for GRBs and X-ray flashes (Yamazaki et al. 2003) and the observed spectral lag in long GRBs (Norris 2002; Salmonson \\& Galama 2002). We present both analytic and numerical analysis for jet precession driven by a neutrino-cooled disc around a spinning black hole in sections 2 and 3. Simplifying the jet emission surface as a point source, we demonstrate the profile and evolution of a GRB pulse in section 4. Conclusions and discussion are shown in section 5. ", "conclusions": "We have suggested that the differential rotation of the outer part of a neutrino dominated accretion disc may result in precession of the central black hole and the inner part of the disc, hence may power a precessed jet via neutrino annihilation around the inner part of the disc. Both analytic and numeric results are present. Our calculations show that for a black hole-accretion disk system with $M \\simeq 3.66 M_\\odot$, $\\dot{M} \\simeq 0.54 M_\\odot \\rm s^{-1}$, $a=0.9$ and $\\alpha=0.01$ may drives a precessed jet with $P=1$ s and $L=10^{51}$ erg s$^{-1}$, corresponding to the scenario for long GRBs. A precessed jet with $P=0.1$s and $L=10^{50}$ erg s$^{-1}$ may be powered by a system with $M \\simeq 5.59 M_\\odot$, $\\dot{M} \\simeq 0.74 M_\\odot \\rm s^{-1}$, $a=0.1$, and $\\alpha=0.01$, possibly being responsible for the short GRBs. These results are generally consistent with simulations for long and short GRB productions from collapsars and from mergers of compact stars. Both temporal and spectral features observed in GRB pulses may be explained with our model. The correlation between $E_{\\rm iso}$ (or $L_{\\rm iso}$) and $E_p$ in the burst frame (Amati et al. 2002; Liang et al. 2004) are difficult to be explained in the framework of internal shock scenarios. Our model suggests an $E_p$-tracing-flux behavior within a GRB pulse due to the on-axis/off-axis effect for a given observer, similar to that proposed by Yamazaki et al. (2004). The $E_p$-tracing-flux behavior would give rise to the observed correlations between $E_{\\rm iso}$ ($L_{\\rm iso}$) and $E_p$ in the burst frame." }, "1003/1003.1142_arXiv.txt": { "abstract": "Most proposed dark matter candidates are stable and are produced thermally in the early Universe. However, there is also the possibility of unstable (but long-lived) dark matter, produced thermally or otherwise. We propose a strategy to distinguish between dark matter annihilation and/or decay in the case that a clear signal is detected in gamma-ray observations of Milky Way dwarf spheroidal galaxies with gamma-ray experiments. The sole measurement of the energy spectrum of an indirect signal would render the discrimination between these cases impossible. We show that by examining the dependence of the intensity and energy spectrum on the angular distribution of the emission, the origin could be identified as decay, annihilation, or both. In addition, once the type of signal is established, we show how these measurements could help to extract information about the dark matter properties, including mass, annihilation cross section, lifetime, dominant annihilation and decay channels, and the presence of substructure. Although an application of the approach presented here would likely be feasible with current experiments only for very optimistic dark matter scenarios, the improved sensitivity of upcoming experiments could enable this technique to be used to study a wider range of dark matter models. \\vspace{3ex} \\noindent {\\bf Keywords:} dark matter theory, gamma-ray theory, dwarf galaxies ", "introduction": "\\label{introduction} Many different astrophysical and cosmological observations have found evidence for the existence of non-luminous, non-baryonic dark matter (for reviews see, e.g., Refs.~\\cite{Jungman:1995df, Bergstrom:2000pn, Bertone:2004pz}) and indicate that it constitutes about 80\\% of the mass content of the Universe~\\cite{Dunkley:2008ie}. Despite the precision with which the cosmological dark matter density has been measured, little is known about the origin and properties of the dark matter particle, such as its mass, spin, couplings, and its distribution on small scales. Several candidate dark matter particles have been proposed with masses from the electroweak scale to superheavy candidates at the Planck scale (see, e.g., Refs.~\\cite{Bertone:2004pz, Bergstrom:2009ib} for a comprehensive list). Light particles have also been considered as possible dark matter candidates: axions~\\cite{axions}, sterile neutrinos with masses in the keV range~\\cite{Dodelson:1993je} and light scalars with MeV--GeV masses~\\cite{LDM, Boehm:2003bt, MeVDM, Boehm:2006mi}. However, the most popular candidates are weakly interacting massive particles (WIMPs) which arise in extensions of the Standard Model such as supersymmetric (e.g., Ref.~\\cite{Jungman:1995df}), little Higgs (e.g., Ref.~\\cite{Birkedal:2006fz}), or extra-dimensions models (e.g., Ref.~\\cite{Hooper:2007qk}). Most proposed WIMPs are stable and are produced thermally in the early Universe with an annihilation cross section (times relative velocity) of $\\langle\\sigma v\\rangle \\sim 3 \\times 10^{-26}$~cm$^3$~s$^{-1}$. However, dark matter may be unstable but long-lived; the only requirement in order for it to be a thermal relic and present today is that it has a lifetime $\\tau_\\chi$ much longer than the age of the Universe $t_{\\rm U} \\simeq 4 \\times 10^{17}$~s. It is also possible that dark matter is not a thermal relic, which would allow it to have a larger annihilation cross section than the canonical value for WIMP thermal relics. Although the case of a non-thermal~\\cite{Kaplinghat:2000vt, unitarity} or unstable~\\cite{decaymodels} dark matter candidate was considered several decades ago, recently a great deal of interest in these scenarios has been generated by the rise in the positron fraction in the tens of GeV range measured by the PAMELA experiment~\\cite{PAMELA}. One possibility to explain the PAMELA data is by the injection of positrons by annihilation~\\cite{Cirelli:2008pk, Barger:2008su, PAMELAannihilationexotic, Grasso:2009ma}, decay~\\cite{PAMELAdecay, Nardi:2008ix} or both~\\cite{Cheung:2009si} of dark matter in the solar neighborhood (see also Ref.~\\cite{He:2009ra} and references therein). In the case of annihilation in the smooth halo, enhancements to the annihilation cross-section of the order of $10$-$10^5$ are required (with respect to thermal dark matter with no Sommerfeld or Breit-Wigner enhancements and assuming there is no nearby dark matter clump)~\\cite{Cirelli:2008pk}. However, note that the contribution from substructure could be significant~\\cite{Brun:2009aj} and interestingly, the cumulative effect from distant subhalos could modify the electron and positron spectra at our galactic position~\\cite{Cline:2010ag}. Indirect dark matter searches look for the products of dark matter annihilation or decay, which include not only antimatter, as in the case of PAMELA, but also neutrinos~\\cite{Rott:2009hr, indirectnusun} and photons. In particular, targets for indirect searches in gamma-rays include dark matter in extragalactic structures~\\cite{cosmoindirect}, the Galactic Center~\\cite{GC, Albert:2005kh, Aharonian:2006wh, Baltz:2008wd}, the Milky Way halo~\\cite{Baltz:2008wd, MWhalo, Pieri:2009je, Abdo:2010nc}, its subhalos~\\cite{Baltz:2008wd, Pieri:2009je, subhalos, Kuhlen:2008aw} and known dwarf galaxies~\\cite{Baltz:2008wd, dG, Evans:2003sc, Strigari:2007at, Essig:2009jx, MAGICdG, HESSdG, Driscoll:2008zz, Wagner:2009wp, Abdo:2010ex}. Different approaches have been proposed to constrain dark matter properties by using indirect or direct measurements or their combination~\\cite{direct, indirectnuprop, indirectgamma, Hensley:2009gh, Beltran:2008xg}. To extract the properties of the dark matter particle from the detection of an indirect signal requires several pieces of information. While the energy spectrum of the signal is dependent on the dark matter properties (mass $m_{\\chi}$, annihilation cross section $\\langle\\sigma v\\rangle$, lifetime $\\tau_{\\chi}$, and annihilation and/or decay channels), sufficient degeneracies exist to prevent accurate reconstruction of the dark matter properties from the energy spectrum alone. In particular, the sole measurement of the energy spectrum would make it impossible to know if the indirect signal from dark matter is produced by annihilation or decay. The spectrum of the former is characterized by a cutoff at an energy equal to the dark matter mass, while the cutoff in the spectrum from the latter is at an energy equal to half of the dark matter mass. In this work we address the following two questions: (1) in the case that an indirect dark matter signal is detected in gamma rays, can annihilation and/or decay be identified as the origin of the signal? and (2) what information about the particle properties can be obtained from the indirect measurement? If dark matter is unstable and produces an observable signal from decay, an annihilation signal may also be present. In principle, there is a range of parameters for which the two signals would be comparable, which would present a challenge for distinguishing the cases of annihilation, decay, and those where both annihilation and decay are significant contributors to the measured signal. In order to tackle this problem, we propose a strategy to distinguish between these scenarios using gamma-ray observations of Milky Way dwarf spheroidal galaxies with current or future gamma-ray telescopes. Current missions include the Large Area Telescope (LAT) aboard the Fermi Gamma-ray Space Telescope (\\emph{Fermi})~\\cite{Atwood:2009ez}, which observes gamma-rays in the range from 20~MeV to greater than 300~GeV, and atmospheric \\v{C}erenkov telescopes (ACTs) such as HESS~\\cite{Hinton:2004eu}, VERITAS~\\cite{Holder:2008ux}, and MAGIC~\\cite{Lorenz:2004ah}, which observe emission above $\\sim$100~GeV, and the planned ground-based ACT arrays CTA~\\cite{Wagner:2009cs} and AGIS~\\cite{Buckley:2009zza}. We demonstrate that, in the case that a gamma-ray signal is clearly detected, the origin could be identified as decay, annihilation, or both by examining the dependence of the intensity and energy spectrum on the angular distribution of the emission. Furthermore, if annihilation and decay each contribute significantly to the signal, we show how these observations could be used to extract information about the dark matter mass, lifetime, annihilation cross section, and dominant annihilation and decay channels. In addition, as a byproduct of this analysis, one might also establish or limit the contribution to the signal from substructure in the dwarf galaxy's halo. The paper is organized as follows. We outline our proposed observing strategy in \\S\\ref{method}. In \\S\\ref{dwarfs} we summarize the properties of the dwarf galaxies we consider and describe our approach to modeling their dark matter halos and subhalos. The gamma-ray spectra from dark matter annihilation and decay are discussed in \\S\\ref{spectra}. In \\S\\ref{results} we demonstrate the proposed method by presenting example results for various dark matter scenarios for selected dwarf galaxies. We summarize our results and conclude in \\S\\ref{conclusions}. ", "conclusions": "\\label{conclusions} In this work we have outlined a strategy to constrain dark matter properties in the event of the clear detection of an indirect signal from gamma-ray observations of dwarf galaxies. We addressed the question of how scenarios of dark matter annihilation, decay, or both could be distinguished, and what information could be obtained about the intrinsic properties of the dark matter particle and its small-scale distribution from this type of indirect measurement. In principle, the indirect detection of dark matter in gamma-ray experiments would provide the energy spectrum of the signal. The spectrum from dark matter annihilation has a maximum energy equal to the dark matter particle mass, while the spectrum from decay has a cutoff at half the mass. Consequently, spectral information alone is insufficient to identify the process that produced the signal. However, in addition to the difference in the endpoint of the energy spectrum, annihilation and decay give rise to different angular distributions of the intensity of the emission. Using this information, we demonstrated that if a dark matter signal is clearly detected from a dwarf galaxy, an analysis of the energy spectrum of the emission as a function of the angular distance from the center of the object could provide the necessary information to distinguish the cases of annihilation, decay, or both. The technique we propose, which combines spectral and angular information, is particularly important due to the uncertainties in the presence and properties of substructure in a dwarf galaxy halo. In particular, whereas the annihilation rate scales as the square of the dark matter density, the decay rate depends linearly on the density. However, the angular distribution of the annihilation (or decay) signal from dark matter substructure also scales approximately linearly with the smooth halo density in the outer regions of the object, so its angular distribution roughly mimics that from dark matter decay in the smooth halo. Furthermore, since the amplitude of the signal from annihilation in substructure depends sensitively on the properties of the subhalo population and is not fixed by the smooth halo density profile, the relative amplitude of the annihilation signals from substructure and the smooth halo is effectively independent (e.g., Ref.~\\cite{Kuhlen:2008aw}). Thus, identifying annihilation in the smooth halo based on the emission profile in the central regions of the object is not sufficient to distinguish between annihilation in substructure or decay as the origin of emission from the outer regions of the halo. In order to break this degeneracy, we have studied the energy spectrum as a function of the angular distance from the center of the object (see Figs.~\\ref{fig:annuli} and~\\ref{fig:inout}). If a flattening in the radial distribution of the intensity of the signal is observed at an angular distance $\\psi_{\\rm cross}$ (c.f. Eq.~\\ref{eq:taucross}), this would point either to a significant contribution from both dark matter annihilation and decay or to only dark matter annihilation with an important contribution from substructure. We have shown that a change in the energy spectrum along with the change of slope in the angular distribution could provide the necessary information to confirm or reject the presence of a signal from dark matter decay. If dark matter decay is established by an observed spectral change as a function of angle, the signal in the innermost parts could be studied to provide information about the annihilation contribution, and that in the outermost regions (beyond $\\psi_{\\rm cross}$) would provide information about the decay signal. In principle, determining $\\psi_{\\rm cross}$ and the intensity of the signal in the inner (annihilation-dominated) and outer (decay-dominated) regions, could help to constrain $\\langle \\sigma v \\rangle$, $\\tau_\\chi$ and $m_\\chi$. In this case, limits on the properties of substructure could also be placed. On the other hand, if beyond $\\psi_{\\rm cross}$ the energy spectrum does not change, we expect annihilation in substructure to be the cause of the flattening in the radial distribution and attribute the signal to annihilation. In this case, we could determine to some extent the contribution of substructure and begin to constrain its properties, such as the mass function slope, minimum subhalo mass, and structural properties. The indirect measurement would then yield the particle mass and annihilation cross-section, and the absence of a strong decay signal would enable further limits on decaying dark matter to be placed. Finally, if decay (annihilation) in the smooth halo can be established as the dominant contribution due to a very shallow (steep) emission profile at all angles, then the mass and lifetime (annihilation cross-section) could be determined, and upper (lower) bounds placed on the annihilation cross-section (lifetime). In addition, as it is known and we have shown in \\S\\ref{spectra}, to a first approximation the annihilation and decay spectra can be classified as soft (due to hadronic or gauge boson channels) or hard (due to leptonic channels). Hence, once the origin of the signal in the different regions is established (annihilation in the smooth halo, annihilation in the smooth halo and substructure, or decay), further information can be obtained about the annihilation and decay channels by studying the energy spectra. On general grounds, in this situation the distinction between dominantly leptonic or dominantly hadronic and gauge boson channels could be achieved. Our example scenarios focused on the case of a signal from Milky Way dwarf spheroidal galaxies. However we note that, in principle, this method could also be applied to the case of our own galaxy. In that case an appropriate treatment of the secondary photons produced by the prompt electrons and positrons via inverse Compton scattering off the ambient photon background is necessary (see Ref.~\\cite{Boehm:2010qt} for a very recent related study). In the case of our galaxy, this contribution is very important (and it could be the dominant one) due to the large light emission by stars and the infrared light as a result of the scattering, absorption and re-emission of absorbed starlight by dust. In addition, in order to recover any information from an observation, the different backgrounds for a potential signal would have to be properly addressed. The galactic center is a complex region, which makes the modeling of these backgrounds a difficult task. All in all, bearing in mind these differences, a similar application of the methodology described here could be possible in the case of the Milky Way. We have demonstrated the proposed method for different scenarios that could also potentially explain the rise in the positron fraction observed in the PAMELA data~\\cite{PAMELA} (Fig.~\\ref{fig:inoutPAMELA}). In the case that these dark matter models produce a detectable gamma-ray signal from dwarf galaxies, we see that if the dark matter interpretation is correct, the observation of the gamma-ray energy spectrum at different angular distances from the center of the dwarf galaxies could help to establish if the origin of the observed positron excess is due to dark matter annihilation, decay, or both. Finally, in Figs.~\\ref{fig:example1},~\\ref{fig:example2}, and~\\ref{fig:example3} we have shown the range of dark matter parameters for which a transition between annihilation and decay would occur within $2^\\circ$ of the center of Draco (similar results are obtained for Ursa Minor and Sagittarius). We have presented the results for different annihilation-decay combinations and see that the dark matter lifetime must be in the range $\\sim (10^{25}- 10^{31}$~s) $(3 \\times 10^{-26}$~cm$^3$~s$^{-1} / \\langle \\sigma v \\rangle$), the actual (narrower) range depending on the combination of annihilation and decay channels. The effect of substructure in the dwarf galaxy halo is shown to be important only for $\\psi_{\\rm cross} \\gtrsim 1^{\\circ}$. The choice of energy threshold, however, strongly affects the range of $\\tau_\\chi$ and $\\langle \\sigma v \\rangle$ probed by this technique (Fig.~\\ref{fig:example3}), suggesting that, in the event of a detection, observations covering a large energy range may be able to explore a substantial region of the dark matter annihilation and decay parameter space. In order to apply the technique proposed in this study, the firm detection of a gamma-ray signal of dark matter origin would be required. Although the main idea of this work is to show the different potential features of a dark matter signal and to describe the method in a detector-independent way, we think it is worthwhile to add a short discussion along these lines. The main point is to have access to two types of features, those in the angular distribution and those in the energy spectrum. In order to reconstruct the energy spectrum in a given angular bin, a minimum number of photons is necessary. The integrated number of photons in each angular bin scales with the solid angle $\\Omega$ of the bin, i.e., $\\Omega \\sim \\pi (r_{\\rm out}^2 - r_{\\rm in}^2)$, where $r_{\\rm out}$ and $r_{\\rm in}$ are the angular radii of the outer and inner edge of each bin. For the equal width annuli used here, $\\Omega$ increases linearly with the radius of the bin center, i.e., the annulus centered on 0.85$^\\circ$ encompasses 17 times as much solid angle as the annulus centered on 0.05$^\\circ$. Let us assume, for simplicity, that the difference in the total intensity between inner and outer annuli is of the order of 10 to 100 (cf.~Fig.~\\ref{fig:annuli}). The photon flux from within the outer annulus is therefore typically a factor of a few smaller than from the innermost annulus, if the spectral shapes in each annulus are similar. Although the diffuse backgrounds would also scale with solid angle, in the case that the backgrounds are small, sensitivity to a signal a factor of only a few smaller than the signal in the innermost bin would be required to measure the angular dependence of the integrated flux. If the backgrounds are large, sensitivity to signals a factor of 10 to 100 smaller than that in the innermost bin would be necessary. While the spectra of the inner and outer annuli depend on the specific combination of channels, very roughly we can say the energy spectrum goes as $E^{-2}$ over most of the relevant energy range (see, e.g., Fig.~\\ref{fig:annuli}), and hence the integrated photon yield over some range scales approximately as $E^{-1}$. The examples shown in this study use the energy spectrum over 1 to 2 decades in energy, so for equal logarithmic bins in energy, the flux at high energies is $\\sim$ 1 - 10\\% of the flux at low energies. To reconstruct the energy spectrum, one could require, e.g., a few signal photons in the highest energy bin of the outermost annulus. This optimistically implies that $\\mathcal{O}(1000)$ total photons (angle- and energy-integrated) within a 1 degree radius of the source are needed. On general grounds, to get angular information a factor of a few more statistics is necessary than are needed to detect the dark matter signal clearly (i.e., with spectral info) in the first place. Thus, if a signal in current experiments were detected just beyond the current limits, in the future they might have the ability to do a study of this kind, as the signal will improve with exposure. While this coarse estimate does not include a proper treatment of backgrounds, which depend strongly on energy and on the experiment under consideration, we note that the expected value of the photon counts measured by \\emph{Fermi} due to the Galactic diffuse gamma-ray background above 1~GeV in a region of 1~degree around the center of the dwarfs considered here after five years is $\\sim$~100 - 1000. In summary, we have shown that a dark matter particle with an annihilation cross-section and lifetime just beyond the limits currently established could produce a clear spectral change on an angular scale within the reach of future experiments. Ongoing observations by \\emph{Fermi}, HESS, VERITAS, and MAGIC and future observations by the planned CTA and AGIS experiments will continue to improve the prospects for detecting and mapping a dark matter signal in the coming years. \\ack We thank Joakim Edsj\\\"o and Torbj\\\"orn Sojstrand for help with PYTHIA and Jorge Pe\\~narrubia for useful communications. SPR thanks the CCAPP, where this work started, and JSG thanks the CFTP for hospitality. SPR is partially supported by the Portuguese FCT through CERN/FP/83503/2008, CERN/FP/109305/2009 and CFTP-FCT UNIT 777, which are partially funded through POCTI (FEDER), and by the Spanish Grant FPA2008-02878 of the MICINN\\@. JSG acknowledges partial support from NSF CAREER Grant PHY-0547102 (to John Beacom)." }, "1003/1003.1374_arXiv.txt": { "abstract": "{Strong inflows of gas from the outer disk to the inner kiloparsecs are induced during the interaction of disk galaxies. This inflow of relatively low-metallicity gas dilutes the metallicity of the circumnuclear gas. This process is critical for the galaxy evolutions. We have investigated several aspects of the process as the timing and duration of the dilution and its correlation with the induced star formation. We analysed major (1:1) gas-rich interactions and mergers, spanning a range of initial orbital characteristics. Star formation and metal enrichment from SNe are included in our model. Our results show that the strongest trend is between the star formation rate and the dilution of the metals in the nuclear region; i.e., the more intense the central burst of star formation, the more the gas is diluted. This trend comes from strong inflows of relatively metal-poor gas from the outer regions of both disks, which fuels the intense star formation and lowers the overall metallicity for a time. The strong inflows happen on timescales of about 10$^8$ years or less (i.e., on an internal dynamical time of the disk in the simulations), and the most intense star formation and lowest gas phase metallicities are seen generally after the first pericentre passage. As the star formation proceeds and the merger advances, the dilution reduces and enrichment becomes dominant -- ultimately increasing the metallicity of the circumnuclear gas to a level higher than the initial metallicities of the merging galaxies. The ``fly-bys'' -- pairs that interact but do not merge -- also cause some dilution. We even see some dilution early in the merger or in the ``fly-bys'' and thus do not observe a strong trend between the nuclear metallicities and separation in our simulations until the merger is well advanced. We also analyse the O and Fe enrichment of the ISM, and show that the evolution of the $\\alpha/Fe$ ratios, as well as the dilution of the central gas metallicity, can be used as a clock for ``dating'' the interaction.} ", "introduction": "Understanding how galaxy interactions and mergers shape the ensemble population of galaxies in the local universe is crucial for our understanding of galaxy evolution. It has been recognised for several decades that mergers may play an important role in the evolution of galaxies, especially for the early types \\citep[e.g.,][]{toomre77}. While there have been many observations and theoretical studies of ``first-order'' effects in interactions and mergers (such as morphological and kinetical evolutiona and star formation), much less work has been done on 2$^{nd}$ order effects such as the evolution of metallicity. Recently, we have undertaken the study of the evolution of metallicity in major dry (i.e., gas poor) mergers, showing that they can lead to metallicity gradients in agreement with those of local ellipticals \\citep{dimatteo09}. However, in gas-poor mergers, the evolution of the local galaxy metallicity is driven only by the mixing of stars during the coalescence process. Gas can alter the properties of the final remnant significantly, affecting both the final gas phase and stellar metal abundance and ratios. In particular, gas phase metallicity can play an important role in constraining our understanding of the complex interaction between gas flows and star formation in interactions and mergers and allow us to constrain the relative time scales of the different epochs (i.e., date the merger/interaction). If the progenitor disks have strong gas metallicity gradients, as observed in many galaxies in the local universe \\citep{shields90, diner96}, one would expect that interaction-induced gas inflows will drive a noticeable amount of gas into the central regions from the outer disk and thus initially lower the circumnuclear gas phase metallicity. Outflows of gas ejected in tidal tails may contribute to modifying the metallicity profile of the final merger remnant. Subsequent star formation (and outflows) may increase (decrease) the gas phase circumnuclear metallicity. The exact timing of the inflows versus the initiation of intense star formation and subsequent metal enrichment may provide robust constraints on the underlying physical mechanisms that determine the rate and relative timing of gas flows and star formation within the models. Comparing the evolution of the metallicity and its relationship to both inflows and star formation will allow us to constrain the dissipation time scales, the star formation evolution, the time delay between the onset of inflows, sufficient to dilute the metallicity, etc. In the past few years, a number of studies have investigated the role played by interactions and environment in determining the gas phase metallicities of galaxies. \\citet{donzelli00} show that merging galaxies have on average higher excitation in their optical emission line spectra than interacting pairs, and attribute this difference to lower gas metallicity in the mergers. \\citet{marquez02} studied a sample of more than one hundred spiral galaxies, ranging from isolated to interaction with strong morphological distortions, and show that the [NII]/H$\\alpha$ ratios, used as a metallicity indicator, indicate a clear trend from the metallicity to morphological type, with earlier type spirals showing higher ratios, while they found no trend with the status of the interaction. Recently, \\citet{kewley06} have derived the luminosity-metallicity relation for a sample of local galaxy pairs and compared it with that of nearby field galaxies. They found that pairs with small projected separations (s $< 20\\; \\mathrm{kpc \\; h^{-1}}$) have systematically lower metallicities than either isolated galaxies or pairs with larger separations, for a given luminosity. They also found a correlation between gas metallicity and burst strengths -- all galaxies in their interacting sample with strong central bursts having close companions and metallicities lower than the comparable field galaxies or pairs with wider separations. \\citet{rupke08}, in a study of strong interactions with high star formation rates (ultra-luminous infrared galaxies), found that the metal abundance in these intense starbursts is a factor of two lower than that of galaxies of comparable luminosity and mass. These results have been generally explained by gas inflows induced by the interaction, which dilute the pre-existing nuclear gas to produce a lower metallicity than that observed for wider separated pairs or isolated systems. More recently, \\citet{peeples08, peeples09} have studied, respectively, a sample of low-mass, high-metallicity and high-mass, low-metallicity outliers from the mass-metallicity relation of star-forming galaxies from selected from the Sloan Digital Sky Survey (SDSS).They showed that the low-mass, high-metallicity outliers are usually isolated galaxies, with no evident companion or strong interactions. On the other hand, the high-mass, low-metallicity outlierts typically consist of systems that have high star formation rates and evidence for disturbed morphologies. However, \\citet{cooper08} found a strong metallicity-density relation for star forming galaxies in the local universe, with the more metal-rich galaxies apparently favouring regions of higher galaxy overdensity \\citep[see also][]{ellison09}. They conclude that the discrepancy found with the outer studies (including those just discussed) is due to the fact that the number of close pairs (s $< 100\\; \\mathrm{kpc\\;h^{-1}}$) in the SDSS sample constitutes only a tiny fraction of the whole sample (less than 1\\% of galaxies). In this case, close pairs with low central metallicity could not contribute significantly to the scatter found in the mass-metallicity relation. Of course it can be difficult to discern trends in the gas phase metallicity during a merger when studying large samples of galaxies in SDSS as other factors may dominate the overall mass-metallicity relationship and its scatter. A confirmation that the mass-metallicity relation is affected by interactions only for close pairs showing signs of strong disturbances has been found by \\citet{dansac08}, who pointed out that, in such pairs, the gas metallicity depends on the mass ratio of the two interacting systems: less massive members are systematically enriched, while a galaxy in interaction with a comparable stellar mass companion shows a metallicity lower than that of a galaxy in isolation. So while mergers and interactions alone may not drive the overall scatter in the mass-metallicity relationship, it may be a contributing factor. All these studies suggest that the dilution and enrichment of the ISM in the central regions of disk galaxies strongly depend on the exact timing of the different processes at play: gas inflows, interaction-driven star formation, gas consumption, feedback and subsequent enrichment. While a number of numerical studies have investigated the response of the gaseous component of galaxies during tidal interactions and the subsequent star formation such interactions induce \\citep{iono04, mihos94a, springel00, cox06, cox08, kapferer05, dimatteo07, dimatteo08}, little attention has been given to the detailed evolution of the metal content during galaxy encounters. Using a galaxy pair catalogue from cosmological simulations, \\citet{perez06} have shown that the O/H abundance ratio in the central regions of close galaxy pairs (s $< 50\\; \\mathrm{kpc\\;h^{-1}}$) shows a lower level of enrichment than the mean O/H abundance ratio of a control sample, thus confirming the role played by gas inflow in diluting the metal content of the nuclear regions. Recently, \\citet{rupke10} have analysed simulations of major galaxy mergers, studying the dilution of the gas metallicity in the nuclear regions due to gas inflow. They found a dilution of about 0.1 - 0.3 dex, happening shortly after the first pericentre passage between the two galaxies. However, their models do not include either star formation prescriptions or metal enrichment, so that while it has been possible to give predictions about the strength of the dilution, nothing is still known about the role played by interaction-induced star formation in the metal dilution and then subsequent enrichement. Moreover, the exact timing of the dilution peak and its correlation with the increase in the amplitude of the star formation have not been studied yet in any detail. ", "conclusions": "In this work, we have studied the dilution and subsequent enrichment of the interstellar medium during interactions and mergers of disk galaxies of comparable mass. Our models, which include star formation and chemical enrichment, have made it possible to follow, for the first time, the evolution of the circumnuclear gas phase metallicity from the initial interaction until the coalescence phase, thus significantly improving the existing simulations and leading to insight between star formation, gas inflows and dilution, and metallicity evolution in mergers and interactions. Our main results can be summarised as follows: \\begin{itemize} \\item The dilution of the gas initially in circumnuclear regions of the merger pairs starts just after the first pericentre passage, as the disk is destabilised and gas flows inwards. For mergers, usually a second dilution peak is observed in the final phase of coalescence. A significative dilution can take place also in fly-bys, just after their close passage, mimicking to some extent what we observe in our merger simulations. Thus even flybys may be important for the star formation and evolution of galaxies. \\item On average, the amplitude of the metallicity dilution we find (0.2-0.3 dex) is in agreement with observations \\citep{kewley06,rupke08}. \\item The strongest correlation we found is between the maximum in the SFR and the strength of the circumnuclear dilution -- pairs experiencing the strongest bursts also show the strongest circumnuclear dilution. This can be explained in terms of interaction-driven gas inflows from the outer disk regions into the galaxy centre. \\item The dilution phase (defined as $z/z_{iso} <$0.8) is very rapid, half of the sample sustains this high dilution phase for less than $5\\times 10^8$ years. \\item The gas inflow causes concomitant enhanced star formation; such enhanced star formation releases a significant amount of enriched material in the ISM. This enrichment results in merger remnants which have higher final circumnuclear metallicities than the corresponding galaxies evolved in isolation. The same can happen for fly-bys which experience enhanced star formation rates. \\item Type Ia and II SNe formed during the interaction release O and Fe in the ISM, thus changing the average mass ratio of these two elements. In particular, we note that, if the merger takes place on short time scales (lower than the typical delay time of SNe Ia, which unfortunately is not well constrained), the stellar population formed during the last SF burst forms from an $\\alpha-$enriched ISM, while, for mergers with longer time scales, SNeIa can have the time to release a sufficient amount of iron in the gas. The resulting stellar population has thus lower values of $[\\alpha/Fe]$ ratios. \\item The circumnuclear metallicity, as well as the [$\\alpha/Fe$] ratios, can be used as an indicator of the timing of merger/interaction state, thus helping in disentangle projection effects, i.e., galaxies which appears close in projection, but that are not strongly interacting. The remnants of these mergers appear qualitatively consistent with what is known about the bulge of our galaxy, other nearby galaxies, and early type galaxies generally. Thus, in agreement with many authors, it appears when considering chemical evolution that mergers can reproduce early type systems. \\item If mergers play a significant role in the galaxies that have had their metallicities determined at high-redshift, then it is possible that dilution is signficantly affecting these estimates. \\end{itemize}" }, "1003/1003.3237_arXiv.txt": { "abstract": "We present the Smoothed Hessian Major Axis Filament Finder (SHMAFF), an algorithm that uses the eigenvectors of the Hessian matrix of the smoothed galaxy distribution to identify individual filamentary structures. Filaments are traced along the Hessian eigenvector corresponding to the largest eigenvalue, and are stopped when the axis orientation changes more rapidly than a preset threshold. In both N-body simulations and the Sloan Digital Sky Survey (SDSS) main galaxy redshift survey data, the resulting filament length distributions are approximately exponential. In the SDSS galaxy distribution, using smoothing lengths of $10$\\hmpc\\ and $15$\\hmpc, we find filament lengths per unit volume of $1.9 \\times 10^{-3}$~$h^2$~Mpc$^{-2}$ and $7.6 \\times 10^{-4}$~$h^2$~Mpc$^{-2}$, respectively. The filament width distributions, which are much more sensitive to non-linear growth, are also consistent between the real and mock galaxy distributions using a standard cosmology. In SDSS, we find mean filament widths of $5.5$\\hmpc\\ and $8.4$\\hmpc\\ on $10$\\hmpc\\ and $15$\\hmpc\\ smoothing scales, with standard deviations of $1.1$\\hmpc\\ and $1.4$\\hmpc, respectively. Finally, the spatial distribution of filamentary structure in simulations is very similar between $z=3$ and $z=0$ on smoothing scales as large as $15$\\hmpc, suggesting that the outline of filamentary structure is already in place at high redshift. ", "introduction": "Observational evidence for filamentary structures in the large-scale distribution of galaxies was first presented in galaxy redshifts surveys \\citep[e.g.][]{TG78,Davis82,CfaSurvey,Sathy98, 2dF, SloanGreatWall}. When similar structures were seen in cosmological N-body simulations of the dark matter distribution \\citep[e.g.][]{CosmicWeb, Sathy96, MMF, Hahn07a}, a picture of a vast `cosmic web,' in which filaments skirted the boundaries of voids and were connected by galaxy clusters, began to emerge. These filaments are thought to provide pathways for matter to accrete onto galaxy clusters \\citep[e.g.][]{CL0016} and to torque dark matter halos to align their spin axes \\citep{Hahn07a,Hahn07b,Hahn09}. Filaments also produce deep potential wells and will give rise to a gravitational lensing signal on the largest scales \\citep{LensingFils2, LensingFils}. A number of authors have claimed detections of filaments using weak lensing \\citep[e.g.][]{Kaiser98,LensingFils2,LensingFils}, but simulations predict that structure along the line of sight should produce shear comparable to that of the target filaments \\citep{Dolag06} and the evidence remains far from conclusive. In addition, the formation of filaments is accompanied by gravitational heating, which gradualy raises the temperature of the intergalactic medium over time and produces the so-called warm-hot intergalactic medium (WHIM) by $z=0$ \\citep[e.g.][]{CO99}. Perhaps the simplest and most effective means of identifying {\\it clusters} in discretely sampled fields, such as redshift surveys and N-body simulations, is the friends-of-friends algorithm \\citep[FOF,][]{FOF}, in which particle groups are assembled based on the separation of nearest neighbors. These FOF structures can then be quantified with `Shapefinders,' statistics which measure the length, breadth, and thickness of structures and are related to the Minkowski functionals \\citep{Shapefinders}. \\citet{SURFGEN} have developed an algorithm for computing the Shapefinders on structures at an arbitrary density threshold. Many of those found in data and simulations are indeed filamentary, but FOF algorithms are optimized for structures that lie above a set density threshold, a condition approximately met by clusters at the present epoch. Filaments and walls, however, are not bound and a strict density cut alone would not provide clean samples of such structures. Another algorithm, called the Skeleton \\citep{Skeleton,SkeletonData,Skeleton3D}, identifies filaments by searching for saddle points in a density field and then following the density gradient along the filament until it reaches a local maximum. Although it appears to be effective at making an outline of the cosmic network, it lacks an intuitive definition of filament ends. \\citet{SpineWeb} also lacks such definition, but has been successful at tracing the filament network in cosmological simulations using watershed segmentation \\citep[see also][]{Platen07} and a Delaunay tessellation density estimator \\citep{Schaap00}. If we wish to analyze filament length distributions or their spatial relationship to clusters, it is important to separate individual filaments in the cosmic web. Structure-finding techniques that only detect filaments between galaxy cluster pairs \\citep[e.g.][]{Pimbblet05a,Colberg05,GP09} would present a biased view of the filament-cluster relationship. An early technique for identifying filaments in two-dimensional data was developed by \\citet{GottFil} that works on a similar principle to the algorithm described in this paper. It divides the density field into a pixelized grid and identifies as filament elements any grid cell that has a larger density than its immediate neighbors along two of the four axes (including the two coordinate axes and two axes at $45^\\circ$~angles to the grid) through the grid cell. The algorithm was run on the Shane-Wirtanen galaxy count catalogue \\citep{ShaneWirtanen}, but has not been developed further. A later algorithm, presented by \\citet{Dave97}, works on a similar principle, identifying ``linked sequences'' using the eigenvectors of the inertia tensor. The authors found that the algorithm was poor at discriminating between cosmological models using CfA1-like mock galaxy catalogs, but primarily because of the small number of galaxies in the catalogs. In Paper~$1$, we used the distribution of the Hessian eigenvalues of the smoothed density field ($\\lambda$-space) on a grid to study three types of structure: clumps, filaments, and walls. Filaments were found in the $\\lambda$-space distributions at a variety of smoothing scales, ranging at least from $5-15$\\hmpc, in both N-body simulations and the galaxy distribution measured by the Sloan Digital Sky Survey \\citep[SDSS,][]{SDSS}. Furthermore, filaments were found to dominate the large-scale distribution of matter using smoothing scales of $10-15$\\hmpc, giving way to clumps with $\\sim 5$\\hmpc\\ smoothing. The fact that the eigenvalues of the Hessian can be used to discriminate different types of structure in a particle distribution is fundamental to a number of structure-finding algorithms \\citep[e.g.,][]{CPS00,Hahn07a,MMF,FR08}. However, the relationship between $\\lambda$-space and a particular structure is not always trivial. For example, one might think that a filamentary grid cell would have two positive and one negative eigenvalue. This will be true near the centre of a filament connecting two overdense filament ends, but in the vicinity of the overdensities or in the case that the filament ends at an underdensity, all three eigenvalues will become negative. In addition, when working with a smoothed density field, these criteria select regions that are near clumps and do not necessarily lie along the filament. Finally, these criteria disregard the structure's width -- for example, the regions away from the centre of the filament may have positive values of $\\lambda_2$. In this paper, we will describe a procedure to identify filaments in the three-dimensional galaxy distribution using an algorithm called the Smoothed Hessian Major Axis Filament Finder ({\\small SHMAFF}), and compare their properties in cosmological N-body simulations to those in the SDSS galaxy redshift survey. We describe our methodology, which uses the eigenvalues and eigenvectors of the smoothed Hessian matrix \\citep[see ][hereafter, Paper~$1$]{paper1}, in \\S~\\ref{subsec:Method}. In \\S~\\ref{subsec:FilPars}, we run the code with a range of possible input parameters and justify our choices for each. We discuss the behavior of the algorithm when used on Gaussian random fields in \\S~\\ref{subsec:FilGauss}, allowing us to distinguish those features of the large-scale distribution of matter that are a direct consequence of the non-linear growth of structure. In \\S~\\ref{subsec:MockFils}, we use mock galaxy catalogues to estimate the incompleteness and contamination rates of filament samples and then use these quantities to interpret the distribution of filaments found in the SDSS (\\S~\\ref{subsec:FilData}). In \\S~\\ref{subsec:Results} we summarize our results and discuss the implications of our findings. ", "conclusions": "\\label{subsec:Results} This paper develops and uses an algorithm called the Smoothed Major Axis Filament Finder to identify individual filaments in large-scale structure. In short, it uses the local eigenvectors of the density second-derivative field to define the filament axis and trace individual filaments. Filament ends are defined as points at which the rate of change of the axis of structure exceeds a specified threshold (see \\S~\\ref{subsec:Method}). In a $\\Lambda$CDM cosmological simulation, this definition produces filament samples that are consistent with our visual impression of structure on a particular scale, are complete with few duplicate detections (\\S~\\ref{subsubsec:Ct}), and are robust to sparse sampling (\\S~\\ref{subsec:SparseFils}). In addition to the smoothing scale, the filament finder takes the input parameters $C$, the maximum angular rate of change of the filament axis, and $K$, the width of filament removal in units of the smoothing length. Using Gaussian smoothing, the `best' values of these input parameters are $C=30$, $40$, and $50$\\ctunit~on $5$, $10$, and $15$\\hmpc~smoothing scales, respectively, and $K=1$ for all smoothing scales. After we collapse the fingers-of-god, contamination and completeness in filament samples found in the mock $M_r<-20.5$ galaxy distribution are $\\sim 26$ and $\\sim 81$~per~cent, respectively. Galaxy clusters are important for defining large-scale filaments and should not be removed before running the filament finder. In redshift space and on smoothing scales above $\\sim 10$\\hmpc, collapsing fingers-of-god to their mean position produces mock filament samples comparable to those in real space. In this paper, we presented two volume-limited subsamples from the northern portion of the SDSS spectroscopic survey (using the NYU-VAGC catalogue) and computed their filament distributions on $10$ and $15$\\hmpc\\ smoothing scales. These distributions were then directly compared to those found in a series of redshift-space mock galaxy catalogues generated from a cosmological simulation using the concordance cosmology. The filament length distributions found in SDSS data are very similar to those found in mock catalogues and are consistent with being drawn from an underlying exponential distribution. The width distributions of filament elements are also very similar between the SDSS data and mock catalogues, suggesting that real filaments are consistent with those in a $\\Lambda$CDM universe having $\\sigma_8=0.85$, $\\Omega_\\Lambda=0.71$, $\\Omega_m=0.29$, and $h=0.69$. Tests on a range of cosmological simulations are needed before this can be turned into a cosmological constraint. We also generated filament distributions at six redshifts in the output of a $\\Lambda$CDM cosmological N-body simulation, from $z=3$ to $z=0$. The orientation of the filament network is stable out to $z=3$ on comoving smoothing scales at least as large as $15$\\hmpc. Most of the filaments detected on $15$\\hmpc~scales at $z=0$ can be detected at $z=3$. In addition, on a given comoving smoothing scale, filament width distributions shift to smaller widths as the filaments continue to collapse. Narrower filaments will collapse more rapidly, so this also leads to a broadening of the width distributions. We have demonstrated that our filament finder is able to locate and follow real structures, perhaps most strikingly in \\S~\\ref{subsec:FilEvolve}, in which we showed that many of the same structures could be seen in a cosmological simulation at both $z=3$ and $z=0$. There is some subjective freedom in deciding what constitutes the end of a filament, as no single physical threshold stands out as a discriminator. Nevertheless, we demonstrated in \\S~\\ref{subsubsec:Ct} that the total length of the cosmic network is insensitive to the choice of $C$ above a certain scale-dependent threshold (once double detections are removed). The minimum value of $C$ needed to probe the entire filament network may be telling us about the {\\it intrinsic} clumpiness of filamentary structure, and may therefore be able to distinguish models of warm and cold dark matter. In this paper, we fully developed the SHMAFF algorithm and applied it to the low-redshift galaxy distribution, but there is much that can still be learned from its application to redshift surveys. The filament evolution seen in cosmological simulations (see \\S~\\ref{subsec:FilEvolve}) can be tested in the DEEP2 galaxy survey \\citep{DEEP2} at $z \\sim 1$, and the results of this comparison have already been presented in \\citet{EnaFil}. On $l=5$\\hmpc\\ and $l=10$\\hmpc\\ scales, they confirm a shift in the filament width distribution to smaller widths from $z \\sim 0.8$ to $z \\sim 0.1$, as well as a broadening of the filament width distribution. A possible extension of this work is a careful test of the $\\Lambda CDM$ cosmological model, including precision constraints on cosmological parameters, such as $\\sigma_8$, and tests for primordial non-Gaussianity using the length distribution of filamentary structures. In addition, it would be useful to elaborate on the relationship of large-scale filaments to galaxy clusters and to explore the properties of galaxies in filaments relative to the general galaxy population. Finally, it would be interesting to conduct a careful search for walls in SDSS. Paper~$1$ hinted at their presence in the data, but they were only present at low contrast and the $\\lambda$-space distributions were not optimal for identifying individual wall-like structures." }, "1003/1003.0528_arXiv.txt": { "abstract": "One goal of helioseismology is to determine the subsurface structure of sunspots. In order to do so, it is important to understand first the near-surface effects of sunspots on solar waves, which are dominant. Here we construct simplified, cylindrically-symmetric sunspot models, which are designed to capture the magnetic and thermodynamics effects coming from about 500~km below the quiet-Sun $\\tau_{5000}=1$ level to the lower chromosphere. We use a combination of existing semi-empirical models of sunspot thermodynamic structure (density, temperature, pressure): the umbral model of Maltby et al. (1986) and the penumbral model of Ding and Fang (1989). The OPAL equation of state tables are used to derive the sound speed profile. We smoothly merge the near-surface properties to the quiet-Sun values about 1~Mm below the surface. The umbral and penumbral radii are free parameters. The magnetic field is added to the thermodynamic structure, without requiring magnetostatic equilibrium. The vertical component of the magnetic field is assumed to have a Gaussian horizontal profile, with a maximum surface field strength fixed by surface observations. The full magnetic field vector is solenoidal and determined by the on-axis vertical field, which, at the surface, is chosen such that the field inclination is 45$^\\circ$ at the umbral-penumbral boundary. We construct a particular sunspot model based on SOHO/MDI observations of the sunspot in active region NOAA 9787. The helioseismic signature of the model sunspot is studied using numerical simulations of the propagation of f, p$_1$, and p$_2$ wave packets. These simulations are compared against cross-covariances of the observed wave field. We find that the sunspot model gives a helioseismic signature that is similar to the observations. ", "introduction": "The subsurface structure of sunspots is poorly known. Previous attempts to use helioseismology to determine the subsurface properties of sunspots have mainly been done under the assumptions that the sunspot is non-magnetic (it is usually treated as an equivalent sound-speed perturbation) and that it is a weak perturbation to the quiet-Sun. Neither of these two assumptions is justifiable. The helioseismology results have been\\,---\\,perhaps unsurprisingly\\,---\\,confusing and contradictory \\citep[e.g.,][]{gizon09}. The effects of near-surface magnetic and structural perturbations on solar waves are strong is more easily dealt with in numerical simulations. Examples of such simulations of wave propagation through prescribed sunspot models include \\cite{cameron08, hanasoge08, khomenko08, parchevsky09}. Necessary for all these attempts are appropriate model sunspots. There are numerous sunspot models available, some of which have been used in helioseismic studies. For recent reviews of sunspot models see, e.g., \\citet{Solanki2003}, \\citet{Thomas2008}, and \\citet{Moradi10}. Various authors, e.g., \\citet{khomenko08a}, \\citet{Moradi2008}, and \\citet{Moradi2009ApJ}, have constructed magnetohydrostatic parametric sunspot models for use in helioseismology. In a previous paper \\citep{cameron08}, we considered a self-similar magnetohydrostatic model. Although the deep structure of sunspots is of the utmost interest, it is likely to be swamped in the helioseismic observations unless we are able to accurately model and remove the surface effects \\citep[see, e.g.,][]{Gizon10}. The aim of this paper is to construct a simple parametric sunspot model, which captures the main effects of the near-surface layers of sunspots on the waves. In principle, the 3D properties (magnetic field, pressure, density, temperature, Wilson depression, flows) of sunspots can be inferred in the photosphere and above using spectropolarimetric inversions \\citep[e.g.,][]{mathew03}. However, today, these inversions are unavailable for most sunspots. In such circumstances, we choose to construct sunspot models that are based upon semi-empirical models of the vertical structure of typical sunspots. In producing our models, we treat separately the thermodynamic structure and the magnetic field: we do not require magnetostatic equilibrium. There are several reasons for not requiring the model to be hydrostatic. The first is that we are much more interested in the waves then in whether the background model is magnetodydrostatic. For this reason we are more interested in geting for example, the sound speed, density, and magnetic field to be close to those which are observed. A second reason is that we are not including non-axisymettric structure in our model, which certainly affects the force-balance. A third reason is that sunspots have both Evershed and moat flows, the existence of which implies a net force and so indicates that the sunspot is not strictly magnetohydrostatic. In this paper, we take the umbral and penumbral models to match those of existing semi-empirical models. In the absence of horizontal magnetic field measurements we assume that the field inclination at the umbra/penumbra boundary is approximately $45^{\\circ}$. We do not consider the effects of the Evershed or moat flows, although they could be included in the framework we set out. The surface model of the sunspot needs to be smoothly connected to the quiet-Sun model below the surface, and details will be provided in Section 2. Since we are not including the chromosphere in our simulations, we have chosen to smoothly transition the sunspot model above the surface to the quiet-Sun model. The description for constructing models including the surface structure of sunspots, which will be fleshed out in Section 2, is intended to be generic. For illustrative purposes we choose parameter values of the sunspot model that are appropriate for the sunspot in NOAA 9787, which was observed by SOHO/MDI in January 2002 \\citep{gizon09}. In Section~3, we describe the setup of the numerical simulation of the propagation of f, p$_1$, and p$_2$ wave packets through the model sunspot. We use the SLiM code \\citep{cameron07}. We briefly compare the simulations and the SOHO/MDI observations in Section~4. In Section~5, we conclude that this simple sunspot model, which is intended to be a good description of the sunspot's surface properties, leads to a good helioseismic signature and provides a testbed for future studies. ", "conclusions": "We have outlined a method for constructing magnetic models of the near-surface layers of sunspots by including available observational constraints and semi-empirical models of the umbra and penumbra. Applying this type of model to the sunspot of AR 9787, we showed that the observed helioseismic signature of the sunspot model is reasonably well captured. Our approach was to compare numerical simulations of wave propagation through the model sunspot and the observed SOHO/MDI cross-covariances of the Doppler velocity. Possible improvements of the simulation include a better treatment of wave attenuation (esp. in plage), improved initial conditions, and inclusion of the moat flow. It should also be noted that improved observations of the surface vector magnetic field by SDO/HMI will help tune the sunspot models. The dominant influence of the sunspot's surface layers on helioseismic waves means that it is necessary to model it very well in order to extract information from the much weaker signature of the sunspot's subsurface structure. The sunspot model used here accounts for most of the helioseismic signal, although there are substantial differences that remain to be explored. In any case, the model provides a testbed which is sufficiently similar to a real sunspot to be used for numerous future studies in sunspot seismology. Finally, we remark that our numerical simulations of linear waves could potentially be used to interpret other helioseismic observations than the cross-covariance (cf. helioseismic holography or Fourier-Hankel analysis), as well as other magnetic phenomena \\citep[e.g., small magnetic flux tubes, see][]{Duvall2006}." }, "1003/1003.0002_arXiv.txt": { "abstract": "Recent claims of a gamma-ray excess in the diffuse galactic emission detected by the Fermi Large Area Telescope made use of spatial templates from the interstellar medium (ISM) column density and the 408~Mhz sky as proxies for neutral pion and inverse Compton (IC) gamma-ray emission, respectively. We identify significant systematic effects in this procedure that can artificially induce an additional diffuse component with a morphology strikingly similar to the claimed gamma-ray haze. To quantitatively illustrate this point we calculate sky-maps of the ratio of the gamma-ray emission from neutral pions to the ISM column density, and of IC to synchrotron emission, using detailed galactic cosmic-ray models and simulations. In the region above and below the galactic center, the ISM template underestimates the gamma-ray emission due to neutral pion decay by approximately 20\\%. Additionally, the synchrotron template tends to under-estimate the IC emission at low energies (few GeV) and to over-estimate it at higher energies (tens of GeV) by potentially large factors that depend crucially on the assumed magnetic field structure of the Galaxy. The size of the systematic effects we find are comparable to the size of the claimed ``Fermi haze'' signal. We thus conclude that a detailed model for the galactic diffuse emission is necessary in order to conclusively assess the presence of a gamma-ray excess possibly associated to the WMAP haze morphology. ", "introduction": "\\label{sec:introduction} One of the most exciting yet observationally challenging scientific objectives of the Large Area Telescope (LAT) on board the {\\em Fermi Gamma-ray Space Telescope} \\citep{Atwood:2009ez}, is the indirect detection of particle dark matter \\citep{2008JCAP...07..013B}. However, limited gamma-ray statistics make diffuse signals arising from the pair-annihilation of dark matter difficult to differentiate from astrophysical processes. The limitation of using a diffuse signal to search for non-standard emission stems from difficulties in controlling the instrumental background and formulating a rigorous model for the astrophysical diffuse foregrounds. \\begin{figure*} \\plotone{f1.eps} \\caption{Normalized line-of-sight ratio of emission due to $\\pi^0$ decay divided by the input line-of-sight gas density plotted on a linear scale (left) and logarithmic scale (right). Normalization is done on an equal area projection for $|$b$|$~$>$~5$^\\circ$. } \\label{pi0divgas} \\end{figure*} An intriguing excess of microwave radiation in the WMAP data has been uncovered by \\citet{Finkbeiner:2004us} and \\citet{Dobler:2007wv}. The morphology and spectrum of the WMAP haze indicates a hard electron-positron injection spectrum spherically distributed around the galactic center. While the origin of this haze need not be related to {\\em new} particle physics, the possibility that the WMAP haze corresponds to synchrotron radiation of stable leptons produced by dark matter has been explored in several studies \\citep[see e.g.][]{Hooper:2007kb}. A potentially conclusive way to determine whether the WMAP haze originates from a population of energetic leptons is to observe gamma-rays produced by inverse Compton up-scattering (IC) of photons in the interstellar galactic radiation field (ISRF). Recently, \\citet{Dobler:2009xz} (hereafter D09) examined the LAT gamma-ray sky and reported an excess emission morphologically similar to the WMAP haze. D09's observations suggest a confirmation of the {\\em haze hypothesis}: that excess microwave emission stems from relativistic electron synchrotron with a spherical source distribution and a hard injection spectrum. In the ``Type 2\" and ``Type 3\" fits of D09, the excess was claimed over a best-fit background using spatial templates which employed the gas map of \\citet{1998ApJ...500..525S} (SFD) to trace gamma-ray emission from $\\pi^0$ decay, and the 408~Mhz Haslam synchrotron map \\citep{1982A&AS...47....1H} to trace IC emission from galactic cosmic ray electrons. The spatial templates (plus an isotropic component obtained by mean-subtracting the residual skymap) were used to fit the observed gamma-ray sky in energy bins spanning 2-100~GeV. This analysis uncovered a residual gamma-ray emission above and below the galactic center with a morphology and spectrum similar to that found in the WMAP dataset \\citep{2004ApJ...614..186F}. In this $Letter$, we test the following assumptions used in D09 for the removal of astrophysical foregrounds at gamma-ray energies: \\begin{itemize} \\item[(1)] that line of sight ISM maps are adequate tracers for the morphology of $\\pi^0$ emission, and \\item[(2)] that the 408 Mhz synchrotron map \\citep{1982A&AS...47....1H} is an adequate tracer for the morphology of the galactic IC emission. \\end{itemize} Assumption (1) entails neglecting the morphology of galactic cosmic-ray sources, since the observed $\\pi^0$ emission results from the line-of-sight integral of the gas density (``target'') times the cosmic-ray density (``beam''). Assumption (2) neglects the difference between the morphology of the ISRF and the galactic magnetic fields. On theoretical grounds, we expect that any detailed galactic cosmic-ray model would predict {\\em systematic deviations} from the templates used in D09. Utilizing the galactic cosmic-ray propagation code {\\tt Galprop}, we find that the procedure based on spatial templates creates deviations comparable to the amplitude of the D09 residual. Furthermore, we find that these deviations are morphologically similar to the Fermi haze. We thus conclude that the determination of an excess gamma-ray diffuse emission cannot reliably be assessed from the spatial template proxies used in the ``Type 2\" and ``Type 3\" fits of D09. We stress that our results do not claim that there is no ``haze'' in the Fermi data. In particular, the systematic effects we study here are not relavent to explain the puzzling excess emission in the ``Type 1'' fit of D09, which employes Fermi-LAT data in the 1-2 GeV range as a proxy for the morphology of the $\\pi^0$ component. We comment on this ``Type 1'' approach in Section~\\ref{sec:discussion}. \\begin{figure*} \\plotone{f2.eps} \\caption{Normalized ratio of emission due to IC at four energy values (3.55~GeV (top left), 7.08~GeV (top right), 14.1~GeV (bottom left), 35.5~GeV (bottom right)) divided by synchrotron emission at 408 MHz. Normalization is done on an equal area projection for $|$b$|$~$>$~5$^\\circ$.} \\label{icsdivsync} \\end{figure*} ", "conclusions": "\\label{sec:discussion} \\begin{figure*} \\plotone{f4.eps} \\caption{Intensity of the Interstellar Radiation function integrated across the line-of-sight, for energies of 12.4eV (top left), 1.24eV (top right), 0.124eV (bottom left), and 0.0124eV(bottom right). Plots are shown with arbitrary normalization.} \\label{fig:isrf} \\end{figure*} While a complete study of the diffuse emission detected by Fermi-LAT in terms of foreground templates is outside the scope of this $Letter$, work is currently ongoing that will also include the possible contribution to gamma radiation of nearby structures such as Loop I \\citep{jeanmarc} which are not modeled within {\\tt Galprop}. Whether improved spatial templates could conclusively pinpoint an excess diffuse emission in the Fermi-LAT data is hard to assess, as changes in the spatial morphology of each input will influence the best fit intensities of each component. However, we note that if the gas and synchrotron templates used in D09 were appropriate matches to $\\pi^0$ and IC emission, we would expect ratios very close to unity across the entire skymap. We instead find significant deviations with an intensity comparable to the Fermi haze at low energies. Moreover, these deviations are not randomly distributed across the sky, but instead contain a structure resembling the bivariate gaussian reported by D09. Thus, our results point to the possibility that the Fermi haze determined by the ``Type 2'' and ``Type 3'' templates of D09 is the result of systematic effects in the spatial template fitting procedure as opposed to the existence of a new source class. We have shown that the use of the gas map would create a best match bivariate gaussian intensity approximately 17\\% as strong as the $\\pi^0$ decay amplitude. Since we expect $\\pi^0$ decay to dominate the diffuse gamma-ray sky below 10~GeV, this error can explain a large fraction of the D09 haze. Similarly, IC due to ordinary galactic cosmic rays is expected to show an excess of approximately 20\\% in the haze region, with large variances depending on the assumed magnetic field model. At high energies ($>$~10~GeV), $\\pi^0$ decay is weaker, and thus a 15\\% error is not expected to dominate astrophysical IC or haze signals. Furthermore, some magnetic field models create synchrotron templates which would overestimate IC emission in this region. Thus, the remaining amplitude of a D09 haze may in fact be larger or smaller than reported in this region. We note additional errors may be present at high energies both due to low photon counts as well as instrumental effects such as cosmic ray contamination. In addition, nearby cosmic rays such as those in the giant radio feature Loop I can fake a diffuse gamma-ray emission in the direction of the galactic center unrelated to a WMAP haze counterpart. D09 additionally adopts another template, dubbed ``Type I'', where they use the Fermi-LAT data in the 1-2 GeV range as a proxy for the morphology of $\\pi^0$ emission. Using this template, D09 finds an excess emission that becomes more prominent at higher energies, indicating a residual with a harder spectrum than $\\pi^0$-decay emission. They note that this excess has a morphology comparable with those found by their astrophysical template fits, except with a more peaked structure that may be due to the subtraction of the 1-2 GeV components of the ICS and bremsstrahlung maps. This emission template does not claim to subtract any $\\gamma$-ray foregrounds except for those due to $\\pi^0$-decay, but notes that the residual has a morphology which is peculiar for models such as ICS of high energy astrophysical leptons. As this template does not make use of either the synchrotron or SFD maps, it falls outside the scope of our analysis. We agree that the emission is likely not due to $\\pi^0$, and it is difficult to construct an astrophysical source distribution which is more pronounced towards high latitudes than near the galactic plane. Any galactic electron sources should lie close to the galactic plane, and would primarily upscatter starlight to GeV energies. In Figure~\\ref{fig:isrf}, we show the ISRF used in our {\\tt Galprop} models integrated over the line-of-sight. We see that the ISRF is also strongest along the galactic plane, extended in longitude rather than latitude by a ratio from 5-4 at 12.4eV to 9-1 at 0.0124eV. We further note that the ISRF dims by between 63-72\\% between 10$^\\circ$-30$^\\circ$ latitude. This decay is itself slightly stronger than the decay in diffuse emission used in the D09 gaussian template. This ISRF is morphologically identical to the IC morphologies obtained by comvolving the input ISRF with a isotropic and monochromatic input electron spectrum. Since this spectrum is much broader than the haze determined by D09, and falls slightly more quickly at high galactic latitudes, the input source class would have to be significantly peaked above and below the galactic plane, with very little extent along the plane. Furthermore, the source class must have a flux at 4.25~kpc above the galactic center, which is almost equvalent to the flux at 1.5~kpc above the galactic center, in order that the product of the lepton flux and ISRF fall off by a ratio comparable to the gaussian haze. The characteristics of this input source spectrum is not similar to those of expected galactic sources. This problem night be overcome by several ad hoc changes in the ISRF around the galactic center region, or by changes in the convection currents and diffusion constants away from the galactic plane \\citep{2009arXiv0910.2027G}. However, these same changes will greatly alter the assumed morphology of astrophysical ICS as well, possibly eliminating the need for an extra diffuse component. In summary, we showed that significant systematic effects make it difficult to reliably assess a diffuse gamma-ray emission in the region associated to the WMAP haze. A fully self-consistent galactic cosmic-ray model is necessary to model the astrophysical diffuse emission from the Galaxy and to compare it to the Fermi-LAT data." }, "1003/1003.0833_arXiv.txt": { "abstract": "{The instrument BLAST (Balloon-borne Large-Aperture Submillimeter Telescope) performed the first deep and wide extragalactic survey at 250, 350 and 500 $\\mu$m. The extragalactic number counts at these wavelengths are important constraints for modeling the evolution of infrared galaxies.} {We estimate the extragalactic number counts in the BLAST data, which allow a comparison with the results of the P(D) analysis of Patanchon et al. (2009). } { We use three methods to identify the submillimeter sources. 1) Blind extraction using an algorithm when the observed field is confusion-limited and another one when the observed field is instrumental-noise-limited. The photometry is computed with a new simple and quick point spread function (PSF) fitting routine (FASTPHOT). We use Monte-Carlo simulations (addition of artificial sources) to characterize the efficiency of this extraction, and correct the flux boosting and the Eddington bias. 2) Extraction using a prior. We use the \\textit{Spitzer} 24 $\\mu$m galaxies as a prior to probe slightly fainter submillimeter flux densities. 3) A stacking analysis of the \\textit{Spitzer} 24 $\\mu$m galaxies in the BLAST data to probe the peak of the differential submillimeter counts. } { With the blind extraction, we reach 97~mJy, 83~mJy and 76~mJy at 250~$\\mu$m, 350~$\\mu$m and 500~$\\mu$m respectively with a 95\\% completeness. With the prior extraction, we reach 76~mJy, 63~mJy, 49~mJy at 250~$\\mu$m, 350~$\\mu$m and 500~$\\mu$m respectively. With the stacking analysis, we reach 6.2~mJy, 5.2~mJy and 3.5~mJy at 250 $\\mu$m, 350~$\\mu$m and 500~$\\mu$m respectively. The differential submillimeter number counts are derived, and start showing a turnover at flux densities decreasing with increasing wavelength.} { There is a very good agreement with the P(D) analysis of Patanchon et al. (2009). At bright fluxes ($>$100~mJy), the Lagache et al. (2004) and Le Borgne et al. (2009) models slightly overestimate the observed counts, but the data agree very well near the peak of the differential number counts. Models predict that the galaxy populations probed at the peak are likely $z\\sim 1.8$ ultra-luminous infrared galaxies.} ", "introduction": "Galaxy number counts, a measurement of the source surface density as a function of flux density, are used to evaluate the global evolutionary photometric properties of a population observed at a given wavelength. These photometric properties mainly depend on the source redshift distribution, spectral energy distribution (SED), and luminosity distribution in a degenerate way for a given wavelength. Even though this is a rather simple tool, measurements of number counts at different observed wavelengths greatly help in constraining those degeneracies. Backward evolution models, among these \\citet{Chary2001,Lagache2004,Gruppioni2005,Franceschini2009,Le_Borgne2009,Pearson2009,Rowan2009,Valiante2009} are able to broadly reproduce (with different degrees of accuracy) the observed number counts from the near-infrared to the millimeter spectral ranges, in addition to other current constraints, like such as measured luminosity functions and the spectral energy distribution of the Cosmic Infrared Background (CIB) \\citep{Puget1996,Fixsen1998,Hauser1998,Lagache1999,Gispert2000,Hauser2001,Kashlinsky2005,Lagache2005,Dole2006}. In the details, however, the models disagree in some aspects like the relative evolution of luminous and ultra-luminous infrared galaxies (LIRG and ULIRG) and their redshift distributions, or the mean temperature or colors of galaxies, as is shown for instance in \\citet{Le_Floch2009} from Spitzer 24~$\\mu$m deep observations. One key spectral range lacks valuable data to get accurate constraints as yet: the sub-millimeter range, between 160~$\\mu$m and 850~$\\mu$m, where some surveys were conducted on small areas. Fortunately this spectral domain is intensively studied with the BLAST balloon experiment \\citep{Devlin2009} and the \\textit{Herschel} and \\textit{Planck} space telescopes. This range, although it is beyond the maximum of the CIB's SED in wavelength, allows us to constrain the poorly-known cold component of galaxy SED at a redshift greater than a few tenths. Pioneering works have measured the local luminosity function \\citep{Dunne2000} and shown that most milli-Jansky sources lie at redshifts $z>2$ \\citep{Ivison2002,Chapman2003a,Chapman2005,Ivison2005,Pope2005,Pope2006}. Other works showed that the galaxies SED selected in the submillimeter range \\citep{Benford1999,Chapman2003b,Sajina2003,Lewis2005,Beelen2006,Kovacs2006,Sajina2006,Michalowski2009} can have typically warmer temperatures and higher luminosities than galaxies selected at other infrared wavelengths. Data in the submillimeter wavelength with increased sensitivity are thus needed to match the depth of infrared surveys, conducted by Spitzer in the mid- and far-infrared with the MIPS instrument \\citep{Rieke2004} at 24~$\\mu$m, 70~$\\mu$m and 160~$\\mu$m \\citep{Chary2004,Marleau2004,Papovich2004,Dole2004,Frayer2006a,Frayer2006b,Rodighiero2006,Shupe2008,Frayer2009,Le_Floch2009,Bethermin2010a} as well as the near-infrared range with the IRAC instrument \\citep{Fazio2004a} between 3.6~$\\mu$m and 8.0~$\\mu$m \\citep{Fazio2004b,Franceschini2006,Sullivan2007,Barmby2008,Magdis2008,Ashby2009}. Infrared surveys have allowed the resolution of the CIB by identifying the contributing sources -- directly at 24~$\\mu$m and 70~$\\mu$m, or indirectly trough stacking at 160~$\\mu$m \\citep{Dole2006,Bethermin2010a}. Although large surveys cannot solve by themselves all the unknowns about the submillimeter SED of galaxies, the constraints given by the number counts can greatly help in unveiling the statistical SED shape of submillimeter galaxies as well as the origin of the submillimeter background. The instrument BLAST (Balloon-borne Large-Aperture Submillimeter Telescope, \\citet{Pascale2008}) performed the first wide and deep survey in the 250-500~$\\mu$m range \\citep{Devlin2009} before the forthcoming \\textit{Herschel} results. \\citet{Marsden2009} show that sources detected by \\textit{Spitzer} at 24 $\\mu$m emit the main part of the submillimeter background. \\citet{Khan2009} claimed that only 20\\% of the CIB is resolved by the sources brighter than 17~mJy at 350~$\\mu$m. \\citet{Patanchon2009} has performed a $P(D)$ fluctuation analysis to determine the counts at BLAST wavelength (250~$\\mu$m, 350~$\\mu$m and 500~$\\mu$m). In this paper we propose another method to estimate the number counts at these wavelengths and compare the results with those of \\citet{Patanchon2009}. \\begin{figure} \\centering \\includegraphics[width=8cm]{13910f1.eps} \\caption{\\label{fig:figext} Position of sources brighter than the 95\\% completeness flux at 250~$\\mu$m in deep zone. \\textit{Top}: initial map. \\textit{Bottom}: residual map. The area out of the mask are represented darker. These 1$^\\circ \\times $1$^\\circ$ map are centered on the coordinates (RA,Dec) = (3h32min30s,-27$^\\circ$50'). The horizontal axis is aligned with the right ascension.} \\end{figure} ", "conclusions": "Our analysis provides new stacking counts, which can be compared with the \\citet{Patanchon2009} P(D) analysis. We have a good agreement between the different methods. Nevertheless, some methods are more efficient in a given flux range. The blind extraction and the extraction using a prior give a better sampling in flux and slightly smaller error bars. The P(D) analysis uses only the pixel histogram and thus looses the information on the shape of the sources. The blind extraction is a very efficient method for extracting the sources, but lots of corrections must be applied carefully. When the confusion noise totally dominates the instrumental noise, the former must be determined accurately, and the catalog flux limit must take this noise \\citep{Dole2003} into account. Estimating the counts from a catalog built using a prior is a good way to deal with the flux boosting effect. This method is based on assumptions however. We assume that all sources brighter than the flux cut at the studied wavelength are present in the catalog extracted using a prior. We also assume a flux distribution at the studied wavelength for a selection at the prior wavelength to correct for the Eddington bias. Consequently an extraction using a prior must be used in a flux range where the blind extraction is too affected by the flux boosting to be accurately corrected. P(D) analysis and stacking counts estimate the counts at flux densities below the detection limit. These methods have different advantages. The P(D) analysis fits all the fluxes at the same time, where the stacking analysis flux depth depends on the prior catalog's depth (24 $\\mu$m \\textit{Spitzer} for example). But the P(D) analysis with a broken power-law model is dependent on the number and the positions of the flux nodes. The uncertainty due to the parameterization was not evaluated by \\citet{Patanchon2009}. The stacking counts on the other hand are affected by biases due to the color dispersion of the sources. The more the prior and stacked wavelength are correlated, the less biased are the counts. A way to overcome this bias would be to use a selection of sources (in redshift slices for example), which would reduce the color dispersion, and the induced bias; we did not use this approach here because of a low signal-to-noise ratio. The stacking and P(D) analysis are both affected by the clustering in different ways. For the stacking analysis this effect depends on the size of the PSF. This effect is small for BLAST and will be smaller for SPIRE. The clustering broadens the pixel histogram. \\citet{Patanchon2009} show that it is negligible for BLAST. Clustering will probably be an issue for SPIRE. The cirrus can also affect the P(D) analysis and broaden the peak. \\citet{Patanchon2009} use a high-pass filtering that reduces the influence of these large scale structures. The methods used in this paper will probably be useful to perform the analysis of the \\textit{Herschel} SPIRE data. The very high sensitivity and the large area covered will reduce the uncertainties and increase the depth of the resolved source counts. Nevertheless, according to the models (e.g. \\citet{Le_Borgne2009}), the data will also be quickly confusion-limited and it will be very hard to directly probe the break of the counts. The P(D) analysis of the deepest SPIRE fields will allow us to constrain a model with more flux nodes and to better sample the peak of the normalized differential number counts. The instrumental and confusion noise will be lower, and a stacking analysis per redshift slice will probably be possible. These analyses will give stringent constraints on the model of galaxies and finally on the evolution of the infrared galaxies." }, "1003/1003.0658_arXiv.txt": { "abstract": "HM Cancri is a candidate ultracompact binary white dwarf with an apparent orbital period of only 5.4 minutes, as suggested by X-ray and optical light-curve modulations on that period, and by the absence of longer-period variability. In this Letter we present Keck-I spectroscopy which shows clear modulation of the helium emission lines in both radial velocity and amplitude on the 5.4-minute period and no other. The data strongly suggest that the binary is emitting \\ions{He}{i} 4471 from the irradiated face of the cooler, less massive star, and \\ions{He}{ii} 4686 from a ring around the more massive star. From their relative radial velocities, we measure a mass ratio \\q. We conclude that the observed 5.4-minute period almost certainly represents the orbital period of an interacting binary white dwarf. We thus confirm that \\obj\\ is the shortest-period binary star known: a unique test for stellar evolution theory, and one of the strongest known sources of gravitational waves for the \\emph{Laser Interferometer Space Antenna (LISA)}. ", "introduction": "Two interacting binary stars have been discovered which appear to have orbital periods shorter than ten minutes: V407 Vul, with a period of 9.5 minutes \\citep{Mot96,Cro98}, and HM Cnc, with a period of 5.4 minutes \\citep{Isr02}. The uniquely short period of 5.4 minutes, if it is the orbital period, implies that \\obj\\ must have formed from two white dwarfs, driven together as a result of gravitational-wave radiation. It may currently be experiencing stable mass transfer through Roche-lobe overflow. Because it is potentially so extreme and unique, substantial effort has been put into unveiling \\obj's true nature, but as yet without conclusive results. The key observational data show that: (1) there is no evidence for variability on periods other than 5.4 and 9.5 minutes in \\obj\\ and V407 Vul, respectively \\citep{Ram00,Ram02a}; (2) the optical flux maxima lead the X-ray maxima by about 90 degrees in both systems \\citep{Bar07}; and (3) the periods are decreasing in both systems \\citep{Str02,Str03,Str04,Str05,Hak03,Hak04,Isr04}. Three competing models have been proposed for V407 Vul and \\obj. One of them, the `Intermediate Polar (IP)' model, predicts that these systems are not in fact ultracompact binaries but have rather mundane orbital periods of several hours. The ultrashort periods then represent the spins of magnetic white dwarfs \\citep{Nor04}. The X-rays as well as the variable optical flux originate from the accretion flow crashing onto the magnetic poles of the magnetic white dwarf during part of the accretor's spin cycle. The spin-up of the X-ray and optical periods may be expected since the magnetic white dwarf is accreting matter of high specific angular momentum. The absence of variability on the much longer orbital period could be explained if the orbital planes of both systems are viewed exactly face-on. The observed phase offsets between the optical and X-ray light-curves are difficult to explain \\citep{Nor04}, and the emission lines are unusually weak for IPs; in V407 Vul there appear to be no emission lines at all \\citep{Ram02b,Ste06}. The second, `Unipolar Inductor (UI)' model, is essentially a more energetic version of the Jupiter--Io system \\citep{Wu02,Dal06,Dal07}. A magnetic white dwarf is orbited by another, non-magnetic one; the magnetic field induces an electrical potential across the non-magnetic white dwarf which sets up currents along magnetic flux tubes connecting the two stars. Ohmic dissipation of these currents in the flux tubes' footpoints on the magnetic star gives rise to the X-rays. In this model the 5.4-minute period is the orbital period, but the two stars are detached. Since the binary loses angular momentum due to gravitational-wave radiation, it is expected to evolve towards shorter orbital periods, as observed. The main problem is the offset between the optical and X-ray flux, which requires the footpoints on the magnetic star's surface to be almost 90 degrees ahead of the orbiting non-magnetic star in both systems \\citep{Bar05,Bar07}. The third, `AM CVn' model, has a Roche-lobe filling white dwarf losing mass to a more massive white dwarf (\\citealt{Mar02,Ram02b}, or \\citealt{Cro98} for a magnetic version). The orbital period is 5.4 minutes in the case of \\obj, and the accretion stream hits the accretor directly without forming a disk. The phase offset between optical and X-ray flux is naturally accounted for, since the accretion stream will deflect and hit the accretor off-axis. The absence of longer periods is also expected, but the observed decrease of the periods in \\obj\\ and V407 Vul is considered problematic for this model \\citep{Str02,Str04}, although solutions to this problem have been put forward (\\citealt{Del06,DAn06}; Roelofs \\& Deloye in preparation). This Letter describes our search for kinematic evidence, using phase-resolved spectroscopy of \\obj, for or against these different models. Section \\S\\,\\ref{sec:obs} describes our observations and data reduction. We present our results in \\S\\,\\ref{sec:results} and discuss their implications in \\S\\,\\ref{sec:conclusion}. ", "conclusions": "\\label{sec:conclusion} \\begin{figure} \\begin{center} \\includegraphics[width=84mm]{f3.pdf} \\caption{`AM CVn' model of \\obj, with the donor star on the left, transferring mass to the accretor. The impact of the accretion stream into the accretor causes the X-rays. The irradiated face of the donor star is the source of the \\ions{He}{i} 4471 emission line, and the main source of the modulated optical emission. An equatorial belt or disk around the accretor is the source of the \\ions{He}{ii} emission lines. Masses $M_2=0.27\\,M_\\odot$, $M_1=0.55\\,M_\\odot$ are assumed (see text); the center-of-mass is at the origin. The orbital phase (from Eq.~(\\ref{eq:ephemeris})) would be $\\sim$0.35 as viewed from the bottom of this page.} \\label{fig:model} \\end{center} \\end{figure} \\begin{figure} \\begin{center} \\includegraphics[width=84mm]{f4.pdf} \\caption{Doppler tomograms of \\ions{He}{i} 4471 (gray-scale) and \\ions{He}{ii} 4686 (contours). The (assumed) irradiation-induced \\ions{He}{i} 4471 emission from the secondary star has been aligned with the positive $K_Y$-axis. The gray-scale indicates the flux overdensity as a percentage of the continuum. \\ions{He}{ii} 4686 is emitted in a ring centered on the expected location of the primary star, for mass ratios \\q. Solid and dashed contours indicate where the \\ions{He}{ii} 4686 flux rises above and drops below the average in the map, respectively. A bright spot occurs on the left side of the ring, approximately 110 degrees ahead of the secondary star.} \\label{fig:back} \\end{center} \\end{figure} \\subsection{The Nature of \\obj} The observed radial-velocity variations in the spectral lines strongly suggest that \\obj\\ has a 5.4-minute orbital period. We favor the semi-detached binary white dwarf (or AM CVn) model, shown in Figure \\ref{fig:model}. It is the only model that can account for the spectroscopy presented in this Letter. The IP model predicts an orbital period of hours rather than minutes, and the observed line kinematics do not match predictions by \\citet{Nor04} for spectral line variability on the accretor's spin period (if any). We see the broad, fairly constant (and in the case of \\ions{He}{ii} 4686, double-peaked) \\ions{He}{ii} emission as a clear signature of accretion, which is not predicted by the UI model. We conclude that \\obj\\ is by far the shortest-period binary star known. Its direct progenitor was probably a detached binary white dwarf, since evolutionary models with initially non-degenerate donor stars cannot reach such short orbital periods (e.g.\\ \\citealt{Yun08,Slu05}). The observed decrease of \\obj's 5.4-minute period, considered the main obstacle for the AM CVn model, may be expected naturally if it is a semi-detached binary white dwarf with this orbital period, due to the relatively long mass-transfer turn-on timescales predicted for fairly low-mass helium donors (\\citealt{Del06}; Roelofs \\& Deloye in preparation). \\subsection{System Parameters} Based on the radial velocities of the \\ions{He}{i} 4471 and \\ions{He}{ii} 4686 emission lines moving in anti-phase, we can estimate the ratio of donor to accretor mass $q\\equiv M_2/M_1$ in \\obj. The \\ions{He}{i} 4471 line seems to originate from the irradiated face rather than the exact mass center of the donor star, and so we must estimate a `K-correction' (e.g.\\ \\citealt{Mun05}) to obtain the donor's actual projected orbital velocity from the observed emission-line velocity. Without a K-correction we get an upper limit of \\qhigh. Assuming that the \\ions{He}{i} line originates from the inner Lagrange point (the maximum conceivable K-correction) gives a lower limit \\qlow. The true value will be in between; assuming a flat probability density distribution between our upper and lower limits gives \\q. Figure~\\ref{fig:back} shows a combined Doppler tomogram of the \\ions{He}{i} 4471 and \\ions{He}{ii} 4686 lines, showing the (putative) irradiated donor star and ring-like emission centered on the (proposed) accreting star. The mass ratio derived above can easily be estimated from this figure. A further interesting feature is the `bright spot' in the \\ions{He}{ii} 4686 emission, occurring in the lower left quadrant of the Doppler tomogram. The velocity vectors from the center-of-mass to the donor star and to this \\ions{He}{ii} 4686 bright spot make a 110 degree angle approximately. Kinematically, this bright spot matches with a spot approximately on the side of the accreting star as seen from the donor, if the material in the bright spot is moving at roughly half the Keplerian (break-up) velocity near the surface of the accretor (but depending on its mass). An accretion stream impact spot at this location was proposed by \\citet{Bar07} to explain the relative phases of the optical and X-ray light-curves. Assuming $M_2=0.13\\,M_\\odot$, which is the minimum donor star mass corresponding to a Roche-lobe-filling degenerate helium object, an accretor mass $M_1\\approx 0.55\\,M_\\odot$ is required to get an accretion stream impact spot sufficiently on the side of the accretor, as shown by \\citet{Bar07}. Our mass ratio \\q\\ implies that both the donor and the accretor have to be more massive than this to still produce an impact spot sufficiently on the side (but see \\citealt{Woo09} for a possible relaxation of this constraint). The measured rate of change of the orbital period in \\obj\\ (Eq.~\\ref{eq:ephemeris}) also suggests that the donor star has to be significantly more massive than its fully-degenerate value, if one assumes that the rate of change matches the secular rate set by gravitational-wave radiation (neglecting mass transfer). Given $q\\approx0.50$, masses $M_2=0.27$, $M_1=0.55$ would be required to yield the observed period derivative. We note that these are quite typical masses for binary white dwarfs according to theoretical models \\citep{Nel01}. Assuming instead $M_2=0.13\\,M_\\odot$, which is again the minimum mass of a Roche-lobe filling helium donor, an accretor mass $M_1=1.3\\,M_\\odot$ would be required. This is incompatible with our mass ratio measurement. Note that mass transfer reduces the orbital period change rate relative to the case of a detached binary, so that higher masses would be required; but in practice, it is likely that the mass transfer is not yet sufficiently developed to have a big influence \\citep{Del07}. We can conclude that the donor star is probably at least twice as massive as a fully-degenerate helium donor would be. For masses $M_2=0.27\\,M_\\odot$, $M_1=0.55\\,M_\\odot$ and assuming that the radial-velocity semi-amplitude of the wobbling \\ions{He}{ii} 4686 line represents the projected orbital velocity of the accretor, we obtain an inclination $i\\approx38^\\circ$ of the orbital plane. \\subsection{\\obj\\ as a Gravitational Wave Source} With the system parameters derived in the previous section, we can calculate the gravitational-wave strain amplitude at Earth. The distance to \\obj\\ is the largest remaining uncertainty; it is probably $\\sim$5 kpc based on the expected sizes and measured temperatures of the white dwarf components \\citep{Bar07}. This gives a dimensionless gravitational-wave strain amplitude $h\\simeq 1.0\\times 10^{-22}$. This, together with the design sensitivities of space-borne gravitational wave detectors like the \\emph{Laser Interferometer Space Antenna (LISA)}, makes \\obj\\ one of the easiest detectable sources of gravitational waves known (see \\citealt{Nel09})." }, "1003/1003.0685_arXiv.txt": { "abstract": "We present a new high-resolution N-body/SPH simulation of an encounter of two gas-rich disk galaxies which closely matches the morphology and kinematics of the interacting Antennae galaxies (NGC 4038/39). The simulation includes radiative cooling, star formation and feedback from SNII. The large-scale morphology and kinematics are determined by the internal structure and the orbit of the progenitor disks. The properties of the central region, in particular the starburst in the overlap region, only match the observations for a very short time interval ($\\Delta t \\approx$ 20 \\Myr) after the second encounter. This indicates that the Antennae galaxies are in a special phase only about 40 $\\Myr$ after the second encounter and 50 $\\Myr$ before their final collision. This is the only phase in the simulation when a gas-rich overlap region between the nuclei is forming accompanied by enhanced star formation. The star formation rate as well as the recent star formation history in the central region agree well with observational estimates. For the first time this new model explains the distributed extra-nuclear star formation in the Antennae galaxies as a consequence of the recent second encounter. The proposed model predicts that the Antennae are in a later merger stage than the Mice (NGC 4676) and would therefore lose their first place in the classical Toomre sequence. ", "introduction": "\\begin{figure*} \\centering \\plotone{./p1.eps} \\caption{Simulated gas properties projected on top of \\HI kinematic data by \\citet{HibbardEtAl2001AJ} at the time of best match. Simulated gas particles are displayed in blue (NGC 4038) and red (NGC 4039). Yellow points represent the observational data. {\\it Upper left panel:} Projected positions in the plane-of-the-sky (x$^\\prime$-y$^\\prime$ plane). {\\it Upper right and lower left panel:} Declination (y$^\\prime$) and Right Ascension (x$^\\prime$) against line-of-sight velocity. Similarly to the observations, we apply a column density threshold of $N_\\mathrm{gas} = 10^{20} \\cm^{-2}$ to the simulated gas distribution.} \\label{pic1:PV} \\end{figure*} \\label{Intro} In the local universe ($z < 0.3$) about $\\sim5-10\\%$ of all galaxies are interacting and merging (e.g. \\citealp{2008ApJ...672..177L,2010ApJ...709.1067B}). Mass assembly via this mechanism was more important at earlier cosmic times when major mergers were more frequent \\citep[e.g.][]{2002ApJ...565..208P,ConseliceEtAl2003AJ....126.1183C} and also more gas-rich \\citep[e.g.][]{2010Natur.463..781T}. Major mergers dramatically affect the formation and evolution of galaxies. By inducing tidal torques they can efficiently transport gas to the centers of the galaxies \\citep{BarnesHernquist1996ApJ, 2006MNRAS.372..839N}, trigger star formation \\citep{Mihos&Hernquist1996ApJ, 2000MNRAS.312..859S, 2008MNRAS.384..386C}, feed super-massive black holes \\citep{2005ApJ...630..705H, SpringelDiMatteoHernquist2005MNRAS, JohanssonEtAl2009ApJ} and convert spiral galaxies into intermediate-mass ellipticals \\citep{1992ApJ...393..484B, NaabBurkert2003ApJ,2004AJ....128.2098R,2009ApJ...690.1452N}. The \\object{Antennae} galaxies (NGC 4038/39) are the nearest and best-studied example of an on-going major merger of two gas-rich spiral galaxies. The system sports a beautiful pair of elongated tidal tails extending to a projected size of $\\sim20 \\arcmin$ (i.e. $106 \\kpc$ at an assumed distance of 22 Mpc), together with two clearly visible, still distinct galactic disks. The latter has been assumed to be an indication of an early merger state, putting the system in the first place of the \\citet{Toomre1977egsp.conf..401T} merger sequence of 11 prototypical mergers. Due to their proximity and the ample number of high-quality observations covering the spectrum from radio to X-ray \\citep[e.g.][]{NeffUlvestad2000AJ, WangEtAl2004ApJS, WhitmoreEtAl1999AJ, 2005ApJ...619L..87H, ZezasEtAl2006ApJS..166..211Z} the Antennae provide an ideal laboratory for understanding the physics of merger-induced starbursts through comparison with high-resolution simulations. At the center of the Antennae galaxies, HST imaging has revealed a large number of bright young star clusters ($\\gtrsim 1000$) which plausibly have formed in several bursts of star formation induced by the interaction \\citep{WhitmoreEtAl1999AJ}. The spatial distribution and the age of these clusters are correlated: the youngest clusters are found in the overlap region ($\\tau < 5 \\Myr$), while the young starburst is generally located in the overlap and a ring-like configuration in the disk of NGC 4038 ($\\tau \\lesssim 30 \\Myr$). An intermediate-age population ($\\tau = 500-600 \\Myr$) is distributed throughout the disk of NGC 4038 \\citep{WhitmoreEtAl1999AJ,ZhangFallWhitmore2001ApJ}. Of particular interest is the spectacular nature of an extra-nuclear starburst observed in the dusty overlap region between the merging galactic disks \\citep{MirabelEtAl1998A&A,WangEtAl2004ApJS}. The Antennae seem to be the only interacting system where an off-center starburst is outshining the galactic nuclei in the mid-IR \\citep{XuEtAl2000ApJ} and among only a few systems which show enhanced inter-nuclear gas concentrations \\citep{1999ApJ...524..732T}. To date, this prominent feature has not been reproduced in any simulation of the Antennae system \\citep[see][]{BarnesHibbard2009AJ}. Thus, the question remains whether this feature cannot be captured by current sub-grid modeling of star formation or whether the previous dynamical models (e.g. initial conditions) were not accurate enough. A first simulation of the Antennae galaxies was presented by \\citet{Toomre&Toomre1972ApJ}, reproducing the correct trends in the morphology of the tidal tails. \\citet{Barnes1988ApJ} repeated the analysis with a self-consistent multi-component model consisting of a bulge, disk and dark halo component. \\citet{MihosBothunRichstone1993ApJ} included gas and star formation in their model and found the star formation to be concentrated in the nuclei of the disks, thus, not reproducing the overlap star formation. In this Letter, we present the first high-resolution merger simulation of NGC 4038/39 with cosmologically motivated progenitor disks galaxy models. We are able to match both the large-scale morphology and the line-of-sight kinematics, as well as important key aspects of the distribution and ages of newly-formed stars at the center of the Antennae, being a direct consequence of the improved merger orbit. ", "conclusions": "\\label{discussion} The new numerical model for the Antennae galaxies presented in this Letter improves on previous models in several key aspects. We find an excellent morphological and kinematical match to the observed large-scale morphology and \\HI velocity fields \\citep{HibbardEtAl2001AJ}. In addition, our model produces a fair morphological and kinematical representation of the observed central region. A strong off-center starburst naturally develops in the simulation - in good qualitative and quantitative agreement with the observed extra-nuclear star-forming sites \\citep[e.g.][]{MirabelEtAl1998A&A,WangEtAl2004ApJS}. This is a direct consequence of our improved merger orbit. All previous studies using traditional orbits failed to reproduce the overlap starburst \\citep [see e.g.][] {KarlEtAl2008AN....329.1042K}. The exact timing after the second encounter shortly before the final merger ensures that the galaxies are close enough for the efficient tidally-induced formation of the overlap region. The formation of the extra-nuclear starburst is likely to be supported by compressive tidal forces which can dominate the overlap region in Antennae-like galaxy mergers during close encounters \\citep{2008MNRAS.391L..98R,2009ApJ...706...67R}. Energetic feedback from supernovae prevents the depletion of gas by star formation at earlier merger stages and ensures that by the time of the second encounter enough gas is left over to fuel the starburst. Simulating the system with an identical orbit, but now employing an isothermal EQS without feedback from supernovae ($q_\\mathrm{EQS}=0$) resulted in most of the gas being depleted by star formation at earlier phases of the merger, i.e. during the first encounter. Our model predicts that the observed off-center starburst is a transient feature with a very short lifetime ($ \\approx 20$ \\Myr) compared to the full merger process ($\\approx 650 \\Myr$ from first encounter to final merger). This fact serves as a plausible explanation for why such features are rarely observed in interacting galaxies \\citep{XuEtAl2000ApJ}. However, the observed puzzling gas concentration between the two nuclei of the \\object{NGC 6240} merger system might be of a similar origin (\\citealp[Engel in prep. 2010]{1999ApJ...524..732T}) suggesting that the Antennae overlap region, although rare, is not a unique feature. In addition, our improved model can serve as a solid basis and testbed for further theoretical studies of the enigmatic interacting NGC 4038/39 system. For example, the overlap region in the Antennae is dominated by molecular gas, which we do not model in the simulation presented here. Given that we now have a dynamically viable method for forming the overlap region, detailed investigations of the molecular gas formation process can be undertaken using improved theoretical models \\citep[e.g.][]{2008ApJ...680.1083R,2009ApJ...707..954P}. In a first application using this new orbital configuration we have been able to qualitatively and quantitatively reproduce the magnetic field morphology of the Antennae galaxies \\citep{2009arXiv0911.3327K}. Finally, accurate modeling of nearby interacting systems also provide unique insights into the merger dynamics and timing of observed merger systems. The Antennae galaxies are traditionally in the first place in the classical Toomre sequence which orders galaxies according to their apparent merger stage \\citep{Toomre1977egsp.conf..401T} with the Mice (\\object{NGC 4676}) being between their first and second pericenter \\citep{Barnes2004MNRAS} and thus in second place behind the Antennae. According to our proposed model the Antennae galaxies are in a later merger phase, after the second pericenter. As a consequence the Antennae would lose their first place, and thus requiring a revision of the classical Toomre sequence." }, "1003/1003.5296_arXiv.txt": { "abstract": "We employ the metric of Schwarzschild space surrounded by quintessential matter to study the trajectories of test masses on the motion of a binary system. The results, which are obtained through the gradually approximate approach, can be used to search for dark energy via the difference of the azimuth angle of the pericenter. The classification of the motion is discussed. ", "introduction": "Several astrophysical observational data have shown that our universe is undergoing an era of accelerated expansion \\cite{trj1,trj2,trj3}. Therefore, in order to explain this bizarre phenomenon, various models of cosmology have been put forward, ranging from the simplest one of a cosmological constant to the scalar field theory of dark energy and modified gravitational theory as well \\cite{trj4,trj5,trj6}. On the other hand, taking dark energy for the presumed cosmological component is also used, combined with the Einstein field equations, to deal with some local gravitational issues \\cite{trj7,trj8}, such as the relevant features of black holes. Recently, Kiselev has brought forward new static spherically symmetric exact solutions of the Einstein equations, either for a charged or uncharged black hole surrounded by quintessential matter or free from it, which satisfy the condition of the additivity and linearity \\cite{trj9}. Otherwise, using a linear post-Newtonian approach, Kerr et al. considered the orbital differential equations for test bodies of a binary system in the Kerr-de Sitter spacetime and gave the elliptically orbital solution \\cite{trj10}. The solutions for parabolic and hyperbolic orbits can be obtained via the formulae of \\cite{trj11}. In this paper, we study the trajectories of test masses in a binary system under the general metric mentioned in \\cite{trj9}. Four discrete theoretical values of the state parameter $\\omega$ from $0$ to $-1$ are used to solve the orbital equations, respectively. Similar to the classic tests of Einstein's general relativity \\cite{trj12}, the formulae gotten here can be used to explore whether dark energy exists or not along with the improvement upon observational techniques in astronomy. By introducing the effective potential, we also discuss the classification of the motion. The metric for a black hole surrounded by quintessence is laid out in Section II. We calculate trajectories of test masses in detail in Section III. In Section IV, we display the classification of the motion and a brief conclusion follows. We use the metric signature $(+,-,-,-)$ and make $G$, $\\hbar$ and $c$ equal to unity. ", "conclusions": "\\begin{figure} \\includegraphics[width=3.5 in]{potential.eps}\\\\ \\caption{The effective potential $V^2$ for timelike geodesics for variable $L$ with $\\lambda=0$ and $M=1$. Solid line for $L>4$, dash line for $L=4$, dash-dot line for $L=2\\sqrt{3}$.} \\label{fig1} \\end{figure} Here, we turn to a discussion of the classification of the motion. First, we introduce the relativistic effective potential \\begin{equation} V^2=\\left(1-\\frac{2M}{r}-\\frac{\\lambda}{r^{3\\omega+1}}\\right)\\left(1+\\frac{L^2}{r^2}\\right).\\label{eq40} \\end{equation} Thus, the radial equation derived from (\\ref{eq14}) can be rewritten as $\\left(\\frac{dr}{d\\tau}\\right)^2 = E^2-V^2$, which implies that the possible motions only occur when $E^2 \\geqslant V^2$. Expand Eq. (\\ref{eq40}), \\begin{equation} V^2 = 1 - \\frac{2M}{r} + \\frac{L^2}{r^2} - \\frac{2ML^2}{r^3} - \\frac{\\lambda}{r^{3\\omega+1}} - \\frac{\\lambda L^2}{r^{3\\omega+3}}. \\end{equation} Because the effect of $\\lambda$ on $V^2$ is very small, the curves of $V^2$ are almost the same as that in the Schwarzschild space \\cite{trj16}. So we illustrate the effective potential $V^2$ for variable $L$ with $\\lambda=0$ in Fig. \\ref{fig1}. From Fig. \\ref{fig1}, we see that for $L>4$, there is a potential barrier whose peak value is larger than 1 when $r$ is very small, while there is also a potential well when $r$ is large. Thus, the possible motions can be classified as three kinds \\begin{eqnarray} E^2 < 1, \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\text{the bound state}; \\nonumber \\\\ 1 \\leqslant E^2 < V_\\textrm{max}^2, \\ \\ \\ \\ \\text{the scattering state}; \\nonumber \\\\ E^2 \\geqslant V_\\textrm{max}^2, \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\text{the absorbing state}. \\nonumber \\end{eqnarray} With decreasing $L$, the central potential barrier decreases in altitude. In the range of $2\\sqrt{3} \\leqslant L \\leqslant 4$, $V_\\textrm{max}^2 \\leqslant 1$, which indicates that the appearance of the scattering state is impossible, with the bound and absorbing states being left over. If $L < 2\\sqrt{3}$, both the peak and the hollow of the effective potential will disappear, and only the absorbing state is left over. But note that the above analysis is based on one assumption that the radius of the gravitation source is so small that the exterior of the source is applicable for very small $r$. In summary, we have studied the trajectories of test bodies on the motion of a binary system in the presence of quintessence. There are four cases of the state parameter $\\omega$ of quintessence, and we obtain four corresponding orbital equations, where it is assumed that the test mass moves round the center of force with eccentricity $e<1$. The difference of the azimuth angle of the pericenter can be cast in formulae in the two former cases, while the other two complicated ones can hardly be done, considering their different forms. Moreover, the effect caused by dark energy is too small to be detectable in the solar system. But the common feature of our four cases is that the resultant solutions reduce to the outcomes in Schwarzschild space naturally as the parameter $\\lambda$ vanishes." }, "1003/1003.0827_arXiv.txt": { "abstract": "{Clusters of galaxies are effective gravitational lenses able to magnify background galaxies and making it possible to probe the fainter part of the galaxy population. Submillimeter galaxies, which are believed to be star-forming galaxies at typical redshifts of 2 to 3, are a major contaminant to the extended Sunyaev-Zeldovich (SZ) signal of galaxy clusters. For a proper quantification of the SZ signal the contribution of submillimeter galaxies needs to be quantified.} {The aims of this study are to identify submillimeter sources in the field of the Bullet Cluster (1E~0657-56), a massive cluster of galaxies at $z\\simeq 0.3$, measure their flux densities at 870~$\\mu \\mathrm{m}$, and search for counterparts at other wavelengths to constrain their properties.} {We carried out deep observations of the submillimeter continuum emission at $870~\\mu$m using the Large APEX BOlometer CAmera (LABOCA) on the Atacama Pathfinder EXperiment (APEX) telescope. Several numerical techniques were used to quantify the noise properties of the data and extract sources.} {In total, seventeen sources were found. Thirteen of them lie in the central 10 arcminutes of the map, which has a pixel sensitivity of 1.2~mJy per $22''$ beam. After correction for flux boosting and gravitational lensing, the number counts are consistent with published submm measurements. Nine of the sources have infrared counterparts in Spitzer maps. The strongest submm detection coincides with a source previously reported at other wavelengths, at an estimated redshift $z\\simeq 2.7$. If the submm flux arises from two images of a galaxy magnified by a total factor of 75, as models have suggested, its intrinsic flux would be around 0.6~mJy, consistent with an intrinsic luminosity below $10^{12} L_\\odot$.} {} \\date{Received / Accepted 19/02/2010 } ", "introduction": "\\label{sec:introduction} \\begin{figure*}[t] \\centering \\includegraphics{13833f1_1.eps} \\includegraphics{13833f1_2.eps} \\caption{\\emph{Left}: layout of the bolometers on the LABOCA~array. Each circle represents one bolometer that was active at the time of the observations. The two rectangles illustrate the sizes of the two raster+spiral scanning pattern for one bolometer (solid line--small pattern, dashed line--large pattern). \\emph{Right}: movement pattern for one bolometer during the large scanning pattern. This pattern has a $2\\times 2$ raster setup with 100\\arcsec~between the raster points (the patterns are described in detail in Sect.~\\ref{sec:submillimeter-laboca}). The smaller scanning pattern is similar, but because the distance between raster points is considerably smaller, it is harder to illustrate. A comparison between the two patterns can instead be made from the two rectangles in the left panel. For clarity, we plot also the dashed rounded rectangle of the left panel in the righ panel.} \\label{fig:orbits} \\end{figure*} The large concentrations of mass (up to $10^{15} M_\\odot$) on angular scales of a few arcminutes in galaxy clusters act as natural gravitational lenses capable of magnifying background galaxies that would be too dim to be detectable otherwise. In the mm and submm wavebands, lensing by galaxy clusters makes it possible to probe the fainter part of the brightness distribution of the so-called submillimeter galaxies (SMGs), which are believed to be dusty high-redshift star-forming galaxies \\citep{1997MNRAS.290..553B}. % Pioneering observations of SMGs at 450 and $850~\\mu$m were done using SCUBA on the James Clerk Maxwell Telescope. The first observations toward two massive clusters at $z \\simeq 0.35$ resulted in the detection of a total of six sources above the noise level of 2 mJy/beam at $850~\\mu$m \\citep{SmailIvison:1997aa}. % The authors estimated the surface density of the sources to be three orders of magnitude larger than the expectation from a non-evolving model using the local IRAS $60~\\mu$m luminosity function. Those observations provided evidence for the presence of a large number of actively star-forming galaxies at high redshift, which might be the counterparts of the luminous and ultraluminous infrared galaxies observed in the local universe \\cite[e.g.][]{SandersMirabel:1996aa}. At redshifts beyond one, the flux density of a redshifted infrared-luminous galaxy is largely redshift-independent~: its decrease with an increasing distance is compensated by the steep rise in the mm and submm due to the redshifted spectral energy distribution (\\citealt{1993MNRAS.264..509B}). % During the last decade, several hundreds of SMGs have been discovered using bolometer arrays, mostly SCUBA \\cite[e.g.][]{Blain:1998aa, BorysChapman:2003aa, CoppinChapin:2006aa}, % and more recently MAMBO at 1.2~mm and AzTEC~at 1.1~mm \\citep{BertoldiCarilli:2007aa, ScottAustermann:2008aa,AustermannDunlop:2010aa}. The mm/submm galaxy population is the subject of many multi-wavelength studies (see the review by \\citealt{BlainSmail:2002aa}). The median redshift of SMGs with known redshifts is around $2-3$ \\citep{SmailIvison:2000aa}. So far, only a handful of massive galaxy clusters have been mapped in the submm, and most of them are clusters in the northern hemisphere observed with SCUBA. Observation of seven massive clusters with a sensitivity of 2 mJy/beam provided a catalogue of 17 submm sources brighter than the 50\\% completeness limit \\citep{1998ApJ...507L..21S}. % Nine other cluster fields in the redshift range 0.2 -- 0.8 were observed with a similar sensitivity, resulting in the detection of 17 new submm sources \\citep{2002MNRAS.330...92C}. % Deeper observations with a 3-sigma limit of 1.5--2 mJy/beam of three massive clusters probed the sub-mJy number counts, because of the gravitational magnification of the clusters. \\citep{CowieBarger:2002aa}. % \\cite{Knudsenvan-der-Werf:2008aa} targeted twelve clusters and the New Technology Telescope Deep Field. They detected 59 sources (some of them being multiple images of the same galaxy), and determined that seven of them have sub-mJy lensing-corrected flux densities. The LABOCA~bolometer camera on APEX has been used to survey the 870~$\\mu \\mathrm{m}$~emission in a protocluster at $z\\simeq 2.4$ \\citep{BeelenOmont:2008aa} and the Extended Chandra Deep Field South \\cite{WeisKovacs:2009ab}. % \\cite{2009A&A...506..623N} % observed the Sunyaev-Zeldovich (SZ) increment toward the massive cluster Abell~2163 and noted one bright point source. This is a good candidate for an SMG lensed by the cluster. The Bullet Cluster (1E~0657-56) at $z\\simeq 0.3$ is one of the most massive galaxy clusters known to date \\citep[see][]{Markevitch:2002lr,SpringelFarrar:2007aa}. A bright millimeter source was recently discovered in the Bullet Cluster field and identified as the lensed image of a background galaxy at a redshift of about 2.7 (\\citealt{2008MNRAS.390.1061W}, hereafter W08). The source happens to lie close to a critical line of the lens, causing a large flux amplification. The observations were performed with the AzTEC~bolometer camera on the 10-meter ASTE telescope in the Atacama desert in Chile, which provides an angular resolution of $30''$ at 1.1~mm wavelength. A doubly lensed source at the same location had been previously identified in Hubble Space Telescope ($HST$) maps and was used, together with other multiply lensed galaxies and a large number of weakly lensed sources, to obtain a calibrated map of the projected mass distribution of the Bullet Cluster \\citep{Bradac:2006fj}. Recently, \\cite{GonzalezClowe:2009aa} identified a third image by analyzing maps in the optical ($HST$), and in the near- and mid-infrared (Magellan and Spitzer). By fitting the Spectral Energy Distribution (SED) of a starburst galaxy to the observations, they also inferred a redshift of about 2.7. Their lensing model gave a magnification of 10--50 for the three images. In this paper we present results of observations of the Bullet Cluster field at a wavelength of 870~$\\mu \\mathrm{m}$~using the LABOCA~bolometer camera. At that wavelength, the emission is a combination of extended signal due to the Sunyaev--Zeldovich~effect by the hot intracluster gas and of point sources, which are potential high-redshift star-forming galaxies. Recovery of the extended SZ signal from the LABOCA~data (the SZ increment) requires a different data reduction and will be presented in a subsequent paper. The SZ~decrement from the Bullet Cluster~has been mapped by e.g. \\cite{2009ApJ...701...42H} with the APEX-SZ instrument, operating at 2~mm. This paper is organized as follows: the observations are presented in Sect.~\\ref{sec:observations} and the data reduction in Sect.~\\ref{sec:data-reduction}; the results are presented and discussed in Sects.~4 and \\ref{sec:analysis}. Throughout the paper, we adopt the following cosmological parameters: a Hubble constant $H_0 = 70$~km~s$^{-1}$~Mpc$^{-1}$, a matter density parameter $\\Omega_0 = 0.3$, and a dark energy density parameter $\\Omega_{\\Lambda0} = 0.7$. The redshift $z=0.296$ of the Bullet Cluster corresponds to an angular-diameter distance of 910~Mpc and a scale of 4.41~kpc/arcsec. ", "conclusions": "Continuum observations at 870~$\\mu \\mathrm{m}$~of the Bullet Cluster have been presented. The data were filtered to remove large-scale signal such as the Sunyaev--Zeldovich~increment from the cluster. The main results are summarized below: \\begin{itemize} \\item Seventeen submm sources with signal-to-noise ratios larger than 4 were detected in the map. Their measured fluxes densities range from 4.6 to 48~mJy. For each source but one, which lies in a noisy area, we calculated the value of the flux density corrected for the flux boosting due to confusion noise and for the gravitational magnification by the cluster. \\item The brightest submm source coincides with a previously reported galaxy at an estimated redshift of 2.7--2.9 detected by Spitzer and the AzTEC~and BLAST experiments. With its flux density of 48~mJy it is one of the brightest submm galaxy ever detected. After correction for gravitational magnification $|\\mu|$, the intrinsic flux of the source is about $0.64 \\, (|\\mu|/75)^{-1}$~mJy. \\item We found reliable infrared counterparts for nine of the submm sources. An infrared color-color analysis suggests that they have starburst-dominated spectral energy distributions. \\item The cumulative number counts derived from the observations agree well with those from other surveys. \\item The observations resolve 14\\% of the cosmic infrared background radiation at 870~$\\mu \\mathrm{m}$, in a sky area of $\\sim 78.5\\, \\mathrm{arcmin}^2$. \\end{itemize} We plan to apply a similar analysis to several other cluster fields for which we have LABOCA~data. The results presented here will also be used to remove the contribution of the submm sources and obtain a map of the SZ increment at 870~$\\mu \\mathrm{m}$~in the Bullet Cluster~with sub-arcminute angular resolution. Although the SZ decrement of the Bullet Cluster has been observed using several instruments (e.g., SEST by \\citealt{Andreani:1999bl}, APEX-SZ by \\citealt{2009ApJ...701...42H}, ACT by \\citealt{HincksAcquaviva:2009aa}, and recently the South Pole Telescope by \\citealt{PlaggeBenson:2009aa}) the only observation of the increment was done using ACBAR at 275~GHz, with an angular resolution of 4.5\\arcmin~\\citep{Gomez:2004dz}. Removal of lensed background sources is not a trivial task, since it can also bias the SZ measurement \\citep{LoebRefregier:1997aa}. In the case of the Bullet Cluster, the extremely bright lensed submm source close to the center of the cluster field has an integrated flux density comparable to that of the SZ flux from the central region of the cluster. A careful analysis must be performed to recover the SZ increment in such a system, and may require a joint analysis using measurements in other wavebands." }, "1003/1003.0016_arXiv.txt": { "abstract": "\\noindent We analyze the size evolution of HII regions around 27 quasars between $z=5.7$ to 6.4 ('quasar near-zones' or NZ). We include more sources than previous studies, and we use more accurate redshifts for the host galaxies, with 8 CO molecular line redshifts and 9 MgII redshifts. We confirm the trend for an increase in NZ size with decreasing redshift, with the luminosity normalized proper size evolving as: $\\rm R_{NZ,corrected} = (7.4 \\pm 0.3) - (8.0 \\pm 1.1) \\times (z-6)$ Mpc. While derivation of the absolute neutral fraction remains difficult with this technique, the evolution of the NZ sizes suggests a decrease in the neutral fraction of intergalactic hydrogen by a factor $\\sim 9.4$ from $z=6.4$ to 5.7, in its simplest interpretation. Alternatively, recent numerical simulations suggest that this rapid increase in near-zone size from $z=6.4$ to 5.7 is due to the rapid increase in the background photo-ionization rate at the end of the percolation or overlap phase, when the average mean free path of ionizing photons increases dramatically. In either case, the results are consistent with the idea that $z \\sim 6$ to 7 corresponds to the tail end of cosmic reionization. The scatter in the normalized NZ sizes is larger than expected simply from measurement errors, and likely reflects intrinsic differences in the quasars or their environments. We find that the near-zone sizes increase with quasar UV luminosity, as expected for photo-ionization dominated by quasar radiation. ", "introduction": "Observations have set the first constraints on the epoch of reionization (EoR), corresponding to the formation epoch of the first luminous objects (see the review by Fan, Carilli, \\& Keating 2006a). Studies of Gunn-Peterson (GP) absorption in quasar spectra (Gunn \\& Peterson 1965), and related phenomena, suggest a qualitative change in the state of the intergalactic medium (IGM) at $z \\sim 6$, indicating a rapid increase in the neutral fraction of the IGM, from $\\rm f(HI) < 10^{-4}$ at $z \\le 5.5$, to $\\rm f(HI) > 10^{-3}$ (by volume) at $z \\ge 6$. Conversely, the large scale polarization of the cosmic microwave background (CMB) is consistent with instantaneous reionization at $z \\sim 10.4 \\pm 1.2$ (Komatsu et al. 2010). These data suggest that reionization is less an event than a process, beginning at $z \\ge 10$, and with the `percolation', or 'bubble overlap' phase ending at $z \\sim 6$ to 7. Unfortunately, current methods for measuring the neutral fraction of the IGM have fundamental limitations, as GP absorption saturates at low neutral fractions ($\\rm f(HI) > 10^{-3}$), and CMB polarization measurements are an integral measure of the Thompson scattering optical depth back to $z = 1000$. Indeed, Mesinger (2010) has pointed out that there is no definitive evidence for (or against) a largely ionized IGM, even as late as $z \\sim 6$, while recent measurements of the temperature of the high $z$ Ly$\\alpha$ forest suggest fairly late reionization ($z < 9$; Bolton et al. 2010). Luminous quasars at the end of cosmic reionization will generate the largest ionized regions in the Universe during this epoch. These have been alternatively called Cosmological Str\\\"omgren Spheres or quasar near-zones (NZ), and we adopt the latter to be consistent with the recent literature. The near-zone behavior has been described by numerous authors (Shapiro \\& Giroux 1987; Madau \\& Rees 2000; Cen \\& Haiman 2000). A straight-forward calculation (Haiman 2002) shows that, when recombination is not important, and when the quasar lifetime is much less than the age of the Universe, the physical size of the expected NZ behaves as: $$ \\rm R_{NZ} = 8.0~ f(HI)^{-1/3}(\\dot{N_Q}/6.5\\times 10^{57} {\\rm s^{-1}})^{1/3} (t_Q/2\\times 10^7 {\\rm yr})^{1/3} [(1+{\\it z}_Q)/7]^{-1} \\rm Mpc $$ \\noindent where $\\rm R_{NZ}$ is the NZ proper radius, $\\dot{\\rm N_Q}$ is the rate of ionizing photons from the quasar, $z_Q$ is the quasar host galaxy redshift\\footnote{Note that we make the reasonable assumption that the quasar is at rest relative to its host galaxy.}, and $\\rm t_Q$ is the quasar age. This calculation ignores clumping, and assumes the mean cosmic density over the large scales being considered. Spectra of $z \\sim 6$ quasars show evidence for large ionized regions around bright quasars on proper scales $\\sim 5$ Mpc. The evidence comes in the form of excess transmission on the blue wing of the Ly$\\alpha$ emission line, prior to the onset of full GP absorption. This excess emission has been interpreted as being due to the local ionizing effect of the quasar. A number of studies have attempted to use these NZ to derive constraints on the IGM neutral fraction, using the equation above and assuming mean quasar lifetimes derived from demographics (White et al. 2003, 2005; Mesinger \\& Haiman 2004; Wyithe, Loeb, \\& Carilli 2005; Yu \\& Lu 2005; Gnedin \\& Prada 2004; Furlanetto, Hernquist \\& Zaldariaga~2004; Wyithe \\& Loeb~2004; Walter et al. 2003). The most detailed treatment of quasar near-zones associated with $z \\sim 6$ quasars to date is that by Fan et al. (2006b). They point out numerous factors that make the absolute measurement of the neutral fraction using NZ radii difficult, including: (i) uncertainties in the mean quasar lifetimes and ionizing photon rates, (ii) large scale structure in biased regions and pre-ionization by local galaxies, and (iii) clumpiness in the IGM (see also Lidz et al. 2007; Maselli et al. 2007; Bolton \\& Heahnelt 2007a,b). However, Fan et al. also point out that over the narrow redshift range being considered, the systematics are likely to be comparable for all sources, and hence, `if there is an order of magnitude evolution in the IGM ionization, the size of the HII regions should show strong evolution, providing a reliable {\\sl relative} measurement of the neutral fraction (with redshift).' More recently, Wyithe et al. (2008) and Bolton et al. (2010) have argued, based on numerical simulations, that the dominant effect on quasar near-zones during the end of reionization, or the percolation stage, when the overall neutral fraction is low, may be a rapidly increasing contribution to photo-ionization from the cosmic ionizing background with cosmic time, acting in concert with the quasar photons to lift the GP trough near the edge of the near-zone. Fan et al. (2006b) emphasize that the measurements of the NZ sizes are uncertain due to a number of factors, including the determination of where the transmission drops to zero, ie. the on-set of the GP effect ($z_{GP}$), and the redshift of the host galaxy ($z_Q$). To mitigate the former, they define the ionization zone, $\\rm R_{NZ}$, as the region in which the transmitted flux ratio (relative to the extrapolated continuum) is above 0.1 when smoothed to a resolution of 20 $\\AA$. This is well above the average GP transmission for quasars at this redshift ($< 0.04$), and hence mitigates affects of a fluctuating IGM, while still providing a good relative measurement of NZ sizes over the narrow redshift range considered. Willott et al. (2008) have considered this technique for two $z \\sim 6$ quasars, and point out that the transmission can rise above 0.1 just one or two resolution elements further from the quasar, raising the question of where to draw the line for $\\rm R_{NZ}$? In this paper, we use the strict and quantifiable definition of Fan et al. (2006b), namely, the first redshift below the quasar redshift at which the transmission drops under 0.1, in order to be consistent with previous literature. We refer the reader to the extensive discussion in Fan et al. 2006b and Willott et al. (2008) for details. One major uncertainty in the Fan et al. (2006b) analysis was the host galaxy redshift, $z_Q$. In particular, most of the host galaxy redshifts in the Fan et al. sample were based on a combination of Ly$\\alpha$+NV and high-ionization lines such as CIV and Si IV. It is well documented that broad Ly$\\alpha$ emission is a poor indicator of the host galaxy redshift due to absorption on the blue-side of the emission line, while the high ionization lines can have large velocity offsets, of order 1000 km s$^{-1}$ ($\\Delta z \\sim 0.02$ at $z = 6$), with respect to low ionization lines such as MgII (Richards et al. 2002). We have undertaken an extensive study of the radio through submm properties of $z > 5.7$ quasars to probe the dust, molecular gas, and star formation properties of the host galaxies (Wang et al. 2010; Wang et al. 2008; Wang et al. 2007; Carilli et al. 2007; Walter et al. 2009; Walter et al. 2004; Bertoldi et al. 2003; Bertoldi et al. 2003). One key result from our program is the detection of CO emission from the host galaxies of 8 quasars (Wang et al. 2010). High resolution imaging shows that the CO is centered within 0.2\" of the quasar, at least in one source (Walter et al. 2004; Riechers et al. 2009), and CO almost certainly provides the host galaxy redshift to very high accuracy ($\\Delta z < 0.002$). In parallel, we have been using near-IR spectra to detect MgII emission from $z \\sim 6$ quasars (Kurk et al. 2007; Jiang et al. 2007), which also provides much more accurate host galaxy redshifts ($\\Delta z \\le 0.007$). ", "conclusions": "Using a larger sample and improved quasar host redshifts, we have confirmed a statistically significant increase in luminosity normalized NZ sizes from $z \\sim 6.4$ to 5.7, first noted by Fan et al. (2006b). In the simplest physical model, the factor 2.3 change in mean NZ radius is consistent with a factor of 9.4 decrease in mean neutral fraction from $z \\sim 6.4$ to 5.7. Fan et al. (2006b) derive a volume averaged neutral fraction of $9.3\\times 10^{-5}$ at $z = 5.7$ from their analysis of GP absorption troughs and gaps. The NZ measurements then suggest that the neutral fraction has increased to $\\sim 9\\times 10^{-4}$ at $z = 6.4$. Wyithe et al. (2008) and Bolton et al. (2010) have proposed an alternative interpretation for this evolution in NZ sizes. Based on numerical simulations, they argue that the dominant effect on quasar near-zones during the end of reionization, or the percolation stage, may be a rapidly increasing contribution to photo-ionization from the cosmic ionizing background with cosmic time. The increasing mean free path allows the background photons to contribute substantially to the ionization around the edge of the NZ, thereby effectively increasing the observed sizes of the NZ. In either case, the results are consistent with the idea that $z \\sim 6$ to 7 corresponds to the tail end of cosmic reionization. Our improved $z_Q$ measurements imply that the intrinsic scatter in the NZ sizes is large at any given redshift: $\\rm \\sigma_{NZ} = 2$ Mpc. This large scatter must reflect real physical differences between quasars, perhaps some combination of differences in SEDs, UV escape fractions, ages, and/or local environments. Bolton \\& Haehnelt (2007b) have argued that the quasar age is less a factor when f(HI) $< 0.1$, since the size is then most likely dictated by absorption (and recombination) in denser regions of residual neutral gas within the sphere itself. The small NZ sizes seen for the BAL quasars in Table 1 suggest a dependence on the immediate local environment, although see the cautionary note in Section 2 concerning these sources. We also find an increase in $\\rm R_{p,NZ}$ with quasar UV luminosity. This correlation is qualitatively consistent with photo-ionization dominated by quasar radiation. Lastly, we find no correlation between $\\rm M_{1450}$ and redshift over the relatively narrow redshift range explored in this sample. We re-emphasize that the use of the near-zone technique has major pit-falls when trying to estimate the absolute neutral fraction at a given redshift, as detailed in Fan et al. (2006b) and Bolton \\& Heahnelt (2007a,b). In this paper, we have improved substantially one of the observational uncertainties in the calculation, namely, the quasar host galaxy redshifts, and we have focused on the evolution of the NZ sizes and neutral fraction, not the absolute neutral fraction. Still, as Bolton \\& Heahnelt point out, a larger sample with accurate measurements of the spectral regions around both Ly$\\alpha$ and Ly$\\beta$ is necessary to set meaningful constraints on the absolute neutral fraction of the IGM using the NZ technique." }, "1003/1003.5319_arXiv.txt": { "abstract": "{}{The relation connecting the emitted isotropic energy and the rest-frame peak energy of the $\\nu$F$\\nu$ spectra of Gamma-Ray Bursts (the Amati relation), strictly depends on the cosmological model, so we need a method to obtain an independent calibration of it.} {Using the Union Supernovae Ia catalog, we obtain a cosmographic luminosity distance in the $y$-redshift and demonstrate that this parametrization approximates very well the fiducial standard comsomlogical model $\\Lambda$CDM. Furthermore, by this cosmographic luminosity distance $d_l$, it is possible to achieve the Amati relation independent on the cosmological model} {The cosmographic Amati relation that we obtain agrees, in the errors, with other cosmological-independent calibrations proposed in the literature.}{ This could be considered a good indication in view to obtain standard candles by Gamma-Ray Bursts} ", "introduction": "Supernovae $Ia$ (SNeIa) are considered accurate and reliable standard candles, \\citep{Phillips}. In recent years, their use as cosmological distance indicators have led to the puzzling discovery that the Universe is in a phase of accelerated expansion, \\citep{Riess,Perlmutter}. This feature has also led to the revision of the standard cosmological model, leading to what is known today as the $\\Lambda$CDM concordance model, see e.g. \\citep{Ostriker}. However it is not possible to observe these objects very far in the Universe. The most distant Supernova $Ia$ was observed at a redshift of z $\\sim$ 1.7 \\citep{Benitez}. For this reason, several cosmological analysis made by using the various compiled sample of SNeIa, like the Union Catalog, \\citep{Kowalski}, are not able to investigate higher redshift regions of the Universe. If we had distance indicators at higher redshifts, then we could extend our knowledge at these unexplored regions. One of the possible solutions to this problem could come from the Gamma-Ray Bursts (GRBs) assumed as cosmological indicators \\citep{Piran,Meszaros}. GRBs are the most powerful explosions in the Universe: this feature allows them to be observed at extremely high redshift. The most distant GRB observed up to now is at a redshift of $\\sim$ 8.1, \\citep{Tanvir,Salvaterra}. However, GRBs are not standard candles, since they have no known and well-defined luminosity relation. Due to this lack, we have to find another way to use GRBs as cosmological beacons. A possible solution could consist in finding correlations between photometric and/or spectroscopic properties of GRBs themselves. In the scientific literature there are several of these relations, \\citep{Schaefer}. One of these is the Amati relation, \\citep{Amati}, which relates the isotropic energy emitted by a GRB with the peak energy in the rest-frame of the $\\nu$F($\\nu$) electromagnetic spectrum of a GRB. This relation has already been widely used to constraining the cosmological density parameter \\citep{Amati2}, with quite remarkable results. However, there is still not a physical link between this correlation and the mechanisms underlying the production and the emission of a GRB. The basic emission process of a GRB is very likely not unique, so it is not easy to explain, from a physical point of view, such a relation. Recently it has been suggested that the Amati relation could depend strongly on the satellite measurements used for detection and the observation of each GRB \\citep{Butler07}. However this hypothesis has been rejected recently, \\citep{Amati2}, since the relation seems to be verified regardless of the satellite considered for the observations and detection. Although not supported by self-consistent physical motivations, it is a phenomenological relation which could be extremely useful for cosmological considerations. However, a problem related with such a relation is that it must be calibrated independently of the considered cosmological model. In order to compute the energy emitted from an astrophysical object at a certain redshift $z$, we need, as a matter of fact, a measurement of the bolometric flux and the distance of the same object. For the first quantity, we follow the idea outlined by \\citep{Schaefer} : one can obtain a very precise measurement of the bolometric fluence emitted by a GRB from the observed fluence, the integrated flux in the observation time and the spectral model that best fits the spectral energy distribution of each GRB. However, the distance depends on the considered cosmological model. People usually adopt the standard $\\Lambda$CDM model, with fixed values of the density parameter $\\Omega_i$. This procedure leads to the so-called {\\it circularity problem} when the Amati relation is used to standardize GRBs. For this reason we need a cosmology-independent calibration of the relation. Recently, it was released a calibration with SNeIa data by using different numerical interpolation methods \\citep{Liang}; the results seem very reliable to address cosmological issues by GRBs. In this work we shall take into account a similar analysis: by taking into account SNeIa data from the cosmographic point of view (for a detailed description see e.g. \\citep{Weinberg,Visser}), it could be possible to obtain a calibration of the Amati relation. We will use results obtained from a cosmographic fit of a sample of SNeIa extended up to very high redshift with the GRBs. The use of the cosmography to deduce the cosmological parameters from SNeIa was widely discussed in the literature, \\citep{Visser2}, and the results are very close to that attained by other and more accurate analysis. Recently applications of cosmographic methods have taken into account galaxy clusters \\citep{Capozz} and GRB, \\citep{Izzo,Vitagliano} but their reliability drastically fails at high redshifts. Indeed, the estimates of the deceleration parameter $q_0$ and of the jerk parameter $j_0$ are usually achieved only at very low redshift and then any extrapolation could led to shortcomings and misleading results as soon as they are extended. However by an appropriate parameterization of the redshift parameter, one can circumvent the problem introducing a new redshift variable ranging from 0 and 1 \\citep{Visser}. Let us consider the following quantity as the new redshift variable: \\begin{equation} y = \\frac{z}{1+z}\\,, \\end{equation} we obtain that the range of variation is between 0 and 1. In this way, we can derive a luminosity distance by which we can obtain the Amati relation suitable for cosmography. The layout of the paper is the following: in Sect. 2 we tackle the cosmographic analysis considering the SNeIa Union sample. Results will be used to derive the luminosity distance for each GRB and then we will fit the cosmographic Amati relation. In Sect. 3 a discussion on how to extend the same relation is reported. We add further 13 GRBs (as of December 2009), computing the bolometric fluence and the peak energy for each of them and after we calculate the cosmographic parameters using the new relation. Finally, we calculate the isotropic energy for each GRB and then compute the best fit for the considered sample of data (Sect. 4). Discussion and conclusions are reported in Sec. 5. ", "conclusions": "The issue to extend the cosmic scale ladder up to medium-high redshift is an important questions of modern cosmology. A possible way to achieve this goal is to take into account GRBs, the most powerful explosions in the Universe. The energy emitted by these objects spans about six orders of magnitude. However, they cannot be assumed as standard candles in a proper sense. Dispite of this lack, the existence of several correlations between spectroscopic and photometric observable quantities of GRBs allow us to solve in part this problem. The fundamental pre-requisite to obtain such relations is to estimate the emitted energy in a way independent of the cosmological model. In this paper, we have considered a relation for the luminosity distance $d_l$ that is independent on the dynamics of the Universe, but, in principle, could work only at small redshift. Although we have use a parameterization for the redshift which allows to transform the variable $z$ in a new variable $y$, ranging in a limited interval, we have seen that the obtained luminosity distance at high-redshift differs slightly from the fiducial model $ \\Lambda$CDM at high redshift, see Fig.\\ref{fig:no3}. Nevertheless, since we obtained the curve $ d_l(y) $ by an analysis of the SNeIa Union survey, that extends up to a redshift of $ \\sim $ 1.7, we achieved an independent estimate at slightly higher redshift\\footnote{ Estimates of the Baryonic Acoustic Oscillations (BAO) performed by forthcoming surveys of clusters at intermediate redshift ($z \\approx$ 2.5 - 3.5) will give better approximations of the curve $ d_l(y) $.}. By the way using the $ d_l(y) $ obtained with the cosmographic fit of the SNeIa, we have constrained a sample of GRBs in a cosmology-independent way so that we have fitted a cosmographic Amati relation for GRBs. The results are similar to those obtained from other analysis performed using other methods, \\citep{Schaefer,Liang,Amati}. It is important to stress the independence from cosmology and the calibration obtained by SNeIa. In our opinion, this characteristic is relevant, from one side, to constrain cosmological models, in particular, dark energy models, and, from another side, to check the physical validity of the Amati relation. \\paragraph{Acknowledgements.} We warmly thank L. Amati for providing us the GRB sample in \\citep{Amati09}. LI warmly thanks also R. Benini for useful discussion and help with the Mathematica package data analysis." }, "1003/1003.3015_arXiv.txt": { "abstract": "We investigate the local supermassive black hole (SMBH) density function and relative mass accretion rates of all active galactic nuclei (AGNs) identified in a volume-limited sample of infrared (IR) bright galaxies ($\\Lir > 3 \\times 10^9 \\Lsun$) to $D<15$~Mpc (Goulding \\& Alexander 2009). A database of accurate SMBH mass ($\\Mbh$) estimates is compiled from literature sources using physically motivated AGN modeling techniques (reverberation mapping, maser mapping and gas kinematics) and well-established indirect $\\Mbh$ estimation methods (the M-$\\sigma_*$ and $\\Mbh$--$L_{\\rm K,bul}$ relations). For the three sources without previously published $\\Mbh$ estimates, we use 2MASS $K$-band imaging and {\\sc galfit} to constrain the bulge luminosities, and hence SMBH masses. In general, we find the AGNs in the sample host SMBHs which are spread over a wide mass range ($\\Mbh \\approx (0.1$--$30) \\times 10^7 \\Msun$), but with the majority in the poorly studied $\\Mbh \\approx 10^6$--$ 10^7 \\Msun$ region. Using sensitive hard X-ray (2--10 keV) and mid-IR constraints we calculate the bolometric luminosities of the AGNs ($L_{\\rm Bol,AGN}$) and use them to estimate relative mass accretion rates. We use these data to calculate the volume-average SMBH growth rate of galaxies in the local Universe and find that the AGNs hosting SMBHs in the mass range $\\Mbh \\approx 10^6$--$10^7 \\Msun$ are dominated by optically unidentified AGNs. These relatively small SMBHs are acquiring a significant proportion of their mass in the present-day, and are amongst the most rapidly growing in the local Universe (SMBH mass doubling times of $\\approx$~6~Gyrs). Additionally, we find tentative evidence for an increasing volume-weighted AGN fraction with decreasing SMBH mass in the $\\Mbh \\approx 10^6$--$10^8 \\Msun$ range. Overall, we conclude that significant mass accretion onto small SMBHs may be missed in even the most sensitive optical surveys due to absent or weak optical AGN signatures. ", "introduction": "It is now well established that all massive galaxies ($M_* \\approx 10^{10}$--$10^{12} \\Msun$) in the local Universe harbour central super-massive black holes (SMBHs), with masses proportional to those of their stellar spheroids (hereafter, bulge; e.g.,\\ \\citealt{kormendy95}; \\citealt{magorrian98}). Comparisons between the SMBH mass density in the local Universe and the total energy produced by active galactic nuclei (AGNs) across cosmic time have shown that these SMBHs were primarily grown through mass-accretion events (e.g.,\\ \\citealt{soltan82}; \\citealt{rees84}; \\citealt{marconi04}). The space density of high-luminosity AGNs appears to have peaked at higher redshifts than lower-luminosity AGNs, suggesting that the most massive SMBHs ($\\Mbh \\approx 10^8$--$10^9 \\Msun$) grew first, a result commonly referred to as `AGN cosmic downsizing' (e.g.,\\ \\citealt{cowie03}; \\citealt{ueda03}; \\citealt{mclure04}; \\citealt{hasinger05}; \\citealt{alonso08}). Extrapolation of these results imply that the most rapidly growing SMBHs in the nearby Universe should be of comparatively low mass ($\\Mbh \\ll 10^{8} \\Msun$). To determine the characteristic masses of these growing SMBHs requires a complete census of AGN activity and SMBH masses in the local Universe. Using data from the Sloan Digital Sky Survey \\citep[SDSS; ][] {sdss_tech} in conjunction with the well established SMBH--stellar velocity dispersion relation (hereafter, M--$\\sigma_*$; e.g., \\citealt{gebhardt00}; \\citealt{tremaine02}), \\citeauthor{heckman04} (2004; hereafter, H04) deduced that relatively low mass SMBHs ($\\Mbh \\approx 3 \\times 10^7 \\Msun$) residing in moderately massive bulge-dominated galaxies host the majority of present-day accretion onto SMBHs. However, the space density of SMBHs derived from the M--$\\sigma_*$ relation in the optical survey of H04 was limited by the spectral resolution of the SDSS ($ \\sigma_* > 70 \\kmps $) to SMBHs of $\\Mbh \\goa 3 \\times 10^6 \\Msun$ (assuming the M--$\\sigma_*$ relation of Gebhardt et~al. 2000). Furthermore, due to attenuation of optical emission by dust, source selection and AGN classification at optical wavelengths will be biased against gas-rich, dust-obscured objects. These surveys are unlikely to include galaxies hosting the smallest bulges, and consequently the lowest mass SMBHs, and may therefore be missing a significant proportion of SMBH growth in the local Universe. Indeed, the nearby Scd galaxy, NGC~4945, hosting a low-mass SMBH ($\\Mbh \\approx 1.4 \\times 10^6 \\Msun$; \\citealt{greenhill97}) only displays evidence for AGN activity in X-ray \\citep{iwasawa93} and mid-infrared (mid-IR) observations \\citep{GA09}. By contrast, the AGN in NGC~4945 (accreting at $\\goa 30$ percent of the predicted Eddington limit; \\citealt{itoh08}) is completely hidden at optical wavelengths, and classified as a starburst galaxy. Clearly, using optical data alone, the intrinsic AGN properties of sources similar to NGC~4945 cannot be derived. While optical emission-line diagnostics alone cannot reliably characterise the properties of a non-negligable fraction of the AGN population, they are readily identified at obscuration independent wavelengths (e.g., X-ray; mid-IR). Hence, the identification of AGNs made at X-ray and mid-IR wavelengths complements traditional UV/optical methods to yield a more complete census of AGN activity. Indeed, using the high resolution mid-IR spectrograph on-board the NASA \\spitzer \\ Space Telescope (\\spitzer-IRS), \\citeauthor{GA09} (2009; GA09) found using the first complete volume-limited sample of all ($\\approx 94$ percent) local ($D<15$ Mpc) bolometrically luminous galaxies ($\\Lir > 3 \\times 10^9 \\Lsun$), that $\\approx 50$ percent of local AGNs are not identified in sensitive optical surveys. At least $30$ percent of these AGNs were previously identified as pure optical starburst galaxies, similar to NGC~4945 (i.e., not even otherwise known to be transition-type objects as defined by \\citealt{kauff03b}). Furthermore, $\\approx 30$ percent of the optically unidentified AGNs were found to reside in late-type spiral galaxies (Sc--Sd; e.g., similar to NGC~4945). Complimentary to this, from a heterogenous sample of Palomar galaxies, \\citet{sat07} and \\citet{sat08} have also concluded that optically unidentified AGNs exist in some late-type spiral galaxies. With the inclusion of these new optically unidentified AGNs, it is natural to ask, what are the masses of local active SMBHs, what are their Eddington ratios, and hence, how rapidly are active SMBHs growing in the local Universe? In this paper, we investigate the growth rates and space density of actively accreting SMBHs using the 17 AGNs identified in the volume-limited survey of GA09. Whilst the source statistics considered here are significantly smaller than those studies using the SDSS, this work compliments that of H04 by including a relatively large number (given the considered small volume) of optically unidentified AGNs (10) which would not be reliably identified or characterised in the SDSS survey.\\footnote{Following GA09, throughout this paper we define an optically unidentified AGN as an object which is not unambiguously identified as a Seyfert galaxy at optical wavelengths using solely traditional optical emission-line diagnostics. This has the advantage that we may directly compare statistics from the sample considered here to those derived from large {\\it N} surveys such as the SDSS.} Furthermore, by including a significant population of bolometrically luminous (but dust-obscured) late-type spiral galaxies (Sc--Sd) we are able to extend the SMBH density function to $\\Mbh < 3 \\times 10^6 \\Msun$. As many of the late-type spiral galaxies host small galactic bulges, and hence lower mass SMBHs, particular attention is paid to obtaining accurate mass estimates for these SMBHs. Given their proximity, many of the sources in GA09 are well-studied and have multiple estimates of SMBH mass ($\\Mbh$) from a variety of methods (i.e., reverberation mapping techniques; mapping of water maser spots; gas kinematical estimates; the M--$\\sigma_*$ relation; correlation of $\\Mbh$ with the luminosity of the galactic bulge); below we discuss the relative accuracy of each SMBH mass estimate technique. Furthermore, to determine the relative mass accretion rates and hence average growth times of the SMBHs in our sample we require the best available estimates of the AGN bolometric luminosity ($L_{\\rm Bol,AGN}$). Here we use two approaches: 1) for the AGNs with currently published data, we use high-quality well-constrained sensitive hard X-ray (2--10 keV) luminosities to directly measure $L_{\\rm Bol,AGN}$; and 2) we accurately infer $L_{\\rm Bol,AGN}$ using a well-constrained hard X-ray to high-ionisation mid-IR emission line relation. In \\S2 we outline the construction and basic reduction analysis of the AGN sample derived from GA09. In \\S3 we present the SMBH mass estimates. For a minority of objects (three out of 17 AGNs) without published $\\Mbh$ estimates we outline the use of a bulge/disc decomposition method with 2MASS {\\it K}-band images, and following Marconi \\& Hunt (2003), we use the $\\Mbh$--$\\Lbul$ relation to estimate their SMBH masses. In \\S4 we use hard (2--10 keV) X-ray measurements and high-ionisation mid-IR emission to estimate the intrinsic luminosity of the AGNs considered in our sample. Using our well-defined estimates for SMBH mass and AGN bolometric luminosity, we investigate the relative mass accretion rates of our sample of active SMBHs in \\S5. We use these estimates to provide new constraints on the volume-average SMBH growth rates in the local Universe. We further compare these results to the previous works of H04 and \\citet{gre_ho07} by producing a local AGN population density function. Finally, in \\S6 we present our conclusions. ", "conclusions": "From a volume-limited sample of 64 bolometrically luminous galaxies to $D<15$ Mpc ($\\approx 94$ percent complete) GA09 unambiguously identified seventeen ($\\approx 27^{+8}_{-6}$ per cent) sources to be hosting AGN activity in galaxies with $\\Lir > 3 \\times 10^9 \\Lsun$, using [NeV] $\\lambda 14.32 \\um$ emission as a robust AGN indicator. Using the SMBH mass and AGN bolometric luminosity estimates derived here, we discuss the relative mass accretion rates of these seventeen AGNs and use them to derive the average present-day growth times of SMBHs in the very nearby Universe. Furthermore, we evaluate the unique contribution that our new optically unidentified AGNs make to the space density of active SMBHs in the local Universe, which have until now been previously derived from large-scale optical surveys (e.g., H04; Greene \\& Ho 2007). \\subsection{Derived AGN Properties and Relative Mass Accretion Rates} In Fig.~\\ref{fig_2}, we plot $L_{\\rm Bol,AGN}$ against our adopted $\\Mbh$ estimates (the associated 1-$\\sigma$ errors for $\\Mbh$ measurements are described in section 3) for the 17 AGNs in our volume-limited sample. $L_{\\rm Bol,AGN}$ is inferred from either accurate intrinsic high-quality hard X-ray (2--10 keV) constraints (where available) or AGN-produced [OIV] $\\lambda 25.80 \\um$ emission (see \\S4.1 and 4.2). The $L_{\\rm Bol,AGN}$ 1-$\\sigma$ errors for the sources with hard X-ray constraints are the result of combining the uncertainty in the $L_{X,2-10keV}$ measurement with that of the mean spread in the bolometric correction factor employed from \\citet{marconi04}. For those AGNs with $L_{\\rm Bol,AGN}$ derived from $L_{\\rm [OIV]}$, the error is derived from the uncertainty in $L_{\\rm [OIV]}$ as quoted in GA09 combined in quadrature with the intrinsic scatter of the empirical [OIV]--$L_{\\rm bol,AGN}$ relation (equation 4). Seven objects within the sample have been classified as optical AGNs from previous surveys (see GA09 and Table~1) using typical optical emission-line diagnostics (e.g., the Baldwin-Phillips-Terlevich diagnostic diagrams; \\citeauthor{bpt} 1981); however, all have detected \\nev emission, and thus are unambiguously identified to host AGNs at mid-IR wavelengths (GA09).\\footnote{We note that NGC~3627 is ambiguously classified as T2/S2 from the optical spectroscopy presented in \\citet{ho97a}, and thus may host an active central source. However, \\citet{roberts01} suggest from its optical emission-line ratios that NGC~3627 is most likely a LINER/HII composite.} We find that with the exception of NGC~5128 (Centaurus A), our sample is dominated by AGNs with SMBHs in the mass range $\\Mbh \\approx (0.1$--$5) \\times 10^7 \\Msun$ (median of $\\Mbh \\approx 7 \\times 10^6 \\Msun$). Due to the irregular structure of one of the galaxies in the sample (NGC~5195), $\\Mbh$ is poorly determined; in Fig.~\\ref{fig_2} we plot $\\Mbh$ estimates from both the M-$\\sigma_*$ and $\\Mbh$--L$_{\\rm K,bul}$ relations (connected blue-dashed line). \\begin{figure} \\begin{center} \\includegraphics[width=1.05\\linewidth]{d15_growth_time.ps} \\caption{$\\Mbh$ is plotted against the characteristic mean mass doubling time ($t_{2M}$) of a SMBH in units of Hubble-time ($t_H$) for the $D<15$~Mpc AGNs. Growth time errors are calculated from the log-normal standard deviations of the sample. For comparison, the growth time function of H04 is shown (dashed line). We find good agreement with H04 over the region $\\Mbh \\approx 0.3$--$10 \\times 10^7 \\Msun$ and further extend the growth time constraints to lower SMBH masses ($\\Mbh \\approx 10^6 \\Msun$).} \\label{fig_4} \\end{center} \\end{figure} We find the AGNs in our sample are spread over a wide range of bolometric luminosities, $L_{\\rm Bol,AGN} \\approx 10^{40}$--$10^{45} \\ergps$. To assess the relative mass-accretion rates of the sample ($L_{\\rm Bol,AGN} / L_{\\rm Edd} \\sim \\eta$), we over-plot lines of constant Eddington ratios ($\\eta \\approx 10^{-3},10^{-1},1.0$; derived following \\citealt{rees84}) and their associated mass-doubling times ($t \\approx 30,0.3,0.03$~Gyrs, respectively). Given the large range in bolometric luminosities, it is not surprising that the AGNs in the sample are found to be accreting at rates covering over 5 orders of magnitude ($\\eta \\approx 10^{-5}$--1). With the exception of a few AGNs, the observed range in Eddington ratios is found to be roughly consistent with those found by H04 for active galaxies (solid contours in Fig.~\\ref{fig_2}). As our work is not limited by the spectral resolution of the SDSS (i.e., with a limit of $\\Mbh \\ga 3 \\times 10^6 \\Msun$), we show in Fig.~\\ref{fig_2} that significant accretion, $\\eta > 10^{-3}$ (i.e., radiatively efficient accretion systems; e.g., thin discs) occurs onto SMBHs with $\\Mbh \\approx (1$--$3) \\times 10^6 \\Msun$. The majority of these low-mass, rapidly-accreting SMBHs are hosted in late-type, disc-dominated spiral galaxies (Sc--Sd). By contrast, it is generally assumed that gas-rich late-type spirals are preferentially inactive galaxies and that a large bulge may be a necessary component for the existence of a SMBH, and thus a luminous AGN. Furthermore, of the four AGNs within the sample with SMBHs consistent with $\\Mbh \\approx 10^6 \\Msun$, we find that three sources are not identified as AGNs in sensitive optical surveys. This indicates that significant SMBH accretion may be missed by statistically-large optical surveys such as H04 even if the spectral resolution was sufficient to identify SMBHs down to $\\Mbh \\approx 10^6 \\Msun$. For the subset of our AGN sample which host SMBHs with $\\Mbh \\goa 3 \\times 10^6 \\Msun$, we find that many of the optically unidentified AGNs are accreting at relatively low Eddington ratios ($\\eta \\loa 10^{-3}$), and are unlikely to make a significant additional contribution to the present-day growth of SMBHs. However, these same AGNs may form part of a separate, underlying population of radiatively inefficient accretion systems such as advection dominated accretion flows (ADAFs; e.g., \\citealt{adaf}) or those which contain optically-thick slim-discs. Further spectral analysis of the X-ray data may distinguish between these particular accretion systems, but is beyond the scope of these analyses (see Goulding et~al. in preparation). \\subsection{The Present-Day Growth of SMBHs} \\begin{figure*} \\begin{center} \\includegraphics[width=0.95\\textwidth]{new_bh_mass_density_fn.ps} \\caption{(Upper panel) Comparison of volume-weighted space densities of active SMBHs in the local Universe, $\\Phi$ in units of number ${\\rm Mpc}^{-3} {\\rm \\ log \\ M}^{-1}_{{\\rm BH}}$. Mid-IR active SMBH function (filled squares; Goulding \\& Alexander 2009) is compared to the optically identified NL AGN function (dotted curve) of H04, the BL AGN function (filled circle) of Greene \\& Ho (2007), and total SMBH function (active+inactive galaxies; solid curve; Marconi et~al. 2004). Sample selection bias is analysed using a robust Monte-Carlo simulation (shaded region; see Appendix). (Lower panel) Ratio of mid-IR active SMBHs to the total local SMBH mass function. The total SMBH mass function is extrapolated by 0.3 dex to $\\Mbh < 10^6 \\Msun$. For comparison the volume-weighted AGN fraction of H04 is also shown (dashed line). We estimate a mean volume-weighted local AGN fraction of $\\approx 25^{+29}_{-14}$ percent over the range $\\Mbh \\approx (0.5$--$500) \\times 10^6 \\Msun$.} \\label{fig_3} \\end{center} \\end{figure*} Using the relative mass accretion rates estimated for our sample (Fig.~\\ref{fig_2}), we can infer the volume-averaged growth time of SMBHs in the local Universe. Assuming a mean Kerr spin parameter ($a$) for our sample of $a \\approx 0.67$ (e.g., \\citealt{treister06,hopkins07}), i.e., an accretion efficiency ($\\epsilon$) of $\\approx 0.1$, the characteristic mass doubling time ($t_{2M}$) of a SMBH accreting matter at the Eddington limit is $t_{2M} \\approx 30$~Myrs (Rees 1984). Under the further assumption that $a$, and hence $\\epsilon$, does not vary significantly for changes in $\\Mbh$ \\citep{king08}, we assess the present-day growth rate of SMBHs.\\footnote{We note that the spin variation and spin directionality of SMBHs in AGNs is currently an ongoing area of research, and a consensus between groups has yet to be reached for an average value of the Kerr spin parameter; for example see \\citet{brenneman06}, \\citet{king08} and \\citet{fabian09}.} Following H04, we calculate and extend to lower masses ($\\Mbh < 3 \\times 10^6 \\Msun$) the integrated growth of SMBHs. Growth time errors are calculated from the log-normal standard deviations of the sample. We note here that we also include the optically unidentified AGNs which would not be detected in the SDSS. In Fig.~\\ref{fig_4}, we find that the mean growth time for low-mass SMBHs ($\\Mbh \\approx 10^6 \\Msun$) is $\\approx 6^{+6}_{-3}$~Gyrs, which is consistent with these AGNs growing on time-scales similar to that of the age of the Universe. Our results are found to be broadly consistent with a simple extrapolation of the growth times calculated by H04 to $\\Mbh \\approx 10^6 \\Msun$ (dashed-line in Fig.~\\ref{fig_4}). Thus, the AGNs hosting SMBHs in the mass range $\\Mbh \\approx 10^6$--$10^7 \\Msun$, which are dominated by optically unidentified AGNs (see Fig.~\\ref{fig_2}), are acquiring a significant proportion of their mass in the present-day, and are amongst the most rapidly growing SMBHs in the local Universe. Furthermore, we find our derived growth times of SMBHs with $\\Mbh \\goa 3 \\times 10^6 \\Msun$ are in good agreement with those presented in H04, with mean growth times of $t_{2M} \\approx 47^{+29}_{-18}$ and $\\approx 198^{+195}_{-90}$~Gyrs for AGNs with $\\Mbh \\approx (0.3$--3) and (3--$30) \\times 10^7 \\Msun$, respectively.\\footnote{We note that our sample contains only one galaxy with $\\Mbh \\goa 10^8 \\Msun$ (NGC~5128) and thus may not be representative for high $\\Mbh$ systems.} \\subsection{Space-Density of AGNs in the local Universe} An accurate active SMBH mass function, especially for lower mass SMBHs ($\\Mbh \\approx 10^6 \\Msun$), is crucial for extending our understanding of the role played by accretion in the growth of SMBHs across cosmic time. In this section we calculate the space density of active SMBHs for our sample and compare it to complimentary optical studies of local narrow-line (NL; H04) and broad-line (BL; Greene \\& Ho 2007) AGNs, and the total mass function of local SMBHs by Marconi et~al. (2004). Following Greene \\& Ho (2007), in the top panel of Fig.~\\ref{fig_3} we plot the volume-weighted space density, $\\Phi$ against $\\Mbh$ in mass bins of 0.5 dex. The volume, $V$, encompassed by the GA09 sample to $D<15$~Mpc is $V \\approx 1.3 \\times 10^4$~Mpc$^3$.\\footnote{The RBGS includes all IR-bright galaxies detected by {\\it IRAS} with $f_{60 \\mu m}>5.24$ Jy at $|b|> 5 \\degr$.} As we note in \\S2, given the luminosity limit imposed in our volume-limited survey ($L_{\\rm IR} \\goa 3 \\times 10^9 \\Lsun$), we do not include the dwarf like systems which are likely to host the very smallest SMBHs ($\\Mbh < 10^5 \\Msun$). With an adjustment for the distance model adopted in this work, an examination of the Palomar survey \\citep{ho97a} shows there are possibly three optical Seyferts (NGC~185; NGC~1058; NGC~4395) with $\\Mbh < 10^5 \\Msun$ which are not included in our sample due to our lower luminosity limit. However, we also note that NGC~185 has since been re-classified as an HII galaxy \\citep{ho01b}. The derived volume-weighted space density for our active SMBHs (filled squares), which is dominated by NL-AGNs, is found to be significantly greater (a factor of $\\approx 100$) than the SMBH density of BL-AGNs (filled circles) presented in Greene \\& Ho (2007) in the mass region $\\Mbh \\approx (0.9$--$90) \\times 10^{6} \\Msun$. The significant increase in active SMBH density when compared to the BL-AGN density of Greene \\& Ho (2007) is partially to be expected due to the greater relative sensitivity of our {\\it Spitzer}-IRS observations coupled with the greater abundance of observed Seyfert 2 to Seyfert 1 galaxies identified in the local Universe. However, this still may not be a good indicator of the intrinsic Seyfert~1:Seyfert~2 ratio. \\citet{tommasin10} find in a large sample of 81 Seyfert galaxies that only six sources do not contain significant \\nev emission in their mid-IR spectroscopy, the majority of which are Seyfert~1s. It is likely that the identification of low equivalent-width emission lines (such as [NeV] or [OIV]) in BL-AGNs is further complicated by a strong IR continuum emission which dominates the mid-IR regime. Hence, it is possible that by requiring the detection of \\nev to infer AGN status, we may be rejecting broad-line objects, and thus finding a lower Seyfert~1:Seyfert~2 ratio than is representative in the local Universe. In comparison to the active SMBH mass function containing the optically identified NL-AGNs of H04 (dotted line), we also find a significantly larger space-density of SMBHs. We find that the space-density of active SMBHs identified in the mid-IR is roughly constant in the mass region $\\Mbh \\approx (0.9$--$90) \\times 10^{6} \\Msun$ with a value of $\\Phi \\approx 6.3 \\times 10^{-4} {\\rm \\ Mpc}^{-3} {\\rm \\ log M}^{-1}_{\\rm BH}$. This space-density of AGNs is a factor of $\\approx 10$ greater than that estimated by H04 over the same $\\Mbh$ range. Since we find only two of the 17 AGNs in our sample are sufficiently luminous/unobscured to be detected in the SDSS survey, we determine that this is consistent with our results. We further suggest that the space density derived here may still be a lower-limit for the number of NL-AGNs in the local Universe. A further examination of the (distance-model adjusted) Palomar survey suggests that at least four further NL-AGNs are not included in our volume-limited survey. Of these, two (NGC~3486; NGC~4565) lack high-resolution {\\it Spitzer}-IRS spectroscopy of the central region (as noted in Table 2 of GA09), one (NGC~3031; $\\Lir \\approx 2.8 \\times 10^9 \\Lsun$) lies fractionally below our luminosity limit for this survey, and NGC~4258 is not included in the RBGS due to its extremely large angular size. Whilst the main focus of this paper is to compare SMBH statistics derived from mid-IR and optical detection techniques, it is prudent to note that the majority of AGN space densities calculated at higher redshifts are typically derived from sources detected in wide-field X-ray surveys. Hence, we now establish whether our mid-IR active space density may be missing a significant fraction of X-ray detected AGNs. Recently, a comparison between high-resolution {\\it Spitzer}-IRS spectroscopy and X-ray detected AGNs, was made by \\citet{dudik09} for a large sample of optically classified LINERS. Dudik et~al. (2009) reported inconsistencies between [NeV] non-detections and the presence of hard X-ray nuclear emission in a subset of their sample. However, they conclude that the limited sensitivity of their mid-IR observations may be driving their observed result. Using the relations between high-ionisation mid-IR emission and hard X-ray luminosity (M08; GA09; Goulding et al. in preparation), we suggest that a detected [NeV] luminosity of $L_{\\rm [NeV]} \\approx 10^{38} \\ergps$ in an AGN (i.e., the limiting luminosity in Dudik et~al. 2009) would be equivalent to a hard X-ray luminosity of $L_{\\rm X,2-10keV} \\approx 5 \\times 10^{40} \\ergps$; indeed, almost all of the X-ray detected AGNs which lack significant [NeV] emission in Dudik et~al. (2009) are below this threshold. We thus conclude that with sufficiently sensitive high-resolution mid-IR spectroscopy, there is currently no conclusive evidence to suggest that X-ray detected Type-2 AGNs lack significant [NeV] emission in the mid-IR. Hence, with the exception of some Seyfert 1 galaxies (as discussed previously), it is unlikely that our derived mid-IR space density lacks significant numbers of X-ray detected AGNs with [NeV] emission below our sensitivity limit. Given our comparatively small volume to that considered by using the SDSS (e.g., Greene \\& Ho 2007; H04), we further validate our derived space density of active SMBHs by robustly testing our results to find if: 1) our sample is over-dense, and thus strongly subject to cosmic variance; or 2) given the modest errors associated with our $\\Mbh$ estimations, the derived space-density is strongly subject to scattering of objects in our defined binning structure. We discuss these analyses in the Appendix. Briefly, we find that our sample is broadly representative of galaxies to $z \\sim 0.3$, and further show that even in our most pessimistic case, we find an increase in our derived space-density of at least a factor of $\\approx 2$ (maximum increase by a factor of $\\approx 11$) at $\\Mbh \\approx 3 \\times 10^6 \\Msun$ over the optical NL space-density of H04. \\subsection{The volume-weighted local AGN fraction} The ratio of the space densities of the active SMBH to total SMBH mass function (i.e., the volume-weighted local active SMBH fraction) is shown in the lower panel of Fig. \\ref{fig_3}. We calculate an overall active SMBH fraction of $\\approx 25^{+29}_{-14}$ percent for SMBHs of $\\Mbh \\approx (0.5$--$500) \\times 10^6 \\Msun$ down to our [NeV] completeness limit ($L_{\\rm [NeV]} \\ga 10^{38} \\ergps$). We find that this fraction is consistent with being constant throughout this $\\Mbh$ range. However, given our detection sensitivity limit, we are unable to probe lower Eddington ratios for AGNs hosting smaller SMBHs. Instead, we consider the effect of the AGN fraction for a fixed value of Eddington ratio (e.g., $\\eta > 10^{-3}$; i.e., thin-disc accretion systems) and find tentative evidence that the AGN fraction ($\\approx 16^{+9}_{-6}$ and $\\approx 8^{+10}_{-5}$ percent) may increase with decreasing SMBH mass ($\\Mbh \\approx 10^6$--$10^7 \\Msun$ and $\\Mbh \\approx 10^7$--$10^8 \\Msun$ bins, respectively). For the lowest-mass SMBHs ($\\Mbh \\approx (5$--$30) \\times 10^5 \\Msun$), we estimate an overall non-negligible volume-weighted AGN fraction of $18^{+12}_{-7}$ percent, potentially showing that a considerable proportion of small bulge (and pseudo-bulge) galaxies (i.e., late-type spiral galaxies; Sc--Sd) host AGN activity. It has been previously suggested by H04 and Greene \\& Ho (2007) that the AGN fraction may peak at $\\Mbh \\approx (0.7$--$2) \\times 10^7 \\Msun$. However, with the inclusion of the additional low-mass optically unidentified AGNs (see Fig.~\\ref{fig_2}), we find that the AGN fractions are consistent with remaining constant or even increasing for $\\Mbh < 10^7 \\Msun$. As noted in \\S5.3, the space-density of AGNs, and hence the local AGN fraction derived in this work, may only be a lower-limit given the nature of our volume-limited survey which by definition does not include IR-faint systems. To improve upon these current source statistics, a larger sample of late-type spiral galaxies would be required to investigate our findings further. With the greater sensitivity and resolving power proposed for the next generation of space-based mid-IR spectrographs, for example the {\\it Space Infrared Telescope for Cosmology and Astrophysics (SPICA)}\\footnote{See http://www.ir.isas.jaxa.jp/SPICA/} and the mid-IR instrument (MIRI) on-board the {\\it James Webb Space Telescope (JWST)}\\footnote{See http://www.roe.ac.uk/uktac/consortium/miri/}, surveys such as these can be continued and extended to study more distant (i.e., greater volumes) and heavily obscured AGNs." }, "1003/1003.4259_arXiv.txt": { "abstract": "We introduce an objective method to assess the probability of finding extreme events in the distribution of cold dark matter such as voids, overdensities or very high mass haloes. Our approach uses an ensemble of N-body simulations of the hierarchical clustering of dark matter to find extreme structures. The frequency of extreme events, in our case the cell or smoothing volume with the highest count of cluster-mass dark matter haloes, is well described by a Gumbel distribution. This distribution can then be used to forecast the probability of finding even more extreme events, which would otherwise require a much larger ensemble of simulations to quantify. We use our technique to assess the chance of finding concentrations of massive clusters or superclusters, like the two found in the two-degree field galaxy redshift survey (2dFGRS), using a counts-in-cells analysis. The Gumbel distribution gives an excellent description of the distribution of extreme cell counts across two large ensembles of simulations covering different cosmologies, and also when measuring the clustering in both real and redshift space. We find examples of structures like those found in the 2dFGRS in the simulations. The chance of finding such structures in a volume equal to that of the 2dFGRS is around 2 \\%. ", "introduction": "The discovery of extreme objects, such as large voids or highly overdense regions called superclusters, in which many galaxy clusters are found close together, has often been presented as a challenge to the hierarchical structure formation paradigm. However, the main drawback of using the presence of unusual structures to rule out a particular model is that it is not always clear how to assess the probability of finding such objects. In this paper we introduce a new methodology to address this problem in which we use N-body simulations and extreme value theory to provide a quantitative assessment of the probability of finding rare structures in a given cosmology. Claims of rare structures are common in the literature. Frith et~al. (2003) argued that the shape of the local galaxy counts implies a underdense volume in the southern sky $\\sim 300h^{-1}$Mpc across (see also Busswell et~al. 2004; Frith et~al. 2005, 2006). Cruz et~al. (2005) found a cold spot in the cosmic microwave background radiation that is much bigger than expected in a Gaussian distribution. Rudnick, Brown \\& Williams (2007) suggested that this cold spot is a secondary anisotropy, coinciding with the angular position of a void in a survey of radio galaxies. Swinbank et~al. (2007) found a large association of galaxy clusters, a supercluster of galaxies, at $z\\sim 0.9$ in the UK Infrared Deep Sky Survey. Sylos Labini, Vasilyev \\& Baryshev (2009a,b) argue that large scale density fluctuations are present in the local galaxy surveys which cannot be explained in current structure formation models. There are two common problems with the interpretation of such results. Firstly, what is the selection function of these objects, which would allow us to fix the frequency of finding such structures? And, secondly, what exactly are we looking at? For example, in the case of an overdensity of galaxies have we seen one massive cluster or are we looking at a projection of smaller structures along the line of sight? How should we compare the observations to theoretical predictions? In this paper we assess how common the superclusters found in the two-degree field galaxy redshift survey (2dFGRS; Colless et~al. 2001, 2003) are in the CDM cosmology. These structures were identified as ``hotspots'' in the distribution of galaxy counts-in-cells (Baugh et~al. 2004; Croton et~al. 2004). One structure is in the NGP part of the 2dFGRS at a redshift of $z = 0.08$ and a right ascension of 13 hours, and the other is in the SGP region at $z = 0.11$ at a right ascension of 0.5 hours. The higher order moments of the counts are strongly influenced by the presence of these structures (Croton et~al. 2004; Nichol et~al. 2006). A subsequent analysis of galaxy groups in the 2dFGRS revealed that these regions contain a surprisingly large fraction of all the massive clusters in the survey (Eke et~al. 2004a). Of the 94 groups in the full flux limited 2dFGRS out to $z \\sim 0.15$ with 9 or more members and estimated masses above $5 \\times 10^{14} h^{-1} M_{\\odot}$, 20 percent reside in these superclusters (Padilla et~al. 2004). The supercluster in the NGP region of the 2dFGRS is part of the ``Sloan Great Wall'' (Gott et~al. 2005). The 2dFGRS superclusters are by no means the largest superclusters in the local universe (for a list of superclusters, see Einasto et~al. 2001). The Shapley supercluster, for example, contains more Abell clusters than either of the 2dFGRS structures (Raychaudhury et~al. 1991; Proust et~al. 2006; Munoz \\& Loeb 2008). However, not all of the clusters contained within Shapley and similar mass concentrations have measured redshifts. Many of the member clusters have been identified in projection, and so their actual size is open to debate (Sutherland \\& Efstathiou 1991). The advantage of focusing on the 2dFGRS structures is that they have been identified from an unbiased redshift survey which was designed to map a particular volume of the Universe, rather than by targetting known structures. The volume of space in which the superclusters are found is therefore well defined. Furthermore, through the construction of the 2dFGRS Percolation Inferred Galaxy Group (2PIGG) catalogue (Eke et~al. 2004a), there is a clear, objective way to connect the observed properties of the galaxy groups which make up the superstructures to dark matter haloes in N-body simulations. In this paper we use extreme value theory to assess the probability of finding structures like the 2dFGRS superclusters in the CDM cosmology. Previous attempts to address the probability of finding such structures have used small numbers of simulations, and so have not been able to make definitive statements. For example, Croton et~al. (2004) measured the moments of the galaxy cell count distribution in the 22 mock catalogues whose construction was described by Norberg et~al. (2002). None of these mocks displayed higher order moments that looked like those measured in the 2dFGRS, giving a probability of less than 5\\% that such a structure could arise in a CDM model. One possible way around this problem is to generate estimates of the error on a measurement from the data itself (see Norberg et~al. 2009). Around 50 such estimates are required to get an accurate estimate of the variance on a measurement in the case of Gaussian statistics, and this method is clearly not applicable to a structure which appears once or twice in the dataset. The method we describe in this paper is calibrated against N-body simulations and can be extrapolated to very low probabilities, without requiring any assumption about the detailed form of the underlying distribution, just its asymptotic behaviour. The layout of this paper is as follows. In Section 2 we first give a very brief overview of extreme value theory, before discussing how we connect dark matter haloes from an N-body simulation to galaxy groups in the 2PIGG catalogue. Finally in Section 2 we describe the ensembles of N-body simulations that we use to measure the counts-in-cells distribution of massive haloes. The results of the paper are presented in Section 3, in which we show the impact of mass errors on the distribution of halo counts-in-cells and demonstrate how well extreme value theory describes the probability of finding a ``hot'' cell. Our conclusions are given in Section 4. ", "conclusions": "We have introduced a new objective methodology for assessing the likelihood of finding extreme structures in hierarchical structure formation models. Quite often the probability of finding an unusual structure such as a large void or an overdensity is estimated using a Gaussian distribution, because the smoothing scale in question is large. This is a good approximation for events which represent small departures from the mean density. However, for extreme events this is a poor assumption. The probability distribution of the density contrast on a particular smoothing scale, although assumed to be initially Gaussian in most models, rapidly evolves away from this form due to gravitational instability (e.g. Gazta\\~{n}aga, Fosalba \\& Elizalde 2000). Assuming a Gaussian probability distribution rather than the true distribution would lead to a misestimation of the probability of finding a cell with an extreme density by many orders of magnitude. The advantage of our approach is that we do not need to specify the actual form of the probability distribution of cell counts. We have shown that the distribution of extreme cell counts is well described by a Gumbel distribution in a range of different situations: real-space, redshift-space and both with and without errors in halo mass. The simulations give the mean and variance of the Gumbel distribution. The analytic form can then be extrapolated into the tails to give the probability of events which would require hundreds or even thousands of realizations of N-body simulations to find. By using N-body simulations, we can assess the probability of events which cannot be calculated analytically, such as the largest Einstein ring expected in a CDM model (e.g. Oguri \\& Blandford 2009). Cells with the number of massive haloes seen in the 2dFGRS can be found within our simulations, provided that the clustering of these haloes is measured in redshift-space and the mass errors introduced by the group finding algorithm are taken into account. However, if we consider a volume of the size of the 2dFGRS $L_*$ sample, which is 300 times smaller than our simulation volume, then we expect to find such an overdensity of cluster mass haloes in $\\approx 2$ out of a hundred cases. Norberg et al. (2010) have carried out a related analysis using the final release of the SDSS. These authors applied a different technique to identify overdense regions. They split the galaxy distribution into zones, as would be done to carry out a jackknife estimation of the error on clustering statistics (Norberg et~al. 2009). By comparing the distribution of two and three point correlation functions measured from the jackknife resamplings, a zone whose omission produced an outlying estimate of the clustering was found. However, when applying the same analysis to the same ensemble of N-body simulations used in this paper, Norberg et~al. (2010) found that such outliers were quite common. There are some key differences between our analyses. The SDSS volume limited sample is an order of magnitude larger than the 2dFGRS sample considered in this paper. Norberg et al. found one ``unusual'' structure in the SDSS volume limited sample for $L_*$ galaxies. Also, the method for quantifying unusual structures is different from the one we use, and will pick up a very different type of structure. The zones used by Norberg et~al. sample conical volumes of space, covering a large baseline in radial distance. The superstructure in their case could be a projection of independent structures along the line of sight. In our case, we use compact cells. Norberg et~al. (2010) conclude that in the larger volume of the SDSS $L_*$ sample, the structures found by their correlation function study are consistent with those found in CDM. Our results do not contradict this; we have applied a different test to search for overdense regions with a different structure." }, "1003/1003.3479_arXiv.txt": { "abstract": "We present an analysis of the far-infrared (FIR) spectral energy distributions (SEDs) of two massive K-selected galaxies at $z =$ 2.122 and $z =$ 2.024 detected at 24$\\micron$, 70$\\micron$, 160$\\micron$ by Spitzer, 250$\\micron$, 350$\\micron$, 500$\\micron$ by BLAST, and 870$\\micron$ by APEX. The large wavelength range of these observations and the availability of spectroscopic redshifts allow us to unambiguously identify the peak of the redshifted thermal emission from dust at $\\sim$ 300$\\micron$. The SEDs of both galaxies are reasonably well fit by synthetic templates of local galaxies with L$_{IR}$ $\\sim$ 10$^{11}$L$_{\\odot}$ -- 10$^{12}$L$_{\\odot}$ yet both galaxies have L$_{IR}$ $\\sim$ 10$^{13}$L$_{\\odot}$. This suggests that these galaxies are not high redshift analogues of the Hyper-LIRGs/ULIRGs used in local templates, but are instead \"scaled up\" versions of local ULIRGs/LIRGs. Several lines of evidence point to both galaxies hosting an AGN; however, the relatively cool best fit templates and the optical emission line ratios suggest the AGN is not the dominant source heating the dust. For both galaxies the star formation rate determined from the best-fit FIR SEDs (SFR(L$_{IR}$)) agrees with the SFR determined from the dust corrected H$\\alpha$ luminosity (SFR(H$\\alpha$)) to within a factor of $\\sim$ 2; however, when the SFR of these galaxies is estimated using only the observed 24$\\micron$ flux and the standard luminosity-dependent template method (SFR(24$\\micron$)), it systematically overestimates the SFR by as much as a factor of 6. A larger sample of 24 K-selected galaxies at $z \\sim$ 2.3 drawn from the Kriek et al. (2008) GNIRS sample shows the same trend between SFR($24\\micron$) and SFR(H$\\alpha$). Using that sample we show that SFR($24\\micron$) and SFR(H$\\alpha$) are in better agreement when SFR($24\\micron$) is estimated using the log average of local templates rather than selecting a single luminosity-dependent template, because this incorporates lower luminosity templates. The better agreement between SFRs from lower luminosity templates suggests that the FIR SEDs of the BLAST-detected galaxies may be typical for massive star forming galaxies at $z \\sim 2$, and that the majority are scaled up versions of lower luminosity local galaxies. ", "introduction": "It is well known that in the local universe the obscuration towards star forming regions is correlated with star formation rate (SFR; e.g., Wang \\& Heckman 1996; Hopkins et al. 2001; Buat et al. 2005, 2007), with the most active galaxies frequently being the most dust obscured. In the most extreme galaxies, the star forming regions can become optically-thick (e.g., Genzel et al. 1998), which makes mid- and far-infrared (MIR, FIR) data critical for a complete metric of their SFR. At $z \\sim$ 2, where the average massive galaxy is forming stars at rates $\\sim$ 100 - 1000 times higher than in the local universe (e.g., P\\'{e}rez-Gonz\\'{a}lez et al. 2005; Juneau et al. 2005; Damen et al. 2009), MIR- and FIR-determined SFRs are critical for understanding the buildup of today's massive galaxies. \\newline\\indent The advent of the MIPS instrument onboard $Spitzer$ has facilitated the first deep photometry at 24$\\micron$ and numerous studies have used these data to infer the star formation history of the universe up to $z \\sim$ 2 (e.g., Le Floc'h et al. 2005; P\\'{e}rez-Gonz\\'{a}lez et al. 2005; Caputi et al. 2007; Reddy et al. 2008; Magnelli et al. 2009). These studies suggest that obscured star formation dominates in massive galaxies at $z \\sim$ 2; however, the conversion from observed 24$\\micron$ flux to a total infrared luminosity (L$_{IR}$), and subsequently a SFR at $z \\sim$ 2 requires significant extrapolation based on local FIR templates. \\newline\\indent The most common method for determining the L$_{IR}$ of a galaxy from 24$\\micron$ photometry is to artificially redshift local luminosity-dependent MIR/FIR templates to the redshift of the galaxy and then choose the L$_{IR}$ of the template that predicts a 24$\\micron$ flux closest to the observed 24$\\micron$ flux. By applying this method, or more sophisticated but related versions of it, some authors have found that the local templates produce estimates of L$_{IR}$ and SFR(L$_{IR}$) for distant galaxies that are in agreement with other SFR indicators (e.g., Elbaz et al. 2002, Marcillac et al. 2006, Reddy et al. 2006; Daddi et al. 2007a; Papovich et al. 2009; Kartaltepe et al. 2010), while others have argued that this method tends to systematically overestimate the L$_{IR}$ and hence the SFR (e.g., Papovich et al. 2007; Rigby et al. 2008; Murphy et al. 2009). Some of these differences can probably be reconciled by the different selection criteria (e.g., Reddy et al. 2010), redshift range (e.g., Murphy et al. 2009) and luminosities (e.g., Papovich et al. 2007) of the above samples; however, the reliability of observed 24$\\micron$ fluxes as a metric of the L$_{IR}$ and SFR over the full range of galaxy masses and SFRs at $z \\sim$ 2 remains unclear. \\newline\\indent The next step will be to directly observe the cold dust SED, rather than extrapolate it from the rest-frame 6 -- 10$\\micron$ region, which can be complicated by a superposition of Polycyclic Aromatic Hydrocarbon (PAH) emission lines and silicate absorption features. The peak of the cold dust emission in FIR luminous galaxies in the local universe occurs between 40$\\micron$ to 150$\\micron$ which corresponds to 120$\\micron$ to 450$\\micron$ at $z \\sim$ 2. In the next few years the $Herschel$ telescope will provide unprecedented amounts of data at these wavelengths; however, the recent flight of the BLAST telescope (Devlin et al. 2009) has already provided some of the first deep observations of distant galaxies at 250$\\micron$, 350$\\micron$ and 500$\\micron$. \\newline\\indent In this paper we take advantage of the BLAST observations of the Extended Chandra Deep Fields South field (ECDFS) to directly measure the FIR SEDs of $z \\sim$ 2 galaxies. Two star forming galaxies in the Kriek et al. (2008) NIR spectroscopic sample are located in that field and are detected or have upper limits at 24$\\micron$, 70$\\micron$, and 160$\\micron$ in the FIDEL (Magnelli et al. 2009) and SWIRE (Lonsdale et al. 2003) surveys, and are also detected by BLAST (Dye et al. 2009). The ECDFS was recently surveyed by the LABOCA instrument on the APEX telescope, and both galaxies were detected at 870$\\micron$ (Weiss et al. 2009). By combining these data sets we have produced some of the first FIR SEDs that sample the cool dust bump of $z \\sim$ 2 galaxies with good resolution and allow us to constrain their L$_{IR}$ directly from the FIR SED. Throughout this paper we assume a H$_{0}$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m}$ = 0.3, $\\Omega_{\\Lambda}$ = 0.7 cosmology when computing luminosity distances. ", "conclusions": "\\subsection{Dust Temperatures and Evidence for \"Scaled Up\" Cool Galaxies} It is remarkable that despite an L$_{IR}$ $\\sim$ 10$^{13}$ L$_{\\odot}$, both of the BLAST-detected $z \\sim$ 2 galaxies are better fit with local templates that have PAH features and a cool \"cold\" dust bump. In the nearby universe, galaxies with L$_{IR}$ $\\sim$ 10$^{13}$ L$_{\\odot}$ tend to have both the warmest \"cold\" dust (T$_{dust}$ $\\sim$ 10 -- 70K, emission peaking in the FIR) and significant amounts of \"hot\" dust (T$_{dust}$ $\\sim$ 70 -- 500K, emission peaking in the MIR). Although both of these dust components can come from deeply embedded star formation (e.g., Tran et al. 2001), it is thought that the dominant radiation source in most, if not all nearby galaxies with L $\\sim$ 10$^{13}$ L$_{\\odot}$ is an AGN (e.g., Genzel et al. 1998; Lutz et al. 1998). \\newline\\indent By contrast, the best fit templates of the $z \\sim$ 2 galaxies are much cooler ULIRG and LIRG templates which are most frequently associated with star forming galaxies in the local universe (e.g., Armus 2009). These templates also tend to have stronger PAH features that are more prevalent because the \"hot\" dust component from the AGN no longer dominates the MIR emission. \\newline\\indent The first claims that IR luminous galaxies at $z \\sim$ 2 may be scaled up versions of lower luminosity cooler galaxies came from 24$\\micron$ observations of submillimeter-selected galaxies (SMGs). Pope et al. (2006) showed that the SEDs of SMGs peak at longer wavelengths than expected based on an extrapolation of their 24$\\micron$ flux using local templates, arguing that SMGs are cooler at a given luminosity than the local templates. Subsequent MIR spectroscopic observations of SMGs from $Spitzer$ by Lutz et al. (2005), Pope et al. (2008), and Menendez-Delmestre et al. (2009) all found strong PAH features and relatively weak MIR continuum in the SMGs suggesting much less \"hot\" dust than is usually found in local galaxies of comparable L$_{IR}$. Similarly, Papovich et al. (2009) found a weaker than expected rest-frame 24$\\micron$ luminosity for a lensed SMG, arguing for the lack of a \"hot\" dust component, and hence a scaled up cool galaxy. \\newline\\indent With access to the FIR we can determine a dust temperature (T$_{d}$) independent of the MIR SED and compare this to the local T$_{d}$-L$_{IR}$ relation. We fit the 160$\\micron$ $\\rightarrow$ 870$\\micron$ data (50$\\micron$ $\\rightarrow$ 300$\\micron$ rest-frame) to modified blackbody curves of the form S($\\nu$) = A$\\nu^{\\beta}$B($\\nu$,T$_{d}$), where A is a normalization, B is the Planck function, and $\\beta$ accounts for frequency-dependent dust emissivity. Even with detections in numerous bands the S/N of the data is too low to constrain both $\\beta$ and T$_{d}$ simultaneously so we have assumed $\\beta$ = 1.5, a typical value for submillimeter galaxies. The 70$\\micron$ points are omitted in the fitting because they correspond to rest-frame 23$\\micron$ and are too far into the MIR to include in a single-component blackbody fit. We note that including the 70$\\micron$ data actually produces identical best-fit temperatures; however, the $\\chi^2$ of fits are substantially larger because no single-temperature model can describe the 23$\\micron$ $\\rightarrow$ 300$\\micron$ rest-frame simultaneously. \\newline\\indent We find dust temperatures of 40$^{+2}_{-1}$K and 41$^{+5}_{-7}$K for ECDFS-4511 and ECDFS-12514, respectively. These correspond to rest-frame monochromatic IRAS color measurements, C $\\equiv$ log$_{10}$(f$_{60\\micron}$/f$_{100\\micron}$), of -0.05$^{+0.05}_{-0.03}$ and -0.03$^{+0.09}_{-0.20}$, respectively. Comparing these to the C-L$_{IR}$ relation measured by Chapin et al. (2009) in the nearby universe using IRAS shows that these galaxies have temperatures typical of nearby galaxies with L$_{IR}$ = 10$^{11}$ -- 10$^{12}$ L$_{\\odot}$. They are also similar to the T$_{d}$ of local LIRGS/ULIRGS measured by Clements et al. (2010). Therefore, with the combined Spitzer, BLAST and APEX data we can now confirm that some massive galaxies at $z \\sim$ 2 not only have PAH-dominated MIR SEDs similar to lower luminosity local galaxies, but they also have \"cold\" dust SEDs consistent with being scaled up versions of lower luminosity galaxies. This suggests that star formation may be occurring in a different environment within galaxies at $z \\sim$ 2 compared to $z \\sim$ 0. \\newline\\indent Local galaxies with SFRs $>$ 100 M$_{\\odot}$ yr$^{-1}$ tend to have compact nuclear starbursts which are the cause of the warmer dust temperatures (e.g., Sanders \\& Mirabel 1996). The $z \\sim$ 2 galaxies have SFRs $\\sim$ 1000 M$_{\\odot}$ yr$^{-1}$, but similar dust temperatures which suggests that the star formation may be occurring in multiple large starburst regions throughout the galaxy, rather than concentrated in a single compact starburst which would presumably have a substantially hotter dust temperature. Indeed, a lensed $z > 2$ galaxy with precisely such characteristics has recently been observed by Swinbank et al. (2010a) using interferometric submillimeter observations. \\newline\\indent Kriek et al. (2009a) presented SINFONI-IFU H$\\alpha$ spectra of both these galaxies. Their maps show that the H$\\alpha$ emission is resolved on scales of a few kpc, and is also consistent with the extended star formation hypothesis. The greater availability of gas at high redshift may permit more spatially extended star formation at a high rate, something that is uncommon in local galaxies. \\subsection{Morphologies} In the inset of Figure 1 we show 3$\"$$\\times$3$\"$ NIC2 F160W images of the galaxies from Kriek et al. (2009a). At $z \\sim$ 2 these correspond to rest-frame V-band morphologies and span an angular size of $\\sim$ 25 kpc. Like most of the high redshift ULIRGs and HLIRGs previously observed with HST (e.g., Chapman et al. 2003; Dasyra et al. 2008; Swinbank et al. 2010b) the galaxies clearly have disturbed morphologies. Kriek et al. (2009a) measure effective radii (R$_{e}$) of 5 and 3 kpc for ECDFS-4511 and ECDFS-12514, respectively, which means that both galaxies are comparable in size to local galaxies of similar mass. Rest-frame optical morphologies make it difficult to identify the location of the star formation (which is presumably highly obscured); however, the disturbed nature of both galaxies shows that these are complex systems, possibly \"stream-fed\" galaxies (Dekel et al. 2009) or mergers. Unlike the FIR SEDs, the morphologies do not appear to be scaled up versions of local star forming disk galaxies. \\subsection{AGN Contribution to the MIR and FIR Emission} Despite SEDs that are associated with local star forming galaxies, part of the MIR and FIR emission in both galaxies almost certainly comes from an AGN. Using emission line ratios and the BPT diagram (Baldwin et al. 1981), Kriek et al. (2007) found that both galaxies were in the region of the diagram that Kewley et al. (2006) considered to have \"composite\" spectra with contributions from both star formation and an AGN. \\newline\\indent ECDFS-4511 is detected in the 2 Ms $Chandra$ map of the CDFS (Luo et al. 2008), and has an L$_{x}$ $\\sim$ 2$\\times$10$^{44}$ erg s$^{-1}$, implying it hosts a Seyfert-luminosity AGN. ECDFS-12514 is located in the 250 ks map of the ECDFS (Virani et al. 2006), but is undetected. Although undetected in the X-rays, ECDFS-12514 is formally classified as an IRAC powerlaw galaxy (F$_{\\nu}$ $\\propto$ $\\nu^{\\alpha}$, $\\alpha$ $<$ -0.5; e.g., Donley et al. 2007, see Figure 1) which makes it a candidate for an obscured AGN. \\newline\\indent Without more data we cannot constrain the precise fraction of the MIR and FIR light that comes from the AGN; however, the consistency of the MIR/FIR SED with local templates that are star formation dominated, and the \"composite\" emission line ratios suggest that young stars are probably the dominant energy source heating the dust. Pope et al. (2008) and Menendez-Delmestre et al. (2009) showed that the MIR SEDs of most SMGs (the BLAST-detected galaxies are formally SMGs) are still dominated by PAHs and that AGN only contribute $\\sim$ 30\\% of the total MIR emission. Putting the evidence together the most plausible scenario is that both galaxies host a moderate-luminosity AGN, but the AGN does not dominate the MIR/FIR SED. Nevertheless, we note that both our SFR(L$_{IR}$) and SFR(H$\\alpha$) should be formally considered upper limits as the AGN will contribute to both of these. \\subsection{Implications for SFRs at $z \\sim$ 2 and MIR-excess Galaxies} The results of our FIR SED fitting show that caution is required when determining and interpreting the SFRs of $z \\sim$ 2 galaxies based on limited data. Using only 24$\\micron$ photometry to infer SFRs at $z \\sim$ 2 may work well for lower luminosity, or less obscured galaxies (e.g., Papovich et al. 2007; Reddy et al. 2010); however, the local templates can overestimate the SFR(L$_{IR}$) of the most FIR luminous galaxies by as much as a factor of $\\sim$ 6 (see also Papovich et al. 2007; Murphy et al. 2009). Likewise, the Balmer decrement measurements for these galaxies seem to indicate extra extinction toward the star forming regions. It appears that on average this relation is similar to that suggested by Calzetti et al. (1997); however, the lack of secure H$\\beta$ detections in our sample precludes quantitative statements. To correctly determine the SFR(H$\\alpha$) of individual galaxies, rather than ensemble averages requires high quality Balmer decrement measurements. \\newline\\indent The extra extinction towards the star forming regions suggested by the Balmer decrements also indicates that the UV SFRs for some of these IR luminous galaxies may not be reliable because the A$_{v}$ is large enough that parts of the star forming regions are optically-thick. The SFRs estimated from a fit to the stellar populations (SFR(SED); effectively a dust corrected UV SFR) by Muzzin et al. (2009) for ECDFS-4511 and ECDFS-12514 are 168$^{+41}_{-63}$ M$_{\\odot}$ yr$^{-1}$ and 145$^{+160}_{-0}$ M$_{\\odot}$ yr$^{-1}$, and underestimate the SFR(L$_{IR}$) of the galaxies by factors of 5.2 and 3.5, respectively. If we compare the SFR(SED) to the SFR(24$\\micron$) as in e.g., Daddi et al. (2007b) we find that these galaxies have MIR SFRs 18 -- 25 times greater than the UV SFRs, and are extremely strong \"MIR excess\" galaxies. \\newline\\indent The purpose of the Daddi et al. (2007b) classification was to identify candidate obscured AGN, and given that it is likely that both galaxies contain an AGN it appears that method successfully identifies such sources. However, MIR excesses of factors of 18 -- 25 would suggest that the AGN is the dominant MIR energy source for both galaxies, something that is unlikely given the shape of the MIR/FIR SEDs and the emission line ratios. This calls into question whether galaxies with weaker inferred MIR excesses, such those near the cutoff of SFR(UV$_{corr}$)/SFR(MIR+UV) $\\sim$ 3 suggested by Daddi et al. (2007b) have a significant contribution from an AGN. \\subsection{Comparison to Models} Our results suggest that a substantial population of massive $z \\sim$ 2 galaxies are \"scaled up\" versions of lower luminosity galaxies. Clearly such a population needs to be understood within the context of modern galaxy formation and evolution models. Recently, Hopkins et al. (2010) have used their merger-driven evolutionary model to predict the IR luminosity functions of galaxies and quasars between 0 $< z <$ 6. In their model they argue that there is a threshold in L$_{IR}$ above which disks, merger-driven star formation, and obscured AGN are the dominant IR energy source of galaxies. At $z = 0$ these limits are log(L$_{IR}$) $\\geq$ 11.5 for merger-driven star formation, and log(L$_{IR}$) $\\geq$ 12.5 for obscured AGN. Based on the evolution of the merger rate and gas fraction to $z \\sim$ 2 they suggest that these thresholds evolve in log(L$_{IR}$) to log(L$_{IR}$) $\\geq$ 12.5 for merger-driven star formation and log(L$_{IR}$) $\\geq$ 13.0 for obscured AGN. With log(L$_{IR}$) $\\sim$ 13.0, the BLAST galaxies lie at an L$_{IR}$ that is near the suggested transition point between merger-driven star formation dominated or AGN dominated. \\newline\\indent Interestingly, the FIR SEDs of these galaxies do not look like those of the AGN-dominated HLIRGs detected by $ISO$ (Genzel et al. 1998), or of the average nearby ULIRG, which are mostly merger-driven starbursts (e.g., Sanders et al. 1998). Instead, these galaxies look like scaled up versions of even lower luminosity systems, which are not universally associated with mergers or AGN. Taken at face value, it appears they do not fit well within the Hopkins et al. (2010) description. \\newline\\indent On the other hand, the NIC2 morphologies of the galaxies are clearly disturbed enough to be considered candidates for undergoing mergers (Kriek et al. 2009b). If these galaxies are merger-driven starbursts, it may be that they are missing the centrally-concentrated starburst which produces the much hotter dust seen in these galaxies in the local universe. As Hopkins et al. (2010) point out, the efficiency of a merger at funneling gas and stars into the core of the galaxy is proportional to the gas fraction, with higher gas fraction mergers being less efficient. If both galaxies have high gas fractions most of the star formation could be occurring on much larger physical scales than similar mass starbursts in the local universe, hence the cool dust SED. \\newline\\indent Galaxies with similar stellar masses (M $\\sim$ 10$^{11}$ M$_{\\odot}$) and gas fractions as high as 80\\% have been observed in the distant universe (e.g., Tacconi et al. 2010) and could be analogues of the BLAST-detected galaxies. If that is the case, it begs the question of what are the descendants of extended, massive, gas-rich star forming galaxies at $z \\sim$ 2? Local spheroids seem like an obvious candidate; however, some fraction of the progenitors of those systems are clearly small and quiescent at $z \\sim$ 2 (e.g., van Dokkum et al. 2010; Hopkins et al. 2010; Bezanson et al. 2009; van der Wel et al. 2008). If the large and active BLAST-sources are also progenitors of massive local spheroids it suggests there may be very different evolutionary paths which arrive at the same galaxies. \\newline\\indent These galaxies also appear to be consistent with the \"stream-fed\" galaxies proposed by Dekel et al. (2009) and the scaled up \"turbulent disks\" discussed by Genzel et al. (2008). In their model, Dekel et al. (2009) suggest that massive galaxies may grow quickly through rapid cold gas accretion and clumpy star formation which funnels stars into a central bulge. Given that these cold flows often separate into multiple large star forming clumps, it might be expected that these clumps produce cooler dust temperatures at fixed SFR than centrally concentrated starbursts. \\newline\\indent The SFRs of the BLAST galaxies are higher than the 100 -- 200 M$_{\\odot}$ yr$^{-1}$ predicted for \"stream-fed\" galaxies by Dekel et al. (2009). The BLAST galaxies are the most IR luminous of the K-selected sample of Kriek et al. (2008), so if they are \"stream-fed\" galaxies, they may just represent the most active tail of the distribution. \\newline\\indent High quality kinematic data could potentially distinguish between merger-dominated or accretion-dominated systems. These galaxies do have SINFONI IFU observations; however, these data lack the S/N to reliably determine V/$\\sigma$. Therefore, at present based on the MIR/FIR properties and morphologies we cannot distinguish between evolutionary models such suggested by both Hopkins et al. (2010) and Dekel et al. (2009)." }, "1003/1003.2858_arXiv.txt": { "abstract": "We report on the near-infrared selected AGN candidates extracted from 2MASS/ROSAT catalogues and discuss their properties. First, near-infrared counterparts of a X-ray source in ROSAT catalogues (namely, Bright Source Catalogue (BSC) and Faint Source Catalogue (FSC)) were extracted by positional cross-identification of $\\leq 30''$. Because these counterparts would contain many mis-identifications, we further imposed near-infrared colour selection criteria and extracted reliable AGN candidates (BSC: 5,273, FSC: 10,071). Of 5,273 (10,071) candidates in the BSC (FSC), 2,053 (1,008) are known AGNs. Near-infrared and X-ray properties of candidates show similar properties with known AGNs and are consistent with previous studies. We also searched for counterparts in other wavelengths (that is, optical, near-infrared, and radio), and investigated properties in multiwavelength. No significant difference between known AGNs and unclassified sources could be seen. However, some unclassified sources in the FSC showed slightly different properties compared with known AGNs. Consequently, it is highly probable that we could extract reliable AGN candidates though candidates in the FSC might be spurious. ", "introduction": "Active galactic nuclei (AGNs) are luminous in wide wavelength because of the vast amount of energy produced by accretion onto a central supermassive blackhole. Accordingly, they are observable at nearly entire wavelength despite locating at great distances, and have been searched in various wavelengths: e.g., optical \\citep{Schneider2007-AJ,Veron2006-AA}, infrared \\citep{Low1988-ApJ,Low1989-ApJ}, or radio \\citep{Frayer2004-AJ}. X-ray emission is especially a characteristic property of AGNs. The vast majority of X-ray sources are AGNs and all classes of AGNs appear in X-ray surveys. There are many studies that collect AGN samples using X-ray data \\citep[e.g., ][]{Kim1999-ApJ,Watanabe2004-ApJ,Polletta2007-ApJ}. Therefore, whether an object is a X-ray source can be a criterion to select AGNs. Colour selection is a powerful technique in extracting AGN candidates. A classical method is known as the $UV$-excess \\citep[UVX; ][]{Sandage1965-ApJ,Schmidt1983-ApJ,Boyle1990-MNRAS}, which extract bluer quasars. \\citet{Richards2002-AJ} selected quasars via their nonstellar colours using SDSS photometry. Other selections are such as red quasar survey using optical and near-infrared combined colours \\citep{Glikman2007-ApJ}, or mid-infrared selected AGNs using Spitzer data \\citep{Lacy2004-ApJS,Stern2005-ApJ}. Colour selection using near-infrared photometry has also performed by some previous studies. \\citet{Cutri2001-ASPC,Cutri2002-ASPC} extracted obscured AGN candidates using the colour selection of $(J-K_\\textnormal{\\tiny S})>2.0$. The $KX$ method using the excess in K-band, proposed by \\citet{Warren2000-MNRAS}, was used for extracting quasars \\citep{Jurek2008-MNRAS,Maddox2008-MNRAS,Smail2008-MNRAS,Nakos2009-AA}. However, these extracted only peculiar AGNs (in the former case) or extracted AGNs using combined with optical photometry (in the latter case). Because the optical light suffers more extinctions than the infrared light and an AGN is surrounded by a dust torus, a selection using optical and near-infrared combined colours may miss AGNs (especially obscured AGNs), because of lack of optical detection. However, \\citet{Kouzuma2010-AA} proposed colour selection criteria to extract AGNs using only near-infrared colours. They demonstrated by both observed and simulated colours that AGNs are differentiated from several types of objects in a $(H-K_\\textnormal{\\tiny S})$-$(J-H)$ colour-colour diagram (CCD). This enables us to extract AGN candidates using only near-infrared photometry. In this paper, we first extract bright sources in both near-infrared and X-ray by a cross-identification between the Two Micron All Sky Survey (2MASS) and ROSAT all-sky survey catalogues, and select AGN candidates on the basis of the near-infrared colour selection criteria proposed by \\citet{Kouzuma2010-AA}. In addition, we investigate properties of candidates using not only near-infrared and X-ray data but also photometric data at other wavelengths derived by cross-identifications with some catalogues. In Section \\ref{DATA}, we introduce the 2MASS and ROSAT. In Section \\ref{EXTRACTION}, we describe the method to extract AGN candidates (including the criteria of both cross-identification and colour selection) and the results. In Section \\ref{PROPERTIES}, properties of AGN candidates are investigated by photometric data at near-infrared, X-ray, and other wavelengths. ", "conclusions": "We have cross-identified the 2MASS PSC with ROSAT catalogues, and have extracted AGN candidates on the basis of the near-infrared colour selection. Of 5,273 (10,701) AGN candidates in the BSC (FSC), 3,220 (9,693) are unclassified sources. We investigated their properties using near-infrared, X-ray, optical, and radio data. Overall, most unclassified sources in the BSC have similar properties with known AGNs and their properties are also consistent with previous studies, whereas some unclassified sources, especially those in the FSC, have properties different from known AGNs. It is highly probable that most unclassified sources are AGNs because of the following reasons. First, our candidates are X-ray sources. The fact that a object is a X-ray source supports the object is an AGN because it is believed that the vast majority of X-ray sources are AGNs. Second, our candidates satisfy the near-infrared colour selection criteria proposed by \\citet{Kouzuma2010-AA}, that is, they have properties in near-infrared colours similar to those of AGNs. Third, some candidates have been already known AGNs and not only unclassified sources show properties similar to those of known AGNs but also their properties in multiwavelength are consistent with previous studies. However, it should be noted that there may be relatively large number of contamination in unclassified FSC sources. To confirm that they are real AGNs, spectroscopic observations are required. Similar works using a cross-identification with a ROSAT catalogue were performed by \\citet{Boller1992-AA2,Boller1998-AAS}. \\citet{Boller1992-AA2} extracted a sample of 14,708 extragalactic IRAS sources on the basis of a supervised selection whose selection quality was assessed by \\citet{Boller1992-AA}, and they cross-identified between the sample and the first processing of the ROSAT all-sky survey data. The cross-identified 244 IRAS galaxies comprised infrared luminous nearby galaxies, spirals and ellipticals, Seyfert galaxies, and QSOs. They investigated properties of the cross-identified sources using X-ray, far-infrared, and optical fluxes, and found some correlations between them. The work in \\citet{Boller1998-AAS} is also similar to that in \\citet{Boller1992-AA2}, which cross-indentified the 14,315 IRAS galaxies with the second processing of the ROSAT all-sky survey. These studies are similar to our study in that a ROSAT catalogue is cross-identified with an infrared catalogue. However, whereas these studies investigated extragalactic sources extracted by a mid- and far-infrared colour selection, we investigated AGN candidates extracted by the near-infrared colour selection. The number of sample is also totally different. In addition, we discussed not only properties in infrared and X-ray but also near-infrared variability and radio properties, which are not discussed in \\citet{Boller1992-AA2,Boller1998-AAS}. This paper extracted several thousands of AGN candidates across the entire sky. The use of only near-infrared colours enables us to search for AGN candidates across the entire sky. When using this technique on other deep near-infrared surveys such as DENIS, UKIDSS, and future surveys, many AGNs, which have not been detected in the optical wavelength, can be extracted." }, "1003/1003.5680_arXiv.txt": { "abstract": "\\begin{center} {\\large\\bf Abstract} \\end{center} We formalize a classification of pair interactions based on the convergence properties of the {\\it forces} acting on particles as a function of system size. We do so by considering the behavior of the probability distribution function (PDF) $P(\\bF)$ of the force field $\\bF$ in a particle distribution in the limit that the size of the system is taken to infinity at constant particle density, i.e., in the ``usual'' thermodynamic limit. For a pair interaction potential $V(r)$ with $V(r \\rightarrow \\infty) \\sim 1/r^\\gamma$ defining a {\\it bounded} pair force, we show that $P(\\bF)$ converges continuously to a well-defined and rapidly decreasing PDF if and only if the {\\it pair force} is absolutely integrable, i.e., for $\\gamma > d-1$, where $d$ is the spatial dimension. We refer to this case as {\\it dynamically short-range}, because the dominant contribution to the force on a typical particle in this limit arises from particles in a finite neighborhood around it. For the {\\it dynamically long-range} case, i.e., $\\gamma \\leq d-1$, on the other hand, the dominant contribution to the force comes from the mean field due to the bulk, which becomes undefined in this limit. We discuss also how, for $\\gamma \\leq d-1$ (and notably, for the case of gravity, $\\gamma=d-2$) $P(\\bF)$ may, in some cases, be defined in a weaker sense. This involves a regularization of the force summation which is generalization of the procedure employed to define gravitational forces in an infinite static homogeneous universe. We explain that the relevant classification in this context is, however, that which divides pair forces with $\\gamma > d-2$ (or $\\gamma < d-2$), for which the PDF of the {\\it difference in forces} is defined (or not defined) in the infinite system limit, without any regularization. In the former case dynamics can, as for the (marginal) case of gravity, be defined consistently in an infinite uniform system. ", "introduction": "Interactions are traditionally classified as long-range (or short-range) with respect to the non-additivity (or additivity) of the {\\it potential energy} in the usual thermodynamic limit, i.e., when the number of particles $N$ and volume $V$ are taken to infinity at constant particle density. This is the property which determines the way in which standard instruments of statistical mechanics are applied to determine {\\it equilibrium} properties (see e.g. \\cite{ruelle, Dauxoisetal, assisi}). Indeed in the case of long-range interactions, these instruments are applied using an appropriately generalized thermodynamic limit, in which the coupling or density are also scaled with system size. Such an analysis gives rise generically to features at equilibrium which are qualitatively different from those in short-range systems --- inhomogeneous statistical equilibria, non-equivalence of statistical ensembles, negative specific heat in the microcanonical ensemble (see e.g. \\cite{Dauxoisetal, assisi}). Most of these unusual features were first noted and studied in the context of the study of gravitating systems in astrophysics (see e.g. \\cite{thirring_book, chavanis_phase+transitions_2006} for reviews), and it has been realized in recent years that they are more generic in long-range interacting systems. This thermodynamic analysis extended to long-range systems is believed to determine, however, the behavior of such systems only on time scales which diverge as some power of $N$ (when expressed in terms of the characteristic dynamical time scales). On shorter times scales --- usually those of interest in practical applications --- study of several such systems (see e.g. \\cite{yawn+miller_2003, yamaguchi_etal_04, antoniazzi1, antoniazzi2, yamaguchi_1dgravLB_2008} and references therein) shows that they appear generically, like in the well-documented case of gravity, to relax from almost any initial conditions to (almost) time-independent states --- referred to variously as ``meta-equilibria'', ``quasi-equilibria'' or ``quasi-stationary states'' (QSS). The physics of these states, which are generically very different from those at thermal equilibrium, is understood to be the result of evolution in the collisionless regime described by Vlasov equation (usually referred to as the ``collisionless Boltzmann equation'' in the astrophysical literature \\cite{binney}). Both the genesis of these states and their long-time relaxation are poorly understood, and are the subject of active study (see e.g. \\cite{campa_etal2007, Baldovin_etal2009, chavanis_kEqns_2010, teles_etal2010, gupta+mukamel_2010}). In this article we consider a simple classification of pair interactions different to this usual thermodynamic one. Instead of considering the convergence properties of {\\it potential energy} in the usual thermodynamic limit, we consider those of the {\\it force}, in the same limit. The resulting classification can, like the usual one, be understood easily from simple considerations. To see this let us consider, as illustrated schematically in Fig.~1, \\begin{figure}[h!] \\begin{center} \\includegraphics[width=7cm]{figure.eps} \\label{figure1} \\caption{An (approximately) uniform system consisting of two sub-systems A and B.} \\end{center} \\end{figure} a uniform system of particles interacting by a pair potential $V(r \\to \\infty) \\sim 1/r^\\gamma$, and divided into two pieces, $A$ and $B$. For the usual thermodynamic classification one can consider that when the potential $V(r)$ is integrable at large $r$, i.e., $\\gamma > d$ where $d$ is the spatial dimension, the potential energy of a typical particle comes essentially from its interaction with particles in a finite region about it. The energy of a particle (e.g. $P$ or $P'$) in $A$ is thus insensitive to whether $B$ is present or not (and thus the total energy is equal, up to surface effects, to the sum of the energies of the subsystems): from an energetic point of view a particle ``does not care'' what the size of the system is, and the interaction is in this sense short-range. The distinction we consider is the analogous one deduced when one reasons in terms of {\\it force} (or acceleration) rather than potential energy, and since forces are the primary physical quantities in dynamics, we refer to the corresponding classification as one of {\\it dynamical} range. It is straightforward to see that in such a system, if the pair force is absolutely integrable, i.e., $\\gamma > d-1$, the force acting on a typical particle is due essentially to its interaction with particles in a finite neighborhood around it, while if $\\gamma < d-1$ this is not the case. Thus in the former case a particle in $A$ ``does not care'' whether the sub-system $B$ is present or not, and in this sense the interaction is ``dynamically short-range\". The classification differs from the standard one for interactions with $d < \\gamma < d-1$: for such interactions the potential energy ``sees the bulk'', but the force, which is its derivative, does not. While the principle motivation for defining such a classification is that it may be relevant to understanding the qualitative behaviors of the out of equilibrium dynamics of such systems, it is {\\it not} the aim of this article to establish that this is the case. We will limit ourselves in this respect to some brief remarks in our conclusions below. Our goal here is to provide a precise formulation of such a classification of the range of interactions based on the convergence properties of forces. While in the usual thermodynamic classification case one considers (see e.g.\\cite{ruelle}) the mathematical properties of essential functions describing systems at equilibrium in the limit $N \\to \\infty$, $V \\to \\infty$ at fixed particle density $n_0=N/V$ (i.e. the usual thermodynamic limit), we will consider the behavior of functions characterising the forces in this same limit. More specifically we consider, following an approach introduced by Chandrasekhar for the case of gravity \\cite{chandra43, book}, the definedness of the probability distribution function (PDF) of the force field in statistically homogeneous particle distributions as the size of the system becomes arbitrarily large. Such distributions and this limit may be described mathematically using the language of stochastic point processes, considering the class of such processes which have a well defined positive mean density when the infinite size limit is considered. To avoid any confusion we will refer to the usual thermodynamic limit in this context simply as the {\\it infinite system limit}. Indeed the existence or non-existence of the quantities we are studying in this limit has no direct relation here to the determination of properties at thermal equilibrium. Further, in the context of the literature on long-range interactions the term ``thermodynamic limit\" is now widely associated with the generalized such limit, which involves adopting a different scaling of $V$ (or possibly coupling constants) with $N$ (for a discussion see e.g. \\cite{balescu}). In this article we also discuss, in Sec.~\\ref{Definedness of dynamics in the infinite system limit}, a further (and different) classification which can be given of pair interactions according to their range. This is relevant when one addresses more generally, for any given pair interaction, a question which arises for Newtonian gravity in a cosmological setting: can a consistent dynamics be defined in an infinite system with non-zero density? A rigorous approach to the same question and the connection with the possibility of defining a {\\em statistical mechanical} state for the system has been developed in \\cite{lanford1,lanford2} in the particular case of a short range non-negative pair potential with finite support. Our conclusion, which generalizes a previous discussion given by two of us in \\cite{1dPoisson_2010}, is that the answer to this question is that a necessary and sufficient condition for such a dynamics is not the integrability of the pair force, but instead the integrability of its gradient. This means that one requires $\\gamma \\geq d-2$, with gravity in any dimension (i.e. $\\gamma=d-2$, the interaction potential solving the appropriate $d$-dimensional Poisson equation) being the marginal case in which such an infinite system limit may be defined. The reason is simply that, in an infinite system without any preferred point (i.e. when this limit is defined respecting statistical translational invariance), the physically meaningful quantity is the relative position of particles (as there is no meaning to absolute position). It is thus the convergence of relative forces on particles with system size which matters. In terms of the schema given above the distinction arises thus when one considers two close-by points (e.g. $P$ and $P'$ in Fig.~1) in sub-system $A$, say, and asks whether their relative forces --- and thus relative motions --- depend on the presence of $B$ or not (or, equivalently, on the size of the system). The answer is that this difference of forces does not essentially depend on $B$ if the gradient of the pair force is absolutely integrable, i.e., $\\gamma > d-2$, as in this case this difference is dominated by the contribution from particles in a finite neighborhood around them. Thus for the case that $d-2 < \\gamma \\leq d-1$ the forces acting on two such particles become ill defined as the size of the system is extended to infinity, but their difference remains finite. Indeed, as has been discussed in \\cite{1dPoisson_2010} in the context of gravity in one dimension, the diverging component of the force on a particle represents a force on their centre of mass, which has no physical relevance in an infinite system without a preferred origin. The paper is organized as follows. In the next section we recall the essential properties of stochastic point processes of relevance to our considerations, and then consider the general analyticity properties of the PDF of the total force at an arbitrary spatial point in such a particle distribution. We show that, for any pair force which is bounded, this PDF in the infinite volume limit is either well defined and rapidly decreasing, or else vanishes pointwise, i.e., the total force is an ill defined stochastic quantity. This means that it suffices, when studying pair potentials with different possible behaviors at large scales, to show that some chosen moment of the PDF converges to a finite value in this limit (or diverges) in order to establish that the whole PDF itself is well-defined (or ill defined). In Sec.~\\ref{PDF-F-mean-variance} we give a general and formal expression for the variance of the total force PDF in a generic infinite uniform stochastic process in terms of the pair force and the two-point correlation properties of the SPP. From this we then deduce our principal result that the force PDF exists strictly in the infinite system limit if and only if the pair force is absolutely integrable at large separations (i.e. $\\gamma \\leq d-1$) , while it can be defined only in a weaker sense, introducing a regularization, when the pair force is not absolutely integrable. In the following section we discuss the physical relevance of the use of such a regularization, which is the generalization of a simple formulation given by Kiessling \\cite{kiessling} of that originally introduced by Jeans for the case of gravity \\cite{jeans}, often misleadingly referred to as the ``Jeans swindle\" \\cite{binney, kiessling}. By analyzing the evolution of density perturbations in an infinite system, we show that the physical relevance of such a regularization of the forces requires also a constraint on the behavior of the PDF of total {\\it force differences} as a function of system size. This leads to the conclusion that $\\gamma \\leq d-2$ is the necessary and sufficient condition in order for it to be possible to have a well defined infinite system limit at constant density for dynamics under a given pair interaction. In the conclusions we review briefly the the relation of our results to previous work in the literature, and comment a little more on the possible relevance of our principle classification of interactions into {\\it dynamically short-range} and {\\it dynamically long-range} to the study of the out of equilibrium dynamics of such systems. ", "conclusions": "In conclusion we make some brief remarks on how the results derived here relate to previous work in the literature on force PDFs. In this context we also discuss the important assumption we made throughout the article, that the pair force considered was {\\it bounded}. Finally we return briefly to the question of the relevance of the classification dividing interactions according to the integrability properties of the pair force, concerning which we have reported initial results elsewhere \\cite{range_PRL2010}. The first and most known calculation of the force PDF is that of Chandrasekhar \\cite{chandra43}, who evaluated it for the gravitational pair interaction in an infinite homogeneous Poisson particle distribution (in $d=3$). This results in the so-called {\\em Holtzmark distribution}, a probability distribution belonging to the Levy class (i.e. power law tailed with a diverging second moment) with $P(\\bF)\\sim F^{-9/2}$ at large $F$. According to our results here, a well defined PDF may be obtained for such a force law, which is {\\it not} absolutely integrable at large separations, only by using a prescription for the calculation of the force in the infinite system limit. In his calculation Chandrasekhar indeed obtains the force on a point by summing the contributions from mass in {\\it spheres} of radius $R$ centered on the point considered, and then taking $R \\rightarrow \\infty$ (with $n_0$ fixed). This prescription is a slight variant of the one we have employed (following Kiessling \\cite{kiessling}): instead of the smooth exponential screening of the interaction, it uses a ``spherical top-hat\" screening so that the force may be written formally as in Eq.~(\\ref{force-full-regularised}) with the replacement of $e^{-\\mu |\\bx-\\bx^\\prime|}$ by a Heaviside function $\\Theta ( \\mu^{-1} - |\\bx-\\bx^\\prime|)$. It is straightforward to verify that the result of Chandrasekhar is unchanged if the smooth prescription Eq.~(\\ref{force-full-regularised}) is used instead. As the Poisson distribution corresponds to an SF $S(k\\rightarrow 0) \\sim k^n$ with $n=0$, the general condition (\\ref{condition-longrange}) for the existence of the PDF we have derived, which gives $n > -1$ for gravity in $d=3$, is indeed satisfied. The fact that the PDF is power-law tailed (and thus {\\it not} rapidly decreasing) arises from the fact that the calculation of Chandrasekhar does not, as done here, assume that the singularity in the gravitational interaction is regularized. Indeed it is simple to show explicitly \\cite{book} that this power law tail arises from the divergence in the pair force at zero separation. This can be done by considering the contribution to the total force on a system particle due to its nearest neighbor particle, which turns out to have a power law tail identical, both in exponent and amplitude, to that of the full $P(\\bF)$. Our analysis shows that it is true in general that well defined, but power-law tailed force PDFs, can arise only when there are singularities in the pair force: for a bounded force we have seen that the PDF is necessarily rapidly decreasing when it exists. More specifically, returning to the analysis of Sec.~\\ref{Analyticity properties of the force PDF}, it is straightforward to see that the crucial property we used of ${\\cal Q}_N(\\{ {\\bf f}_i\\})$, that it have {\\it compact support}, is no longer valid when the pair force has singularities. The analyticity properties which lead to a rapidly decreasing PDF may then not be inferred. We note that this is true at finite $N$, and has nothing to do with the infinite volume limit, i.e., the appearance of the associated power-law tail arises from the possibility of having a single particle which give an unbounded contribution rather than from the combination of the contribution of many particles which then diverges in the infinite system limit. The exponent in such a power-law tail will depend on the nature of the divergence at small separation. More specifically, for a central pair force as considered above and now with a singularity $f(x \\to 0) \\sim 1/x^{a}$, a simple generalization of the analysis for the case of gravity (see \\cite{book}) of the leading contribution to the total force coming from the nearest neighbor particle leads to the conclusion that $P (F \\to \\infty) \\sim F^{-d-\\frac{d}{a} }$ (where $F = |\\bFo|$). This implies that the variance diverges (i.e. the PDF becomes fat-tailed) for $a > d/2$. Force PDFs have been calculated in various other specific cases. Wesenberg and Molmer \\cite{wesenberg+molmer_PDFdipoles_2004} derived that of forces exerted by randomly distributed dipoles in $d=3$, corresponding to a pair force with $\\gamma=2$. According to our results this is the marginal case in which a summation prescription is required for the force, and indeed a prescription using spheres, like that used by Chandrasekhar for gravity, is employed. We note that \\cite{wesenberg+molmer_PDFdipoles_2004} focusses on the power-law tails associated with the singularity at zero separation of the force, which lead in this case (as can be inferred from the result summarized above) to the divergence of the first moment of the force PDF. One of us (AG) has given results previously \\cite{andrea_1dforcesPDF_2005} for the PDF for a generic power-law interaction in $d=1$ for $\\gamma > -1$ in our notation above. The conditional force PDF is then derived for the case of an infinite ``shuffled lattice'' of particles, i.e., particles initially on an infinite lattice and then subjected to {\\it uncorrelated} displacements of finite variance, and using again, as Chandrasekhar, a ``spherical top-hat\" prescription for the force summation (for $\\gamma \\leq 0$, when the pair force is not absolutely integrable). It is simple to show \\cite{book} that such a distribution has an SF with $n=2$ at small $k$, and thus the existence of the force PDF in these cases is again in line with the constraint (\\ref{condition-longrange}) derived. Power-law tails are again observed in these cases, and their exponents related explicitly to the singularity in the assumed power-law force at zero separation. The calculation of Chandrasekhar has been generalized in \\cite{PDF_SL_2006} to the case of particles on an infinite shuffled lattice. This leads again, in line with condition (\\ref{condition-longrange}), to a well defined PDF, again with or without power-law tails according to whether the singularities in the pair force are included or not. Chavanis \\cite{chavanis-forcePDF-poisson_2008} considers, on the other hand, the generalization of Chandrasekhar calculation (for the PDF of gravitational forces in a Poisson distribution) to $d=2$ and $d=1$. The condition (\\ref{condition-longrange} for gravity ($\\gamma=d-2$) gives $n > -d+2$, which implies that the force PDF is not defined in the infinite system limit we have considered for $d \\leq 2$, and indeed in \\cite{chavanis-forcePDF-poisson_2008} well defined PDFs are obtained in $d=2$ and $d=1$ by using a different limiting procedure involving in each case an appropriate rescaling of the coupling with $N$. The physical meaning of such a procedure is discussed in \\cite{1dPoisson_2010}, which considers in detail the calculation of the force PDF for gravity in $d=1$ in a Poisson distribution (as in \\cite{chavanis-forcePDF-poisson_2008}). An exact calculation of the force PDF of the {\\it screened} gravitational force in the infinite system limit is given, which allows one to see in this case exactly how the general result given here is verified in this specific case: all moments of the PDF diverge simultaneously as the screening length is taken to infinity, giving a PDF which converges point-wise to zero. The force PDF for gravity in $d=1$ for a class of infinite particle distributions generated by perturbing a lattice has been derived recently by three of us in \\cite{1dgrav-sl}. It is straightforward to show that one of the conditions imposed on the perturbations to obtain the PDF, that the variance of the perturbations be finite, corresponds in fact to the condition $n>1$ which coincides precisely with the more general condition (\\ref{condition-longrange}) derived here. Unlike in the other specific cases just discussed, it turns out that in this case (gravity in $d=1$) it is in fact necessary to use the smooth prescription Eq.~(\\ref{force-full-regularised}). As explained in detail in \\cite{1dgrav-sl}, the top-hat prescription does not give a well defined result in this case, because surface contributions to the force which do not decay with distance in this case are not regulated by it. We underline that the general result given in the present article are for this specific prescription Eq.~(\\ref{force-full-regularised}). Further analysis would be required to derive the general conditions in which a top-hat prescription also gives the same (and well-defined) PDF. Finally let us comment on why we anticipate the classification of pair interactions according to their ``dynamical range'', formalized here using the force PDF, should be a useful and relevant one physically in the study of systems with long-range interactions. The reason is that this classification reflects, as we have explained, the relative importance of the mean field contribution to the force on a particle, due to the bulk, compared with that due to nearby particles. Now it is precisely the domination by the former which is understood to give the regime of {\\it collisionless} dynamics which is expected to lead to the formation of QSS states, which are usually interpreted to be stationary states of the Vlasov equations describing such a regime of the dynamics (see e.g. \\cite{balescu}). In a recent article \\cite{range_PRL2010} by three of us, we have reported a numerical and analytical study which provides strong evidence for the following result, very much in line with this naive expectation: systems of particles interacting by attractive power law pair interactions like those considered here can always give rise to QSS; however when the pair force is {\\it dynamically short-range} their existence requires the presence of a sufficiently large soft core, while in the {\\it dynamically long-range} case QSS can occur independently of the core, whether hard or soft, provided it is sufficiently small. In other words only in the case of a pair force which is ``dynamically long-range\" can the occurrence of QSS be considered to be the result only of the long distance behavior of the interaction alone. This finding is very consistent with what could be anticipated from the preceding (naive) argument: the effect of a ``soft core'' is precisely to reduce the contribution to the force due to nearby particles, which would otherwise dominate over the mean field force in the case of a pair force which is absolutely integrable at large distances. Indeed the meaning of ``sufficiently large'' specified in \\cite{range_PRL2010} is that the size of the soft core must increase in an appropriate manner with the size of the system as the limit $N \\rightarrow \\infty$ is taken, while we have always implicitly assumed it to be fixed in units of the interparticle distance here. Further work on these issues will be reported elsewhere. \\vskip 0.5 cm We thank M. Kiessling and T. Worrakitpoonpon for useful conversations, and P. Viot for useful comments on the manuscript. \\appendix" }, "1003/1003.5822_arXiv.txt": { "abstract": "Haumea, a rapidly rotating elongated dwarf planet ($\\sim 1500$ km in diameter), has two satellites and is associated with a ``family\" of several smaller Kuiper Belt objects (KBOs) in similar orbits. All members of the Haumea system share a water ice spectral feature that is distinct from all other KBOs. The relative velocities between the Haumea family members are too small to have formed by catastrophic disruption of a large precursor body, which is the process that formed families around much smaller asteroids in the Main Belt. Here we show that all of the unusual characteristics of the Haumea system are explained by a novel type of giant collision: a graze-and-merge impact between two comparably sized bodies. The grazing encounter imparted the high angular momentum that spun off fragments from the icy crust of the elongated merged body. The fragments became satellites and family members. Giant collision outcomes are extremely sensitive to the impact parameters. Compared to the Main Belt, the largest bodies in the Kuiper Belt are more massive and experience slower velocity collisions; hence, outcomes of giant collisions are dramatically different between the inner and outer solar system. The dwarf planets in the Kuiper Belt record an unexpectedly large number of giant collisions, requiring a special dynamical event at the end of solar system formation. ", "introduction": "The four largest dwarf planets in the Kuiper Belt form a distinct population of bodies with high albedos and volatile-rich surfaces \\citep{Schaller2007,Stansberry2008}. A significant history of collisions is suggested by the abundance of satellites in this group, which is much higher than expected for the Kuiper Belt as a whole \\citep{Brown2006}. Three of the four have known satellites: Pluto has three, Haumea (formerly 2003 EL$_{61}$) has two, Eris (2003 UB$_{313}$) has one, and Makemake (2005 FY$_9$) has no substantial satellite \\citep{Brown2008}. The size and orbits of these satellites are different from those found around smaller (100-km size) Kuiper Belt Objects (KBOs). To date, most known satellites around smaller KBOs are thought to have formed via a still-debated capture mechanism \\citep{Knoll2008}. Hence, a different satellite formation process is needed for the dwarf planets, and the most promising mechanism is collisions. Recently, numerical simulations support a giant collision origin for Pluto's massive satellite, Charon \\citep{Canup2005}. However, the formation of the smaller satellites on the other dwarf planets has not been studied in detail. Haumea, a $\\sim 1500$ km diameter classical belt object with a semi-major axis of 43 AU, is a particularly puzzling case as it is also associated with several smaller KBOs with diameters between 70 and 365 km. The smaller KBOs share similar orbits and surface properties. The associated KBOs have been likened to collisionally-produced dynamically and compositionally associated ``families\" that are observed in the asteroid belt \\citep{Brown2007}. We collectively refer to Haumea, its satellites and proposed family members as the Haumea system. Haumea has the only known family in the Kuiper Belt. The Haumea family members share a deep water spectral feature and neutral color \\citep{Brown2007,Ragozzine2007,Schaller2008}. The water feature is unique in the Kuiper Belt \\citep{Brown2007} and indicative of unusually carbon-free water ice \\citep{Pinilla2009}. Haumea also has the distinction of being the only known highly elongated dwarf planet though its precise shape is not known. It has a spin period of only 3.9 hours \\citep{Rabinowitz2006}, the fastest of all the major and dwarf planets. The surface of Haumea is nearly homogenous with the exception of a red spot or faint red hemisphere \\citep{Lacerda2008,Lacerda2009}; hence, the light curve is primarily a reflection of the non-spherical shape. Using the observed light curve and rotation period, \\citet{Rabinowitz2006} fit a density of 2.6 g~cm$^{-3}$ assuming an equilibrium fluid body (a Jacobi ellipsoid). Although the dimensions of Haumea are not yet uniquely constrained \\citep{Lockwood2009}, the observations require a tri-axial shape (see Table~\\ref{tab:sims}) \\citep{Rabinowitz2006}. The derived density is greater than the average of $\\sim 2$ g~cm$^{-3}$ for the largest KBOs \\citep{Brown2008}, although the density may be smaller with some internal friction \\citep{Holsapple2007}. Based on the relatively clean water ice surface and higher than average bulk density, \\citet{Brown2007} argue that Haumea is differentiated with a rocky core and icy mantle. They propose that the family members and satellites are collisionally-derived fragments that originated primarily from the icy mantle. The satellites and family members are orders of magnitude less massive than Haumea, and the family members have a minimum velocity at infinity ($V_{\\infty}$) of about 150 m~s$^{-1}$ \\citep{Ragozzine2007,Ragozzine2009}. The velocity dispersion is much less than expected if the Haumea family formed as the result of a catastrophic impact, as in the formation of asteroid belt families \\citep[e.g., ][]{Nesvorny2006,Michel2004}. In a catastrophic disruption event, a parent body is disrupted and dispersed such that the largest remnant is less than or equal to half the original mass. In the gravity regime, the fragments have initial velocities relative to the largest remnant comparable to the escape speed ($V_{\\rm esc}$) of the disrupted parent body \\citep{Benz1999}. Thus, the observed magnitude of the velocity dispersion ($V_{\\infty}$) of asteroid belt families is significant with respect to $V_{\\rm esc}$ of the largest remnant; in other words, $V_{\\rm esc}$ of the parent body was much greater than $V_{\\rm esc}$ of the largest remnant in gravity dominated disruption events. Unlike most asteroid belt families, the velocity dispersion among Haumea family members is a small fraction of the escape velocity from Haumea ($V_{\\rm esc} \\sim 900$~m~s$^{-1}$). Based on the current models of family formation via catastrophic disruption, the Haumea family could not have formed by catastrophic disruption of a much larger parent body. Two impact scenarios have been proposed for the formation of the Haumea system. \\citet{Brown2007} proposed an impact event that falls in the catstrophic disruption category, which does not agree with the observed velocity dispersion among family members. To explain the small velocity dispersion, \\cite{Schlichting2009} suggest the breakup of a single large moon in orbit around Haumea. However, they do not provide an explanation for the initial state: a large moon in close orbit around a fast-spinning, elongated planet. To date, no known impact scenario explains all of the unusual characteristics of the Haumea system. In this work, we quantitatively model the formation of the Haumea system. We propose that the Haumea family formed via a novel type of giant collision: a {\\it graze and merge} impact between two comparably sized bodies resulting in high angular momentum, which spun off icy fragments that became satellites and family members. The analytic and numerical methods are described in \\S \\ref{sec:method}. The results are presented in \\S \\ref{sec:results}, and the implications for giant impacts in the Kuiper Belt are discussed in \\S \\ref{sec:disc}. ", "conclusions": "In this work, we show that a graze and merge collision event between nearly equal mass bodies is able to produce a symmetrically-elongated, fast-spinning primary, a family of collisional fragments, and multiple bound satellites. The satellites and family members are derived from the icy mantle of the merged, differentiated primary. The family members have a small relative velocity with respect to the primary. This scenario matches all of the observed characteristics of the Haumea system. Our results predict that observations of future family members should all have relative velocities less than $\\sim 0.5 V_{\\mathrm{esc}}$ of Haumea and that the family members should not be isotropically distributed because they formed in a plane. We now have the tools to read the record of giant collisions throughout the solar system. At present, there are several 10's of Pluto-sized bodies known in the outer solar system \\citep{Brown2008}. The new Pan-STARRS observatory will detect up to an order of magnitude more bodies in the Kuiper Belt \\citep{Trujillo2008}, which will complete the Haumea system and test the predictions of the graze and merge scenario. Models of the dynamical evolution of the Kuiper Belt indicate that the population must have been much larger in the past (by a factor of e.g. $\\sim 1000$) \\citep{Morbidelli2008}. Satellite formation was a common outcome of giant impacts, and the abundance of satellites around dwarf planets indicates that giant impacts were frequent. We find that the types of collision events that formed the observed satellites and families in the outer solar system are distinctly different from the Earth's moon formation and families in the asteroid belt. The narrow range of impact parameters that formed the Pluto and Haumea systems place strong constraints on the dynamical history of the largest bodies in the Kuiper Belt." }, "1003/1003.5013_arXiv.txt": { "abstract": "{% We apply ring-diagram analysis and spherical harmonic decomposition methods to compute 3-dimensional power spectra of magnetograms obtained by the Global Oscillation Network Group (GONG) during quiet periods of solar activity. This allows us to investigate the power distribution in acoustic waves propagating in localized directions on the solar disk. We find evidence of the presence of five-minute oscillations in magnetic signals that suggests a non-homogeneous distribution of acoustic power. In this paper, we present our results on the asymmetry in oscillatory power and its behaviour as a function of frequency, time and magnetic field strength. These characteristics are compared with simultaneous velocity measurements.} ", "introduction": "The nature of the propagation of acoustic waves in the solar atmosphere is important for understanding the interaction of these waves with magnetic fields that modify the surface amplitude of the propagating waves. The absorption of acoustic waves by sunspots has been discussed using various data sets and techniques (e.g. Braun, La Bonte and Duvall 1987, Lites et al. 1998, Norton et al. 1999, Ulrich 1996). The studies using the techniques of local-helioseismology have also shown a significant decrease in oscillatory power in localized regions of strong magnetic field (Chou et al. 2009, Howe et al. 2004, Rajaguru, Basu and Antia 2001). The observed modulation in power is also affected by neighboring regions and this ``neighborhood effect'' is discussed in detail by Nicholas, Thompson and Rajaguru (2004). Observations further suggest that the amplitude is suppressed at low-frequencies while it is enhanced around magnetic field regions above 5 mHz (Jain \\& Haber 2002). In this paper, we present characteristics of the oscillations observed in magnetic and velocity signals during quiet solar activity using simultaneous full-disk measurements from Global Oscillations Network Group (GONG). The origin of five-minute oscillations in magnetograms is believed to be cross talk between Doppler velocity and Zeeman splitting. Since the presence of strong magnetic fields modulates the power, quiet periods provide clues to the me-chanism of how the waves propagate in the absence of the field. For a comprehensive comparison of results, two different techniques, ring-diagram analysis and spherical harmonic decomposition, are applied. ", "conclusions": "Using simultaneous Dopplergrams and magnetograms obtained from GONG, we find evidence of the presence of 5-minute oscillations in magnetic signals. The oscillations in these magnetograms are believed to arise due to cross talk between Doppler velocity and Zeeman splitting. The important findings in this work are summarized below; \\begin{enumerate} \\item The signal-to-noise ratio in magnetic rings is low compared to the velocity rings. \\item The anisotropy in magnetic rings appear in two quadrants that correspond to waves propagating in the retrograde direction. Thus, in the quiet sun the effect is mainly due to solar rotation which shifts the frequency of one half of the ring, thereby decreasing its amplitude. \\item The asymmetry parameter varies with frequency and the maximum asymmetry is obtained in 5-min oscillation band. \\item Our analysis do not show any significant correlation between asymmetry parameter and magnetic field. \\item We do find anisotropy in the velocity rings, but it is much less prominent due to high signal-to-noise ratio. \\end{enumerate}" }, "1003/1003.2633_arXiv.txt": { "abstract": "We discuss a novel mechanism of dust acceleration which may dominate for particles smaller than $\\sim0.1~\\mu$m. The acceleration is caused by their direct electrostatic interactions arising from fluctuations of grain charges. The energy source for the acceleration are the irreversible plasma processes occurring on the grain surfaces. We show that this mechanism of charge-fluctuation-induced acceleration likely affects the rate of grain coagulation and shattering of the population of small grains. ", "introduction": "Dust is an important constituent of interstellar medium (ISM), molecular clouds and accretion disks (see Whittet 2003, Draine 2009). It gets involved in many key processes, for instance, it controls heating and cooling of the ISM (see Draine 2003, Tieliens 2005), reveals magnetic fields through grain alignment (see Lazarian 2007 for a review) and interferes with the attempts to measure properties of CMB radiation (see Lazarian \\& Finkbeiner 2003, Fraisse et al. 2009). Small, i.e., less than $\\sim10^{-5}$~cm grains, are an important component of the interstellar dust population, with a notable fraction of very small grains -- Polycyclic Aromatic Hydrocarborn (PAH) particles, which are essentially large molecules\\footnote{The information about these grains is being obtained through the transient heating of grains by UV photons. Indeed, after absorbing a photon a grain gets heated to rather high temperature, inducing emission of infrared photons (see Allamandola et al. 1989). The PAHs were invoked by Draine \\& Lazarian (1998) to explain the anomalous emission of observed in the range of 10-100 GHz (see also Hoang, Draine \\& Lazarian 2010).} (see Leger \\& Puget 1984). In what follows, due to the reasons that are explained in \\S 5, we do not directly address PAH particles, but our approach to the small grain acceleration may be extended to this important population of grains. Most properties of grains, including light extinction, electron photoemission, and chemical activity depend not only on grain chemical composition, but also on their sizes. In astrophysical media, these sizes are affected by grain-grain collisions. The minimal velocities of grains are determined by their Brownian motion corresponding to the temperature of the ambient gas. Large-scale hydrodynamic motions associated with turbulence can make grains move faster (see Draine 1985). Since most astrophysical media are magnetized and grains are charged the hydrodynamic treatment of acceleration is frequently not adequate. A proper treatment of the grain acceleration through the interaction of charged grains with magnetohydrodynamic (MHD) turbulence has been developed recently (Lazarian \\& Yan 2002, Yan \\& Lazarian 2003, Yan, Lazarian \\& Draine 2004, Yan 2009). This treatment makes extensive use of the advances of compressible MHD turbulence (see Cho \\& Lazarian 2002, 2003) and provides the mathematical formalism of the second-order Fermi acceleration of charged grains interacting with MHD turbulence. However, the acceleration mechanisms based on the MHD interaction of turbulence and charged grains exhibit acceleration rates that decrease as dust particles get smaller. This decrease arises from the fact that the Larmor radius of charged grains becomes smaller with the decrease of grain mass and, correspondingly, grains have to interact with smaller, i.e., less powerful, turbulent fluctuations. In addition, compressible fluctuations, i.e., fast modes, which were identified in Yan \\& Lazarian (2003) with the most efficient acceleration, get suppressed at the small scales due to plasma damping, while the Alfvenic mode gets inefficient for acceleration at small scale due to anisotropy (see Yan \\& Lazarian 2003 for more discussion). The acceleration of grains with sizes less than $\\sim10^{-5}$~cm becomes rather inefficient for most media discussed in Yan \\& Lazarian (2003), Yan, Lazarian \\& Draine (2004, henceforth YLD04). Are there other mechanisms of dust acceleration which dominate for grains smaller than $10^{-5}$~cm? One can expect that dust-plasma interactions may be important for such grains. This paper presents a novel promising acceleration mechanism based on grain-grain Coulomb collisions in the presence of grain charge fluctuations. This mechanism utilizes intrinsic non-equilibrium nature of the dusty astrophysical plasmas. In addition to acting on the small grains, the mechanism may drive the acceleration of larger grains whenever the MHD acceleration mechanism is suppressed (e.g., when MHD turbulence is damped). In what follows, we introduce major timescales characterizing dynamics of charged interstellar grains in \\S2, present the mechanism of small grain acceleration associated with charge fluctuations during their Coulomb collisions in \\S3, calculate the acceleration for grains of a given size distribution in \\S4, analyze implications of the effect for interstellar phases in \\S5, discuss the importance of our results in \\S6, and summarize them in \\S7. ", "conclusions": "We showed that there exists a novel powerful mechanism of charge-fluctuation-induced acceleration operating in astrophysical plasmas. This mechanism, which occurs to be more efficient for smaller grains, is based on the charge fluctuations that grains experience during the mutual Coulomb collisions. As the result, the kinetic energy of grains is not conserved in the collisions, which causes the stochastic heating analogous to the second-order Fermi acceleration. This acceleration mechanism is a generic plasma process which can operate in very different environments ranging from the ISM to laboratory gas discharges. It is a remarkable manifestation of the intrinsic non-equilibrium (energetic openness) of dusty plasmas (Tsytovich 1997, Fortov 2005): Even if electrons and ions themselves were in detailed equilibrium, their absorption and the subsequent recombination on dust grains cannot be balanced by the corresponding inverse process (of the ionization and emission from grains) due to its apparent inefficiency -- temperatures necessary for that should be enormously high. Therefore, the process of dust charging goes only in one direction, resulting in the imbalanced plasma flux on the grain surface. The dust dynamics gets coupled to the charging process due to charge fluctuations (which are inevitable due to the discreteness of the elementary charge). Thus, in terms of the energy balance, the mechanism of charge-fluctuation-induced acceleration is based on the conversion of the energy flux (associated with the plasma flux on grains) into the kinetic energy of dust. One should point out that for typical ISM conditions the situation is very different from the equilibrium one: The plasma temperature is much larger that the dust material temperature ($T_{d{\\rm m}}$) due to efficient radiative cooling of the dust. Moreover, the ions and electrons are usually produced by the ISM UV and/or cosmic ray ionization. Grains themselves emit photoelectrons interacting with UV photons. All these non-equilibrium effects can alter charge fluctuations and therefore additionally contribute to our acceleration mechanism. In particular environments, e.g., within quiescent dark clouds and protoplanetary disks with suppressed turbulence, the mechanism we discussed in the paper may be dominant for grains substantially larger than $10^{-5}$~cm. However, we notice the rapid decrease of the mechanism efficiency with the grain size. The treatment presented in this paper is limited in two respects. First, the Fokker-Planck approach presumes relatively small energy variation (\\ref{25}) occurring in each Coulomb collision due to charge fluctuations. If this variation becomes large, such approach (describing the evolution of the velocity distribution function) is no longer applicable. On the other hand, the mean energy evolution -- which is essential for this paper -- is still described by Eq. (\\ref{0}). Second, the direct mechanical encounters of grains cannot be treated within the formalism. When such collisions start dominating (at large $T_d$) the additional dissipation associated with grain deformation, shattering, etc. should be taken into account. In future, we plan to extend the mechanism for the PAH population of grains, explicitly taking into account the substantial variation of grain charge in the process of individual collision with an electron or ion (remind that in this paper we assumed small charge fluctuations, which allowed us to obtain analytically tractable results). The fact that the efficiency of grain acceleration increases with the decrease of grain size makes the process very efficient. Extrapolating our results to PAH particles we may expect their collisions to be very frequent. We note that the additional motivation for the PAH studies comes from the fact that they are responsible for a component of CMB foreground radiation (Draine \\& Lazarian 1998, see also Hoang et al. 2010). Furthermore, we plan to analyze the impact of the charge-fluctuation-induced acceleration on the dust coagulation processes occurring in the ISM and other astrophysical environments, e.g. protostellar accretion disks. It is already clear that the new mechanism can induce coagulation and shattering which have not been considered in the literature before. Together with the MHD turbulence acceleration, our mechanism testifies that the astrophysical grains achieve velocities much larger than those arising from the Brownian motion." }, "1003/1003.3635_arXiv.txt": { "abstract": "We present integral field spectroscopic observations of the central region of the active galaxy NGC 4258 obtained with the fibre IFU system INTEGRAL. We have been able to detect cold neutral gas by means of the interstellar NaD doublet absorption and to trace its distribution and kinematics with respect to the underlying disc. The neutral gas is blue-shifted with projected velocities in the 120--370 km/s range. We have also detected peculiar kinematics in part of the ionized gas in this region by means of a careful kinematic decomposition. The bipolar spatial distribution of the broader component is roughly coincident with the morphology of the X-ray diffuse emission. The kinematics of this gas can be explained in terms of expansion at very high (projected) velocities of up to 300 km/s. The observations also reveal the existence of a strip of neutral gas, parallel to the major kinematic axis, that is nearly coincident with a region of very high [SII]/H$\\alpha$ ratio tracing the shocked gas. Our observations are consistent with the jet model presented by \\cite{wilsonetal01} in which a cocoon originating from the nuclear jet is shocking the gas in the galaxy disc. Alternatively, our observations are also consistent with the bipolar hypershell model of \\cite{Sofue80} and \\cite{SofueandVogler01}. On balance, we prefer the latter model as the most likely explanation for the puzzling features of this peculiar object. ", "introduction": "\\label{intro} The nearby spiral galaxy NGC 4258 (M106) is a particularly interesting galaxy. It presents two peculiarities that have been the cause of many studies of this object: the confirmation of the existence at its centre of a massive black hole, detected in 1995 through measurements of high speed motions in the water masers (\\cite{miyoshi95}), and the existence of the mysterious {\\em anomalous arms} discovered in the optical by \\cite{courtes61} and in radio by \\cite{vdkruit72}. As we shall see below, it is indeed not unlikely that both phenomena are closely related. These {\\em anomalous arms} are two large arm-like gaseous structures of uncertain origin that extend through most of the visible disc of the galaxy. There have been several possible explanations for the origin of these arms, most of which can be traced back to three distinct possibilities: (a) expelled gas from the nucleus in the plane, either in a ballistic (\\cite{vdkruit72}, \\cite{valbada78}, \\cite{valbadavhulst82}) or a jet model (\\cite{fordetal86}, \\cite{ceciletal92}), (b) a bar-induced shock front (\\cite{coxdownes96}), and (c) gas expelled out of the plane either in a jet, hypershell, etc. (\\cite{Sofue80}, \\cite{Sanders82}, \\cite{SofueandVogler01}). There are indeed many observational features that have to be explained by any suitable model (see for example the excellent works by \\cite{ceciletal00} and \\cite{wilsonetal01} including a plethora of optical, radio and X-ray data). The morphology of the anomalous arms can be greatly influenced by projection effects. Although it is pretty well established that the inner parts of these arms are in the disc plane, it is not known for certain whether the arms are three dimensional structures seen in projection (and enhanced by a limb-brightening effect), or if they are indeed mostly contained in a plane (and if this is the case, the orientation of this plane). On the other hand, the highly distorted kinematics that appear in the innermost region of the galaxy also imposes restrictions on the possible scenarios. The galaxy NGC 4258 has an inclination of about 60--64$^\\circ$, and a position angle of 146--160$^\\circ$ (cf. \\cite{jarrettetal03}, \\cite{ceciletal92}). The strong absorption below the major axis shows clearly that the region located SW of the major axis is the near side of the galaxy, the one on the NE being the far side. The most relevant observational facts are: a) the anomalous arms are made up of hot gas that has been shock excited; b) the radio emission of the innermost region shows a clear jet (nearly in the N--S direction), perpendicular to the maser disc around the black hole (The maser disc is actually tipped down some 8$^\\circ$ from edge on [see \\cite{herrnsteinetal97}], meaning that the northern jet is pointing toward us while the southern one is pointing away from us); c) the NW (SE) arm lies behind (in front of) the galaxy disc; and d) the inner section of the anomalous arms is roughly straight and it is roughly aligned with the major axis of the galaxy. From the several scenarios that have been proposed to explain the observational facts, the model that seems to be favoured by the latest observations in X-rays (\\cite{ylwr07}) is the one proposed by {\\cite{wilsonetal01}) in which a jet moving away from the disc induces strong shocks in the disc plane giving rise to the anomalous arms. We have made new observations of the inner region of NGC 4258 with an integral field spectrograph. These new observations will provide new kinematic and morphological information that will permit us to put the available models to the test. ", "conclusions": "We have performed two-dimensional spectroscopic observations in the central region of NGC 4258 and have used these new data to test the available models for this peculiar object. The main results derived from the analysis of these new 2D spectroscopy observations of NGC 4258 are: \\begin{enumerate} \\item The interstellar NaD absorption lines are strongly blueshifted. Comparison of the 2D velocity fields of the NaD with the stellar velocity field reveal that the blueshifted velocities are quite high, in the range $-350$ km/s to $-120$ km/s, and occur everywhere in the detected neutral gas on the near side of the galaxy disc (that is, the one that we can observe in absorption). \\item The velocity map of the ionized gas appears to be strongly distorted. There is a component of ionized gas of bipolar morphology with peculiar, almost counter-rotating, kinematics. \\item The [SII]/H$\\alpha$ and OI/H$\\alpha$ emission line ratio maps show a thin ring-like region of high values indicative of shocks surrounding the galaxy nucleus. The border of the bipolar ionized gas component overlaps this front shock on the near side of the galaxy. \\end{enumerate} The geometry of the bipolar ionized gas component and its partial overlapping with the front shock could fit into the model very well, in that hot gas driven by the jet blasts the ISM disc, giving rise to the anomalous arms. However, the radial motions of the ISM seem hard to explain by this model alone. This outflow, could be naturally interpreted in terms of the presence of the bipolar hypershell model (\\cite{Sofue80} and \\cite{SofueandVogler01}) in that the anomalous arms would be the enhanced emission from the edges of the shells. The lack of symmetry of the bubbles (lack of presence of the anomalous arms in the NE and SW) would be explained by the disruption caused by the jet in the regions where it intersects the bubbles." }, "1003/1003.3918_arXiv.txt": { "abstract": "Kerrr in the title is not a typo. The third ``r'' stands for \\textit{regular}, in the sense of \\textit{pathology-free} rotating black hole. We exhibit a long search-for, exact, Kerr-like, solution of the Einstein equations with novel features: i) no curvature ring singularity; ii) no ``anti-gravity'' universe with causality violating timelike closed world-lines ; iii) no ``super-luminal'' matter disk. \\\\ The ring singularity is replaced by a classical, circular, rotating string with Planck tension representing the inner engine driving the rotation of all the surrounding matter.\\\\ The resulting geometry is \\textit{regular} and \\textit{smoothly} interpolates among inner Minkowski space, borderline deSitter and outer Kerr universe. The key ingredient to cure all unphysical features of the ordinary Kerr black hole is the choice of a ``noncommutative geometry inspired'' matter source as the input for the Einstein equations, in analogy with spherically symmetric black holes described in earlier works. ", "introduction": "Among the several black hole solutions of the Einstein equations, the Kerr geometry is without any doubts the most appropriate to fit the observational data showing that collapsed objects exhibit high angular momenta. Therefore, the complete and through understanding of its properties is crucial for correct description of astrophysical objects. Furthermore, recent expectations are related to the possible outcome of LHC experiments, including production of micro black holes. \\\\ On the other hand, the history of Kerr solution is studded with technical difficulties in solving Einstein equations, accompanied with a complete ignorance of the appropriate matter source. The textbook procedure is based on the so-called ``vacuum solution'' method consisting in assuming an ``ad-hoc`` symmetry for the metric and solving field equations with no source on the r.h.s. Integration constants are then determined comparing the weak-field limit of the solution with known Newtonian-like forms. While mathematically correct this approach is physically unsatisfactory especially in General Relativity, where basic postulate is that geometry is determined by the mass-energy distribution. Furthermore, insisting on the vacuum nature of the solution leads to the presence of curvature singularities where the whole classical theory, i.e. General Relativity, fails. In the Kerr geometry, there are further complications such as: an anti-gravity region and causality violating closed time-like curves. These pathologies should not be present in a physically meaningful gravitational field. A simple way out, is to replace the pathological vacuum region with a regular matter source. In a series of recent papers we have presented a spherically symmetric, regular, matter distribution leading to both neutral and charge black hole solutions with no curvature singularity \\cite{Nicolini:2005zi,Nicolini06,Ansoldi07,Spallucci:2008ez,Nicolini:2008aj,Casadio:2008qy}, \\cite{Nicolini:2009gw,Bleicher:2010qr,Rizzo:2006zb,Gingrich:2010ed} and traversable wormholes \\cite{Garattini:2008xz}. Global structure and inner horizon stability for such a kind of geometries are currently under investigation \\cite{Batic:2010vm,Arraut:2010qx}. The regularity of the metric follows from the presence of a \\textit{minimal length} providing a universal cut-off for short-distance physics. The idea that there should be a minimal distance is supported by many results in different approaches to quantum gravity \\cite{sny,dw,Yoneya:1976pb}, \\cite{Padmanabhan:1986ny,Padmanabhan:1987au,Padmanabhan:1988se}, \\cite{Ashtekar:1992tm,Rovelli:1992vv}, \\cite{Garay:1996sf,Garay:1994en,Fontanini:2005ik,Nicolini:2009dr,Modesto:2009qc}, \\cite{Rinaldi:2009ba,Rinaldi:2010zu}. This new parameter enters the Einstein equations through the energy-momentum tensor, and represents the degree of delocalization of the matter distribution \\cite{Smailagic:2003yb,Smailagic:2003rp,Smailagic:2004yy,Spallucci:2006zj}.\\\\ In this paper we are going to apply the same approach to the axially symmetric problem attempting to remove not only the curvature singularity but all the pathologies quoted above. ", "conclusions": "In this paper we presented the first example of a smooth matter distribution which leads to a pathology-free Kerr solution. The form of the source is a generalization of the corresponding Gaussian mass/energy distribution we introduced for spherically symmetric sources to the case of a rotating object. For both solutions, the same mechanism is at work: the curvature singularity is replaced by a deSitter vacuum domain. In the spherically symmetric case it is an inner deSitter core, while in the Kerrr solution it turns out to be a Saturn-like belt of rotating deSitter vacuum, surrounding an empty Minkowskian disk. The novel feature of the Kerrr solution is that the Minkowski disk joins the deSitter belt through a a rotating string with Planckian (finite!) tension. \\\\ Beside removing the nasty ring singularity the gaussian cloud of matter eliminates the negative $R$ sheet of the Kerr black hole by ensuring analyticity of the metric across the disk. Positivity of $R$ forbids the presence of closed time-like curves.\\\\ To keep the length of the paper short enough to fit the journal format, we must postpone a detailed study of Kerrr black hole thermodynamics to a next article. We anticipate that as in the case of spherically symmetric regular black holes, we find that the Hawking temperature is not unbounded but reaches a maximum value and then drops to zero as the extremal configuration is approached.\\\\ In a forthcoming paper we shall present the extension of the present work to the Kerr-Newman black hole." }, "1003/1003.5241.txt": { "abstract": "We present new mid- to far-infrared images of the two dwarf compact elliptical galaxies that are satellites of M31, NGC~185 and NGC~147, obtained with the Spitzer Space Telescope. Spitzer's high sensitivity and spatial resolution enable us for the first time to look directly into the detailed spatial structure and properties of the dust in these systems. The images of NGC~185 at 8 and 24 micron display a mixed morphology characterized by a shell-like diffuse emission region surrounding a central concentration of more intense infrared emission. The lower resolution images at longer wavelengths show the same spatial distribution within the central 50\\arcsec\\ but beyond this radius, the 160\\mum\\ emission is more extended than that at 24 and 70\\mum. On the other hand, the dwarf galaxy NGC 147 located only a small distance away from NGC~185 shows no significant infrared emission beyond 24 micron and therefore its diffuse infrared emission is mainly stellar in origin. For NGC~185, the derived dust mass based on the best fit to the spectral energy distribution is $1.9 \\times 10^3 \\rm{M}_\\odot$, implying a gas mass of $3.0 \\times 10^5 \\rm{M}_\\odot$. These values are in agreement with those previously estimated from infrared as well as CO and HI observations and are consistent with the predicted mass return from dying stars based on the last burst of star formation $1 \\times 10^9$~yr ago. Based on the 70 to 160\\mum\\ flux density ratio, we estimate a temperature for the dust of $\\sim$~17K. For NGC~147, we obtain an upper limit for the dust mass of $4.5 \\times 10^{2} \\rm{M}_\\odot$ at 160\\mum\\ (assuming a temperature of $\\sim$~20K), a value consistent with the previous upper limit derived using ISO observations of this galaxy. In the case of NGC~185, we also present full $5-38$\\mum\\ low-resolution (R$\\sim$100) spectra of the main emission regions. The IRS spectra of NGC~185 show strong Polycyclic Aromatic Hydrocarbons (PAH) emission, deep silicate absorption features and H$_2$ pure rotational line ratios consistent with having the dust and molecular gas inside the dust cloud being impinged by the far-ultraviolet radiation field of a relatively young stellar population. Therefore, based on its infrared spectral properties, NGC~185 shows signatures of recent star formation (a few $\\times 10^8$ years ago), although its current star formation rate is quite low. ", "introduction": "It was by resolving into stars the two dwarf elliptical galaxies NGC~147 and NGC~185, along with their formidable large companion the Andromeda galaxy, that Baade refined his concept of two stellar populations in the 1940s. Of NGC~147, Baade (1944) wrote ``In contrast with the uncertain information provided by the blue-sensitive plates, the 4-hour red exposure of NGC~147 is truly revealing. It shows that the nebula is a large star cloud, ellipsoidal in structure, and of a density gradient so low that even the central region is fully resolved''. He later on remarked that NGC~147 is an example of a dust-free galaxy containing only stars of Population II (Baade 1951). As part of the same work, Baade's notes on NGC~185 revealed ``NGC~185, considerably brighter than NGC~147, has been classified by Hubble as Ep, the peculiarity being the abnormally slow increase in intensity toward the center. It is intermediate in this respect to NGC~205, similarly classified by Hubble as Ep, and NGC~147. NGC~185 is one of the few elliptical nebulae in which patches of obscuring material are conspicuous. Two such dark clouds are near the center of NGC~185''. Based on his observation of bright young blue stars in NGC~185, he also noted the apparent relationship between these Population I (metal-rich) stars and the presence of dust. Based on our current understanding of stellar evolution, we now know that it is the dying stars that inject gas and metal-enriched dust back into the interstellar medium (ISM) (Faber \\& Gallagher 1976). The presence of planetary nebulae in NGC~185 (Ford, Jenner \\& Epps 1973), for example, confirms that mass loss is occurring in this dwarf elliptical galaxy. From this enriched material is formed a new population of luminous hot and young metal-rich stars. Indeed, the ISM in NGC~185, as seen in the form of atomic (HI) and molecular gas (CO), appears spatially concentrated near the present-day star-forming region (Young 2001; Welch, Mitchell \\& Yi 1996). However, whereas NGC~185 contains gas, dust and young stars with ages of about 100~Myr (in the central 150~pc by 90~pc region; Martinez-Delgado, Aparicio, \\& Gallart 1999) to 400~Myr (Butler \\& Martinez-Delgado 2005), it is unknown why NGC~147 appears to be presently dust- and gas-free (Young \\& Lo 1997; Sage, Welch, \\& Mitchell 1998) and had its most recent star forming episode $\\lesssim$~1~Gyr ago (Han et al.\\ 1997). Either an efficient removal mechanism of some kind is at play (external sweeping, formation of condensed objects, gas ejection by violent nuclear activity, or sweeping by a hot galactic wind) or NGC~147 is at a completely different evolutionary stage than NGC~185. Although there are some claims that these two dwarf galaxies form a stable binary system (van den Bergh 1998), the latter explanation implies that these two galaxies did not form at the same time and/or in the same environment despite their physical proximity. In this paper, we explore the nature of these two systems from a completely new direction, using a different wavelength regime. We present new mid- to far-infrared images of NGC~147 and NGC~185 obtained with the {\\it Spitzer Space Telescope} (hereafter {\\it Spitzer}). The higher sensitivity and spatial resolution of {\\it Spitzer} over previous infrared observatories enable us for the first time to measure {\\it directly} the detailed structure and composition of the dust in these two dwarf galaxies. The paper is organized as follows. In Section~2, we describe the {\\it Spitzer} observations and the data reduction. The {\\it Spitzer} images are presented in Section~2, followed by a description of the mid- and far-infrared emission morphology and a comparison with other tracers of the ISM in Section~3. Foreground point source contamination is discussed in Section~4. Dust mass measurements and gas mass estimates are derived in Section~6 based on the spectral energy distribution of the dwarf galaxies presented in Section~5. In Section~7, we present the results of the analysis of the first set of {\\it Spitzer} spectroscopic observations of one of M31's dwarf galaxies. We discuss the spectral line measurements and spectral properties of NGC~185's dust clouds. We conclude our paper by presenting a summary of our observations and results in Section~8. ", "conclusions": "New IRAC, MIPS and IRS observations of NGC 147 and NGC~185 give a better assessment of the dust content and properties in these prototypical local dwarf galaxies. Spitzer's high sensitivity and spatial resolution enable us for the first time to look directly into the detailed spatial structure and properties of the dust in these systems. The images of NGC~185 at 8 and 24 micron display a mixed morphology characterized by a shell-like diffuse emission region surrounding a central concentration of more intense infrared emission. The lower resolution images at longer wavelengths show the same spatial distribution within the central 50\\arcsec\\ but beyond this radius, the 160\\mum\\ emission is more extended than that at 24 and 70\\mum. On the other hand, the dwarf galaxy NGC 147 located only a small distance away from NGC~185 shows no significant infrared emission beyond 24 micron and its diffuse infrared emission is mainly stellar in origin. For NGC~147, we obtain an upper limit for the dust mass of $4.5 \\times 10^{2} \\rm{M}_\\odot$, a value consistent with the previous upper limit derived using ISO observations of this galaxy. For NGC~185, the derived dust mass based on the best fit to the ``Total'' SED is $1.9 \\times 10^3 \\rm{M}_\\odot$, implying a gas mass of $3.0 \\times 10^5 \\rm{M}_\\odot$ (assuming a standard gas-to-dust mass ratio of 160). These values are in agreement with those previously estimated from infrared as well as CO and HI observations. The gas estimate is also consistent with the predicted mass return from dying stars, based on the last burst of star formation, $1 \\times 10^9$~yr ago. Based on simply the 70 to 160\\mum\\ flux density ratio, the estimated temperature for the dust is $\\sim$~17K. The fact that NGC~147 resembles more a ``typical'' dust and gas-free elliptical galaxy than NGC~185 (and NGC~205, for that matter, see Marleau et al.\\ 2006) remains puzzling in the context of the possible binary system scenario. In the case of NGC~185, we also presented full $5-38$\\mum\\ low-resolution (R$\\sim$100) spectra of the main emission regions. The IRS spectra of NGC~185 show strong PAH emission, deep silicate absorption features and H$_2$ pure rotational line ratios consistent with having the dust and molecular gas inside the dust cloud being impinged by the far-ultraviolet radiation field of a relatively young stellar population. Although the current rate of star formation is quite low ($\\sim 10^{-10} \\rm{M}_\\odot$/yr), this suggests that the star formation history of NGC~185 is complex, perhaps as much as that of the more active NGC~205 (Monaco et al.\\ 2009)." }, "1003/1003.1892_arXiv.txt": { "abstract": "Calibration of a sensor array is more involved if the antennas have direction dependent gains and multiple calibrator sources are simultaneously present. We study this case for a sensor array with arbitrary geometry but identical elements, i.e. elements with the same direction dependent gain pattern. A weighted alternating least squares (WALS) algorithm is derived that iteratively solves for the direction independent complex gains of the array elements, their noise powers and their gains in the direction of the calibrator sources. An extension of the problem is the case where the apparent calibrator source locations are unknown, e.g., due to refractive propagation paths. For this case, the WALS method is supplemented with weighted subspace fitting (WSF) direction finding techniques. Using Monte Carlo simulations we demonstrate that both methods are asymptotically statistically efficient and converge within two iterations even in cases of low SNR. ", "introduction": "In this paper we study the calibration of the direction dependent response of sensor array antennas, excited by simultaneously present calibrator sources. The antenna array has arbitrary geometry but identical antennas. The calibration involves the complex gain of the antenna elements towards each source. The source powers are known but we will allow for small deviations in apparent source locations to account for, e.g., refractive propagation paths. This problem is one of the main challenges currently faced in the field of radio astronomy. For low frequency observations ($<$ 300 MHz) this community is building or developing a number of new instruments, for example the low frequency array (LOFAR) \\cite{Bregman2005-1}, the Murchison wide field array (MWA) \\cite{Lonsdale2008-1} and the primeval structure telescope (PaST) \\cite{Peterson2005-1}, which are all large irregular phased arrays. These difficulties arise due to the influence of the ionosphere on the propagation of radio waves, which can qualitatively be categorized into the following four regimes depending on the field of view (FOV) of the individual receptors and the baselines between them \\cite{Lonsdale2004-1}: \\begin{figure} \\begin{center} \\input fig1 % \\end{center} \\caption{Graphical representation of regime 3: all lines of sight towards a single source sample the same ionosphere, but the ionospheric delay differs per source due to the large FOV and short baselines (after \\cite{Lonsdale2004-1}).} \\label{fig:lonsdale3} \\end{figure} \\begin{enumerate} \\item All antennas and all lines of sight sample the same ionospheric delays, thus the ionosphere causes no distortion of the array manifold (small FOV, short baselines). \\item Lines of sight from different antennas sample different ionospheric patches, but all sources in the antenna FOV experience the same delay; the ionosphere thus causes an antenna based gain effect (small FOV, long baselines). \\item All lines of sight towards a single source sample the same ionosphere but the delay differs per source. This regime requires source based direction dependent calibration (large FOV, short baselines, see Fig. \\ref{fig:lonsdale3}). \\item The ionospheric delays differ per station and per source (large FOV, long baselines). \\end{enumerate} The first two regimes can be handled under the self-calibration assumption used in astronomy that all errors are telescope based thereby reducing the estimation problem to a single direction independent gain factor per telescope \\cite{Cornwell1994-1}. The problem of estimating these direction independent gains based on a single calibrator source is treated in \\cite{Boonstra2003-1}, while the multiple source case has been discussed in \\cite{Wijnholds2006-1}. The problem of finding a complex gain per antenna per source, the fourth regime, is not tractable without further assumptions which allow to parameterize the behavior of the gains over space, time and/or frequency \\cite{Tol2007-1}. In this paper we will focus on the third regime, which is sketched in Fig. \\ref{fig:lonsdale3}, thus filling the gap in the available literature. It allows for the calibration of individual closely packed groups of antennas such as a LOFAR station or a subarray of the MWA and PaST telescopes and thus forms a valuable step towards the calibration of the whole array. We will state our problem in general terms, since the problem plays a role in a range of applications varying from underwater acoustics to antenna arrays. This explains the persistent interest in on-line calibration, autocalibration or self-calibration of sensor arrays \\cite{Flanagan1999-1, Astely1999-1, Pesavento2002-1, See2004-1, Tol2007-1}. In most applications the main driver for studies on array calibration is to improve the DOA estimation accuracy. Many studies in this field therefore try to solve for the DOAs and a number of array parameters. In \\cite{Weiss1989-1, See1995-1, Flanagan1999-1} self-calibration schemes are presented which solve for the direction independent gains and the sensor positions, i.e. the directional response of the sensors is assumed known. Other studies assume a controlled environment to calibrate the array by measuring the calibrator sources one at a time \\cite{Pierre1991-1, Ng1996-1} or exploit the array geometry, e.g. a uniform linear array (ULA) having a Toeplitz matrix as array covariance matrix \\cite{Astely1999-1, Li1999-1}. Weiss and Friedlander \\cite{Weiss1995-1} have presented a technique for almost fully blind signal estimation. Their work, however, focuses on estimation and separation of source signals, not on characterizing the array itself. We thus feel that the problem at hand also forms an interesting addition to the literature available on sensor array calibration in general. In the next section we introduce the data model and provide a mathematical formulation of the problem. In Sec.\\ \\ref{sec:analysis} the Cram\\`er-Rao bound for this estimation problem is discussed. Section \\ref{sec:algorithms} presents an alternating least squares (ALS) and a weighted alternating least squares (WALS) approach optimizing subsets of parameters iteratively. These algorithms are validated using Monte Carlo simulations in Sec.\\ \\ref{sec:simulations}. The conclusions are drawn in Sec.\\ \\ref{sec:conclusions}. \\emph{Notation}: Overbar $\\overline{(\\cdot)}$ denotes conjugation, the transpose operator is denoted by $^T$, the complex conjugate (Hermitian) transpose by $^H$ and the pseudo-inverse by $^\\dagger$. An estimated value is denoted by $\\expect \\{ \\cdot \\}$. $\\odot$ is the element-wise matrix multiplication (Hadamard product), $\\oslash$ is the element-wise matrix division, $\\otimes$ denotes the Kronecker product and $\\circ$ is used to denote the Khatri-Rao or column-wise Kronecker product of two matrices. $\\diag(\\cdot)$ converts a vector to a diagonal matrix with the vector placed on the main diagonal, $\\vecdiag(\\cdot)$ produces a vector from the elements of the main diagonal of its argument and $\\vec(\\cdot)$ converts a matrix to a vector by stacking the columns of the matrix. We exploit many properties of Kronecker products, including the following (for matrices and vectors of compatible dimensions): \\begin{eqnarray} \\vec(\\mathbf{A} \\mathbf{B} \\mathbf{C}) &=& (\\mathbf{C}^T \\otimes \\mathbf{A}) \\vec(\\mathbf{B}) \\label{eq:prop1} \\\\ \\vec(\\mathbf{A} \\diag(\\mathbf{b}) \\mathbf{C}) &=& (\\mathbf{C}^T \\circ \\mathbf{A}) \\mathbf{b} \\label{eq:prop2} \\\\ (\\mathbf{A} \\circ \\mathbf{B})^H (\\mathbf{C} \\circ \\mathbf{D}) &=& \\mathbf{A}^H \\mathbf{C} \\odot \\mathbf{B}^H \\mathbf{D} \\label{eq:prop3} \\end{eqnarray} \\section {Data model} \\label{sec:datamodel} We consider an array of $p$ elements located at a single site (``Regime 3''). Denote the complex baseband output signal of the $i$th array element by $x_i(t)$ and define the array signal vector $\\mathbf{x}(t) = [x_1(t), x_2(t), \\cdots, x_p(t)]^T$. We assume the presence of $q$ mutually independent i.i.d.\\ Gaussian signals $s_k(t)$ impinging on the array, which are stacked in a $q \\times 1$ vector $\\mathbf{s}(t)$. Likewise the sensor noise signals $n_i(t)$ are assumed to be mutually independent i.i.d.\\ Gaussian signals and are stacked in a $p \\times 1$ vector $\\mathbf{n}(t)$. If the narrow band condition holds \\cite{Zatman1998-1}, we can define the $q$ spatial signature vectors $\\mathbf{a}_k$ ($p \\times 1)$, which describe for the $k$th source the phase delays at each antenna due only to the geometry. The $q$ sources are considered calibrator sources, thus we assume that their powers, their nominal positions, and the locations of the antennas (hence $\\mathbf{a}_k$) are known. Refractive effects caused by ionospheric phase gradients may shift the apparent locations of the sources, thus we will also consider cases where the $\\mathbf{a}_k$ are only parametrically known. The sensors are assumed to have the same direction dependent gain behavior towards the $q$ source signals received by the array. This can include the antenna pattern and ionospheric phase effects. They are described by gain factors $g_{0k}$ and are collected in a matrix $\\mathbf{G_0} = \\diag([g_{01}, g_{02}, \\cdots, g_{0q}])$. The direction independent gains and phases (individual receiver gains) can be described as $\\bgamma = [\\gamma_1, \\gamma_2, \\cdots, \\gamma_p]^T$ and $\\bphi = [\\e^{\\jcmplx\\phi_1}, \\e^{\\jcmplx\\phi_2}, \\cdots, \\e^{\\jcmplx\\phi_p}]^T$ respectively with corresponding diagonal matrix forms $\\bGamma = \\diag(\\bgamma)$ and $\\bPhi = \\diag(\\bphi)$. With these definitions, the array signal vector can be described as \\begin{equation} \\mathbf{x}(t) = \\bGamma\\bPhi \\left ( \\sum_{k=1}^q \\mathbf{a}_k g_{0k} s_k(t) \\right ) + \\mathbf{n}(t) = \\mathbf{G} \\mathbf{A} \\mathbf{G}_0 \\mathbf{s}(t) + \\mathbf{n}(t)\\label{eq:sigvec} \\end{equation} where $\\mathbf{A} = [\\mathbf{a}_1, \\cdots, \\mathbf{a}_k]$ (size $p \\times q$) and $\\mathbf{G} = \\bGamma \\bPhi$; for later use we also define $\\bg = \\bgamma \\odot \\bphi$. The signal is sampled with period $T$ and $N$ sample vectors are stacked into a data matrix $\\mathbf{X} = [\\mathbf{x}(T), \\mathbf{x}(2T), \\cdots, \\mathbf{x}(NT)]$. The covariance matrix of $\\mathbf{x}(t)$ is $\\mathbf{R} = \\expect \\{ \\mathbf{x}(t) \\mathbf{x}^H(t) \\}$ and is estimated by $\\widehat\\mathbf{R} = N^{-1} \\mathbf{X} \\mathbf{X}^H$. Likewise, the source signal covariance $\\bSigma_s = \\diag(\\bsigma_s)$ where $\\bsigma_s = [\\sigma_{s1}^2, \\sigma_{s2}^2, \\cdots, \\sigma_{sq}^2]^T$ and the noise covariance matrix is $\\bSigma_n = \\diag(\\bsigma_n)$ where $\\bsigma_n = [\\sigma_{n1}^2, \\sigma_{n2}^2, \\cdots, \\sigma_{np}^2]^T$. Then the model for $\\mathbf{R}$ based on \\eqref{eq:sigvec} is \\begin{equation} \\mathbf{R} = \\mathbf{G} \\mathbf{A} \\mathbf{G}_0 \\bSigma_s \\mathbf{G}_0^H \\mathbf{A}^H \\mathbf{G}^H + \\bSigma_n.\\label{eq:R} \\end{equation} In this model, $\\bSigma_s$ is known. Since $\\mathbf{G}_0$ and $\\bSigma_s$ are diagonal matrices, we can introduce \\begin{eqnarray} \\bSigma & = & \\mathbf{G}_0 \\bSigma_s \\mathbf{G_0}^H \\nonumber\\\\ & = & \\diag([ |g_{01}|^2\\sigma_{s1}^2, \\cdots, |g_{0q}|^2\\sigma_{sq}^2]) = \\diag(\\bsigma) \\,. \\end{eqnarray} Since the direction dependent gains $g_{0k}$ are not known, the real valued elements of $\\bsigma = [\\sigma_1^2, \\sigma_2^2, \\cdots, \\sigma_q^2]^T$ are not known either. This implies that our problem is identical to estimating an unknown diagonal signal covariance matrix without any DOA dependent errors. We may thus restate \\eqref{eq:R} as \\begin{equation} \\mathbf{R} = \\mathbf{G} \\mathbf{A} \\bSigma \\mathbf{A}^H \\mathbf{G}^H + \\bSigma_n\\label{eq:R2} \\end{equation} and solve for $\\mathbf{G}$, $\\bSigma$ and $\\bSigma_n$ under the assumption that $\\mathbf{A}$ is known (or parametrically known). The data model described by \\eqref{eq:R2} is commonly used in papers on sensor array calibration (e.g. \\cite{Fuhrmann1994-1, Boonstra2003-1}). Flanagan and Bell \\cite{Flanagan1999-1}, Weiss and Friedlander \\cite{Weiss1989-1} and See \\cite{See1995-1} effectively use the same model, but focus on position calibration of the array elements and are therefore more explicit on the form of $\\mathbf{A}$. If the source positions and locations of the sensors within the array are known, an explicit formula for $\\mathbf{a}_k$ can be used to compute the nominal spatial signature vectors. Estimation of source locations is required to account for the refractive effects produced by an ionospheric phase gradient. The $i$th element of the array is located at $\\mathbf{r}_i = [x_i, y_i, z_i]^T$. These positions can be stacked in a matrix $\\mathcal{R} = [\\mathbf{r}_1, \\mathbf{r}_2, \\cdots, \\mathbf{r}_p]^T$ (size $p \\times 3$). The position of the $k$th source can be denoted as $\\mathbf{l}_k = [l_k, m_k, n_k]^T$. The source positions can be stacked in a matrix $\\mathcal{L} = [\\mathbf{l}_1, \\mathbf{l}_2, \\cdots, \\mathbf{l}_q]^T$ (size $q \\times 3$). The spatial signature matrix $\\mathbf{A}$ can thus be described by \\begin{equation} \\mathbf{A} = \\exp \\left ( -\\jcmplx \\frac{2\\pi}{\\lambda} \\mathcal{R} \\mathcal{L}^T \\right ) \\end{equation} where the exponential function is applied element-wise to its argument. In the remainder of this paper we will specialize to a planar array having $z_i = 0$ for convenience of presentation but without loss of generality. From \\eqref{eq:R2} we observe that $\\mathbf{G}$ and $\\bSigma$ share a common scalar factor. Therefore we may impose the boundary condition $\\sigma_1^2 = 1$. To solve the ambiguity in the phase solution of $\\mathbf{G}$ we will take the first element as phase reference, i.e. $\\phi_1 = 0$ is imposed.\\footnote{ In \\cite{Wijnholds2006-2} it is shown that $\\sum_{i=1}^p \\phi_i = 0$ is the optimal constraint for this problem. This constraint has the disadvantage that the location of the phase reference is not well defined. Furthermore, the choice for the constraint used here simplifies our analysis in combination with the constraints required to uniquely identify the source locations and the apparent source powers.} When solving for source locations a similar problem occurs: a single rotation of all DOA vectors can be compensated by the direction independent gain phase solution. We will therefore fix the position of the first source. In this paper we will address four related sensor array calibration problems based on this data model. These are summarized below where the parameter vectors adhering to the aforementioned boundary conditions are stated explicitly. \\begin{enumerate} \\item The sensor noise powers are the same for all elements, i.e.\\ $\\sigma_{n1}^2 = \\sigma_{n2}^2 = \\cdots = \\sigma_{nq}^2$ such that $\\bSigma_n = \\sigma_n^2 \\mathbf{I}$ where $\\mathbf{I}$ is the identity matrix. In this scenario the parameter vector to be estimated is $\\btheta = [\\bgamma^T, \\phi_2, \\cdots \\phi_p, \\sigma_2^2, \\cdots \\sigma_q^2, \\sigma_n^2]^T$. \\item The sensor noise powers are allowed to differ from one element to the other, i.e.\\ $\\bSigma_n = \\diag(\\bsigma_n)$. In this case the parameter vector is $\\btheta = [\\bgamma^T, \\phi_2, \\cdots, \\phi_p, \\sigma_2^2, \\cdots \\sigma_q^2, \\sigma_{n1}^2, \\cdots \\sigma_{np}^2]^T$. \\item $\\bSigma_n = \\sigma_n^2 \\mathbf{I}$ and $\\mathbf{A} = \\mathbf{A}(\\mathcal{L})$, i.e.\\ similar to the first scenario but with unknown source locations. In this case $\\btheta = [\\bgamma^T, \\phi_2, \\cdots, \\phi_p, \\sigma_2^2, \\cdots, \\sigma_q^2, \\sigma_n^2, \\mathbf{l}_2^T, \\cdots, \\mathbf{l}_q^T]^T$. \\item $\\bSigma_n = \\diag(\\bsigma_n)$ and $\\mathbf{A} = \\mathbf{A}(\\mathcal{L})$, giving $\\btheta = [\\bgamma^T, \\phi_2, \\cdots, \\phi_p, \\sigma_2^2, \\cdots, \\sigma_q^2, \\sigma_{n1}^2, \\cdots, \\sigma_{np}^2, \\mathbf{l}_2^T, \\cdots, \\mathbf{l}_q^T]^T$. \\end{enumerate} ", "conclusions": "\\label{sec:conclusions} In this paper we have developed a weighted alternating least squares (WALS) algorithm to solve for direction independent complex gains, apparent source powers and receiver noise powers simultaneously in a snapshot observation by a sensor array. Our solution for finding the direction independent gains extends and improves the available methods in the literature in the case of multiple calibrator sources and provides robustness to small entries in the array covariance matrix and small gain values (due to e.g.\\ failing elements). Although we have assumed independent sensor noise powers, our gain estimation method is easily applied to cases in which some pairs of elements experience correlated noise. We also found that unbiased estimation of the receiver noise powers requires weighting with the best available array covariance matrix model instead of weighting with the measured array covariance matrix, which is common practice in signal processing. The statistical performance was compared with an unweighted alternating least squares (ALS) algorithm and the CRB and found to provide a statistically efficient estimate thereby outperforming the ALS method. The computational complexity of ALS and WALS scale with $4\\frac{1}{2}p^3$ and about $7p^3$ respectively assuming that the number of calibrator sources is much smaller than the number of array elements, i.e.\\ the WALS method comes with only a minor increase in computational burden. Simulations indicate that only one or two iterations are sufficient to reach the CRB and that every iteration adds about one significant digit, the WALS method converging slightly faster than the ALS method. We extended these methods with source location estimation using weighted subspace fitting. Simulations indicate that this extension to the WALS algorithm provides a statistically efficient simultaneous estimate of omnidirectional gains, apparent source powers, source locations and receiver noise powers. Again, one to two iterations proved to be sufficient to reach the CRB. However, the convergence of the parameter is blocked at some level lower than the CRB by the noise in the measurement, at which point the algorithm keeps jumping from one local optimum to another." }, "1003/1003.4872_arXiv.txt": { "abstract": "The surface brightness profile of H$\\alpha$ emission in galaxies is generally thought to be confined by a sharp truncation, sometimes speculated to coincide with a star formation threshold. Over the past years, observational evidence for both old and young stellar populations, as well as individual H II regions, has demonstrated that the outer disk is an actively evolving part of a galaxy. To provide constraints on the origin of the aforementioned H$\\alpha$ truncation and the relation of H$\\alpha$ emission in the outer disk to the underlying stellar population, we measure the shape of the outer H$\\alpha$ surface brightness profile of 15 isolated, edge-on late-type disk galaxies using deep, long-slit spectroscopy. Tracing H$\\alpha$ emission up to 50\\% beyond the optical radius, $R_{25}$, we find a composite H$\\alpha$ surface brightness profile, well described by a broken-exponential law, that drops more steeply in the outer disk, but which is not truncated. The stellar continuum and H$\\alpha$ surface brightness both exhibit a break at $\\sim0.7 R_{25}$, but the H$\\alpha$ emission drops more steeply than the stellar continuum beyond that break. Although profiles with truncations or single exponential laws correctly describe the H$\\alpha$ surface brightness profiles of some individual galaxies, flexible broken-exponentials are required in most cases and are therefore the more appropriate generic description. The common existence of a significant second surface brightness component beyond the H$\\alpha$ break radius disfavors the hypothesis that this break is a purely stochastic effect. ", "introduction": "\\label{sec:intro} The outer disks of galaxies are increasingly being recognized as a diverse environment whose properties may aid us in understanding the formation of galaxies. Since the radio observations of the 1950s \\citep{vandehulst,dieter}, we have known that the baryonic component of spiral galaxies, as traced by neutral hydrogen, is significantly larger than suggested by classical size indicators such as $R_{25}$ (the radius where the surface brightness drops to 25 mag arcsec$^{-2}$; we use measurements of the $B$-band surface brightness obtained from the literature throughout this analysis). Optical studies, however, have focused on the inner parts of galaxies, where, among other aspects, the radial distribution of H$\\alpha$ emission has been subject of intense scrutiny for many years. For example, \\citet{martinkennicutt} used narrow-band imaging of nearby galaxies to construct H$\\alpha$ surface brightness profiles. Such profiles usually exhibit a ``truncation\" in the H$\\alpha$ surface brightness at relatively large radii, although usually still within the optical radius $R_{25}$. This truncation is also reflected in the scarcity of H$\\alpha$ rotation curves that extend to radii much beyond $\\sim 0.7 R_{25}$ in comprehensive long-slit surveys (e.g., \\citet{vogt}, being one of the deepest such surveys, observe no rotation curves beyond $R_{25}$ in a sample of 329). A widely accepted hypothesis for such a truncation is the existence of a threshold in the surface gas density, below which star formation becomes inefficient \\citep{kennicutt89,martinkennicutt}. Alternate explanations include the hypothesis that the truncation is indicative of an actual break in the mass distribution that is related to the initial formation conditions of the disk \\citep{vanderkruit87}. However, ``truncation\" may be too strong a descriptive statement because it is now known that low-level star formation occurs at large galactocentric radii, well beyond the radius corresponding to the critical threshold. \\citet{blandhawthorn,ferguson1998a,ferguson1998b} have detected individual H$\\alpha$ emission regions in the outer disks of spiral galaxies as far as 2 $R_{25}$, and UV emission appears relatively common in outer disks \\citep{thilker,zaritskychristlein,thilker2}. Furthermore, in a study of NGC 3814 using deep two-band optical imaging, \\citet{herbertfort} have found statistically significant overdensities of marginally resolved sources in the outer disk that are likely to be star clusters. This outer-disk star formation has raised new interest in the hypothesis of a star formation threshold. The absence of evidence for a break in the surface brightness profiles of UV emission, which is also a star formation indicator, led \\citet{boissier,boissier2} to suggest that the star formation rate does not exhibit a break, and that the H$\\alpha$ truncation is a stochastic effect. As the expectation number of star formation regions with stars massive enough to generate a Stroemgren sphere drops below unity, the H$\\alpha$ emission profile goes to zero, while low level(mass) SF continues as measured with the UV. Other, perhaps more exciting, possible explanations include a change in the initial mass function \\citep{meurer}. These arguments depend critically on reliable measures of H$\\alpha$ and continuum profiles out to large radii. At the same time, studies of faint optical continuum emission have shown that stellar disks can also be traced far into the outer disk. While some galaxies exhibit a single exponential surface brightness profile to the largest measurable radii \\citep{blandhawthorn2005}, others are more accurately described by broken exponential profiles with a characteristic break radius, beyond which the stellar surface brightness profile may be shallower or steeper than in the inner disk \\citep{pohlen,erwin}. These profiles are sometimes described as exponential, sub-exponential, and super-exponential \\citep{vlajic2009}; the exponential and sub-exponential types are also variously referred to as Freeman Types I and II \\citep{freeman}, and the designation ``Type III\" has come into usage to describe the super-exponential (up-bending) shape. How these outer stellar structures are related to possible structures in the gaseous disk, and in particular, H$\\alpha$ emission, is a crucial question in understanding the origin of the outer disk stellar populations. Is the truncation in H$\\alpha$ emission a real indicator of a truncation in star formation, or, as \\citet{boissier,boissier2} suggest, merely a stochastic effect? If, on the other hand, the break in the star formation surface density is real, is it also responsible for the characteristic break in the stellar continuum profile? In that case, is there significant star formation beyond the break, and is it enough to have created the outer disk stellar content in situ? Or must other processes be invoked to populate the outer disks with stars? To answer these questions, a systematic and quantitative study of the distribution of H$\\alpha$ emission at large radii is required. The traditional technique of narrow-band imaging with subsequent subtraction of broad-band continuum emission is generally insufficiently sensitive to probe to large radii. The limitation lies not only in the achievable signal-to-noise ratio (which is limited because typical narrow-band filters are much wider than the H$\\alpha$ emission line), but also in the stellar continuum subtraction. At large radii, because the spectral flux density at the peak of the H$\\alpha$ line is of the same order of magnitude as the stellar continuum, the wavelength dependence of the stellar continuum spectral energy is sufficient to significantly degrade the subtraction. For example, absorption troughs, which the H$\\alpha$ emission line is often embedded in, can render it undetectable. The way around these problems is to use higher spectral resolution. This can involve very narrow filter bandpasses, which is now possible over a wide range of redshifts due to the increasing availability of tunable filters \\citep{bh98,cepa}, or traditional spectroscopy. In a pioneering effort in the study of gaseous outer disks, \\citet{blandhawthorn} detected H$\\alpha$ emission at $\\sim 1.25$ R$_{25}$, beyond the truncation radius of the neutral hydrogen disk, using the Fabry-Perot staring technique. However, true spectroscopic observations provide improvements in the subtraction of the stellar continuum, of the [NII] emission lines, and of H$\\alpha$ absorption troughs, allowing one to reach sensitivity limits fainter than 10$^{-18}$ erg s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$ (see also \\citet{madsen}). We have observed late-type, edge-on disk galaxies with multi-hour, long-slit spectroscopy. The long-slit technique, although suffering from the much lower throughput in comparison to the Fabry-Perot technique, allows us, for edge-on geometry, to take advantage of fields of view that cover the entire galaxy and greatly facilitates sky and continuum subtraction. We typically detect H$\\alpha$ emission to galactocentric radii of 1.5 R$_{25}$, and in certain cases up to 2 R$_{25}$. In a previous paper \\citep{christleinzaritsky}, we discuss the kinematic properties of these outer H$\\alpha$ disks and find the kinematics to be disk-like, with generally no indication of kinematic anomalies or higher velocity dispersions as one approaches the outer edge of the disk. A subsample of galaxies with known optical warps has been studied separately \\citep{christleinbh} to determine if kinematic anomalies, such as breaks in the rotation curve, were associated with the onset of warps that could indicate ongoing accretion processes of compact HI clouds or satellite galaxies as causes of the warp. We use data from both of these samples in this paper. We now turn our attention to the surface brightness profile of H$\\alpha$ emission in the outer disk. Our aim is to measure the distribution of H$\\alpha$ emission in the outer disk and determine if the H$\\alpha$ emission profiles can be categorized as done for the stellar continua. We will determine whether there is indeed a break in the H$\\alpha$ surface brightness profile, and if so, whether this break can be associated with that in the stellar surface brightness. If there is a break, is it steep enough to constitute a truncation? And how is the stellar continuum, a measure of integrated star formation, distributed in comparison to the current star formation, as indicated by H$\\alpha$? We describe our sample in \\S \\ref{sec_data}. In \\S \\ref{sec_composite}, we discuss and compare the properties of H$\\alpha$ and stellar continuum emission based on a composite of our 15 individual galaxy spectra, while in \\S \\ref{sec_individual}, we discuss what constraints we can place on the profile shapes of individual galaxies. ", "conclusions": "We have presented a study of H$\\alpha$ emission and stellar continuum in 15 low-redshift, late-type, edge-on galaxies. The use of deep long-slit spectroscopy, rather than narrow-band imaging, allows us to probe H$\\alpha$ emission to very faint levels, $\\sim$10$^{-18}$ erg s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$. H$\\alpha$ emission is traced out to 50\\% beyond the R$_{25}$ radius, roughly twice as far as typical H$\\alpha$ rotation curves in the literature. It is thus possible to probe the outer galactic disks, which have hitherto primarily been a domain of radio astronomy, in H$\\alpha$ emission, and thus simultaneously gain information about kinematics, star formation, metallicity, and stellar continuum with arcsecond-scale spatial resolution. In the present paper, we have focused on the properties of the surface brightness profile of H$\\alpha$ emission and compared them to those of the stellar continuum. Our conclusions are: \\begin{itemize} \\item In this sample, which is dominated by Freeman Type II (``sub-exponential\") profiles in the stellar continuum (a broken exponential profile with a steeper slope in the outer than the inner disk), H$\\alpha$ emission in the composite profile also follows a Type II profile. For the H$\\alpha$ composite profile, we can rule out both an unbroken single exponential profile as well as a truncated profile, for which the flux drops to zero at a given break radius. \\item There is a well-defined break radius in the composite H$\\alpha$ surface brightness profile at $r_0\\approx0.7$ R$_{25}$. \\item The H$\\alpha$ and stellar continuum distributions may be consistent with a single break radius $r_0$, but the stellar continuum profile drops more slowly in the outer disk. \\item The presence of a clearly defined break in the H$\\alpha$ surface brightness profile, despite the fact that our observations of the outer disk slope are clearly not limited by Poisson noise, indicates that any apparent deficit of H$\\alpha$ emission in the outer disk is not, in general, a stochastic effect arising from low number statistics. The break in H$\\alpha$ is real, but it is not a truncation, as evidenced by the common presence of low-level outer-disk H$\\alpha$ emission to $\\sim 1.5 r_0$. \\item For most objects in our sample, a truncated exponential as well as an unbroken exponential can be ruled out even on an individual basis, thereby requiring a broken exponential profile. All but one of our galaxies are individually consistent with a broken exponential profile, with a fixed H$\\alpha$ break radius at $r_0=0.7 R_{25}$, showing that the conclusions drawn from the composite profile describe not just a meaningless average, but are representative of at least a significant fraction of individual galaxies. \\end{itemize}" }, "1003/1003.1720_arXiv.txt": { "abstract": "We present photometric and spectroscopic follow-up of a sample of extragalactic novae discovered by the Palomar 60-inch telescope during a search for ``Fast Transients In Nearest Galaxies'' (P60-FasTING). Designed as a fast cadence (1-day) and deep ($g$\\,$<$21\\,mag) survey, P60-FasTING was particularly sensitive to short-lived and faint optical transients. The P60-FasTING nova sample includes 10 novae in M\\,31, 6 in M\\,81, 3 in M\\,82, 1 in NGC\\,2403 and 1 in NGC\\,891. This significantly expands the known sample of extragalactic novae beyond the Local Group, including the first discoveries in a starburst environment. Surprisingly, our photometry shows that this sample is quite inconsistent with the canonical Maximum Magnitude Rate of Decline (MMRD) relation for classical novae. Furthermore, the spectra of the P60-FasTING sample are indistinguishable from classical novae. We suggest that we have uncovered a sub-class of faint and fast classical novae in a new phase space in luminosity-timescale of optical transients. Thus, novae span two orders of magnitude in both luminosity and time. Perhaps, the MMRD, which is characterized only by the white dwarf mass, was an over-simplification. Nova physics appears to be characterized by quite a rich four-dimensional parameter space in white dwarf mass, temperature, composition and accretion rate. ", "introduction": "Since the discovery of classical novae, astronomers have pursued their use as standard candles to determine distances (see \\citealt{h29}). \\citet{z36} first noticed some regularity in nova light curves and termed this the ``life-luminosity'' relation. \\citet{a56} undertook a comprehensive search for novae in M31, discovering thirty novae in 290 nights, and found a clear relation --- luminous novae evolve faster than less luminous novae. The modern name for this observation is the maximum-magnitude rate-of-decline relation (MMRD relation). The MMRD relation has attracted considerable theoretical attention (e.g. \\citealt{l92}). The basic idea is that the relation is entirely due to the mass of the accreting white dwarf. The more massive the white dwarf, the higher the surface gravity, the higher the pressure at the base of envelope and stronger the thermonuclear runaway (and hence, higher the peak luminosity). Also, the more massive the white dwarf, the smaller the envelope mass to attain the critical pressure for thermonuclear runaway (TNR) and hence, faster the decline. In more recent times, \\citet{dl95} used a sample of novae in M31 and LMC to propose an arctangent relation between the peak luminosity and rate of decline. \\citet{dd00} used a sample of Galactic novae to propose a linear relation between the same two parameters. \\citet{dbk+06} used a score of novae in M31 from POINT AGAPE survey and claimed their observations were consistent with the \\citet{dl95} formulation of the MMRD. In comparison to supernovae, classical novae are not very luminous. Hence, searches have traditionally focussed only on the Milky Way and its nearest neighbors (Andromeda and the Large Magellanic Cloud). \\citet{hsg+08} looked into archival data and found 49 nova candidates{\\footnote{We use the term candidate where the light curve is very sparse and/or there is no spectrosopic confirmation}} in M\\,81 in the past 20 years --- unfortunately, these candidates neither have light curves nor spectra. \\citet{fcj03} undertook a search for novae using 24 orbits of the Hubble Space Telescope and found nine nova candidates in M49. Even with their sparsely sampled light curves for nine novae, they concluded that novae are not good standard candles. Another survey, CFHT-COVET{\\footnote{{\\bf C}anada {\\bf F}rance {\\bf H}awaii {\\bf T}elescope {\\bf CO}ma {\\bf V}irgo {\\bf E}xploration for {\\bf T}ransients}} (aimed at finding transients in the gap between novae and supernovae) found a dozen nova candidates in many galaxies in the Virgo supercluster, including some in the far outskirts of galaxies (Kasliwal et al 2010, in prep). Here, we report on novae discovered in high cadence monitoring observations of a representative collection of galaxies with distance less than that of the Virgo cluster. The original motivation of this search -- P60-FasTING{\\footnote{{\\bf P}alomar {\\bf 60}-inch {\\bf Fas}t {\\bf T}ransient {\\bf I}n {\\bf N}earest {\\bf G}alaxies}} was to explore rapid transients (those which last less than a couple of nights) in the nearest galaxies. A strong spectroscopic follow-up effort was a part of P60-FasTING. The survey was capable of finding novae in the major galaxies out to 4\\,Mpc: M\\,31, M\\,81, the star-burst M\\,82 and NGC\\,2403. We present our sample of 21 transients, which although spectroscopically indistinguishable from classical novae, photometrically occupy a new region of phase space. The paper is organized as follows: \\S\\,\\ref{sec:obs} describes the discovery, photometric and spectroscopic follow-up observations of this nova sample, \\S\\,\\ref{sec:analysis} describes the data analysis, \\S\\,\\ref{sec:discussion} discusses the implications and \\S\\,\\ref{sec:conclusion} presents our conclusion. \\begin{table*}[!hbt] \\begin{center} \\caption[]{\\bf Novae Discovered by P60-FasTING} \\begin{tabular}{lllllllll} \\hline \\hline Nova & Host & Discovery Date & RA(J2000) & DEC(J2000) & Offset from Host & Reference\\cr % \\hline P60-NGC2403-090314 & NGC2403 & 2009 Mar 14.160 & 07:36:35.00 & +65:40:20.8 & 101.0\"W, 252.0\"N & \\cite{k+09c} \\cr % P60-M82OT-090314 & NGC3034 & 2009 Mar 14.496 & 09:56:12.60 & +69:41:32.3 & 104.2\"E, 48.2\"N & \\nodata \\cr P60-M81OT-090213 & NGC3031 & 2009 Feb 13.404 & 09:55:35.96 & +69:01:51.0 & 15\"E, 124\"S & \\cite{k+09b} \\cr P60-M31OT-081230 (2008-12b) & NGC224 & 2008 Dec 30.207 & 00:43:05.03 & +41:17:52.3 & 233.4\"E,103.8\"N & \\cite{k+09a} \\cr P60-M81OT-081229 & NGC3031 & 2008 Dec 29.373 & 09:55:38.15 & +69:01:43.6 & 26.7\"E, 131.4\"S & \\cite{rks09} \\cr P60-M81OT-081203 & NGC3031 & 2008 Dec 3.303 & 09:55:16.92 & +69:02:17.7 & 87.2\"W, 97.4\"S & \\cite{k+08f} \\cr P60-M82OT-081119 & NGC3034 & 2008 Nov 19.536 & 09:55:58.39 & +69:40:56.2 & 29.5\"E, 10.4\"N &\\cite{k+08e} \\cr % P60-M81OT-081027 & NGC3031 & 2008 Oct 27.402 & 09:55:36.11 & +69:03:22.0 & 15.8\"E, 33.1\"S & \\cite{k+08d} \\cr P60-M81OT-080925 & NGC3031 & 2008 Sep 25.49 & 09:55:59.35 & +69:05:57.1 & 2.35'E, 2.03'N & \\cite{k+08c} \\cr P60-M31OT-080915 (2008-09c) & NGC224 & 2008 Sep 15.36 & 00:42:51.42 & +41:01:54.0 & 1.34'E, 14.24'S & \\cite{k+08h} \\cr P60-M31OT-080913 (2008-09a) & NGC224 & 2008 Sep 13.18 & 00:41:46.72 & +41:07:52.1 & 10.8'W, 8.3'S & \\cite{k+08g} \\cr P60-NGC891OT-080813 & NGC891 & 2008 Aug 13.45 & 02:22:32.70 & +42:21:56.1 & 8\"W,59\"N & \\cite{k+08b}\\cr P60-M31OT-080723 (2008-07b) & NGC224 & 2008 Jul 23.33 & 00:43:27.28 & +41:10:03.3 & 8.1'E,6.1'S & \\cite{k+08i}\\cr P60-M82OT-080429 & NGC3034 & 2008 Apr 29.24 & 09:55:21.00 & +69:39:42.0 & 165\" W, 64\" S & \\cite{k+08a}\\cr P60-M81OT-071213 & NGC3031 & 2007 Dec 13.40 & 09:55:25.98 & +69:04:34.8 & 40\"W,40\"N & \\cite{k+07} \\cr \\hline \\hline \\label{tab:allnovae} \\end{tabular} \\end{center} \\end{table*} \\begin{table*}[!hbt] \\begin{center} \\caption[]{\\bf Additional M31 Novae} \\begin{tabular}{llll} \\hline \\hline Nova Name & Discovery Date & Classification & Reference\\cr \\hline 2007-10a & 54380.606 & Fe II & \\cite{pbs+07,gr07} \\cr % 2007-11f & 54433.716 & \\nodata & \\cite{ovk+07} \\cr % 2007-12b & 54444.528 & He/N & Nakamo,Hornoch \\cite{lrr+07,bds+09} \\cr % 2008-08c & 54708.127 & \\nodata & \\cite{vol+08};Hornoch \\cr % 2008-10b & 54759.698 & Fe II & \\cite{hpb+08,dco+08,bfh+08} \\cr % 2008-11a & 54774.438 & Hybrid & Nishiyama;Hornoch \\cite{scb+08} \\cr % \\hline \\hline \\label{tab:addm31novae} \\end{tabular} \\end{center} \\end{table*} ", "conclusions": "\\label{sec:conclusion} We conclude that P60-FasTING has uncovered classical novae in a new region in the luminosity-timescale phase space of optical transients. Classical novae span at least two orders of magnitude in time and two orders of magnitude in luminosity. Future surveys would have a large enough sample to meaningfully constrain the relative populations of classical novae in the different areas of phase space. P60-FasTING was designed as a pilot project, to begin to set the stage for future projects such as Palomar Transient Factory (PTF\\footnote{http://www.astro.caltech.edu/ptf}, \\citealt{lkd+09}, \\citealt{rkl+09}, \\citealt{gsv+08}), PanSTARRS (PS1\\footnote{http://pan-starrs.ifa.hawaii.edu}) and Large Synoptic Survey Telescope (LSST\\footnote{http://www.lsst.org}). Both PTF and PS1 are now underway. PTF is looking at several nearby galaxies with a similar depth and cadence as P60-FasTING. Among nearby galaxies, PS1 day-cadence fields only cover M31 but are a couple of magnitudes deeper. LSST will be both deeper and faster cadence and cover the visible sky. P60-FasTING is only the trailblazer for the uncovering of a wealth of information about classical novae by near-future synoptic surveys. \\bigskip" }, "1003/1003.4175.txt": { "abstract": "{Using a semiclassical approach to Gravitoelectromagnetic Inflation (GEMI), we study the origin and evolution of seminal inflaton and electromagnetic fields in the early inflationary universe from a 5D vacuum state. We use simultaneously the Lorentz and Feynman gauges. Our formalism is naturally not conformal invariant on the effective 4D de Sitter metric, which make possible the super adiabatic amplification of electric and magnetic field modes during the early inflationary epoch of the universe on cosmological scales. This is the first time that solutions for the electric field fluctuations are investigated in a systematic way as embeddings for inflationary models in 4D. An important and new result here obtained is that the spectrum of the electric field fluctuations depend with the scale, such that the spectral index increases quadratically as the scale decreases.} ", "introduction": "The origin of cosmological scales magnetic fields is one of the most important, fascinating and challenging problems in modern cosmology. Many scenarios have been proposed to explain them. Magnetic fields are known to be present on various scales of the universe\\cite{3}. Primordial large-scale magnetic fields may be present and serve as seeds for the magnetic fields in galaxies and clusters. Until recently the most accepted idea for the formation of large-scale magnetic fields was the exponentiation of a seed field as suggested by Zeldovich and collaborators long time ago. This seed mechanism is known as galactic dynamo. However, recent observations have cast serious doubts on this possibility. There are many reasons to believe that this mechanism cannot be universal. This is why the mechanism responsible for the origin of large-scale magnetic fields is looked in the early universe, more precisely during inflation\\cite{turner}, which should be amplified through the dynamo mechanism after galaxy formation. In principle, one should be able to follow the evolution of magnetic fields from their creation as seed fields through to dynamo phase characteristic of galaxies. It is believed that magnetic fields can play an important role in the formation and evolution of galaxies and their clusters, but are probably not essential to our understanding of large-scale structure in the universe. However, an understanding of structure formation is paramount to the problem of galactic and extragalactic magnetic fields\\cite{1,2}. It is natural to look for the possibility of generating large-scales magnetic fields during inflation with strength according with observational data on cosmological scales: $< 10^{-9}$ Gauss\\cite{giova}. However, the FRW universe is conformal flat and the Maxwell theory is conformal invariant, so that magnetic fields generated at inflation would come vanishingly small at the end of the inflationary epoch. The possibility to solve this problem relies in produce non-trivial magnetic fields in which conformal invariance to be broken. On the other hand, the five dimensional model is the simplest extension of General Relativity (GR), and is widely regarded as the low-energy limit of models with higher dimensions (such as 10D supersymmetry and 11D supergravity). Modern versions of 5D GR abandon the cylinder and compactification conditions used in original Kaluza-Klein (KK) theories, which caused problems with the cosmological constant and the masses of particles, and consider a large extra dimension. In particular, the Induced Matter Theory (IMT) is based on the assumption that ordinary matter and physical fields that we can observe in our 4D universe can be geometrically induced from a 5D Ricci-flat metric with a space-like noncompact extra dimension on which we define a physical 5D apparent vacuum. The vacuum we shall consider is very restrictive in the sense that we shall not consider any kind of charges, matter or currents on the 5D spacetime. In a relativistic framework, it can be expressed by the 5D null geodesic equations, which are only valid for massless test particles in 5D. However, observers that move with frames $U^4 \\equiv {d x^4 \\over dS}=0$ (described by a constant foliation on the extra dimension), can see the physics described by the effective 4D energy-momentum tensor embedded in the 5D apparent vacuum, which is geometrically described by a 5D Ricci-flat spacetime. From the mathematical point of view, the Campbell-Magaard theorem\\cite{campbell} serves as a ladder to go between manifolds whose dimensionality differs by one. This theorem, which is valid in any number of dimensions, implies that every solution of the 4D Einstein equations with arbitrary energy-momentum tensor can be embedded, at least locally, in a solution of the 5D Einstein field equations in vacuum. Because of this, matter, charge and currents may be 4D manifestations of the topology of space. Gravitoelectromagnetic Inflation (GEMI)\\cite{gemi} was proposed recently with the aim to describe, in an unified manner, electromagnetic, gravitational and the inflaton fields in the early inflationary universe, from a 5D vacuum. It is known that conformal invariance must be broken to generate non-trivial magnetic fields. A very important fact is that in this formalism conformal invariance is naturally broken. Other conformal symmetry breaking mechanisms have been proposed so far\\cite{tur}. However, most of these are developed in the Coulomb gauge. In order to simplify the equations of motion for $A^{\\nu}$, in this paper we use simultaneously the Lorentz and Feynman gauges, to calculate the electric and magnetic spectral indices for the spectrums of these fluctuations taking into account the induced currents. The main contribution of this paper is the study for the spectrum of the electric field fluctuations in a systematic way as embeddings for inflationary models in 4D. This topic has been ignored in the literature. The paper is organized as follows: in Sect. II we introduce the 5D vacuum of the fields on a generic 5D Ricci flat metric, to obtain the equations for the vector fields using simultaneously the generalized Lorentz and Feynman gauges. Also, we impose a semiclassical approach to the vector fields. In Sect. III we study the particular case of a 5D Ricci flat space-time for an extended de Sitter expansion. In Sect. IV describe the dynamics of the vector fields on an effective 4D de Sitter space-time, when we make a static foliation on the noncompact extra dimension, which is considered as space-like: $\\psi=\\psi_0$. We develop the equations of motion for the fields using a particular Lorentz gauge on the effective 4D de Sitter space-time. After it, we describe the dynamics of the classical and quantum fields, to finally calculate the evolution and spectrums of the inflaton, electric and magnetic fields. The conclusions are developed in the Sect. V. Finally, we have included two appendixes where we have developed respectively the details of the calculations for the modes of the electric field fluctuations, and the spectrum for these fluctuations. ", "conclusions": "" }, "1003/1003.6100_arXiv.txt": { "abstract": "{Complex molecules such as ethanol and dimethyl ether have been observed in a number of hot molecular cores and hot corinos. Attempts to model the molecular formation process using gas phase only models have so far been unsuccessful.} {To demonstrate that grain surface processing is a viable mechanism for complex molecule formation in these environments.} {A variable environment parameter computer model has been constructed which includes both gas and surface chemistry. This is used to investigate a variety of cloud collapse scenarios.} {Comparison between model results and observation shows that by combining grain surface processing with gas phase chemistry complex molecules can be produced in observed abundances in a number of core and corino scenarios. Differences in abundances are due to the initial atomic and molecular composition of the core/corino and varying collapse timescales.} {Grain surface processing, combined with variation of physical conditions, can be regarded as a viable method for the formation of complex molecules in the environment found in the vicinity of a hot core/corino and produce abundances comparable to those observed.} ", "introduction": "A number of complex molecules have been discovered in the interstellar medium. First was methanol, CH$_{3}$OH, (Ball \\cite{Ball}) followed by dimethyl ether, CH$_{3}$OCH$_{3}$, (Snyder et al. \\cite{Snyder}) and ethanol, CH$_{3}$CH$_{2}$OH, (herein denoted C$_{2}$H$_{5}$OH) (Zuckerman et al. \\cite{Zuckerman}). A variety of other isomers, isotopic variants and similar compounds have also been seen. Charnley et al. (\\cite{Charnley95}) review the distribution of complex molecules while Ikeda et al. (\\cite{Ikeda01} \\& \\cite{Ikeda02}) present more recent observations. CH$_{3}$OH is seen in cold dark clouds with a relative abundance of 10$^{-9}$ (with respect to H$_{2}$) while in hot cores this can be as large as 10$^{-6}$. The highest abundances are observed in regions where grain mantles are likely to have recently evaporated, strongly suggesting mantle processing either directly or in the formation of precursors. A variety of computer models have been constructed to investigate molecule formation. Predominantly these are based on chemical reaction networks and use a \\lq\\lq flat\\rq\\rq\\ (fixed) parameter space. As with the models presented here most are \\lq\\lq single point\\rq\\rq, where a single representative point in a cloud is modelled. This is satisfactory as chemistry at another point with different physical parameters could be modelled simply by changing the parameters. Over time the size and complexity of the reaction networks has increased as more reaction rate data is published in the literature and greater computer power becomes available to process more complex networks. (See Millar (\\cite{Millar90}) for a review of model development.) Millar et al. (\\cite{Millar91b}) model the production of CH$_{3}$OH by a fixed parameter gas-phase-only model and are unable to produce abundances in excess of 10$^{-7}$ lending further weight to the involvement of grain mantle processing. C$_{2}$H$_{5}$OH is seen in fewer sources than CH$_{3}$OH and all known sources are dense, hot core star-forming regions, exactly the places where grain mantles are likely to evaporate. Abundances are in the range 10$^{-9}$-10$^{-8}$ (with respect to H$_{2}$). CH$_{3}$OCH$_{3}$ is often seen in the same sources as C$_{2}$H$_{5}$OH and has recently been detected in hot corinos (Ceccarelli et al. \\cite{Ceccarelli04}; Bottinelli et al. \\cite{Bottinelli04}). It has an abundance range of 10$^{-8}$-10$^{-6}$ and the CH$_{3}$OCH$_{3}$/C$_{2}$H$_{5}$OH abundance ratio differs greatly between apparently similar sources. (Abundances with respect to H$_{2}$ for selected sources are listed in Table \\ref{Observed fractional abundances of complex molecules}.) \\begin{table}[ht] \\caption{Observed fractional abundances of complex molecules} \\label{Observed fractional abundances of complex molecules} \\centering \\begin{tabular} {lrrr} \\hline Source - Hot Core & CH$_3$OH & C$_2$H$_5$OH & CH$_3$OCH$_3$ \\\\ \\hline \\object{NGC6334F} & $2.0\\times10^{-7}$ & $9.0\\times10^{-9}$ & $4.0\\times10^{-8}$ \\\\ \\object{G327.3-0.6} & $1.0\\times10^{-7}$ & $1.0\\times10^{-8}$ & $3.0\\times10^{-8}$ \\\\ \\object{G31.41+0.31} & $9.0\\times10^{-8}$ & $2.0\\times10^{-8}$ & $2.0\\times10^{-8}$ \\\\ \\object{G34.3+0.2} & $9.0\\times10^{-8}$ & $6.0\\times10^{-9}$ & $1.0\\times10^{-8}$ \\\\ \\object{G10.47+0.03} & $2.0\\times10^{-7}$ & $1.0\\times10^{-8}$ & $3.0\\times10^{-8}$ \\\\ \\object{Sgr B2 (N)} & $2.0\\times10^{-8}$ & $1.0\\times10^{-9}$ & $7.0\\times10^{-10}$ \\\\ \\object{DR21(OH)} & $1.0\\times10^{-8}$ & & \\\\ \\object{W3(H$_2$O)} & $4.0\\times10^{-8}$ & & \\\\ \\object{W51 e1/e2} & $3.0\\times10^{-7}$ & $9.0\\times10^{-9}$ & \\\\ \\object{Orion Compact Ridge} & $2.0\\times10^{-7}$ & & \\\\ \\object{Orion Hot Core} & $1.0\\times10^{-6}$ & $2.0\\times10^{-8}$ & \\\\ \\hline Source - Hot Corino \\\\ \\hline \\object{IRAS16293-2422} & $1.0\\times10^{-7}$ & & $2.4\\times10^{-7}$ \\\\ \\object{NGC1333-IRAS4A} & $<1.0\\times10^{-8}$& & $<2.8\\times10^{-8}$ \\\\ \\object{NGC1333-IRAS4B} & $7.0\\times10^{-7}$ & & $<1.2\\times10^{-6}$ \\\\ \\object{NGC1333-IRAS2A} & $3.0\\times10^{-7}$ & & $3.0\\times10^{-8}$ \\\\ \\hline \\end{tabular} Data from Ikeda et al. \\cite{Ikeda01}; Ikeda et al. \\cite{Ikeda02}; Bottinelli et al. \\cite{Bottinelli07} \\end{table} \\pagebreak Even under very favorable conditions fixed parameter gas-phase-only chemical models produce maximum abundances for C$_2$H$_5$OH and CH$_3$OCH$_3$ of the order of 10$^{-11}$ (Herbst \\& Leung \\cite{Herbst}; Millar et al. \\cite{Millar91a}; Charnley et al. \\cite{Charnley92}), between two and five orders of magnitude below that observed. Again this suggests a scenario in which mantles form on dust grains at some point in a cloud's lifetime when conditions are favorable. The mantles are active and complex molecules (or their precursors) form on them more efficiently than in the gas phase. Later the cloud evolves becoming warmer and denser and the contents of the mantles pass back into the gas phase. Several other gas-phase-only models (for example Charnley et al. \\cite{Charnley92}; Caselli et al. \\cite{Caselli}) have demonstrated that the chemistry in hot cores is active, with particular gas phase reaction channels initiated with simple species thought to originate from mantles, a possible explanation for observed complex molecule abundances. A number of more in-depth models have included surface reaction chemistry. One of the first has been developed by Hasegawa et al. (\\cite{Hasegawa92}) and Hasegawa \\& Herbst (\\cite{Hasegawa93}). This model has 274 chemical species with 2928 reactions and includes 1\\% (by mass) dust grains in the gas cloud. Molecules can freeze out onto grains and later pass back into the gas phase. Once on a grain heavier species molecules are held stationary at lattice binding sites while lighter species, most commonly hydrogen atoms, can migrate around the grain. When mobile light species encounter fixed heavy species the two can combine to produce progressively larger molecules. While large molecules can thaw off the grain surface there are a number of other possible desorbtion mechanisms that can act as well. (See Section \\ref{Desorption mechanisms}.) One additional benefit of this approach is that grain surface catalysis of hydrogen molecules can be modelled directly. Shalabiea \\& Greenberg (\\cite{Shalabiea95b}) then take the next logical step in interstellar cloud modelling by including variations in physical parameters. They use a gas phase reaction rate network and interchange between gas and surface similar to Hasegawa et al. (\\cite{Hasegawa92}) and Hasegawa \\& Herbst (\\cite{Hasegawa93}). However certain physical parameters (eg. density) can also change. Shalabiea \\& Greenberg (\\cite{Shalabiea95b}) model these by including differential equations for them and solving for the relevant parameter(s) at each point in time in much the same way as species abundances are solved for. They also coin the terms \\lq\\lq pseudo-time-dependent\\rq\\rq, where the model chemistry evolves over time while the parameters are static and \\lq\\lq real-time-dependent\\rq\\rq\\ where both chemistry and physical parameters evolve. Their paper provides comparison between the two types, with the gas and grain model including 218 chemical species and 2075 reactions. With the more recent identification of \\lq\\lq hot corinos\\rq\\rq\\ (Table \\ref{Observed fractional abundances of complex molecules} \\& Section \\ref{Calculating parameters}) where the temperature range is optimal for the grain surface production of larger molecules and their subsequent return to the gas phase some specific modelling of these type of objects has been done. Garrod \\& Herbst (\\cite{Garrod06}) use a gas and grain reaction network with 655 species and 6509 reactions. The same model was used by Aikawa et al. (\\cite{Aikawa08}). Subsequently Garrod et al. (\\cite{Garrod08}) produced an extended and more generalized version. Comparison of the output of these models with the results presented here is made in Section \\ref{Conclusion}. As is discussed further in Section \\ref{Desorption mechanisms} some mechanism(s) must be returning material from grain mantles back to the gas phase since the alternative is depletion of molecular species heavier than H$_2$ in a timescale shorter than the lifetime of known clouds, which is not seen. More recent modelling work has focused on the investigating the effectiveness of several proposed desorption mechanisms. Willacy \\& Millar (\\cite{Willacy98}) provide a set of gas and surface models which include different desorption mechanisms and compare their effectiveness. These models use 282 species and 4864 reactions. Garrod et al. (\\cite{Garrod07}) specifically model formation of CH$_{3}$OH in a quiescent cloud and point out that if grain surface processes are invoked to explain the abundance of CH$_{3}$OH then at least one non-thermal desorbtion mechanism must be active. Tielens \\& Charnley (\\cite{Tielens97}) have argued that, at least in some cases, the reaction network model cannot be justified and that a Monte Carlo simulation approach is more appropriate. This has led to some debate in the current literature about the relative merits of the two methods. Willacy \\& Millar (\\cite{Willacy98}) discuss this and conclude that currently there are computational impracticalities in using the Monte Carlo approach in real-time-dependent models, and further that the two different approaches may give fairly similar results for systems of increasing complexity. A direct comparison of the two approaches by Garrod et al. (\\cite{Garrod09}) demonstrates considerable similarity between them. Certainly previous gas phase only reaction rate models have yielded results comparable with observations for many species. The model used here builds on its predecessors. Termed the \\lq\\lq Gas/Surface\\rq\\rq\\ model it is a single point, gas and surface phase, chemical reaction network model with 279 species and 2968 reactions. It includes the grain surface mechainsm used by Hasegawa et al. (\\cite{Hasegawa92}) and Hasegawa \\& Herbst (\\cite{Hasegawa93}) as well as the desorption mechanisms discussed by Willacy \\& Millar (\\cite{Willacy98}). It is also fully \\lq\\lq real-time-dependent\\rq\\rq. Instead of the differential equation approach taken by Shalabiea \\& Greenberg (\\cite{Shalabiea95b}) each separate time step has its own set of physical parameters which are fed into the model from a storage file at the beginning of each time step. This approach has several advantages. Firstly it reduces computational strain on the model since there is no additional equation solving necessary. This is particularly important in situations where parameters are changing very quickly and there are major differences between adjacent time steps. (For example in the final stages of a cloud collapse where density increases drastically in a short time.) Secondly, this approach allows the same set of software to be used to model multiple different situations with no reprogramming at all. Only the input parameter file has to be changed. This allows great flexibility in the scope of scenarios that could be investigated. Further, in certain situations where there are major, abrupt changes in physical parameters (eg. shocks), the differential equation approach is highly prone to breakdown as generating a numerical solution for a given parameter requires some level of continuity between one time step and the next, wheras the parameter file approach avoids this problem completely. The Gas/Surface model is used here to investigate the scenario of complex molecule formation in grain mantles and their subsequent release back into the gas phase. We consider clouds at a variety of initial densities with different exposure to photoionization. As these clouds collapse to become denser and darker, their chemistry becomes more complex and their temperature drops allowing grain mantles to form. Simple gas-phase dominated chemistry at the beginning leads to the deposition of a variety of basic molecules onto an initially bare grain surface. Once on the grain, surface processes allow the build up of complex molecules in a frozen mantle. Later, as the cloud heats up, the mantle evaporates introducing the complex molecules into the gas phase with a variety of chemical consequences significantly different from gas phase only processing. ", "conclusions": "\\label{Conclusion} Overall the Gas/Surface model as used here reproduces to a reasonable degree the observed complex molecule abundances seen in hot corinos. In particular abundance values produced for methanol and dimethyl ether are comparable to those observed in both gas and surface phases. The surface phase molecular abundances produced by the model reflect the starting conditions, particularly the nature of the initial hydrogen - atomic or molecular. This determines the degree of processing in the initial gas phase and hence the abundance of species frozen on to the dust and available for surface chemistry. This in turn allows an understanding of the differences in grain mantle composition observed in different regions. {The model results clearly indicate that both the starting conditions and particularly the timescale of collapse are major contributors to the overall chemistry. A number of other workers in the area have reached similar conclusions. Garrod \\& Herbst (\\cite{Garrod06}) specifically investigated methyl formate (HCOOCH$_3$). They concluded that its formation required the grain surface production of precursor molecules which were then released into the gas phase to produce methyl formate itself. This tallies with the similar behaviour seen here although Garrod \\& Herbst were able to produce abundances closer to observed values. They suggest a similar pattern for dimethyl ether, exactly as the Gas/Surface model shows. Later applications of the same model produced comparable results. Aikawa et al. (\\cite{Aikawa08}) consider a collapse case where again large organic species (or their immediate precursors) are formed on grain mantles before evaporation into the gas phase. Observed abundances can be both well modelled or over and under produced depending on starting conditions and timescale. Garrod et al. (\\cite{Garrod08}) apply their model to more general star formation cases and note specifically that the longer the time available for grain surface chemistry to progress the more complex the eventual gas phase chemistry becomes. In this case agreement with observations is seen for abundances of ethanol and dimethyl ether though not for some other complex molecules. Hassel et al. (\\cite{Hassel08}) specifically model L1529, a \\lq\\lq lukewarm\\rq\\rq\\ corino, which may be the prototype of a new class of such objects. Their analysis demonstrates that, at least in the case of L1529, there is some difficulty in establishing definitvely how a particular chemical state came to evolve and a number of scenarios are possible depending on the length of time the corino spent in each particular phase. In conclusion then grain surface processing, combined with the variation of physical conditions modelled here, can be regarded as a viable method for the formation of complex molecules, particularly ethanol and dimethyl ether, in the environment found in the vicinity of a hot corino and produce abundances comparable to those actually observed." }, "1003/1003.0456_arXiv.txt": { "abstract": "{ {\\small We study the impact of the cosmological parameters uncertainties on the measurements of primordial non-Gaussianity through the large-scale non-Gaussian halo bias effect. While this is not expected to be an issue for the standard $\\Lambda$CDM model, it may not be the case for more general models that modify the large-scale shape of the power spectrum. We consider the so-called local non-Gaussianity model, parametrized by the $f_{\\rm NL}$ non-Gaussianity parameter which is zero for a Gaussian case, and make forecasts on $f_{\\rm NL}$ from planned surveys, alone and combined with a Planck CMB prior. In particular, we consider EUCLID- and LSST-like surveys and forecast the correlations among $f_{\\rm NL}$ and the running of the spectral index $\\alpha_s$, the dark energy equation of state $w$, the effective sound speed of dark energy perturbations $c^2_s$, the total mass of massive neutrinos $M_\\nu=\\sum m_\\nu$, and the number of extra relativistic degrees of freedom $N_\\nu^{rel}$. Neglecting CMB information on $f_{\\rm NL}$ and scales $k > 0.03 h$/Mpc, we find that, if $N_\\nu^{\\rm rel}$ is assumed to be known, the uncertainty on cosmological parameters increases the error on $f_{\\rm NL}$ by 10 to 30\\% depending on the survey. Thus the $f_{\\rm NL}$ constraint is remarkable robust to cosmological model uncertainties. On the other hand, if $N_\\nu^{\\rm rel}$ is simultaneously constrained from the data, the $f_{\\rm NL}$ error increases by $\\sim 80\\%$. Finally, future surveys which provide a large sample of galaxies or galaxy clusters over a volume comparable to the Hubble volume can measure primordial non-Gaussianity of the local form with a marginalized 1--$\\sigma$ error of the order $\\Delta f_{\\rm NL} \\sim 2-5$, after combination with CMB priors for the remaining cosmological parameters. These results are competitive with CMB bispectrum constraints achievable with an ideal CMB experiment.}} ", "introduction": "\\label{Introduction} Tests of deviations from Gaussian initial conditions offer an important window into the very early Universe and a powerful test for the mechanism which generated primordial perturbations. While standard single-field slow-roll models of inflation lead to small departures from Gaussianity, non-standard scenarios allow for a larger non-Gaussianity (NG) level (e.g. \\cite{BKMR04, bartolofnl05, Chenreview}, and refs. therein). In particular, large NG can be produced if any of the conditions below is violated: {\\it a)} single field, {\\it b)} canonical kinetic energy {\\it c)} slow roll and {\\it d)} adiabatic (Bunch-Davies) initial vacuum state. The type of NG arising in standard inflation reads \\cite{SalopekBond90, Ganguietal94,VWHK00, KS01} \\begin{equation} \\Phi=\\phi+f_{\\rm NL}\\left(\\phi^2-\\langle \\phi^2 \\rangle\\right)\\,, \\label{eq:fnl} \\end{equation} where $\\Phi$ denotes Bardeen's gauge-invariant potential, which, on sub-Hubble scales reduces to the usual Newtonian peculiar gravitational potential up to a minus sign, and $\\phi$ denotes a Gaussian random field. The NG parameter $f_{\\rm NL}$ is often considered to be a constant, yielding NG of the {\\it local} type with a bispectrum which is maximized for squeezed configurations \\cite{Creminellishapes}. NG of the local type is generated in standard inflationary scenarios (where $f_{\\rm NL}$ is expected to be of the same order of the slow-roll parameters) as well as in multi-field inflationary scenarios\\footnote{Note that Eq.~(\\ref{eq:fnl}) is not general, i.e. there is a plethora of possible deviations from Gaussianity arising in the different inflationary scenarios proposed in the literature.}. The standard observables to constrain NG are the Cosmic Microwave Background (CMB) and the Large-Scale Structure (LSS) of the Universe. Traditionally, the most popular method to detect primordial NG has been to measure the bispectrum or the three-point correlation function of the CMB \\cite{VWHK00, komatsuetal05, yadav}, while the LSS bispectrum has been shown to be sensitive to primordial NG only at high redshift \\cite{VWHK00, Scocc, sefusattikomatsu, Coor06, PPM07}. Other powerful techniques to measure NG are based on weak lensing tomography \\cite{fedeli09}, Integrated Sachs-Wolfe effect (ISW) \\cite{CVM08,AfshordiTolley08}, abundance \\cite{MVJ00, VJKM01, Loverdeetal07, RB00, RGS00} and clustering \\cite{GW86, MLB86} of rare events such as density peaks, since they trace the tail of the underlying distribution. Refs. \\cite{DDHS07} and \\cite{MV08} (hereafter MV08) showed that primordial NG affects the clustering of dark matter halos inducing a scale-dependent large-scale bias. This effect, which goes under the name of non-Gaussian halo bias, is particularly promising, yielding already stringent constraints from existing data \\cite{Slosaretal08, AfshordiTolley08}. Forthcoming constraints on NG exploiting the non-Gaussian halo bias are expected to be similar to those achievable from an ideal CMB survey \\cite{CVM08}. These predictions have been confirmed by N-body simulations \\cite{DDHS07,Grossietal09,Desjacques,Pillepich2010}. Forecasts for $f_{\\rm NL}$ constraints from the halo-bias have been carried out so far assuming perfect knowledge of all other cosmological parameters. While for a $\\Lambda$CDM model this is expected to be a reasonable assumption, for more general models one may expect $f_{\\rm NL} $ to be degenerate with other parameters and thus to have a larger marginal error. Here we study the degeneracies among the large-scale non-Gaussian halo bias (for NG of the local type) and the cosmological parameters which affect the large-scale halo power spectrum, focusing on dark energy perturbations, massive neutrinos, number of relativistic species, and running spectral index, which can produce large deviations of the underlying cosmology from the minimal $\\Lambda$CDM scenario. The paper is organized as follows. In \\S \\ref{Non-Gaussian halo bias} we briefly review the analytic expressions of the non-Gaussian halo power spectrum generalized to redshift dependent transfer functions. The redshift dependence is due to the presence of both dark energy perturbations and massive neutrinos. In \\S \\ref{Methodology} we summarize the Fisher matrix formalism applied to the observed halo power spectrum. In \\S \\ref{Model parameters} we describe the assumed fiducial cosmology. Finally, in \\S \\ref{Results} and \\S \\ref{Conclusions} we discuss the results and draw our conclusions. ", "conclusions": "\\label{Conclusions} Deviations from non-Gaussianity, usually parameterized by the parameter $f_{\\rm NL}$, offer a powerful tool to identify the mechanism which generates the seeds for the structures we observe currently in our Universe. Here we study the impact of the uncertainties of the cosmological parameters on the $f_{\\rm NL}$ errors expected for the case of local non-Gaussianity for the large-scale non-Gaussian halo bias effect. We forecast the correlations among $f_{\\rm NL}$ and the remaining cosmological parameters (including the running of the spectral index $\\alpha_s$, and the dark energy parameters $w$ and $c^2_s$) within two possible cosmological models. The first model contains massive neutrinos (hypothesis robustly confirmed by neutrino oscillation data), where the total neutrino mass is a parameter to be constrained by the cosmological data. The second model assumes massless neutrinos (or neutrinos with a mass too small to be relevant for the cosmological observations considered here) and allows for extra relativistic degrees of freedom $N_\\nu^{rel}$, which could be induced by the presence of sterile neutrinos, non minimally coupled quintessence fields, or even by the violation of the spin statistics theorem in the neutrino sector. We follow here a conservative approach, assuming that $f_{\\rm NL}$ is constrained exclusively from the very large scale halo power spectrum (i.e. we neglect CMB information on $f_{\\rm NL}$), and restrict ourselves to scales $k \\leqslant 0.03 h$/Mpc, without exploiting information e.g., from BAOs, which will further reduce degeneracies and forecasted errors. We present first the Fisher matrix forecasts for $f_{\\rm NL}$ assuming EUCLID- and LSST-like surveys for the two model cosmologies considered here. Then, we add the Planck Fisher forecasts for the remaining cosmological parameters to study the impact on the $f_{\\rm NL}$ correlations. The combined errors on $f_{\\rm NL}$ do not change significantly in the presence of a dark energy equation of state, massive neutrinos, running of the spectral index, or clustering of dark energy perturbations, which are the parameters we have particularly focused on, since they are expected to affect the matter power spectrum on large scales, and represent the main deviations from a minimal $\\Lambda$CDM model. However, the errors on $f_{\\rm NL}$ are highly affected in the presence of extra relativistic degrees of freedom $N_\\nu^{rel}$. We find that if $N_\\nu^{\\rm rel}$ is assumed to be fixed, the effect of the uncertainties on the other cosmological parameters increases the error on $f_{\\rm NL}$ only by 10 to 30\\% depending on the survey. If $N_\\nu^{\\rm rel}$ is considered as a parameter to be simultaneously constrained from the data, then the uncertainty in the underlying cosmology increases the $f_{\\rm NL}$ error by $\\sim 80\\%$. We thus conclude that, except for the effect of $N_\\nu^{\\rm rel}$, the halo-bias $f_{\\rm NL}$ constraints are remarkably robust to uncertainties in the underlying cosmology. One important point to discuss is the effect of the (Gaussian) halo bias, as its value boosts the effect of $f_{\\rm NL}$ on the halo power spectrum shape, and in our analysis it has been assumed to be known. The (Gaussian) halo bias depends strongly on the type of halos selected by the survey --whether they correspond to extremely high and rare peaks in the initial fluctuation field--, and on their accretion history. Errors on $f_{\\rm NL}$ may be improved -at least in principle- by up to a factor of two by optimizing the choice of tracers. The bias factor itself will need to be estimated from the survey, at the same time as the other cosmological parameters; the signal comes from scales much smaller than those used here, where the NG effect on halo bias is completely negligible. We estimate that the error on the (Gaussian) halo bias will be of the same order (in \\%) as the error on the linear growth factor $f$ as a function of redshift, which is forecasted to be $< \\sim 10\\%$ \\cite{Licia2df,White:2008jy}. Such a residual uncertainly will therefore increase the $f_{\\rm NL}$ errors reported here by at most 10\\%. Let us recall that the purpose of this work is to show the main correlations between $f_{\\rm NL}$ and the other cosmological parameters, and to understand if these degeneracies can degrade dramatically the $f_{\\rm NL}$ errors. We have shown that, after the combination with Planck constraints on parameters different from $f_{\\rm NL}$, the degeneracies get mostly broken, independently on the particular cosmological parameter, even without adding information from smaller scales corresponding to $k>0.03h$/Mpc. Therefore we conclude that the $f_{\\rm NL}$ constraints are very robust against underlying cosmology assumptions. Finally, future surveys which provide a large sample of galaxies or galaxy clusters over a volume comparable to the Hubble volume (LSST, EUCLID) will measure primordial non-Gaussianity of the local form with a marginalized 1--$\\sigma$ error of the order $\\Delta f_{\\rm NL} \\sim 2-5$, after combination with CMB priors for the remaining cosmological parameters. These results are competitive with CMB bispectrum constraints achievable with an ideal CMB experiment $\\Delta f_{\\rm NL} \\sim$few \\cite{YKW07, LiguoriRiotto08}." }, "1003/1003.5790_arXiv.txt": { "abstract": "We present the first successful application of the method of Matched Expansions for the calculation of the self-force on a point particle in a curved spacetime. We investigate the case of a scalar charge in the Nariai spacetime, which serves as a toy model for a point mass moving in the Schwarzschild black hole background. We discuss the singularity structure of the Green function beyond the normal neighbourhood and the interesting effect of caustics on null wave propagation. ", "introduction": "Extreme Mass Ratio Inspirals, such as a stellar-mass compact object orbiting around a supermassive black hole, can be understood within perturbation theory of General Relativity as the smaller body's motion deviating from the background geodesic due to a {\\it self-force}~\\cite{Poisson:2003}. The calculation of the self-force is necessary in order to model the orbital motion and obtain accurate gravitational-wave templates. The expression for the self-force in a general spacetime is formally similar for the cases of an orbiting point scalar charge, electrical charge and a point mass. In all these cases, the self-force is given by some local terms which are relatively easy to calculate, plus a {\\it `tail'} term which, in the case of a scalar charge $q$ following a worldline $z(\\tau)$, is given by \\begin{equation} \\label{eq:S-F} q^2\\int_{-\\infty}^{\\tau^-} \\nabla_\\mu G_{ret} (z(\\tau), z(\\tau')) d \\tau' \\end{equation} where $G_{ret}(x, x')$ is the {\\it `retarded' Green function} of the scalar field wave equation. As the `tail' term depends on the {\\em entire} past history of the particle's worldline, it is usually quite difficult to calculate and one has to do so numerically. ", "conclusions": "" }, "1003/1003.5273_arXiv.txt": { "abstract": "Solar sub-surface fluid topology provides an indirect approach to examine the internal characteristics of active regions (ARs). Earlier studies have revealed the prevalence of strong flows in the interior of ARs having complex magnetic fields. Using the Doppler data obtained by the Global Oscillation Network Group (GONG) project for a sample of 74 ARs, we have discovered the presence of steep gradients in meridional velocity at depths ranging from 1.5 to 5 Mm in flare productive ARs. The sample of these ARs is taken from the Carrington rotations 1980--2052 covering the period August 2001-January 2007. The gradients showed an interesting hemispheric trend of negative (positive) signs in the northern (southern) hemisphere, i.e., directed toward the equator. We have discovered three sheared layers in the depth range of \\mbox{0--10 Mm}, providing an evidence of complex flow structures in several ARs. An important inference derived from our analysis is that the location of the deepest zero vertical vorticity is correlated with the remaining life time of ARs. This new finding may be employed as a tool for predicting the life expectancy of an AR. ", "introduction": "\\label{sec:introduction} Solar surface flows have been studied over the last three decades using mostly the surface observations such as the photospheric magnetic tracers and Doppler measurements. These tracers revealed that the magnetized regions rotate faster than the surrounding medium of field-free plasma \\citep{1978ApJ...219L..55G,1993SoPh..147..207K,1996ARA&A..34...75H}. Doppler measurements have shown poleward meridional flows at the solar surface \\citep{1979SoPh...63....3D,1982SoPh...80..361L,1988SoPh..117..291U}. More recently, it has become possible to study sub-surface structures and flows subsequent to the advent of helioseismology \\citep{1984ARA&A..22..593D,1991ARA&A..29..627G}, which probes the solar interior using acoustic modes of oscillations \\citep{1970ApJ...162..993U,1971ApL.....7..191L}. Helioseismic studies have revealed that sunspots are rather shallow, near-surface phenomena \\citep{2000SoPh..192..159K,2004ApJ...610.1157B,2006ApJ...640..516C}, and their rotation rate with depth depends on the stage of evolution and age. \\citet{2009A&A...506..875S} have reported that a majority of ARs displays sudden decrease in the rotation speed, compatible with dynamic disconnection of sunspots from their parental magnetic roots. Local helioseismology has further revealed the sunspots to be the locations of large flows near the surface \\citep{2002ApJ...570..855H, 2004ESASP.559..337B, 2004ApJ...603..776Z}. Earlier studies have found that magnetic helicity of ARs follows a hemispheric rule, i.e., positive (negative) in southern (northern) hemisphere \\citep{2001ApJ...549L.261P}. Such a hemispheric trend is also reported for flows in ARs \\citep{2005ApJ...631..636K, 2006SoPh..236..227Z, 2007ApJ...667..571K}. Zonal flows exhibit larger amplitudes in southern hemisphere especially at higher latitudes, coinciding with larger magnetic activity in that hemisphere. This may be attributed to the inclination of the solar rotation axis \\citep{2005ApJ...621L.153B}. Also, ARs advect poleward at nearly the same rate as the quiet regions (QRs) \\citep{2009ApJ...698.1749H} and show convergent horizontal flows combined with cyclonic vorticity, counter clockwise in the northern hemisphere \\citep{2004ApJ...605..554K, 2004ApJ...603..776Z}. Activity related variations are also reported in solar surface flows \\citep{ 2002Sci...296..101V, 2003ApJ...585..553B, 2004ESASP.559..293A, 2005ApJ...634.1405H, 2009ApJ...706L.235M}. It is believed that sheared flows in sub-surface layers cause sunspot motions that may lead to unstable magnetic topologies, causing reconnection of magnetic field lines required for flares. \\citet{2003ApJ...585..553B} have found that the maximum meridional velocity of surface flows is smaller when the Sun is more active. The maximum unsigned zonal and meridional vorticities of ARs are correlated with the total X-ray flare intensity \\citep{2006ApJ...645.1543M}. Furthermore, steep meridional velocity gradients are found in flaring ARs at the depth range of \\mbox{4--5 Mm} \\citep{2004ESASP.559..293A, 2009ApJ...706L.235M}, which decreased after flares. Although internal flows in ARs and QRs have been studied \\citep{2006ApJ...645.1543M, 2009SoPh..258...13K, 2010mcia.conf..516M}, an understanding of their distinctive characteristics in flare productive ARs as compared to that in dormant ARs requires further investigation. In this letter, we report on the properties of internal flows in ARs of varying levels of flare productivity and magnetic complexity observed during the Solar Cycle 23, including large ARs such as NOAA 10030, 10484, 10486, 10070, etc., termed as super-active regions (SARs). These SARs are found to possess distinctly different characteristics of flows in their interiors as compared to the less productive or dormant ARs. We have addressed important issues on the relationship of flow characteristics with the life time of ARs, magnetic and flaring activities and hemispheric trends. ", "conclusions": "\\label{S-conc} From our study of the 74 ARs observed during August 2001-January 2007 of Cycle 23, we have discovered statistically significant relationship among the sub-surface flow topology, energetics and the remaining life time of ARs. We have found the following important results: \\begin{enumerate} \\item Three sheared layers in the depth range 0--10 Mm were found to exist in 44 ARs which revealed the complexity of flow structures beneath the surface of more complex ARs. The two extrema in $u_y$ and $d'u_y$ profiles of these ARs were found to be located at the depths of 1.92$\\pm$0.15 and 4.69$\\pm$0.30 Mm. \\item The ARs having two extrema as compared to the ARs having only a single extremum in $u_y$ were found to be more active as they possessed as large as twice the mean magnetic field (MI), mean GOES X-ray flux and mean life time. \\item The extrema of meridional velocity gradients were found to follow a hemispheric trend, viz., in the northern hemisphere, 24 (70\\%) out of 34 ARs had negative gradients while in the southern hemisphere, 29 (74\\%) out of 39 ARs had positive gradients. \\item ARs of larger MI possessed steeper gradient in meridional velocity profiles. \\item Flaring activity of an AR is found to be associated with depth of the first extremum of vertical vorticity. \\item ARs having zero vertical vorticity at deeper layers are expected to last longer. \\end{enumerate} In summary, we have discovered a new hemispheric trend involving the meridional velocity gradient of sub-surface flows. More importantly, life time of ARs appears to be correlated with the depth of the deepest zero vertical vorticity. This inference may be useful in predicting the expected life time of ARs." }, "1003/1003.5759_arXiv.txt": { "abstract": "We present new measurements of the energy spectra of cosmic-ray (CR) nuclei from the second flight of the balloon-borne experiment CREAM (Cosmic Ray Energetics And Mass). The instrument (CREAM-II) was comprised of detectors based on different techniques (Cherenkov light, specific ionization in scintillators and silicon sensors) to provide a redundant charge identification and a thin ionization calorimeter capable of measuring the energy of cosmic rays up to several hundreds of TeV. The data analysis is described and the individual energy spectra of C, O, Ne, Mg, Si and Fe are reported up to $\\sim 10^{14}$ eV. The spectral shape looks nearly the same for all the primary elements and can be expressed as a power law in energy $E^{-2.66 \\pm 0.04}$. The nitrogen absolute intensity in the energy range 100-800 GeV/n is also measured. ", "introduction": "CREAM is a balloon-borne experiment designed to directly measure the elemental composition and the energy spectra of cosmic rays from H to Fe % in the energy range 10$^{11}$-10$^{15}$ eV. CREAM aims to experimentally test % astrophysical models proposed to explain the acceleration mechanism of cosmic rays and their propagation in the Galaxy \\cite{1}.\\\\ Since 2004, four instruments were successfully flown on long-duration balloons in Antarctica. The instrument configurations varied slightly in each mission, due to various detector upgrades. In this paper, we describe the procedure used to analyze the data collected during the second flight and reconstruct the energy spectra of the major primary CR nuclei and nitrogen. Results are compared with other direct observations and discussed.\\\\ ", "conclusions": "The CREAM-II instrument carried out measurements of high Z cosmic ray nuclei with an excellent charge resolution and a reliable energy determination.\\\\ The energy spectra of the major primary heavy nuclei from C to Fe were measured up to 10$^{14}$ eV and found to agree well with earlier direct measurements. A new measurement of the nitrogen intensity in an energy region thus far experimentally unexplored indicates a less steep power-law trend in the spectrum with respect to lower energies. \\footnotesize" }, "1003/1003.4455_arXiv.txt": { "abstract": "Aiming to perform a study of the warm dust and gas in the luminous blue variable star \\G\\ and its associated nebula, we present infrared Spitzer imaging and spectroscopy, and new CO $J$\\,=\\,2\\,$\\rightarrow$\\,1 and 4\\,$\\rightarrow$\\,3 maps obtained with the IRAM 30m radio telescope and with the Submillimeter Telescope, respectively. We have analyzed the nebula detecting multiple shells of dust and gas connected to the star. Using Infrared Spectrograph--Spitzer spectra, we have compared the properties of the central object, the nebula, and their surroundings. These spectra show a rich variety of solid-state features (amorphous silicates, polycyclic aromatic hydrocarbons, and CO$_2$ ices) and narrow emission lines, superimposed on a thermal continuum. We have also analyzed the physical conditions of the nebula, which point to the existence of a photo-dissociation region. ", "introduction": "Luminous blue variable (LBV) stars are massive objects which, as they evolve from the main sequence, undergo a short period of extremely high mass loss (up to 10$^{-3}$\\,\\Myr), sometimes accompanied by so-called giant eruptions, such as the well-known nineteenth-century outburst of \\objectname{$\\eta$ Car} \\citep{Smith06b}. This mass-loss strongly influences the further stellar evolution and leads to the formation of extended circumstellar nebula \\cite[see][for a review]{Humphreys94}. \\cite{Clark05} reported a total of 12 galactic LBV stars and 23 LBV candidates. From the whole sample, however, only 20 have nebular circumstellar emission. Such nebulae have both gaseous and dusty components, and two kinds of geometry, spherical and bipolar, have been recognized. Since LBV nebulae are strong emitters in the mid- and far-infrared, the physical and chemical conditions of the material surrounding LBV stars can be traced studying their spectral characteristics in this wavelength range. Thus, a few of these nebulae were spectroscopically studied with the Infrared Space Observatory ({\\it ISO}) satellite, revealing the presence of solid-state features characteristic of polycyclic aromatic hydrocarbons (PAHs) and silicate dust in both forms, amorphous and crystalline \\citep[e.g.][]{Lamers96a,Voors99,Voors00b}. Together with these, numerous forbidden lines on top of a smooth continuum were detected \\citep[e.g.][]{Lamers96b}. With the advent of {\\it Spitzer Space Telescope} \\citep{Werner04}, unprecedented high sensitivity and resolution imaging and spectroscopy have been possible in the infrared, making feasible the study of LBV nebulae with much great detail. Thus, recently new works have been published \\citep{Morris08, Smith07, Umana09, Gvaramadze09} based on these new data, which have represented an important step forward in the knowledge of this brief evolutionary phase of high-mass stars. The molecular gas in LBV nebulae is also a subject of interest. It may trace radiatively affected and/or shocked regions, which may tell us about the evolution of the progenitor star. In addition, the chemical evolution and the formation of circumstellar molecular gas may also be determined by molecular studies. Gaseous CO was firstly detected around the \\objectname{AG Car} nebula \\citep{Nota02}, and after then in some other objects \\citep{Rizzo08a}. Highly excited ammonia was also detected in the \\objectname{Homunculus Nebula} \\citep{Smith06a}. Concerning the molecular counterpart, the most studied source is \\objectname{G79.29+0.46} \\citep[][hereafter Paper\\,I]{Rizzo08b}. In this object, the detection of the CO $J$\\,=\\,3\\,$\\rightarrow$\\,2 line surrounding the nebula unveiled the presence of moderately dense (10$^4$\\,--\\,10$^5$\\,cm$^{-3}$) CO gas, probably affected by a C-shock at a velocity of $\\sim$\\,15\\,\\kms. This CO-ring-like structure has a diameter of $\\sim$\\,3\\farcm2, and it is strikingly homogeneous in shape. The origin of the nebula of \\G\\ is not clear. From Very Large Array ({\\it VLA}) observations of radio continuum and \\ion{H}{1}, \\cite{Higgs94} showed the thermal nature of the continuum emission and suggested that the nebula is an ionized shell of $\\sim$\\,15\\,\\Msun\\ of swept-up interstellar material. However, after examining the {\\it IRAS} high-resolution images of \\G, \\cite{Waters96} concluded that it is a detached shell due to an epoch of high mass loss ($\\sim$\\,5$\\times$10$^{-4}$\\,\\Myr) during the red supergiant or LBV phase, followed by a less intense mass-loss period. Concerning the central object, it has been spectroscopically studied in the optical and in the near-infrared wavelength range \\citep{Waters96, Voors00b}. Its LBV classification is supported by the high luminosity ($>$\\,10$^{5}$\\,\\Lsun), the moderate stellar wind (94\\,--\\,110\\,\\kms), and the current high mass-loss rate (10$^{-6}$\\,\\Myr) \\citep{Waters96}. Unfortunately, its effective temperature (T$_{eff}$) is not well-known because it is highly model dependent. The detection of a \\ion{He}{1} emission line near Br$\\alpha$ \\citep{Waters96} suggests a T$_{eff}$ well above 10,000\\,K. On the other hand, the LBV nature of the central star may limit the T$_{eff}$ to a maximum of 30,000\\,K. In Paper\\,I, a careful discussion about the distance to \\G\\ was done. We will assume a distance of 1.7\\,kpc throughout all the paper. The fingerprint of stellar evolution onto the circumstellar/interstellar medium (ISM) can provide a valuable input to learn about the processes, which govern the star itself and also affect the evolution of the surrounding gas and dust. The aim of this paper is to trace the distribution of warm dust, PAHs, and gaseous CO in relation with \\G\\ and its associated nebula. To do this, we present infrared Spitzer imaging and spectroscopy, and new submillimeter CO $J$\\,=\\,2\\,$\\rightarrow$\\,1 and 4\\,$\\rightarrow$\\,3 maps obtained with the IRAM 30m radio telescope and the Submillimeter Telescope ({\\it SMT}), respectively. In comparison to Paper\\,I, this new data set improves the angular resolution and the analysis of the physical conditions. ", "conclusions": "In order to shed light on the distribution of warm dust, PAHs and CO gas in \\G, we have analyzed infrared Spitzer data and observed the CO $J$\\,=\\,2\\,$\\rightarrow$\\,1 and 4\\,$\\rightarrow$\\,3 lines using the IRAM 30m radio telescope and the SMT. An analysis of the high-sensitivity MIPS 24\\,\\mic\\ image has led us to the discovery of a second outer dusty shell with a spherical shape, previously overlooked. In addition, the analysis of the SED of the central object has revealed the possible existence of a third extremely young dust shell closer to the LBV star, unresolved in the Spitzer images. Furthermore, the new CO data shown in this paper have revealed two separated gaseous shells located between the two dusty shells. The dynamic time of these gaseous components would indicate that the various shells have been formed during the current evolutionary stage. The presence of multiple shells around a LBV star is of great interest because its implications for stellar evolution theories. This would indicate that the mass-loss history has been more complex than previously assumed, with periods of enhanced and dimmed stellar winds instead of a continuous steady one. Using the IRS low-resolution spectra, we have compared three regions in the field of \\G: the central object, the shell, and the IRDC. The three studied regions present a clear similarity at the longest wavelengths ($>$\\,25\\,\\mic), due to the presence of a large amount of foreground/background dust at $\\sim$\\,65\\,K. A rich variety of solid-state features has been detected. Amorphous silicate features in absorption are present in the spectra of the central object and the IRDC, probably due to the foreground ISM and/or to the presence of amorphous silicate dust in a likely inner nebula very close to the central LBV star. PAHs features in emission are present in the spectrum of the shell, and a tentative detection of amorphous silicate in absorption is reported as well. Furthermore, a CO$_{2}$ ice feature in absorption is present in the spectrum of the IRDC. The IRS high-resolution spectra of the central object are full of fine structure lines, similar to those detected in other LBV stars. We have also detected pure rotational lines of H$_{2}$ arising from optically thin quadrupole transitions in two regions of \\G, the central object and the shell, which is the first detection of H$_2$ in the spectrum of a LBV star. Using different flux line ratios, we have analyzed the physical conditions of \\G\\ and its surroundings. In the case of the shell, indications of slow shocks have been found. We have estimated an incident UV field in the inner dusty shell of the order of 10$^{4}$ in units of Habing field. This high value of $G_0$, together with other evidence such as moderately high H$_2$ densities, the presence of PAHs in the shell, the detection of low-excitation fine structure lines, and the CO mid-J emission, suggests the existence of a PDR. This study of the multiple phase material directly associated with \\G\\ may help to understand its recent past and consequently the evolutionary mechanisms of this kind of object. In this sense, \\G\\ becomes a paradigm of LBV nebula and deserves further studies. In particular, a follow-up study of complex molecules should also be carried out, because they can provide information about the physical parameters in detail. In addition, similar studies to the one shown in this paper for other LBVs and evolved massive stars will also help to understand this short, but crucial, evolutionary phase of high-mass stars." }, "1003/1003.4987_arXiv.txt": { "abstract": "Data products from the Advanced Camera for Surveys Virgo Cluster Survey are used to understand the bulge star formation history in early-type galaxies at redshifts $z$~$\\ga$~$2$. A new technique is developed whereby observed high-redshift age-metallicity relationships are utilized to constrain the typical formation epochs of metal-rich or ``bulge'' globular clusters. This analysis supports a model where massive Virgo galaxies underwent an extremely intense mode of bulge globular cluster formation at $z\\sim3.5$ that was followed by an era of significant bulge growth and little globular cluster production. Intermediate-mass galaxies showed a less-intense period of globular cluster formation at $z\\sim2.5$ that was synchronized with the bulk of bulge star growth. The transition between the massive and intermediate-mass galaxy star formation modes occurs at a galaxy stellar mass of \\ms $\\sim3\\times10^{10}$ \\msun, the mass where many other galaxy properties are observed to change. Dwarf early-type galaxies in Virgo may have experienced no significant period of bulge globular cluster formation, thus the intense star bursts associated with globular cluster formation may be difficult to directly observe at redshifts $z\\la4$. Though the above conclusions are preliminary because they are based upon uncertain relationships between age and metallicity, the technique employed will yield more stringent constraints as high-redshift galaxy observations and theoretical models improve. ", "introduction": "Globular star cluster formation only occurs during strong star formation events \\citep{harris_supergiant_1994,elmegreen_universal_1997,larsen_mass_2009}. Most GCs found in early-type galaxies formed at redshifts $z\\ga2$ \\citep[see references in ][]{brodie_extragalactic_2006}, thus globular cluster (GC) systems are valuable observational tools to help understand the nature of major star formation events in early Universe \\citep{ashman_formation_1992,forbes_origin_1997,ct_formation_1998,harris_globular_2001,brodie_extragalactic_2006}. High-redshift galaxy observations continue to provide more and more detailed galaxy properties (e.g. ages, metallicities, sizes, star-formation rates) at the redshifts of GCs formation \\citep[e.g. ][]{hopkins_normalization_2006,bouwens_uv_2007,reddy_multiwavelength_2008,franx_significant_2003,van_dokkum_spectroscopic_2003,glazebrook_high_2004,daddi_passively_2005,cimatti_gmass_2008,maiolino_amaze._2008,van_dokkum_growth_2010}. Though much progress has been made in this field, interpretations are generally limited due to the challenging nature of such observations. By combining constraints from near-field GC system observations with far-field galaxy observations, certain limitations can be overcome and unique predictions for this important period of galaxy formation can be made \\citep[e.g. ][]{shapiro_star-forming_2010}. One of the strongest predictions resulting from GC system work is that many galaxies experienced two modes or epochs of intense star formation sometime before $z\\sim2$. This follows from observations that early-type galaxies with stellar masses of \\ms~$\\ga$~$10^{10}$~\\msun~generally host two GC metallicity subpopulations \\citep[e.g. ][]{strader_globular_2006,peng_acs_2006}. The metal-poor or ``halo'' GC subpopulation is thought to have formed within the numerous metal-poor proto-galaxies that likely dominated the early Universe. The metal-rich or ``bulge'' GC subpopulation likely formed later and shows similar properties to the host galaxy's bulge \\citep[e.g.~metallicities, spatial distributions][]{harris_globular_2001,forbes_connection_2001,harris_halo_2002,dirsch_wide-field_2005,bassino_large-scale_2006,forte_quantitative_2007,forte_globular_2009}. Thus by directly comparing the properties of a GC subpopulation to the galactic component that they are associated with, the relative contribution of e.g. bulge stellar mass produced during the intense MR GC formation events can be understood. Furthermore, this type of comparison can constrain more general galaxy formation models by also considering observational and theoretical analysis of high-redshift galaxies. The present work conducts such an analysis and compares properties of bulge GC subpopulations to their host galaxy's bulge stars. A model of bulge star formation is developed for redshifts $z\\ga2$, which is compared to existing constraints from direct, high-redshift galaxy observations. The GC system observations come from a imaging survey of 100 early-type galaxies in the nearby ($\\sim15$Mpc) Virgo Galaxy Cluster using the Advance Camera for Surveys mounted on the Hubble Space Telescope (HST). The sample size, its homogeneity and broad range of high-level data product resulting from the Advance Camera for Surveys Virgo Cluster Survey \\citep[ACSVCS; ][]{ct_acs_2004} allow for a detailed investigation into the relationship between galaxies and their GC systems, hence constraints on the first few gigayears of bulge star formation in Virgo early-type galaxies. \\subsection{Globular Cluster System Observations}\\label{review} Spectroscopy work to age-date individual GCs in elliptical galaxies has shown that the bulk of GCs are very old, with typical ages of $\\ga 10$ Gyrs or $z\\ga2$ \\citep[e.g. ][]{larsen_keck_2002,strader_extragalactic_2005,puzia_vlt_2005,beasley_globular_2006,conselice_keck_2006,sharina_ages_2006,pierce_gemini/gmos_2006,pierce_gemini/gmos_2006-1,cenarro_stellar_2007,beasley_2df_2008}. Because GCs are dominated by old stellar populations, the colour distribution of a galaxy's GC system can be interpreted as its intrinsic metallicity distribution \\citep[see ][]{strader_globular_2007,kundu_bimodal_2007,spitler_extendingbaseline_2008}. Observations show that most extragalactic GC systems have a bimodal colour distribution \\citep{strader_globular_2006,peng_acs_2006}, which implies they host two {\\it metallicity} subpopulations: metal-poor (MP) GC subpopulations with typical metallicities of [m/H] $\\sim-2$ to $-1$ and metal-rich (MR) GC subpopulations ranging [m/H] $\\sim-1$ to $0$. MR GCs show roughly similar chemical properties to the their host galaxy's bulge \\ms \\citep[e.g. ][]{harris_globular_2001,forbes_connection_2001,harris_halo_2002,forte_quantitative_2007,forte_globular_2009} and the spatial distributions of MR GCs resemble the host galaxy's stellar distribution \\citep[e.g. ][]{bassino_large-scale_2006}, though the GCs generally show a shallower distribution in the central regions of the galaxy possibly due to GC destruction \\citep[e.g. ][]{dirsch_wide-field_2005}. These observations suggest the formation of MR GCs and galaxy bulges are closely linked. MP GCs show metallicities more resembling galaxy halo stars \\citep[e.g. ][]{searle_compositions_1978} and are thus thought to have formed during a ``pre-galactic'' era before galaxies started to assembly in a $\\Lambda$CDM universe. Because the intrinsic metallicity distribution of early-type galaxy stellar metal-poor halos have only been studied in a few galaxies \\citep[e.g. ][]{elson_red_1997,harris_halo_2002,harris_leo_2007,norris_gemini/GMOS_2008,weijmans_stellar_2009,foster_metallicity_2009}, the relationship between MP GCs and their host galaxy is more difficult to discern and hence constraints from MP GC subpopulations will not be considered here. ", "conclusions": "\\label{disc} \\begin{figure} \\center \\includegraphics[scale=0.4,angle=0]{fig5} \\caption[Schematic diagram illustrating bulge formation histories derived from the analysis of MR GC subpopulations]{Schematic diagram illustrating bulge formation histories derived from the analysis of MR GC subpopulations in early-type galaxies of the Virgo Cluster. As a rough function of the host mass (see Fig.~\\ref{figmasshist}), these galaxies show three distinct bulge formation histories. The ``star-formation intensity'' is a intentionally ambiguous term that is meant to demonstrate the fact that the exact mechanism (e.g. local star-formation rate, gas turbulence, etc.) that determines the GC formation efficiency is unknown. The horizontal dotted line represents the threshold for GC production.}\\label{figsfh} \\endcenter \\end{figure} This section brings together the 3 topics of the preceding sections into a preliminary model of bulge star formation in Virgo early-type galaxies. Interpretations from direct high-redshift galaxy observations are compared to this model, followed by caveats and predictions. Fig.~\\ref{figsfh} shows the implied star-formation history of early-type galaxies in Virgo as developed from the preceding analysis on their MR GC subpopulations. The quantity on the y-axis represents the factor(s) that contribute to GC formation. It is intentionally left ambiguous because the detailed physics of GC formation are not well-understood. Instead the generic term ``star-formation intensity'' (SFI) is used as a proxy the mechanisms (e.g. star-formation rates, gas turbulence) that dictate GC formation efficiency within a galaxy. Given on each panel of Fig.~\\ref{figsfh} is a SFI threshold above which GC formation occurs. The scenario in the top panel of the Fig.~\\ref{figsfh} applies to very low-mass galaxies with no significant MR GC subpopulation. Here the GC formation epochs are unconstrained and the SFI never breaches the threshold for GC formation. Detecting strong star bursts in galaxies of \\ms $\\la 10^9 M_{sun}$ at redshifts $z\\la4$ may prove challenging. In the middle panel of Fig.~\\ref{figsfh}, the SFI history of intermediate-mass galaxies is given. These galaxies show similar bulge and MR GC metallicities, hence the formation epochs of the bulge stars and MR GCs are coeval. The peak of SFI in theses galaxies is represented as a Gaussian with mean redshift of $z\\sim2.5$ and spanning $z\\sim1.5-4$~or~$\\sim1$ Gyrs. This age-range matches inferences made from the observed MR GC metallicity distribution spread and the observed age-metallicity relationships (see Fig.~\\ref{figmm}). The implied peak of GC and bulge star formation apparently coincides with direct observations of the cosmic star-formation density peak (at redshifts $z\\sim2-3$), which some believe is dominated by star formation in low-mass galaxies \\citep[e.g. ][]{hopkins_normalization_2006,bouwens_uv_2007,reddy_multiwavelength_2008}. Here the SFI passes the GC formation threshold, to indicate GCs formed in these galaxies. For massive galaxies with offset bulge star and MR GC mean metallicities, the inferred SFI scenario is presented in the bottom panel of Fig.~\\ref{figsfh}. The metallicity offset implies that the bulge continued to grow after the peak of MR GC production, thus a second distribution in the SFI history is given at lower redshifts. The extent of this bulge formation epoch is unconstrained from the present analysis, but is chosen to not overlap significantly with the MR GC formation and drop off relatively fast to match observations that a significant number of quiescent massive galaxies appear to be passively evolving at $z\\sim2-3$ \\citep[e.g. ][]{franx_significant_2003,van_dokkum_spectroscopic_2003,glazebrook_high_2004,daddi_passively_2005,van_dokkum_confirmation_2008,cimatti_gmass_2008,bezanson_relation_2009,naab_minor_2009}. Although the SFI was low and few MR GCs formed, the bulge likely doubled in total \\ms and its global metallicity increased by a factor of 2 during this second mode of bulge growth. The main GC SFI distribution reaches very high values of SFI to signify the ultra-efficient period of MR GC production supported by the observation (see Fig.~\\ref{figtred}). The GC and bulge epochs overlap to illustrate the fact that some very enriched MR GCs exist. In massive galaxies, the peak of MR GC formation occurred at $z\\sim3-4$. This redshift is somewhat earlier than the observed peak in the number density of submillimeter galaxies \\citep[$z\\sim2.4$; ][]{chapman_redshift_2005}, whose properties can be explained by very intense star-formation rates of $\\sim1000$ \\msun~yr$^{-1}$ \\citep{hughes_high-redshift_1998,blain_submillimeter_2002,swinbank_properties_2008}. These strong star formation events were possibly induced by major galaxy mergers \\citep{tacconi_submillimeter_2008} or disk instabilities \\citep[e.g. ][]{escala_stability_2008,shapiro_star-forming_2010} from rapid, cold-gas accretion \\citep{dekel_cold_2009}. In the model developed here, this epoch is followed by a mode of bulge growth that is not as efficient at GC production. Observationally, this period may manifest itself as the heavily dust-obscured massive galaxies found at $z\\sim2$ \\citep[e.g. ][]{dey_significant_2008,bussmann_hubble_2009}, which just finished a submillimeter-bright phase \\citep[see discussion in ][]{dey_significant_2008}. At roughly the same time, approximately $z\\sim2.5$ or $\\sim11$ Gyrs ago, intermediate-mass galaxies were experiencing their peak of bulge growth {\\it and} MR GC formation. The implied SFIs may resemble those in massive galaxies at the same epoch. The average age of bulges in both massive and intermediate-mass galaxies should therefore be roughly identical, despite a relatively large difference in their mean metallicities. \\subsection{Caveats} The above scenario depends on the following: \\begin{itemize} \\item That there is no significant age difference between the galaxy bulge stars and MR GCs. Even though obviously young galaxies are removed from the analysis, transforming galaxy broadband colours into metallicities can be complicated by undetected age variations between the galaxies and MR GCs. \\item The assumption that the destruction of GCs, be it from disrupting tidal forces of the host galaxy or otherwise, does not influence the above results significantly. \\item The use of uncertain galaxy age-metallicity relationships from high-redshift observations. Environmental and morphological differences could translate into large systematic uncertainties on the attempt to age-date the peak of GC formation. \\item The assumption that the apparent offsets in bulge star and MR GC metallicities translates into an offset in time. Gas accretion, mixing and outflows contribute to inhomogeneities in the spatial metallicity distributions of a galaxy and will all sabotage the simple model presented here. \\end{itemize} The last two points can be better understood by incorporating detailed predictions from theoretic simulations of cosmological metal-enrichment and galaxy formation \\citep[e.g. ][]{hernquist_analytical_2003,dav_enrichment_2007,kobayashi_simulations_2007}. These issues will also benefit from more detailed observational analysis of galaxies at high-redshift \\citep[e.g. ][]{calura_evolution_2009}. \\subsection{Predictions} Presented below are a number of predictions resulting from the model outlined above. \\begin{figure} \\center \\includegraphics[scale=0.4]{fig6} \\caption[Galaxy \\ms histograms for old ACSVCS early-type galaxies]{Galaxy \\ms histograms for old ACSVCS early-type galaxies. Dotted histograms show the total \\ms distribution of the sample. Solid histograms in upper, middle and lower panel correspond to galaxies with no MR GCs, with no significant and a significant offset between the typical enrichment of its MR GCs and its bulge, respectively. These are classified according to the observed metallicity offset between their MR GCs and bulge stars, as shown in Fig.~\\ref{figcolordiff}. At \\ms $\\sim10^{10}$ \\msun, the 3 classes overlap.}\\label{figmasshist} \\endcenter \\end{figure} {\\bf Understanding the mode of bulge star formation around the transition galaxy stellar bulge mass.} The three bulge formation histories presented in Fig.~\\ref{figsfh} roughly correlate with galaxy \\ms, as shown in Fig.~\\ref{figmasshist}. At \\ms $\\sim10^{10}$ \\msun the three modes overlap. This bulge \\ms corresponds to the characteristic mass where many galaxy properties transition \\citep[e.g. ][]{dekel_galaxy_2006}. This work presents predictions for the bulge star formation histories of galaxies below, at and above this transition bulge \\ms. {\\bf GC and galaxy ages.} The derived MR GC formation redshifts in massive and intermediate-mass galaxies suggests an age offset of $\\sim1$ Gyr should exist in their MR GC subpopulations. Typical theoretical modelling and measurement uncertainties mean that absolute ages of individual extragalactic GCs are difficult to constrain with such precision. However, in principle, two very large samples of at least $1000$ MR GC ages (one for each galaxy-mass class) derived from standard Lick absorption-line analysis can be used to determine whether or not such a {\\it relative} age offset exists. This minimum sample size is derived assuming a mean MR GC age of 11.0 Gyrs (or log age $=1.04$ in Gyrs) and 11.7 Gyrs (log age $=1.07$) in intermediate and massive galaxies, respectively. Typical uncertainties are assumed to be $\\pm0.15$ on the logarithm of the age (in Gyrs) of individual GCs \\citep[e.g. ][]{proctor_keck_2008}. Note, systematic uncertainties from stellar population models and the age-metallicity degeneracy may complicate such attempts. Furthermore, the spread of metallicity may imply a spread of age, thus making this prediction more challenging to measure directly. Another interesting test of this model is to age-date the bulge stars of a massive galaxy (such as Virgo's Messier~87) and determine whether the mean age of 1000s of its MR GCs are offset from the bulge's average age by $\\sim1$ Gyr. A similar offset is expected in the Milky Way. {\\bf GC formation efficiency.} The discovery of an offset between the mean metallicity of MR GCs and their host bulge could mean that a significant fraction of the bulge \\ms in massive galaxies formed after their typical MR GCs. If the relative MR GC numbers observed in galaxies (see Fig.~\\ref{figtred}) were instead normalized by only the galaxy \\ms that formed in conjunction with the MR GCs and it is assumed that GC formation efficiency does not depend on metallicity, an extremely MR GC efficient formation in massive galaxies is implied. The apparent prevalence of ultra-compact dwarfs (which are sometimes considered as massive MR GCs; e.g. see \\citealt{forbes_uniting_2008}) around massive galaxies may be a natural by-product of an ultra-efficient GC formation epoch \\citep[see discussion in ][]{larsen_mass_2009}. This would mean that the local star-formation intensity history, not environment, dictates where such massive compact star clusters will be found \\citep{hau_ultra-compact_2009}." }, "1003/1003.3872_arXiv.txt": { "abstract": "We present a series of numerical simulations aimed at understanding the nature and origin of turbulence in coronal loops in the framework of the Parker model for coronal heating. A coronal loop is studied via reduced magnetohydrodynamics simulations in Cartesian geometry. A uniform and strong magnetic field threads the volume between the two photospheric planes, where a velocity field in the form of a 1D shear flow pattern is present. Initially the magnetic field that develops in the coronal loop is a simple map of the photospheric velocity field. This initial configuration is unstable to a multiple tearing instability that develops islands with $X$ and $O$ points in the plane orthogonal to the axial field. Once the nonlinear stage sets in the system evolution is characterized by a regime of MHD turbulence dominated by magnetic energy. A well developed power law in energy spectra is observed and the magnetic field never returns to the simple initial state mapping the photospheric flow. The formation of $X$ and $O$ points in the planes orthogonal to the axial field allows the continued and repeated formation and dissipation of small scale current sheets where the plasma is heated. We conclude that the observed turbulent dynamics are not induced by the complexity of the pattern that the magnetic field-lines footpoints follow but they rather stem from the inherent nonlinear nature of the system.\\\\[.5em] \\textit{Key words:} magnetohydrodynamics (MHD) --- Sun: corona --- Sun: magnetic topology --- turbulence\\\\[.5em] \\textit{Online-only material}: animations at \\href{http://www.df.unipi.it/~rappazzo/shear/}{http://www.df.unipi.it/$\\scriptstyle{\\sim}$rappazzo/shear/} (\\textit{non permanent link})\\\\ ", "introduction": "In recent papers \\citep{rved07,rved08} we have described reduced magnetohydrodynamics (RMHD) simulations of the Parker problem \\citep{park72,park88,park94} for coronal loops in Cartesian geometry. We have shown that the system develops small scales, organized in current sheets elongated in the direction of the DC magnetic field, through an MHD turbulent cascade, and that a well defined power law spectrum is developed for total energy. The energy spectra develop steep slopes [seen also in the similar simulations of a line-tied RMHD system by \\cite{dgm03}] with spectral indices going from the classical $-5/3$ Kolmogorov spectrum up to almost~$-3$. The energy spectral indices (slopes) have a bearing on the heating rate of the boundary-forced system. The heating rates are significantly increased for realistic values of the magnetic field intensity and loop length compared to the scaling laws with fixed indices \\citep{dg99}, the last being recovered when a Kolmogorov-like spectrum develops (i.e.\\ for weak fields or long loops). In the published simulations we have always used photospheric velocity patterns made up of large spatial scale projected convection cells (\\emph{large-scale eddies}), mimicking disordered photospheric motions. As a consequence the magnetic field developing in the corona was not in equilibrium, but dynamical evolution occurred from the outset with a time-scale set by the interplay of forcing and nonlinearity. In this paper, in order to clarify the origin of the MHD turbulent dynamics found in our previous work, we explore the dynamics of coronal loops system when a shear velocity flow stirs the footpoints of the magnetic field-lines. It is commonly thought that the topology of the photospheric driver should strongly influence the dynamics of a coronal loop, and that the magnetic field-lines anchored to the photospheric planes should \\emph{passively} follow their footpoints motions. In this picture the electric currents should develop along neighboring field-lines whose footpoints have a relative shear motion. In particular it might be argued that in our previous simulations the turbulent dynamics of the coronal medium originated from the ``complexity'' of the photospheric forcing patterns, with their large-scale eddies and stagnation points. The \\emph{simple} unidimensional shear forcing velocity used here allows a very \\emph{clear-cut numerical experiment}: if the field-lines were to passively follow the imposed footpoint motion only a sheared magnetic field would develop inside the volume. The dynamics of these current layers are subject to tearing instabilities \\citep{fkr63}. If this were the physical process at work the topology of the magnetic field would remain a \\emph{mapping} of the forcing velocity pattern, periodically disrupted by tearing-like instabilities when the shear grew beyond a threshold amount. Simulations show that initially the magnetic field is sheared and a tearing instability develops, but afterwards turbulent dynamics similar to those with more \\emph{complex} boundary patterns develop, clarifying that turbulence does not stem from a direct influence of the photospheric velocity pattern but it is due to the inherent nonlinear properties of the system. The Parker Scenario for coronal heating has been the subject of intense research. Both analytical \\citep{su81,vb86,ant87,ber91,gff92,cls97,nb98,uz07,low09, lj09,bat09,aa10} and numerical \\citep{msvh89,ls94,hvh96,gn96,hz09, hbz10} investigations have been carried out. Given the complexity of such boundary forced field-line tangling systems simplified models have been developed, including 2D incompressible MHD models with magnetic forcing \\citep{evpp96,dg97,gve98,dgd98,ev99} and shell models \\citep{nmcv04,bv07}. Insights gleaned from these models include the important result that dissipation in forced MHD turbulence occurs in the form of ``bursty'' events with well defined power law distributions in energy release, peak dissipation, and duration, in a way reminiscent of, and consistent with, the distribution of flares in the solar corona. An analytical model of a forced system very similar to the simulation presented here was proposed by \\cite{hp92}, and recently extended to the anisotropic turbulence regime by \\cite{bgp08}. They started with same MHD system threaded by a strong axial magnetic field in cartesian geometry and apply at the top and bottom boundaries two 1D velocity fields of opposite direction and assumed that the sheared structure that develops in the corona then dissipates via an effective ``turbulent resistivity'' provided by a cascade, so that a dissipative equilibrium is set up in which shearing is balanced by slippage provided by the turbulence. This amounts essentially to a \\emph{one-point closure} model of MHD turbulence \\citep{bis03}, where turbulence acts only on very small-scales while the large-scales remain laminar and indeed with the same large-scale magnetic structure. They then use the eddy-damped quasi-normal Markovian approximation (EDQNM) \\citep{pfl76} to estimate the effective cascade and dissipation for the given driving shear, and this allow them to develop a heating theory in which the only free parameter is the equivalent Kolmogorov constant. In contrast, our simulations will show that nonlinearity cannot be neglected even at the large-scales, so that the turbulent self-consistent state has little resemblance to the imposed shear flow. As a result \\cite{hp92} overestimate the heating rate as laminar dynamics would lead to a higher energy injection (Poynting flux). More recently \\cite{dka05,dlkn09} have proposed the so-called ``secondary instability'' as a leading mechanism operating in the Parker Scenario, responsible for the rapid release of energy. In their view, disruption on ideal time-scale must arise after some time while slow quasi-steady reconnection allows magnetic energy to continue to accumulate in the system. In their view the system evolution may be described by a sequence of equilibria destabilized by magnetic reconnection. The bulk of numerical simulations performed by \\cite{evpp96,dg97,gve98,dg99,ev99,dgm03,rved07,rved08} has proven that the system does not evolve through a sequence of equilibria, rather more complex dynamics develop. The initial setup of the simulation presented in \\citep{dlkn09} is also very similar to the one implemented in this paper but, besides the lower resolution and therefore the higher influence of numerical diffusion, the time interval for which they advance the equations is too short compared with coronal loops and active region time-scales. These leads them to claim as representative of the dynamics what is actually a \\emph{transient} event taking place only during the early stage of the dynamics, in our case the multiple tearing mode current sheet collapse. This evolution is not generic, but is only representative of a small class of very symmetric boundary velocity patterns, those which admit coronal equilibria at all times. We will return to this question and a more detailed discussion in the conclusion. The paper is organized as follows. In \\S~\\ref{sec:gebc} we describe the basic governing equations and boundary conditions, as well as the numerical code used to integrate them. In \\S~\\ref{par3} we discuss the initial conditions for our simulations and briefly summarize the linear stage dynamics more extensively detailed in \\cite{rved08}, while in \\S~\\ref{sec:ed} we outline the main points of \\cite{hp92} relevant to this work. The results of our numerical simulations are presented in \\S~\\ref{sec:ns}, while the final section is devoted to our conclusions and discussion of the impact of this work on coronal physics. ", "conclusions": "\\label{sec:con} In this paper we have investigated the dynamics of the Parker problem for the heating of coronal loops when the footpoints of the magnetic field-lines are stirred by a 1D shear velocity pattern at the photosphere-mimicking boundary, and compared these results with those previously obtained when a more complex ``vortex-like'' velocity pattern was imposed \\citep{rved08}. This very simple forcing is ideal to investigate the origin of turbulence in coronal loops and the influence of the boundary velocity forcing on the dynamics of the system. We will also compare our results with those of \\cite{hp92} and of the more recent simulations of \\cite{dka05,dlkn09}. In summary, the main results presented in this paper are the following: \\newcounter{ctr} \\begin{list}{\\arabic{ctr}.}{\\leftmargin=1.2em} \\usecounter{ctr} \\item Initially the sheared velocity forcing induces a sheared perpendicular magnetic field inside the volume. The resulting current layers are known to be unstable to tearing modes \\citep{fkr63}. In fact when the system transitions from the linear to the nonlinear stage it is due to a multiple tearing instability, as shown in Figure~\\ref{fig3}. But once the system has become fully nonlinear the dynamics are fundamentally different. As the nonlinear terms no longer vanish they now do transport energy from the large to the small scales where in correspondence of the X-points \\emph{nonlinear} magnetic reconnection takes place, without going through a series of equilibria disrupted by tearing-like instabilities. Similarly to the case with disordered vortical boundary forcing velocities \\citep{rved07, rved08} in the fully nonlinear stage the system is highly dynamical and chaotic (and increasingly so at higher Reynolds numbers). For this we do not observe secondary tearing of the current sheets as in 2D high-resolution simulations of decaying MHD turbulence \\citep{bw89}, as now at the small scales fast turbulent dynamics take place. \\item The dynamics of the Parker model do not depend strongly on the pattern of the velocity forcing that mimics photospheric motions, as far as they are constant in time (we defer the investigation of time-dependent boundary forcing to a future work). The shear forcing [eq.~(\\ref{eq:f0})] applied only at the top plate is a very simple and ordered one-dimensional forcing. We have shown that the resulting dynamics are very similar to those developed when a more complex and disordered ``vortex-type'' forcing velocity is applied. We conclude that the turbulent properties of the system are not induced by the \\emph{complexity} of the path that the footpoints follow. It is rather the system itself to be intrinsically turbulent, and turbulence develops as we continuously inject energy at the scale of photospheric motions ($\\sim 1,000\\, km$). \\item The system reaches a statistically steady state where, although the footpoints of the field-lines are continuously dragged by the forcing shear, this does not induce a sheared magnetic field in the computational box. In fact the topology of \\emph{the magnetic field is not a mapping of the forcing velocity field}. Nonlinear interactions are able to redistribute the energy that is injected only at the wavenumber $\\mathbf{n_{_\\perp}} = 4\\, \\mathbf{\\hat{e}_x}$ also to perpendicular wavenumbers and to smaller wavenumbers through an MHD turbulent cascade. In physical space this corresponds to the magnetic field being organized in magnetic islands, to small-scales formation (current sheets elongated along the axial direction) and to magnetic reconnection taking place. \\item Kinetic and Magnetic energies develop an inertial range where spectra exhibit a power-law behavior. Fluctuating magnetic energy dominates over kinetic energy. Spectra and integrated quantities, like energies and total dissipative rates, have values similar to those obtained with a vortex-type forcing velocity. In particular the total dissipation rate of the same system simulated with different Reynolds number appear to overlap each other beyond $Re = 200$, suggesting that this is independent of the Reynolds number beyond a threshold. \\end{list} As shown in Figures~\\ref{fig1}, \\ref{fig2} and \\ref{fig3} initially the system until time $t \\sim 79\\, \\tau_A$ follows the linear curves~(\\ref{eq:lin2}) and (\\ref{eq:diff1}). Up to this point the shear velocity at the top boundary induces a sheared magnetic field in the volume. As discussed in \\S~\\ref{sec:runa} we have introduced a perturbation mimicking those naturally present in the corona. With no perturbation the system would relax over the resistive diffusive timescale $\\tau_R$ ($\\sim 25\\, \\tau_A$ for run~A) in a saturated diffusive equilibrium as described in (\\ref{eq:lin2}) and (\\ref{eq:diff1}). While the simulation presented here used a very small amplitude for the perturbation ($\\epsilon = 10^{-16}$), we have performed shortest simulations with different values for the amplitude. As expected for higher values of $\\epsilon$ the instability develops sooner and for smaller values later, always following the linear curves until the instability transitions to the nonlinear stage. The more complete and systematic analysis of \\cite{rom04,rom09} in 2D confirms this behavior. \\cite{dlkn09} have performed a similar simulation with a lower resolution and with a fixed value for the perturbation and for a time interval that covers only the initial stage of our simulations. They in fact stop right after the first big dissipative peak, that in our Figures~\\ref{fig1}, \\ref{fig2} and \\ref{fig3} corresponds at $t \\sim 100\\, \\tau_A$. Continuing the simulation, and using a higher numerical resolution, the system reaches a statistically steady state where magnetic energy consistently fluctuates around a mean value and the shear is not recreated in the topology of the orthogonal magnetic field. Their analysis is then limited to a \\emph{transient} event taking place only during the early stages of the dynamics, and that afterward does not repeat. As shown in \\cite{rved08} during the linear stage the system is able to accumulate energy well beyond the average value maintained in the nonlinear stage only if the boundary forcing velocity satisfies the condition that its \\emph{vorticity is constant along the streamlines}. The sheared profiles used in this paper satisfy this condition as well the profile used by \\cite{dlkn09} (a linear combination of 6 sheared profiles). These profiles are a very small subset of all the possible forcing profiles, and while they are very useful to get insight into the origin of turbulence in coronal loops they are not representative of the disordered photospheric motions, for which the strong stress buildup required for secondary instability to develop does not take place. The significance of their conclusions is then strongly diminished. Furthermore the Parker angle for this system cannot be defined as the relative angle between magnetic field-lines at which the system becomes unstable. This is not a \\emph{well-posed} definition. In fact for given initial conditions the angle or equivalently the time (as the linear equation (\\ref{eq:lin2}) and (\\ref{eq:diff1}) imply) at which the instability develops depend on the value of the amplitude of the perturbation that we add to the system. Depending on the value of the perturbation the Parker angle so defined \\emph{is not unique}. On the other hand in the fully nonlinear stage the average magnetic field line magnitude fluctuates around a mean value. It is then possible to give a unique value for the Parker angle, defined now as the average inclination of the magnetic field-lines respect to the axial direction as done in \\cite{rved07,rved08}, and as originally introduced by \\cite{park88}. As summarized in \\S~\\ref{sec:ed} the one-point closure model developed by \\cite{hp92} splits the domain into large and small scales. They conjecture that the large-scale fields evolve into a stationary laminar regime, the field magnitudes determined by the effective diffusion coefficients. These laminar regimes correspond to our linear saturated diffusive regimes computed in \\S~\\ref{par3}. In Figure~\\ref{fig8} the dotted lines show such diffusive curves for different values of the Reynolds numbers. In their model the large-scale fields computed in this way are used to obtain $S$, the energy flowing into the system for unit time at the boundary (the power) due do the work done by photospheric motions on the magnetic field-lines footpoints. They also calculate, through an EDQNM approximation, the value of the spectral energy flux $\\epsilon$ flowing along the inertial range at the small scales. Both $S$ and $\\epsilon$ are functions of the effective diffusion coefficients, and the solution of the problem results requiring balance between the two powers $S = \\epsilon$ [$S$ and $\\epsilon$ have both the dimension of a power, i.e.\\ energy over time, as $S$ is the Poynting flux integrated over the boundary surface and $\\epsilon$ is integrated over the whole volume as in \\cite{rved07}]. As shown in our simulations the large-scale fields are not laminar, and they are stationary only statistically. Nevertheless it is useful to use \\cite{hp92} model in order to understand why it is not applicable. From Figure~\\ref{fig8} we can estimate that the \\emph{effective} Reynolds number for which the diffusive regime dissipation matches the dissipation of the simulated system is $R_{eff} = 150$. Unfortunately this values is too low for their model to work. In fact for $R = 150$ the dynamics are so diffusive that only a few modes of the order of the injection scale ($\\sim 1,000\\, km$) are not suppressed but only reduced in magnitude. Therefore there is no flux of energy at the small scales $\\epsilon = 0$. At a more fundamental level the idea to split the domain into large and small scales does not work because nonlinearity cannot be confined only at the small-scales. As shown by our simulations \\emph{nonlinearity} is at work at all scales, and unfortunately this fundamental aspect \\emph{cannot be circumvented}. Finally the use of RMHD equations is valid as far as the magnetic field fluctuations $\\mathbf{b_{_\\perp}}$ are small compared to the axial magnetic field $B_0$. This seems particularly apt to describe the dynamics of long-lived slender loops that apparently show no dynamics while shining bright at the resolution scale ($\\sim 800\\, km$) of current state-of-the-art X-ray and EUV imagers onboard Hinode and Stereo. Clearly these results do not apply to highly dynamical active regions where dynamics cannot be modeled as fluctuations about an equilibrium configuration. The series of simulations that we have performed proves that dragging the footpoints of magnetic field-lines in the Parker problem quickly triggers nonlinear dynamics for small values of the orthogonal magnetic fields, and that these small magnetic field fluctuations are able to transport a considerable amount of energy toward the small scales with the overall energy flux $\\sim 1.6 \\times 10^6\\, erg\\, cm^{-2}\\, s^{-1}$ \\citep{rved08} in the lower range of the observed constraint $\\sim 10^7\\, erg\\, cm^{-2}\\, s^{-1}$. This prevents the orthogonal magnetic fluctuations to grow to an arbitrarily high value, self-consistently limiting the dynamics of the Parker problem to small fluctuations if the initial conditions are given by a uniform strong axial magnetic field." }, "1003/1003.4513_arXiv.txt": { "abstract": "We derive the evolution equations describing a thin axisymmetric disk of gas and stars with an arbitrary rotation curve that is kept in a state of marginal gravitational instability and energy equilibrium due to the balance between energy released by accretion and energy lost due to decay of turbulence. Rather than adopt a parameterized $\\alpha$ prescription, we instead use the condition of marginal gravitational instability to self-consistently determine the position- and time-dependent transport rates. We show that there is a steady-state configuration for disks dominated by gravitational instability, and that this steady state persists even when star formation is taken into account if the accretion rate is sufficiently large. For disks in this state we analytically determine the velocity dispersion, surface density, and rates of mass and angular momentum transport as a function of the gas mass fraction, the rotation curve, and the rate of external accretion onto the disk edge. We show that disks that are initially out of steady state will evolve into it on the viscous timescale of the disk, which is comparable to the orbital period if the accretion rate is high. Finally, we discuss the implications of these results for the structure of disks in a broad range of environments, including high redshift galaxies, the outer gaseous disks of local galaxies, and accretion disks around protostars. ", "introduction": "In the past few years, observational and theoretical advances in many areas have led to intense study of the role of accretion and gravitational instability in determining the structure and rates of transport through disks. In the high redshift universe, clumpy star-forming galaxies at redshifts $z\\sim 2-3$ \\citep[e.g.][]{elmegreen04b, elmegreen05a, genzel08a, forster-schreiber09a, cresci09a} appear to be undergoing rapid accretion, and also have velocity dispersions that are much larger than those present in local galaxies. Numerical simulations \\citep[e.g.][]{dekel09b, ceverino09b, agertz09a, bournaud09a, mcnally09a} suggest that the large velocity dispersion and the massive clump morphology are both produced by a combination of gravitational instability and rapid external accretion. Around active galactic nuclei, radiative cooling pushes thin accretion disks into a state of gravitational instability \\citep{shlosman90a, goodman03a}, and in this state their accretion rates and structures may be determined by gravitationally-driven turbulence \\citep{gammie01a}. Closer to home, accretion disks around protostars of mass $\\ga 1$ $\\msun$ are expected to experience strong gravitational instability for a significant part of their lives \\citep{kratter08a} due to a combination of rapid accretion and strong radiative cooling. Numerical simulations indicate that the non-circular motions produced by this instability provide the dominant mechanism for mass and angular momentum transport in the disk \\citep[e.g.][]{krumholz07a, krumholz09c}. Dozens of simulations of gravitational instability in disks have been published, both for disks undergoing external accretion \\citep[e.g.]{vorobyov07a, vorobyov08a, vorobyov09a, kratter10a, machida10a} and for those in isolation (e.g.\\ \\citealt{lodato04a, lodato05a, kim07, cai08a, cossins09a, agertz09a}; see \\citealt{durisen07a} for a review of earlier work). Based on these, several authors have presented one-dimensional time-dependent disk evolution models in which the effects of gravitational instability are approximated by an $\\alpha$ prescription, with $\\alpha$ obtained by fits to simulation results or by general energy arguments \\citep{goodman03a, hueso05a, kratter08a, rice09a, rice10a}. Similar order-of-magnitude energy arguments have been extended to the case of galactic disks by \\citet{dekel09a}, \\citet{klessen09b}, and \\citet{elmegreen10a}. In the realm of purely analytic work, \\citet{bertin99a} present steady-state solutions for self-gravitating disks with decaying turbulence, while \\citet{rafikov09a} and \\citet{clarke09a} derived steady-state accretion rates for disks in balance between radiative cooling and accretion-driven heating in protostellar disks. While the analytic and one-dimensional models have provided a good understanding of the basic mechanism of gravitationally-driven turbulence and transport in disks, they also suffer from significant weaknesses. No analytic models published to date consider the case of gravitational instability-dominated disks that are time-dependent rather than in steady state. Most previous work has been limited to a particular rotation curve (e.g.\\ Keplerian or flat), to a disk of pure gas without stars, and to a disk that is vertically supported by thermal pressure rather than supersonic turbulence. As a result of these limitations it is not clear under what circumstances disks that are not in equilibrium can be expected to evolve into it, and it is not even clear what the equilibrium state is for a star-forming, supersonically turbulent disk such as a galactic disk. Our goal is to improve this situation by developing a first-principles theory for the evolution of a thin, supersonically turbulent disk of star-forming gas in a state of marginal gravitational stability, driven by a specified rate of external accretion. The only assumptions we make are (1) that the disk maintains a state of marginal gravitational instability ($Q=1$) at all times, and (2) the rate of energy loss due to radiative cooling can be parameterized as a certain fraction of the energy per crossing time of a disk scale height. We do not assume that the disk is in steady state or that it is characterized by any particular rotation curve. From these assumptions, in Section \\ref{sec:eqns} we derive equations describing the instantaneous, position-dependent rates of mass, energy, and angular momentum transport, and the time evolution of the disk surface density and velocity dispersion. In Section \\ref{sec:steady} we show that these equations admit an exact steady-state solution, and we derive the steady-state profiles of surface density, velocity dispersion, and transport of mass, energy, and angular momentum. In Section \\ref{sec:nonsteady} we show that disks that are not in the steady state will evolve toward it, and that for high accretion rates this evolution occurs on an orbital timescale. Finally, in Section \\ref{sec:discussion} we discuss the implications of our findings, and we summarize in Section \\ref{sec:conclusion}. ", "conclusions": "\\label{sec:discussion} \\subsection{Cosmological Evolution of the Velocity Dispersion of Galactic Disks} \\label{sec:diskevol} \\begin{figure} \\plotone{halosigma} \\caption{ \\label{fig:halosigma} Disk velocity dispersion $\\sigma$ versus redshift $z$ for halos of mass $M_{h,12} = 0.1$, $0.3$, and $1.0$ (blue, red, and black lines) and gas fraction $f_g = 1/2$ or $1$ (dashed and solid lines). The plot uses our fiducial $\\eta = 3/2$ and assumes $f_{b,0.18} = 1/2$, i.e.\\ half the infalling baryons are gas and half are stars. The $f_g = 1$ case is appropriate for systems where there no stars or where $\\sigma_* \\gg \\sigma$, such as present-day galactic disks, while the case $f_g=1/2$ is appropriate for high-redshift disks where $\\sigma_* \\approx \\sigma$. } \\end{figure} We have now shown that disks dominated by gravitationally-driven turbulence rapidly converge to an equilibrium state in which their velocity dispersions are determined by their gas fractions and accretion rates. Since this convergence happens on an orbital timescale, most galactic disks should be found near their equilibrium state. We can use this result, coupled to a simple model for how galaxy halos accrete mass, to study the evolution of disk velocity dispersions over cosmic time. Using Press-Schechter fits to dark matter simulations, \\citet{neistein08a} estimate the mean dark matter accretion rate onto halos of a given mass at a given redshift. \\citet{bouche09a} extend this to give an estimate of the gas accretion rate onto the disk at the center of the halo, which we adopt: \\begin{equation} \\label{mdotgas} \\dot{M}_g = 7.0\\epsilon_{\\rm in} f_{b,0.18} M_{h,12}^{1.1} (1+z)^{2.2}\\, \\msun\\mbox{ yr}^{-1}, \\end{equation} where $z$ is the redshift, $f_{b,0.18}$ is the gas mass fraction of the infall divided by 0.18, the universal baryon fraction, $M_{h,12}$ is the halo mass in units of $10^{12}$ $\\msun$, and $\\epsilon_{\\rm in}$ is the fraction of gas entering the halo that reaches the galactic disk rather than being shock-heated and joining the halo. This is approximately given by \\begin{equation} \\epsilon_{\\rm in} = \\left\\{ \\begin{array}{ll} 0.7 f(z), \\quad & M_{h,12} < 1.5 \\\\ 0, & M_{h,12} > 1.5 \\end{array} \\right., \\end{equation} where $f(z)$ is a function that is linear in time and varies from unity at $z=2.2$ to 0.5 at $z=1$.\\footnote{We compute the time as a function of redshift, and all other cosmology-dependent quantities, using $\\Omega_m = 0.28$, $\\Omega_{\\Lambda} = 0.72$, and $h=0.70$.} Inserting $\\dot{M}_g$ from equation (\\ref{mdotgas}) for $\\dot{M}_{\\rm ext}$ in equation (\\ref{sigmaeq}), we are able to evaluate the expected velocity dispersion of gravitational instability-dominated galactic disks as a function of halo mass and redshift. We do so in Figure \\ref{fig:halosigma}. Examining the plot, we see that for a Milky Way-like halo ($M_{h,12} = 1$, \\citealt{xue08a}) where $\\sigma_* \\gg \\sigma$ (so that the $f_g=1$ case applies), we predict a typical velocity dispersion of $10.7$ km s$^{-1}$. While this is very slightly higher than the value of $\\sigma \\simeq 8$ km s$^{-1}$ observed in typical Milky Way-like disks today (\\citealt{blitz04}, \\citealt{dib06a}, and references therein), the agreement is quite good given our purely analytic model. Our results are quite similar to the numerical ones obtained by \\citet{kim07} and \\citet{agertz09a}. Moreover, as \\citeauthor{agertz09a}\\ point out, gravitationally-driven turbulence has the advantage that it can operate even in the outer H~\\textsc{i} disk where there is very little star formation, so mechanisms such as supernovae that are invoked to explain turbulence in the inner disk \\citep[e.g.][]{de-avillez07a, joung09a} are unavailable. We also predict lower velocity dispersion in smaller halos, and this appears to be consistent with the somewhat lower H~\\textsc{i} velocity dispersions seen in dwarf galaxies \\citep{walter08a, chung09a}. Finally, however, we do note that there are alternative models to explain outer disk turbulence, including magnetorotational instability \\citet{sellwood99a, piontek07a} and accretion of clumpy gas \\citet{santillan07a}. Within the same framework we are able to explain the large velocity dispersions of $20-80$ km s$^{-1}$ found in galactic disks found at redshifts $\\sim 1.5-3$ \\citep{cresci09a}. The observed galaxies likely correspond to $\\sim 10^{12}$ $\\msun$ halos. For redshifts in this range and $f_g \\sim 1/2$, typical of galaxies at that redshift \\citep{daddi09a, tacconi10a} we predict typical velocity dispersions of $30-50$ km s$^{-1}$, with fluctuations at the factor of $\\sim 1.5$ level, corresponding to the expected factor of $\\sim 3$ level variations in the accretion rates of halos at the same mass and redshift. This is in good agreement with the observations. \\begin{figure} \\plotone{halovsigma} \\caption{ \\label{fig:halovsigma} Ratio of disk maximum circular velocity $V_{\\rm max}$ to velocity dispersion $\\sigma$ as a function of redshift $z$ for halos of mass $M_{h,12} = 0.1$, $0.3$, and $1.0$ (blue, red, and black lines, top to bottom) and gas fraction $f_g = 0.1$, $0.5$, or $1$ (dot-dashed, dashed, and solid lines). All parameters are the same as for Figure \\ref{fig:halosigma}. } \\end{figure} It is also instructive to compare the velocity dispersions we predict to the expected rotation velocities of galactic disks. We compute the approximate virial velocity of a halo as a function of mass and redshift following the approximation given in Appendix A2 of \\citet{dekel06a}, and we take the maximum circular velocity to be $1.2$ times this based on fitting the zero point of the Tully-Fisher relation \\citep{dutton07a}. With this approximation, we plot $V_{\\rm max}/\\sigma$ in Figure \\ref{fig:halovsigma}. We see that accretion-driven turbulence naturally produces the transition from disks with $V_{\\rm max}/\\sigma\\sim 5$ found at redshifts $\\ga 2$ to disks with $V_{\\rm max}/\\sigma \\sim 20-25$ found today. Interestingly, we find that there is little dependence of $V_{\\rm max}/\\sigma$ on halo mass. Instead, the primary dependence is in $f_g$, the gas mass fraction; analytically, $V_{\\rm max}/\\sigma \\propto f_g^{-1/3}$. Thus the most gas-dominated systems (or old galaxies that have $\\sigma_* \\gg \\sigma$) have the largest $V_{\\rm max}/\\sigma$, while gas-poor systems have smaller $V_{\\rm max}/\\sigma$. This suggests that the range of $V_{\\rm max}/\\sigma$ seen for galaxies at $z\\sim 2$ by the SINS survey \\citep{cresci09a} represents a sequence in gas fraction. The dispersion-dominated galaxies should on average be comparatively gas poor, while rotation-dominated ones should be gas rich. Of course fluctuations in accretion rate can also cause changes in $\\sigma$, so detecting this effect will require samples large enough for this noise source to be averaged out. Nonetheless, it seems likely that data to test this prediction will become available in the next few years. \\subsection{High-Redshift Galaxies} \\label{sec:highz} It is particularly interesting to apply our models to the $z\\sim 2-3$ galaxies observed by \\citet{elmegreen04b, elmegreen05a}, \\citet{genzel08a}, \\citet{forster-schreiber09a}, \\citet{cresci09a}, and others, since these are thought to be examples of strongly gravitational instability-dominated disks. We first note that, in the redshift range $z=2-3$ for halos of mass $M_h = 10^{12}$ $\\msun$, thought to be typical of the observed systems, the models shown in Figures \\ref{fig:halosigma} and \\ref{fig:halovsigma} give $\\chi = 8\\times 10^{-3} - 1.1\\times 10^{-2}$. For these values of $\\chi$ and gas fractions $f_g = 1/2$, using equations (\\ref{tacc}) and (\\ref{tsf}) the ratio of star formation time to accretion time $t_{\\rm SF} / t_{\\rm acc} = 1.8 - 2.6$, so the star formation rate is roughly $1/2 - 1/3$ of the total accretion rate. Given the uncertainties in this model and the dispersion in expected accretion rates, for simplicity we can simply adopt $\\dot{M}_* \\approx \\dot{M}_{\\rm ext}$. Since the star formation rates are observed (and have typical values $\\sim 100$ $\\msun$ yr$^{-1}$), we can plug them into our model in place of $\\dot{M}_{\\rm ext}$ in order to predict disk properties. Doing so, we find that the redshift $2-3$ disks should have velocity dispersions (from equation \\ref{sigmaeq}) \\begin{equation} \\sigma \\approx 47\\mbox{ km s}^{-1}\\, f_g^{-1/3} \\dot{M}_{*,100}^{1/3}, \\end{equation} where $\\dot{M}_{*,100} = \\dot{M}_* / 100$ $\\msun$ yr$^{-1}$, independent of their maximum rotation velocities $V_{\\rm max}$. Thus galaxies of similar star formation rate and gas fraction should have the same $\\sigma$ independent of $V_{\\rm max}$. The viscous accretion timescale required for the gas at the edge of one of these disks to reach the center is (from equation \\ref{tvisc}) \\begin{equation} t_{\\rm visc} \\approx 600\\mbox{ Myr } f_g^{2/3} R_{10} V_{200}^{-1} \\dot{M}_{*,100}^{-2/3}, \\end{equation} where $R_{10} = R/10$ kpc and $V_{200} = V_{\\rm max}/200$ km s$^{-1}$, and the gas mass is (from equation \\ref{coldeneq}) \\begin{eqnarray} \\nonumber M_g & = & R v_{\\phi} \\left(\\frac{f_g^2 \\dot{M}_*}{\\eta}\\right)^{1/3} \\\\ & = & 3\\times 10^{10}\\,\\msun\\; f_g^{2/3} R_{10} V_{200} \\dot{M}_{*,100}^{1/3}. \\end{eqnarray} The ratio of baryonic to dynamical mass within the disk region $R$ is \\begin{equation} \\frac{M_{\\rm bar}}{M_{\\rm dyn}} = 0.3\\, f_g^{-1/3} V_{200}^{-1} \\dot{M}_{*,100}^{1/2}. \\end{equation} To the extent that these quantities have been observed, they are in good agreement with the results of our model. It is also useful to verify that our approximation that $\\sigma_* \\approx \\sigma$ is valid for these galaxies. Once stars form, transient spiral structures will dynamically heat them until the stellar disk becomes stable against further spiral patterns \\citep{sellwood84a, carlberg85a}. The characteristic timescale for this heating is $[(Q_{\\rm lim} - Q_*)/\\tau] t_{\\rm orb}$, where $Q_{\\rm lim}\\approx 2$ is the limiting value at which the disk becomes stable against spiral perturbations, $Q_*$ is the current Toomre $Q$ parameter for the stars, and numerical experiments show that $\\tau\\sim 4-5$ for $t_{\\rm orb}$ evaluated at one disk scale length. If we assume that the stars are born at $Q_* = 1$, then the characteristic timescale over which the heat is \\begin{equation} t_{\\rm heat} \\approx 750\\mbox{ Myr } R_{10} V_{200}^{-1}, \\end{equation} where we have taken the radial scale length to be half of $R$. In contrast, the time required to double the stellar mass is \\begin{equation} t_{\\rm *} = \\frac{1-f_g}{f_g}\\left(\\frac{M_g}{\\dot{M}_*}\\right) = 300\\mbox{ Myr } \\frac{1-f_g}{f_g^{1/3}} R_{10} V_{200} \\dot{M}_*^{-2/3}. \\end{equation} Thus we see that the time required for stars to increase their velocity dispersion via spiral structure is generally comparable to or longer than the time required for a new generation of stars to form with the same velocity dispersion as the gas. Our approximation that the stellar population has the same velocity dispersion as the gas in these galaxies is therefore reasonable. \\subsection{Effects of Stellar Feedback} \\label{sec:feedback} In our idealized models, we have neglected the influence of stellar feedback by setting $\\calG = 0$. This is obviously reasonable if we are concerned with the outer parts of a galactic disk where there is no star formation, or a protostellar disk where at most a few stars will form. It is not reasonable for the centers of present-day galactic disks, where supernovae are clearly important. It is questionable whether stellar feedback is important in ULIRGs or the high surface density galaxies found in the early universe. Supernovae are not effective in such environments \\citep{thompson05, joung09a}, but stellar radiation pressure may be. Whether it can actually drive the observed velocity dispersions in high redshift galaxies is a matter of debate \\citep{murray09a, krumholz09d}. In those situations where feedback is significant, we can qualitatively see how it would change our results by noting that adding a non-zero $\\calG$ to our equations would have roughly the same effect as lowering $\\eta$. Physically, if star formation injects turbulence into the ISM at an appreciable rate, this is equivalent to reducing the rate at which turbulence decays -- we effectively increase the ``cooling time\" of the disk. Consulting equations (\\ref{sigmaeq}) -- (\\ref{alphaeq}), we see that the effect of this is to increase the velocity dispersion and surface density in the equilibrium state, while reducing the radial velocity and the rate of angular momentum transport. We caution that this analysis is only valid as long as the feedback is not too strong. In particular, we require that $\\calL$ remains larger than $\\calG$ for a $Q=1$ disk, and that the turbulent stresses created by the feedback mechanism are significantly weaker than the stresses induced by gravitational instability-driven turbulence. If the first requirement is not met, then feedback will drive the velocity dispersion up to the point where $Q>1$, and the gravitational instability will shut off. If the latter condition fails, then gravitational instability will continue, but our calculation of the transport rate will not be correct because we have not included stresses induced by feedback. Even if both requirements are met, our analysis of feedback effects should be regarded as qualitative rather than quantitative. Energy injection $\\calG$ appears on the right hand side of the torque equation with the opposite sign as $\\calL$, but their functional dependence on other disk parameters (surface density, velocity dispersion, etc.) are almost certainly different. The exact effects of feedback will depend on how energy injection varies with these quantities, which will in turn depend on the type of feedback and the physics of the ISM. \\subsection{Protostellar Disks} \\label{sec:protostellar} Although we have focused our discussion thus far on galactic disks, our model applies for arbitrary rotation curves, gas fractions, and infall rates, so we can apply it equally well to protostellar disks. To understand the expected levels of turbulence in protostellar disks, we use the parameterization of infall due to \\citet{kratter08a, kratter10a}, who introduce the dimensionless numbers: \\begin{equation} \\xi = \\frac{G \\dot{M}_{\\rm ext}}{c_{s,d}^3} \\qquad\\qquad \\Gamma = \\frac{\\dot{M}_{\\rm ext}}{M_{*d} \\Omega_{k,\\rm in}}, \\end{equation} where $c_{s,d}$ is the sound speed in the disk, $M_{*d}$ is the total mass of the disk and star and $\\Omega_{k,\\rm in}$ is the Keplerian orbital period of the infalling material. Physically, $\\xi$ represents the ratio of the external accretion rate to the maximum rate ($\\sim c_{s,d}^3/G$) at which a stable disk can process material, while $\\Gamma$ represents (up to a factor of $2\\pi$) the fraction by which the disk plus star mass changes per outer disk orbit. Indeed, since $v_{\\phi}(R) = R\\Omega_{k,\\rm in} = \\sqrt{G M_{*d}/R}$, with a little algebra it is easy to show that, in the case of a Keplerian disk consisting entirely of gas, our $\\chi$ simply reduces to \\citeauthor{kratter10a}'s $\\Gamma$ parameter. With this understanding, we can explain the observation by \\citet{kratter10a} that, in their simulations, the typical velocity dispersion of disks that do not fragment is comparable to the disk thermal sound speed (see their Figure 8). For a purely gaseous Keplerian disk, our model gives $s = [3\\chi/(4\\eta)]^{1/3}$, and \\citeauthor{kratter10a} show that the disk sound speed is related to the Keplerian velocity at the disk edge by $c_{s,d}/v_\\phi(R) = (\\Gamma/\\xi)^{1/3}$ (their Equation 18). Combining these two results, the expected Mach number of the accretion-driven turbulence is \\begin{equation} \\label{machdisk} \\mathcal{M} = \\frac{\\sigma}{c_{s,d}} = s \\frac{c_{s,d}}{v_\\phi(R)} = \\left(\\frac{3\\xi}{4\\eta}\\right)^{1/3}. \\end{equation} Since fragmentation is avoided only for disks with $\\xi$ of no more than a few, we can take $\\xi \\sim 1$, and it immediately follows that the expected Mach number $\\mathcal{M} \\sim 1$. We can apply a similar analysis to real protostellar disks: the Mach number of the turbulence in these disks should follow equation (\\ref{machdisk}). This means that disks accreting with $\\xi \\sim 1$, corresponding to $\\dot{M}_{\\rm ext} \\sim 10^{-5}$ $\\msun$ yr$^{-1}$ for typical out disk temperatures $T\\sim 50$ K, should have disks whose turbulent velocity dispersions are roughly transonic. This state should prevail during the majority of the main accretion phase. Once the main accretion phase ends and the accretion rate drops, the turbulent velocity dispersion should drop to subsonic values. This represents another prediction from our analysis: class 0 and class I protostars should have disks with transsonic turbulent velocity dispersions, while class II and class III sources should have subsonic turbulent velocity dispersions. As ALMA comes online in the next few years and provides resolved molecular line maps of protostellar disks at a variety of stages in their evolution \\citep[e.g.][]{krumholz07d}, we will be able to test this prediction. \\subsection{On the Validity of a Local Viscous Approximation for Gravitational Instability-Induced Transport} \\label{locality} The central approximation we make in our model is that transport of mass, angular momentum, and energy produced by gravitationally-driven turbulence can be represented with a local viscous stress tensor. The validity of this approximation has been the subject of great debate in the past decade. \\citet{balbus99a} show that self-gravitating disks cannot in general be modeled with a viscous formalism, but that such an approximation may be reasonable for disks near $Q=1$, the condition that we adopt throughout this work, and that appears to apply to the galactic and protostellar disks we are interested in studying. Based on a combination of analytic arguments and local simulations, \\citet{gammie01a} argues that a local prescription is applicable to $Q=1$ disks that are sufficiently thin, $s \\la 0.12$, and more recent global simulations \\citep{lodato04a, lodato05a, boley06a, cossins09a} generally support this result. \\citeauthor{gammie01a}'s condition for a local transport approximation to apply is well-satisfied for galactic disks at redshifts $z \\la 2$ (\\S~\\ref{sec:diskevol}) and for non-fragmenting protostellar disks (\\S~\\ref{sec:protostellar}). It is marginally violated for the observed disks at $z\\sim 2$ (\\S~\\ref{sec:highz}), suggesting that our model should be considered with some caution for them. At a minimum the thickness of these disks likely produces different fragmentation behavior than a standard thin disk analysis would suggest \\citep{begelman09a}." }, "1003/1003.0937_arXiv.txt": { "abstract": "We suggest a novel discretisation of the momentum equation for Smoothed Particle Hydrodynamics (SPH) and show that it significantly improves the accuracy of the obtained solutions. Our new formulation which we refer to as relative pressure SPH, {\\em rpSPH}, evaluates the pressure force in respect to the local pressure. It respects Newtons first law of motion and applies forces to particles only when there is a net force acting upon them. This is in contrast to standard SPH which explicitly uses Newtons third law of motion continuously applying equal but opposite forces between particles. {\\em rpSPH} does {\\em not} show the unphysical particle noise, the clumping or banding instability, unphysical surface tension, and unphysical scattering of different mass particles found for standard SPH. At the same time it uses fewer computational operations. and only changes a single line in existing SPH codes. We demonstrate its performance on isobaric uniform density distributions, uniform density shearing flows, the Kelvin--Helmholtz and Rayleigh--Taylor instabilities, the Sod shock tube, the Sedov--Taylor blast wave and a cosmological integration of the Santa Barbara galaxy cluster formation test. {\\em rpSPH} is an improvement these cases. The improvements come at the cost of giving up exact momentum conservation of the scheme. Consequently one can also obtain unphysical solutions particularly at low resolutions. ", "introduction": "The smoothed particle hydrodynamics method was invented by \\cite{1977AJ.....82.1013L} and \\cite{1977MNRAS.181..375G}, both with interests in astrophysical applications. Besides an enormous literature of successful application also many shortcomings of it have been presented in the literature \\citep{1993A&A...268..391S, 1994MmSAI..65.1013H, Swegle1995123, 2002ApJ...569..501I, 2007MNRAS.380..963A, 2009arXiv0906.0774R} often suggesting a fix to the reported problem \\citep[][to name but a few]{ Monaghan1994399, Cummins1999584, 2000PThPS.138..609R, 2007A&A...464..447A, Hu2007264, 2008IJNMF..56.1261G, 2008JCoPh.22710040P, 2009ApJ...701.1269B, Rafiee20092785, Xu20096703} Similarly there have been troubling news of how seemingly small differences in the initial setup led to very unexpected results \\citep[e.g.][]{Lombardi1999687}. This is all somewhat surprising given that in many of the cases where large inaccuracies have been found the only relevant equation (besides moving the particles $\\dot{\\vec{r}}=\\vec{v}$) stems from the pressure gradient accelerations \\be \\rho \\frac{D\\vec{v}}{Dt} = - \\nabla p, \\ee where $D/Dt$ denotes the Lagrangian derivative, and $p$ the pressure. In what follows we describe a new discretization of the momentum equation that avoids essentially all of the previously known problems of SPH. We will refer to this new method as ``relative pressure SPH'' or abbreviated as {\\em rpSPH}. We will first describe it, discuss implementation details and then present results for relevant tests highlighting the superior performance of this new approach. All the simulations shown here are carried out with Gadget-2~\\citep{2005MNRAS.364.1105S} (version 2.0.4) with only most minor changes explained in the text. The appendix describes how to convert Gadget to our {\\em rpSPH} formalism. ", "conclusions": "We have presented a novel discretization of the pressure equation for the smoothed particle hydrodynamics, which we call {\\em rpSPH} that removes the local pressure from the scheme and only considers pressure gradients. This methodology avoids the clumping and banding instability, artificial surface tension, unphysical particle noise, dramatically reduces inherent shear viscosity and numerical dissipation, and allows to realistically evolve density distributions sampled even with disparate particle masses. We have discussed a large number of test in all of which our new discretization outperforms the traditional SPH results. While our approach is not manifestly momentum conserving and easy to break wiht self-gravity and or low resolutions it clearly seems much more accurate than previous approaches to Lagrangian hydrodynamics using SPH. Since our formulation is more accurate and requires as little as one line of code to be changed in previous implementations we do believe it to likely to be useful. The caveat of {\\em rpSPH} remains that if one knows that one cannot afford to resolve the pressure gradient in ones initial conditions that because it is not momentum conserving it can give very wrong results. Fortunately, a resolution and convergence study can reveal whether one is in this limit. In summary, some of the biggest shortcomings of SPH can in some circumstances can be overcome if one gives up the idea of applying equal but opposite forces to particle pairs. While the latter is what happens physically to the atoms or molecules making up the fluid it is simply incorrect for Lagrangian fluid elements the particles are meant to represent. Physically it also does not make sense to introduce repulsive forces for two spatially separated points even when there is no pressure gradient between them. To require such symmetry between particles neglects that they are spatially separated and that the gradient of the pressure field is different at the two locations in general. Our new discretization avoids these unphysical forces and allows the SPH particles to behave as Lagrangian volume elements recovering fluid behaviour in a large number of tests. We have successfully used a fifth order spline kernel giving smaller errors on the uniform shear problem and the Rayleigh-Taylor problems discussed above. Consequently, we believe that further improvements to {\\em rpPSPH} should be possible in the future. We have also studied multiple forms of discretising the specific internal energy equation \\be \\frac{d\\epsilon}{dt} = - \\frac{P}{\\rho} \\nabla \\cdot \\vec{v} \\label{equ:ie}.\\ee The simplest version that we successfully applied to some of our test problems is given by \\be \\frac{d\\epsilon}{dt} \\approx \\sum_{j=1}^N \\frac{m_j}{\\rho_j} \\frac{P_j}{\\rho_j} \\frac{(\\vec{x}_i - \\vec{x}_j) \\cdot (\\vec{v}_i-\\vec{v}_j)}{|\\vec{x}_i - \\vec{x}_j|} {\\nabla_i W_{ij}(h_i)}. \\ee While we prefer the entropy formulation this form here may be useful for codes that start from an internal energy formulation. While standard SPH is conservative it fails to correctly capture fluid instabilities and shows large non-Newtonian viscosity. {\\em rpSPH}, on the other hand is more accurate, but is not inherently momentum or energy conserving. Consequently it is a useful modification to the SPH algorithm when one is studying problems where one can afford to resolve the relevant pressure gradients and the density field." }, "1003/1003.0089_arXiv.txt": { "abstract": "The new EAS Cherenkov array Tunka-133 with about 1 km$^2$ geometric acceptance area, is installed in the Tunka Valley (50 km from Lake Baikal). The array will permit a detailed study of cosmic ray energy spectrum and mass composition in the energy range of 10$^{15}$ - 10$^{18}$ eV with a uniform method. The array consists of 19 clusters, each composed of 7 optical detectors with 20 cm PMTs. Since November 2008, the first 12 clusters are in operation, commissioning of the whole array is planned for September 2009\\footnote{At the time of submission of this paper to the electornic arXiv(February 2010) the comleted Tunka-133 array is already taking data}. We describe the array construction and DAQ, preliminary results and plans for the future development: deployment of radio-antennas and muon detectors network. ", "introduction": "The elaborate study of primary mass composition in the energy range $10^{15} - 10^{18}$ eV is of crucial importance for the understanding of the origin and propagation of cosmic rays in the Galaxy. Since 2006 the work on creation of the EAS Cherenkov array Tunka-133 with a geometric area of 1 km$^2$ (\\cite{Tunka1}, \\cite{Tunka2}) is carried out in the Tunka Valley. Such an array will allow the investigation of cosmic rays in the energy range from 10$^{15}$ to 10$^{18}$ eV by a uniform method. This energy range includes the knee in the energy spectrum at 3$\\cdot$10$^{15}$ eV, and other features of the spectrum probably connected with the transition from galactic to extragalactic cosmic rays. For one year operation (400 hours) the array will register more than 300 events with energy above 10$^{17}$eV and the core position inside the array geometry. The use of algorithms of reconstruction of EAS parameters developed for the Tunka-25 array \\cite{Tunka25} will provide an accuracy of the measurement of EAS maximum depth $X_{max}$ $\\sim$ 25 g$\\cdot$ cm$^{-2}$. The Tunka-133 array has been detecting Cherenkov light from EAS over the last two years, with a steadily increasing number of detectors (7 detectors -- 2006, 28 detectors -- 2007 and 84 detectors -- 2008) and array effective area. The final array of 133 detectors will be completed by autumn 2009, and the internal effective area will reach 1 km$^2$. ", "conclusions": "The deployment of the Tunka-133 array is approaching its completion. More than 75\\% of detectors are installed and operating since the last winter season. Commissioning of the complete array is planned in autumn 2009.\\\\ The first results show that the array parameters are in a good agreement with the expected ones. Recording of the EAS Cherenkov light pulse waveform for each detector provides a more reliable measurement of $X_{max}$ and opens the possibility to reconstruct EAS with their core outside the array. Works for upgrading Tunka-133 with day-time operating detectors (radio-antennas and muon scintillation counters) have started." }, "1003/1003.3222.txt": { "abstract": "{We present high resolution large scale observations of the molecular and atomic gas in the Local Group Galaxy M33. The observations were carried out using the HEterodyne Receiver Array (HERA) at the 30m IRAM telescope in the \\mbox{CO(2--1)} line achieving a resolution of $12\\arcsec \\times 2.6\\kms$, enabling individual Giant Molecular Clouds (GMCs) to be resolved. The observed region is 650 square arcminutes mainly along the major axis and out to a radius of 8.5 kpc, and covers entirely the $2\\arcmin \\times 40\\arcmin$ radial strip observed with the HIFI and PACS Spectrometers as part of the {\\tt HERM33ES} Herschel key program. The achieved sensitivity in main beam temperature is \\mbox{20--50~mK} at $2.6~\\kms$ velocity resolution. The \\mbox{CO(2--1)} luminosity of the observed region is $1.7\\pm0.1\\times10^{7}~{\\rm K\\kms pc^2}$ and is estimated to be $2.8\\pm0.3\\times10^{7}~{\\rm K\\kms pc^2}$ for the entire galaxy, corresponding to H$_2$ masses of $1.9\\times10^8~{\\rm M_{\\sun}}$ and $3.3\\times10^8~{\\rm M_{\\sun}}$ respectively (including He), calculated with a $\\ratio$ twice the Galactic value due to the half-solar metallicity of M33. \\ion{H}{i}~21~cm VLA archive observations were reduced and the mosaic was imaged and cleaned using the multi-scale task in the CASA software package, yielding a series of datacubes with resolutions ranging from $5\\arcsec$ to $25\\arcsec$. The \\ion{H}{i} mass within a radius of 8.5 kpc is estimated to be $1.4 \\times 10^9{\\rm M_{\\sun}}$. The azimuthally averaged CO surface brightness decreases exponentially with a scale length of $1.9\\pm0.1$~kpc whereas the atomic gas surface density is constant at $\\Sigma_{\\ion{H}{i}}=6\\pm2~{\\rm M_{\\sun}}$pc$^{-2}$ deprojected to face-on. For a $\\ratio$ conversion factor twice that of the Milky~Way, the central kiloparsec H$_2$ surface density is ${\\rm \\Sigma_{H_{2}}=8.5\\pm0.2~{\\rm M_{\\sun}}pc^{-2}}$. The star formation rate per unit molecular gas (SF~Efficiency, the rate of transformation of molecular gas into stars), as traced by the ratio of CO to H$_{\\alpha}$ and FIR brightness, is constant with radius. The SFE, with a $\\ratio$ factor twice galactic, appears 2--4 times greater than of large spiral galaxies. A morphological comparison of molecular and atomic gas with tracers of star formation is presented showing good agreement between these maps both in terms of peaks and holes. A few exceptions are noted. Several spectra, including those of a molecular cloud situated more than 8~kpc from the galaxy center, are presented.} %\\abstract{We present high resolution large scale observations of the molecular and atomic gas in the Local Group Galaxy M33. The observations were carried out using the HERA multibeam receiver at the 30m IRAM telescope in the \\mbox{CO(2--1)} line achieving a resolution of $12\\arcsec \\times 2.6\\kms$, enabling individual Giant Molecular Clouds (GMCs) to be resolved. The observed region is 650 square arcminutes mainly along the major axis { and out to a radius of 8.5 kpc,} and covers entirely the $2\\arcmin \\times 40\\arcmin$ radial strip observed with the HIFI and PACS Herschel instruments as part of the {\\tt HERM33ES} Herschel key program. The achieved sensitivity in main beam temperature is 10-30~mK. The CO(2--1) luminosity of the observed region is $1.7\\pm0.1\\times10^{7}~{\\rm K\\kms pc^2}$ and is estimated to be $2.8\\pm0.3\\times10^{7}~{\\rm K\\kms pc^2}$ for the entire galaxy (or M(H$_2$) $= 3.3\\times10^8~{\\rm M_{\\sun}}$ for $\\ratioo$ ratio twice the Galactic value). \\ion{H}{i} 21~cm VLA archive observations where reduced and the mosaic was imaged and cleaned using the multi-scale task in CASA, yielding a series of datacubes with resolutions ranging from $5\\arcsec$ to $25\\arcsec$. The \\ion{H}{i} mass within a radius of 8.5~kpc is estimated to be $1.4 \\times 10^9{\\rm M_{\\sun}}$. The azimuthally averaged CO surface brightness decreases exponentially with a scale length of $1.9\\pm0.1$~kpc whereas the atomic gas surface density is constant at $\\Sigma_{\\ion{H}{i}}=6\\pm2~{\\rm M_{\\sun}}$pc$^{-2}$. For a $\\ratioo$ conversion factor twice the that of the Milky~Way in order to compensate for the half-solar metallicity of M33, the central kiloparsec H$_2$ surface density is ${\\rm \\Sigma_{H_{2}}=8.5\\pm0.2~{\\rm M_{\\sun}}pc^{-2}}$. The star formation rate per unit molecular gas (SF Efficiency), as traced by the ratio of CO to FIR brightness, is constant with radius. The SFE, the rate of transformation of molecular gas into stars, with a $\\ratio$ factor twice galactic, appears { 2-4 times} greater that of large spiral galaxies. A morphological comparison of molecular and atomic gas with tracers of star formation is presented showing good agreement between these maps both in terms of peaks and holes. A few exceptions are noted. Several spectra, including those of a molecular cloud situated more than 8~kpc from the galaxy center, are presented. } %PG \\abstract{We present high resolution large scale single dish observations of molecular gas in the Local Group Galaxy M33. The observations were carried out using the HERA multibeam receiver at the 30m IRAM telescope in the \\mbox{CO(2--1)} line achieving a resolution of $12\\arcsec \\times 2.6\\kms$. The observed region is 650 square arcminutes mainly along the major axis and covers entirely the $2\\arcmin \\times 20\\arcmin$ radial strip observed with the HIFI and PACS Herschel instruments as part of the {\\tt HERM33ES} Herschel key program. The achieved sensitivity in main beam temperature is 10-30~mK. The estimated CO(2--1) luminosity for the entire galaxy is $2.8\\pm0.3\\times10^{7}~{\\rm K\\kms pc^2}$ and the associated molecular gas mass $9.3\\times10^7~{\\rm M_{\\sun}}$. \\ion{H}{i} 21~cm VLA archive observations where reduced and imaged yielding a series of datacubes with resolutions ranging from 5 to $25\\arcsec$. A morphological comparison of molecular gas and star formation tracers is presented showing a good agreement between these maps. The radial distribution of the CO surface brightness, molecular and atomic gas mass surface density and star formation efficiency is studied showing an exponential decrease of the molecular gas with a scale length of $1.87\\pm0.09$~kpc, a constant $\\Sigma_{\\ion{H}{i}}=11\\pm2~{\\rm M_{\\sun}}$ atomic gas surface density, and a star formation efficiency nearly constant with radius. Several spectra, including those a molecular cloud situated more than 8~kpc from the galaxy center, are presented} ", "introduction": "\\begin{table} \\begin{minipage} {88mm} \\caption{\\label{tab.M33prop}Adopted parameters for M33} \\begin{tabular*} {88mm}{@{\\extracolsep{\\fill}}lr} \\hline\\hline\\noalign{\\smallskip} $\\alpha$(J2000) & $1^\\mathrm{h}33^\\mathrm{m}50\\fs9$\\\\ $\\delta$(J2000) & $+30\\degr39\\arcmin39\\arcsec$\\\\ Distance & 840~kpc\\,\\,\\footnotemark[1]\\\\ Optical Radius {\\rm $R_{25}$} & $30.8\\arcmin$\\,\\, \\footnotemark[2]\\\\ Inclination & $56\\degr$\\,\\, \\footnotemark[2]\\\\ Position angle & $22.5\\degr$\\,\\, \\footnotemark[2]\\\\ Central Oxygen abundance & $12 + \\log(O/H) = 8.4$\\,\\, \\footnotemark[3]\\\\ %1500 \\AA& $5\\arcsec$ & GALEX FUV \\footnotemark[4]\\\\ \\noalign{\\smallskip}\\hline \\end{tabular*} \\footnotetext[1]{\\citet{Galleti.2004}} \\footnotetext[2]{{\\tt HYPERLEDA} \\citep{Paturel.2003}} \\footnotetext[3]{\\citet{Magrini.2009}} \\end{minipage} \\end{table} The Local Group galaxies span a broad range in mass, luminosity, morphology, and metallicity. Two large spirals (the Milky Way and M31) are the centers of two galaxy sub-groupings, each being surrounded by a large number of dwarf galaxies. In addition, M31 --- the Andromeda Galaxy --- has a small spiral companion, M33 (the Triangulum Galaxy); their separation is approximately 15 degrees, corresponding to 200 kpc (assuming a common distance of 840 kpc; \\citet{Galleti.2004}). Gaseous streams are observed between them, indicating tidal interaction \\citep{Putman.2009}. %Of these, with the Milky~Way, M31 is the only other large spiral and M33 is the only small spiral, the others being irregular dwarf or spheroidal galaxies. M33 provides a means of observing a galaxy morphologically similar to our own but with a mass only a tenth of the Milky~Way and factor two lower metallicity { \\citep{Rosolowsky.2008,Magrini.2009}}. Further evidence for the difference between M33 and the Milky~Way is the large gas fraction and blue stellar colors of the former relative to the latter. M33 thus represents an environment in which to study the interstellar medium (ISM) and star formation (SF) that cannot be replaced by Galactic observations and where individual GMCs can be resolved to probe their star formation activity. It may also be possible to apply what we learn by studying M33 to the physics of early-universe objects, which share many of the characteristics of M33. In this article we present sensitive and high-resolution mapping observations of the CO $J = 2 \\rightarrow 1$ transition in M33 in order to study the morphology and dynamics of the molecular component. The total mapped area covers 650 square arcminutes, mainly along the major axis of the galaxy. %In a companion paper, a large sample of Giant Molecular Clouds (GMC) detected in these observations will be described (\u00c9) A $2\\arcmin \\times 40\\arcmin$ wide strip along the major axis (see Fig.~\\ref{fig.FUV_COcover}) was observed to a particularly low noise level of 25~mK at $2.6~\\kms$ velocity resolution to compare with the sensitive \\ion{C}{ii} Herschel/HIFI and Herschel/PACS spectroscopy observations which will be obtained as part of the {\\tt HERM33ES} Herschel Key Program \\citep{Kramer.2010}. While the most sensitive and among the highest resolution, these are not the first maps of M33 in the CO lines. \\citet{Engargiola.2003} observed the whole of the inner disk (up to about 5kpc along the major axis) with the BIMA array at $13\\arcsec$ resolution; \\citet{Heyer.2004} observed the inner disk and a small major axis strip at 50$\"$ resolution with FCRAO; \\citet{Rosolowsky.2007a} combined the BIMA$+$FCRAO$+$NRO data to improve the sensitivity and resolution of the previous maps; and \\citet{Gardan.2007} observed a rectangle at high sensitivity and $15\\arcsec$ resolution extending from NGC~604 to the R$_{25}$ radius. { Table \\ref{tab.surveys} summarizes the characteristics of previous molecular and atomic gas surveys in M33}. This work extends the \\citet{Gardan.2007} work further North and to the South at higher resolution and sensitivity. Earlier studies of GMCs in M33 include \\citet{Wilson.1997} and \\citet{Rosolowsky.2003} and studies similar to our own of other Local Group galaxies have been made by e.g. \\citet[][]{Fukui.2008,Israel.2003,Leroy.2006,Nieten.2006}. As mentioned by \\citet{Blitz.2006} for M33 and \\citet{Leroy.2006} for IC~10 and discussed more extensively by \\citet{Gardan.2007}, the Star Formation Rate (SFR) per unit H$_2$ mass or Star Formation Efficiency (SFE $=$ SFR${\\rm /M(H_2)}$ in~yr$^{-1}$) was found to be up to an order of magnitude higher in these small galaxies than in large spirals. This appears to be the case in distant galaxies as well, given the factor 10 increase in (commoving) SFR density \\citep[e.g.][]{Madau.1996, Wilkins.2008}. Are there local universe analogs of these distant objects? Is M33 one of them? One of the obvious questions is whether the H$_2$ mass has not been underestimated in these subsolar metallicity objects. The articles using the data presented here and as part of the {\\tt HERM33ES} project will attempt to answer that issue clearly. %Assuming the H$_2$ mass estimate is appropriate, what is the cause of the quick conversion of H$_2$ into stars? A further possibility might be that the stellar initial mass function (IMF) in M33 is top heavy, causing the SFR (and SFE) to be overestimated -- but then what changes the IMF? %The metallicity gradient in M33 is very weak, unlike most spirals, and we assume $12 + log(O/H) = 8.4-0.03 R_{~kpc}$ \\citep{Rosolowsky.2008, Magrini.2009} when necessary. Metallicities were lower in the past and H$_2$ production is believed to take place on grain surfaces -- therefore \\ion{H}{i} to ${\\rm H_2}$ conversion is expected to be less efficient in low metallicity systems \\citep{Krumholz.2009a}, and this is indeed observed in many systems \\citep[e.g.][]{Leroy.2007}. %therefore \\ion{H}{i} to H$_2$ conversion is expected to be less efficient in younger systems and indeed in the local universe the low-Z objects have lower H$_2$/\\ion{H}{i} fractions. The conversion of \\ion{H}{i} to H$_2$ in the intermediate redshift systems would have to be much more efficient than today, generating unusually high molecular fractions in these distant galaxies (contrary to expectations) with less stellar gravity, in order to have a similar efficiency in converting H$_2$ to stars. Here, we will present new maps of CO and HI. In this first paper, we restrict ourselves to a study of the radial distribution of the molecular gas and infrared surface brightness, the molecular and atomic gas surface densities and the star formation efficiency and to a qualitative comparison between the maps of star formation rate tracers, i.e. the dust maps, and the gas maps. In addition, we will discuss CO spectra in a few selected regions. A series of articles will follow, focussing on at least ($i$) cloud populations, life cycle, and mass spectrum ($ii$) dynamics of the molecular gas and the role of spiral arms ($iii$) diffuse CO emission, after subtraction of the clouds identified and ($iv$) a more detailed comparison of the star formation rate -- gas surface density relation. %Rather, we think that it is more likely that the initial conditions when the H$_2$ is formed are such that the subsequent collapse can happen more quickly and/or create stars with a higher-mass IMF (the tracers of star formation are biased towards high-mass stars). \\begin{figure} [tbp] \\begin{flushleft} \\includegraphics[angle=0,width=8.8cm]{figures/lores_FUV_COcover} \\caption{\\label{fig.FUV_COcover}The Local Group galaxy M33. This color image shows the GALEX FUV data which trace young stars and dust in the disk through attenuation. Overlaid are the observed fields (a) using the IRAM 30m Pico Veleta telescope in the CO(2--1) line (thick white outline) and (b) by HIFI/PACS instruments as part of the {\\tt HERM33ES} Herschel Key Program (thin yellow stripe). %In white: area covered by our CO(2--1) map overlaid on \\emph{GALEX} FUV data. In yellow: 2$\\arcmin$ wide strip along the major axis that will be observed by HIFI instrument as part of the {\\tt HERM33ES} Herschel Key Program } \\end{flushleft} \\end{figure} ", "conclusions": "This work presents high-resolution maps of the atomic and molecular gas in the disk of M33 via observations of the \\ion{H}{i}~21~cm and \\mbox{CO(2--1)} lines. The whole disk out to 8.5 kpc is covered in the \\ion{H}{i} line and about 60\\% of the emission in CO. Assuming the $\\ratio$ factor to be twice that of the Galaxy, because of the sub-solar metallicity of M33, we estimate a molecular gas mass of $3.3\\times10^8~{\\rm M_{\\sun}}$ roughly 20\\% of the $1.4\\times10^9~{\\rm M_{\\sun}}$ detected in the inner 8.5~kpc in \\ion{H}{i}. Azimuthally averaging, the \\ion{H}{i} surface density is close to constant with radius but the H$_2$ decreases exponentially with a scale length of 1.9kpc. The H$_2$/\\ion{H}{i} mass ratio decreases from about unity to 1\\%. The correspondence between the peaks and holes in the distributions of molecular and atomic gas is excellent and follows the peaks and troughs in the FIR, MIR, and H$\\alpha$ images. The SFE is approximately constant with radius, suggesting that molecular gas is transformed into stars at a similar rate (assuming a similar IMF) at all galactocentric radii. However, the SFE in the small, gas-rich, low-metallicity, blue spiral M33 appears {2-4 times} higher than what is observed in large nearby spirals. The sensitivity of the survey is such that CO emission is detected far out in the disk of M33 although few clouds are present and the lines are much weaker in intensity. %Lonely cloud still rather unique\\\\ %Excellent morphological correspondance between CO and the tracers of SF" }, "1003/1003.3236.txt": { "abstract": "From a deep multi-epoch \\CHANDRA\\ observation of the elliptical galaxy NGC 3379 we report the spectral properties of eight luminous LMXBs (\\LX$>1.2\\times10^{38}$\\ergps). We also present a set of spectral simulations, produced to aid the interpretation of low-count single-component spectral modeling. These simulations demonstrate that it is possible to infer the spectral states of X-ray binaries from these simple models and thereby constrain the properties of the source. Of the eight LMXBs studied, three reside within globular clusters, and one is a confirmed field source. Due to the nature of the luminosity cut all sources are either neutron star binaries emitting at or above the Eddington luminosity or black hole binaries. The spectra from these sources are well described by single-component models, with parameters consistent with Galactic LMXB observations, where hard-state sources have a range in photon index of 1.5$-$1.9 and thermally dominant sources have inner disc temperatures between $\\sim0.7-1.55$ keV. %The luminosity range of these sources is $1.5-10.1\\times10^{38}$\\ergps. The large variability observed in the brightest globular cluster source (\\LX$>4\\times10^{38}$\\ergps) suggests the presence of a black hole binary. At its most luminous this source is observed in a thermally dominant state with \\kt=1.5 keV, consistent with a black hole mass of $\\sim$4\\Msol. This observation provides further evidence that globular clusters are able to retain such massive binaries. % We also observed a source transitioning from a bright state (\\LX$\\sim 1\\times10^{39}$\\ergps), with prominent thermal and non-thermal components, to a less luminous hard state (\\LX=3.8$\\times10^{38}$\\ergps, $\\Gamma=$1.85). In its high flux emission this source exhibits a cool-disc component of $\\sim$0.14 keV, similar to spectra observed in some ultraluminous X-ray sources. Such a similarity indicates a possible link between `normal' stellar mass black holes in a high accretion state and ULXs. % %The $L-T$ relation has been investigated for all the sources in a TD state, with the parameter space of these LMXBs being consistent with Galactic observations, albeit at the bright end of the $L\\propto T^4$ relation. Source 42, the GC=BHB, was the only source to be observed in the TD state in multiple pointings. The $L-T$ relation of this object is flatter than the typical relation, with luminosity increasing with temperature as $\\sim~T^{1.25}$. Flattening has been observed in Galactic black hole binaries (Kubota \\& Makishima 2004) and it has been suggested that this could be due to the standard Shakura-Sunyaev accretion disc evolving into a slim disc. However, as the emission of this object is below Eddington, the transition to a slim disc is not expected to occur and instead, this flattening could be explained by a change in the disc color temperature relation. ", "introduction": "The discovery with \\CHANDRA\\ of several low-mass X-ray binary (LMXB) populations in early-type galaxies, and the associations of these LMXBs with either Globular Clusters (GCs) or the stellar field, have provided new impetus to the study of the formation and evolution of LMXBs in GCs and to the possible relation of field LMXBs to the GC population (Grindlay \\& Hertz 1985; Verbunt \\& van den Heuvel 1995; see Fabbiano 2006 and refs. therein). Given the characteristics of these data, most of this work has been based on population studies (e.g., Kim \\etal\\ 2006: Kundu \\etal\\ 2007; Sivakoff et al 2007; Voss \\& Gilfanov 2007). However, in the few cases of detection of luminous sources in deep enough observations, detailed spectral and variability studies can be pursued, to provide more direct constraints on the nature of the X-ray sources. One such example is the recent discovery of a variable luminous GC source in NGC 4472, with temporal and spectral characteristics supporting a stellar BH binary (Maccarone et al 2007; Shih et al 2008). In this paper we report the results of the spectral analysis of eight luminous sources (\\LX $> 1.2\\times 10^{38}$ \\ergps\\ in the 0.3$-$8.0 keV band), detected in the nearby elliptical galaxy NGC 3379 (in the poor group Leo, D=10.6 Mpc, Tonry \\etal\\ 2001\\footnote{Note that Jensen \\etal\\ (2003) report a distance of 9.82 Mpc. However, the distance of Tonry \\etal\\ (2001) is adopted here for consistency with Brassington \\etal\\ (2008). This choice does not affect our conclusions in any way.}) with \\CHANDRA\\ ACIS-S (Weisskopf et al 2000). These sources were observed at five different epochs, as part of a monitoring campaign with \\CHANDRA\\ (PI: Fabbiano) providing the rare opportunity of long-term spectral monitoring of LMXBs in an elliptical galaxy. These sources are part of the sample of 132 sources detected in NGC 3379 from these observations with luminosities greater than a few 10$^{36}$ \\ergps\\ (Brassington \\etal\\ 2008 $-$ hereafter B08). In section \\ref{sec:selection} we describe the observations and the properties of the sources under study from the B08 catalog. In section \\ref{sec:results} we describe our spectral analysis and report the results. In section \\ref{sec:sims} we present spectral simulations that we have performed to aid in the interpretation of our results and in section \\ref{sec:discussion} we discuss the spectral analysis and compare the results to our simulations. Our conclusions are summarized in section \\ref{sec:conclusions}. ", "conclusions": "\\label{sec:conclusions} In this paper we have investigated the spectral properties of eight bright (\\LX$>1.2\\times^{38}$ \\ergps) sources within NGC 3379, selected from the B08 catalog. From deep multi-pointing \\CHANDRA\\ observations we have been able to determine the source properties at different epochs, and thereby characterize source spectral variability. A set of simulations used to infer the spectral states of these sources has also been presented. Due to the luminosity cut imposed on this sample, all of the eight sources are either NS emitting at, or above the Eddington luminosity, or BHBs, with source luminosities ranging between $1.6-10.8\\times10^{38}$\\ergps. From optical correlations three sources have been determined to be BH-GCs, one is detected in the field, while the remaining four sources have insufficient optical data to classify a correlation. In all cases (bar S102) single component PO and DBB models have been applied to the data and have provided statistically acceptable fits. To aid the interpretation of these single component models, simulations of low-count (250, 500 and 1000) data have been presented. This work demonstrates that it is possible to determine if a source is in a hard, thermally dominant, or intermediate state by comparing the best-fit values of \\NH\\ from the two spectral models. In cases where the line-of-sight absorption in the PO fit is significantly higher than that of the Galactic value it is likely that a strong thermal component is present. If it is then determined from the DBB model, that this same source has a best-fit \\NH\\ consistent with the Galactic value it points to the source being in a TD state. If instead the DBB best-fit \\NH\\ tends to 0, it is likely that the source has both strong thermal and non-thermal components and can be said to be in an intermediate state. Although we add the caveat that whilst an elevated value of \\NH\\ from the single-component PO model is indicative of a thermally dominant state, intrinsic absorption, arising from the disc could also explain this result in some cases. We have also provided simulations of cool-disc spectra, allowing us to identify when a source is in a high accretion state with a cool-disc component. % A summary of how the best-fit values from single-component models can be interpreted is provided in Figure \\ref{fig:flow}. From these simulations it has been determined that when a source has prominent disc emission as well as a significant non-thermal component, the luminosity derived from the single-component DBB model should be taken as a lower limit of \\LX. The value inferred from the single-component PO model will be much greater than the true source value and should be disregarded. However, when a canonical PO model ($\\Gamma$=1.7 and \\NH\\ frozen to the Galactic absorption value) is applied to the data, the derived \\LX\\ value will provide the closest value to the true source luminosity. When the source is in a TD state the X-ray luminosity from the DBB should be adopted. We have also performed these simulation for a range of different absorption values, confirming that the results are not significantly dependant on the value of Galactic absorption, demonstrating that these simulations can be used to interpret a variety of extra-galactic binary spectra. Further, different absorption models have also been been used in a sub-set of simulations, and these results have confirmed that with low-count data the choice of absorption model does not affect the interpretation of spectral state. Four of the eight sources presented in this paper have been determined to exhibit spectral variability and observations have been subsequently been grouped to reflect these variations. By comparing the spectral properties of these groupings to our simulations, spectral transitions have been identified. The most luminous source detected within a GC, S42, has been determined to show significant variability (\\LX$>4\\times10^{38}$\\ergps) and can therefore be confirmed as a strong BHB candidate, predominantly residing in the thermally dominant state. The best-fit inner-disc temperature of this source is 1.5 keV, and we estimate a black hole mass of $\\sim$4\\Msol, providing further evidence (in addition to Maccarone \\etal\\ 2007 and Irwin \\etal\\ 2010) that GCs do retain black hole binaries. If we assume that the two further GC-LMXB sources in this paper are also individual black holes, this indicates that 30\\% of the GC-LMXB correlations in NGC 3379 contain a black hole binary, which constitutes a significant fraction of the GC-LMXB population. A transition from a luminous steep power law state to a hard state has been identified in S102 where we have identified a binary with a cool-disc component (\\kt=0.14 keV and $\\Gamma$=1.6) emitting around the Eddington luminosity of a 10\\Msol\\ black hole (\\LX$\\sim 1\\times10^{39}$\\ergps). Taken at face value, the temperature of this cool-disc could suggest an IMBH. This source then declines over the three months between pointings to a source in a hard state, with $\\Gamma \\sim1.85$ and \\LX=3.8$\\times10^{38}$\\ergps, consistent with NT state Galactic black hole binaries (RM06). The bright state of this source is similar to spectra observed in some ULXs (Soria \\etal\\ 2007), and indicates the similarities between these bright sources and more `normal' stellar mass black hole binaries in high accretion states. The spectral properties of the eight sources are largely consistent with the parameters that have been observed in Galactic LMXBs, with sources in a NT state exhibiting a range of $\\Gamma$=1.5$-$1.9 and sources in a TD state with \\kt=0.7$-$1.55 keV. We have also shown that this population is consistent with the general trend of increasing luminosity as sources become softer, transitioning from a NT to TD state. % The $L-T$ relation has been investigated for all the sources in a TD state, with the parameter space of these LMXBs (excepting the luminous cool-disc state of S102) being consistent with Galactic observations, albeit at the bright end of the $L\\propto T^4$ relation. S42, the BH-GC, was the only source to be observed in the TD state in multiple pointings. The $L-T$ relation of this object is flatter than the typical relation, with luminosity increasing with temperature as $\\sim~T^{1.25}$. Flattening has been observed in Galactic black hole binaries (Kubota \\& Makishima 2004) and it has been suggested that this could be due to the standard Shakura-Sunyaev accretion disc evolving into a slim disc, although this flattening could also be explained by a change in the disc color temperature relation. %%%!!!%%% Facility keyword add by GJS %{\\it Facility:} \\facility{CXO (ACIS-S)}" }, "1003/1003.5737_arXiv.txt": { "abstract": "The Jovian regular satellite system mainly consists of four Galilean satellites that have similar masses and are trapped in mutual mean motion resonances except for the outer satellite, Callisto. On the other hand, the Saturnian regular satellite system has only one big icy body, Titan, and a population of much smaller icy moons. We have investigated the origin of these major differences between the Jovian and Saturnian satellite systems by semi-analytically simulating the growth and orbital migration of proto-satellites in an accreting proto-satellite disk. We set up two different disk evolution/structure models that correspond to Jovian and Saturnian systems, by building upon previously developed models of an actively-supplied proto-satellite disk, the formation of gas giants, and observations of young stars. Our simulations extend previous models by including the (1) different termination timescales of gas infall onto the proto-satellite disk and (2) different evolution of a cavity in the disk, between the Jovian and Saturnian systems. We have performed Monte Carlo simulations and show that in the case of the Jovian systems, four to five similar-mass satellites are likely to remain trapped in mean motion resonances. This orbital configuration is formed by type I migration, temporal stopping of the migration near the disk inner edge, and quick truncation of gas infall caused by Jupiter opening a gap in the Solar nebula. The Saturnian systems tend to end up with one dominant body in the outer regions caused by the slower decay of gas infall associated with global depletion of the Solar nebula. The total mass and compositional zoning of the predicted Jovian and Saturnian satellite systems are consistent with the observed satellite systems. ", "introduction": "The four Galilean satellites around Jupiter have similar masses and the inner three bodies are trapped in mutual mean-motion resonances. They exhibit a trend of decreasing mean bulk density with increasing orbital radius, and a diversity of axial moments of inertia (Table 1). These properties are likely caused by a progressive increase in ice-to-rock ratio with an orbital radius and different states of differentiation. The inferred undifferentiated interior of the outermost satellite, Callisto, would require a relatively long accretion timescale, $5\\times 10^5$ years or more (Stevenson et al. 1986; Schubert et al. 2004; Barr \\& Canup 2008). On the other hand, Saturn has only one big icy satellite, Titan, which is located relatively far from Saturn. Recent Cassini data indicates that Titan also has an incompletely di\ufb00erentiated interior (Iess. et al. 2010) (Table 1). The origin of the pronounced difference between the Jovian and Saturnian systems is an intriguing question. Two important models for circum-planetary proto-satellite disks have been proposed. One is the ``solids enhanced minimum mass\" (SEMM) model (Mosqueira \\& Estrada 2003a, 2003b; Estrada et al. 2009) and the other is an actively-supplied gaseous accretion disk (CW) model (Canup \\& Ward 2002, 2006, 2009). The SEMM model postulates a massive quiescent disk with a peak surface density near $10^5$ g cm$^{-2}$ and a temperature profile determined by the luminosity of the gas giant planet. The Jovian/Saturnian satellites in this model form from solid materials supplied by ablation and capture of planetesimal fragments passing through the massive disk (Mosqueira et al. 2010). The CW model postulates a low mass, viscously evolving disk with a peak surface density near 100 g cm$^{-2}$, that is continuously supplied by mass infall from the circum-stellar proto-planetary disk onto the circum-planetary proto-satellite disk. The disk temperature profile in the CW model is dominated by viscous heating of the evolving disk and the luminosity of the giant planet plays a lesser role. The ``satellitesimals\" in this model are assumed to form immediately from dust grains that are supplied by gas infall (Canup \\& Ward 2006). Since the motions of the ``satellitesimals\" are decoupled from gas accretion, they are retained while the disk gas is actively replenished. As a result, a high dust-to-gas ratio is realized in the disk. The satellites that have grown massive enough are disposed through type I migration into the planet due to satellite-disk interaction (e.g., Ward 1986; Tanaka et al. 2002) and new generations of satellites are repeatedly accreted from the supplied dust grains. The finally surviving satellites are those formed at the very end of the host planet's accretion. In the present paper, we will focus on the latter model. Canup \\& Ward (2006) performed N-body simulations of accretion of satellites from small bodies and their type I migration in decaying accretion disks. They showed that the equilibrium total mass of satellites resulting from a balance between disposal by type I migration and repeated satellitesimal accretion is universally $\\sim 10^{-4} M_p$, where $M_p$ is the host planet mass. The proto-satellite disk gas is eventually depleted due to the decline of infall gas due to global depletion of the proto-planetary disk, then type I migration stops and satellites with total mass $\\sim 10^{-4} M_p$ survive. The result is consistent with actual Jovian and Saturnian systems. They suggested that the difference between Jovian and Saturnian systems is caused by stochastic timing between the repeated formation/migration and depletion of disk gas. The disk evolution in their model is regulated by the mass infall rate from the proto-planetary disk onto the proto-satellite disk (section 2.1) and the infall is responsible for the final stage gas giant planet formation. Although Canup \\& Ward (2006) assumed the same proto-satellite disk evolution for both Jovian and Saturnian systems, recent theories of gas giant formation and observational data of extrasolar planets raise the possibility that the disk evolution can be different between Jovian and Saturnian systems as follows. Unless the growth of gas giants is truncated or at least significantly slowed down at some critical mass, only a single gas giant should exist in each planetary system, because gas accretion onto a planet is a runaway process (e.g., Bodenheimer \\& Pollack 1986; Pollack et al. 1996; Ikoma et al. 2000). Cores for gas giants do not generally form at the same time, and the rate of gas mass accretion in the proto-planetary disk that can supply gas to giants is observationally $\\sim$ O$(10^{-5}) M_J$ year$^{-1}$ (where $M_J$ is a Jupiter mass). However, the observed data of extrasolar planets show that there are many systems with multiple gas giants. Our Solar system also has two gas giants. Furthermore, the mass distribution of extrasolar gas giants is centered at $M_{p,ave} \\sim \\mbox{a few} M_{\\rm J}$ with an upper cut-off of $M_{p,max} \\sim 10 M_{\\rm J}$. The observationally inferred proto-planetary disk masses have a mean value of $\\sim 10 M_J$ and a maximum value of $\\sim 100 M_J$ (e.g., Beckwith \\& Sargent 1996). These mean and maximum values are one order greater than $M_{p,ave}$ and $M_{p,max}$, respectively. These facts suggest that the giants did not accrete most of disk gas, so they are consistent with existence of the rapid truncation (e.g., Lin \\& Papaloizou 1985) or severe decline (e.g., Lubow \\& D'Angelo 2006; Tanigawa \\& Ikoma 2007) of gas infall onto the planets due to gap opening in the proto-planetary disks. Since the proposed gap opening conditions show that the critical mass for gap opening is larger in outer regions (section 2.2), outer gas giant planets, in principle, tend to be larger than inner ones. However, the outer one can be smaller, if the disk gas is depleted before the outer planet can complete its formation. Since Saturn is three times less massive than Jupiter, it is most likely that Jupiter opened up a gap to halt its growth while Saturn did not and its growth was terminated by global depletion of the proto-planetary disk (the Solar nebula). Although reduced gas infall can still continue after the gap opening (Lubow \\& D'Angelo 2006; Tanigawa \\& Ikoma 2007), the severe reduction in infall rate makes the Jovian disk evolution significantly different from the Saturnian one (see section 2). Here, we explore a possible path to produce the pronounced different architectures of the Jovian and Saturnian satellite systems, by introducing the different evolution of proto-satellite disks due to the gap opening to the actively-supplied gaseous accretion proto-satellite disk model (Canup \\& Ward 2002, 2006). The purpose of the present paper is to demonstrate how this diversity in gas giant formation can profoundly affect the regular satellite formation process. Since we survey a wide range of parameters for initial and boundary conditions, we adopt a semi-analytical model to simulate accretion and migration of satellites from satellitesimals that has been developed by Ida \\& Lin (2004, 2008) for modeling sequential planet formation. While the semi-analytical calculations inevitably introduce approximations, N-body simulations have to employ unrealistic initial and/or boundary conditions because of the computational limitations (e.g., Kokubo \\& Ida (1998) started from planetesimals of $\\sim (10^{-4}-10^{-3}) M_{\\oplus}$, or Canup \\& Ward (2006) replaced infalling dust grains by embryos with the isolation mass of an accreting satellite). We show in section 4.1 that our calculation produces results consistent with Canup \\& Ward (2006)'s N-body simulation, although inconsistency may exist in detailed features (we need more careful comparison to reconcile the inconsistency). In section 2, we explain our disk evolution models for Jovian and Saturnian systems in details, highlighting their differences. In section 3, semi-analytical treatments to simulate accretion of satellites from satellitesimals, orbital migration, and resonant trapping are described. The simulation results are presented in section 4. We successfully explain the pronounced difference between Jovian and Saturnian systems in the framework of our model. We also discuss implications for the formation of Saturn's rings. Section 5 is devoted to a summary of the major results. ", "conclusions": "We simulated growth and orbital migration of proto-satellites in an accreting circum-planetary disk that is modified from the disk model by Canup \\& Ward (2002, 2006), in order to address the different architectures between Jovian and Saturnian satellite systems: \\begin{enumerate} \\item Jovian system has four similar-sized satellites (Galilean satellites) locked in mean-motion resonances except the outermost one (Callisto), while Saturnian system has only one big satellite, Titan, in outer region. \\item Inner two satellites are rocky and outer two are mostly icy for Jovian system, while Titan is mostly icy. \\end{enumerate} We modified the semi-analytical model for the population synthesis model for planet formation developed by (Ida \\& Lin 2004, 2008) to apply the model for satellite formation, introducing the effect of resonant trapping that characterizes Galilean satellites. The semi-analytical model reproduces the results consistent with N-body simulations by Canup \\& Ward (2006), although details such as mass distribution of surviving satellites are not necessarily consistent (it is a future work to reconcile the inconsistency). We considered a coupled system of gas giant planet formation and satellite formation around the planet, although planet formation is not simultaneously simulated. The fact that multiple gas giant systems are found in extrasolar systems in addition to our Solar system suggests that gas accretion onto gas giants is truncated or at least significantly reduced by gap formation in the proto-planetary disk. Since Saturn is smaller than Jupiter and it is theoretically expected that Saturnian core is formed later than that of Jovian core, we consider a working hypothesis that {\\it Jupiter opened up a gap while Saturn did not.} Accordingly, we assume the followings: \\begin{enumerate} \\item Gas infall onto a proto-satellite disk is truncated or significantly reduced at the time of gap opening for Jupiter while that onto Saturnian disk slowly decays as the proto-planetary disk (the Solar nebula) globally decayed on a timescale of $\\tau_{\\rm dep} \\sim 1$-10 Myrs. The Jovian disk decays on its viscous diffusion timescale of $\\tau_{\\rm diff} \\sim 10^3$ years after the gap opening. \\item Based on inference from observations of different spin periods between CTTSs and WTTSs, the proto-satellite disk has an inner cavity in the early, high disk accretion stage and the cavity vanishes in the subsequent, low disk accretion stage. \\end{enumerate} Since typical type I migration timescales are much longer than $\\tau_{\\rm diff}$ and much shorter than $\\tau_{\\rm dep}$, the Jovian satellite system may have been ``frozen'' at the time of gap opening in which disk mass and temperature were relatively high while the Saturnian satellites may be final survivors in the last stage in which disk mass and temperature were relatively low. Thus, we carried out simulations with a cavity and abrupt decay of disk gas for Jupiter and those without a cavity and with slow disk depletion on timescales of 1-10 Myrs for Saturn. Then, we have found the following results: \\begin{enumerate} \\item In Jupiter-condition systems, four or five similar-sized satellites are formed in $80\\%$ of the runs and the inner ones tend to be locked in mean-motion resonances due to type I migration and its stoppage at the disk inner edge. In Saturn-condition systems, predominant bodies repeatedly accrete in outer regions and migrate onto the surface of the host planet, sweeping up inner small satellites. In $70\\%$ of runs, only one big body remains. The big bodies have more than $95\\%$ of total mass of remaining bodies in the system in about half of the runs. \\item Because Jupiter-condition systems are formed in relatively hot disks, the inner two satellites tend to be rocky and outer ones are ice-rich. On the other hand, in Saturn-condition systems, the finally surviving bodies form in the longer lasting cold disks and are therefore ice-rich. Moreover, the surviving massive bodies coresponding to Titan were usually formed on long timescales ($\\sim$ Myrs), which is consistent with the incomplete differentiation of Titan's interior. \\end{enumerate} If we consider a small amount of residual mass infall onto the circum-Jovian disk after the gap opening, the formation timescale of the outermost satellite is longer than Myrs, which is consistent with incomplete differentiation of Callisto's interior. Thus, our simulations produced the Galilean satellites' analogues and the Titan's analogues in each setting with high probability. We have demonstrated that our assumption on gap formation and associated existence of a disk cavity naturally explain the different architectures between Jovian and Saturnian satellite systems, although more detailed fluid dynamical simulations are needed for the radial dependence and its time evolution of infall onto proto-satellite disks and gap formation in proto-planetary disks, and more detailed MHD simulations are necessary for evolution of a cavity. The different architectures may be fossil evidence that Jupiter opened up a clear gap in the proto-planetary disk to terminate its growth. This result gives deep insights into the mass distribution and multiplicity of extrasolar gas giant planets (Ida \\& Lin 2008) as well as into the observed bimodal distributions of spin periods of young stars (Herbst \\& Mundt 2005). Notice that the Jovian/Saturnian satellite systems are analogous to extrasolar super-Earth systems, because super-Earths accrete only solid materials and have masses of $\\sim (10^{-5}$--$10^{-4})M_{\\odot}$. A high fraction of solar-type stars may have close-in super-Earths (e.g., Mayor et al. 2009), while many systems including the Solar system do not have any close-in super-Earths. Our results should also give deep insights into the origin of the diversity of super-Earth systems found around other stars." }, "1003/1003.4507_arXiv.txt": { "abstract": "{Direct-imaging searches for planets reveal wide orbit planets amenable to spectroscopy, and their atmospheres represent an important comparison to the irradiated atmospheres of Hot Jupiters. Using AO integral field spectroscopy of 2M1207\\,b, the shape of the continuum emission over the $J$, $H$, and $K$ bands from the atmosphere of this young, planetary mass companion is measured in order to compare with atmospheric and evolutionary models, and objects of similar temperature in young clusters and the field. The 2M1207\\,b spectrum has the highest spectral resolution (R$\\sim$300--1500) and largest wavelength coverage, including the first $J$-band spectrum, for this benchmark object. The high signal-to-noise of the data allow a clear identification of signatures of low surface gravity, and comparison with a grid of AMES-Dusty models reveals a best-fit effective temperature of \\teff=1600\\,K with a preferred surface gravity of \\logg=4.5. The $J$-band flux is depressed relative to nearly all L-type objects, and the detailed shape of the absorption features across the $H$-band exhibit differences from the model predictions. The possible origins of 2M1207\\,b and its low luminosity are examined with the new data and analysis which suggest that extinction from a disk with large grains is a viable scenario and is preferred over scatttering off an optically thick disk. The 2M1207\\,b spectrum presents an important comparison for the types of features which may be present in upcoming spectra of the atmospheres of planets imaged in orbit around stellar primaries.} ", "introduction": "Over 400 extrasolar planets are currently known\\footnote{www.exoplanet.eu}, with most systems detected indirectly in orbits closer than that of Jupiter. For indirectly detected planets, the rare combination of a favourable geometry resulting in a transit and a very bright host star for high signal-to-noise data is required to obtain a planetary spectrum either in transmission \\citep{c02} or emission \\citep{ch05,d05}. Although only two transiting planets have been investigated spectroscopically, transmission studies have shown the possibility of methane absorption \\citep{sw08} or Rayleigh scattering \\citep{s08} and a haze layer \\citep{p08}, and emission studies have observed the presence \\citep{g08} and absence \\citep{r07} of water band absorption in different systems. The atmospheres of directly-imaged planets represent a complementary case study to the transiting planet spectroscopy, uncontaminated by the physical effects of extreme proximity to the host star. In some cases, spectroscopy of directly imaged planets is feasible. The first confirmed directly-imaged planets orbiting normal stars \\citep{k08,m08} and the first imaged planetary mass companion \\citep{c04} present the opportunity for spectroscopic analysis of young planet atmospheres. For technical reasons, targets observed in imaging planet searches are typically younger than stars searched for planets indirectly, and this selection criterion enables an investigation of the early evolution of planetary atmospheres. The first spectrum of HR\\,8799\\,c in the $L$-band suggests differences in the continuum shape from theoretical models \\citep{j10}, possibly due to non-equilibrium CO/CH$_4$ chemistry and indicating the importance of measuring the atmosphere over an even larger wavelength range. This paper presents the results of the highest resolution $J$, $H$, $K$ spectrum of 2M1207\\,b which should provide a high signal-to-noise comparison spectrum to planets imaged in orbit around stellar primaries. \\begin{table*} \\caption{SINFONI observations} \\label{table:3} \\centering \\begin{tabular}{l c c c c l l l l l} % \\hline\\hline Source & Grating & Scale & NDIT & DIT & $\\#$exp. & Standard (HIP) & \\teff\\ (K) & Dates & Project ID \\\\ \\hline 2M1207A+b & $H$+$K$ & 100 & 1 & 300\\,s & 24 & 54930, 62566 & 29000, 15000 & 28\\,Jan\\,07, 07\\,Feb\\,07 & 078.C-0800(B) \\\\ & $J$ & 100 & 1 & 600\\,s & 6 & 59951 & 16400 & 22\\,Feb\\,07 & 078.C-0800(B) \\\\ \\hline \\end{tabular} \\end{table*} ", "conclusions": "\\label{results} The larger coverage in wavelength and higher signal-to-noise enables an investigation of the atmosphere and origin of this young, planetary mass companion. The flux calibrated spectrum of 2M1207\\,b is compared with previous observations of field and young low mass objects and with theoretical model atmospheres. The overall shape is significantly different to field L dwarfs, with a marked depression in the $J$-band as shown in Fig.\\,\\ref{FieldComp}. The $H$-band continuum exhibits a distinct triangular shape unlike the older, higher surface gravity objects, and the $K$-band peaks at $\\sim$2.3\\micron\\ \\citep[as suggested by][]{m07} rather than at $\\sim$2.15\\micron\\ as with field L dwarfs, and the region associated with the peak is more clearly defined in these higher resolution data. The spectrum of 2M1207\\,b is unlike field L-dwarfs spanning a range of spectral types, but the young age of the TW Hydra association (8 Myr) complicates the comparison. In Fig.\\,\\ref{YoungComp}, the spectrum of 2M1207\\,b is plotted with several spectra of low mass members of Upper Sco \\citep{l08}, with spectral types ranging from M8--L2, and with the borderline brown dwarf/planet companion AB Pic\\,b \\citep{b10}. Even amongst young, low-mass objects in the 5\\,Myr Upper Sco region, the 2M1207\\,b spectrum appears distinct in the shape and relative fluxes of the continuum across the different bands. No Upper Sco member has a lower $J$-band relative to $H$-band and the triangular shape of the $H$-band is only beginning to develop by the L2 spectral type. The spectral sequence shown in Fig.\\,\\ref{YoungComp} suggests that 2M1207\\,b is one of the coolest young companions imaged. The most analogous object to 2M1207\\,b is the companion AB Pic\\,b which also shows some of the key features present in 2M1207\\,b. Since the $J$-band data required binning, the spectral resolution is not yet high enough to identify and measure absorption lines, though future observations with more signal-to-noise could investigate the depths of the \\ion{Na}{I} and \\ion{K}{I} lines. A series of spectral indices based on ratios of the flux of different regions of the continuum have been developed to characterise the spectral type of cool objects \\citep{m03,s04,a07}. We measured these for 2M1207\\,b and summarise the results in Table\\,\\ref{spectral_indices} for the \\citet{a07} H$_2$O and gravity sensitive indices, the \\citet{s04} H$_2$O and FeH indices, and the \\citet{m03} H$_2$O indices. These suggest a spectral type in the broad range M8.5--L4, while the gravity sensitive index is consistent with that found for other young brown dwarfs. The discrepancy in inferred spectral types is not surprising given the unusual overall shape of the spectrum and highlights the importance of obtaining larger wavelength coverage of flux calibrated spectra in assigning spectral types and their uncertainties. \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{14173fg3.pdf} \\caption{The $J$,$H$,$K$ spectrum of 2M1207\\,b (black line) and several comparison spectra of field L dwarfs with spectral types L1 (blue), L4.5 (green), and L8 (red). All spectra are normalised over the 1.60-1.65\\micron\\ range.} \\label{FieldComp} \\end{figure} The flux calibrated 2M1207\\,b spectrum with R$\\sim$300 at $J$ and R$\\sim$1500 at $HK$ is shown in Fig.\\,\\ref{2M1207b-spectra} along with the best fit to the model DUSTY spectrum \\citep{a01}. Available models with temperatures higher and lower by 100\\,K are also plotted in Fig.\\,\\ref{2M1207b-spectra}. The best fit, estimated by a least-squares fit to the grid of DUSTY models, and confirmed by visual examination is shown in the middle panel of Fig.\\,\\ref{2M1207b-spectra} with \\teff=1600 K and \\logg=4.5. The uncertainty in temperature is clearly less than $\\pm$100\\,K, but the ability to discriminate surface gravity variations with the broad morphology and flux level of the spectrum is less than for effective temperature. Consequently, \\logg=4.5 is only marginally favoured over \\logg=3.5 or 5.5 due to the uncertainty on the flux level of the $J$-band spectrum resulting from the photometry uncertainty of 0.2$^{\\rm{m}}$. While the location of the peak in the $H$-band is consistent with the models, the observed slopes on either side are not as steep as the model predicts, possibly due to difficulties modelling the H$_2$ collisionally-induced absorption. If the precision of the $J$-band photometry could be improved, the model variations would allow surface gravity to be constrained to $\\sim$0.5 dex, important to test one formation model. These physical parameters are consistent with those measured previously \\citep{m07}, though the inclusion of the $J$-band and the higher signal-to-noise of these data makes the fitting more secure. We did not compare the data to the COND models since they do not adequately reproduce the previous lower resolution $H$- and $K$-band spectrum \\citep{m07}. In addition to the treatment of dust condensation and settling, the inclusion of non-equilibrium CO/CH$_4$ chemistry can significantly alter the shape of the continuum emission \\citep{f08}, particularly in the $K$- and $L$-bands. A suitable parameterisation of the eddy diffusion coefficient for such models may reproduce the shape of the observed spectrum, however, the flux level for the \\teff=1400\\,K (the hottest shown), K$_{zz}$=4 non-equilibrium atmosphere shown in \\citet{f08} is still significantly higher than 2M1207\\,b; for the range of K$_{zz}$, \\logg, and metallicity presented for the \\teff=1400\\,K case, the $K$-band magnitude varies by up to 0.7\\,mag. \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{14173fg4.pdf} \\caption{The $J$,$H$,$K$ spectrum of 2M1207\\,b (bottom) and several comparison spectra of young (5--30Myr) objects in Upper Sco of spectral types (from the top): M8, L0, L2, and $\\sim$L0--L1 (AB Pic b). The spectra have been flux calibrated as in Figure 1, and then offset by an integer for clarity. } \\label{YoungComp} \\end{figure} \\begin{table} \\begin{minipage}[t]{\\columnwidth} \\caption{Spectral indices from Allers et al. (2007), Slesnick et al. (2004), and McLean et al. (2003).} \\label{spectral_indices} \\centering \\renewcommand{\\footnoterule}{} % \\begin{tabular}{llllll} % \\hline\\hline Index & H$_2$O & grav. & H$_2$O-1 & H$_2$O-2 & FeH \\\\ Value & 1.17 & 1.01 & 0.660 & 0.831 & 0.877 \\\\ SpT & L0 & low grav. & L0 & L1.5 & M8.5 \\\\ \\hline Index & H$_2$O-A & H$_2$O-B & H$_2$O-C & H$_2$O-D & \\\\ Value & 0.593 & 0.724 & 0.733 & 0.817 & \\\\ SpT & L2.5 & L3 & L0 & L4 & \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table} \\begin{figure*} \\centering \\includegraphics[width=\\textwidth]{14173fg5.pdf} \\caption{The spectrum of 2M1207\\,b covering the $J$, $H$, and $K$ passbands (black lines). {\\it Middle panel:} the best-fit DUSTY model with \\teff=1600 K and \\logg=4.5 in red along with the 1500\\,K (green line) and 1700\\,K (blue line) \\logg=4.5 models. {\\it Top and bottom panels:} DUSTY models of the same effective temperature as the middle panel, but with surface gravity of \\logg=3.5 and \\logg=5.5, respectively.\\rm The observed spectrum has been flux calibrated using the apparent magnitudes of \\citet{c04} and \\citet{m07}. While the DUSTY models reproduce the general morphology, the flux scaling necessary implies a radius of $\\sim$0.052\\,$R_{\\sun}$ at a distance of 53\\,pc, significantly smaller than any predicted radius of an ultracool dwarf. Discussion of this unphysical radius, or equivalently low luminosity is given in Sect.\\,\\ref{results}.} \\label{2M1207b-spectra} \\end{figure*} As noted previously, to match the predicted spectral model fluxes to the observed photometry, either 2M1207\\,b is underluminous or has a smaller than expected radius. The required radius to explain the photometry is $\\sim$0.052\\,$R_{\\sun}$, about one third of the radius of $\\sim$0.16\\,$R_{\\sun}$ from the DUSTY evolutionary model prediction for an object of \\teff=1600\\,K at 5--10\\,Myr, and significantly smaller than any predicted radius for a partially degenerate object above 1\\mjup\\ \\citep{c09}. The origin of 2M1207\\,b and its underluminous flux is uncertain, and different scenarios have been proposed for the specific case of 2M1207\\,b: an edge-on disk \\citep{m07} and the result of a collision of protoplanets \\citep{mam07}. More generally, the impact on the luminosity due to initial conditions and the early accretion history of young objects with ages comparable to 2M1207\\,b, has been investigated \\citep{mar07, b09}. The protoplanet collision hypothesis predicts that the object is much smaller and should have a surface gravity of 3.0 rather than 4.5. The upper panel of Fig.\\,\\ref{2M1207b-spectra} shows the data with DUSTY spectral models with three different effective temperatures for a lower surface gravity of \\logg=3.5 (since \\logg=3.0 is not available). The data favour the higher surface gravity \\logg=4.5 value, though increasing the $J$-band flux by one sigma would make the data consistent with the \\logg=3.5 model. A $J$-band magnitude difference with smaller uncertainties would be required to entirely rule out the collision scenario. For an initial investigation into the edge-on disk possibility, we compared the photometry and spectra with synthetic spectra generated by the TORUS and MCFOST codes \\citep{p09}. A central object simulated by the \\teff=1600 K DUSTY atmosphere was encircled with a disk 10 AU in radius, consistent with the physical separation of the 2M1207 pair. The disk was constructed with canonical values for flaring, surface density power law, dust size distribution and chemistry. The inclination was allowed to vary from edge-on to the $\\sim$60 degree limit suggested for the geometry of the primary disk. Within the limited range of disk parameters searched, scattering alone could not explain the flux level of the central object since the predicted scattered flux was only a few per-cent, rather than the required $\\sim$10\\%. For comparison, the T Tauri stars HH 30 and HK Tau with imaged edge-on disks \\citep{s98} are fainter by 3--4 magnitudes relative to the photospheric levels, consistent with a scattered flux level of a few per-cent or less. Thus, among disk scenarios, grey extinction from a disk composed of larger grains than considered in the simulations is favoured over scattering off the surface of an optically thick disk; grain sizes larger than 4\\micron\\ have been suggested as a possibility \\citep{m07}. Additionally, extinction of a hotter object is not able to explain the wavelength dependence of the photometry over the entire 0.9--3.6\\micron\\ range. In summary, the SINFONI spectrum of 2M1207\\,b represents the highest resolution, largest wavelength coverage spectrum of a young, planetary mass companion. The overall shape is distinct from field and young cluster L-dwarfs, with substantially different ratios of flux across the bands, a triangular shape to the $H$-band continuum, and a peak at longer wavelengths in the $K$-band. Spectral indices developed for cool objects give conflicting values, indicating the unusual nature of this object. The spectrum is best fit by a DUSTY model atmosphere with \\teff=1600\\,K and \\logg= 4.5, though it is not yet possible to exclude the \\logg= 3.0 atmosphere predicted by one formation model. We find that the apparent underluminosity of 2M1207\\,b is consistent with extinction from large dust grains in a nearly edge-on disk, but more precise $J$-band and longer wavelength photometry and spectroscopy will help to determine the nature of this important object. This spectrum of 2M1207\\,b will serve as a key comparison for upcoming near-infrared spectroscopy of planets imaged around stellar hosts." }, "1003/1003.3327_arXiv.txt": { "abstract": "We study the self-consistent, linear response of a galactic disc to non-axisymmetric perturbations in the vertical direction as due to a tidal encounter, and show that the density distribution near the disc mid-plane has a strong impact on the radius beyond which distortions like warps develop. The self-gravity of the disc resists distortion in the inner parts. Applying this approach to a galactic disc with an exponential vertical profile, Saha \\& Jog showed that warps develop beyond 4-6 disc scalelengths, which could hence be only seen in HI. The real galactic discs, however, have less steep vertical density distributions that lie between a $sech$ and an exponential profile. Here we calculate the disc response for such a general $sech^{2/n}$ density distribution, and show that the warps develop from a smaller radius of 2-4 disc scalelengths. This naturally explains why most galaxies show stellar warps that start within the optical radius. Thus a qualitatively different picture of ubiquitous optical warps emerges for the observed less steep density profiles. The surprisingly strong dependence on the density profile is due to the fact that the disc self-gravity depends crucially on its mass distribution close to the mid-plane. General results for the radius of onset of warps, obtained as a function of the disc scalelength and the vertical scaleheight, are presented as contour plots which can be applied to any galaxy. ", "introduction": "It is a common knowledge now that the galaxies are not isolated structures in the universe, rather galaxy interactions including mergers are common. Hence understanding the dynamics of these interactions and the effects they produce on the structure and evolution within galaxies has been of much interest. One such effect is the generation of asymmetric features due to tidal encounters between galaxies. Spiral galaxies are known to display a variety of non-axisymmetric features, both in the plane and also in the direction perpendicular to the plane. The most common vertical distortion is a warp, a feature of azimuthal wavenumber $m=1$. Most nearby galaxies show such integral-sign or s-shaped warps in their outer parts (e.g., Binney \\& Tremaine 1987). A similar planar distortion commonly seen is the lopsidedness in disc corresponding to azimuthal wavenumber $m=1$ (Jog \\& Combes 2009). The origin of warp is not yet clear despite a long search. A commonly suggested mechanism for the origin of warps is due to the tidal interaction with its neighbours (e.g., Schwartz 1985, Zaritsky \\& Rix 1997). Weinberg (1995) studied the generation of warp in our Galaxy due to perturbation from the neighbouring Large Magellanic Cloud. Vertical distortions other than warps are also commonly observed in spiral galaxies. Small-scale corrugations in external galaxies have been studied (Quiroga 1984) and so have been the distortions in the stellar distributions of galaxies (Florido et. al 1991). Scalloping in HI in the outer regions of our Galaxy has been studied too (Kulkarni, Blitz \\& Heiles 1982). Recently Matthews \\& Uson (2008) find that corrugation in HI is seen in IC2233 even in the inner regions. While warps are seen mostly in the outer parts of a galactic disc, surprisingly little work has been done to discuss the radius at which warps develop. Saha \\& Jog (2006) proposed this as being determined due to the self-consistent disc response to a tidal field. The disc self-gravity resists distortion in the inner parts, and only in the outer parts does the disc begin to respond to the external potential. In this paper, we continue with this approach and study self-consistent vertical distortions for different forms of vertical density distributions in the disc. We are motivated by the fact that in the previous studies, not much attention has been paid to the effect the form of density distribution for the responding disc might have on the overall behaviour of the system. For the purpose of modeling, the vertical distribution was taken to be exponential for mathematical simplicity in an earlier work (Saha \\& Jog 2006). The study of vertical distribution of stars in galactic discs has an interesting history. The vertical distribution was at first deduced to be of type $sech^2$ as resulting for an isothermal disc (Spitzer 1942). However, observations showed a steeper profile closer to a $sech$ or an exponential both for our Galaxy (e.g., Gilmore \\& Reid 1983, Kent et al. 1991) as well as for external galaxies (e.g., Wainscoat, Freeman \\& Hyland 1989, Rice et al. 1996). An exponential profile all the way to the mid-plane is unphysical and hence van der Kruit (1988) proposed a generalized function $sech^{2/n}$ to represent the vertical density distribution. In this scheme, $n=1$ and 2 correspond to a $sech^2$ and a $sech$ distribution and as $n$ tends to $\\infty$ it asymptotically approaches an exponential distribution. Later observers analyzed their data and cast in this format, and have shown that a true density distribution is less steep than an exponential and in most case lies between a $sech$ and an exponential distribution (Barteldrees \\& Dettmar 1994, de Grijs, Peletier \\& van der Kruit 1997). While dust extinction prevents a determination of the stellar density profile close to the mid-plane, the near-infrared bands do not have this limitation and represent the true density profiles representing the old stars. For a sample of 24 galaxies studied in the K$_s$ band, de Grijs et al. (1997) find that, a mean value of $<2/n> = 0.5$ corresponding to the $n$ index = 4. Thus the vertical density profile for stars in a typical galactic disc lies between a $sech$ and an exponential profile, being closer to a $sech$. In this paper, we study the vertical response of an axisymmetric disc to an external imposed perturbation, where the disc density follows such a general $sech^{2/n}$ distribution. We also study the response of an exponential disc for the sake of comparison. The motivation for our work comes from the fact that the matter distribution close to the mid-plane contributes strongly to the vertical self-gravitational force (Banerjee \\& Jog 2007), hence it is plausible that the different vertical density profiles affect the disc response in different ways. We show that indeed the disc response has a strong dependence on the vertical density distribution. This in turn significantly affects the radius for the onset of various non-axisymmetric features along the vertical direction. Section 2 of the paper presents the details of the model and the methods of calculation, while section 3 presents the results. Section 4 presents the discussion and section 5 concludes the paper. ", "conclusions": "\\noindent {\\it 1. Stellar warps and their detection:} Observations show that stellar warps are common and occur in more than 50\\% of spiral galaxies(Sanchez-Saavedra, Battaner \\& Florido 1990, Reshetnikov \\& Combes 1998). These therefore must start within the Holmberg radius. Indeed this distinction though obvious is not often made in the literature- namely, a stellar warp by necessity must start within the optical radius. Recent systematic study of 325 edge-on galaxies Ann \\& Park (2006) confirms this point, with a typical warp radius = 0.7 times the optical radius or $\\sim 3$ disc scalelengths. Interestingly, this observed value agrees well with our typical onset radius of $\\sim 3$ disc scalelengths (Table 1). Thus, we have shown that a realistic, less-steep vertical density distribution of type $sech^{2/n}$ results in the onset of stellar warps within the optical radius as is observed in most galaxies. A similar value is seen for our Galaxy, where the stellar warp is shown to start from 3.1 disk scalelengths based on the COBE/DIRBE data (Drimmel \\& Spergel 2001), and 2.4 disk scalelengths based on the 2MASS data (Lopez-Corredoira et al. 2002). The {\\it Spitzer} observations of ten galaxies show the onset of warps lie within 3-6 disc scalelengths, thus many start from within the optical radius (Saha, de Jong \\& Holwerda 2009). These authors treat an exponential vertical density profile for simplicity, and try to explain the small observed warp onset radii by assuming that the scaleheight increases with radius - while keeping the disc mass constant. Such flaring with radius is observed in some galaxies (de Grijs \\& Peletier 1997) and is expected for a multi-component, coupled, star-gas disc (Narayan \\& Jog 2002). However, the values of flaring they use are ad-hoc, and even this cannot explain the entire range of smaller values of R$_{min}$ that are observed for their sample. Further, they use $z_0/ R_D$ as a single thickness parameter but as we have shown (Fig. 4), the value of R$_{min}$ is not a simple function of this parameter. Instead the small onset radii can be explained naturally as we have done, by using the observed less-steep vertical density distribution. \\medskip \\noindent {\\it 2. Warp onset radius: Dependence on disc mass} At high redshifts, the galaxy size is smaller as seen for the Hubble deep field sample (Elmegreen et al. 2005) with a typical disc scalelength of 1.5 kpc. Such a size variation with redshift is expected in the hierarchical evolution scenario (e.g., Steinmetz \\& Navarro 2002). The warp onset radius for this sample is observed to be smaller $\\sim$ 1.4 disc scalelengths (Reshetnikov et al. 2002). This observed trend agrees exactly with our result (see Fig. 4)- namely that the smaller mass galaxies allow warps to develop from a smaller starting radius. At the opposite end, we predict that massive, nearby disc galaxies would have warp onset at a larger radius and thus are less likely to show a stellar warp. This would be tricky to confirm because the radial range over which optical warps are seen is small (starting at 3 R$_D$ and going up to 4-5 R$_D$), and observational data need to be analyzed to study this point. There is some evidence for this correlation: M31 has a massive disc with a scalelength of 5.4 kpc (Geehan et al. 2006), and it has a stellar warp starting at radii larger than the isophote at 26.8 ${\\mu}_B$ (Innanen et al. 1982). This is beyond the Holmberg radius, whereas the average onset radius of stellar warps is within this radius, see point 1 above. \\medskip \\noindent {\\it 3. Effect of nearby perturbers, and live halo :} A real galaxy is likely to undergo close encounters with satellites less massive than the LMC, while our calculation is meant for a distant encounter with distance large compared to the galaxy size. Further, for simplicity, we have taken the halo to be rigid and the forcing frequency to be zero. The long-standing problem about the tidal origin of warps has been the resulting small amplitude (Hunter \\& Toomre 1968). This is overcome if the halo is live and the wake generated at the resonance points of the frequency of the perturber's motion is included (Weinberg 1998, Tsuchiya 2002, Weinberg \\& Blitz 2006). The Sagittarius dwarf is about 10 times less massive than the LMC but is three times closer, so their direct tidal torques on the Galaxy (being proportional to the mass of the perturber/distance$^3$ to the perturber) are comparable. Hence the Sagittarius dwarf also cannot directly produce the Galactic warp that is observed. However, as Bailin (2003) has argued, the magnitude of the tidal field due to the wake generated in the halo by the Sagittarius dwarf could also explain magnitude of the warp seen in the Galaxy. This needs to be checked by simulations. Yet another pathway to create vertical perturbations would be to have an even smaller mass perturber, like the subhalo, come even closer to $\\sim 5-10$ kpc. This possibility has been studied by Chakrabarti \\& Blitz (2009) via simulations, who show that this can explain the HI amplitudes of warps and the higher order vertical modes as observed by Levine et al. (2006). In the present paper we have not attempted to obtain the actual warp amplitude. However, the concept of negative disc response studied here would still apply in these general cases. That is, due to its strong self-gravity, the inner disc region would resist being distorted by vertical perturbations. \\medskip \\noindent {\\it 4. Dynamical implications: } The present paper shows the surprisingly strong dependence of the resulting warp radius on the vertical density distribution in the disc. This is because the disc self-gravity which decides the warp radius depends crucially on the vertical disc distribution close to the mid-plane. A similar strong dependence on the disc vertical distribution may also affect other dynamical studies such as the vertical heating due to tidal encounters, or the bending instabilities. We plan to look at these in future papers." }, "1003/1003.4993_arXiv.txt": { "abstract": "Numerical-relativity simulations indicate that the black hole produced in a binary merger can recoil with a velocity up to $v_{\\rm max} \\simeq 4,000$ km/s with respect to the center of mass of the initial binary. This challenges the paradigm that most galaxies form through hierarchical mergers, yet retain supermassive black holes at their centers despite having escape velocities much less than $v_{\\rm max}$. Interaction with a circumbinary disk can align the binary black hole spins with their orbital angular momentum, reducing the recoil velocity of the final black hole produced in the subsequent merger. However, the effectiveness of this alignment depends on highly uncertain accretion flows near the binary black holes. In this {\\it Letter}, we show that if the spin $\\textbf{S}_1$ of the more massive binary black hole is even partially aligned with the orbital angular momentum $\\textbf{L}$, relativistic spin precession on sub-parsec scales can align the binary black hole spins with each other. This alignment significantly reduces the recoil velocity even in the absence of gas. For example, if the angle between $\\textbf{S}_1$ and $\\textbf{L}$ at large separations is 10 degrees while the second spin $\\textbf{S}_2$ is isotropically distributed, the spin alignment discussed in this paper reduces the median recoil from 864 km/s to 273 km/s for maximally spinning black holes with a mass ratio of 9/11. This reduction will greatly increase the fraction of galaxies retaining their supermassive black holes. ", "introduction": "\\label{S:intro} Observations suggest that most galaxies host supermassive black holes (SBHs) at their centers whose masses are tightly correlated with properties of their host spheroids \\citep{Magorrian:1997hw,Ferrarese:2000se,Tremaine:2002js}. If galaxies form through hierarchical mergers, their SBHs may form from the merger of the smaller SBHs in their progenitor galaxies. The final stage of these black-hole mergers involves highly curved, dynamical spacetime that can only be simulated with fully nonlinear numerical relativity (NR). Following a major breakthrough in 2005 \\citep{Pretorius:2005gq,Campanelli:2005dd,Baker:2005vv}, numerical relativists can now accurately determine the anisotropic emission of gravitational waves during the final stage of black-hole mergers. When gravitational waves are preferentially emitted in one direction during a merger, conservation of linear momentum requires that the final black hole produced in that merger recoil in the opposite direction. These recoil velocities or ``kicks'' can approach 4,000 km/s for maximally spinning mergers \\citep{Campanelli:2007ew,Gonzalez:2007hi}. Kicks this large exceed the escape velocities of even the most massive galaxies, and would thus eject SBHs from their hosts \\citep{Merritt:2004xa}. Frequent SBH ejections would seem to contradict the tightness of the observed correlations between SBHs and their host galaxies. Kicks would pose an even greater problem at high redshifts, where typical galactic escape velocities decrease while recoils remain a fixed fraction of the speed of light. How might we avoid black-hole mergers that lead to large kicks? To answer this question, we must take a closer look at how the predicted recoils depend on the dimensionless spins $\\xI \\equiv \\Si/m_{i}^2$ and mass ratio $q \\equiv m_2/m_1 \\leq 1$ of the merging black holes. Reliable NR simulations have been performed for $q \\geq 0.1$ and $|\\xI| \\leq 0.9$; in this range the recoils are well described by the fitting formula \\citep{Campanelli:2007ew} \\beq \\label{E:ktot} \\vec{v}(q,\\xa,\\xb) = v_m \\mathbf{\\hat{e}}_1 + v_\\perp (\\cos \\xi~\\mathbf{\\hat{e}}_1 + \\sin \\xi~\\mathbf{\\hat{e}}_2) + v_\\parallel \\ez \\eeq where \\begin{eqnarray} \\label{E:vm} v_m &=& A\\eta^2 \\frac{1-q}{\\qsum}(1 + B\\eta)~, \\\\ \\label{E:vper} v_\\perp &=& H\\eta^2 \\VD^\\parallel \\cdot \\ez~, \\\\ \\label{E:vpar} v_\\parallel &=& K\\eta^2 \\cos(\\Theta - \\Theta_0) |\\VD^\\perp|~. \\end{eqnarray} Here ($\\mathbf{\\hat{e}}_1, \\mathbf{\\hat{e}}_2, \\ez$) are an orthonormal basis with $\\ez$ parallel to the orbital angular momentum $\\LN$, $\\eta \\equiv q/(\\qsum)^2 \\leq 1/4$ is the symmetric mass ratio, and $\\VD^{\\parallel,\\perp}$ are the components of \\beq \\label{E:Delta} \\VD \\equiv \\frac{q\\xb - \\xa}{\\qsum} \\eeq parallel and perpendicular to $\\LN$. $\\Theta$ is the angle between $\\VD^\\perp$ and the separation $\\mathbf{r}$ of the two black holes ``at merger''. NR simulations indicate that the best-fit values for the coefficients appearing in the above formula are $A = 1.2 \\times 10^4$ km/s, $B = -0.93$ \\citep{Gonzalez:2006md}, $H = (6.9 \\pm 0.5) \\times 10^3$ km/s \\citep{Lousto:2007db}, and $K = (6.0 \\pm 0.1) \\times 10^4$ km/s \\citep{Campanelli:2007cga}. The angle $\\xi \\sim 145^\\circ$ for a wide range of quasi-circular configurations \\citep{Lousto:2007db}, while $\\Theta_0$ depends on the mass ratio $q$ but not the spins \\citep{Lousto:2008dn}. The large value of $K$ implies that the merger of equal-mass black holes with maximal spins pointed in opposite directions in the orbital plane generates a recoil of $K/16 = 3,750$ km/s. The most obvious way to avoid these large kicks is to require black holes to be non-spinning: \\beq \\xI = 0 \\to \\VD = 0 \\to v_\\perp = v_\\parallel = 0~. \\eeq In this case, $v_m$ is maximized at the modest value $175 \\pm 11$ km/s for a mass ratio $q \\simeq 0.36$ \\citep{Gonzalez:2006md}. However, theory shows that non-spinning black holes can be spun up to the Kerr limit $|\\boldsymbol{\\chi}| = 1$ by steady accretion after increasing their mass by only a factor of $\\sqrt{6}$ \\citep{Bardeen:1970}. Observations of Fe K$\\alpha$ fluorescence can be used to measure black-hole spins \\citep{Reynolds:1998ie}, and indicate that real SBHs can approach this limit: for the SBH in the Seyfert galaxy MCG-06-30-15 the measured spin is $|\\boldsymbol{\\chi}| = 0.989_{-0.002}^{+0.009}$ at 90\\% confidence \\citep{Brenneman:2006hw,Berti:2009}. If black holes are highly spinning, the recoil can be reduced by aligning the spins and thus $\\VD$ with the orbital angular momentum: \\beq \\xI \\parallel \\LN \\to \\VD^\\perp = 0 \\to v_\\parallel = 0~. \\eeq This spin configuration leads to smaller kicks, since the coefficient $H$ is almost an order of magnitude less than $K$. The merger of equal-mass, maximally spinning black holes with one spin aligned with $\\LN$ and the other anti-aligned generates a recoil of $H/16 = 431$ km/s. Gaseous accretion disks are needed to provide dynamical friction to allow SBHs separated by $r \\simeq 1$ pc to merge in less than a Hubble time \\citep{Begelman:1980vb}. These same accretion disks can exert torques on the SBHs which align their spins and orbital angular momentum with that of the disk, thus producing the desired aligned spin configuration \\citep{Bogdanovic:2007hp}. However, the effectiveness of this alignment mechanism depends on the highly uncertain nature of the accretion flow near the merging black holes. \\cite{Dotti:2009vz} find a residual misalignment of $10^\\circ$ ($30^\\circ$) between the black-hole spins and accretion disk depending on whether the disk is cold (hot). This misalignment could be even greater in a gas-poor merger or one in which accretion onto the SBHs proceeds through a series of small-scale, randomly oriented events \\citep{King:2007nx,Berti:2008}. In this {\\it Letter}, we present a new mechanism to reduce gravitational recoils by aligning black-hole spins {\\it with each other} prior to merger. \\cite{Boyle:2007sz,Boyle:2007ru} showed that the symmetries of binary black-hole systems imply that recoils are only generated by a weighted difference of the two spins. This general result can be seen to hold for the fitting formula of \\cite{Campanelli:2007ew} by noting that a weighted difference of spins appears in the numerator of $\\VD$ in equation (\\ref{E:Delta}). Spin alignment is a consequence of relativistic spin precession as the black holes inspiral due to the loss of energy and angular momentum to gravitational radiation. We begin calculating the inspiral at an initial separation $r_i = 500 R_S$ where spin alignment begins for comparable-mass binaries \\citep{Schnittman:2004vq}, and end at a final separation $r_f = 5 R_S$ near where NR simulations typically begin. Here $R_S = 2GM/c^2$ is the Schwarzschild radius of a non-spinning black hole of mass $M$. Relativists use units in which $G = c = 1$, allowing them to measure distance and time in units of $M$, where $M \\equiv m_1 + m_2$ is the sum of the masses of the merging black holes. We shall do this for the rest of the paper. The spin alignment discussed in this {\\it Letter} occurs for both gas-rich and gas-poor mergers, as gravitational radiation (GR) dominates the dynamics even in the presence of gas at binary separations less than \\beq \\label{E:rGR} r_{\\rm GR} \\sim 3000 M q^{1/4} \\left( \\frac{\\dot{M}}{1 M_\\odot~{\\rm yr}^{-1}} \\right)^{-1/4}~, \\eeq where $\\dot{M}$ is the rate of gas infall \\citep{Begelman:1980vb}. We briefly describe the relativistic dynamics leading to spin alignment in \\S~\\ref{S:align}; readers interested in further details can find them in our longer paper on how spin alignment affects the distributions of black-hole final spins \\citep{Kesden:2010yp}. The most notable effect of spin alignment is to suppress the recoil velocity when the spin of the larger black hole is initially partially aligned with $\\LN$. The magnitude of this suppression for distributions with different mass ratios and initial spins is presented in \\S~\\ref{S:kick}. Some concluding remarks are provided in \\S~\\ref{S:disc}. ", "conclusions": "\\label{S:disc} In this {\\it Letter}, we have shown that as black holes inspiral, spin precession aligns their spins with each other for the spin distributions expected in astrophysical mergers. This spin alignment drastically reduces the recoils expected for the black holes produced in the binary mergers. Spin alignment is most effective for the highly spinning, comparable-mass mergers that are predicted to yield the largest recoils (up to $v_{\\rm max} = 3,750$ km/s according to the fitting formula of \\cite{Campanelli:2007ew}). Aligning the black-hole spins generically suppresses the recoils \\citep{Boyle:2007sz,Boyle:2007ru}; we found a similar suppression with the alternative fitting formula of \\cite{Baker:2008md}. This spin alignment is a purely relativistic effect that will occur for all black-hole mergers, as gravitational radiation will always dominate the inspiral for separations $r < r_{\\rm GR} \\sim 3000 M$. As long as torques at $r > r_{\\rm GR}$ align $\\Sa$ and $\\LN$ such that $\\theta_1 \\leq 30^\\circ$, spin alignment during the final inspiral will nearly eliminate the $v \\gtrsim 1,000$ km/s recoils that are so difficult to reconcile with galaxies keeping their supermassive black holes. While there is still great uncertainty about how merging black holes interact with surrounding gas, the PN spin precession discussed in this paper is inevitable and results from well established physics. We therefore believe that spin alignment must be accounted for in future population studies of merging black holes." }, "1003/1003.3057_arXiv.txt": { "abstract": "On UT 2009 January 16, we observed a white light megaflare on the dM4.5e star YZ CMi as part of a long-term spectroscopic flare-monitoring campaign to constrain the spectral shape of optical flare continuum emission. Simultaneous $U$-band photometric and 3350-9260\\AA$ $ spectroscopic observations were obtained during 1.3 hours of the flare decay. The event persisted for more than 7 hours and at flare peak, the $U$-band flux was almost 6 magnitudes brighter than in the quiescent state. The properties of this flare mark it as one of the most energetic and longest-lasting white light flares ever to be observed on an isolated low-mass star. We present the $U$-band flare energetics and a flare continuum analysis. For the first time, we show convincingly with spectra that the shape of the blue continuum from 3350\\AA$ $ to 4800\\AA$ $ can be represented as a sum of two components: a Balmer continuum as predicted by the Allred et al radiative hydrodynamic flare models and a $T\\sim$10,000K blackbody emission component as suggested by many previous studies of the broadband colors and spectral distributions of flares. The areal coverage of the Balmer continuum and blackbody emission regions vary during the flare decay, with the Balmer continuum emitting region always being significantly ($\\sim$3-16 times) larger. These data will provide critical constraints for understanding the physics underlying the mysterious blue continuum radiation in stellar flares. ", "introduction": "Due to their strong and persistent surface magnetic fields, active M dwarfs can produce frequent flares, sometimes lasting for many hours and reaching luminosities that approach a significant fraction of the star's bolometric luminosity \\citep{Kunkel1969, Moffett1974, Bond1976, Hawley1991, Favata2000, Robinson2005, Osten2007, Kowalski2009}. During both solar and stellar flares, emission is seen as line and continuum radiation from X-ray to radio wavelengths, with a prominent mode of radiative energy release occurring in the blue and near-ultraviolet (NUV) continuum, which is commonly referred to as the white light continuum. Unfortunately, the spectral components which comprise the white light continuum are not well-constrained, and therefore the physics of the underlying flare mechanism that gives rise to such a conspicuous signature remains a mystery. Significant effort has gone into characterizing the white light continuum in stellar flares, beginning with a single component (hydrogen recombination) model \\citep{Kunkel1969, Kunkel1970}. A two-component spectral model consisting of hydrogen recombination and an impulsively heated photosphere was first proposed by \\cite{Kunkel1970} as a possible explanation of the spread in broadband flare colors. Time-resolved multi-channel photometry \\citep{Mochnaki1980} and spectrometry \\citep{Kahler1982} were obtained longward of 4200\\AA$ $ for flares on YZ CMi and UV Ceti; at peak, these flare spectra were fit with a Planck function of $T=7400-9500$K, thereby showing evidence of a hot optically thick component present during the impulsive phases of flares. Instrumental effects shortward of the Balmer jump (3646\\AA) prevented a precise characterization of the Balmer continuum in the data of \\cite{Mochnaki1980}. Higher resolution spectra in the blue/NUV were obtained during large flares on AD Leo \\citep[3560\\AA - 4440\\AA;][]{Hawley1991} and CN Leo \\citep[3050\\AA - 3860\\AA;][]{Fuhrmeister2008} showing a nearly continuous rise into the NUV with a $T\\sim8400-11,300$K blackbody; surprisingly no discontinuity at the Balmer jump was seen in either of these events \\citep[see also ][]{Eason1992, Garcia2002}. The white light continuum has therefore been attributed to a single dominant blackbody component with $T\\sim9000-10,000$K \\citep[e.g.,][]{Hawley2003}. In contrast, flare models predict spectra that are dominated by hydrogen bound-free continua, which provide a poor match to the observations \\citep[ hereafter A06]{Hawley1992, Allred2006}. A multi-component model was reintroduced by \\cite{Zhilyaev2007}, who used high-cadence $UBVRI$ photometric observations of the dM3.5e star EV Lac to suggest that the white light emission consists primarily of blackbody radiation at flare peak and hydrogen continuum during the flare decay. However, significant ambiguity results when using broadband photometry to characterize the white light continuum. A06 showed that a flare model spectrum that included a large Balmer jump, when convolved with broadband $UV+UBVR$ filters, actually resembles the shape of a $T\\sim$9,000K blackbody. To unambiguously identify the spectral components present during stellar flares, we have therefore begun to compile a large catalog of time-resolved optical/NUV flare spectra on several nearby flare stars. In this Letter, we present time-resolved spectra obtained during a large flare on the dM4.5e star YZ CMi and use results from the recent radiative hydrodynamic (RHD) flare models of A06 to show that both Balmer continuum and hot blackbody components are present in the white light continuum emission. An extensive paper describing our flare catalog, including discussion of the flare emission lines and comparison with models will be forthcoming. ", "conclusions": "We observed a white light megaflare on the dM4.5e star YZ CMi in the $U$ band and with simultaneous optical/NUV spectroscopy. The $U$-band energetics and light curve morphology qualify this flare as an extraordinary and rare event, similar to the $\\sim6.6 \\times 10^{34}$ erg flare on YZ CMi seen by \\cite{Andrews1969}, \\cite{Lovell1969}, and \\cite{Kunkel1969}. Following a solar analogy, we speculate that the flaring region on YZ CMi was a complex arcade of sequentially reconnecting magnetic loops. Each reconnecting loop accelerated a beam of nonthermal electrons that impacted the lower atmosphere, producing the observed blue/NUV line and continuum emission. The sum of a large number of individual emitting regions may have enabled this flare to persist for such an unusually long time. Using high time resolution, high signal-to-noise spectra, we have shown that the blue/NUV flare continuum radiation can be explained as a sum of a $T\\sim$10,000K blackbody component and a Balmer continuum component, with only the blackbody emission present from 4000-4800\\AA. The relative filling factors of the two components indicate that the Balmer continuum comes from a larger region, plausibly originating from the flaring loops at chromospheric heights, where the Balmer lines originate. This is consistent with the height of the Balmer continuum emission derived from solar flare data \\citep{Hudson2010}, and also with the height of formation predicted by the RHD model of A06. We found that the blackbody emission arises from a region $\\sim$3-16 times smaller in area than the Balmer continuum emission region. This blackbody emission may possibly originate in concentrated magnetic footpoint regions in the lower atmosphere, similar to the localized areas that emit in white light during solar flares \\citep{Metcalf2003, Fletcher2007, Isobe2007, Jess2008}, which are often spatially and temporally coincident with impulsive heating by nonthermal electrons inferred from hard X-ray observations \\citep[e.g.,][]{Rust1975, Hudson1992, Neidig1993, Fletcher2007}. In accordance with this scenario, a large complex of photospheric hot spots may have been created during the first impulsive events in the YZ CMi $U$-band light curve. As the nonthermal electron beams weakened (in energy and/or in spatial extent) so did the areal coverage of these hot spots. This gradual decay extends into the time covered by our spectral observations and is seen in the overall decline of the blackbody areal coverage in the upper panel of Figure \\ref{fig:figs}d. During our spectral observations, a few new hot spots may have been created at the footpoints of newly reconnected loops, causing the transient increases seen in the $U$-band and in the effective area of the blackbody emitting region. We see strong evidence through all of our observations for a blackbody continuum emission component with an approximate temperature of $10,000$K. The persistence of the hot blackbody emission indicates that a continual source of particle acceleration and plasma heating likely still exists during the decay phase of the flare. Yet, the A06 models, which employ a nonthermal electron beam as is seen on the Sun, predict the photosphere of an M dwarf to be heated by at most $\\sim$1200K during a large flare. The physical mechanism which generates the strong blackbody emission (presumably from a more strongly heated photosphere) therefore remains unknown. The possible anti-correlation between the Balmer continuum and blackbody emission components may provide an important clue to the nature of the heating mechanism that is responsible for the blackbody emission." }, "1003/1003.5784_arXiv.txt": { "abstract": "Aims: Our aim is to investigate the resistive relaxation of a magnetic loop that contains braided magnetic flux but no net current or helicity. The loop is subject to line-tied boundary conditions. We investigate the dynamical processes that occur during this relaxation, in particular the magnetic reconnection that occurs, and discuss the nature of the final equilibrium. \\\\ Methods: The three-dimensional evolution of a braided magnetic field is followed in a series of resistive MHD simulations. \\\\ Results: It is found that, following an instability within the loop, a myriad of thin current layers forms, via a cascade-like process. This cascade becomes more developed and continues for a longer period of time for higher magnetic Reynolds number. During the cascade, magnetic flux is reconnected multiple times, with the level of this `multiple reconnection' positively correlated with the magnetic Reynolds number. Eventually the system evolves into a state with no more small-scale current layers. This final state is found to approximate a non-linear force-free field consisting of two flux tubes of oppositely-signed twist embedded in a uniform background field. ", "introduction": "The braiding of magnetic loops in the solar corona, via convective motions at the solar surface, has long been suggested as a potential mechanism for heating the corona. \\cite{parker1972} proposed that in response to arbitrary footpoint motions at the photosphere, the coronal magnetic field will ideally relax towards a force-free equilibrium containing tangential discontinuities, corresponding to current sheets. Therefore, it was argued, as such a singular state is approached during the relaxation, the diffusivity of the plasma must always become important, leading to magnetic reconnection and plasma heating. Since this idea was first proposed, there have been many arguments for and against its validity \\citep[e.g.][]{vanballegooijen1985,longcope1994,ng1998}. In Parker's original `topological dissipation' model, an initially uniform magnetic field is taken between two perfectly conducting parallel plates. Then random (smooth) motions are applied on the perfectly conducting plates. A number of attempts have been made to simulate such a scenario numerically, using a variety of numerical approaches. \\cite{craigsneyd2005} employed an ideal relaxation scheme which includes a fictitious frictional term in the equation of motion \\citep{craig1986}. They applied various complex deformations at the boundaries of the domain. They found that regardless of the nature or extent of the deformation, no tangential discontinuities developed in the relaxed state, with the current concentrations remaining large-scale in all cases. \\cite{mikic1989} employed a slightly different approach. They began with a uniform field and then sequentially applied large-scale shear flows of random orientation on one boundary (while the field at the other boundary was held fixed). In the time between each boundary shear, they solved the ideal MHD equations { (neglecting the pressure and advective terms in the equation of motion)}, including a large spatially-uniform viscosity in order to relax the magnetic field towards equilibrium. They found, once again, that no discontinuities were formed during any of these relaxation processes. However, progressively smaller scale current structures were found to develop as the number of shear disturbances increased. The implication is therefore that after a sufficient length of time, the scales would reach those appropriate for dissipation in the corona. They found, furthermore, that the current density in the domain increased exponentially in time, which is consistent with the earlier analysis of \\cite{vanballegooijen1985}, who predicted thin non-singular current layers rather than tangential discontinuities. {\\cite{galsgaard1996} employed a further different approach to those discussed above, solving the full set of resistive MHD equations without employing any artificial force (such as enhanced viscosity) to inhibit the plasma dynamics}. They applied a similar sequence of large-scale shearing motions as \\cite{mikic1989}, this time at both driving boundaries. The amplitude and orientation of the shearing were chosen at random (from a normal distribution). The results were similar to those of \\cite{mikic1989}, in that after only a few Alfv{\\' e}n crossing times, multiple small-scale current filaments were found to form in the domain. However, with the shearing applied on both boundaries, exponentially growing currents were already obtained {after the second shear motion due to the interlocking of the field lines, with the tension force setting up a stagnation flow}. Due to the finite resistivity in their simulations, magnetic reconnection and Joule dissipation occurred in these filaments. This energy release was found to be intermittent or `bursty' when the time scale for energy input (via the boundary shearing) was long compared with the Alfv{\\' e}n crossing time. The work was extended by \\cite{galsgaard2002} with the inclusion of a stratified atmosphere. {Furthermore, the generation and properties of turbulence in Parker's model have been studied in the framework of reduced MHD by a number of authors. In the initial studies, two-dimensional simulations were performed, with imposed `forcing terms' taking the place of the boundary driving \\citep[e.g.][]{einaudi1996,dmitruk1997,georgoulis1998}. More recently, three-dimensional simulations of Parker's model with the reduced MHD equations have been performed \\citep[e.g.][]{dmitruk2003,rappazzo2008}, and the resulting energy spectra and heating event statistics investigated.} In the present work, we approach the topological dissipation problem discussed above from a different angle to previous studies. {We take a braided magnetic field that is close to force-free as an initial condition, and concentrate on the details of the subsequent evolution, including current sheet formation. In general, braided force-free fields could arise in the solar corona either through the emergence of braided flux from the interior or through random footpoint motions at the solar surface, accompanied by some intermediate relaxation processes. The ideal relaxation of this field towards a force-free state was considered in \\citet{wilmotsmith2009a} where only large-scale current features were found in the end state. } We now follow the evolution of the system in a resistive MHD simulation. In a previous paper (\\cite{wilmotsmith2010}, hereafter referred to as Paper I) we described how during the early evolution an instability occurs which moves the system away from equilibrium, leading to the formation of current sheets and thus the onset of magnetic reconnection. In the present paper we address the following key questions: \\begin{enumerate} \\item As the plasma seeks a new equilibrium, what is the nature of the resistive relaxation process? \\item In particular what are the properties of the magnetic reconnection processes that take place? \\item What is the nature of the final state of the relaxation? \\item What is the dependence of each of the above on the magnetic Reynolds number? \\end{enumerate} In Section \\ref{numsec} we introduce the numerical method and summarise the results of our previous investigations. In Section \\ref{evsec} we describe qualitatively the evolution of the system, and in Section \\ref{finsec} we discuss the nature of the final state. In Section \\ref{recsec} we investigate the reconnection in the system in more detail, and then in Section \\ref{etasec} discuss the dependence on the plasma resistivity. Finally in Section \\ref{concsec} we present our conclusions. ", "conclusions": "\\label{concsec} We have described a series of numerical experiments in which we considered the resistive MHD evolution of a braided magnetic field between two perfectly conducting parallel plates. Although the magnetic field in the simulations is initially approximately force-free, it was described in Paper I that the field experiences an instability on some characteristic time-scale. Here we considered the subsequent evolution of the system, which is best described as a resistive relaxation. The route taken to find a new equilibrium involves the formation of a complex array of current sheets which are scattered throughout the domain. The current sheets have a ribbon-like appearance, tending to be highly elongated in the direction along the loop, i.e.~parallel to the strong axial field. {The formation of this {myriad} % of current layers suggests that a turbulent cascade may develop for higher Reynolds numbers than we have been able to use. Due to the absence of boundary forcing, this turbulence would fall under the heading of decaying turbulence. It would be of interest to push the spatial resolution of our simulations (thus allowing the reduction of $\\eta$) to investigate whether a regime of fully-developed turbulence arises, by examining spectral properties of the magnetic and velocity fields. This is beyond the scope of this paper, and we leave it to a future investigation.} The relaxation process as described above results in a simplification of the magnetic field structure as demonstrated by the mapping of field lines from one line-tied boundary to the other (Figure \\ref{fig:twofluxtubes}). The magnetic field lines are untangled, such that no three field lines are braided about one another any longer in the final state. The final state in fact consists of two (weakly) twisted flux tubes, of oppositely signed twist, embedded in an approximately uniform field. This final state approximates a force-free field, $\\nabla \\times {\\bf B}=\\alpha\\BB$, in which field lines in one twisted flux tube have positive $\\alpha$ and in the other negative $\\alpha$. Thus the { system approaches} a {non-linear} force-free field, which is not consistent with the Taylor relaxation picture put forward by \\cite{heyvaerts1984}. In principle since the net current in the system is zero, a Taylor relaxation would lead to a final state with $\\alpha=0$ on every field line, i.e.~the homogeneous potential field. Clearly the single constraint (global helicity conservation) of the Taylor hypothesis is not sufficient to describe this relaxation. Extra constraints on the relaxation which prevent some of the helicity cancellation must be present. {Although it is always dangerous to place too much trust in the extrapolation of simulation results to the extremely high Reynolds number coronal plasma,} it appears unlikely that the cause could be related to the `turbulence' {in the system} not being sufficiently developed. In fact we retain {\\it more} twist in our flux tubes {in the final state} for lower $\\eta$, where the turbulence is better developed. In fact, an additional constraint on the relaxation has recently been discovered \\citep{yeates2010}. The `unbraiding' of the magnetic flux during the relaxation to a non-linear force-free field occurs by magnetic reconnection. This reconnection occurs in the absence of magnetic nulls or closed field lines (note the presence of a strong background field throughout the domain). Furthermore, reconnection occurs in a multitude of regions that are spread throughout the volume. In order to determine the efficiency of the reconnection process, we must therefore sum the reconnection rate over all reconnection diffusion regions within the volume. The outcomes of performing such an analysis are as follows: \\begin{enumerate} \\item The global reconnection rate continues to grow for some time after the peak current in the domain begins to fall. That is, during the intermediate stages of the simulations, the reconnection rate grows {\\it not} because of an intensification of the current or an increase of the rate in any {\\it one} region, but rather because of an increase in the fragmentation of the volume in which the reconnection processes take place. \\item The peak value of the global reconnection rate is at most weakly dependent on the resistivity. \\item The number of identifiable reconnection regions increases as the resistivity is decreased. \\item The total quantity of magnetic flux that is reconnected is greater than the total poloidal flux present. This implies that the magnetic flux is `multiply-reconnected' in the complex array of reconnection regions. \\item As the resistivity is decreased -- {resulting in an increase in the complexity of the % {multitude} of current layers} -- the average number of reconnections for each unit of flux increases. \\end{enumerate} In summary, the loss of stability and subsequent relaxation of braided coronal loops results in a {complex array} of current layers which permit a dissipation of the stored magnetic energy and the attainment of a lower energy non-linear force-free field. These results yield a new and intriguing potential resolution to the long-standing debate over the validity of Parker's model of coronal heating by topological dissipation. Specifically, while current singularities are not an inevitable result of random field line tangling, once the field line mapping becomes sufficiently complex, an instability will be triggered, causing the electric current structures to collapse to small scales. The complexity that is inherent in the field then results in {the formation of a fragmented system of current layers} (the sum the associated reconnection processes being `fast' in the sense that the rate of energy release is only weakly dependent on the magnetic Reynolds number). The global result is a turbulent dissipation of the excess magnetic energy stored in the field, ultimately in the form of heat. There are many interesting aspects of the above study that warrant further investigation, such as determining what parameters and physical quantities govern the nature of the final state, the possible turbulence properties of the relaxation, and the spatial and temporal distribution of the resulting heating in the loop." }, "1003/1003.6114_arXiv.txt": { "abstract": "We present the results of a pilot wide-field radial velocity and metal abundance survey of red giants in ten fields in the Small Magellanic Cloud (SMC). The targets lie at projected distances of 0.9 and 1.9 kpc from the SMC centre ($m-M=18.79$) to the North, East, South and West. Two more fields are to the East at distances of 3.9 and 5.1 kpc. In this last field we find only a few to no SMC giants, suggesting that the edge of the SMC in this direction lies approximately at 6 kpc from its centre. In all eastern fields we observe a double peak in the radial velocities of stars, with a component at the classical SMC recession velocity of $\\sim 160$ km s$^{-1}$ and a high velocity component at about 200 km s$^{-1}$, similar to observations in H{\\small I}. In the most distant field (3.9 kpc) the low velocity component is at 106 km s$^{-1}$. The metal abundance distribution in all fields is broad and centred at about [Fe/H] $\\sim -1.25$, reaching to solar and possibly slightly supersolar values and down to [Fe/H] of about $-2.5$. In the two innermost (0.9 kpc) Northern and Southern fields we observe a secondary peak at metallicities of about $\\sim -0.6$. This may be evidence of a second episode of star formation in the centre, possibly triggered by the interactions that created the Stream and Bridge. ", "introduction": "The Small Magellanic Cloud (hereafter SMC) is, together with the Large Magellanic Cloud (LMC) the nearest dwarf irregular galaxy to our own, and provides an invaluable laboratory to study star formation and chemical evolution in low mass galaxies. There is recent evidence that the LMC and SMC are on their first pass around the Milky Way and that the SMC may not be bound to the LMC \\citep{kallivayalil06a, kallivayalil06b,besla07}. The SMC may be a rare example of a comparatively isolated dwarf galaxy and possibly even a surviving fragment from the era of reionization. However, the SMC has also been interacting with the LMC during the past few Gyrs and these interactions have modulated the recent star formation history of both galaxies (e.g., \\citealt{bekki05,bekki09} and references therein). The SMC is best modelled as an old dwarf spheroidal galaxy possessing a gaseous disk \\citep{bekki09} that has been distorted by star formation and tidal stresses, giving the galaxy its present irregular appearance (e.g., \\citealt{harris04,cioni06}). The distribution and chemical abundances of field stars in the SMC thus provide clues to its star formation history. Open questions include: whether there is an `edge' to the SMC, the metallicity distribution of its field stars, the presence of a metal abundance gradient and whether a metal-poor halo exists around the SMC or other dwarf galaxies as it does around the Milky Way and other giants. Stars belonging to the SMC have been found along the Magellanic Bridge; an old and intermediate age population out to 5$^{\\circ}$ but only a young population at $\\sim 6.5^{\\circ}$ \\citep{harris07}. \\cite{noel07} explored three fields to the South of the SMC identified SMC stellar sequences belonging to the intermediate age population out to 6.5 kpc from the SMC centre. In other galaxies, \\cite{munoz06} observed LMC stars as far as 23$^{\\circ}$ from its centre. Extended stellar envelopes are also detected around other dwarfs (e.g., \\citealt{minniti96,vansevicius04,hidalgo09}, but at least in some cases, these are actually tidal in origin \\citep{munoz06}. Although stars are proved to exist at large projected distances in many nearby dwarfs, these objects may not represent a classical metal-poor halo as is encountered in the Milky Way or M31. For instance, in the LMC stars studied by \\cite{munoz06}, the metallicity distribution is broad and centred around [Fe/H] $\\sim -1$, with a large range of ages \\citep{gallart04}, unlike the largely old and metal-poor stars that are believed to populate the outer halos of giant galaxies. The SMC itself appears to have formed stars quickly at early epochs reaching a metallicity of $\\sim -1$ and to have then suffered a series of star formation episodes over the past 3 Gyrs, after a period of quiescence, which have produced younger stellar populations and more metal rich stars \\citep{harris04}. In the innermost regions of the SMC, \\cite{carrera08} found an average metallicity of [Fe/H] $\\sim -1$, in agreement with previous studies, but also claimed to have detected a metal abundance gradient (richer inward), arguing that this is related to an age gradient, with younger (and more metal rich) stars towards the SMC centre. While this agrees with the earlier work of \\cite{piatti07a,piatti07b}, the study of SMC clusters and field giants (in the proximity of clusters) by \\cite{parisi08,parisi10}, as well as the work by \\cite{cioni09} on the C/M ratio of AGB stars in the SMC, do not support the existence of the metal abundance gradient claimed by \\cite{carrera08}. In this {\\it Letter} we report on a pilot program for an extensive radial velocity and Calcium Triplet survey of the SMC, based on data collected during a similar survey of the Galactic Bulge. The observations and data reductions are described in the next section, while we present the main results and our discussion in the following sections. We adopt the most recent distance modulus of 18.79 for the SMC \\citep{ sz09}. ", "conclusions": "The data we present here show that, while the SMC is detected to large distances (about 6 kpc) along the Magellanic Bridge \\citep{harris04} and to the South \\citep{noel07}, we appear to have approached the edge of the SMC in our easternmost fields. We estimate that the SMC `edge' in this direction lies at about 6 kpc from its centre. The shape of the SMC has however been tidally distorted by interactions with the LMC and it has been elongated along the N-S direction \\citep{kunkel00}. Exploration of the radial and azimuthal behavior at larger distances will be one of the outcomes expected from a wider-field spectroscopic survey. The kinematics of stars in our fields is complex. There is evidence for the presence of two components in some fields, particularly to the East and the South, i.e. the regions most affected by interactions with the LMC and the Magellanic Bridge: a low velocity one around 160 km s$^{-1}$ and a high velocity component at about 210 km s$^{-1}$. Although the evidence is weak, while the high velocity component is at the same position in our 3.9kpc field to the East, the low velocity component appears to have lower velocity. This is reminiscent of the claims for multiple peaks in the H{\\small I} velocity distribution \\citep{ mathewson84,stanimirovic04}. Similar bifurcations are also observed in tidal streams. The H{\\small I} features are often attributed to multiple and overlapping gas shells, but their presence in the stellar distributions, especially in the zones to the East and South closer to the Magellanic Bridge, may favor multiple components models such as those of \\cite{mathewson84}. An intriguing possibility is that we are detecting stars from the LMC in the SMC Eastern fields: \\cite{munoz06} find the presence of LMC stars as far 23$^{\\circ}$ from the LMC centre (which of course lies to the East of the SMC), while \\cite{bekki08} has argued for the existence of a common halo encompassing the LMC and SMC. Based on the recent study of LMC kinematics by \\cite{vandermarel02}, we would expect LMC stars to lie at the position of the second velocity peak we observe in the Eastern fields. This seems somewhat less likely because we find that the secondary peaks in our data contain about the same number of stars as the primary velocity peaks, and therefore appear more likely to be associated with the SMC velocity structure than the LMC. Since our easternmost fields are closer to the LMC than the 0.9 kpc eastern field, we would expect the LMC contribution (if it causes the secondary peak) to increase `outward' from the SMC, unlike the observations. It is of course very likely that some LMC stars are actually superposed over the SMC, but they can probably be securely separated out only by chemical tagging. One striking feature we observe in our data is a broad metallicity distribution centered on [Fe/H] $\\sim -1.2$, extending from $-2.5$ to solar or even slightly supersolar values. This is very similar to what observed in the LMC by \\cite{munoz06} and in Sagittarius by \\cite{monaco05} and may suggest a very similar chemical evolution pattern in most dwarf galaxies. In fact the abundance distribution we observe in the SMC is also very similar to that measured for the M31 ``giant stream'' population, with a peak near [Fe/H] $\\sim -1$ and tails to high and low metallicities \\citep{koch08}. The ingestion of massive galaxies such as the SMC has been invoked to explain the wide metallicity distribution and the presence of metal-rich stars in the M31 halo (see e.g., \\citealt{koch08}). In the case of the SMC there is the question of how a galaxy massive enough to host such metal-rich stars could have been accreted by the SMC without more significant disruption of the SMC (but see \\citealt{tsujimoto09}). The broad metallicity distribution may instead imply the presence of multiple stellar generations. It is known that the SMC has undergone recent star formation, possibly induced by encounters with the LMC, after a long period of quiescence \\citep{harris04}. \\cite{carrera08} claim that there is a metal abundance gradient in the SMC and suggest that this is due to the presence of younger stars in the centre of this galaxy. We find that the main population of the SMC does not exhibit a metal abundance gradient \\citep{parisi08,parisi10}, but that in the inner fields to the North and South there is a contribution from more metal-rich stars, with peak metallicity around [Fe/H] $\\sim -0.6$. \\cite{noel07} find evidence for an intermediate age population in their fields to the South out to 6.5 kpc, while a younger stellar population is detected by \\cite{harris04} along the Magellanic Bridge. The presence of more metal-rich stars forming a separate peak in the inner fields resembles the picture of \\cite{carrera08} where a recent burst of star formation has led to self-enrichment in the inner regions. The approximate North-South trend is roughly in the directions of the Bridge and Stream features and it is tempting to speculate that the interactions that created these gaseous features are also responsible for the star formation episodes. A wider and larger spectroscopic survey will allow us to clarify the structure and kinematics of the SMC, explore the existence of metallicity gradients, search for a metal poor halo and detect the presence of streams." }, "1003/1003.1734_arXiv.txt": { "abstract": "We report on the discovery of a molecular cavity in the Norma near arm in the general direction of Westerlund 1 (Wd1), but not associated with it. The cavity has a mean radial velocity of $-91.5$~km\\,s$^{-1}$, which differs by as much as $\\sim$40~km\\,s$^{-1}$ from the mean radial velocity of the Wd1 stars. The cavity is surrounded by a fragmented molecular shell of an outer diameter of about 100\\,pc and 10$^{6}$M$_\\odot$, which is expanding at velocities of 6 to $8$~km\\,s$^{-1}$. The amount of kinetic energy involved in the expanding shell is $\\sim10^{51}$ erg. Inside this cavity the atomic HI gas surface density is also the lowest. Structure of the extended Very High Energetic (VHE) $\\gamma$-ray emission, recently reported by the H.E.S.S. collaboration \\citep{ohm09}, coincides with the cavity. The observed morphology suggests that the inner wall of the molecular shell is the zone of the $\\gamma$-ray emission, and not the dense gas surrounding massive stars of Wd1 as had been speculated by the H.E.S.S. collaboration. A likely candidate responsible for creating the observed cavity and the $\\gamma$-ray emission is the pulsar PSR J1648-4611. ", "introduction": "Recently, the H.E.S.S. collaboration carried out a search for Very High Energy (VHE, E$>$100~Gev) $\\gamma$-ray emission from Galactic young stellar clusters and found positive detections in the direction of Westerlund 1 (Wd1) and Westerlund 2 (Wd2) (Ohm et al. 2009; Ohm09 henceforth). The motivation behind these searches comes from the fact that the stellar clusters are potential acceleration sites of VHE particles, since they host a variety of energetic sources such as Supernova Remnants (SNRs) and pulsar wind nebulae. Extended and point like emission was detected surrounding Wd2 \\citep{aharonian07}, whereas only extended emission off-centered from the cluster was detected in the case of Wd1 (Ohm09). The emission surrounding Wd2 was associated with the cluster by \\cite{dame07}, who found structural coincidences between the CO map and the $\\gamma$-ray source. Ohm09, searched for structural coincidences between the $\\gamma$-ray and HI maps of Wd1, and suggested a possible association of Wd1 with the $\\gamma$-ray emission. One way of confirming this association is to compare the observed $\\gamma$-ray emission structure with the CO gas, which is a well-known tracer of the $\\gamma$-ray emission zones, apart from tracing the spiral arms better than HI in the inner disk (e.g. \\cite{damThad08}). In order to study the Wd1 region in CO, we need to establish the velocity of the Wd1 cluster. \\citet{K&D07}, analyzed the neutral gas environment around the cluster and concluded that Wd1 lies at a radial velocity of $-50$~km~s$^{-1}$, thus locating it in the external part of the Scutum-Crux (SCx) arm. Recent radial velocity measurements of stellar members of Wd1 yield a mean velocity of $-50~$km~s$^{-1}$ \\citep{2009Ap&SS.324..321M, ritchie09}, thus confirming the velocity derived from the HI analysis. We have searched for molecular clouds towards Wd1, in the Columbia-Calan CO survey data cube \\citep{bronf89}. This survey covers the entire southern MW and has a spatial resolution of 0.125$^\\circ$, a velocity resolution of 1~km\\,s$^{-1}$ and a sensitivity of 0.1~K. Our analysis covers large angular scale ($2^\\circ\\times2^\\circ$ centered at l=339.55$^{\\circ}$, b=$-0.4^{\\circ}$), allowing us to study the large scale structure of the interstellar medium around the cluster. The aim of the present work is to search for the molecular clouds associated with the H.E.S.S. $\\gamma$-ray emission, and to investigate whether Wd1 is the source of the observed $\\gamma$-ray emission. With this aim, we first analyzed the morphology and kinematics of the molecular gas at the established radial velocity of Wd1 ($-50$~km\\,s$^{-1}$) in the SCx arm. The resulting poor association of the CO structures with the H.E.S.S map, prompted us to carry out the analysis of the CO morphology of the Norma arm at $-90$~km\\,s$^{-1}$, which is the next major arm in the line of sight (LOS) to Wd1. In this latter arm, we discovered a ring, which could be interpreted as an expanding shell, with very good structural correspondence with the observed $\\gamma$-ray emission. In this {\\it Letter}, we compare the CO maps at both velocities with the $\\gamma$-ray observations, and discuss in some detail the CO structure of the Norma arm. In \\S2, the CO molecular emission in the direction of Wd1 is presented. The association between molecular gas and VHE emission is presented in \\S3, and finally in \\S4, we discuss the consequences of this correspondence. \\begin{figure} \\plotone{fig1.eps} \\caption{{\\bf Top:} Reproduction of the longitude-velocity diagram of \\citet{bronf00} indicating the position of spiral arms along at l=339.75$^\\circ$ (solid line). The flat rotation curve are marked with small squares. {\\bf Bottom:} The $^{12}$CO emission profiles in the direction of Wd1 in a narrow ($15^\\prime\\times15^\\prime$; solid line), and wide ($2^\\circ\\times2^\\circ$; dashed line scaled by a factor of 1/30) beams. The range of radial velocities for 3 of the arms are indicated by the shaded boxes. The arrows point to the radial velocities obtained by fitting Gaussian profiles to the CO emission peaks relevant to this study.} \\end{figure} ", "conclusions": "We analyzed the morphology and kinematics of the molecular environment in the general direction of Wd1, with the aim of looking for the molecular gas that may be associated with the recently reported H.E.S.S. $\\gamma$-ray emission. We discovered an expanding molecular ring at the radial velocity of $-91.5$~km\\,s$^{-1}$, with the inner contours of the ring coincident with the observed $\\gamma$-ray emission structure. The observed morphologies of the $\\gamma$-ray emission and the CO emission along with their kinematics, can be understood in terms of an expanding fragmented shell of molecular gas. Thus, the 2 components are physically associated and both lie in the Norma arm. Consequently, the molecular structure surrounding the Wd1 ($-50$~km\\,s$^{-1}$) cluster does not correspond to the observed $\\gamma$-ray emission map. The $\\gamma$-ray emission is produced by accelerated energetic particles when they interact with dense gas. Recently, \\citet{fujita09} proposed a model of generating VHE $\\gamma$-rays by a SNR located in a cavity surrounded by high density molecular gas in order to explain the observed photon spectra for the hidden SNR in the open cluster Wd2, and the old-age mixed-morphology SNR W28. In this model, the particles are accelerated at the end of the Sedov phase, when the SN shock reaches and collides with the surrounding spherical high-density molecular gas (shell). This interaction produces $\\gamma$-ray emission in the inner wall of the molecular shell, exactly the configuration that produces our observed CO morphology. Hence, the presence of a stellar remnant inside the molecular cavity can give rise to the observed $\\gamma$-ray emission. There are three stellar remnant candidates in the field of view. They are a binary system GX340+0 containing a neutron star \\citep{schultz93}, a magnetar (CXO J164710.2-455216), and the pulsar PSR J1648-4611, with the third one the most likely precursor. The first two sources are seen geometrically inside the molecular cavity, but lie at distances different from that of the cavity according to their present distance estimations. The magnetar is associated with Wd1 following the analysis by \\citet{muno06}, based on statistical grounds. However, it could as well be located further out, in the molecular cavity. For the neutron star in the binary system GX340+0, a maximum distance of 11$\\pm$3~kpc has been estimated based on its luminosity, assuming that it is radiating at the Eddington limit \\citep{penninx93}. Yet, it is conceivable that the neutron star is not emitting with the limiting luminosity, in which case it could be located at the distance of our cavity. A third stellar remnant (PSR J1648-4611) is located half a degree South-East of the ring center and is seen just at the outer border of the molecular ring (see Fig.~4). A distance by dispersion measure yields a value of 5--5.7~kpc, depending on the electron density distribution model adopted \\citep{Torres03}. That is at the same distance of the cavity. This makes this pulsar a strong candidate for producing the $\\gamma$-ray emission at the inner boundaries of the molecular shell. The location of the pulsar off-center of the cavity, could be due to a large proper motion of the pulsar. following the SN explosion. It requires only a velocity as little as 8~km\\,s$^{-1}$ for the pulsar to move from the center of the shell to its present position, if the explosion occurred 6~Myr ago, and the pulsar is moving in the plane of the sky. Recently, the LAT instrument on board of the Fermi satellite detected unpulsed emission from a region coincident with this pulsar \\citep{abdo09}. The model proposed by \\citet{fujita09} expects emission in the Fermi band coincident with the $\\gamma$-ray emission, but such emission is not yet detected by Fermi observations. However, the observed configuration of a cavity surrounded by a molecular shell, with associated VHE $\\gamma$-rays emission, and a bright Fermi source in the vicinity, are common to both, Wd2 and the expanding shell discussed in our work." }, "1003/1003.4327_arXiv.txt": { "abstract": "{In this work we study Herbig-Haro objects located in the region around the head of the cometary globule CG~30. Two sets of optical images are presented. The first set was obtained with the 3.5~m New Technology Telescope in 1995 in three emission lines: H$\\alpha$, S~II$\\lambda\\lambda$6731,6716~\\AA\\ and [O~II]$\\lambda$3729~\\AA. The second set is an H$\\alpha$ image of the CG~30/31/38 complex obtained in 2006 with the 8~m Subaru telescope. A proper motion study of the HH objects in the region was performed using the H$\\alpha$ images from both epochs. Due to the high resolution of our images we were able to, for the first time, resolve the HH~120 object into ten knots and measure proper motions for some of them. We discover several new HH objects and a large bipolar jet, HH~950, emerging from the head of CG~30. We suggest that two previously known submillimeter sources are the driving sources for the HH~120 and HH~950 flows.} ", "introduction": "\\label{sec:Introduction} Cometary globules (CGs) are compact interstellar molecular clouds usually associated with H~II regions and OB stars. A particularly fine group is located in the Gum Nebula (Hawarden \\& Brand 1976; Zealey et~al. 1983; Reipurth 1983). They show compact and bright-rimmed heads and faint tails that extend from the heads and point away from the nearby bright young photoionizing stars (Fig.~1). Their sizes range between 0.1-0.8~pc, they exhibit high densities, 10$^4$-10$^5$~cm$^{-3}$ and temperatures around 10 K. Their typical masses range between 10-100~M$_\\odot$. Cometary globules are the sites of star formation (see Reipurth 1983). In some cases, bipolar Herbig-Haro (HH) flows have been observed emerging from the heads of cometary globules (Schwartz 1977; Reipurth 1983; Pettersson 1984). The largest known optical H~II region in the Galaxy is the Gum Nebula (Gum, 1952). It is located in the Galactic plane ($l$=258$^\\circ$, $b$=-2$^\\circ$) at a distance of about 450~pc, and is excited by $\\zeta$ Pup, $\\gamma^2$ Vel, and in the past by the progenitor to the Vela pulsar. Its linear diameter is 250~pc. With its apparent diameter of 36$^\\circ$\\ it covers a large part of the southern sky in the constellations Vela, Puppis, Pyxis, Canis Major and Carina. One of the most studied groups of CGs in the Gum Nebula is the CG~30/31/38 complex. \\subsection{CG~30 and HH~120} In this paper we present two sets of very deep images of CG~30 (see \\S~2). CG~30 has been studied by various authors at different wavelengths. It has been recognized as a region of star formation due to its association with the Herbig-Haro object HH~120, first noted by Westerlund (1963) and Reipurth (1981). The earliest spectroscopic measurements of HH~120 were made by Pettersson (1984), who obtained spectra of two different parts of HH~120, one being its brightest condensation (knot A). Its mean radial velocity of -42~km~s$^{-1} \\pm$12~km~s$^{-1}$ is consistent with the value -39~km~s$^{-1}$ obtained by Schwartz \\& Greene (2003). The electron temperature T$_e$ and the electron density n$_e$ are found by Pettersson (1984) to be 9100~K$\\pm$400~K and 1700~cm$^{-3} \\pm$600 ~cm$^{-3}$, respectively. Also five IR sources associated with CG~30 were observed. It was argued that one of them, CG30-IRS4 (also known as IRAS~08076-3556), is an energy source for HH~120 and its associated reflection nebula. Graham \\& Heyer (1989) imaged HH~120 in the R (0.6~$\\mu$m), J (1.2~$\\mu$m), H (1.6~$\\mu$m) and K (2.2~$\\mu$m) bands. Like Pettersson (1984) they also detected a source CG30-IRS4 in H and K band images. Scarrott et~al. (1990) performed optical polarization studies of a nebula seen against the central dark region of CG~30 that contains HH~120. It was confirmed that the so called CG~30 Nebula is predominantly a reflection nebula illuminated by the IR source CG30-IRS4. Persi et~al. (1994) observed HH~120 and its energy source CG30-IRS4 in the near infrared, mm continuum, and an ammonia line. It was recognized that the IRAS 08076-3556 source is a very young low mass Class~I source embedded in a dense core of CG~30. The measurements of the 1.3~mm luminosity and the bolometric luminosity show that their ratio is actually closer to that of the Class~0 sources. The authors stated that the strong 1.3~mm emission is probably due to a circumstellar dust disk around the IRAS source. Hodapp \\& Ladd (1995) reported the discovery of eight shocked objects in infrared images in the H$_2$ 1-0 S(1) emission line. On the basis of their relative positions the authors concluded that they form parts of two outflows propagating in approximately perpendicular directions. None of the supposed flows could be directly associated with the optical HH~120 object. Nielsen et~al. (1998) observed the cometary globules CG~30/31/38 in $^{13}$CO, $^{12}$CO and H$_2$ lines. They proposed that CG~30 with a mass of $\\sim$10~M$_\\odot$ is associated with another globule along the line-of-sight having a mass close to 2~M$_\\odot$ and moving with a velocity of 4~kms$^{-1}$ with respect to the local standard of rest. Also a dense molecular outflow associated with CG~30 was detected in the $^{12}$CO~(J=1-0) line and its total mass was estimated to be 0.28~M$_\\odot$. They reported the maximum flow velocity to be 9~kms$^{-1}$ and its dynamical age to be 1.7$\\times$10$^{4}$~years. This flow is propagating in a direction perpendicular to the tail of CG~30 and is originating from the position of the CG30-IRS4 source. Bhatt (1999) discussed the role of magnetic fields in cometary globules. The author found that in the case of the CG~30/31 complex, the light from the stars in the region is polarized in the range from $\\sim$0.1 to $\\sim$4 per cent. The polarization vectors seem to be perpendicular to the direction of the tails of CG~30 and 31, but, interestingly, almost aligned with the direction of the molecular flow detected by Nielsen et~al. (1998). Bhatt also claimed that if the polarization is caused by dust grains aligned by the magnetic field, then the polarization vectors must be parallel to the projected magnetic field in the region. The tails of CG~30 and 31 are much shorter and more diffuse than the tail of CG~22, where the magnetic field is found to be parallel to the globules's tail. According to Bhatt, CG morphology depends on the relative orientation of the cloud magnetic field and the radius vector of the CG head from the central source of radiation and the winds that produce the cometary tails. Long and narrow tails are to be observed when the magnetic field is parallel to the radius vector, while short and diffuse tails develop if those two are perpendicular. In their search for young Solar System analogues, Zinnecker et al. (1999) observed four regions containing low luminosity sources associated with extended reflection nebulosities, among which was also CG~30. The images were obtained in near-IR (J, H and K) broad-bands. In their images they observed the confines of the $\\sim$0.3~pc diameter dark globule with the very red CG30-IRS4 source at the center and the bluer nebulosity just above it. While the widely accepted distance of the Gum Nebula is 450~pc, Knude et~al. (1999, 2000) and Nielsen et~al. (2000) suggested a distance of 200-250~pc on the basis of color-magnitude ((V - I) - V) diagrams and $uvby\\beta$ photometry of the stars that appear to be located in the CG~30/31/38 region. The distance to the globules remains under debate, in this paper we adopt 450~pc. In their survey of Bok Globules, Launhardt et~al. (2000) observed CG~30 at submm wavelengths (850~$\\mu$m). They discovered two sources lying along a north-south direction and separated by a projected distance of $\\sim$9000~AU. The northern condensation was identified as a possible driving source for the HH~120 jet (and with the CG30-IRS4 source) and the southern condensation was proposed to be associated with the larger of the IR flows. The two sources were observed again by Wu et~al. (2007) at 350~$\\mu$m. Nisini et~al. (2002) obtained spectra in the 1-2.5~$\\mu$m wavelength span. The most prominent features observed were [Fe~II] and H$_2$ lines. On the basis of H$_2$ emission they found that HH~120 consists of multiple temperature components probably due to a slow, J type shock. Kim et~al. (2005) studied the low-star mass formation in the CG~30/31/38 complex. They obtained X-ray, optical and near-IR photometry of the stars in the region and found 14 new pre-main sequence (PMS) stars in addition to the 3 previously known stars in the region. According to the authors, these PMS stars belong to two groups: one group having ages of $\\leq$5~Myr at $d =$200~pc, with spectral classes K6 - M4, and the other group of F - G type stars with ages of $<$100~Myr and $d \\sim $2~kpc. They conclude that there were at least two episodes of star formation - ongoing star formation such as in the head of the CG~30 cloud triggered by UV radiation from OB stars, and a formation episode that may have been triggered $<$5~Myrs ago by preexisting O stars, such as the progenitor of the Vela SNR and $\\zeta$~Pup. Chen et~al. (2008a) observed CG~30 in the 3~mm dust continuum, in N$_2$H$^+$~(1 - 0) emission line and at 3 - 8~$\\mu$m wavelengths. The authors detected two sub-cores inside CG~30. From the millimeter continuum observations they derived the gas masses of the two sub-cores to be 1.1~M$_\\odot$ (northern sub-core) and 0.33~M$_\\odot$ (southern sub-core). The authors classified the northern source as a Class~I object and the southern source as a Class~0 protostar. The IR observations revealed two perpendicular collimated bipolar jets that coincide with the knots previously discovered by Hodapp \\& Ladd (1995). The N$_2$H$^+$~(1-0) emission maps revealed two cores spatially associated with the mm continuum dust cores. In a subsequent study, Chen at al. (2008b) studied CG~30 in the $^{12}$CO(2-1) line and 1.3~mm dust continuum. The $^{12}$CO(2-1) observations showed the existence of two bipolar molecular flows propagating in almost perpendicular directions. The authors suggested that one of the flows is associated with the northern compact source and the other one with the southern source. The northern flow exhibits the projected length of 27,000~AU and the position angle P.A. $\\sim$128$^\\circ$. The projected length of the southern flow is 20,000~AU and its direction of propagation (P.A.) is $\\sim$57$^\\circ$. The velocities of the flows are low ($\\lesssim$12~kms$^{-1}$). In this work we study the kinematics of the HH objects in CG 30 by measuring their proper motions. We combine our observations with the existing data at IR and submm wavelengths in order to understand the global outflow properties in the globule. Table~2 lists the coordinates of all the HH objects that appear in our images. We have discovered a large bipolar HH flow, which we catalogue as HH~950. This flow is extending $\\sim$12'2 in the northeast-southwest direction, escaping from the interior of CG~30, and displaying a complex structure with various working surfaces. This paper is organized as follows. In Section~2 we present our data and the reduction techniques applied. In Section~3 we present our methodology and the results of this work. The latter are discussed in Section~4, and the conclusions are given in Section~5. ", "conclusions": "In this work we have studied Herbig-Haro objects associated with CG~30. We find that most of the HH objects belong to two flows: HH~120 and HH~950. The proper motions of the knots in HH~120 suggest that this object is actually composed of at least two outflows. The candidate for the driving source of HH~120 is the northern submm source CG30~SMM-N discoverd by Launhardt et~al. (2000), which must therefore be a binary or a multiple system. The southern submm source CG30~SMM-S is located exactly on the axis defined by six IR objects discovered by Hodapp \\& Ladd (1995). This leaves no doubt that these objects really do form one flow and that this submm source is its driving source (as was already suggested by Launhardt et al. 2000). We propose that this flow forms a part of the large HH~950 flow that appears in our optical images. HH~950 consists of two lobes - while its northeastern lobe is almost featureless (the only visible feature is its bright rim), the southwestern lobe shows a complex structure made of fine filaments and knots. Interesting is the fact that the HH knots in this lobe do not all lie on the same axis. It seems that the axis swiftly changed toward the south at the time the knots B and C were produced. We argue that CG30~SMM-S could be a binary system. At least one member of the system is a young star emitting the HH~950 jet. The interaction of both objects causes the change of direction of the HH~950 axis. {\\bf Acknowledgments}\\\\ We are grateful to an anonymous referee for a very careful and helpful report. We thank Ralf Launhardt for providing the accurate positions of the submillimeter sources and Klaus Hodapp for providing us with the original infrared image of the CG~30. Primo\\v{z} Kajdi\\v{c} acknowledges the Direcci\\'on General de Estudios de Posgrado of the UNAM for a scholarship supporting his graduate studies. This study has been supported by the NSF through grant AST0407005. This material is based upon work supported by the National Aeronautics and Space Administration through the NASA Astrobiology Institute under Cooperative Agreement No. NNA04CC08A issued through the Office of Space Science. \\begin{table}\\centering \\caption{Coordinates of the observed HH objects in the region around CG~30} \\begin{tabular}{ccccc} \\hline \\hline Feature name & R.A. (hh:mm:ss.s) & Dec. (dd:mm:ss) \\\\ \\hline \\hline \\multicolumn{3}{c}{HH~948}\\\\ A & 08:09:30.6 & -36:04:25\\\\ B & 08:09:30.6 & -36:04:19\\\\ C & 08:09:31.8 & -36:04:16\\\\ \\hline \\multicolumn{3}{c}{HH~949}\\\\ A & 08:09:36.8 & -36:04:29\\\\ B & 08:09:37.2 & -36:04:27\\\\ C & 08:09:37.4 & -36:04:26\\\\ D & 08:09:36.8 & -36:04:25\\\\ E & 08:09:37.2 & -36:04:25\\\\ F & 08:09:36.5 & -36:04:35\\\\ G & 08:09:35.7 & -36:04:50\\\\ H & 08:09:37.7 & -36:04:20\\\\ \\hline \\multicolumn{3}{c}{HH~120}\\\\ A & 08:09:32.5 & -36:04:55 \\\\ B & 08:09:32.9 & -36:04:56 \\\\ C & 08:09:32.6 & -36:04:56 \\\\ D & 08:09:32.8 & -36:04:56 \\\\ E & 08:09:33.1 & -36:04:56 \\\\ F & 08:09:33.0 & -36:04:55 \\\\ G & 08:09:32.3 & -36:04:58 \\\\ H & 08:09:32.3 & -36:04:55 \\\\ I & 08:09:32.5 & -36:04:53 \\\\ J & 08:09:32.5 & -36:04:51 \\\\ K & 08:09:32.4 & -36:04:44\\\\ \\hline \\multicolumn{3}{c}{HH~950}\\\\ A & 08:09:14.6 & -36:07:15 \\\\ B & 08:09:16.5 & -36:07:12 \\\\ C & 08:09:17.8 & -36:07:27 \\\\ D & 08:09:22.1 & -36:07:07 \\\\ E & 08:09:24.6 & -36:06:40 \\\\ F & 08:09:28.6 & -36:06:10 \\\\ \\hline \\multicolumn{3}{c}{The submm sources (from Wu et~al. (2007)).}\\\\ North & 08:09:32.87 & -36:04:56.3 \\\\ South & 08:09:32.55 & -36:05:16.6 \\\\ \\hline \\end{tabular} \\label{tab-feathh120} \\end{table} \\begin{table} \\caption{Proper motions: velocity and angle} \\begin{center} \\begin{tabular}{cccc} \\hline \\hline Feature & Proper motion ( ''/century) & kms$^{-1}$ & P.A. (degrees)\\\\ \\hline \\hline HH~120A, B, C, D & 1.3 &26 & 332\\\\ HH~120E & 2.1 & 45& 309\\\\ HH~948B & 1.4 & 31&308\\\\ HH~950A & 3.8 & 82&253\\\\ HH~950C & 4.2 & 89&240\\\\ HH~950D & 1.8 & 38&250\\\\ HH~950E & 5.0 & 107&244\\\\ \\hline \\end{tabular} \\end{center} \\label{tab-pm} \\end{table} \\begin{figure} \\centering \\includegraphics[width=1.0\\textwidth]{Figure01.eps} \\caption{The H$\\alpha$ image of the CG~30/31/38 complex obtained with the Subaru telescope and Subaru Prime Focus Camera on January 4, 2006, composed of five six-minute exposures. This is the deepest image of the complex ever obtained. The angular size of the image is 21'6$\\times$23'5. North is up and east is left.} \\label{fig-cg30} \\end{figure} \\begin{figure} \\centering \\begin{tabular}{cc} \\includegraphics[width=0.5\\textwidth]{Figure02a.eps}& \\includegraphics[width=0.5\\textwidth]{Figure02b.eps}\\\\ \\includegraphics[width=0.5\\textwidth]{Figure02c.eps}& \\end{tabular} \\caption{The three images were obtained with ESO NTT telescope and ESO Multi Mode Instrument, using three narrow band filters: H$\\alpha$ (top), [O~II]~$\\lambda$3729~\\AA\\ (medium) and [S~II]~$\\lambda$6730+6716~\\AA\\ (bottom). Each image is a 30 minute exposure. The angular size of the images is 5'8$\\times$5'7. North is up an east is left.} \\label{fig-colors} \\end{figure} \\begin{figure} \\centering \\begin{tabular}{c} \\includegraphics[width=1.0\\textwidth]{Figure03a.eps}\\\\ \\includegraphics[width=1.0\\textwidth]{Figure03b.eps} \\end{tabular} \\caption{The detailed Subaru H$\\alpha$ images of HH~120. In addition to the knot A which had already been observed, we resolve ten other knots: B, C, D, E, F, G, H, I, J and K. Bottom: Logarithmically scaled countours are plotted on top of the HH~120 image. Only knots A, E and G are resolved by the contours.} \\label{fig-sub} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=1.0\\textwidth]{Figure04.eps} \\caption{The three NTT images have been combined into a single image in order to have an overview of all of the HH objects in the region. The HH knots located around HH~120 are best visible in the [S~II] image and so are the newly discovered objetcs (HH~949, HH~948A, B and C, HH~120K) The HH~950 knots are best resolved in the [O~II] image (the HH~950 knots C and D only appear on this image). The field of view (FOV) of this image is 5'8$\\times$5'7. North is up and east is left.} \\label{figure-combine} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=1.0\\textwidth]{Figure05.eps} \\caption{The previous image was enlarged to show the structure of the HH~949. Also, the logarithmically scaled contours were superposed in order to better resolve the individual knots of this flow. Eight knots can be resolved. Six of them (A to F) have previously been labeled by Hodapp \\& Ladd (1995) as IR knot 7, one of them (H) as IR knot 8 and G has been known as IR knot 5. The angular size of this image is 42''$\\times$34''.} \\label{hh949} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=1.0\\textwidth]{Figure06.eps} \\caption{This image includes a hand made sketch in order to make it easier for the reader to recognize different features in the image. The shape of the HH~950 jet and all the HH objects are marked. The FOV of the image is 8'9$\\times$7'6. North is up and east is left.} \\label{fig-draw} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=1.0\\textwidth]{Figure07.eps} \\caption{Proper motion vectors for individual features in the CG~30 complex. The boxes indicate regions used for cross correlation. The values for the vectors are given in Table~\\ref{tab-pm}. The FOV is 8'6$\\times$9'1. North is up and east is left.} \\label{fig-pm} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=1.0\\textwidth]{Figure08.eps} \\caption{The infrared image from Hodapp \\& Ladd (1995) is superposed onto the large Subaru image, so we can directly see the positions of the eight infrared shocked objects and put them in the context of our optical data. The two crosses mark the positions of the two submm sources found by Launhardt et~al. (2000). For a detailed discussion, see \\S~4.The angular size of the image is 8'6$\\times$7'7.} \\label{fig-krizi} \\end{figure}" }, "1003/1003.3488_arXiv.txt": { "abstract": "We present precise Doppler measurements of four stars obtained during the past decade at Keck Observatory by the California Planet Survey (CPS). These stars, namely, HD\\,34445, HD\\,126614, HD\\,13931, and Gl\\,179, all show evidence for a single planet in Keplerian motion. We also present Doppler measurements from the Hobby-Eberly Telescope (HET) for two of the stars, HD\\,34445 and Gl\\,179, that confirm the Keck detections and significantly refine the orbital parameters. These planets add to the statistical properties of giant planets orbiting near or beyond the ice line, and merit follow-up by astrometry, imaging, and space-borne spectroscopy. Their orbital parameters span wide ranges of planetary minimum mass ($M$\\,sin\\,$i$\\,=\\,0.38--1.9\\,\\mjupe), orbital period ($P$\\,=\\,2.87--11.5\\,yr), semi-major axis ($a$\\,=\\,2.1--5.2\\,AU), and eccentricity ($e$\\,=\\,0.02--0.41). HD\\,34445\\,b ($P$\\,=\\,2.87\\,yr, $M$\\,sin\\,$i$\\,=\\,0.79\\,\\mjupe, $e$\\,=\\,0.27) is a massive planet orbiting an old, G-type star. We announce a planet, HD\\,126614\\,Ab, and an M dwarf, HD\\,126614\\,B, orbiting the metal-rich star HD\\,126614 (which we now refer to as HD\\,126614\\,A). The planet, HD\\,126614\\,Ab, has minimum mass $M$\\,sin\\,$i$\\,=\\,0.38\\,\\mjup and orbits the stellar primary with period $P$\\,=\\,3.41\\,yr and orbital separation a\\,=\\,2.3\\,AU. The faint M dwarf companion, HD\\,126614\\,B, is separated from the stellar primary by 489\\,mas (33\\,AU) and was discovered with direct observations using adaptive optics and the PHARO camera at Palomar Observatory. The stellar primary in this new system, HD\\,126614\\,A, has the highest measured metallicity (\\feh\\,=\\,+0.56) of any known planet-bearing star. HD\\,13931\\,b ($P$\\,=\\,11.5\\,yr, $M$\\,sin\\,$i$\\,=\\,1.88\\,\\mjupe, $e$\\,=\\,0.02) is a Jupiter analog orbiting a near solar twin. Gl\\,179\\,b ($P$\\,=\\,6.3\\,yr, $M$\\,sin\\,$i$\\,=\\,0.82\\,\\mjupe, $e$\\,=\\,0.21) is a massive planet orbiting a faint M dwarf. The high metallicity of Gl\\,179 is consistent with the planet-metallicity correlation among M dwarfs, as documented recently by Johnson \\& Apps. ", "introduction": "\\label{sec:intro} The distributions of the masses and orbits of jovian-mass exoplanets offer key tests of planet formation theory. Most theories predict that giant planets form beyond the ``snow line'' and migrate inward on a time scale that competes with the lifetime of the protoplanetary disk \\citep{Thommes08, Ida_Lin08_v, Rice05, Alibert05, Trilling02}. Among various theories for formation, core accretion has been shown efficient at producing, within $\\sim$3\\,Myr, planets of Neptune to Jupiter mass, orbiting within 5\\,AU \\citep{Benz08, Dodson-Robinson08}. These models of giant planet formation and orbital evolution may be directly tested against observations of giant planets found by the Doppler method. The models predict a clearing of gaps in the protoplanetary disks, establishing the distribution of planet masses, allowing a direct comparison with Doppler observed minimum masses (\\msinie). According to the models, the disk dissipates after the planets have undergone some inward migration, leaving them at their current orbital distances. The resulting predicted distribution of semimajor axes can be compared with the observed orbits of giant planets. There remains potential value in both enhancing the sophistication of planet formation theory and in observing a large enough statistical sample of giant planets to permit robust and informative tests of the theory. Moreover, planet-planet interactions among giant planets must be predicted and compared with the distributions of orbital elements (especially eccentricity) for systems containing multiple giant planets, e.g.\\ \\cite{Wright09,Ford_Chiang07,Ford07,Juric08}. Giant planets also gravitationally interact with the dust in their planetary system to shape, on timescales of only years, the dust evolution of the planetary system \\citep{Lisse07,Beichman07,Payne09}. As of May 2009, 350 exoplanets have been discovered, with remarkable properties including close-in orbits, large orbital eccentricities, multi-planet systems, and orbital resonances \\citep{MayorUdry08, Marcy2008}. The hot jupiters have received the most attention, observationally and theoretically, yielding extraordinary information about their chemical composition, internal structure, atmospheric behavior. However most known gas giant planets orbit beyond 1\\,AU, realizing the population that formed beyond the ice line as predicted by theory, and as seen in our Solar System. In 1997, we began a Doppler search for giant planets in Jovian orbits at Keck Observatory. We monitor over 1800 stars within 50\\,pc, with special attention given to those within 20\\,pc. We have acquired a Doppler time baseline of well over 8\\,yr for nearly all of them. The detected long-period exoplanets reveal the distribution of their masses, semimajor axes, and orbital eccentricities for the general population of planetary systems. Remarkably, the exoplanets exhibit a sharp rise in occurrence beyond 1\\,AU \\citep{Johnjohn07,Cumming08}, indicating a great population of giant planets resides there. Many of these planets remain undiscovered even after 10\\,yr of Doppler monitoring, as the amplitudes of a few meters per second require high Doppler precision and a clear indication of Keplerian motion, which is challenging for orbital periods comparable to the duration of observations. Nonetheless, the analysis of Doppler completeness shows that 15--18\\% of all nearby stars have giant planets between 3--20\\,AU \\citep{Cumming08}. Surely, these giant planets offer strong statistical information on the formation and subsequent dynamical evolution of gas giants in general. Unfortunately, with only a decade of data orbits beyond 4\\,AU are just coming into our Doppler field of view. The recently announced Jupiter-analog orbiting HD\\,154345 with $a$\\,=\\,5.0\\,AU (and a circular orbit) exhibited nearly one full orbital period only after 10 full years of Doppler data were collected \\citep{Wright_154345}. But the number of giant planets known orbiting beyond 1\\,AU remains so small that extraordinary statistical efforts are required to extrapolate the true underlying properties \\citep{Cumming08}. Thus, there remains a need for enlarging the observed population of giant planets, especially with the prospect of follow-up observations by such instruments as the James Webb Space Telescope (JWST), the Gemini Planet Imager (GPI), and the Spectro-Polarimetric High-contrast Exoplanet Research (SPHERE) instrument for the Very Large Telescope (VLT). In the future, knowledge of giant planets around the nearest stars will be crucial for detecting Earth-sized planets, e.g.\\ by the Space Interferometry Mission \\citep{Unwin08}, as the giant planets add a ``noise'' to the astrometric signal \\citep{Traub09}. One concern is that multiple giant planets orbiting beyond 2\\,AU will cause curved astrometric motion, with the linear part being (erroneously) absorbed into the inferred proper motion. The resulting residuals will have a time scale of $\\sim$1\\,yr, constituting an astrophysical ``noise'', compromising the detection of the terrestrial planets. Thus, future astrometric missions will benefit greatly from $\\sim$15\\,yr of radial velocity (RV) data. The characterization of both giant and rocky planets will be important for future missions that image and take spectra of planets, such as the Terrestrial Planet Finder and Darwin \\citep{Kaltenegger06, Lawson08}. As a result we continue to survey nearby stars that are likely targets for such surveys \\citep{Kaltenegger08}. Here we describe Doppler measurements from Keck Observatory and from the Hobby-Eberly Telescope (HET) for four stars that shows signs of harboring a planet beyond 1\\,AU. ", "conclusions": "\\label{sec:discussion} We present the detection of four new extrasolar planets. These planets add to the statistical properties of giant planets in long-period orbits near the ice line. Long observational time baselines (10--12\\,yr) were necessary to accurately measure the high-eccentricity, low-amplitude signals of HD\\,34445b and HD\\,126614\\,Ab, as well as the intermediate-amplitude, but long-period signals of HD\\,13931\\,b and Gl\\,179\\,b. HD\\,34445\\,b is a massive planet ($M$\\,sin\\,$i$\\,=\\,0.79\\,\\mjupe) in a mildly eccentric ($e$\\,=\\,0.27), long-period ($P$\\,=\\,2.87\\,yr) orbit around an old G0 star. We clearly detect this planet in the Keck and HET data sets individually, and their combination allows for more precise orbit determination. The relatively large residuals to the one-planet fit (rms = 6--8\\,\\mse) hint at a second, unresolved planet in the system. Underestimated jitter would also explain the large residuals, but we deem this explanation less likely given the metallicity, color, and modest evolution of HD\\,34445. HD\\,126614\\,Ab is a massive planet ($M$\\,sin\\,$i$\\,=\\,0.38\\,\\mjupe) in a long-period ($P$\\,=\\,3.41\\,yr), eccentric ($e$\\,=\\,0.41) orbit around an extremely metal-rich star. At \\feh\\,=\\,+0.56\\,$\\pm$0.04, HD\\,126614\\,A has the highest metallicity of the 1040 stars in the SPOCS catalog \\citep{Valenti05}. It also has the highest \\feh\\ of the $\\sim$250 stars with known planets and measured metallicities. We confirmed the high metallicity of HD\\,126614\\,A by running a separate iodine-free HIRES spectrum through the same SME pipeline used for the SPOCS catalog. We found \\feh\\,=\\,+0.51\\,$\\pm$0.04, consistent with the SPOCS catalog value. Other authors also find an extremely high metallicity, including \\cite{Castro1997} and \\cite{Cenarro07}, who both find \\feh\\,=\\,+0.55. These measurements, along with the detection of HD\\,126614\\,Ab, add statistical weight to the strong positive correlation between giant planet occurrence and metallicity \\citep{Fischer05}. Indeed, HD\\,126614\\,A has been part of the planet-metallicity correlation story for some time. In an early paper discussing the host star properties of some of the first extrasolar planets, \\cite{Gonzalez99} suggested that two bright, high-metallicity stars, namely HD\\,99109 and HD\\,126614, should be searched for Doppler variations. HD\\,99109 is known to host a planet with minimum mass $M$\\,sin\\,$i$\\,=\\,0.5\\,\\mjup and orbital period $P$\\,=\\,1.2\\,yr \\citep{Butler06}, and HD\\,126614\\,A now joins the list of high-metallity stars with planets. In addition to the planet orbiting HD\\,126614\\,A, we detected a faint M dwarf companion using adaptive optics and the PHARO camera at Palomar Observatory. This previously undiscovered star, HD\\,126614\\,B, has an estimated mass of 0.32\\,M$_{\\odot}$ and is separated from HD\\,126614\\,A by 489\\,$\\pm$\\,1.9\\,mas at position angle 56.1\\,$\\pm$\\,0.3\\,deg. This corresponds to a projected separation of 33\\,AU. HD\\,13931\\,b is reminiscent of Jupiter in orbital period ($P$\\,=\\,11.5\\,yr), eccentricity ($e$\\,=\\,0.02), and to a lesser extent mass ($M$\\,sin\\,$i$\\,=\\,1.88\\,\\mjupe). The host star, HD\\,13931, is also similar to the Sun in mass ($M_{\\star}$\\,=\\,1.02\\,$M_{\\sun}$) and metallicity (\\feh\\,=\\,+0.03). HD\\,13931\\,b is one of only 4 known RV-detected planets with orbital periods longer than 10\\,yr. The other such planets---55\\,Cnc\\,d \\citep{Fischer08}, HD\\,217107\\,c \\citep{Vogt05,Wright09} and HD\\,187123\\,c \\citep{Wright07,Wright09}---are all in multi-planet systems. Gl\\,179\\,b is a Jovian-mass ($M$\\,sin\\,$i$\\,=\\,0.82\\,\\mjupe) planet in a long-period ($P$\\,=\\,6.3\\,yr) orbit. The host star, Gl\\,179, is one of only $\\sim$10 M dwarfs currently known to host a planet and is among the faintest ($V$\\,=\\,11.96) stars with a planet discovered by RV measurements. This planet is detected in the Keck velocities alone, but without the HET measurements, the orbital parameters, especially eccentricity and minimum mass, would be determined more poorly. We note with interest that Gl\\,179 is an almost identical twin to Gl\\,876, an M dwarf known to host two Jovian planets locked in resonant orbits and a super-Earth in a $P$\\,=\\,1.9\\,d orbit \\citep{Marcy_876,Rivera05}. The stars are similar in effective temperature, mass, and age, as traced respectively by $V-K$ color (5.00 for Gl\\,179 and 5.15 for Gl\\,876), $M_K$ (6.49 for Gl\\,179 and 6.67 for Gl\\,876), and $S_\\mathrm{HK}$ (0.96 for Gl\\,179 and 1.02 for Gl\\,876). The high metallicity of Gl\\,179 is also strikingly similar to the metallicity of Gl\\,876. \\cite{Johnjohn09a} estimate \\feh\\ = $+$0.3 and $+$0.37 for Gl\\,179 and Gl\\,876, respectively. Based on their analysis of M dwarfs with planets, Johnson \\& Apps find that the planet-metallicity correlation holds for stars at the bottom of the main-sequence. Gl\\,179 and its planet add statistical weight to this finding. The planet-bearing stars presented here are good candidates for follow-up astrometric and direct imaging observations. The astrometric perturbations from these planets on their host stars (10's to 100's of $\\mu$as) will be quite detectable with the expected sub-$\\mu$as sensitivity of NASA's Space Interferometry Mission (SIM) or its proposed variants. Direct detection is also plausible using coronagraphs/interferometers on space-borne and next-generation ground-based observatories, including GPI and SPHERE. HD\\,13931\\,b may be the best candidate with a maximum projected angular separation of 120\\,mas. Indeed, SIM observations of HD\\,13931, even over a fraction of the orbital period (SIM has a planned mission duration of 5 years), combined with the RV measurements presented here, would completely determine the 3-dimensional orbit, giving accurate predictions of the best times for direct imaging observations." }, "1003/1003.3677_arXiv.txt": { "abstract": " ", "introduction": "The minimal supersymmetric extension of the Standard Model (MSSM) has over 100 free parameters, most of them associated with the breaking of supersymmetry~\\cite{mssm,Baer:book,Drees:book}. Three classes of soft supersymmetry-breaking parameters are generally considered, the scalar masses $m_0$, the gaugino masses $m_{1/2}$ and the trilinear scalar couplings $A_0$, which are often assumed to be universal at some high input scale. Universality before renormalization of the $m_0$ parameters for different sfermions with the same electroweak quantum numbers is motivated by the upper limits on flavour-changing neutral interactions, and specific Grand Unified Theories (GUTs) suggest universality relations between squarks and sleptons~\\cite{EN}. Simple GUTs also favour universality before renormalization for the gaugino masses $m_{1/2}$, and universality is also a property of minimal supergravity (mSUGRA) models, which, however, also predict additional relations that we do not discuss here~\\cite{BIM,bfs,vcmssm}. We refer to the scenario with universal $m_0$, $m_{1/2}$ and $A_0$ as the constrained MSSM (CMSSM), and its parameter space has been explored extensively~\\cite{funnel,cmssm,efgosi,cmssmnew,cmssmmap,like1,like2}. There is, however, one important question: at what renormalization scale $M_{in}$ are $m_0$ and $m_{1/2}$ actually universal? The obvious possibility, and the one has been studied most frequently, is that universality applies at the same GUT scale, $\\mgut$, as coupling constant universality. In this case, the density of cold dark matter (assumed here to be composed mainly of the lightest neutralino, $\\schi_1$, hereafter called $\\chi$)~\\cite{EHNOS} is larger than the range favoured by WMAP~\\cite{WMAP} and other experiments in generic regions of the $(m_{1/2}, m_0)$ plane, and is compatible with WMAP only in narrow strips that are either close to the boundary where $\\chi$ ceases to be the lightest sparticle -- the stau~\\cite{stauco} or stop~\\cite{dm:stop} coannihilation strips -- or close to the LEP2 chargino bound~\\cite{LEPsusy} -- the bulk region~\\cite{EHNOS,cmssm} -- or where there is no electroweak symmetry breaking -- the focus-point region~\\cite{focus} -- or in a rapid-annihilation funnel (or A-funnel)~\\cite{funnel}. However, it is not necessarily the case that $M_{in} = \\mgut$, since supersymmetry breaking might arise at some scale either below or even above $\\mgut$, and both possibilities have been studied in the literature. For example, as $M_{in}$ is decreased below $\\mgut$, the differences between the renormalized sparticle masses diminish and the regions of the $(m_{1/2}, m_0)$ planes that yield the appropriate density of cold dark matter move away from the boundaries~\\cite{EOS06}. Eventually, for small $M_{in}$, the coannihilation and focus-point regions of the conventional GUT-scale CMSSM merge. Finally, for very small $M_{in}$ they disappear entirely, and the relic $\\chi$ density falls before the WMAP range everywhere in the $(m_{1/2}, m_0)$ plane, except for very large values of $m_{1/2}$ and $m_0$. What happens to the supersymmetric parameter space and sparticle phenomenology when $M_{in} > \\mgut \\sim 2 \\times 10^{16}$~GeV? Here, we consider values of $M_{in}$ ranging up to the reduced Planck mass $\\mplr \\equiv \\mpl /\\sqrt{2\\pi} \\sim 2.4 \\times 10^{18}$~GeV. Generically, increasing $M_{in}$ increases the renormalization of the sparticle masses which tends in turn to increase the splittings between the physical sparticle masses~\\cite{pp}. As we discuss in more detail below, this in turn has the effect of increasing the relic density in much of the $(m_{1/2}, m_0)$ plane. As a consequence, the coannihilation strip is squeezed to lower values of $m_{1/2}$~\\footnote{For a previous example of this phenomenon, see Fig. 5 of~\\cite{Calibbi} or Fig.~3 of~\\cite{CEGLR}.}, particularly for $\\tan \\beta \\sim 10$, and even disappears as $M_{in}$ increases. At the same time, the focus-point strip often moves out to ever larger values of $m_0$. There are also changes in the impacts of important constraints such as $g_\\mu - 2, b \\to s \\gamma$ and $m_h$, which we also discuss below. The general conclusion is that the supersymmetric landscape would look rather different for $M_{in} > 10^{17}$~GeV from the CMSSM in which the universality scale $M_{in} = \\mgut$. The allowed region of parameter space that survives longest is the rapid-annihilation funnel at large $m_{1/2}$ and $\\tan \\beta$, which is compatible with the $m_h$ and $b \\to s \\gamma$ constraints. In the CMSSM, the funnel region also requires large $m_0$ and would make a contribution to $g_\\mu - 2$ that is too small to explain the experimental discrepancy with Standard Model calculations based on low-energy $e^+ e^-$ data. However, as we shall show, for large $M_{in}$, the funnel region extends to low $m_0$ (including $m_0 = 0$) and in some cases will be compatible with the $g_\\mu - 2$ measurements. We underline that there are some potential ambiguities in these conclusions. We use for our analysis above $\\mgut$ the particle content and the renormalization-group equations (RGEs) of the minimal SU(5) GUT~\\cite{pp,others}, primarily for simplicity and so as to minimize the number of additional parameters to be explored: for a recent review of this sample model and its compatibility with experiment, see~\\cite{Senjanovic:2009kr}. Even in this simplest GUT, there are two couplings in the SU(5) superpotential that make potentially significant contributions to the RGE running but are poorly constrained. We explore their impacts on our results, and find that one of the couplings could have noticeable effects in the focus-point region, if it is large. Secondly, we are aware that the minimal SU(5) model is surely inadequate; for example, it does not include neutrino masses. The effects of a minimal seesaw sector on the GUT RGEs~\\cite{casa} and on the relic density~\\cite{Calibbi,CEGLR,kov} has been considered elsewhere: they also may be small if the neutrino Dirac Yukawa couplings are not large. However, the minimal SU(5) GUT also has issues with proton decay via dimension-five operators, which may be alleviated if the GUT triplet Higgs particles are relatively heavy, as would happen if the associated SU(5) superpotential coupling were large~\\cite{protdec}. We therefore consider this option as the default in our analysis. One could in principle consider non-minimal GUT models in which dimension-five proton decay is suppressed by some other mechanism, but the exploration of such models would take us too far from our objective here. The layout of this paper is as follows. In Section~\\ref{sec:su5}, we recall the superpotential of the minimal SU(5) GUT and the corresponding RGEs for the soft supersymmetry-breaking parameters. Here, we give some examples of the RGE running, assuming universality at some high scale $M_{in} > \\mgut$. We also give examples of the dependences of physical sparticle masses on $M_{in}$, illustrating features that are important for understanding qualitatively the dependences of features in the $(m_{1/2}, m_0)$ plane. Several of these are displayed in Section~\\ref{sec:plane}, for representative values of $M_{in}$ and default values of the unknown SU(5) superpotential parameters. As already mentioned, two of the striking features in these planes are the disappearance of the coannihilation strip and the movement of the focus-point strip to larger $m_0$ as $M_{in}$ increases. We discuss in Section~\\ref{sec:tanb} the sensitivities of these features to the choices of the SU(5) superpotential parameters, showing that the disappearance of the coannihilation strip is relatively model-independent, whereas the movement of the focus-point strip is more model-dependent. Finally, in Section~\\ref{sec:concl} we summarize our results and draw some conclusions for the generalization of the standard CMSSM with $M_{in} = \\mgut$ to different values of $M_{in}$. ", "conclusions": "{tanbeta}} \\label{sec:tanb} As we have discussed above, the $\\stau$ coannihilation region tends to disappear as $M_{in}$ increases, and the focus-point region tends to recede to larger $m_0$ as $\\lambda$ increases. So far, we have shown these effects only for $\\tan \\beta = 10$ and 55. Here we summarize how these effects vary for intermediate values of $\\tan \\beta$. We see in the left panel of Fig.~\\ref{fig:Summary} the region of the $(M_{in}, \\tan \\beta)$ plane where there is a coannihilation/rapid-annihilation strip. For choices lying below the contour, all the points where coannihilation or rapid-annihilation brings the relic density into the WMAP range are excluded by other constraints. For $\\tan \\beta \\lappeq 20$, the region below the curve has $m_h$ lower than the LEP bound, while for larger $\\tan \\beta$, $g_\\mu -2$ is too large in this region~\\footnote{The kink in the boundary contour is an artifact of our approximation to the $m_h$ constraint: incorporating the uncertainty in the theoretical calculation of $m_h$ and the experimental likelihood function would smooth it out. Recall that for very low $\\tan \\beta$, there is some tension between the Higgs mass and $g_\\mu -2$ constraints.}. We see that, as $\\tan \\beta$ increases above 10, the coannihilation/rapid-annihilation strip persists up to progressively larger values of $M_{in}$. However, only for $\\tan \\beta > 47$ does it persist for $M_{in} = \\overline{M_P}$. This information is potentially a useful diagnostic tool, if supersymmetry is discovered. For example, if experiments determine that $m_0$ is (essentially) universal and $m_0 \\ll m_{1/2}$ with $\\tan \\beta \\sim 20$, then we can infer from the left panel of Fig.~\\ref{fig:Summary} that $M_{in} < 10^{17}$~GeV. \\begin{figure}[ht] \\epsfig{file=sum1ns.eps,height=8.0cm} \\epsfig{file=QvM8pfps.eps,height=8.0cm} \\caption{\\it Left: the coannihilation/rapid-annihilation strips are compatible with other experimental constraints only for values of $(M_{in}, \\tan \\beta)$ above the diagonal contour. Right: the focus-point strip has $m_0 < 5$~TeV at $m_{1/2} = 300$~GeV only for values of $M_{in}, \\lambda$ below the red (blue) line for $\\tan \\beta = 10 (50)$.} \\label{fig:Summary} \\end{figure} We have also seen earlier that the focus-point region is sensitive to the value of $\\lambda$. The right panel of Fig.~\\ref{fig:Summary} displays the regions of the $(M_{in}, \\lambda)$ plane where the focus-point strip has $m_0 < 5$~TeV for $m_{1/2} = 300$~GeV. We see immediately that these regions are rather similar for $\\tan \\beta = 10$ and 50, lying below the red and blue lines, respectively. This information could be used to infer a constraint on $\\lambda$, which otherwise does not impact significantly low-energy phenomenology. For example, if experiment indicates that Nature is described by a focus-point model with $m_{1/2} = 300$~GeV, $m_0 < 5$~TeV and $M_{in} > 10^{17}$~GeV, then we can infer from the right panel of Fig.~\\ref{fig:Summary} that $\\lambda < 0.6$." }, "1003/1003.1358_arXiv.txt": { "abstract": "The normalized radial distribution of young stellar populations (and cold gas) in nearby galactic disks is compared between AGN host galaxies and starforming galaxies (both with Hubble types between S0/a and Scd) by using type II supernovae (SNe) as tracers. A subset of 140 SNe\\,II with available supernova position measurements are selected from the SAI-SDSS image catalog by requiring available SDSS spectroscopy data of their host galaxies. Our sample is finally composed of 46 AGNs and 94 starforming galaxies. Both directly measured number distributions and inferred surface density distributions indicate that a) the SNe detected in starforming galaxies follow an exponential law well; b) by contrast, the SNe detected in AGN host galaxies significantly deviate from an exponential law, which is independent of both morphological type and redshift. Specifically, we find a detection deficit around $R_{\\mathrm{SN}}/R_{25,\\mathrm{cor}}\\sim0.5$ and an over-detection at outer region $R_{\\mathrm{SN}}/R_{25,\\mathrm{cor}}\\sim0.6-0.8$. This finding provides a piece of evidence supporting that there is a link between ongoing star formation (and cold gas reservoir) taking place in the extended disk and central AGN activity. ", "introduction": "It is now generally believed that active galactic nuclei (AGNs) play an important role in galaxy formation and evolution. The growth of the central supermassive black hole (SMBH) is suggested to be related to the formation of the bulge of the host galaxy where the SMBH resides. This evolutionary scenario is supported by the well-established Magorrian relationship (e.g., Magorrian et al. 1998; Tremaine et al. 2002; Ferrarese et al. 2006), and by the fact that both star formation and AGN activity show similar evolutions from z$\\sim1$ to the current epoch (e.g., Ueda et al. 2003; Silverman et al. 2008). So far, two kinds of mechanisms have been proposed to explain the co-evolution of AGNs and their host galaxies. One possible mechanism is that both AGN activity and formation of the bulge are triggered by a merger of two gas rich galaxies (e.g., Granato et al. 2004; Springel et al. 2005; Hopkins et al. 2005). Reichard et al. (2009) recently found that more active AGNs with younger circumnuclear stellar populations are on average associated with more lopsided host galaxies. An alternative possibility is the gas inflow caused by the large scale gravitational asymmetry of the host galaxies, such as a bar structure. Both mechanisms can produce an inflow of gas by transporting the angular momentum out of the gas. The falling gas not only forms stars at the central region, but also fuels the central SMBH. The feedback of AGNs onto their host galaxies will likely regulate the growth of the bulge by heating and expelling the surrounding gas through strong radio jets or other AGN-driven outflows (e.g., Croton et al. 2006; Hopkins \\& Hernquist 2006; Di Matteo et al. 2005). The distribution of cold gas in AGN host galaxies is therefore crucial to the study of the co-evolution issue. In addition to directly detecting the gas distribution by the HI line emission, the gas distribution can be approximately (and reasonably) substituted by the spatial distribution of young stellar populations. Although young stellar populations are frequently identified in the host galaxies of some local AGNs (e.g., Cid Fernandes et al. 2001; Gonzalez Delgado et al. 2001; Zhou et al. 2005; Wang et al. 2004; Wang \\& Wei 2006; Mao et al. 2009), their spatial distribution in the host galaxies is still poorly understood. Combining the GALEX near-UV survey with the SDSS survey, Kauffmann et al. (2007) recently found that in the local universe, although the AGN activity is strongly correlated with the age of the stars in the bulges (see also in Wang \\& Wei 2008; Kewley et al. 2006; Wild et al. 2007), the most active AGNs are always associated with the bluest outer disks. However, not all the galaxies with blue outer disks have an active AGN. This result therefore motivates the authors to believe that it could be understood if the amount of gas transported inward from disk is a variable. In this paper, we investigate the co-evolution issue by comparing the radial distribution of the core-collapse supernovae (cc-SNe) detected in AGN host galaxies with the similar distribution of the cc-SNe detected in starforming galaxies. Because cc-SNe are generally accepted to be produced by the explosion of massive stars ($\\geq8-10M_\\odot$) at the end of their lifetime $\\la 10^{7.5}$yr (e.g., Woosley et al. 2002), the radial distribution of cc-SNe reasonably represents the distribution of young stellar populations. The advantage of this approach is that the result does not strongly depend on the spatial resolution of the observations. ", "conclusions": "Because SNe\\,II are generated from the explosion of massive stars ($\\geq8M_\\odot$), the SNe\\,II radial distribution in their host galaxies reasonably reflects not only the radial distribution of young stellar populations, but also the radial distribution of cold gas, assuming a uniform supernova rate. By comparing the radial distributions of the SNe\\,II detected in AGNs and the similar distribution of the SNe\\,II detected in starforming galaxies, we find that the supernova radial distribution in AGN host galaxies deviates greatly from the exponential model that can describe the radial distribution in starforming galaxies well. Both directly measured number distribution and inferred surface density indicate that SNe detected in AGN host galaxies show a bimodal distribution as a function of radius. The comparison of the radial distributions of the SNe\\,II detected in the two types of supernova host galaxy allows us to argue that the existence of the AGN activities is connected with the gas reservoir located in the extended galactic disk, which agrees with the previous studies. Kauffmann et al. (2007) identified a UV-light excess in the extended disk for local AGN host galaxies. Hunt et al. (1999) suggested that the AGN host galaxies show a larger gas fraction in their disk than non-active galaxies. Using deep imaging from Spitzer and GALEX, Zheng et al. (2009) suggested that the star formation in massive galaxies at $z<1$ mainly takes place in the isolated disks. Moreover, stellar rings in the disks are more frequently identified in AGN host galaxies than in starforming galaxies (Hunt \\& Malkan 1999). By examining the spatially resolved stellar populations of 8 AGNs at $z\\sim1$, Ammons et al. (2009) arrived at a conclusion that the strong type II AGNs are associated with extended star formation activities. The star formation occurring in non-active galaxies and star formation associated with SMBH accretion appear to be different events with different origin. The radial distribution of the SNe\\,II in the starforming galaxies can be well modelled by an exponential model, which means the gas distribution in these galactic disks is not significantly disturbed. Some particular dynamical mechanisms are necessary in AGN host galaxies to redistribute gas to trigger both large-scale star formation occurring in the outer disk and central SMBH activity (and also associated circumnuclear star formation). So far, several mechanisms have been proposed to link the central AGNs with the outer parts of their host galaxies. Kauffmann et al. (2007) proposed that the gas distributed in the outer disk of AGN host galaxies could stem from the accretion of gas from an external source. In addition, the redistribution of gas could be resulted from minor merger or major merger of two galaxies (e.g., Martini 2004 and references therein). Recent numerical simulations indicated that the stars and gas could survive to re-form a disk in the merger of two gas rich galaxies (e.g., Hopkins et al. 2009; Hammer et al. 2005). In the merger process, the gas within some characteristic radius loses its angular momentum quickly, and sinks into the galactic central region by the gravitational attraction. The gas that survives outside of the characteristic radius will descend to form a new disk if the strong AGN feedback is taken into account. The multiple disks produced by interactions are indeed observed in individual local Seyfert galaxies, e.g., Mark 315, a Seyfert 1.5 galaxy (Ciroi et al. 2005). Reichard et al. (2009) recently reported a connection between AGN activity and lopsidedness of their host galaxies. An outer loop and an arc with blue colors were observed in Seyfert 1.8 galaxy Mark 334 at an radius $r\\sim$20-30\\arcsec from the center (Smirnova \\& Moissev 2009). The blue colors suggest the existence of young stellar populations ($\\sim$0.5-1 Gyr) that are formed in the interaction process. Adopting the characteristic radius of the galaxy $R_{25}\\approx50$\\arcsec estimated from the Figure 4 in Smirnova \\& Moissev (2009), the relative distance of the outer loop and arc is inferred to be $\\sim0.4-0.6$. Besides the merger scenario, it is now generally believed that the bar-driven gas inflow is related to the formation of the gas rings (Buta \\& Combes 1996). Theoretical and N-body simulation studies indicated that the gravitational asymmetry caused by the bars transports gas angular momentum. The migration of the angular momentum results in a gas inflow within the corotation radius and an outflow of gas out of the corotation radius (e.g., Sellwood \\& Wilkinson 1993; Athanassoula 2003). The gas redistribution fuels central AGNs and circumnuclear starbursts, destroys the bars (e.g., Bournaud \\& Combes 2002), and regulates the gas into a stable configuration (e.g., rings) by itself." }, "1003/1003.0844_arXiv.txt": { "abstract": "We present theory and algorithms to perform an all-sky coherent search for periodic signals of gravitational waves in narrow-band data of a detector. Our search is based on a statistic, commonly called the $\\F$-statistic, derived from the maximum-likelihood principle in Paper I of this series. We briefly review the response of a ground-based detector to the gravitational-wave signal from a rotating neuron star and the derivation of the $\\F$-statistic. We present several algorithms to calculate efficiently this statistic. In particular our algorithms are such that one can take advantage of the speed of fast Fourier transform (FFT) in calculation of the $\\F$-statistic. We construct a grid in the parameter space such that the nodes of the grid coincide with the Fourier frequencies. We present interpolation methods that approximately convert the two integrals in the $\\F$-statistic into Fourier transforms so that the FFT algorithm can be applied in their evaluation. We have implemented our methods and algorithms into computer codes and we present results of the Monte Carlo simulations performed to test these codes. ", "introduction": "\\label{Sec:Intro} Periodic gravitational-wave signals like those originating from rotating neutron stars are an important class of sources that can be detected by currently operating ground-based detectors. Several methods were developed to search for such sources and several searches were performed. This paper continues the series of papers \\cite{JKS98,JK99,JK00,ABJK02} devoted to studies of data analysis tools and algorithms needed to perform an all-sky coherent search for quasiperiodic gravitational waves. The search presented in the current paper is based on the maximum-likelihood statistic called the $\\F$-statistic that we have derived in the Paper I \\cite{JKS98} of this series. It is known that the coherent search for long observation time needed to detect weak gravitational-wave signals from rotating neutron stars are computationally prohibitive (see \\cite{BCCS98} and Paper III of this series \\cite{JK00}). Promising strategies are hierarchical semi-coherent methods. In these methods data is broken into short segments. In the first stage each segment is analyzed using the $\\F$-statistic and in the second stage the $\\F$-statistics from the short segments are combined using a certain algorithm. There are several methods proposed for the second stage: search for coincidences among candidates from short duration segments \\cite{LIGO07Fstat,LIGO09EHa}, stack-slide method \\cite{BC99}, power flux method \\cite{LIGO08powerflux,LIGO09powerflux}, Hough transform method \\cite{PAFS97,PS99,KSPSFP04,LIGO05Hough,LIGO08powerflux}. Recently an optimal method for the second stage has been found, the {\\em global correlation coordinate} method \\cite{P08,AP09}, which exploits global parameter space correlations in the coherent detection statistic. In our paper we shall present methods to optimize the first, coherent stage of a hierarchical method. The techniques presented in this paper were used in the analysis of NAUTILUS bar detector data \\cite{Astone2008} and are presently used in the analysis of the VIRGO data. Alternative techniques for the coherent stage based on the $\\F$-statistic and their application to the real data can be found in Refs.\\ \\cite{Astone2005,LIGO07Fstat,LIGO09EHa,LIGO09EHb} The paper is organized as follows. In Sec.\\ \\ref{Sec:Pio} we present the noise-free response of a ground-based detector to a gravitational-wave signal from a rotating neutron star. This response was derived and discussed in detail in Papers I \\cite{JKS98} and IV \\cite{ABJK02} of our series. In Sec.\\ \\ref{Sec:Kro1} we present data analysis tools to perform coherent search of the data for a gravitational-wave signal given in Sec.\\ \\ref{Sec:Pio}. In Sec.\\ \\ref{sSec:Kro1a} we present the $\\F$-statistic that was derived in Paper I. We limit ourselves to the case when the observation time is an integer multiple of one sidereal day. This simplifies some general formulas considerably. In Sec.\\ \\ref{sSec:Kro1b} we introduce a simplified approximate model for a periodic gravitational-wave signal. This approximate signal has the constant amplitude and its phase is parameterized in such a way that it is a linear function of the parameters. For such a signal the Fisher matrix is constant and consequently it is independent of the values of the signal's parameters. In Sec.\\ \\ref{sSec:false} we briefly review calculation of the false alarm probability. Section \\ref{Sec:Mac} is devoted to construction of the grid of templates in the parameter space. The grid solves a certain covering problem with a constraint. Our constraint is that the nodes of the grid coincide with the Fourier frequencies. This allows to use the fast Fourier transform (FFT) algorithm to compute the $\\F$-statistic at grid nodes, what greatly accelerates the calculation. In Sec.\\ \\ref{Sec:Kaz} we describe our package \\texttt{Top2Bary} that is used to calculate the position and the velocity of the detector located on the Earth with respect to the solar system barycenter. In Sec.\\ \\ref{sSec:Kaza} we introduce various concepts and definitions used in the astrometry and in Sec.\\ \\ref{sSec:Kazb} we describe the content of our package which is a set of \\textsc{fortran} routines. In Sec.\\ \\ref{Sec:Kro} we present various approximations that we use in the calculation of the $\\F$-statistic in order to speed up computations. In Sec.\\ VI~A we discuss resampling of the time series to the barycenter that we need to perform before we can apply the FFT. We develop two algorithms: one slow and very accurate and the other fast but less accurate. We compare the two algorithms using the signal from Sec.\\ \\ref{Sec:Pio}. In Sec.\\ VI~B we describe interpolation of the FFT in the Fourier domain. This interpolation method allows to obtain efficiently an FFT that is twice as fine as the FFT of original data. In Sec.\\ VI~C we describe the Nelder-Mead algorithm that we use to find accurately the maximum of the $\\F$-statistic. In Sec.\\ \\ref{Sec:Kro3} we perform a number of Monte Carlo simulations of the computer code where we have implemented the methods and algorithms from Secs.\\ III--VI. In our simulations we investigate how well we estimate the parameters of the signal in comparison to the Cram\\'er-Rao bound. ", "conclusions": "" }, "1003/1003.0769_arXiv.txt": { "abstract": "An explicit model of $F(R)$ gravity with realizing a crossing of the phantom divide is reconstructed. In particular, it is shown that the Big Rip singularity may appear in the reconstructed model of $F(R)$ gravity. Such a Big Rip singularity could be avoided by adding $R^2$ term or non-singular viable $F(R)$ theory\\cite{Nojiri:2009xw} to the model because phantom behavior becomes transient. ", "introduction": "It is observationally supported that the current expansion of the universe is accelerating.\\cite{Spergel:2003cb,Peiris:2003ff,Spergel:2006hy,Komatsu:2008hk,Perlmutter:1998np,Riess:1998cb,Astier:2005qq,Riess:2006fw} A number of scenarios to account for the current accelerated expansion of the universe have been proposed (for reviews, see Refs.~\\refcite{Peebles:2002gy,Sahni:2005ct,Padmanabhan:2002ji,Copeland:2006wr,review,Nojiri:2008nk,N-O-2,S-F,Lobo-08,C-F}). Approaches to explain the current accelerated expansion of the universe fall into two broad categories. One is the introduction of some unknown matter, which is called ``dark energy'' in the framework of general relativity. The other is the modification of the gravitational theory, e.g., ``$F(R)$ gravity'', where $F(R)$ is an arbitrary function of the scalar curvature $R$ (for reviews, see Refs.~\\refcite{review,Nojiri:2008nk,N-O-2,S-F,Lobo-08,C-F}). Recent various observational data\\cite{Alam:2004jy,Nesseris:2006er,Wu:2006bb,J-B-P} imply that the effective equation of state (EoS), which is the ratio of the effective pressure of the universe to the effective energy density of it, may evolve from larger than $-1$ (non-phantom phase) to less than $-1$ (phantom one), namely, cross $-1$ (the phantom divide). Various investigations to realize the crossing of the phantom divide have been executed in the framework of general relativity: Scalar-tensor theories with the non-minimal gravitational coupling between a scalar field and the scalar curvature or that between a scalar field and the Gauss-Bonnet term, one scalar field model with non-linear kinetic terms or a non-linear higher-derivative one, phantom coupled to dark matter with an appropriate coupling, the thermodynamical inhomogeneous dark energy model, multiple kinetic k-essence, multi-field models (two scalar fields model, ``quintom'' consisting of phantom and canonical scalar fields), and the description of those models through the Parameterized Post-Friedmann approach, or a classical Dirac field or string-inspired models, non-local gravity, a model in loop quantum cosmology and a general consideration of the crossing of the phantom divide (for a detailed review, see Ref.~\\refcite{Copeland:2006wr}). However, explicit models of modified gravity realizing the crossing of the phantom divide have hardly been examined, although there were suggestive and interesting related works.\\cite{review,abdalla,Amendola:2007nt} In the present paper, we review our results in Ref.~\\refcite{Bamba:2008hq} and reconstruct an explicit model of $F(R)$ gravity in which a crossing of the phantom divide can be realized by using the reconstruction method proposed in Refs.~\\refcite{Nojiri:2006gh,Nojiri:2006be} (for more detailed references, see references in Refs.~\\refcite{Bamba:2008hq,Bamba:2009ay,Bamba:2009kc,Bamba:2009vq,Bamba:2009dk}). It is demonstrated that the Big Rip singularity may appear in the reconstructed model of $F(R)$ gravity. ", "conclusions": "We have studied a crossing of the phantom divide in $F(R)$ gravity. We have reconstructed an explicit model of $F(R)$ gravity in which a crossing of the phantom divide can occur by using the reconstruction method.\\cite{Nojiri:2006gh,Nojiri:2006be} As a result, we have shown that the Big Rip singularity may appear in the reconstructed model of $F(R)$ gravity. We finally mention that by adding $R^2$ term (as it was first proposed in Ref.~\\refcite{abdalla}) to the model or by adding non-singular theory,\\cite{Nojiri:2009xw} $R^2 \\left(R^n+c_1\\right)/\\left(R^n+c_2\\right)$, where $n$, $c_1$ and $c_2$ are constants, Big Rip singularities could be avoided because phantom behavior becomes transient.\\cite{B-N-O,N-O-PRD78-046006-08}" }, "1003/1003.0419_arXiv.txt": { "abstract": "We consider models in which a dark-matter particle decays to a slightly less massive daughter particle and a noninteracting massless particle. The decay gives the daughter particle a small velocity kick. Self-gravitating dark-matter halos that have a virial velocity smaller than this velocity kick may be disrupted by these particle decays, while those with larger virial velocities will be heated. We use numerical simulations to follow the detailed evolution of the total mass and density profile of self-gravitating systems composed of particles that undergo such velocity kicks as a function of the kick speed (relative to the virial velocity) and the decay time (relative to the dynamical time). We show how these decays will affect the halo mass-concentration relation and mass function. Using measurements of the halo mass-concentration relation and galaxy-cluster mass function to constrain the lifetime--kick-velocity parameter space for decaying dark matter, we find roughly that the observations rule out the combination of kick velocities greater than 100~km~s$^{-1}$ and decay times less than a few times the age of the Universe. ", "introduction": "\\label{sec:intro} There is a good consensus from many types of observations that dark matter makes up $\\sim 25\\%$ of the mass-energy density of the Universe \\cite{kessler2009,breid2009,vikhlinin2009b,komatsu2010,rozo2010}. However, the nature of dark matter is unknown. The most popular class of candidates is the weakly interacting massive particle (WIMP), since such particles may be produced thermally in the early Universe at the right abundance and with behavior consistent with a large set of cosmological observations \\cite{jungman1996,appelquist2001,cheng2002,servant2002,hubisz2005}. The canonical WIMP is electrically neutral, and it is stable. Once produced in the early Universe, it interacts with itself and with ordinary matter only gravitationally. If the WIMP is the dark matter, then the dark halos of galaxies and galaxy clusters are described by a collisionless gas of self-gravitating WIMPs. However, there is both theoretical and observational room for other types of dark-matter candidates. While observations are \\emph{consistent} with WIMPs, the observations do not \\emph{require} the dark matter to be WIMPs. Moreover, some observations---e.g., of the inner mass distribution in galaxies and/or the abundance subhalos in dark-matter halos \\cite{moore1999a}---have inspired theoretical searches for dark-matter candidates with properties beyond those of the canonical WIMP. There is a large amount of literature on dark-matter particles with additional physical properties beyond those of the canonical stable collisionless WIMP. Examples include (but are by no means limited to) particles with small electric charges \\cite{Davidson:2000hf} or dipoles \\cite{Sigurdson:2004zp}, self-interacting particles \\cite{spergel2000}, or particles with long-range forces \\cite{Gubser:2004uh,Kesden:2006vz,Kesden:2006zb,ackerman2009,feng2009} (see Refs.~\\cite{D'Amico:2009,feng2010} for recent reviews). In the spirit of these lines of investigation, we consider in this work a neutral dark-matter particle $X$ of mass $M_X$ that decays with lifetime $\\tau$ to a slightly less massive neutral particle $Y$ of mass $M_Y =M_X(1-\\epsilon)$ with $\\epsilon \\ll1$ and an effectively massless particle $\\zeta$ which is itself assumed to be noninteracting. We imagine that such a scenario may arise in some implementations of models of inelastic dark matter \\cite{TuckerSmith:2001hy}. When the $X$ particle decays, the daughter $Y$ particle receives a nonrelativistic velocity kick $\\vke \\simeq \\epsilon c$. Now suppose that these particles make up a self-gravitating halo of virial velocity $\\vvire$. If $\\vke \\ll \\vvire$, then the halo may be heated slightly, and its mass distribution thus rearranged slightly, by the velocity kicks imparted to the $Y$ particles once most of the $X$ particles have decayed. If, on the other hand, $\\vke \\gg \\vvire$, then these halos will be completely disrupted after most of the $X$ particles have decayed. There should thus be no halos with $\\vvire \\ll \\vke$ if these particles decay with lifetimes $\\tau$ small compared with the age $t_\\mathrm{H}$ of the Universe, and this has been postulated as a possible explanation for the low number of dwarf galaxies relative to that expected from dissipationless dark-matter simulations \\cite{sanchez2003,abdelqader2008}. Alternatively, if $\\vke \\gg \\vvire$ but $\\tau \\gg t_\\mathrm{H}$, then only a small fraction ($\\sim \\tau/t_\\mathrm{H}$) of the halo particles will be kicked out of the halo. The resulting halo mass and mass distribution may then be affected slightly without being completely disrupted. The canonical-WIMP halo is, of course, recovered in the limits $\\tau \\to \\infty$ and/or $\\epsilon \\to 0$. In this paper, we report on simulations of self-gravitating halos with decaying particles. While the results of these decays can be understood in the limits $\\tau\\gtrsim t_\\mathrm{H}$ and $\\vke\\gg \\vvire$ \\cite{peter2010} using the adiabatic-expansion model \\cite{flores1986,cen2001}, detailed evolution of the halo over the full range of the $\\tau$-$\\vke$ parameter space requires numerical simulation. We use the simulations and observations of the galaxy-cluster mass function and the halo mass-concentration relation to constrain the decay parameter space. Our central results are presented in Fig.~\\ref{fig:summary}, which shows that (roughly) the combination of decay times less than a few times the age of the Universe and kick velocities greater than 100~km~s$^{-1}$ are ruled out. The outline of the paper is as follows. In Sec.~\\ref{sec:sims}, we describe our simulations. In Sec.~\\ref{sec:results}, we characterize the evolution of dark-matter halos as a function of the decay parameters using the simulations, and show how the simulations, in conjunction with observational and theoretical (in the context of WIMPs) determinations of the cluster mass function and the mass-concentration relation, allow us to constrain the decay parameter space. In Sec.~\\ref{sec:discussion}, we discuss various aspects of our findings, including resolution effects (with an eye towards the requirements for cosmological simulations) and observational biases. We summarize our work in Sec.~\\ref{sec:conclusion}. ", "conclusions": "\\label{sec:conclusion} In this work, we have used simulations of two-body dark-matter decay in isolated halos initially in equilibrium to characterize the response of halos as a function of the center-of-mass kick speed \\vk, the decay time $\\tau$, and time. We find that for the union of the parameter space $\\vke/\\vvire \\lesssim 0.2$ and $\\tau/\\tdyne \\gtrsim 100$, dark-matter halos are essentially unchanged; the mass loss is $\\lesssim 10\\%$ over a Hubble time, and the shape of the density profile does not change, although the concentration declines somewhat for $\\vke/\\vvire > 1$. For the union of $\\vke/\\vvire \\gtrsim 1$ and $\\tau/\\tdyne \\lesssim 10$, there is severe mass loss and a radical change to the shape of the dark-matter density profile. For the rest of the decaying-dark-matter parameter space, there is moderate mass loss and significant deviations of the shape of the dark-matter density profile relative to the initial halo properties. We used our simulations in conjunction with the observed mass-concentration relation and the galaxy-cluster mass function to constrain the two-parameter decay parameter space. We find that the mass-concentration relation yields the stronger constraint on $\\tau$ for a broad range of \\vk~, constraining $\\tau \\gtrsim 30-40$ Gyr for $\\vke \\gtrsim 200 \\hbox{ km s}^{-1}$. Constraints on the decay parameters should improve greatly in the next 5-10 years with blossoming wide-field Sunyaev-Zel'dovich and optical surveys, and with cosmological simulations of decaying dark matter. \\newline \\vskip 1.5cm" }, "1003/1003.2136_arXiv.txt": { "abstract": "The third US Naval Observatory (USNO) CCD Astrograph Catalog, UCAC3 was released at the IAU General Assembly on 2009 August 10. It is the first all-sky release in this series and contains just over 100 million objects, about 95 million of them with proper motions, covering about R = 8 to 16 magnitudes. Current epoch positions are obtained from the observations with the 20 cm aperture USNO Astrograph's ``red lens\", equipped with a 4k by 4k CCD. Proper motions are derived by combining these observations with over 140 ground- and space-based catalogs, including Hipparcos/Tycho and the AC2000.2, as well as unpublished measures of over 5000 plates from other astrographs. For most of the faint stars in the Southern Hemisphere the Yale/San Juan first epoch plates from the SPM program (YSJ1) form the basis for proper motions. These data are supplemented by all-sky Schmidt plate survey astrometry and photometry obtained from the SuperCOSMOS project, as well as 2MASS near-IR photometry. Major differences of UCAC3 data as compared to UCAC2 include a completely new raw data reduction with improved control over systematic errors in positions, significantly improved photometry, slightly deeper limiting magnitude, coverage of the north pole region, greater completeness by inclusion of double stars and weak detections. This of course leads to a catalog which is not as ``clean\" as UCAC2 and problem areas are outlined for the user in this paper. The positional accuracy of stars in UCAC3 is about 15 to 100 mas per coordinate, depending on magnitude, while the errors in proper motions range from 1 to 10 mas/yr depending on magnitude and observing history, with a significant improvement over UCAC2 achieved due to the re-reduced SPM data and inclusion of more astrograph plate data unavailable at the time of UCAC2. ", "introduction": "The US Naval Observatory (USNO) operated the 8-inch (0.2 m) Twin Astrograph from 1998 to 2004 for an all-sky astrometric survey. About 2/3 of the sky was observed from the Cerro Tololo Inter-American Observatory (CTIO) while the rest of the northern sky was observed from the Naval Observatory Flagstaff Station (NOFS). The average number of completed fields per year was a factor of 2.0 larger at CTIO than at NOFS. A 4k by 4k CCD with 9 $\\mu$m pixel size was used in a single bandpass (579 to 643 nm) providing a flat field of view (FOV) of just over 1 square degree, taking advantage of only a tiny fraction of the FOV delivered by the optical system of the Twin Astrograph's ``red lens.'' A 2-fold overlap pattern of fields span the entire sky. Each field was observed with a long (about 125 sec) and a short (about 25 sec) exposure, thus each star should appear on at least 2 different CCD exposures, and stars in the mid-magnitude range (about 10 to 14) should have 4 images. UCAC3 contains just over 100 million objects, most of these are stars. It covers the magnitude range of about R = 8 to 16 (Fig.~1) with positional precision at mean epoch ranging from 15 to 100 mas, depending on magnitude (Fig.~2). Mean position errors are shown per 1/10 mag bin with stars up to magnitude 13 excluded whenever the formal error in either one of the coordinates exceeds 100 mas. For fainter stars no such outlier exclusion was adopted, which explains the discontinuity in Fig.~2 and also shows what effect such a restriction has on the derived mean formal position errors. The distribution of proper motions is shown in Fig.~3, and the proper motions errors as a function of magnitude are presented in Fig.~4. The large increase of the formal proper motion errors for stars at magnitude 8 and brighter is caused by the saturation of the CCD data with associated large, formal positional errors. The weighted mean epoch of UCAC3 data for most stars is in the range of 1980 to 2002 (Fig.~5), depending on magnitude as consequence of the observing history of stars and the positional precisions at various epochs. The released catalog is based on all applicable, regular survey field observations, excluding the CCD exposures taken on extragalactic link fields and most calibration fields. Observations of minor planets have been extracted and will be published separately from UCAC3. The released UCAC3 is a compiled catalog, similar to UCAC2. No individual epoch observations are given, nor are the pixel data publicly available at this point. The Tycho-2 catalog \\citep{tycho2} was used as reference star catalog to obtain UCAC3 positions on the Hipparcos System \\citep{hipcat}, which is the current optical realization of the International Celestial Reference Frame (ICRF). Most stars in UCAC3 have proper motions which were derived from the astrograph CCD data combined with various earlier epoch data, including all ground-based catalogs used also for the Tycho-2 project, unpublished new measurements of other astrograph plates, the Southern Proper Motion (SPM) first epoch plates, and Schmidt plate data through the SuperCOSMOS project. A final UCAC4 release is planned which will utilize the new reductions of the Northern Proper Motion (NPM) program, supplementing the SPM data, which then would allow us to derive proper motions for all UCAC stars without the use of Schmidt plate data. This goal could not be achieved for UCAC3 due to a production deadline and lack of time to complete the NPM work. The 2-Micron All-Sky Survey, 2MASS \\citep{2mass} was used extensively to analyze systematic errors of UCAC3 data and to supplement the UCAC3 catalog with near IR photometry. Optical B,R,I magnitudes were copied from the SuperCOSMOS source catalog (photographic photometry) into UCAC3 for the benefit of the users. The number of UCAC3 objects matched with various catalogs is presented in Table 1. For more details about the observational data and earlier reductions the reader is referred to the UCAC1 \\citep{ucac1} and UCAC2 \\citep{ucac2} papers. Contrary to those papers, which each describe one of the earlier releases in detail, the UCAC3 effort will be documented in a series of papers. This paper gives the introduction aiming at the user of the UCAC3 catalog, describing the released data, limitations, and comparisons to other catalogs. Technical details of the reduction process will be outlined in a paper about the new pixel processing \\citep{pxred} and a separate paper on the astrometric reductions leading to the mean positions at the CCD observing epoch \\citep{ared}. Preliminary results of these were already presented at a recent meeting \\citep{uc3aas}. The Southern Proper Motion data re-reduction will be described elsewhere (Girard et al.). Papers about double stars discovered in UCAC3 and confirmed with speckle observing, mining UCAC3 data for new high proper motion stars, and the extragalactic reference frame link of UCAC are in preparation. ", "conclusions": "UCAC3 is the first all-sky catalog of this series. From the details presented above it appears that the systematic errors of the CCD observations for UCAC3 are corrected even better than they were in UCAC2 (see for example the CTIO / NOFS overlap area, the comparison with SPM2 and PM2000). The magnitude dependent systematic errors seem to be well controlled, with the exception of the very faint end of UCAC3 (around 16th mag and fainter). Comparing Figures 13, 24, 28 and 32 suggests a systematic error in UCAC3 positions (both coordinates) as a function of magnitude for stars around magnitude 16. UCAC2, SPM2 and 2MASS agree, while differences of any of these catalogs with UCAC3 show some systematic deviations. The use of Schmidt Survey data likely caused a problem for the proper motions of faint stars (R $\\ge$ 14), even partly affecting the area covered by the new reductions of the SPM data, and particularly affecting the northern hemisphere. Although the formal errors in proper motions significantly dropped for a large number of stars as compared to UCAC2, systematic errors as function of location on Schmidt plates crept into the UCAC3 catalog, increasing the scatter when compared for example with the 2MASS catalog. A significant improvement of the photometry in UCAC3 was achieved, which is handled properly for the first time. The complete re-reduction of the pixel data also extended the limiting magnitude, providing more and fainter stars in UCAC3 than in earlier releases. At the bright end, the residuals of the Tycho-2 reference stars show some remaining magnitude equations. If we assume the internal calibrations of the UCAC3 CCD observations (utilizing the East/West flip data) are correct this indicates magnitude equations in the Tycho-2 catalog itself of about 1 to 2 mas/mag. Assuming the Tycho space-based observations are free of such errors, this indicates uncorrected errors in the order of 100 to 200 mas in the Astrographic Catalog (AC) whose average epoch is around 1900. The AC is the major ground-based catalog used to obtain the Tycho-2 proper motions. The major remaining steps to be taken to conclude this project are a) utilize overlap conditions of the regular 2-fold center-in-corner pattern of fields observed with the CCD astrograph to reduce coordinate dependent errors introduced by the reference stars, b) include re-reductions of the NPM data for proper motions in the north, eliminating the need to resort to Schmidt plate data, c) check on the extragalactic link by employing the dedicated observations in ICRF fields and their corresponding deep CCD imaging with larger telescopes, and d) fix above mentioned problems. These are the goals for UCAC4." }, "1003/1003.2300_arXiv.txt": { "abstract": "The IceCube observatory is the first cubic kilometre scale instrument in the field of high-energy neutrino astronomy and cosmic rays. In 2009, following five successful deployment seasons, IceCube consisted of 59 strings of optical modules in the South Pole ice, together with 118 air shower detectors in the IceTop surface array. The range of physics topics includes neutrino signals from astrophysical sources, dark matter, exotic particle physics, cosmic rays, and atmospheric neutrinos. The current IceCube status and selected results are described. Anticipated future developments are also discussed, in particular the Deep Core low energy subarray which was recently deployed. ", "introduction": "\\label{sec:intro} At the IceCube observatory, the clear South Pole glacial ice is exploited to detect Cherenkov light emitted by charged particles created in neutrino interactions. An array of digital optical modules (DOMs) equipped with photomultipliers extends throughout a volume of approximately $1\\,{\\rm km}^3$, with the central region more densely intrumented. Scheduled for completion in 2011, IceCube allows the detection of neutrino interactions above a few tens of GeV, with a maximum sensitivity in the TeV-PeV range. Extraterrestrial neutrinos are expected from a variety of sources. Generically, high energy cosmic ray particles will give rise to neutrinos in their interactions with matter or radiation, either in the immediate vicinity of the acceleration sites or while traveling through space. It is thus possible to estimate the neutrino flux that corresponds to the observed flux of cosmic rays~\\cite{WBbound}. Such considerations have dictated sensitivity requirements which translate into a cubic kilometre detector. Owing to their low interaction cross-section and lack of electric charge, neutrinos propagate over vast distances, quite undisturbed by matter, radiation, and magnetic fields. This is in contrast to photons, which can be absorbed by matter, and which cannot travel cosmological distances if their energy is above the TeV range, because of absorption by background radiation. Similarly, protons of the highest energies can be absorbed through the GZK mechanism~\\cite{GZK}, while at lower energies they are deflected by magnetic fields and do not point back to the source. It is therefore possible that there are ``hidden'' sources which can only be detected via their neutrino emission. In cases where emission of gamma rays or high energy cosmic rays is seen, a detection of neutrinos could elucidate the nature and location of cosmic ray accelerators. (Gamma rays from $\\pi^0$ decay would be accompanied by neutrinos from $\\pi^+/\\pi^-$ decay.) Promising candidate acceleration sites are relativistic ejecta from active galactic nuclei (AGNs) or gamma-ray bursts (GRBs). There are also objects within our galaxy, such as supernova remnants, pulsar winds, or microquasars, in which acceleration processes could yield neutrino emission. Even if individual neutrino sources are not seen, their combined {\\em diffuse} flux may stand out above the atmospheric neutrino flux. The study of the atmospheric neutrino flux is therefore an important topic. It provides a useful calibration point and may in addition yield results relevant to the intrinsic properties of neutrinos. There are abundant indications of dark matter in the form of non-relativistic weakly interacting massive particles (WIMPs)~\\cite{DMrev}. These could constitute a class of ``hidden'' sources, detectable via neutrinos, as they annihilate in the centre of the Sun or other dense regions ~\\cite{WIMPcapture,WIMPsusy}. Promising candidate WIMPs are the neutralinos of supersymmetric models~\\cite{WIMPsusy} or Kaluza-Klein excitations in models of extra dimensions~\\cite{HooperProfumoUED}. Section \\ref{sec:detstat} describes the IceCube detector and its performance, and selected analysis results are discussed in section \\ref{sec:results}. Future enhancements are discussed in section \\ref{sec:future}. ", "conclusions": "" }, "1003/1003.4710_arXiv.txt": { "abstract": "Gamma-ray bursts have the potential to produce the particle energies (up to $10^{21}$\\,eV) and the energy budget ($10^{44}\\, \\rm{erg\\, yr^{-1}\\, Mpc^{-3}}$) to accommodate the spectrum of the highest energy cosmic rays; on the other hand, there is no observational evidence that they accelerate hadrons. The Fermi GST recently observed two bursts that exhibit a power-law high-energy extension of the typical (Band) photon spectrum that extends to $\\sim 30$ GeV. On the basis of fireball phenomenology we argue that they, along with GRB941017 observed by EGRET in 1994, show indirect evidence for considerable baryon loading. Since the detection of neutrinos is the only unambiguous way to establish that GRBs accelerate protons, we use two methods to estimate the neutrino flux produced when they interact with fireball photons to produce charged pions and neutrinos. While the number of events expected from the Fermi bursts detected to date is small, we conclude that an event like GRB941017 will be detected by the IceCube neutrino telescope if gamma-ray bursts are indeed the sources of the observed cosmic rays. ", "introduction": " ", "conclusions": "" }, "1003/1003.1135_arXiv.txt": { "abstract": "{We explore the amplification of magnetic seeds during the formation of the first stars and galaxies. During gravitational collapse, turbulence is created from accretion shocks, which may act to amplify weak magnetic fields in the protostellar cloud. Numerical simulations showed that such turbulence is sub-sonic in the first star-forming minihalos, and highly supersonic in the first galaxies with virial temperatures larger than $10^4$~K. We investigate the magnetic field amplification during the collapse { both for Kolmogorov and Burgers-type} turbulence with a semi-analytic model that incorporates the effects of gravitational compression and small-scale dynamo amplification. We find that the magnetic field may be substantially amplified { before the formation of a disk. On scales of $1/10$ of the Jeans length, saturation occurs after $\\sim10^8$~yr. } Although the saturation behaviour of the small-scale dynamo is still somewhat uncertain, we expect a saturation field strength of the order $\\sim 10^{-7} n^{0.5}$~G in the first star-forming halos, with $n$ the number density in cgs units. In the first galaxies with higher turbulent velocities, the magnetic field strength may be increased by an order of magnitude{, and saturation may occur after $10^6-10^7$~yr. In the Kolmogorov case, the magnetic field strength on the integral scale (i.e. the scale with most magnetic power) is higher due to the characteristic power-law indices, but the difference is less than a factor of $2$ in the saturated phase. Our results thus indicate that the precise scaling of the turbulent velocity with length scale is of minor importance. They further imply that magnetic fields will be significantly enhanced before the formation of a protostellar disk, where they may change the fragmentation properties of the gas and the accretion rate.}% } ", "introduction": "The formation of the first stars is generally regarded as a well-defined problem, as the initial conditions at $z\\sim100$ can be derived accurately from CMB data \\citep[e.g.][]{Komatsu09} using linear theory \\citep{Bertschinger98} and the chemistry is primordial and well-understood \\citep[e.g.][]{Abel97,Galli98, Stancil98, Omukai01, Yoshida06, Schleicher08, Glover09}. In addition, it is often assumed that magnetic fields are not yet present and that the hydrodynamical equations are sufficient to describe the star formation process. This assumption is not neccessarily true. Indeed, a variety of mechanisms exist to create strong magnetic fields during inflation, the electroweak or the QCD phase transition \\citep[see e.g.][for a review]{Grasso01}. By means of the inverse-cascade, the magnetic power of these fields may have been shifted to larger scales in case of non-zero helicity \\citep{Brandenburg96, Christensson01,Banerjee04b}. Strong primordial fields would have profound implications concerning the thermodynamics of the post-recombination universe, reionization and the formation of the first stars \\citep{Sethi05, Machida06, TashiroSugiyama06a, Tashiro06b, Schleicher08b, Schleicher09a, Schleicher09prim}. In this paper, we will however explore the limiting case in which extremely weak seed fields have been produced before recombination. In such a case, the dominant contribution to the magnetic field strength comes from astrophysical processes after recombination. Cosmological MHD simulations including an approximate treatment of the Biermann battery term suggest that the Biermann battery could create seed field of the order $10^{-18}$~G in the IGM at $z=20$ \\citep{Xu08}. Additional seed fields may be created via the Weibel instability in shocks \\citep{Schlickeiser03, Medvedev04, Lazar09}. The importance of dynamos in cosmic sheets has early been recognized by \\citet{Pudritz89}. The simulations of \\citet{Xu08} run from cosmological scales to the protostellar collapse phase in a primordial minihalo. In such a situation, the following mechanisms are available to amplify the magnetic field: \\begin{itemize} \\item gravitational compression of the magnetic field, \\item { the small-scale turbulent dynamo which amplifies seed magnetic fields already generated from cosmological processes}, \\item { large-scale dynamos in protostellar and galactic disks}, \\item the magneto-rotational instability (MRI). \\end{itemize} Gravitational compression under spherical symmetry leads to an increase of the magnetic field strength with $n^{2/3}$, where $n$ denotes the number density of the gas. If the collapse proceeds preferentially along one axis, for instance because of rotation or strong magnetic fields, the scaling is closer to $n^{0.5}$. In realistic cases, often intermediate values are found \\citep{Machida06, Banerjee08}. This amplification mechanism has also been identified in the simulation of \\citet{Xu08}. Large-scale dynamos typically require the presence of a galactic or protostellar disk and act on relatively long timescales \\citep[see][for a review]{Brandenburg05}. { In such a disk, an exponential growth may also be obtained from the magnetorotational instability \\citep[MRI, see][]{Balbus91}, which may seed other large-scale dynamos with turbulence \\citep{Tan04, Silk06}. However, the length scale of the fastest growing mode decreases for decreasing field strengths, and in the presence of a viscous cutoff length, such amplification may not be possible. As a result, a minimal field strength is required to drive the MRI in a protostellar disk \\citep[see][for a detailed discussion]{Tan04, Silk06}.} A critical condition for any dynamo growth is that the ideal MHD approximation is applicable. \\citet{Maki04, Maki07} investigated this question using detailed models for magnetic energy dissipation via Ohmic and ambipolar diffusion to show that the magnetic field is frozen into the gas unless it is very strong. An approximate fit to their results yields a critical field strength of about $B\\leq10^{-5}(n/10^3\\ \\mathrm{cm}^{-3})^{0.55}$~G. Due to the subtle effects of lithium chemistry, the ionization degree does not drop exponentially at densities of $\\sim10^9$~cm$^{-3}$, but stays almost constant with increasing density. The more recent study by \\citet{Glover09} finds even higher ionization degrees at these densities. This implies that the ideal MHD approximation can be used during the collapse phase to describe the interaction of magnetic fields with matter. Deviations from this behaviour may however occur on very small scales, where ambipolar and Ohmic diffusion become increasingly important. Estimates based on the non-ideal MHD models of \\citet{Pinto08a, Pinto08b} imply that, even if the magnetic field is in equipartition with the gas, ambipolar diffusion is important only on scales 4 orders of magnitudes smaller than the Jeans length, and Ohmic dissipation occurs only on even smaller scales. As Ohmic and ambipolar diffusion depend on the field strength itself, these scales will be significantly smaller for weaker magnetic fields, so that the ideal MHD approximation can be savely applied. This implies that the magnetic Reynolds number ${\\rm Rm}=vl/\\eta$ varies strongly during the growth of the magnetic field, but always fulfills the condition Rm$\\ggg1$ and Pr$_M=\\nu/\\eta>1$. Detailed calculations concerning the ambipolar and Ohmic diffusion scales will be presented in a companion paper, in which we make use of the ionization degree obtained from a numerical simulation to provide an updated calculation of these scales for different field strengths. Gravitational collapse is generally accompanied by the presence of turbulence \\citep{Klessen09}, which may for instance be described by the theory of \\citet{Kolmogorov41}. Numerical simulations show that primordial star formation during the collapse phase occurs in a self-similar fashion, where the density profile at a given time is always well-described by a Bonner-Ebert sphere with a flat central density core \\citep{Abel02, Bromm03, Yoshida08}. Similar results have been found for present-day star formation \\citep[e.g.][]{Banerjee06b, Banerjee07a}. The gas falling on these central cores leads to weak shocks up to Mach~$1$, which drive turbulence in the central density core. This is reflected in the inhomogeneities in the central core and the sub-Keplerian angular momentum profiles, as reported by \\citet{Abel02, Bromm03, Yoshida08}. Under such conditions, a strong tangled magnetic field may be generated already during the collapse phase by the small-scale dynamo that was originally proposed by \\citet{Kazantsev68}. This dynamo provides a very generic means of amplifying magnetic fields and was also proposed to be important in the large-scale structure of the universe \\citep{Ryu08}. {The field amplification is due to the random stretching and folding of the magnetic field lines in a turbulent random flow. In the kinematic regime, the field grows typically on the eddy turnover time, $t_{\\rm ed} = l/v$ where $l$ is a typical turbulent length scale and $v$ is the turbulent velocity.} { In the context of galaxy formation, \\citet{Beck94} proposed that it is the small-scale dynamo that produces the seeds for galactic large-scale dynamos. As pointed out by \\citet{Arshakian09}, the small-scale dynamo can effectively amplify weak seed magnetic fields by $\\sim {13}$ orders of magnitude on a timescale $\\sim 300$ million years in the first galaxies. Similar results were obtained by \\citet{deSouza10} from a direct solution of Kazantsev's equation.} Capturing such dynamos in numerical simulations of protostellar collapse is extremely challenging, as it requires that the turbulent cascade is well-resolved and well-separated from the scale where MHD turbulence is numerically dissipated. State of the art numerical simulations of { turbulence thus require a spatial resolution of at least $512^3$ for a marginally resolved inertial range \\citep{Federrath08, Federrath10}.} Numerical simulations following protostellar collapse, on the other hand, typically resolve the Jeans length and thus the high density region with about $16$ cells, rendering them unable to capture the potential amplification via the turbulent dynamo. { \\citet{Federrath10} showed that at least 30 grid cells are required to resolve turbulent vortices.} Simulations as performed by \\citet{Xu08} therefore cannot resolve the turbulence in the central core. In this paper, we explore the implications of the small-scale dynamo during the gravitational collapse phase within a semi-analytic framework, applied to the formation of the first stars and galaxies. We first review the theoretical background and numerical evidence for the small-scale dynamo in \\S~\\ref{dynamo}, and present a set of analytic estimates. In \\S~\\ref{collapse}, we develop a quantitative model concerning the small-scale dynamo action during the collapse process. This model is applied in \\S~\\ref{application} both to minihalos and atomic cooling halos, taking into account the amount of turbulence that was found in numerical simulations. Phenomenological consequences from the generation of such magnetic fields are discussed in \\S~\\ref{conclusions}. { In a companion paper \\citep{Sur10}, we present numerical simulations confirming the importance of dynamo amplification during gravitational collapse.} ", "conclusions": "\\ We demonstrated in this paper that magnetic fields are generated rapidly by the small-scale dynamo both in minihalos which form the first stars, as well as in atomic cooling halos which may harbor the first galaxies. In this section, we summarize the main results and discuss the main consequences and open questions which cannot be resolved in this semi-analytic framework. \\subsection{Formation of the first stars} The formation of the first stars was often examined by hydrodynamical simulations that neglected potential effects from magnetic fields \\citep[e.g.][]{Abel02, Bromm04, Yoshida06}. Magnetic fields have however been considered to be important in the protostellar disk in presence of an efficient dynamo \\citep{Tan04, Silk06}. {The impact of a uniformly imposed magnetic field ($B_0$) in the primordial collapse has been studied by \\citet{Machida06, Machida08} using direct numerical MHD simulations. In these simulations, a range of $B_{0} = 10^{-6} - 10^{-9}$G was used for different values of the angular velocity of rotation. A protostellar jet within a radius of about 0.02 AU was found to be launched for an initial $B_{0}\\geq 10^{-9}$~G at $n=10^3$ cm$^{-3}$.} In this paper, we showed that the small-scale dynamo leads to a magnetic field strength much larger than the critical value of $10^{-9}$~G~$(n/10^3$~cm$^{-3})$ derived by \\citet{Machida06} for the formation of jets and outflows. However, as these magnetic fields are more tangled than those of \\citet{Machida06}, their results cannot be directly applied to ours, and additional numerical studies concerning tangled magnetic fields are required. The average Alfv'en velocity $v_A=B/\\sqrt{\\rho}$ is typically smaller than the sound speed, though it may dominate locally because of fluctuations in the magnetic field strength. In addition, our study clarifies that the magnetic field strength in the protostellar disk is much higher than previously anticipated. The conditions for the MRI, as formulated by \\citet{Tan04} and \\citet{Silk06}, are thus fulfilled. This leads to the presence of turbulence in the protostellar disk, which may drive a large-scale $\\alpha\\omega$ dynamo in the presence of some kinetic helicity, making the field stronger and more coherent. But also the MRI itself may further amplify the magnetic field \\citep{Balbus91}. This has important consequences for the fragmentation behaviour of the disk. Detailed numerical studies of the collapse of magnetised molecular cloud cores in the context of present-day star formation \\citep[e.g.][]{HennebelleT08, HennebelleF08, HennebelleC09, Mellon09} indicate that even modest field strengths can suppress binary formation and strongly favour the formation of single stars. Jets and magnetic tower flows are very effective in transporting away angular momentum and thus change structure and dynamics of the protostellar accretion disk. { On the other hand, numerical simulations exploring the interaction of turbulence generated by the MRI with gravitational instabilities indicate the excitation of additional modes and an effective reduction of the accretion rate, as well as the broadening of spiral arms by the MRI turbulence \\citep{Fromang04}. Dedicated numerical studies exploring the combination of such effects in a primordial accretion disk will thus be required to understand the full impact on the stellar masses.% } \\subsection{Formation of the first galaxies} The formation of the first galaxies is currently subject to much larger uncertainties than the formation of the first stars. This is because the initial conditions are not completely clear and the amount of metal enrichtment is not fully understood. The presence of supersonic turbulence has however been convincingly demonstrated by \\citet{Greif08} and \\citet{Wise08a} with cosmological simulations encorporating hydrodynamics and primordial chemistry. Additional physics like supernova feedback may just enhance the amount of turbulence found there. The small-scale dynamo is found to be extremely efficient under these conditions and may magnetise the material during the collapse, with the magnetic pressure only half an order of magnitude below the thermal pressure on average. Locally, the magnetic field may even dominate in some places, as it is expected that highly inhomogeneous fields are generated from the small-scale dynamo. Indeed, as discussed by \\citet{Subramanian99} and \\citet{Brandenburg05}, the magnetic field is highly inhomogeneous, reaching equipartition in about $10\\%$ of the volume. The saturation field strength we adopted above results from a spatial average over the different local values. We thus expect fluctuations of the field strength by at least a factor of $10$ \\citep{Wang09, Dubois09}. Due to the increase of the integral scale, the magnetic field becomes more coherent in these systems, making a stronger case for the putative presence of jets and outflows. As in the case of minihalos, the formation of a disk may lead to the presence of an $\\alpha\\omega$ dynamo that makes the magnetic field more coherent on disk scales. It may similarly play a role by making angular-momentum transport more efficient and thus reducing the amount of fragmentation and suppressing binary formation. As discussed above, the magnetic pressure may locally dominate over the thermal pressure. In this case, the magnetic Jeans mass sets the critical scale for fragmentation. For small-scale turbulent fields, it is defined in analogy to the thermal Jeans mass as \\begin{equation} M_{J,B}=2M_\\odot \\left(\\frac{v_A}{0.2\\ \\mathrm{km/s}} \\right)^3\\left(\\frac{n}{10^3\\ \\mathrm{cm}^{-3}} \\right)^{-1/2}\\propto \\frac{B^3}{{\\rho^2}}.\\label{thJeans} \\end{equation} Here, the Alfv'en speed $v_A=B/\\sqrt{4\\pi\\rho}$ replaces the sound speed $c_s$, as magnetic pressure support propagates with the Alfv'en speed. With \\begin{equation} v_A=2.0\\ \\mathrm{km/s} \\left( \\frac{B}{10^{-5}~\\mathrm{G}} \\right) \\left( \\frac{n}{10^2\\ \\mathrm{cm}^{-3}} \\right)^{-0.5}, \\end{equation} the Alfv'en speed in the saturation phase is thus larger than or comparable to the speed of sound. The presence of such fields thus provides additional stability during the formation of intermediate-mass black holes, which are often considered to form in such systems \\citep[e.g.][]{Eisenstein95, Koushiappas04, Begelman06, Spaans06, Shang09,Schleicher10b}. Thus, even if fragmentation cannot be totally avoided in hydrodynamical simulations, the presence of magnetic fields may still give rise to larger seed masses. The detailed consequences however need to be assessed with numberical simulations. Additional open questions concern the further evolution of the magnetic field on larger scales and the build-up of galactic-scale fields, as discussed by \\citet{Arshakian09}. Their model for the magnetic-field evolution in galaxies yields a number of predictions which can be tested with future radio facilities such as the SKA\\footnote{http://www.skatelescope.org/}, which can thus constrain the formation mechanisms of the small- and large-scales magnetic fields in new born and young galaxies. In this respect, cosmological simulations that include an approximate treatment for the mean-field induction equation, as performed by \\citet{Dubois09}, will be very important. An additional issue that needs to be addressed is the role of magnetic helicity. Magnetic helicity is a conserved quantity and affects magnetic field generation and decay. In the presence of helical fields, the small-scale dynamo also creates correlations on larger scales \\citep{Subramanian98, Subramanian99}. The decay law for helical fields was derived by \\citet{Hatori84}. It is independent from the large scale part of the spectrum (i.e. scales above the integral scale) and is generally less efficient due to helicity conservation. During such decay, magnetic power is shifted from small to large scales via an inverse cascade, thus increasing the typical coherence length \\citep{Christensson01}. Our paper was conservative in the sense that we assumed magnetic fields with zero helicity, yielding a lower limit on the integral scale. In the presence of helicity, magnetic fields may be coherent on larger scales, making it more straightforward to drive large-scale jets and outflows. To assess this issue, the turbulent properties of the first galaxies need to be analyzed and understood in further detail. \\subsection{Further discussion} Based on the estimates performed for this paper, it seems likely that the small-scale dynamo will be very efficient during the formation of the first stars and galaxies. The epoch of first star formation may thus also be the epoch where the first strong magnetic fields formed in the universe. This may be important for our understanding of primordial star formation. We further speculate that this mechanism may not only apply to the very first galaxies, but that the formation of any gravitationally bound structures lead to a sufficient amount of accretion-driven turbulence to amplify magnetic fields. We plan to investigate this proposition further with numerical simulations." }, "1003/1003.3901_arXiv.txt": { "abstract": "Several authors have claimed that the observable Hawking emission from a microscopic black hole is significantly modified by the formation of a photosphere or chromosphere around the black hole due to QED or QCD interactions between the emitted particles. Analyzing these models we identify a number of physical and geometrical effects which invalidate them. In all cases, we find that the observational signatures of a cosmic or Galactic background of black holes or an individual black hole remain essentially those of the standard Hawking model, with little change to the detection probability. ", "introduction": "\\label{aba:sec1} Microscopic black holes are of great interest in astrophysics and cosmology\\cite{C05}. A 4D black hole with a Hawking temperature~\\cite{H} $T_{bh}= 1.06\\ {\\rm GeV}/\\left( M_{bh}/10^{13}\\ {\\rm g} \\right)$ should emit all available particle species which appear non-composite compared with the wavelength of the radiated energy. Once $T_{bh}\\gtrsim \\Lambda_{QCD}\\simeq 200 - 300$ MeV, quarks and gluons should be directly emitted\\cite{MW} and then decay into stable species. It has been claimed that interactions between emitted particles significantly modify the astrophysically observable spectra from such black holes. Most scenarios are based on the Heckler\\cite{HE} model in which two-body bremsstrahlung and pair-production interactions form a QED photosphere at $T_{bh}\\gtrsim 45$ GeV and a QCD chromosphere at $T_{bh}\\gtrsim \\Lambda_{QCD}$. Here we summarize our recent detailed analysis\\cite{MCP,PCM} of interaction models and discuss various points which prevent the development of photospheres and chromospheres. ", "conclusions": "" }, "1003/1003.6009_arXiv.txt": { "abstract": "The Gas Pixel Detector belongs to the very limited class of gas detectors optimized for the measurement of X-ray polarization in the emission of astrophysical sources. The choice of the mixture in which X-ray photons are absorbed and photoelectrons propagate, deeply affects both the energy range of the instrument and its performance in terms of gain, track dimension and ultimately, polarimetric sensitivity. Here we present the characterization of the Gas Pixel Detector with a 1~cm thick cell filled with dimethyl ether (DME) at 0.79~atm, selected among other mixtures for the very low diffusion coefficient. Almost completely polarized and monochromatic photons were produced at the calibration facility built at INAF/IASF-Rome exploiting Bragg diffraction at nearly 45~degrees. For the first time ever, we measured the modulation factor and the spectral capabilities of the instrument at energies as low as 2.0~keV, but also at 2.6~keV, 3.7~keV, 4.0~keV, 5.2~keV and 7.8~keV. These measurements cover almost completely the energy range of the instrument and allows to compare the sensitivity achieved with that of the standard mixture, composed of helium and DME. ", "introduction": "Detectors able to image charged particle tracks in a gas have been developed over the last few years for different applications. One of the most promising is the possibility to resolve the path of photoelectrons emitted in the gas in consequence of a photoelectric absorption. The reconstruction of the initial direction of photoelectron emission opens the way for measuring the state of polarization of the absorbed photons because the former is modulated with respect to the direction of the photon electric field with a $\\cos^2$ dependency. This makes the photoelectric effect a good analyzer of X-ray polarization, and a perfect one for absorption from spherically symmetric shells. Only a few gas detectors can resolve so finely the photoelectron tracks to accurately reconstruct the initial direction of emission \\citep{Bellazzini2006, Black2007}. One of the most sensitive is the Gas Pixel Detector (GPD hereafter), developed by INFN-Pisa and INAF/IASF-Rome \\citep{Costa2001,Bellazzini2007} and currently inserted in the focal plane of several future satellite missions \\citep{Bellazzini2010b, Costa2010}. The gas cell is 1~cm or 2~cm thick and a number of mixtures of helium, neon or argon and dimethyl ether (DME hereafter) at 1~atm or 2~atm have been used, the choice of the gas being of fundamental importance for the polarimetric performance of the detector. A hard limit to the lower energy threshold of the instrument is about twice the binding K-shell energy of the absorbing component because above this threshold the photoelectron track, modulated with polarization, prevails on the isotropic one of the Auger electron. Photoelectron range is determined by density, while the average atomic number fixes the mean free path for scatterings with atomic nuclei, which is the length scale on which polarimetric information is smeared. The diffusion coefficient influences the blurring of the photoelectron track during drift in the gas cell and eventually the possibility to resolve the initial part of the photoelectron path and reconstruct correctly the direction of emission. We developed a Monte Carlo software to easily explore the behavior of the instrument to different mixtures and to subsequently test a subset of the most interesting ones. Recently \\citet{Muleri2008} measured the modulation factor $\\mu$, namely the amplitude of the response of the instrument for completely polarized photons, for the GPD filled with helium 20\\% and DME 80\\% at 2.6~keV, 3.7~keV and 5.2 keV. This data confirmed that measured values are basically consistent with what is expected on the basis of the Monte Carlo software and proved that X-ray polarimetry in Astrophysics with the GPD is feasible. In this paper we characterize the behavior of the GPD with a different gas, i.e. pure DME at 0.79~atm. In particular, we describe the configuration of the GPD and the calibration sources we used in Section~\\ref{sec:Setup}, while spectral capabilities of the instrument are discussed in Section~\\ref{sec:Spectral}. The measurement of the modulation factor between 2.0~keV and 7.8~keV is reported in Section~\\ref{sec:Mod}, together with the comparison with Monte Carlo results and what was reported previously on the He 20\\% and DME 80\\% mixture. Note that this is the first time that the modulation factor of a gas polarimeter is presented at energies as low as 2.0~keV. ", "conclusions": "The characterization of the GPD presented in this paper represents one step forward to the optimization of the instrument as a photoelectric polarimeter. In place of the standard mixture composed of a small fraction of helium (20\\%) and DME (80\\%) at 1~atm, we tested the GPD filled with pure DME at 0.79~atm. Measured modulation factor successfully confirms Monte Carlo simulations at lower energy, while above 3.7~keV there is a discrepancy $\\lesssim$10\\% with respect both the expected value and the sensitivity measured for the (similar) standard mixture. Although this issue does not significantly affect the performance of the GPD, which are mostly determined at low energy, it deserves further investigations to exploit in the best way the whole energy range of the instrument. Conversely the modulation factor measured at 2.0~keV, the lowest energy ever presented for a photoelectric polarimeter, is relatively large, 13.5\\%, and encourages attempts to reduce the threshold of the instrument below this energy value. This could allow to address interesting scientific objectives, like the study of the (possibly) highly polarized thermal emission from the surface of cooling neutron stars. We also discussed, for the first time in a systematic way, the spectral capabilities of the GPD. The spectrum of absorbed photons is obtained by summing the charge content of hit pixels and, in principle, it could be heavily affected by nonuniformities in the response of the pixels. As a matter of fact, the energy resolution is quite good, 24\\% at 5.9~keV, that is very close to that of standard proportional counters and to the requirement (20\\%) defined for future space missions that include the GPD. The dependency with energy, beyond a $1/\\sqrt{E}$ contribution that is naturally explained as due to Poisson fluctuations in production and amplification of primary charges, shows an additional term whose origin will be subject of further investigation. In conclusion, DME appears as an interesting alternative to mixtures of helium and DME. Despite the small reduction of the modulation factor at high energy, performance at low energy is confirmed. Possibly even more important, DME provides very good spectral capabilities, which are a mandatory complement of the polarimetric sensitivity to pursue the scientific objectives of any X-ray polarimetry mission." }, "1003/1003.4526_arXiv.txt": { "abstract": "Holography is expected as one of the promising descriptions of quantum general relativity. We present a model for a cosmological system involving two holographic screens and find that their equilibrium exactly yields a standard Friedmann-Robertson-Walker universe. We discuss its cosmological implications by taking into account higher order quantum corrections and quantum nature of horizon evaporation. We will show that this model could give rise to a holographic inflation at high energy scales and realize a late-time acceleration in a unified approach. We test our model from the SN Ia observations and find it can give a nice fit to the data. ", "introduction": "Einstein's classical general relativity is commonly acknowledged as the theory of gravitational interactions for distance sufficiently large compared to the Planck length. This validity has become the foundation of modern cosmology in describing the dynamics of our universe. However, its quantum effects are expected to become important at high energy scale, namely at very early time of cosmological evolution. Especially, the quantization of Einstein gravity has long been known to be perturbativly nonrenormalizable. Various attempts on solving this issue have been intensively studied in the literature. It is widely believed that the quantization of Einstein gravity is related to the solution to Big Bang singularity of our universe. As early as the discovery of black hole thermodynamics by Bekenstein\\cite{Bekenstein:1973ur} and Hawking\\cite{Hawking:1974sw}, people have realized that a nonperturbative feature of Einstein gravity may be related to the holographic thermodynamics. Especially, 't Hooft proposed the holographic principle as a particular property of quantum gravity which states that the description of a volume of space can be thought of as encoded on a boundary of this system, preferably a light-like boundary like a gravitational horizon\\cite{'tHooft:1993gx}. Subsequently, this issue is extensively discussed in cosmology\\cite{Gibbons:1977mu, Bousso:2002ju} and recently realized in the context of developments in string theory\\cite{Maldacena:1997re}. Therefore, it provides a promising description of quantum general relativity. An extended holographic picture was conjectured by Verlinde and in this scenario Einstein gravity is originated from an entropic force arising from the thermodynamics on a holographic screen\\cite{Verlinde:2010hp}(see also \\cite{Padmanabhan:2009vy, Padmanabhan:2009kr, Padmanabhan:2010xh} and references therein for earlier studies along this direction). In this scenario, however, there exists a key controversial issue whether gravity is fundamental or emergent\\cite{Gao:2010yy, Culetu:2010ua}, and thus relevant reinterpretations of Verlinde's Entropic force was discussed in \\cite{Hossenfelder:2010ih}. Recently, a much explicit formulation of Entropic gravity theory was suggested by Easson, Frampton and Smoot (EFS) , in which the general relativity is still a fundamental theory but including a boundary term. In this picture the holographic entropic force arises from the contribution of boundary terms\\cite{Easson:2010av}. This model was soon applied to realize current acceleration \\cite{Easson:2010av} and inflationary period at early universe \\cite{Easson:2010xf}(see \\cite{Wang:2010jm} for a study of entropic inflation within Verlinde's proposal). In both Verlinde's proposal and EFS one, one should be very careful of whether they can explain the current cosmological observations consistently\\cite{Danielsson:2010uy}. First of all, let us assume gravity is entropic, when applied into cosmology, we should explain why our CMB radiation and the holographic screen are not in thermal equilibrium. This problem is very manifest, since the temperature of our CMB radiation and thus our universe is observed as $T_{CMB}=2.73$K, but the horizon temperature can be easily estimated as $T_H\\sim{O}(10^{-30})$K. In such a thermal system out of equilibrium, the heat transfer ought to be very strong. Therefore, if such a heat transfer occurs today and leads to an acceleration, then our universe is always accelerating in the past without radiation and matter dominations since the temperature gap is always very large at early times (we have $T_H\\sim{T_{CMB}^2}\\sqrt{G}$ and therefore this system can only be in thermal equilibrium at Planck scale). This conclusion explains why the modified Friedmann equation appeared in the EFS paper is so different from the normal one in Einstein gravity. Concerning the above question, in the present work we are interested in how to recover a standard Friedmann equation from an entropic cosmological system. We suggest that there exist two holographic screens with one being the approximate Hubble horizon which is similar to the de-Sitter (dS) horizon, while the other Schwarzschild horizon. Each screen has its own thermodynamics which is independent of that of the other. Consider the classical dynamics of such a cosmological system, we find that its thermal equilibrium corresponds to a standard FRW universe. Therefore this model is able to explain the normal thermal history observed in our universe. Moreover, we consider quantum corrections to the area entropy and obtain an EFS universe a little bit deviating from thermal equilibrium. We study the cosmological implications of this model, and find that it can drive an entropic expansion at early universe and realize the late-time acceleration in a unified approach. The letter is organized as follows. In section II we briefly review the idea of entropic force from the viewpoint of an effective action description. In section III we provide an explicit formulation of the FRW universe from the classical thermal equilibrium state of double holographic screens. Effects from higher order quantum corrections of the holographic entropy and a quantum evaporation of the inner horizon are studied in this model. In section IV we study the cosmological implications of this model involving quantum corrections. Our results show that in this model a holographic inflation could be obtained at high energy scales, and a quantum evaporation process of the inner horizon is related to the realization of the late-time acceleration. We confront this model with the latest observations at the end of this section, and the results are in agreement with observational constraints. Section V presents a summary and discussions of the related works. We take the convention $c=k_B=\\hbar=1$ in this letter. ", "conclusions": "Since awaken by Verlinde, the idea of entropic force has become an important issue and its phenomenological applications were soon considered in an FRW universe\\cite{Cai:2010hk, Shu:2010nv, Gao:2010fw, Zhang:2010hi, Wei:2010wwa, Ling:2010zc, Lee:2010bg}, and is expected to explain the current acceleration of our universe\\cite{Li:2010cj}, and drive an inflationary period at early times\\cite{Wang:2010jm}, and its effects on spherical symmetric spacetime were discussed in \\cite{Smolin:2010kk, Caravelli:2010be, Wang:2010px, Myung:2010jv, Liu:2010na, Cai:2010sz, Tian:2010uy, Jamil:2009qs, Jamil:2010xq}, and see \\cite{Myung:2010rz, Konoplya:2010ak, Paeng:2010dj, Kar:2010uy, Hogan:2010zs, He:2010yf, Lee:2010ew, ChangYoung:2010rz, Banerjee:2010yd, He:2010ct} for relevant discussions. The frontier of modern physics suddenly goes back to pre-Einstein time one hundred years ago. In this Letter, we extended the picture of entropic cosmology and suggested a scenario of double holographic screens to explain the past thermal expansion of our universe. We also studied the quantum signatures of this model motivated from physics of quantum gravity, and specifically we considered the higher order quantum corrections to the holographic entropy and the process of horizon evaporations. We found that the higher order quantum corrections to the entropic force may give rise to an implement of holographic inflation. In the meanwhile, the evaporation of the inner horizon could bring a realization of late-time acceleration. In order to let this acceleration happen at current time, we found there exists a fine-tuning to the coefficients of higher order corrections. We test our model from the SN Ia observations, and find it can give a nice fit to the SN Ia data. The unification of inflation and dark energy era was earlier discussed in Refs. \\cite{Nojiri:2005sr, Nojiri:2005pu} by introducing a phantom degree of freedom. We would like to point out that the model we considered is still a toy model with many detailed clues ignored. Among them the most important issue is the study of primordial perturbations seeded by statistic fluctuations on the holographic screens, since we expect these thermal fluctuations could give rise to a nearly scale-invariant spectrum so that explain the CMB observations. We will perform a much complete and careful study on this issue in near future. At the end of this Letter, we would like to make a few comments on the possibility of the avoidance of big bang singularity in entropic cosmology. In the main text, we have studied the cosmological implications of the model of double holographic screens at early universe by considering the higher order quantum corrections. Moreover, in quantum physics, there could be more arguments supporting the avoidance of the Big Bang singularity. Namely, as the Heisenberg uncertainty principle states, in quantum theory a test particle is described by a wave packet, which moves in the bulk spacetime. Consider the measurement of the absolute position of this particle. It could be anywhere since the particle's wave packet has non-zero amplitude, meaning the position is uncertain. The Heisenberg uncertainty principle requires, $\\delta r \\delta E \\geq \\frac{1}{2}~.$ To combine the above uncertainty relation and Eq. (\\ref{Fe}), one can obtain a minimal length scale for the thermal system \\begin{eqnarray} \\Delta r\\simeq\\sqrt{\\frac{G}{4\\pi Tr}}\\sim l_{pl}~, \\end{eqnarray} which implies that our universe cannot be shrunk into trans-Planckian scale. Therefore, we may get a nonsingular cosmic evolution of the universe at early times. In this case, we expect the model could realize the late-time acceleration and also avoid the initial big bang singularity in a unified approach without quantum instability\\footnote{It is widely noticed that in the frame of standard Einstein gravity, a nonsingular bounce model offer suffers from quantum instability due to a ghost degree of freedom\\cite{Cai:2007qw, Cai:2007zv}, this statement can also be extended into the cyclic cosmology\\cite{Cai:2009zp, Cai:2006dm, Xiong:2007cn, Xiong:2008ic}. Recently, a nonsingular bounce model was achieved in the frame of nonrelativistic gravity theory\\cite{Brandenberger:2009yt, Cai:2009in, Cai:2009hc}(see also \\cite{Calcagni:2009ar, Kiritsis:2009sh}).}. If this model has a matter dominated contraction, it was found that both the thermal \\cite{Cai:2009rd} and quantum \\cite{Wands:1998yp, Finelli:2001sr, Cai:2008qw, Starobinsky:1979ty} fluctuations are able to provide a scale-invariant spectrum with local featured signatures\\cite{Cai:2008qb, Cai:2008ed} and sizable non-Gaussianities\\cite{Cai:2009rd, Cai:2009fn}, which may be responsible for the current cosmological observations. We note the study of entropic force with Heisenberg uncertainty principle in the original Verlinde's model appeared in \\cite{Vancea:2010vf}, and its extended form was analyzed in \\cite{Zhao:2010vt, Lee:2010fg, Kuang:2010gs, Ghosh:2010hz, Munkhammar:2010rg, Modesto:2010rm}." }, "1003/1003.4460_arXiv.txt": { "abstract": "We present results of deep polarization imaging at 1.4 GHz with the Dominion Radio Astrophysical Observatory as part of the DRAO Planck Deep Fields project. This deep extragalactic field covers 15.16 deg$^2$ centered at $\\alpha_{2000} = 16^{{\\rm h}} 14^{{\\rm m}}$ and $\\delta_{2000} = 54\\arcdeg 56\\arcmin$, has an angular resolution of $42\\arcsec \\times 62\\arcsec$ at the field center, and reaches a sensitivity of 55 $\\mu$Jy beam$^{-1}$ in Stokes $I$ and 45 $\\mu$Jy beam$^{-1}$ in Stokes $Q$ and $U$. We detect 958 radio sources in Stokes $I$ of which 136 are detected in polarization. We present the Euclidean-normalized polarized differential source counts down to 400 $\\mu$Jy. These counts indicate that sources have a higher degree of fractional polarization at fainter Stokes $I$ flux density levels than for brighter sources, confirming an earlier result. We find that the majority of our polarized sources are steep-spectrum objects with a mean spectral index of $-0.77$, and there is no correlation between fractional polarization and spectral index. We also matched deep field sources to counterparts in the Faint Images of the Radio Sky at Twenty Centimeters catalogue. Of the polarized sources, $77\\%$ show structure at the arc-second scale whereas only $38\\%$ of the sources with no detectable polarization show such structure. The median fractional polarization is for resolved sources is $6.8\\%$, while it is $4.4\\%$ for compact objects. The polarized radio sources in our deep field are predominantly those sources which are resolved and show the highest degrees of fractional polarization, indicating that the lobe dominated structure may be the source of the highly polarized sources. These resolved radio galaxies dominate the polarized source counts at $P_{\\rm 0} = \\sqrt{Q^2 + U^2} <3\\,$mJy. ", "introduction": "Cosmic magnetism has been the focus of many scientific papers describing efforts to understand the role of magnetic fields in astrophysics. Cosmic magnetic fields are observed through synchrotron radiation at radio wavelengths. The percentage polarization ($\\Pi = P/I$) and plane of polarization of the radiation provide information regarding the strength and direction of the magnetic field, while Faraday rotation provides directly information about the line-of-sight component of the magnetic field. One way of studying cosmic magnetic fields is through statistical studies of the properties of polarized radio sources. Radio source counts provide insight into the evolution of galaxies over a range of frequencies, luminosities, and galaxy types. Total intensity (i.e. Stokes $I$) radio source counts have been well studied down to $\\sim10\\,\\mu$Jy in concentrated areas \\citep{Hopkins2003,Windhorst2003,Owen2008}. \\citet{Windhorst2003} used 1.4 GHz surveys and models to show that evolving giant elliptical galaxies and quasars dominate the radio source counts for flux densities at 1.4 GHz of $S_{1.4} \\ge 0.5\\,$mJy, while starburst and normal galaxies begin to dominate the source counts at $S_{1.4} \\le 0.3\\,$mJy. \\citet{Owen2008} produce radio source counts at 1.4 GHz over a region of the \\textit{Spitzer} Wide-area InfraRed Extragalactic Survey \\citep[SWIRE,][]{Lonsdale2003} down to $15\\,\\mu$Jy. They note that there may be a natural confusion limit for total intensity near $1\\,\\mu$Jy given the finite median source size of $\\sim1.2\\arcsec$. All-sky surveys like the NRAO VLA Sky Survey \\citep[NVSS,][]{Condon1998} and Faint Images of the Radio Sky at Twenty Centimeters \\citep[FIRST,][]{White1997} have provided 1.4 GHz source counts with high accuracy down to $2.5\\,$mJy and $1\\,$mJy, respectively. Polarized radio source counts have only just begun to be studied. The NVSS was used by \\citet{Mesa2002} for a study of polarized source counts by correlating sources at 1.4 GHz from the NVSS with sources from the Green Bank 4.85 GHz catalogue \\citep{Gregory1996}. Sources with $S_{1.4} \\ge 80\\,$mJy and $\\Pi \\ge 1\\%$ were found to show an increase in degree of polarization with decreasing flux density. This trend is particularly evident for steep spectrum sources where the median $\\Pi$ increases from $\\Pi \\simeq1.1\\% $ for $S > 800\\,$mJy to $\\Pi \\simeq1.8\\%$ or $100 < S < 200\\,$mJy. They suggest that the origin of this correlation could be a change in the source population with decreasing flux density. Analyzing the polarization structure of the sources with $S > 500\\,$mJy \\citet{Mesa2002} noted that extended sources were less polarized than compact ones. \\citet{Tucci2004} confirmed the \\citet{Mesa2002} result for steep spectrum sources but not for flat spectrum sources. A possible explanation for this effect is an increase of the mean redshift with decreasing flux density. \\citet{Tucci2004} also noted that the Euclidean-normalized polarized source counts remained constant down to a flux density of $1\\,$mJy. \\citet{Beck2004} extrapolated the Euclidean-normalized polarized source counts below $1\\,$mJy by convolving the Stokes $I$ differential source count distribution \\citep[see also][]{Hopkins2000} with a probability distribution function in fractional polarized intensity, the later constructed using polarized sources from the NVSS with $S \\ge 80\\,$mJy. This model does not include a flux density dependence of fractional polarization, although they note that their extrapolation of the polarized source counts is a lower limit. Actual measurements of polarized radio sources down to $500\\,\\mu$Jy were presented by \\citet{Taylor2007} who found that the Euclidean-normalized polarized radio source counts remain constant down to $500\\,\\mu$Jy, with a continued increase in the fractional polarization of radio sources below 80 mJy. The median $\\Pi$ increases to 4.8\\% for sources with $10 < S_{1.4} < 30$ mJy. Deeper polarized source counts are required to determine if the anti-correlation of $\\Pi$ with flux density continues to lower flux density levels and to explore the polarized source count statistics at the boundary where the dominant population may be expected to change from AGN to starburst galaxies. We present a continuation of the \\citet{Taylor2007} polarization study of the DRAO ELAIS N1 deep field to study the polarized source population in the European Large Area \\textit{ISO} Survey North 1 region \\citep[ELAIS N1,][]{Oliver2000}. Aside from the observations at 15, 90, and $175\\,\\mu$m with \\textit{ISO} \\citep{Oliver2000}, the ELAIS N1 region has been observed at a multitude of other wavelengths. \\citet{Ciliegi1999} used the VLA in C configuration at $20\\,$cm in total intensity and reached a 5$\\sigma$ limit of $135\\,\\mu$Jy over $0.12\\,$deg$^2$, the \\textit{Spitzer} Wide-area InfraRed Extragalactic survey \\citep[SWIRE,][]{Lonsdale2003} used IRAC at 3.6, 4.5, 5.6, and $8.0\\,\\mu$m and MIPS at 24, 70, and $160\\,\\mu$m, ground based optical imaging ($u', g', r', i', z'$) covering the SWIRE fields was done by the Issac Newton Telescope Wide Field Survey \\citep[INT WFS,][]{McMahon2001}, \\citet{Manners2003} observed ELAIS N1 with \\textit{Chandra}, and the Giant Metrewave Radio Telescope (GMRT) was used at 610 MHz \\citep{Garn2008} and 325 MHz \\citep{Sirothia2009}. The ELAIS N1 field is also covered by the Westerbork Northern Sky Survey \\citep[WENSS,][]{wenss}, NVSS, and FIRST. Combining these data produces a large database for extragalactic source studies. By increasing both the sensitivity and coverage area of the \\citet{Taylor2007} study, we are able to observe a larger number of polarized radio sources. The total area of the observations has been increased by a factor of two, allowing us to increase the number of sources i the same flux density range as \\citet{Taylor2007}. Increasing the sensitivity of the observations allows us to investigate the polarized radio sources to the faintest flux densities to date and probe the nature of the faint polarized source population. In section 2, we describe the Dominion Radio Astrophysical Observatory synthesis telescope (DRAO ST) observations and preparation of the polarization images. In section 3, the method of detecting and cataloguing the DRAO Planck Deep Fields radio sources in the ELAIS N1 field is described along with the matching to counterparts in WENSS, NVSS, and FIRST. In section 4, we describe the nature of the polarized sources in our sample. ", "conclusions": "We present the final observations of the DRAO Planck Deep Field in the ELAIS N1 region along with a complete sample of compact polarized radio sources. The 40 field observations at 1.4 GHz reach a sensitivity of $55\\,\\mu$Jy beam$^{-1}$ in Stokes $I$ and $45\\,\\mu$Jy beam$^{-1}$ in Stokes $Q$ and $U$. There are 958 robustly detected radio sources in the 15.16 deg$^2$ of the observations down to a signal-to-noise of 8, 136 of which have detectable polarization. The polarized source counts from the DRAO ELAIS N1 deep field are presented down to a polarized intensity of $400\\,\\mu$Jy. The polarized source counts below $P_{\\rm 0} < 3\\,$mJy are higher than published polarized radio source count models. This increase shows a continuing trend of increased polarization fraction with decreasing flux density, confirming the \\citet{Taylor2007} result. We detect starburst and normal galaxies in the total intensity observations, but we have yet to reach the sensitivity required to detect the polarized emission from these galaxies. Our complete polarized source sample consists primarily of FRI and FRII radio galaxies. The DRAO deep field sources were matched to counterparts in the FIRST, NVSS and WENSS catalogues. We found that a higher fraction of polarized radio sources are detected in WENSS than those with no detectable polarization. The mean spectral index of the polarized sources is $-0.77\\,\\pm\\,0.01$ compared with $-0.87\\,\\pm\\,0.01$ for sources with no detectable polarization. The majority of polarized sources are steep spectrum objects. Moreover, we find no correlation between spectral index and fractional polarization. The DRAO deep field sources with a FIRST counterpart show that a higher fraction of polarized sources show structure at arc-second scales than the sources with no detectable polarization. There is a larger fraction of resolved polarized sources, 77 $\\pm$ 8\\%, than for sources with no detectable polarization, 38 $\\pm$ 3\\%. The median polarization of sources classified as compact is 4.4 $\\pm$ 1.1\\% and 6.8 $\\pm$ 0.7\\% for polarized sources classified as resolved objects in FIRST. One possible explanation for the increase in the polarized source counts at $P_{\\rm 0} < 3\\,$mJy is that the resolved sources are the dominant objects at these flux density levels. However, there is currently no information on the origin and location of the polarized emission of these faint polarized radio galaxies." }, "1003/1003.0463_arXiv.txt": { "abstract": " ", "introduction": " ", "conclusions": "The primary aim of the {\\it Wyoming Survey for H$\\alpha$}, or WySH, is to accurately quantify the \\hal\\ luminosity function via narrowband imaging spanning $\\lesssim 4$ square degrees and multiple cosmic epochs. Important features of the survey include the use of narrowband filter pairs at each epoch for improved stellar continuum subtraction and the spatial overlap with deep ultraviolet and infrared surveys that enable interesting follow-up studies. Buttressed by a total of nearly 1200 \\hal\\ detections, we find a modest evolution in the \\hal\\ luminosity function over $z \\sim 0.16$, 0.24, 0.32, and 0.40. The values of the volume-averaged cosmic star formation rate, found by integrating under the luminosity functions, change by a factor of two over this moderate stretch in redshift. Our results indicate that this evolution is largely driven by changes in the characteristic luminosity $L_*$, which also shows an evolution by a factor of two over these epochs. That the evolution in the cosmic star formation rate density over these intermediate redshifts is mainly influenced by systematic changes in the characteristic luminosity is consistent with the findings of Le Floc'h et al.\\ (2005), P\\'erez-Gonz\\'alez et al.\\ (2005), Magnelli et al.\\ (2009), and Westra et al.\\ (2010), though Ly et al.\\ (2007) find the evolution to be more driven by changes in the source number density. Placing our results in the larger context of the slew of recent emission line surveys for $\\dot{\\rho}_{\\rm SFR}$ over $0 \\lesssim z \\lesssim 1.5$, the evolution in the cosmic star formation rate density is estimated to be $\\dot{\\rho}_{\\rm SFR}(z) \\propto (1+z)^{3.4\\pm0.4}$. Results from a complementary near-infrared narrowband imaging survey of \\hal-emitters will extend this work to redshifts of $z \\sim 0.81$ and 2.2. Finally, the large volume covered by this optical survey enables a measure of the impact of cosmic variance. By separately analyzing the different fields in this survey, in particular the Lockman Hole and ELAIS-N1, we find a variation in $\\dot{\\rho}_{\\rm SFR}$ at the 20\\% level for any given redshift." }, "1003/1003.2520_arXiv.txt": { "abstract": "We have compared far-ultraviolet (FUV), near-ultraviolet (NUV), and \\ha\\ measurements for star forming regions in 21 galaxies, in order to characterise the properties of their discs at radii beyond the main optical radius ($R_{25}$). \\ In our representative sample of extended and non-extended UV discs we find that half of the extended UV discs also exhibit extended \\ha\\ emission. We find that extended UV discs fall into two categories, those with a sharp truncation in the \\ha\\ disc close to the optical edge (R$_{25}$), and those with extended emission in \\ha\\ as well as in the ultraviolet. Although most galaxies with strong \\ha\\ truncations near R$_{25}$ show a significant corresponding falloff in UV emission (factor 10--100), the transition tends to be much smoother than in \\ha, and significant UV emission often extends well beyond this radius, confirming earlier results by \\citet{thilker07a} and others. \\ After correcting for dust attenuation the median fraction of total FUV emission from regions outside of $R_{25}$ is 1.7\\%, but it can be as high as 35\\% in the most extreme cases. The corresponding fractions of \\ha\\ emission are approximately half as large on average. This difference reflects both a slightly lower ratio of \\ha\\ to UV emission in the HII regions in the outer discs, as well as a lower fraction of star clusters showing HII regions. Most HII regions in the extended disc have fluxes consistent with small numbers of ionising O-type stars, and this poor sampling of the upper initial mass function in small clusters can probably account for the differences in the emission properties, consistent with earlier conclusions by \\citet{zaritsky07}, without needing to invoke a significant change in the stellar IMF itself. Consistent \\ha/FUV ratios and brightest HII region to total \\ha\\ fluxes in the inner and extended discs across our whole galaxy sample demonstrate no evidence for a change in the cluster luminosity function or the IMF in the low gas density outer disc. ", "introduction": "One of the most interesting discoveries of the Galaxy Evolution Explorer (GALEX) satellite has been the presence of very extended ultraviolet (UV) emitting discs around many nearby galaxies. \\citet{thilker05} first reported an extended UV disc for the spiral galaxy M83. Since then extended discs have be found in many galaxies, most notably NGC 4625 \\citep{gildepaz05}. Extended discs typically exhibit emission well beyond the classical optical edge of the galaxy, which is usually defined by R$_{25}$, the radius at which the surface brightness in the B band drops below 25 magnitudes arcsec$^{-2}$. In M83 UV knots are found extending to 4 R$_{25}$, and are associated with large scale filamentary HI structures. These structures are now commonly termed extended ultraviolet discs (XUV-discs). Subsequent studies have begun to characterise the frequency of XUV-discs and the properties of the star forming regions in these discs. \\citet{zaritsky07} examined a sample of 11 galaxies and found an excess of blue (FUV -NUV $<$ 1, NUV $<$ 25) sources out to 2 R$_{25}$ for $\\sim 25\\%$ of their sample. Based on and analysis of the GALEX Nearby Galaxies Survey (NGS; \\citet{gildepaz07a}), \\citet{thilker07a} concluded that extended discs are common, and can be divided into two distinct types. Type 1 discs ($\\ga$ 20\\% occurence) show structured UV bright regions beyond the typical star forming threshold; Type 2 discs ($\\approx$10\\% occurrence) display diffuse regions of UV emission, but not reaching extreme galactocentric radii. The vast majority of galaxies do not exhibit any significant emission beyond R$_{25}$, and are not classified as extended, though it should be noted that there were no quantitative definitions to identify extended discs. HII regions are often but not always associated with the extended UV emission. HII regions located well beyond the main star forming disc were identified in one of the first \\ha\\ surveys of galaxies (e.g. \\citet{hodge69, hodge74, hodge83}), and have been studied in detail by \\citet{ferguson98}, \\citet{vanzee98}, and \\citet{lelievre00}. Isolated intergalactic HII regions have also been identified by \\citet{gerhard02}, Sakai et al. (2002) and \\citet{cortese04} in the Virgo and Abell 1367 clusters, and by \\citet{ryanweber04a} for nearby galaxies in groups. The lifetimes of massive ionising stars responsible for \\ha\\ emission are much lower than those of the stars responsible for most of the near-UV emission (of order 10 Myr vs 100 Myr), and this is reflected in the relative numbers of UV and \\ha\\ knots observed by \\citet{zaritsky07}. The formation of massive stars beyond the classical 'edge' of galaxies raises questions about their formation and properties, as well as their possible impacts on their low gas density environment. Spectroscopy of HII regions at extreme radii in M83 and NGC4625 by \\citet{gildepaz07b} showed them to be consistent with regions dominated by a single ionising source star with masses in the range 20-40 \\msol. Ages for most of the UV knots are in the range 0--200 Myr, based on their integrated fluxes and colors (e.g., \\citet{thilker05}, \\citet{gildepaz05}, \\citet{zaritsky07}, \\citet{dong08}). Masses for the regions are more difficult to constrain; \\citet{gildepaz05} report a mass range of $10^{3}-10^{4}$ \\msol\\ for NGC4625, and \\citet{werk08} gives an upper mass limits in the range 600\\msol\\ to 7000\\msol\\ for HII regions associated with NGC1533. \\citet{dong08} used the UV to IR spectral energy distributions to estimate the masses of UV knots and measured a range of $10^{3}-3\\times 10^{6}$ \\msol with a peak $\\approx 10^{4.7}$ \\msol \\ for M83. In a study of the HII regions of NGC628 \\citet{lelievre00} postulated that the cluster mass function of these HII regions may be significantly different to that of the disc. Relatively little is known yet about the chemical abundance properties of the outer discs, but preliminary results show a range of behaviours. Abundance studies of a few systems such as M101 and NGC 628 show a continuous exponential abundance gradient extending into the outermost discs \\citep{kennicutt03,ferguson98}. However a recent study of HII regions in the XUV disc of M83 show chemical properties that are decoupled from those of the brighter inner disc, with a nearly constant oxygen abundance of $\\sim$0.3 (O/H)$_\\odot$ \\citep{bresolin09}. The XUV-discs are also important for understanding star formation thresholds in galaxies. Radial profiles in the \\ha\\ often show sharp turnovers, usually located near the optical R{$_{25}$ radius. \\citet{kennicutt89,martin01} showed that the distribution of gas density appears to be roughly continuous across this truncation in contrast to the massive star formation threshold visible in \\ha\\ profiles. These observation have been interpreted as arising from a threshold surface gas density, which could be attributed to gravitation instability \\citep{kennicutt89,martin01}, gas phase instabilities \\citep{schaye04} or gas cloud fragmentation instability \\citep{krumholz08}. However UV emission, which also traces star formation does not display truncations as pronounced as those seen in \\ha\\ \\citep{thilker05}. This raises questions about the nature and interpretation of the apparent \\ha\\ thresholds. A closely related question is the physical origin of the XUV-discs. \\citet{dong08,bush08} have argued that the XUV -discs can be understood in the context of normal gravitational threshold picture. Although the average gas densities beyond the \\ha\\ thresholds are too low in theory to allow star formation, in localised regions the density may be high enough to prompt star formation in a small volume. Others have suggested that interactions play a role. Interactions may disturb the gas creating regions that collapse and thus form pockets of stars. \\citet{gildepaz05} proposed such an interaction between NGC 4625 and its neighbor NGC 4618 and possibly NGC 4625A as a likely trigger for star formation at large radii. Similar processes may be able to account for the presence of HII regions at very large radii in some interacting and merging systems \\citep{ryanweber04a,oosterloo04,werk08}, but other factors may contribute such as galactic outflows and spiral density waves. \\citet{elmegreen06} modelled star formation in the outer discs, citing the effects of compression and turbulence as well as the propagation of gaseous arms as continual drivers for low level star formation. A common question underlying these studies is whether star formation at large radii represents a simple continuation of star formation with in the inner disc, or is it a separate mode of star formation altogether? Perhaps the most radical interpretation of the XUV discs is that they represent low surface density environments in which the stellar initial mass function is truncated at masses well below the $\\sim$100 \\msol\\ \\citep{kroupa08, meurer09}. A strong preferential suppression of massive star formation in the discs could account for the putative excesses of UV emission relative to \\ha\\ emission, and possibly account for the differences in the outer disc profiles at the respective wavelengths. Before these questions posed above can be answered we need uniform measurements of the UV and \\ha\\ properties of the extended discs for a representative sample of nearby galaxies. Although the published studies to date have highlighted the remarkable properties of a handful of galaxies, most have focussed on case studies of extreme examples such as M83 and NGC 4625. The goal of this study is to address the phenomenon in a broader context, by carrying out detailed photometric measurements in the UV and \\ha\\ for the discs and their individual star-forming regions, for 20 galaxies selected from the Spitzer Infrared Nearby Galaxies Survey (SINGS; \\citet{kennicutt03}) as well as the prototype galaxy M83. Some of our work parallels an important study of UV radial profiles of nearby galaxies by \\citet{boissier07}, but our study differs in emphasis in focussing especially on the relation of the UV and \\ha\\ discs, and on measurements of the individual star-forming knots. The remainder of this paper is organized as follows. In Section 2 we describe the GALEX and groundbased \\ha\\ observations that were used in this study, and in Section 3 we describe the processing and analysis of these data, including the treatment of dust attenuation. In Section 4 we present the observed properties of the extended discs in terms of their UV emission, \\ha\\ emission, and star formation properties, and in Section 5 we describe the properties of the populations of star-forming regions in both the inner and outer discs of the galaxies. In Section 6 we discuss the results in the context of the questions raised above, and in Section 7 we present a brief summary of results and conclusions. \\section[]{The Data} \\subsection{Sample Selection} All the galaxies in our sample were taken from the Spitzer infrared Nearby Galaxies Survey (SINGS) and the GALEX Nearby Galaxy Survey (NGS). SINGS is a multi-wavelength imaging and spectroscopic survey of 75 nearby (D$<$30Mpc) galaxies, including visible, near-infrared, ultraviolet and radio observations. The galaxies were selected to span a range of galaxy types, luminosities, and infrared/optical properties among normal galaxies in this volume. We selected a subset of 20 spiral galaxies, with emphasis on nearby, face-on or moderately inclined objects with major axis diameters in the range $2\\arcmin6$, corresponding to a time less than a tenth of the age of the Universe, roughly one billion years after the Big Bang. Follow-up observations confirmed that at least some of these quasars are powered by super-massive black holes with masses $\\simeq 10^9\\, M_\\odot$ \\citep{Barthetal2003,Willottetal2005}, probably residing in the centers of substantial galaxies. However, these exceptionally bright quasars are most likely just the tip of the iceberg: rare objects-- on the tail of the mass distribution. This implies larger numbers of less exceptional objects, and that MBHs existed in large numbers during the Dark Ages, before or around the time when the first stars formed. We are therefore left with the task of explaining the presence of MBHs when the Universe is less than {\\rm 1 Gyr} old, and of much smaller MBHs lurking in {\\rm 13 Gyr} old galaxies. The outstanding questions concern how and when ``seed'' MBHs formed, the frequency of MBHs in galaxies, and how efficiently MBH seeds grew in mass during the first few billion years of their lives. The formation of MBHs is indeed far less understood than that of their light, stellar mass, counterparts, end-points of stellar evolution for stars more massive than $\\simeq 20\\,\\msun$. The ``flow chart\" presented by \\cite{Rees1978} still stands as a guideline for the possible paths leading to formation of massive MBH seeds in the center of galactic structures. In the following I will review the main physical processes thought to influence MBH formation. I will mostly focus on astrophysical processes that happen in galaxies, and I will consider three possibilities: (i) that MBHs are the remnants of the first generation (PopIII) stars, (ii) that MBHs formation is triggered by gas-dynamical instabilities, (iii) that MBH seeds are formed via stellar-dynamical processes. A simplified scheme that describes the possible routes to MBH formation in high-redshift galaxies is shown in Figure~\\ref{fig:scheme}. I will also briefly mention the possibility that MBHs are related to inflationary primordial black holes. \\begin{figure} \\includegraphics[width=\\textwidth]{diagram1} \\caption{Scheme of possible pathways to MBH formation in high-redshift galaxies. Artwork by B. Devecchi.} \\label{fig:scheme} \\end{figure} %Outline %\\begin{itemize} %\\item{1.} Massive black holes in the local Universe: where from? %\\item{2.} Massive black hole formation %\\begin{itemize} %\\item{a.} PopIII stars remnants %\\item{b.} Gas-dynamical processes %\\item{c.} Stellar-dynamical processes %\\tem{d.} Primordial black holes %\\end{itemize} %\\item{3.} Massive black hole growth: hierarchical evolution of mass + spin %\\begin{itemize} %\\item{a.} Mergers: dynamical evolution and gravitational waves %\\item{b.}Accretion: active galactic nuclei, quasars %\\end{itemize} %\\item{4.} Observational signatures of massive black hole seeds %\\end{itemize} ", "conclusions": "We can trace the presence of `super' MBHs at early cosmic times, as the engines powering the luminous quasars that have been detected at high redshifts, corresponding to about 1 billion years after the Big Bang \\citep{Fanetal2001a}. The {\\it HST}, {\\it Chandra} and {\\it Spitzer} satellites, jointly with 8-m class telescopes and large surveys, have made important breakthroughs, and observational cosmology probed capable of putting constraints on MBHs, when shining as quasars, up to high redshift \\citep{Fanetal2001a, Fanetal2001b, Fanetal2004, Barthetal2003, Willottetal2003, Walteretal2004}. Some constraints on the global accretion history, even at high redshift, can be already put by comparing theoretical models predictions to ultra-deep X-ray surveys \\citep{salvaterra07}. The early evolution of MBHs, and most notably, what physical mechanism is responsible for their formation are however still unknown. We now do know that MBHs {\\it are} there, but we do not know how they {\\it got} there. In this article I focused on three plausible mechanisms of MBH seed formation. Broadly speaking, we can divide them into physically-based categories: `light seeds', forming at very early cosmic times ($M_{BH} \\simeq 100-600 \\msun$, $z\\simeq 20-50$), `heavy seeds', forming later on ($M_{BH} \\simeq 10^4-10^6 \\msun$, $z\\simeq 5-10$), and `intermediate seeds', forming with masses, and at epochs, in between the two previous cases ($M_{BH} \\simeq 10^3 \\msun$, $z\\simeq 10-15$),. Light seeds forming early have a longer time to grow by accretion and MBH-MBH mergers, on the other hand, their accretion rates could be depressed in the shallow potential wells of the (mini-)halos, especially in the presence of radiative feedback from the (mini-)quasar itself \\citep[and references therein]{Milos2009}. Heavy seeds forming later, in more substantial galaxies, are not likely to suffer from the same problems, but have had less time to grow. This would be partly compensated by their larger initial masses (Volonteri \\& Begelman 2010). At the current time observational constraints are too weak to favor one model against the others, but future X--ray missions, such as {\\it IXO}, and near infrared facilities such as {\\it JWST}, will have the technical capabilities to detect accreting MBHs at $z\\gta 6$, giving constraints on the accretion properties of MBHs at early times. If the mass of the seeds is below $\\sim10^5 \\msun$, their flux is too weak for single sources to be detected electromagnetically. Seeds of mass $\\lta 10^5 \\msun$ can nevertheless be directly identified during their mergers, by detecting their emission of gravitational radiation \\citep{Hughes2002,berti2006}. Additionally, gravitational waves produced during the inspirals of compact objects into MBHs -- extreme-mass-ratio inspirals (EMRIs) are expected to provide accurate constraints on the population of MBHs in the $10^4M_{\\odot}$--$10^7M_{\\odot}$ range (Gair et al. 2010), which is the mass range where we can expect `memory' of the initial conditions, as detailed in section 4.4. The combination of electromagnetic and gravitational wave observations in the coming years will improve the currently limited constraints on what route, or routes, lead to MBH seed formation." }, "1003/1003.0688_arXiv.txt": { "abstract": "We present results from a WIYN/OPTIC photometric and astrometric survey of the field of the open cluster NGC 188 ((l,b) = (122.8\\arcdeg, 22.5\\arcdeg)). We combine these results with the proper-motion and photometry catalog of Platais et al. and demonstrate the existence of a stellar overdensity in the background of NGC 188. The theoretical isochrone fits to the color-magnitude diagram of the overdensity are consistent with an age between 6 and 10 Gyr and an intermediately metal poor population ([Fe/H] = -0.5 to -1.0). The distance to the overdensity is estimated to be between 10.0 and 12.6 kpc. The proper-motions indicate that the stellar population of the overdensity is kinematically cold. The distance estimate and the absolute proper motion of the overdensity agree reasonably well with the predictions of the Pe\\~{n}arrubia et al. model of the formation of the Monoceros stream. Orbits for this material constructed with plausible radial-velocity values, indicate that dynamically, this material is unlikely to belong to the thick disk. Taken together, this evidence suggests that the newly-found overdensity is part of the Monoceros stream. ", "introduction": "In the last decade, large-scale surveys have demonstrated the existence of extended tidal streams and overdensities that are believed to be generated by the disruption of massive systems such as dwarf galaxies as they merge with our Galaxy. To date, there are three reasonably well-documented large structures known to reside in the halo and outer disk of the Galaxy: the Sagittarius stream with its known progenitor the Sagittarius dwarf galaxy (Ibata et al. 1994, Ibata et al. 2001, Newberg et al. 2002, Majewski et al. 2003), the Monoceros stream (Newberg et al. 2002, Yanny et al. 2003, Ibata et al. 2003, Rocha-Pinto et al. 2003, Conn et al. 2005, 2007, 2008) with the much-disputed possible progenitor the overdensity in Canis Major, and finally the Virgo Overdensity (Newberg et al. 2002, Vivas \\& Zinn 2003, Juri\\'{c} et al. 2008) with no known progenitor (see however Casetti-Dinescu et al. 2009 for a tentative suggestion). These structures' formation has direct implications on the history of the assembling of the Milky Way and thus on the merger history in a $\\Lambda$CDM cosmology. \\renewcommand{\\thefootnote}{\\arabic{footnote}} \\setcounter{footnote}{0} Here we investigate a stellar overdensity found in the background of the open cluster NGC 188, and provide evidence that is is part of the Monoceros stream. Our data may help better constrain models of the formation and origin of this immense ring-like feature mapped from l = $60\\arcdeg$ to $300\\arcdeg$, above and below the Galactic plane (Conn et al. 2008 and references therein). In this investigation, we make use of the excellent absolute proper-motions and photometry catalog provided by Platais et al. (2003, hereafter P03) over an area of 0.75 square degrees down to $V = 21$. We also add new, deep $VI$ photometry for an area half that of P03 obtained with WIYN 3.5m\\footnote{The WIYN Observatory is a joint facility of the University of Wisconsin-Madison, Indiana University, Yale University and the National Optical Astronomy Observatory.}, that reaches $V \\sim 23.5$. We present our data in Section 2, the color-magnitude diagram (CMD) interpretation in Section 3, the proper-motion analysis in Section 4, and discuss the origin of this newly-found overdensity in Section 5. In Section 6 we summarize our results. ", "conclusions": "\\subsection{The Overdensity as Part of the Monoceros Stream} The overdensity is located at (l,b) = $(122.8\\arcdeg, 22.5\\arcdeg)$. In what follows, we will work with two values for the distance to the overdensity determined in Section 3: 10.0 and 12.6 kpc. For 10 kpc, the overdensity is at $(X,Y,Z) = (13.0,7.8,3.8)$ kpc, while for 12.6 kpc it is at $(X,Y,Z) = (14.3,9.8,4.8)$ kpc, where the Sun is at $(X,Y,Z) = (8,0,0)$ kpc. Besides the properties of the stellar population, another hint that this overdensity may belong to the Monoceros stream is its Galactic location. Of the numerous studies that have sampled this structure, here we mention those that have found positive detections near our region: Rocha-Pinto et al. (2003) analyze M giants from 2MASS to map out this structure; indeed they find positive detections at (l,b) $ \\sim (130\\arcdeg, 26\\arcdeg)$ with an inferred heliocentric distance between 9 and 13 kpc. Conn et al. (2005) trace the main sequence of Mon using the Wide Field Camera on the Isaac Newton Telescope. They find a positive detection at (l,b) = $(118\\arcdeg, 16\\arcdeg)$ and at a heliocentric distance of $\\sim 12$ kpc. Momany et al. (2006) investigate an alternative explanation of the overdensity at the Galactic location of the Conn et al. detection, namely the flared disk. Their model of the warped and flared disk is based on 2MASS red clump and red giant stars. For this Galactic location, Momany et al. provide a map of the scale height of the flared disk as a function of distance from the Galactic plane and from the Galactic center. This map is the closest in direction, of those presented, to our NGC 188 field. In this map our overdensity, for both distance determinations, is located beyond the contour corresponding to $3\\times$ the scale height of the flared disk. Therefore, its position alone makes the overdensity unlikely to belong to the flared disk as recently modeled by Momany et al. (2006). In order to provide further evidence for this tentative identification, we compare our distance estimation and proper motion to the predictions of the Pe\\~{n}arrubia et al. (2005, hereafter P05) model for the Monoceros stream. This model describes the disruption of a satellite on a prograde, low inclination, low eccentricity orbit and is constrained by spatial distribution, distance estimates and radial velocities in overdensities mapped above and below the Galactic plane and between $l = 110\\arcdeg - 240\\arcdeg$. In Figure 9, we show the spatial distribution of the debris from the P05 model (top panel). The location of our field is marked (red symbol), and a sample of model particles selected within $4\\arcdeg\\times4\\arcdeg$ centered around our field are also shown (black squares). The middle and bottom panels show the proper motions in Galactic longitude and latitude as a function of distance for the entire model (grey), for the particle sample coincident with our field (squares) and for our field (red). The model indicates that there are two distance groups, one at $\\sim 9-12$ kpc, and the other at $17-23$ kpc. The near group, which corresponds well with our data divides into two groups according to the proper motion in latitude: one group moves downward, toward the Galactic plane, the other, away from the plane. Our data fits very well with the group moving away from the plane. Located at a distance of 3.8/4.8 kpc from the plane, stars in this overdensity will continue to get further away from the plane as they proceed in a prograde sense. We have found no evidence in the proper-motion data of the other group postulated by the model, i.e. the one moving toward the plane. Likewise, we find no evidence in our data of the more distant group at 20 kpc, which may be due to the magnitude limit of our observations. \\subsection{Plausible Orbits} As we do not have a measurement of the radial velocity (RV) for stars in the overdensity, we explore a range of plausible values in order to have a better understanding of the likely orbits constrained by the absolute proper motion and by the distance estimate. From the Besancon model, stars in the color ($B-V < 1.0$) and magnitude ($17 < V < 22$) ranges occupied by the overdensity have a heliocentric RV distribution centered at $\\sim -70$ km~s$^{-1}$, with a FWHM of 140 km~s$^{-1}$, and a tail toward negative values of $\\sim -300$ km~s$^{-1}$. We therefore choose the following RV values to be explored: 0, -70, -140 and -300 km~s$^{-1}$ and two extreme values of 100 and -400 km~s$^{-1}$. In Table 1 we list the adopted heliocentric RV, and the resulting velocity components in cylindrical coordinates $(\\Pi, \\Theta, W)$. We also list the maximum distance from the plane and the eccentricity of the respective orbits. The adopted peculiar velocity of the Sun is from Dehnen \\& Binney (1998). The velocity of the local standard of rest is $(\\Pi, \\Theta, W)$ = (0, 220, 0) km~s$^{-1}$. We have integrated the orbit in the Johnston et al. (1995) Galactic potential. \\begin{table}[htb] \\caption{Plausible Orbits for the Newly Found Structure} \\begin{tabular}{rrrrrrrrrrrrr} \\tableline \\\\ & \\multicolumn{5}{c}{d = 10.0 kpc} & & & \\multicolumn{5}{c}{d = 12.6 kpc} \\\\ \\tableline \\multicolumn{1}{c}{$V_r$} & \\multicolumn{1}{c}{$\\Pi$} & \\multicolumn{1}{c}{$\\Theta$} & \\multicolumn{1}{c}{W} & \\multicolumn{1}{c}{$z_{max}$} & \\multicolumn{1}{c}{e} & & & \\multicolumn{1}{c}{$\\Pi$} & \\multicolumn{1}{c}{$\\Theta$} & \\multicolumn{1}{c}{W} & \\multicolumn{1}{c}{$z_{max}$} & \\multicolumn{1}{c}{e} \\\\ \\multicolumn{1}{c}{(km~s$^{-1}$)} & \\multicolumn{3}{c}{(km~s$^{-1}$)} & \\multicolumn{1}{c}{(kpc)} & & && \\multicolumn{3}{c}{(km~s$^{-1}$)} & \\multicolumn{1}{c}{(kpc)} & \\\\ \\tableline \\\\ 0 & 44 & 241 & 99 & 7.3 & 0.27 &&& 43 & 250 & 123 & 11.0 & 0.34 \\\\ -70 & -14 & 212 & 72 & 5.8 & 0.03 &&& -17 & 225 & 96 & 9.0 & 0.13 \\\\ -140 & -72 & 184 & 46 & 5.0 & 0.23 &&& -77 & 200 & 69 & 8.0 & 0.19 \\\\ -300 & -205 & 118 & -16 & 4.2 & 0.70 &&& -213 & 142 & 8 & 7.6 & 0.66 \\\\ \\\\ 100 & 126 & 282 & 137 & 11.7 & 0.59 &&& 128 & 286 & 161 & 18.2 & 0.64 \\\\ -400 & -288 & 77 & -54 & 5.2 & 0.89 &&& -298 & 106 & -30 & 9.1 & 0.86 \\\\ \\tableline \\end{tabular} \\end{table} In all cases explored here that cover the entire plausible RV-range, the orbit is prograde. Also, the maximum distance above the plane is rather high, thus making the orbits unlikely to represent material in the warped or flared disk as described by Momany et al. (2006) for instance. For the 12.6 kpc distance estimate, the orbits are even more inconsistent with a thick-disk origin of this overdensity than for the 10-kpc distance estimate, as the material travels quite far from the Galactic plane. The orbit with the lowest height above the plane (i.e. 4.2 kpc, for the 10.0 kpc distance estimate) has a highly eccentric orbit, again inconsistent with material belonging to the thick disk. To achieve an orbit more confined to the Galactic plane, for instance within $\\pm 3$ kpc from the Galactic plane we require a distance to the overdensity of 7 kpc for our measured absolute proper motion. This implies a brighter turnoff by $\\sim 0.8$ mag, which the $BV$ data do not support. However, some of the orbits explored here (RV = 0 to -140 km~s$^{-1}$) are quite compatible with the one proposed by P05 for the model of the Monoceros stream. Therefore, the orbits discussed in combination with the properties of the stellar population of the newly found overdensity argue in favor of its membership to the Monoceros stream." }, "1003/1003.0477_arXiv.txt": { "abstract": "We present a two-day long RXTE observation and simultaneous Swift data of the bright X-ray transient XTE J1752-223. Spectral and timing properties were stable during the observation. The energy spectrum is well described by a broken power-law with a high energy cut-off. A cold disc ($\\sim 0.3$ keV) is observed when Swift/XRT data are considered. The fractional rms amplitude of the aperiodic variability (0.002--128 Hz) is $48.2\\pm0.1\\%$ and it is not energy dependent. The continuum of the power density spectrum can be fitted by using four broad-band Lorentzians. A high frequency ($\\sim 21$ Hz) component and two weak QPO-like features are also present. Time-lags between soft and hard X-rays roughly follow the relation $\\Delta t \\propto \\nu^{-0.7}$, with delays dropping from $\\sim 0.5$ (0.003 Hz) to $\\sim 0.0015$ ($\\geq$10 Hz) seconds. Our results are consistent with XTE J1752-223 being a black-hole candidate, with all timing and spectral components very similar to those of Cyg X-1 during its canonical hard state. ", "introduction": "Black hole X-ray transients (BHT) represent the majority of the black hole binary (BHB) population known so far. These systems spend most of their lives in quiescence, displaying luminosities too low to be detected by X-ray all-sky monitors (see e.g. \\citealt{Garcia1998}). However, they also undergo outburst events in which they become as bright as persistent sources, allowing their discovery. During these episodes, both the spectral and the time variability properties of BHTs vary dramatically, yielding the so-called \\textquoteleft states\\textquoteright. There is still much discussion about how many different states there are, and their correspondence with different physical conditions (see e.g. \\citealt{Belloni2010} for a general description), but the presence of a $hard$ state (historically known as \\textit{low/hard}; LHS) at the begining of the outburst which evolves towards a $soft$ state (\\textit{high/soft}; HSS) is widely accepted. The LHS, also associated with the last part of the outburst, is characterized by a power-law dominated energy spectrum with a power-law index of $\\sim 1.6$ (2--20 keV band). A high energy cut-off ($\\sim$ 60--200 keV; \\citealt{Wilms2006}; \\citealt{Motta2009}) is observed and aperiodic variability with a fractional root mean square amplitude (rms) above 30\\% is seen. The energy spectrum is softer during HSS, being dominated by a thermal disc black body component. However, a hard tail up to $\\sim 1$ MeV is present (\\citealt{Grove1998}). The rms associated with the aperiodic variability drops until 1\\% or less. These two \\textquoteleft canonical\\textquoteright ~states were first proposed to describe the behaviour of the prototypical BHB Cyg X-1.\\\\ XTE J1752-223 was discovered by the \\textit{Rossi X-ray Timing Explorer} (RXTE) on October 23, 2009 (\\citealt{Markwardt2009}). The source showed a 2--10 keV flux of 30 mCrab. Significant similarities with the typical properties of BHT during the LHS were soon noticed by \\cite{Markwardt2009a} and \\cite{Shaposhnikov2009}. A bright optical counterpart was detected (\\citealt{Torres2009}), showing in the optical spectrum a broad \\ha~emission line ($FWHM\\sim750$ \\kms) typical of accreting binaries (\\citealt{Torres2009a}). A radio counterpart with a spectrum consistent with that of a compact jet, as expected for LHS, was also reported by \\cite{Brocksopp2009}. All these properties triggered a daily monitoring by RXTE in order to follow up the outburst evolution. In this paper we present spectral and time variability analysis of XTE J1752-223 using an unusually long, almost continuous observation ($\\sim 116$ ks) performed by RXTE during 26th, 27th and 28th October 2009 (MJD 55130-55132). The quality of this data set allows us to perform a detailed study of this source and compare its general behaviour with that shown by the prototypical BHB Cyg X-1 during LHS. % \\begin{figure} \\centering \\includegraphics[width= 8.8cm,height=6.5cm]{spettro.eps} \\caption{Our best fit to the combined XRT, PCA and HEXTE spectrum of XTE J1752-223. The model used consist of an interstellar absorption component, two Gaussian lines, a multicolour black-body disc, a broken power-law and a high energy cut-off (see text). Top panel: XRT, PCA and HEXTE spectra. Bottom panel: residuals plotted as signed contribution to the chi square.} \\label{fig:1752_sp} \\end{figure} ", "conclusions": "An unusually long RXTE observation of the X-ray transient XTE J1752-223 is presented in this paper. The quality of the data have allowed us to obtain high S/N energy spectrum, PDS, rms spectrum and time-lags for this new source. All the obtained results are consistent with a black hole binary in the hard state. In particular, we find a behaviour similar to that exhibited by Cyg X-1 during hard state, but XTE J1752-223 happens to be in a slightly harder state. However, we note that there are two important differences between these two systems: Cyg X-1 is so far a persistent black hole binary and it harbours a high-mass companion. XTE J1752-223 is a transient and probably harbours a low-mass donor. Future multi-wavelength campaigns will probably provide new clues to the fundamental properties of this new black hole candidate. \\vspace{1cm} \\noindent The research leading to these results has received funding from the European Community's Seventh Framework Programme (FP7/2007-2013) under grant agreement number ITN 215212 \\textquotedblleft Black Hole Universe\\textquotedblright. SM and TB acknowledge support to the ASI grant I/088/06/0. TMD, SM, TB and DP acknowledge hospitality during their visits to IUCAA (Pune)." }, "1003/1003.1067_arXiv.txt": { "abstract": "We have measured the sub-milli-arcsecond structure of 274 extragalactic sources at 24 and 43 GHz in order to assess their astrometric suitability for use in a high frequency celestial reference frame (CRF). Ten sessions of observations with the Very Long Baseline Array have been conducted over the course of $\\sim$5 years, with a total of 1339 images produced for the 274 sources. There are several quantities that can be used to characterize the impact of intrinsic source structure on astrometric observations including the source flux density, the flux density variability, the source structure index, the source compactness, and the compactness variability. A detailed analysis of these imaging quantities shows that (1) our selection of compact sources from 8.4 GHz catalogs yielded sources with flux densities, averaged over the sessions in which each source was observed, of about 1 Jy at both 24 and 43 GHz, (2) on average the source flux densities at 24 GHz varied by 20\\%--25\\% relative to their mean values, with variations in the session-to-session flux density scale being less than 10\\%, (3) sources were found to be more compact with less intrinsic structure at higher frequencies, and (4) variations of the core radio emission relative to the total flux density of the source are less than 8\\% on average at 24 GHz. We conclude that the reduction in the effects due to source structure gained by observing at higher frequencies will result in an improved CRF and a pool of high-quality fiducial reference points for use in spacecraft navigation over the next decade. ", "introduction": "The International Celestial Reference Frame (ICRF) was formally adopted as the fundamental celestial reference frame (CRF) by the International Astronomical Union (IAU) in 1997. The catalog includes precise astrometric positions of over 600 extragalactic compact radio sources distributed uniformly over the sky. These positions were determined from the analysis of thousands of dual-frequency S/X-band (2.3/8.4 GHz) Very Long Baseline Interferometry (VLBI) observational sessions recorded between 1979 and 1995. The frame itself was defined by the 212 highest-quality ``defining\" sources with typical position accuracies of 0.25 milli-arcseconds (mas), while the axes of the frame are accurate to 0.02 mas \\citep{MA:98}. In addition to the 212 defining sources, positions for 294 less observed ``candidate'' sources along with 102 ``other'' sources with excessive position variation, were also given to increase the density of the frame \\citep{MA:98}. Since its adoption, incremental updates to the catalog of sources have occurred in the form of two extensions to the ICRF \\citep{FEY:04} using hundreds of additional sessions along with improved analysis and modeling techniques. However, the positions of the original 212 defining sources have remained constant through these extensions. The compact extragalactic radio sources that comprise the ICRF have been the subject of extensive study since the inception of VLBI techniques. In the standard theory, \\citep[e.g.][]{BK:79}, the jet-like emission from quasars and active galactic nuclei is assumed to be powered by a central engine where energetic phenomena occur. The observed, frequency dependent, intrinsic structure of extragalactic radio sources typically consists of a flat spectrum ($S\\propto\\nu^\\alpha, \\alpha\\approx0$) unresolved core at the base of the jet where the optical depth is near unity ($\\tau \\approx 1$) and extended emission in the form of multiple steep spectrum ($\\alpha\\approx-0.5$ to $-$1.5) jet components. These components often move away from the core along the direction of the jet, sometimes at apparent superluminal speeds. The emission outside of the core has been shown to be extended on scales larger than the accuracy of the astrometric position measurements \\citep{FCF:96}. This extended emission, or intrinsic source structure, contributes to the uncertainty in the measured astrometric positions of sources that comprise the ICRF. In addition to extended emission, frequency dependent opacity effects can contribute to variations in the measured astrometric positions of extragalactic radio sources. In particular, opacity conditions in the region near the base of the jet can cause the measured core position to move inward along the jet direction as a function of increasing radio frequency of the observations. Such core shifts have been measured for some compact sources \\citep[e.g.][]{LOBANOV:98, KOVALEV:08}. Based upon the discussion above, there are potential advantages to observing astrometric sources at radio frequencies higher than the typical S/X band observations. First, the resolution of the observations is increased, thus improving the astrometric accuracy. Second, the effects due to intrinsic source structure may be reduced since the core and jet components have different spectral characteristics and the core is expected to be more dominant at high frequencies. We have undertaken a program to observe a number of extragalactic sources at K band (24 GHz) and Q band (43 GHz) using the 10 stations of the Very Long Baseline Array (VLBA). At these higher frequencies, only the VLBA provides the stability, imaging capabilities and frequency coverage to enable such a program. The long term goals of the program include: 1) developing a high-frequency CRF with a variety of applications including improved deep space navigation, 2) providing the astronomical community with an extended catalog of calibrator sources for VLBI observations at 24 and 43 GHz, and 3) studying the effects of the intrinsic structure of extragalactic sources to improve the astrometric accuracy of future high-frequency reference frames. A detailed discussion of the program and the astrometric results is contained in a companion paper \\citep[][hereafter Paper I]{LANYI:09}. In this paper, we concentrate on the imaging aspects of the program and the effects of observed source structure and variability on astrometric accuracy. Theoretically, an optimal CRF would be composed of sources with strong, non-variable, point-like emission. In reality, however, virtually all sources possess some structure and intrinsic variability in the emission over time. In addition, there is the potential for sudden flaring events even for normally quiescent sources. Thus it is highly desirable to characterize and monitor the nature of the sources used in a CRF through periodic VLBI imaging. The VLBA has previously been used to make simultaneous dual-frequency S/X-band (2.3/8.4 GHz) observations of a total of 389 ICRF sources \\citep{FCF:96,FC:97,FC:00}. To date, approximately $90\\%$ of the ICRF sources north of $-20^\\circ$ declination have been imaged at least once at both 2.3 and 8.4~GHz. Based on the initial work of \\cite{CHARLOT:90}, the database of VLBA X-band and S-band images was analyzed by \\cite{FC:97,FC:00} in order to quantitatively improve our understanding of the relationship between extended source structure and the astrometric positions determined from VLBI. Here we discuss our growing database of high-frequency images of potential extragalactic reference frame sources. We apply similar analysis techniques in an attempt to characterize the impact of extended source structure on the astrometric accuracy of the catalog of source positions obtained from our VLBA high-frequency data (Paper I). ", "conclusions": "In the previous section, we discussed results for several quantities that emerged from our analysis of the VLBA images, namely: the source flux density, the source structure index, the source compactness, and the time variability of each of these quantities. In this section, we compare the source structure index to the compactness in order to verify that the two quantities are related. In addition, we compare both the structure index and the compactness to two astrometric quantities derived in Paper I, namely: the formal position uncertainties and position variability of the sources. These comparisons were made to determine whether the reduced structure effects seen at high frequency correspond to more precise astrometric positions. \\subsection{Structure Index and Source Compactness \\label{SEC:SI_SC}} As an initial test of the source compactness, we compared $\\bar C$ to the source structure index, $SI$, at each of the two frequencies. Shown in Figure~\\ref{FIG:K_COMPACT_SI} are the distributions of the mean source compactness for the 274 sources imaged at K band separated in terms of maximum source structure index. Figure~\\ref{FIG:Q_COMPACT_SI} shows similar distributions for the 132 sources imaged at Q band. The three panels in each figure show $SI$ = 1, 2 and 3 sources, respectively. There are two sources at K band and one source at Q band with $SI=4$ that are not shown in the figures. In both figures, it is evident that the distributions in compactness among the three structure indices are quite different with the distribution for $SI = 1$ being strongly peaked at both frequencies, and the distributions broadening with increasing values of $SI$. Within each panel of Figures~\\ref{FIG:K_COMPACT_SI} and \\ref{FIG:Q_COMPACT_SI} the mean and median of the distribution are given. These values are also summarized in Table~\\ref{TAB:COMPACT_SI}. From the table we see that for both frequencies, the mean and median source compactness is directly related to structure index. The correspondence between $SI$ and $\\bar C$ is not unexpected, since both quantities provide an indication of the source structure as derived from the VLBI imaging. We also compared the compactness variability index to the source structure index and the results are shown in Figure~\\ref{FIG:K_COMPACT_VAR_SI}. Plotted in the figure is $\\sigma_C/\\bar{C}$ at K band as a function of each of the three structure indices, $SI$ = 1, 2 and 3. Recall that the compactness variability index was determined for only those sources observed in more than one session (235 sources at K band and 82 sources at Q band). The equivalent plot for the Q band data was not produced because of the greatly reduced number of sources. As in the case of the source compactness, the variability in the compactness shows a clear trend with $SI$, with the $\\sigma_C/\\bar{C}$ distribution being the most narrow for $SI = 1$ and the broadest for $SI=3$. The mean (median) values for each distribution are 0.04 (0.04), 0.08 (0.07) and 0.15 (0.14) for the $SI$ = 1, 2 and 3 sources, respectively. A variability index $\\sigma_C/\\bar{C} = 0$ indicates no variability in the source compactness from one session to the next. This relationship between the compactness variability index and the structure index suggests that the sources with the most structure as measured by $SI$ exhibit increased variability in their structure as measured by $\\sigma_C/\\bar{C}$. \\subsection{Structure Index, Compactness and Source Position Uncertainty \\label{SEC:SI_ACCUR}} The impact of source structure on the high-frequency CRF can be further studied by comparing structure index with the formal precision of the source positions comprising a potential CRF. The formal position uncertainties were taken from the K-band astrometric solution detailed in Paper I. For this solution, we used the CALC/SOLVE software package maintained by the NASA Goddard Space Flight Center (GSFC) to perform a least-squares astrometric solution for the K-band data. The 10 diurnal K-band experiments encompassed 82,334 measurements of bandwidth synthesis (group) delay and phase delay rate. Geodetic parameters estimated for each session include: station positions, 20-minute piecewise linear continuous troposphere zenith parameters, tropospheric gradients in the east--west and north--south directions, linear in time, estimated once per 24 hr session, quadratic clock polynomials to model the gross clock behavior, and 60-minute piecewise linear continuous clock parameters. Corrections for ionospheric refraction drawn from Global Positioning System (GPS) total electron content (TEC) maps were applied to K-band data as discussed in Paper I. Positions for sources having three or more measurements of group delay were the only global parameters that were estimated. The K-band catalog derived from the astrometric solution is comprised of positions and associated formal uncertainties for 268 sources. Because there were too few sources at Q-band (131) to separate into the four structure index categories, we chose only to compare the K-band uncertainties with structure index and compactness. An initial comparison of all 268 sources showed the position uncertainties plotted as a function of $SI$ category to be sensitive to a few outliers with relatively few observations. We, therefore, decided to restrict the comparison to sources with 100 or more group delay measurements. This same restriction was used by \\cite{FC:00} in a similar study performed at X band. Shown in Figure~\\ref{FIG:POS_UNCER_HIST_100} are the distributions of the formal position uncertainties in a) $\\alpha\\cos\\delta$ and b) $\\delta$ for 193 sources with 100 or more group delay measurements at K band. The distributions are separated into the three $SI$ categories 1, 2, and 3. There were two $SI = 4$ sources with greater than 100 delay measurements that are not shown. There is good agreement between the mean and median values suggesting little or no dependence on outliers. The results show no significant difference in the mean and median position uncertainties from one structure index category to the next. In addition, we find the mean and median position uncertainties in $\\delta$ are roughly twice the uncertainties in $\\alpha\\cos\\delta$ for all three $SI$ categories. In \\cite{FC:00}, it was shown that the mean and median X-band source position uncertainties in both $\\alpha\\cos\\delta$ and $\\delta$ increased regularly as a function of structure index from $SI$=1 to 4. In addition to the 100 group delay measurement restriction, \\cite{FC:00} {used the formal position uncertainties from \\cite{MA:98}, which were adjusted by the standard ICRF inflation factor described therein. In order to accurately compare our K-band values, we applied the same inflation factor (the root sum square (rss) of 1.5 times the formal uncertainty and 0.25 mas) to the K-band formal uncertainties. Table~\\ref{TAB:POS_UNCER} lists the mean and median position uncertainties at both X and K bands as a function of $SI$ category. Again we see that there is no increase in position uncertainty from structure index 1 to 3 at K-band unlike the earlier findings at X-band. More importantly, the table shows that for a similar number of sources, the inflated mean/median position uncertainties are smaller at K band than at X band for all structure index categories. It should be noted, that there have been significant improvements in both the number and quality of the VLBI observations since the construction of the ICRF, and the \\cite{FC:00} values reported in Table~\\ref{TAB:POS_UNCER} do not reflect the accuracy of the current X-band CRF. In addition to the comparison of the $SI$ and the formal uncertainties, we compared the mean source compactness ($\\bar C$) to the formal uncertainties. Recall that the maximum possible value of the source compactness $C = 1.0$ indicates that all of the source flux density is contained within one synthesized beam. The mean compactness is just an average over all of the sessions in which the source was imaged. Figure~\\ref{FIG:POS_COMPACT_SCATTER} shows the formal uncertainties in the position versus the source compactness. The plotted formal position uncertainty is the rss of the uncertainties in $\\alpha\\cos\\delta$ and $\\delta$. The figure also shows the structure index for each source by point color and type, with green circles, blue triangles, and red squares indicating maximum $SI$ values of 1, 2, and 3 respectively. Figure \\ref{FIG:POS_COMPACT_SCATTER} does not show a clear trend in the formal position uncertainty versus compactness. This is to be expected since it was previously shown that the $\\bar C$ and $SI$ are correlated, and no significant relationship was found for the position uncertainty as a function of $SI$ (e.g. Figure~\\ref{FIG:POS_UNCER_HIST_100}). \\subsection{Structure Index, Compactness, and Source Position Stability \\label{SEC:SI_STAB}} Another useful comparison can be made between the stability of the source positions over time and the source structure as traced by either the structure index or the source compactness. Time variation of the astrometric coordinates of CRF sources has previously been attributed to variability in the intrinsic source structure \\citep[c.f.][]{CHARLOT:94, CHARLOT:02, FEK:97, FC:00}. Thus, it is natural to search for any relationship between such variations and the structure index at higher frequencies. A measure of source position stability is the weighted root-mean-square (wrms) position variation. These variations were obtained from a series of K-band astrometric solutions described in Paper I. In these solutions, a fraction of the sources were treated as local parameters (i.e. an estimate of the position was derived for each session in which the source was observed) with the remaining sources treated as global. The estimation of source position stability was limited to 88 sources that were observed in five or more of the 10 VLBA sessions. Five was considered a sufficient number of sessions (position estimates) per source to derive reliable statistics. From these solutions, a time series of source positions was generated for each source, and from these time series the wrms position variations were computed. Shown in Figure~\\ref{FIG:TS_HIST} are the distributions of the wrms position variations in $\\alpha\\cos\\delta$ (a) and $\\delta$ (b), respectively. The figure also lists the mean and median position variations for each distribution. In $\\alpha\\cos\\delta$, the mean and median of the distribution are 0.14 and 0.12 mas, respectively. For $\\delta$ the mean and median are roughly twice as large at 0.30 and 0.25 mas, respectively. As stated in Paper I, the larger wrms position variations in $\\delta$ are likely due to the combination of ionospheric/tropospheric effects and network geometry, specifically the lack of long north--south baselines for the VLBA. It would have been desirable to separate the wrms position variations into structure index categories, and to plot the distributions as a function of $SI$ as was done in $\\S$\\ref{SEC:SI_ACCUR}. Because the structure index is quantized into four categories, there were simply too few sources observed in five or more sessions to determine reliable statistics for the separate $SI$ categories. Instead, in Figure~\\ref{FIG:TS_COMPACT_SCATTER}, we plot the wrms variation against the mean source compactness previously described in $\\S$\\ref{SEC:COMPACT}. The wrms variation plotted is the rss of the variations in $\\alpha\\cos\\delta$ and $\\delta$. In addition, the structure index for each source is shown by point type with green circles, blue triangles, and red squares indicating maximum $SI$ of 1, 2, and 3 respectively. Figure \\ref{FIG:TS_COMPACT_SCATTER} does not suggest a clear trend in the compactness and wrms position stability. While it does appear that there are a number of sources clustered in a region of high compactness and low wrms variations, there are also sources that are less compact with equally low wrms variations. A more quantitative comparison between the $SI$ and the position stability will have to wait for additional high-frequency observations." }, "1003/1003.1298_arXiv.txt": { "abstract": "Motivated by a recent astrophysical measurement of the pressure of cold matter above nuclear-matter saturation density~\\cite{Ozel:2010fw}, we compute the equation of state of neutron star matter using accurately calibrated relativistic models. The uniform stellar core is assumed to consist of nucleons and leptons in beta equilibrium; no exotic degrees of freedom are included. We found the predictions of these models to be in fairly good agreement with the measured equation of state. Yet the {\\sl Mass-vs-Radius} relations predicted by these same models display radii that are consistently larger than the observations. ", "introduction": " ", "conclusions": "" }, "1003/1003.4368_arXiv.txt": { "abstract": "{The star HD~49385 is the first G-type solar-like pulsator observed in the seismology field of the space telescope CoRoT. The satellite collected 137 days of high-precision photometric data on this star, confirming that it presents solar-like oscillations. HD~49385 was also observed in spectroscopy with the NARVAL spectrograph in January 2009.} {Our goal is to characterize HD~49385 using both spectroscopic and seismic data.} {The fundamental stellar parameters of HD~49385 are derived with the semi-automatic software VWA, and the projected rotational velocity is estimated by fitting synthetic profiles to isolated lines in the observed spectrum. A maximum likelihood estimation is used to determine the parameters of the observed p modes. We perform a global fit, in which modes are fitted simultaneously over nine radial orders, with degrees ranging from $\\ell=0$ to $\\ell=3$ (36 individual modes).} {Precise estimates of the atmospheric parameters ($T\\ind{eff}$, [M/H], $\\log g$) and of the \\vsini\\ of \\cible\\ are obtained. The seismic analysis of the star leads to a clear identification of the modes for degrees $\\ell=0,1,2$. Around the maximum of the signal ($\\nu\\simeq1013\\,\\mu$Hz), some peaks are found significant and compatible with the expected characteristics of $\\ell=3$ modes. Our fit yields robust estimates of the frequencies, linewidths and amplitudes of the modes. We find amplitudes of $\\sim5.6\\pm0.8$ ppm for radial modes at the maximum of the signal. The lifetimes of the modes range from one day (at high frequency) to a bit more than two days (at low frequency). Significant peaks are found outside the identified ridges and are fitted. They are attributed to mixed modes.} {} ", "introduction": "In the Sun, oscillations are excited by the turbulent motions in the outer part of the external convective envelope and are further propagated into the interior of the star. The study of these oscillations has yielded constraints on the inner structure of the Sun, allowing us to estimate the sound speed and density profiles (\\citealt{2003ApJ...591..432B}, \\citealt{2001ApJ...555L..69T}), the position of the base of the convective zone (\\citealt{1991ApJ...378..413C}), and the rotation profile (\\citealt{2003ARA&A..41..599T}, \\citealt{2008A&A...484..517M}). However, the very low amplitude of these oscillations (a few ppm in photometry) makes it very challenging to detect and analyze them in other stars than the Sun. Achieving a better understanding of the interiors of solar-like pulsations is one of the main objectives of the space mission \\corot\\ (\\textbf{Co}nvection, \\textbf{Ro}tation and planetary \\textbf{T}ransits). \\corot\\ is a space telescope performing high-precision photometry over quasi-uninterrupted long observing runs (\\citealt{Baglin06}). Solar-like oscillations have already been studied in several other stars with \\corot\\ data. The analyses of these stars have encountered difficulties identifying the degrees of the modes, either because of a too low signal-to-noise ratio (HD~175726: \\citealt{2009A&A...506...33M}, HD~181906: \\citealt{2009A&A...506...41G}), or because of a too short lifetime of the modes, inducing large mode linewidths (HD~49933: \\citealt{2008A&A...488..705A}, HD~181420: \\citealt{HD181420} ). The star HD 49385 is the first G-type solar-like pulsator observed in the seismology field of \\corot. It is cooler than the solar-like pulsators previously analysed with \\corot\\ data, and probably more evolved (at the end of the Main Sequence or shortly after it). The choice of \\cible\\ as a \\corot\\ target has motivated us to lead spectroscopic observations, performed with the NARVAL spectrograph at the Pic du Midi Observatory. The fundamental parameters of \\cible\\ are derived from these observations, as described in Sect. \\ref{sect_param}. The photometric observations with \\corot\\ are presented in Sect. \\ref{sect_obs}. Section \\ref{sect_rot} presents the study of the low-frequency part of the power spectrum, in search of a signature of the stellar rotation. The extraction of p-mode parameters is described in Sect. \\ref{sect_analysis}, and Sect. \\ref{sect_concl} is dedicated to conclusions. ", "conclusions": "} The star HD~49385 was characterized from both \\corot\\ seismic observations and NARVAL spectroscopic data. The atmospheric parameters of the star were derived and the 137-day-long photometric time series was analyzed. A clear series of peaks associated to p modes was detected in the power spectrum around 1 mHz. Up to now, the \\corot\\ solar-like pulsators presented some ambiguities in the identification of ridges corresponding to different degrees $\\ell$. Here three very clear ridges appear in the \\'echelle diagram, which were readily identified as $\\ell=0$, 1 and 2 modes. The $\\ell=2$ ridge appears clearly distinct from the $\\ell=0$ ridge in the \\'echelle diagram. Furthermore, two peaks, part of a fainter ridge, were found to be significant and compatible with the characteristics we expect for $\\ell=3$ modes in \\cible. This probably constitutes the first photometric detection of $\\ell=3$ modes in a solar-like pulsating star (other than the Sun). We performed a global fit over nine radial orders with degrees ranging from $\\ell=0$ to $\\ell=3$ modes (simultaneous fit of 36 individual modes). We obtained precise estimates of the mode frequencies (uncertainty of about $0.2\\,\\mu$Hz for radial modes at the maximum of the signal). The obtained value of the maximum amplitude for $\\ell=0$ modes ($5.6\\pm0.8$ ppm) is consistent with the estimate deduced from \\cite{samadi07}, unlike the other \\corot\\ solar-like pulsators (\\textit{e.g.} \\citealt{2008A&A...488..705A}). We found no evidence of a rotational splitting of the modes. This can be explained either by a small inclination angle or by a low rotational velocity (inducing a small rotational splitting compared to the linewidth of the modes). In passing we stress that even for very slow rotators (rotational splitting as low as a few $\\mu$Hz) the $m\\neq0$ components of non-radial multiplets are not expected to be symmetrical with respect to the $m=0$ component. We proposed a simple way of treating this assymetry in future analyses of solar-like pulsators. The p modes of \\cible\\ were found to have lifetimes ranging from about one day to two days, \\textit{i.e.} somewhat shorter than the mode lifetimes in the Sun, but significantly larger than those of the previously observed \\corot\\ solar-like targets. This explains why the spectrum of \\cible\\ is much clearer than for previous \\corot\\ pulsators, for which the large mode linewidths made it harder to separate the $\\ell=2$ ridge from the $\\ell=0$ ridge. The results obtained for \\cible\\ confirm that the linewidths of the modes in G-type pulsators are smaller than those of F-type pulsators, making their analysis easier. The very high quality of the spectrum also enabled us to detect significant peaks outside the identified ridges. These peaks were found to be compatible with mixed modes, whose presence can be expected in the spectrum of evolved objects such as \\cible. The existence of mixed modes in avoided crossing can explain some specific behaviors we observe in the low-frequency eigenmodes. In particular we found that the $\\ell=1$ ridge is distorted compared to the $\\ell=0$ ridge at low frequency. \\cite{deheuvels10} showed that this type of distortion could be associated with a low-degree mixed mode in avoided crossing. The identification of the mode $\\pi_1$ as an $\\ell=1$ mixed mode would be consistent with the observed pattern. Other features like the $\\ell=2$ mode found to overlap the $\\ell=0$ mode (around 855 $\\mu$Hz) might also result from avoided crossing phenomena. This needs to be further investigated in the seismic interpretation of this star." }, "1003/1003.3254_arXiv.txt": { "abstract": "{The stability of the color flavor locked phase in the presence of a strong magnetic field is investigated within the phenomenological MIT bag model, taking into account the variation of the strange quark mass, the baryon density, the magnetic field, as well as the bag and gap parameters. It is found that the minimum value of the energy per baryon in a color flavor locked state at vanishing pressure is lower than the corresponding one for unpaired magnetized strange quark matter and, as the magnetic field increases, the energy per baryon decreases. This implies that magnetized color flavor locked matter is more stable and could become the ground state inside neutron stars. The mass-radius relation for such stars is also studied.} \\PACS{26.60.-c, 21.65.Qr, 26.60.Kp, 04.40.Dg} ", "introduction": "\\label{sec1} The internal composition of neutron stars as well as the real nature of the ground state of matter moves through interconnected avenues. Bodmer~\\cite{Bodmer:1971we} and Witten ~\\cite{Witten:1984rs} suggested that strange quark matter (SQM) could be a stable phase of nuclear matter. This exciting result continues being a conjecture because presently it is impossible to perform laboratory experiments that confirm it. Nevertheless, this issue has attracted great attention in the astrophysical context and many works have been devoted to study the properties of the equation of state (EoS) of SQM and its connection with strange star or neutron star observables. In particular, the existence of a more compact form of matter could be a plausible explanation for the still unexplained observation of sources of gamma-$\\gamma$ rays bursts. Furthermore, studies of the superconductor phases of the QCD suggest that the ground state of matter could be a superconductor phase, being a compact object the natural scenario of this phase transition. The pioneer studies of the pairing interaction of the dense matter appeared more than thirty years ago~\\cite{Barrois:1977xd,Bailin:1983bm}. Under certain conditions, SQM could undergo a phase transition to a superconductor phase, in particular at high densities and low temperature. The most symmetric phase among these phases is the so-called color flavor locked (CFL) state~ \\cite{Alford:1998mk,Schafer:1999fe,Shovkovy:1999mr,Alford:2001zr,Rajagopal:2000ff,Alford:2002kj,Alford:2007xm,Alford:2004pf}. On the other hand, there is not doubt that the role of the magnetic field in astrophysical scenarios is very important. Pulsars, magnetars, neutron stars, the emission of intense sources of X-rays could be associated to sources with intense magnetic fields around $10^{13}-10^{15}$~G or even higher fields~\\cite{duncan,kouve}. Furthermore, the magnetic field intensity may vary significantly from the surface to the center of the source and theoretical estimates indicate that fields as high as $10^{19}$~G could be allowed~\\cite{dong}. The relevance of the magnetic field in color superconductivity has been studied in refs.~\\cite{Alford:1999pb,Gorbar:2000ms,Ferrer:2005vd,Ferrer:2006vw,Ferrer:2006ie,Ferrer:2007iw,Fukushima:2007fc,Noronha:2007wg,Alford:2010qf}. These papers have tackled, among other aspects, the modification of the pairing pattern by the external field, the formation of a gluon condensate at certain field strengths, and the boost of the applied field due to the back reaction of the color superconductor. These results support the idea that the magnetic field enhances color superconductivity. It has also been shown that magnetic fields in neutron stars with color superconducting cores are stable on time scales comparable with the age of the Universe~\\cite{Alford:1999pb,Alford:2010qf}. Thus, seeking for pulsars which do not diminish their magnetic field as the star spins down could help to find evidences of the existence of a superconductor phase of quark matter inside compact stars. Following this line of research, it is therefore worthwhile to study astrophysical observables, such as the mass-radius relation of quark stars, in a magnetized CFL phase. In ref.~\\cite{Felipe:2008cm}, it was shown that magnetized strange quark matter (MSQM) becomes more stable than unpaired SQM. If the CFL superconductor phase is more stable and bound than unpaired SQM at finite density, one expects stable configurations of quark stars more compact than in the unpaired phase. These objects would be self-bound and their masses would scale with the radius as $M \\sim R^{3}$, in contrast to neutron stars which have masses that decrease with increasing radius ($M \\sim R^{-3}$) and are bound by gravity~\\cite{itoh,Hansel}. Thus, self-bound stars could be consistent with small-radius compact objects~\\cite{Lugones:2002va,Lugones:2002zd,Quan,Nice,Xu:2006qh}. The aim of this paper is to study the role of the magnetic field in the CFL phase within the MIT bag model of confinement. Our intention is to show how the magnetic field can influence the stability of the phase and also its implications for the mass-radius relation generated by configurations where the deconfined matter is in the magnetized CFL phase. For the sake of simplicity, we shall assume that the gluonic contribution to the magnetic field inside the CFL phase is negligible. Due to the mixture of the photon field $A_{\\mu}$ and the eighth component of the gluon field $G_{\\mu}^8$, the `rotated' electromagnetic field is $\\tilde{A_{\\mu}}=A_{\\mu}\\cos \\theta -G_{\\mu}^8 \\sin \\theta$. The corresponding electromagnetic coupling constant is $\\tilde{e}=e\\cos\\theta$, where the mixing angle $\\theta$ depends on the gap structure and is given by $\\cos\\theta=g/\\sqrt{e^2/3+g^2}$ ($g$ is the QCD coupling constant) for the CFL phase~\\cite{Alford:1999pb,Gorbar:2000ms}. Since the rotated photon is massless, the magnetic field $\\tilde{B}$ inside the CFL superconducting state is not screened. Moreover, in the region of interest, $e \\ll g$ so that $\\cos \\theta \\sim 1$ and one can consider that the magnetic field strengths inside and outside the CFL core are approximately equal, i.e. $\\tilde{e}\\tilde{B} \\simeq eB$~\\cite{Fukushima:2007fc}. An important issue when studying the stability of quark matter in compact stars is the theoretical treatment of the neutrality conditions~\\cite{Alford:2002kj,Buballa:2005bv}. Besides electromagnetic neutrality, color neutrality must be enforced. To guarantee the latter in the CFL phase, the chemical potentials $\\mu_3$ and $\\mu_8$ coupled to the color charges $T_3=\\text{diag}\\,(1/2,-1/2,0)$ and $T_8=\\text{diag}\\,(1/3,1/3,-2/3)$ should be chosen such that the $T_{3,8}$ densities vanish~\\cite{Alford:2002kj}. The chemical potential for each quark $(i=u,d,s)$ is then specified by its electric and color charges, $\\mu_i=\\mu_{B}-Q \\mu_e + T_3 \\mu_3 + T_8 \\mu_8$, where $\\mu_B$ is the baryon chemical potential and $Q=\\text{diag}\\,(2/3,-1/3,-1/3)$. As it turns out, for the range of parameters considered here, to wit $eB < \\mu_{B}^2$, one can show that $\\mu_3 \\simeq 0$ and $\\mu_8 \\simeq -m_s^2/(2\\mu_B)$ (cf. fig.~\\ref{chempot} and our discussion below eq.~\\eqref{Nequality}). \\begin{figure}[t] \\begin{center} \\includegraphics[width=8.0cm]{Fig1.eps} \\caption{Chemical potentials in the CFL phase for $B=0$ and $B=5\\times10^{18}$~G. The curves are shown for $\\Delta= 50$~MeV and $B_{\\rm bag}=75$~MeV/fm$^3$.\\label{chempot}} \\end{center} \\end{figure} In this simple framework, we shall study the behavior of the system with the variation of the relevant parameters: the bag parameter $B_{\\rm bag}$, the strange quark mass $m_s$, the baryon density $n_B$, the gap parameter $\\Delta$ and the magnetic field $B$. The possible mass-radius configurations of magnetized CFL stars are then obtained by solving the Tolman-Oppenheimer-Volkoff (TOV) equations. We obtain configurations of stable stars with smaller radii, which are allowed due to the compactness of matter and the presence of a strong magnetic field, since the energy per baryon at vanishing pressure is lower in this case. The paper is organized as follows. In sect.~\\ref{sec2} we briefly review the CFL phase properties in the presence of a magnetic field within the MIT bag model. In sect.~\\ref{sec3}, the stability windows of CFL in the presence of a magnetic field are obtained varying the relevant input parameters of the model. Section~\\ref{sec4} is devoted to the study of the mass-radius relation for CFL matter by numerically solving the TOV equations. Finally, our conclusions are given in sect.~\\ref{sec5}. \\begin{figure*}[t] \\begin{minipage}[t]{0.45\\linewidth} \\includegraphics[width=8.0cm]{Fig2a.eps} \\end{minipage} \\hspace{0.3cm} \\begin{minipage}[t]{0.45\\linewidth} \\includegraphics[width=8.0cm]{Fig2b.eps} \\end{minipage} \\caption{The energy per baryon as a function of the baryon density for the CFL phase without magnetic field (left plot) and for magnetized CFL with $B=5\\times 10^{18}$~G (right plot). We take $B_{\\rm bag}=75$~MeV/fm$^3$ and $\\Delta=50, 100$~MeV. For comparison, the SQM and MSQM cases are also shown. The horizontal dotted line corresponds to $\\left. E/A\\right|_{^{56}\\text{Fe}} \\simeq 930$~MeV.\\label{EN}} \\end{figure*} \\begin{figure*}[t] \\begin{minipage}[t]{0.45\\linewidth} \\includegraphics[width=8.0cm]{Fig3a.eps} \\end{minipage} \\hspace{0.3cm} \\begin{minipage}[t]{0.45\\linewidth} \\includegraphics[width=8.0cm]{Fig3b.eps} \\end{minipage} \\caption{The energy per baryon versus pressure for the CFL phase without magnetic field (left plot) and for magnetized CFL with $B=5\\times 10^{18}$~G (right plot). We take $B_{\\rm bag}=75$~MeV/fm$^3$ and $\\Delta=50, 100$~MeV. The SQM and MSQM curves are also depicted. The horizontal dotted line corresponds to $\\left. E/A\\right|_{^{56}\\text{Fe}} \\simeq 930$~MeV.\\label{EP}} \\end{figure*} ", "conclusions": "\\label{sec5} In the present work we have investigated the stability of the magnetized CFL phase within the phenomenological MIT bag model. The study was performed taking into account the variation of the strange quark mass, the baryon density, the magnetic field, the gap and the bag parameters. We found that a strongly magnetized CFL state of strange matter is indeed more stable than the non-magnetized CFL one. We have also shown that magnetized CFL matter is more stable than unpaired magnetized SQM in the range of strong magnetic fields typically expected in compact stars and for a wide range of values of the gap of the QCD Cooper pairs. As the pairing gap increases, the stability windows of the CFL phase is enlarged, but at the same time the EoS becomes stiffer and, consequently, the maximum mass and radii values of stable stellar configurations get larger~\\cite{Lugones:2002zd}. On the other hand, an increase of the bag parameter $B_\\text{bag}$ would lead to smaller values of $M_\\text{max}$ and $R_\\text{max}$. Although our results have been obtained in a simplified framework of the MIT bag model, it is worth pointing out that relaxing our simplifying assumptions would not significantly change our qualitative conclusions. As illustrated in fig.~\\ref{tov} and tables~\\ref{table2} and~\\ref{table3}, the derived mass-radius relation and the relevant macroscopic parameters of compact stars are indeed modified in the presence of a strong magnetic field. In particular, the derived EoS for the magnetized CFL phase turns out to be softer (when compared to the EoS in the absence of a magnetic field), thus allowing for the existence of very compact stable configurations of strange quark stars composed by deconfined matter in a CFL phase. The magnetic field in these compact objects will be frozen over a time scale comparable with the age of our Universe. This feature enriches their phenomenology and distinguishes canonical neutron stars from those with color superconducting quark cores. The variation of the masses and radii between the non-magnetized and magnetized CFL phases are nevertheless rather small. Thus, it seems difficult to prove the existence of color superconductivity of magnetized quark matter by measurements of the $M(R)$ relation alone. More promising approaches are those based on the effects of the magnetic field in transport properties, such as neutrino emission~\\cite{Reddy:2002xc}, bulk and shear viscosity~\\cite{Madsen:1999ci} and glitches~\\cite{Alford:2000ze}. There are several ways in which our study could be extended further. In our analysis we did not take into account perturbative QCD corrections to the equation of state and we have neglected a possible dependence of the gap parameter on the density and the magnetic field. It would be important to see how robust our conclusions are against these and other possible corrections. Finally, it would also be interesting to extend our study to hybrid compact stars with a nuclear matter crust and a quark matter core in the presence of a strong magnetic field, pursuing for instance a model-independent phenomenological approach as the one taken in ref.~\\cite{Alford:2004pf}. \\begin{acknowledgement} We are grateful to J. Horvath, E.~J.~Ferrer and V.~de la Incera for very useful comments and suggestions. The work of R.G.F. was supported by \\emph{Funda\\c{c}\\~{a}o para a Ci\\^{e}ncia e a Tecnologia} (FCT, Portugal) through the project CFTP-FCT UNIT 777, which is partially funded through POCTI (FEDER). The work of A.P.M. has been supported by \\emph{Ministerio de Ciencia, Tecnolog\\'{\\i}a y Medio Ambiente} under the grant CB0407 and the ICTP Office of External Activities through NET-35. A.P.M. also acknowledges TWAS-UNESCO PCI Programme for financial support at CBPF-Brazil. \\end{acknowledgement}" }, "1003/1003.4018_arXiv.txt": { "abstract": "Models of galaxy formation invoke the major merger of gas-rich progenitor galaxies as the trigger for significant phases of black hole growth and the associated feedback that suppresses star formation to create red spheroidal remnants. However, the observational evidence for the connection between mergers and active galactic nucleus (AGN) phases is not clear. We analyze a sample of low-mass early-type galaxies known to be in the process of migrating from the blue cloud to the red sequence via an AGN phase in the green valley. Using deeper imaging from SDSS Stripe 82, we show that the fraction of objects with major morphological disturbances is high during the early starburst phase, but declines rapidly to the background level seen in quiescent early-type galaxies by the time of substantial AGN radiation several hundred Myr after the starburst. This observation empirically links the AGN activity in low-redshift early-type galaxies to a significant merger event in the recent past. The large time delay between the merger-driven starburst and the peak of AGN activity allows for the merger features to decay to the background and hence may explain the weak link between merger features and AGN activity in the literature. ", "introduction": "\\label{sec:intro} In 1988, \\citeauthor{1988ApJ...325...74S} proposed a direct link between major mergers and quasar activity: specifically, that major mergers between gas-rich galaxies fuel a substantial starburst; eventually some of the gas reaches the black hole, triggering an accretion event; and the resulting energy output from the quasar, such as radiation and outflows, sweeps up the remaining gas, thus suppressing further star formation and creating a passively evolving spheroidal remnant. The hierarchical nature of the $\\Lambda$CDM cosmology suggests that mergers of haloes and galaxies are one of the main drivers of evolution. Inspired by the \\cite{1988ApJ...325...74S} scenario, simulations of individual galaxy mergers show that major mergers can radically transform the progenitors by destroying any inbound disks and fueling major starbursts and heating part of the gas content \\citep[\\textit{e.g.},][]{1992ApJ...393..484B,1996ApJ...471..115B, 2006MNRAS.373.1013C}. More recent works have added the second component by invoking some version of quasar feedback to suppress star formation and establishing various observed scaling relations \\citep[\\textit{e.g.},][]{2005Natur.433..604D,2005ApJ...620L..79S, 2005MNRAS.361..776S, 2006ApJS..163....1H, 2008ApJS..175..390H, 2008ApJS..175..356H, 2008MNRAS.391..481S, 2009ApJ...690..802J}. This scenario, in which mergers trigger starbursts and AGN phases, now underpins the current generation of theoretical models as a framework for the co-evolution of galaxies and black holes. However, the observational evidence connecting these various phases is not clear. Many studies find that incidence of merger features in AGN host galaxies does not seem to be significantly enhanced over the general population \\citep[\\textit{e.g.},][]{1998ApJ...496...93D, 1998ApJS..117...25M, 2001AJ....122.2243S, 2007ApJ...660L..19P, 2009MNRAS.397..623G, 2009ApJ...691..705G, 2010MNRAS.401.1552D}, while a minority find some evidence for such an enhancement \\citep[\\textit{e.g.},][]{1986ApJ...311..526H,1996AJ....111..696K,2008ApJ...679.1047K, 2008ApJ...674...80U}. At the same time, recent observational work has questioned the degree to which mergers are involved in both the triggering of starbursts and the subsequent quenching of star formation and formation of red spheroidal galaxies, as the rate of major mergers accounts for all the mass growth in red galaxies since $z \\sim 1$ \\citep{2007ApJ...665L...5B, 2009ApJ...697.1369B}. \\begin{figure*} \\begin{center} \\includegraphics[angle=270, width=\\textwidth]{Final_Image_Comparison_1_lowres.ps} \\caption{Comparison of regular SDSS images to Stripe 82 images of four early-type galaxies that clearly show features indicating a recent merger in the deeper data. In the \\textit{top} row, we show the regular SDSS $gri$-composite images, while in the \\textit{bottom} row, we show the corresponding deeper Stripe 82 co-added $r$-band images. Features that are readily apparent in the Stripe 82 images are not visible in the regular SDSS images.\\label{fig:sdss_vs_stripe82_merger}} \\includegraphics[angle=270, width=\\textwidth]{Final_Image_Comparison_2_lowres.ps} \\caption{Same as Figure \\ref{fig:sdss_vs_stripe82_merger}, but for early-type galaxies that do not show any merger features in the corresponding deeper Stripe 82 images. \\label{fig:sdss_vs_stripe82_nomerger}} \\end{center} \\end{figure*} We decided to test this evolutionary scenario using a sample of spheroidal galaxies that are known to be in the process of migrating from the blue cloud to the red sequence via an AGN phase. We can therefore identify the same population in various evolutionary phases during the galaxy transformation process. The migration of early-type galaxies has been traced for low-mass early-type galaxies by \\cite{2007MNRAS.382.1415S} using the MOSES sample\\footnote{Morphologically Selected Early-types from Sloan, see \\cite{2007MNRAS.382.1415S}}. The objects along this sequence are an ideal laboratory to test whether merging activity is in any way related to both the shutdown of star formation -- leading to the migration from blue to red -- and the fueling of the central black hole, which may also be involved in creating a red, passively evolving remnant. In this Letter, we test the role of mergers by obtaining significantly deeper imaging data for early-types along the evolutionary sequence of \\cite{2007MNRAS.382.1415S}. The original SDSS images used for the morphological classification by eye performed by \\cite{2007MNRAS.382.1415S} do not reach sufficiently deep surface brightness limits to detect features indicating a recent merger, such as shells, tidal tails and other large-scale disruptions. If deeper imaging data of the same objects reveals such features, a link is established between merging activity on the one hand, and the shutdown of star formation and black hole fueling on the other. We obtain such deeper images from SDSS Stripe 82. ", "conclusions": "The decline in merger fraction with age illustrated in Figure \\ref{fig:f_merger} has strong implications for the connection between mergers on the one hand and starbursts, black hole accretion and galaxy aging, on the other. We now discuss the implications of the results in Figure \\ref{fig:f_merger}: \\subsection{Mergers trigger Starbursts and the Migration from the Blue Cloud to the Red Sequence} At least half of all blue early-type galaxies show evidence for a recent, significant merger in deeper imaging data. Given that the parent sample of blue early-type galaxies from MOSES \\citep{2007MNRAS.382.1415S} was selected for early-type morphology in the shallower regular SDSS images, it is likely that the MOSES sample is missing some objects with young ages ($\\sim100$ Myr or less) because they appear disturbed even in the regular SDSS images. This effect would push the merger fraction well above 50\\% and may even approach 100\\%. This in turn implicates a merger as the trigger for the starburst episode at the beginning of the evolutionary sequence for most, if not all early-types migrating from the blue cloud to the red sequence. The nature of the merger signature seen in low-redshift early-type galaxies is a question that remains to be investigated. Do the features seen in the deeper images in Figure \\ref{fig:sdss_vs_stripe82_merger} necessarily imply a major merger (\\textit{i.e.,} a merger with a mass ratio of less than 3:1)? Or can minor mergers produce a similar appearance? Numerical simulations \\citep[\\textit{e.g.},][]{2008ApJ...684.1062F} suggest that mergers with relatively large mass mass ratios can still produce prominent tidal features. Detailed simulations of mergers and simulated observations mimicking available data are needed to decide whether more information on the nature of the merger can be inferred from our result. \\subsection{Mergers trigger Black Hole Accretion} Since the Seyfert AGN early-types are consistent with being the post-starburs descendants \\citep{2007MNRAS.382.1415S} of the star-forming early-types in the blue cloud, this implies that the merger that triggered the starburst could well have triggered the AGN phase $\\sim$500 Myrs later. The time delay between the merger event and black hole feeding could be due to the time required for material to lose sufficient angular momentum to reach the black hole. Alternatively, it may simply be the post-starburst environment, with supernovae and strong stellar winds, that create conditions favorable for accretion onto the black hole. Our results raise the question of whether the sequence of events that plays out in nature can be accurately described as a ``merger trigger for AGN.\" It is clear that the time lag between the merger event to the AGN phase is substantial. The merger itself need not be not directly responsible for the delayed AGN phase and in that sense, mergers may not trigger AGN phases. Still, it seems clear that the AGN phase is most likely a consequence of the merger, albeit via a number of intermediate steps. The intermediate steps between the merger and accretion do not remove the causal link between the two events. The question of the merger trigger for AGN phases in early-type galaxies may thus be resolved in an indirect way. From a broader perspective, we can also conclude that the evolutionary sequence found by \\cite{2007MNRAS.382.1415S} is triggered by mergers. That is, mergers do trigger the migration of low-mass early-type galaxies from the blue cloud to the green valley where black hole feedback processes may be at work advancing this migration \\citep{2009MNRAS.396..818S}. \\subsection{A Fundamental Limit to tracing Mergers?} Another intriguing result of our study is that the rapid decline of the merger fraction to the background level fundamentally limits our ability to link merger signatures over times longer than $\\sim500$ Myr. Even though the merger features may survive for much longer, they become indistinguishable from the background level due to numerous dry and minor mergers. Paradoxically, deeper imaging may not extend the detectability or merger feature further in time, as the background merger fraction will also be enhanced in deeper images. There may well be an optimal image depth for studying merger features that maximizes the contrast between background and fading merger features. In order to overcome this limit, a reliable method for separating features due to a recent major mergers and feature induced by the background merger rate would be required.\\\\ We note that the results presented here are for early-type galaxies \\textit{only}. The majority (up to 90\\%) of active nuclei in the local Universe are hosted by late-type galaxies which follow an evolutionary pathway different from the early-types \\citep{2010ApJ...711..284S}." }, "1003/1003.2408_arXiv.txt": { "abstract": "\\noindent We consider the coherent state approach to non-commutativity, and we derive from it an effective quantum scalar field theory. We show how the non-commutativity can be taken in account by a suitable modification of the Klein-Gordon product, and of the equal-time commutation relations. We prove that, in curved space, the Bogolubov coefficients are unchanged, hence the number density of the produced particle is the same as for the commutative case. What changes though is the associated energy density, and this offers a simple solution to the transplanckian problem. ", "introduction": "\\noindent In recent years, we have seen many proposals aimed to quantize consistently the gravitational field from fundamental principles. From a phenomenological point of view, a more modest approach consists in introducing reasonable modifications to quantum field theory and look for observable consequences in black holes or inflationary models. For example, one can construct a theory where the dispersion relations depart from linearity above a certain energy scale, thus breaking local Lorentz invariance. This approach is motivated by analogue models of gravity in condensed matter systems \\cite{NuovoCimento}, by deformations of the Lorentz group \\cite{AmelinoCamelia}, or by tensor-vector models of gravity \\cite{Jacob}. Modified dispersion relations were considered in the context of renormalization, and particle production in curved spacetime \\cite{MaxMazz,max1,MaxMazz2,max2}. In the latter case, the common result is that the thermal spectrum seen from an accelerated detector or from an asymptotic observer on a black hole background is only marginally affected by non-linear dispersion relations \\cite{UnruhOrig,UnruhShuz,UnruhMax}. There exists an alternative proposal, based on a new symmetry of the path integral duality \\cite{paddy}. In this case, the modification directly brings a minimal length in the propagator, which becomes finite in the ultraviolet regime. Starting from different hypothesis, the same propagator was reconsidered also in other works \\cite{pepe}. In both cases, the form of the field modes, associated to the modified propagator is unknown, therefore it is difficult to evaluate exactly some effects. In contrast to these proposals, here we would like to study how quantum field theory is modified when spacetime has an intrinsic non-commutative structure. This topic has been intensively investigated by assuming that non-commutative effects in field theory are implemented by replacing the ordinary product among functions with the so-called star-product \\cite{NCclassic,NCclassic2,starproduct,starproduct2,starproduct3}. Instead, in this paper we would like to take onboard an alternative point of view, and consider the coherent state approach to non-commutativity introduced in a series of papers by E.\\ Spallucci and collaborators \\cite{Sma1,Sma2,Sma3,Sma4}. As we will briefly explain below, this model does not need the star-product, since all non-commutative effects are encoded in the Gaussian damping of the field modes \\footnote{On mathematical grounds, in this theory the star-product can be seen as replaced by the so-called Voros product \\cite{voros,voros2}.}. As a result, the field propagator is finite in the ultraviolet limit, but the dispersion relation is the same as the relativistic one. Compared to the wider class of modified theories mentioned above, this proposal has a stronger predictive power, as both field modes and propagators are known. For example, this model has been already studies in connection with the Casimir effect \\cite{Casadio}, and inspired several works on black holes \\cite{Nico,Nico2,Nico3,Nico4,Nico5,Nico6,Nico7,Nico8,Nico9,Nico10}, Unruh effect \\cite{Unruh}, inflation \\cite{MaxInf}, and quantum gravity \\cite{NicoSp}. The plan of the paper is the following: in the next section, we recall the main features of the coherent state approach to non-commutativity, and we construct a massive scalar field living on a two-dimensional Minkowski plane. In section III, we consider the generalization in curved space, and we show that Bogolubov transformations are not affected in the fourth section. In section IV, we look at the transplanckian problem of a black hole and compare with analogous calculations for the Unruh effect. Finally, we conclude the paper with few remarks. ", "conclusions": "\\noindent In this paper we constructed the field theory of a massive scalar field on the non-commutative plane, by mode analysis. The results coincide with the path integral approach, and clarified the relation between the Euclidean and Lorentzian propagator. We then used this field theory to tackle the transplanckian problem for a black hole. The main result is that the fuzziness on the manifold puts an upper limit on the energy density that can be stored near the horizon by a free-falling observer. The quantum backreaction on the geometry of the black hole can in principle be calculated with a suitable effective action, and it represents our next goal. In conclusion, we believe that the coherent state formulation of non-commutativity can offer new and intriguing perspectives on the phenomenology of quantum gravity. In this paper we presented just a glimpse of the potential of this theory, which certainly deserves further investigations." }, "1003/1003.5825_arXiv.txt": { "abstract": "We report the first detection of a linear correlation between rms variability amplitude and flux in the Ultraluminous X-ray source NGC 5408 X-1. The rms-flux relation has previously been observed in several Galactic black hole X-ray binaries (BHBs), several Active Galactic Nuclei (AGN) and at least one neutron star X-ray binary. This result supports the hypothesis that a linear rms-flux relation is common to all luminous black hole accretion and perhaps even a fundamental property of accretion flows about compact objects. We also show for the first time the cross-spectral properties of the variability of this ULX, comparing variations below and above $1$ keV. The coherence and time delays are poorly constrained but consistent with high coherence between the two bands, over most of the observable frequency range, and a significant time delay (with hard leading soft variations). The magnitude and frequency dependence of the lags are broadly consistent with those commonly observed in BHBs, but the direction of the lag is reversed. These results indicate that ULX variability studies, using long X-ray observations, hold great promise for constraining the processes driving ULXs behaviour, and the position of ULXs in the scheme of black hole accretion from BHBs to AGN. ", "introduction": "\\label{sect:intro} X-ray variability studies have played a crucial role in furthering our understanding of accreting black hole systems. Among the more recent discoveries in this area was the linear relationship between the rms amplitude of the variability and its flux, a relation which appears to persist over a wide range of timescales. To date this has been observed of four black hole X-ray binaries (BHBs) and one neutron star X-ray binary \\citep{Uttley01, Gleissner04, Uttley05, Gandhi09} and several Active Galaxies \\citep{Uttley01, Vaughan03b, Vaughan03a}. If present in accretion around stellar mass ($\\sim$10 M$_{\\odot}$) and supermassive black holes ($>$ 10$^{6}$ M$_{\\odot}$) it seems reasonable to hypothesise that the linear rms-flux relation is common to accreting black holes of all masses. In this case we would expect such a relation to be present in Ultraluminous X-ray sources (ULXs). ULXs were first discovered by the \\einstein~ \\citep{Fabbiano89} as point-like sources with luminosities greater than 10$^{39}$ erg~ s$^{-1}$. Timing and energy spectral studies suggest they are similar to accreting compact objects such as black hole binaries (BHBs) and active galactic nuclei (AGN). The inferred isotropic luminosities often exceed the Eddington limit for a stellar mass black hole ($ < 20~M_{\\odot}$). Consequently it has been suggested that this, along with their low inferred disc temperatures \\citep{Miller04}, indicates that they represent a new class of intermediate mass black holes with masses ranging from 100-1000 M$_{\\odot}$. Counter to this are theories that assume a lower mass black hole (e.g $<$ 100 $M_{\\odot}$) accreting at a super-Eddington rate \\citep{Begelman02}, emitting anisotropically \\citep{King01} or through relativistic beaming \\citep{Georgan02}. X-ray timing studies of bright ULXs have shown some to display timing properties common to BHBs such as spectral breaks and quasi-periodic oscillations (QPOs) \\citep{Strohmayer03, Strohmayer07, Strohmayer09, Heil09}. However, it has also been suggested that both energy spectral and timing behaviour within ULXs may differ from those seen from BHBs and AGN \\citep{Stobbart06, Goad06, Gladstone09, Heil09}. The low levels of variability seen in current long observations of ULXs with high count rates taken with \\xmm~ leaves only one promising source with variability which may be strong enough to detect an rms-flux relation: NGC 5408 X-1. This source has already been shown to display the kinds of timing features often seen in BHBs and AGN including quasi-periodic oscillations and power spectral breaks \\citep{Strohmayer07,Strohmayer09}. It has been observed in two separate long \\xmm~ ($>$ 100 ks) observations and is relatively bright ($\\sim$ 0.8$~ct~s^{-1}$). In this letter we analyse both observations for evidence of the rms-flux relation. \\begin{figure*} \\begin{center}$ \\begin{array}{ll} \\includegraphics[width=6.0 cm, angle=90]{NGC5408_obs1_rmsflux_label.ps} & \\includegraphics[width=6.0 cm, angle=90]{NGC5408_obs2_rmsflux_label.ps} \\end{array}$ \\end{center} \\caption{Rms-flux points for each of the three energy bands used in the analysis with 1$\\sigma$ errors. \\emph{Left}. obs. id. 0302900101; \\emph{Right} obs. id. 0500750101.} \\label{fig:energyrms} \\end{figure*} ", "conclusions": "We have shown that the linear rms-flux relation, previously seen in Galactic X-ray binaries and Active Galaxies is also present in the ULX NGC 5408 X-1. This detection is important as it extends the apparent ubiquity of the relation in accreting sources to ULXs. This can be taken as further proof that at least this ULX behaves in a manner similar to other accreting black hole systems. The fact that the rms-flux relation has been observed in the three main types of luminous accreting black hole systems (Galactic Binaries, ULXs and Active Galaxies) suggests that it is either a basic consequence of all luminous accretion onto black holes, or that it only occurs under certain restricted conditions but that these conditions are met for all three types of systems. The relation is clearly present above 1 keV, but at softer energies the lower variability amplitude means a linear relation can neither be confirmed nor rejected. However, the coherence of the two light curves indicates a common origin for the variability in the two bands, which would suggest that the soft band light curve may contain a linear rms-flux relation with a very low gradient. We may attempt to explain the observed rms-flux behaviour, in the most basic manner, in terms of two components behaving in different ways. The first component (A) obeys a linear rms-flux relation with zero intercept and a gradient equal to its fractional rms. The second component (B) is variable but its amplitude is constant with flux. This has the effect of adding constant flux and rms, effectively shifting the origin of the observed relation. The fluxes of components A and B are constrained to be non-negative and the average flux and rms of B can be no higher than the minimum observed flux and rms. The fractional rms of component B can be described by the gradient of the line from the origin to this point. The flux intercept of the linear model ($C$) can be understood in terms of the fractional rms of components A and B; if the fractional rms of B is greater than that of A then the flux intercept will be negative, as observed in the hard-band. The cross-spectral analysis is also indicative of a time lag which compares favourably to an extrapolation of the lag observed from BHBs to lower frequencies \\citep[see e.g.][]{Miyamoto88, Cui97, Nowak99, Pottschmidt00, Miyamoto93}. Although soft lags have been observed less frequently the majority of observations of BHBs do not reach the low energy range analysed here. \\cite{Miyamoto93} have observed soft lags in the very high state of GX 339-4 at low energies (1.2-2.3 keV), which appear similar to those we observe. These results, on both the rms-flux relation and time delays, could be confirmed with long X-ray observations of other bright, variable ULXs. These further observations should help to clarify the explanatory value of our proposed variability components (A and B), and how these and the hard-soft X-ray time lags relate to the different components thought to explain the energy and power spectra of accreting black holes. The present results indicate that such intensive studies of ULX variability, although challenging in terms of the observational demands, hold great promise for constraining the processes driving ULXs behaviour, and the position of ULXs in the scheme of black hole accretion from BHBs to AGN." }, "1003/1003.4823_arXiv.txt": { "abstract": "We report on an update of the test on the rotation of the plane of linear polarization for light traveling over cosmological distances, using a comparison between the measured direction of the UV polarization in eight RG at $z>2$ and the direction predicted by the model of scattering of anisotropic nuclear radiation, which explains the polarization. No rotation is detected within a few degrees for each galaxy and, if the rotation does not depend on direction, then the all--sky--average rotation is constrained to be $\\theta = -0.8^o\\pm2.2^o$. We discuss the relevance of this result for constraining cosmological birefringence, when this is caused by the interaction with a cosmological pseudo-scalar field or by the presence of a Cherns-Simons term. ", "introduction": "The possibility that the propagation of light through our universe might suffer from chiral effects, which could rotate the plane of polarization, arises in a variety of important contexts, such as the presence of a cosmological pseudo-scalar condensate, Lorentz invariance violation and charge parity and time (CPT) violation, neutrino number asymmetry, and the Einstein equivalence principle (EEP) violation (see \\citet{Nio07} for a review). The simplest form for modeling cosmological birefringence - a frequency independent rotation of the plane of linear polarization - is described by the interaction of a pseudo-scalar field $\\phi$ with photons through a term \\citep{Kol90,Raf96}: \\begin{equation} \\label{eq:1} \\mathcal{L}_{int}=-\\frac{g_\\phi}{4}\\phi F_{\\mu\\nu}\\tilde{F}^{\\mu\\nu}\\,, \\end{equation} where $g_\\phi$ is the coupling constant, $F^{\\mu\\nu}$ is the electromagnetic tensor and $\\tilde{F}^{\\mu\\nu}\\equiv\\frac{1}{2}\\epsilon^{\\mu\\nu\\rho\\sigma} F_{\\rho\\sigma}$ its dual. $\\phi$ could be a fundamental field or an effective description for cosmological birefringence due to Lorentz violation \\citep{Car89}. Indeed several efforts have been devoted to look for evidence of rotation of the plane of polarization: since we expect tiny effects on the basis of laboratory experiments, cosmological distances are required to have measurable effects and therefore the obvious approach has been to look for rotation in the most distant sources in the universe. What is required for this test is then a polarized distant source, for which the polarization orientation can be predicted: the predicted orientation is then compared with the measured one, looking for a rotation between the two. Radio galaxies (RG) are very good candidates, since these astrophysical objects are often polarized, both at radio and at UV-optical wavelengths, and are found at very high redshifts \\citep{Mil08}. Since the first successful detection of anisotropies in polarization of the cosmic microwave background (CMB) by DASI in 2002 \\citep{dasi}, also the CMB polarization pattern has become an important test for cosmological birefringence, which could probe the propagation of light back to the recombination surface, i.e. up to a redshift as high as $z \\sim 1100$. Cosmological birefringence was first constrained from RG observations, since these were the first cosmological sources providing information on polarization. \\citet{Car89} have used the fact that the distribution of the difference between the position angle (P.A.) of the radio axis and the P.A. of the E vector of linear radio polarization in distant RG ($0.42$, due to scattering of anisotropic nuclear radiation, excludes that the polarization plane rotates by more than a few degrees while the light travels from the source to us for more than 3/4 of the universe lifetime, confirming previous results at lower redshifts \\citep{Cim93, War97}. The all-sky-average constraint derived on the rotation of the polarization from the set of observations considered in this paper ($\\theta = -0.8^o\\pm 2.2^o$) is independent, but consistent with the constraints derived from CMB observations. We have studied the implications of this constraint on physical models of cosmological birefringence, showing how observations at high redshifts as those of RG are complementary to CMB anisotropies, as already occurs for SN Ia and CMB in measuring the expansion history. In the framework of theoretical models associating the cosmological birefringence with the variation of the Newton constant our results increase our confidence in the validity of the EEP, on which all metric theories of gravity are based. An improvement in both quantity and quality of the measurements of the UV linear polarization in RG at high redshift should be possible in the future with the coming generation of giant optical telescopes \\citep{Gil08,Nel08,Joh08}, and would narrow the constraint on $\\theta$ to a level smaller than what is now possible with RG and CMB." }, "1003/1003.3542_arXiv.txt": { "abstract": "{\\object{NRAO~150} is one of the brightest radio and mm AGN sources on the northern sky. It has been revealed as an interesting source where to study extreme relativistic jet phenomena. However, its cosmological distance has not been reported so far, because of its optical faintness produced by strong Galactic extinction.} {Aiming at measuring the redshift of \\object{NRAO~150}, and hence to start making possible quantitative studies from the source.} {We have conducted spectroscopic and photometric observations of the source in the near-IR, as well as in the optical.} {All such observations have been successful in detecting the source. The near--IR spectroscopic observations reveal strong H$\\alpha$ and H$\\beta$ emission lines from which the cosmological redshift of \\object{NRAO~150} ($z=1.517\\pm0.002$) has been determined for the first time. We classify the source as a flat--spectrum radio--loud quasar, for which we estimate a large super--massive black--hole mass $\\sim5\\times 10^{9} \\mathrm{M_\\odot}$. After extinction correction, the new near-IR and optical data have revealed a high-luminosity continuum-emission excess in the optical (peaking at $\\sim2000$\\,\\AA, rest frame) that we attribute to thermal emission from the accretion disk for which we estimate a high accretion rate, $\\sim30$\\,\\% of the Eddington limit.} {Comparison of these source properties, and its broad--band spectral--energy distribution, with those of \\emph{Fermi} blazars allow us to predict that \\object{NRAO~150} is among the most powerful blazars, and hence a high luminosity --although not detected yet-- $\\gamma$--ray emitter.} ", "introduction": "\\label{Int} \\object{NRAO~150}, first catalogued by \\citet[]{Pauliny66}, is nowadays one of the strongest radio and mm AGN sources in the northern sky \\citep[e.g.][]{Terasranta05,Agudo:subm}. The source has been monitored at cm and mm wavelengths for decades \\citep[e.g., ][ and references therein]{Aller85,Reuter97,Terasranta04}, and has displayed flux densities in the range $[2, 16]$\\,Jy at 2\\,cm\\footnote{{\\tt http://www.astro.lsa.umich.edu/obs/radiotel}} and $[1.5, 9.5]$\\,Jy at 3\\,mm\\footnote{H. Ungerechts, private communication}, with absolute maxima at the beginning of year 2009. At radio wavelengths, on VLBI scales, \\object{NRAO~150} displays a compact core plus a one-sided jet extending up to $r \\apgt 80$~mas with a jet structural position angle (PA) of $\\sim 30^\\circ$ \\citep[e.g., ][]{Fey00}. The first set of ultra-high-resolution mm-VLBI images of \\object{NRAO~150} \\citep{Agudo07} have allowed to report a large misalignment ($>100^{\\circ}$) between the cm-wave and the mm-wave jet, which is, together with the one sidedness of the jet, a clear sign of jet orientation close to the line of sight. More intriguing is the evidence of fast ``jet wobbling'' at $\\sim$11$^{\\circ}/$yr in the plane of the sky; the fastest reported for an AGN so far. The observations by \\citet{Agudo07} together with the cosmological redshift measurement presented in this paper and non-contemporaneous X-ray data have allowed to report the first quantitative estimates of the basic physical properties of the inner jet in \\object{NRAO~150}; i.e., Doppler factor $\\delta \\approx 6$, the bulk Lotentz factor $\\gamma \\approx 4$, the angle subtended between the jet and the line of sight $\\phi \\approx 8^{\\circ}$, and the magnetic field intensity of the flow $B\\approx 0.7$\\,G. Such large $B$ estimate seems to be compatible with the highly non-balistic superluminal motion of the inner jet in the source revealed by the mm-VLBI images ($\\beta_{\\rm{app}}\\approx3$ times the speed of light), which have also shown \\object{NRAO~150} as a prime target to study the origin of the jet wobbling phenomenon \\citep{AgudoVSOP}. \\object{NRAO~150} was first detected in December 1981 in the optical by \\citet{Landau83}. However, no optical classification or distance determination has been reported so far, perhaps due to the difficulties to observe the source in the visible range, where \\object{NRAO~150} is strongly affected by Galactic extinction (Galactic latitude $b\\approx-1.6$\\,$^{\\circ}$). This problem is partially overcome in the near-IR range, where spectroscopic observations from strong lined objects can be performed from Earth. Independently of the Galactic absorption along the line of sight of the source, the near-IR range is the adequate spectral range to detect the strong H$\\alpha$ line from AGN at cosmological redshifts between 1.2 and 3.6 \\citep[e.g., ][]{Babbedge04}. Here, we present the results of our spectroscopic near-IR observations, which were successful on detecting, for the first time, emission lines from \\object{NRAO~150}. In Sect.~\\ref{Obs} such observations and their data reduction procedures are outlined, together with a set of optical and near-IR photometric observations performed in the 2005--2007 time span. The cosmological redshift determination, the classification, the first estimate of the mass of the super-massive compact object in \\object{NRAO~150} and its accretion rate, as well as its broad band spectral energy distribution, are presented and discussed in Sect. \\ref{sc:results} and \\ref{sc:disc}, whereas a summary of our main results and our conclusions are provided in Sect.~\\ref{sc:conclu}. ", "conclusions": "" }, "1003/1003.5919_arXiv.txt": { "abstract": "The dominant component of the (100 MeV - 50 GeV) GRB emission detected by LAT starts with a delay relative to the prompt soft (sub-MeV) gamma-rays and lasts long after the soft component fades. This has lead to the intriguing suggestion that this high energy emission is generated via synchrotron emission of relativistic electrons accelerated by the external shock. Moreover, the limits on the MeV afterglow emission lead to the suggestion that, at least in bright GeV bursts the field is not amplified beyond compression in the shock. We show here that considerations of confinement (within the decelerating shock), efficiency and cooling of the emitting electrons constrain, within this model, the magnetic fields that arise in both the upstream (circum burst) and downstream (ejecta) regions, allowing us to obtain a direct handle on their values. The well known limit on the maximal synchrotron emission, when combined with the blast wave evolution, implies that late photons (arriving more than $\\sim$ 100 s after the burst) with energies higher than $\\sim $10GeV do not arise naturally from external shock synchrotron and almost certainly have a different origin. Finally, even a modest seed flux (a few mJy) at IR-optical would quench, via Inverse Compton cooling, the GeV emission unless the magnetic field is significantly amplified behind the shock. An observation of a burst with simultaneous IR-optical and GeV emission will rule out this model. ", "introduction": "The recent observations of the Large Area Telescope (LAT) on board of Fermi of GeV emission (100MeV - 50GeV) from GRBs revealed an interesting pattern. The GeV emission is delayed relative to the onset of the prompt MeV emission \\citep{Abdo080916C}. It shows a constant power-law decay long after the prompt emission dies out \\citep{Abdo090902B,GhiselliniEtal10}. While surprising at first, one may recall that a ''precursor\" of these observations was made already by EGRET that detected an $18$ GeV photon 90 minutes after the burst in GRB 940217 \\cite{Hurley94} and a rising late GeV spectral component in GRB 941017 \\citep{GonzalezEtal03}. This pattern suggests that the bulk of the GeV emission arises from an external shock afterglow \\citep{KB09a,KB09b,GhiselliniEtal10}. While a detectable high energy external shock emission was expected for a long time \\citep{MR94} and it was noted that external shock synchrotron emission may be the strongest afterglow GeV component (see. e.g. \\citealt{Fanetal08,FanPiran08,ZouEtal09}), the observation that this may be the dominant GeV emission over the whole burst, including the prompt phase, were surprising. Following these observations \\cite{KB09a,KB09b} proposed a revolutionary model in which they revise a critical component of the standard external shock scenario. They suggest that there is no magnetic field amplification beyond the usual shock compression. Namely, the downstream (shocked) magnetic field is just $4 \\Gamma$ (where $\\Gamma$ is bulk Lorentz factor behind the shock) times the upstream circum burst field. In doing so they are able to fit the overall afterglow spectrum (ranging from optical to GeV), as the low magnetic field in the emitting region quenches the lower energy emission. Additionally they get rid of a nagging theoretical problem - how are the fields amplified \\citep{Gruzinov01}? The magnetic field plays a triple role in the synchrotron-shock acceleration mechanism. It accelerates the electrons and confines them to the shock, while they are accelerated and it also controls the synchrotron emission. A weaker magnetic fields poses two challenges to the model: cooling and confinement. First, a comparison of the acceleration and the cooling times sets an absolute limit on the energy of synchrotron photon in the radiating fluid frame. Together with the hydrodynamics of the decelerating blast wave this puts a time dependent limit on the maximal energy of observed synchrotron photons. Photons above this limit are (almost certainly) not emitted by external shock synchrotron. Efficient cooling poses another limit on the model. A significant (though not dominant) amount of energy is emitted in the GeV emission. This implies that the emitting electrons must be fast cooling. Otherwise the system would be inefficient and the energy requirement unreasonable. As the cooling takes place mostly in the downstream region this last condition constrains the magnetic field there. While our original motivation was to examine the ''unamplified\" magnetic field scenario our analysis is more general and we allow for an amplification factor. We show that the observations of a significant GeV flux poses strong limits on the downstream magnetic field. These limits can be translated to limits on the upstream circum burst field (in the case of no amplification) or on the amplification factor. Confinement is most important in the upstream region, where the magnetic field is weakest. Thus, observations of GeV photons limit the upstream magnetic field with a weak dependence on field amplification in the shock. Finally we turn to the influence of Inverse Compton (IC) cooling on the observed GeV emission. Given the strong low energy (IR-optical) radiation fields (from the prompt, reverse shock and the forwards shock itself), the magnetic field density should be strong enough in order that IC cooling won't quench the GeV emission. This sets yet another, independent, limits on the downstream magnetic field. These considerations shed a direct light on the magnetic fields which are among the most elusive parameters of the external shock model. Note that here we derive constraints assuming that the external shock is adiabatic. If it is radiative (as suggested e.g., by \\citealt{GhiselliniEtal10}) then the constraints will be more stringent. We examine in this letter these limits that arise from the GeV emission. We don't attempt to provide a complete solution to the whole multiwavelength afterglow. As such our analysis is very general and it depends only on the assumptions of synchrotron process and the blast wave hydrodynamics. ", "conclusions": "Using the confinement and cooling conditions we have obtained limits on the values of the magnetic fields needed in the downstream and upstream regions in order to produce the observed GRB GeV emission via an external shock synchrotron. These constrains are based on minimal assumptions of synchrotron cooling and blast wave hydrodynamics. Both are essential ingredients of the external shock synchrotron model. The arguments we present allow us to explore directly the magnetic fields in both upstream and downstream regions, which are among the least constrained physical parameters of the model. We find that with no amplification the minimal fields required are on the high end ($\\tilde100 \\mu$ G), unless the external density is very low. The limits are even higher for a radiative solution. It is, of course, possible that this is a condition for GeV emission. However, the detection of GeV emission from all MeV bright GRBs that are within the LAT viewing angle suggests that the emission is generic. In this case at least a modest amplification is probably needed. Finally, we point out two critical predictions of the external shock synchrotron model: (i) No detection of late very energetic ($>10$ GeV) photons and (ii) No simultaneous detection of a bright ($>$mJy) IR-optical (depending on the specific case) signal with the GeV photons unless the upstream magnetic field is strongly amplified in the shock. Continued observations should be compared with these predictions and can provide future tests of this model. This research was supported by an ERC advanced research grant, by the ISF center for excellence for high energy astrophysics, ISF grant No. 174/08, and by an IRG Marie-Curie Grant. We thank Pawan Kumar and Rudolfo Barniol Duran and Boaz Katz for helpful discussions." }, "1003/1003.0639.txt": { "abstract": "We present a detailed investigation of \\sbs, a close binary star hosted by the planetary nebula PN\\,G135.9+55.9 \\citep[\\TS, ][]{paperI}. The nebula, located in the Galactic halo, is the most oxygen-poor one known to date and is the only one known to harbor a double degenerate core. We present XMM-{\\sl Newton} observations of this object, which allowed the detection of the previously invisible component of the binary core, whose existence was inferred so far only from radial velocity and photometric variations. The parameters of the binary system were deduced from a wealth of information via three independent routes using the spectral energy distribution (from the infrared to X-rays), the light and radial velocity curves, and a detailed model atmosphere fitting of the stellar absorption features of the optical/UV component. We find that the {\\sl cool} component must have a mass of $0.54\\pm0.2$\\,\\msun, an average effective temperature, T$_{\\mathrm {eff}}$, of $58\\,000\\pm3\\,000$\\,K, a mean radius of $0.43\\pm0.3$\\,\\rsun, a gravity $\\log g=5.0\\pm0.3$, and that it nearly fills its Roche lobe. Its surface elemental abundances are found to be: 12 + log He/H = 10.95 $\\pm$0.04\\,dex, 12 + log C/H = 7.20$\\pm$0.3\\,dex, 12 + log N/H $<$ 6.92 and 12 + log O/H $<$ 6.80, in overall agreement with the chemical composition of the planetary nebula. The {\\sl hot} component has T$_{\\mathrm {eff}}$ = 160--180\\,kK, a luminosity of about $\\sim 10^4$\\lsun\\ and a radius slightly larger than that of a white dwarf. It is probably bloated and heated as a result of intense accretion and nuclear burning on its surface in the past. The total mass of the binary system is very close to Chandrasekhar limit. This makes \\TS\\ one of the best type Ia supernova progenitor candidates. We propose two possible scenarios for the evolution of the system up to its present stage. ", "introduction": "\\objectname{\\sbs} was identified as a planetary nebula (PN) in \\cite{2001A&A...370..456T} and subsequently designated as \\objectname{\\pn}. More recently, we refer to this object as \\TS\\ \\citep{paperI}. The object has unusually few spectral lines for a PN and is renown for its extremely low oxygen content \\citep{2001A&A...370..456T, 2002AJ....124.3340J, 2005A&A...430..187P,2005IAUS..228..323S, paperI}. It is located above the Galactic plane at a distance of at least a dozen kpc, which places it among a handful of known halo PNe. Direct images obtained on the ground \\citep{2002A&A...395..929R,2002AJ....124.3340J}, and most recently by HST \\citep{2005AIPC..804..173N, paperI} confirm its PN identification. The observed expansion velocity of the nebula \\citep{2003A&A...410..911R} is typical of PNe. But another outstanding feature of this PN is that it harbors a close binary system \\citep{2004ApJ...616..485T}, revealed serendipitously by the displacement of stellar lines with respect to nebular lines within a single observing night. Since only one component of the binary could be observed in the optical and UV, it was suggested that the visible component has a temperature of 110\\,000--120\\,000\\,K. The lower limit is the minimum effective temperature needed to produce the observed [Ne\\,V] nebular emission line, while the upper limit was deduced from the slope of the continuum \\citep{2004ApJ...616..485T}. There was an ambiguity in the determination of the orbital period, although it was clear that the nucleus is a close binary with a period less than 4 hours. The high temperature, coupled with high log\\,$g$, determined from the profiles of absorption lines, led all studies prior to \\citet{paperI} to assume that the observed optical/UV component was the central star of the planetary nebula, i.e. the post-AGB star that lost its envelope and was the source of its ionization. \\citet{2005A&A...430..187P} suggested that, if the ionizing star were even hotter, the deduced oxygen abundance could be increased to a more common level for oxygen-poor PNe. However, a higher temperature would have required a higher reddening to match the observed continuum slope, and \\citet{2004ApJ...616..485T} had already used a higher extinction than would normally be estimated for the direction of \\sbs\\ in order to justify a temperature of 120\\,000\\,K . Next, we obtained photometric light curves of the binary core of \\sbs\\ \\citep{2005AIPC..804..173N}. The orbital period of the system turned out to be 3.92\\,hr and, to explain the double-humped shape of the light curve, we were led to invoke a Roche lobe-filling optical/UV component. It was observed that the depths of the minima in the light curve are uneven, an effect known to occur when the visible component is irradiated by a hotter (more energetic) source. The orbital dynamics required that this invisible component be another compact object of at least 0.85\\,\\msun\\ \\citep{2005AIPC..804..173N}. \\citet{2002AJ....124.3340J} pointed out the possibility that \\objectname{\\sbs} may be associated with the X-ray source \\objectname{1RXS\\,J115327.2+593959}. To detect the invisible source of irradiation and reveal the other component of the close binary, we observed it with the XMM-{\\sl Newton} X-ray observatory. We also conducted new optical spectroscopic observations of the object with the Gemini-North telescope to improve our knowledge of radial velocities of the optical/UV component of the binary and to better fit photospheric line profiles with atmospheric models. We also used the publicly available HST STIS observations of the object in the UV to bridge the optical and X-ray observations discussed here. The ionization state and chemical composition of the planetary nebula are analyzed in a companion paper \\citep{paperI}, while here we present a multifaceted analysis and modeling of the binary system. We analyze the history and the future of the stellar system in the light of evolutionary models for close binary stars. In Section 2, we present our new observations; in Section 3 we determine the physical parameters of the binary; in Section 4 we discuss the evolution of the object from the early stages, when it was a wide system comprised of main sequence stars, to the latest stage of a merging of two white dwarfs (WD) with possible type Ia supernova outcome; and in Section 5 we summarize our main results. ", "conclusions": "After a decade of intense study, we have achieved a good understanding of an object whose discovery spectrum was misidentified and incomprehensible in 1997. Since then, the object has been observed at practically all wavelengths with the help of the most advanced instruments. This paper accompanies \\citet{paperI}, which focuses upon the chemical composition and ionization state of \\TS's nebular shell. Here, we focus on the nature of the close binary nucleus of the PN. %The binary is one of the shortest period systems among the PNe with binary cores, with an orbital period of 3.924 hours. \\TS\\ is one of the shortest period systems among the double-degenerate or pre-double-degenerate systems, with an orbital period of 3.924 hours %Unlike any other binary central star of a PN, it is a double-degenerate system. This fact would not have caused confusion if the older of the %two white dwarfs components were significantly cooler than the core of the star that most recently ejected its envelope to form the current PN. However, observations and analysis clearly demonstrate that \\TS's nucleus is comprised of two compact stars, both extremely hot and thus, both being sources of ionization for the nebula. This unusual phenomenon created confusion and misinterpretation of the object in the past. Nevertheless, the correct understanding of the ionization source does not change the essence of those previous interpretations. \\TS\\ remains a PN with a record low oxygen abundance \\citep{paperI}. %According to our scenario, the object evolved from the wide binary through two common envelope episodes. In the current stage we are observing the second common envelope, which formed the PN. The core of the envelope-shedding post-AGB star is in the process of contraction and heating up. As for the lower mass component, at the present time, it nearly fills its Roche lobe and has an ellipsoidal shape. Recently, the more massive companion, which became a white dwarf earlier, underwent a period during which it accreted mass at a high rate and, since then, has maintained steady nuclear burning at its surface, resulting in a change in its structure and its heating up to the temperatures typical of supersoft X-ray sources. It is believed that a fraction of super soft X-ray sources are symbiotic stars, and \\TS\\ has just emerged from such a state, becoming one of the softest X-ray sources ever, similar to Lin 358 \\citep{2007ApJ...661.1105O}. According to our scenario, TS\\,01\\ evolved through two common envelope episodes. In the current stage we are observing the remainders of the second common envelope as a PN. The core of the envelope-shedding post-AGB star is in the process of contraction and heating up. At the present time, it nearly fills its Roche lobe and has an ellipsoidal shape. Before the last CE episode, the more massive component, which became a white dwarf earlier, underwent a period during which it accreted mass at a high rate and burned hydrogen steady. Since then it stays close to the temperatures range typical for supersoft X-ray sources. Its properties make \\TS\\ one of the softest X-ray sources ever, similar to Lin 358 \\citep{2007ApJ...661.1105O}. The parameters of the binary system were deduced using a wealth of information and via three independent routes. Although, each of these methods requires its own assumptions and each alone produces ambiguous results, in combination, they converge to values with unusual precision. Using the spectral energy distribution, from the far infrared to X-rays, the light and radial velocity curves, and by fitting atmospheric models to the stellar absorption features of the {\\sl cool} component, we find that the {\\sl cool} component has a mass of $0.54\\pm0.2$\\,\\msun, an average T$_{\\mathrm {eff}}$ of $58\\,000\\pm3\\,000$\\,K, a mean radius of $0.43\\pm0.3$\\,\\rsun, and $\\log g=5.0\\pm0.3$. The {\\sl cool} component nearly fills its Roche lobe. The temperature and gravity over the surface of the {\\sl cool} component are not homogeneous. The chemical composition of the {\\sl cool} component from atmosphere model fitting was determined as: 12 + log He/H = 10.95 and 12 + log C/H = 7.20, with an uncertainty of about 0.3\\,dex, and upper limits 12 + log N/H $<$ 6.92 and 12 + log O/H $<$ 6.80. Overall, the agreement with the abundances found in the nebula by \\citet{paperI} is very good, except for the carbon abundance, which is found to be higher in the nebula for a reason yet not understood. The parameters for the {\\sl hot} component are less certain. It is fairly clear that the spectral energy distributions of real stars at such high temperatures depart from that of a black body. The range of temperatures that we determined for the {\\sl hot} component spans 160--200\\,kK. It seems that the real object acts like a 180-200\\,kK blackbody in the X-ray range but appears as a 160\\,kK blackbody in the UV/optical range. Uncertainty in its temperature leads to uncertainty in its size, but it is obvious from our calculations that the {\\sl hot} component is larger than normal for a white dwarf, R$_{\\mathrm hot} > 0.1$\\,\\rsun, and is probably bloated as a result of intense accretion in the recent past. However, we have indirect information on the hot component through photoionization modeling by reproducing the intensities of the lines emitted by the nebula \\citep{paperI}. We estimate the distance to the object as $\\sim 21$ kpc, and our most reasonable luminosity estimate for the X-ray component is $\\sim 10^4$\\,\\lsun, appropriate for a supersoft X-ray source. The total mass of the binary is very close to Chandrasekhar limit. This makes \\TS\\ one of the best of the known candidates for the progenitor of a type Ia supernova." }, "1003/1003.3324.txt": { "abstract": "This review addresses the issue of whether there are physically realistic self-similar solutions in which a primordial black hole is attached to an exact or asymptotically Friedmann model for an equation of state of the form $p=(\\gamma-1)\\rho c^2$. %with $0 \\leq \\gamma \\leq 2$. In the positive pressure case ($1 < \\gamma < 2$), there is no such solution when the black hole is attached to an exact Friedmann background via a sonic point. However, it has been claimed that there is a one-parameter family of asymptotically Friedmann black hole solutions providing the ratio of the black hole size to the cosmological horizon size is in a narrow range above some critical value. There are also ``universal'' black holes in which the black hole has an apparent horizon but no event horizon. It turns out that both these types of solution are only asymptotically {\\it quasi}-Friedmann, because they contain a solid angle deficit at large distances, but they are not necessarily excluded observationally. Such solutions may also exist in the $2/3 \\le \\gamma < \\le 1$ case, although this has not been demonstrated explicitly. %could still be physically plausible if the fluctuations generating the black holes were set up in a preceding inflationary period. In the stiff case ($\\gamma = 2$), there is no self-similar solution in an exact background unless the matter turns into null dust before entering the event horizon, which is a contrived and probably unphysical situation. However, there may be asymptotically quasi-Friedmann solutions without a sonic point which contain universal black holes. In the negative pressure case ($0 < \\gamma < 2/3$), corresponding to a dark-energy-dominated universe, there is a one-parameter family of black hole solutions which are properly asymptotically Friedmann (in the sense that there is no angle deficit) and such solutions may arise naturally in the inflationary scenario. The ratio of the black hole size to the cosmological horizon size must now be {\\it below} some critical value, so the range is more extended than in the positive pressure case and one needs less fine-tuning. If one tries to make a black hole which is larger than this, one finds a self-similar solution which connects two asymptotic regions, one being exactly Friedmann and the other asymptotically quasi-Friedmann. This might be regarded as a cosmological wormhole solution providing one defines a wormhole throat quasi-locally in terms of a non-vanishing minimal area on a spacelike hypersurface. %, which is $0.70$ for $\\gamma=1/3$. %, as well as asymptotically Friedmann solutions which contain contain wormholes or naked singularities. This has the important physical implication that The possibility of self-similar black holes in phantom fluids ($\\gamma < 0$), where the black hole shrinks as the big rip singularity is approached, or tachyonic fluids ($\\gamma >2$) remains unclear. We also consider the possibility of self-similar black hole solutions in a universe dominated by a scalar field. If the field is massless, the situation resembles the stiff fluid case, so any black hole solution is again contrived, although there may still be universal black hole solutions. The situation is less clear if the scalar field is rolling down a potential and therefore massive, as in the quintessence scenario. Although no explicit asymptotically Friedmann black hole solutions of this kind are known, they are not excluded and comparison with the $0 < \\gamma < 2/3$ perfect fluid case suggests that they should exist if the black hole is not too large. This implies that a black hole might grow as fast as the cosmological horizon in a quintessence-dominated universe in some circumstances, supporting the proposal that accretion onto primordial black holes may have played a role in the production of the supermassive black holes in galactic nuclei. % This appears to be the case today and %[We calculate the accretion of a black hole either in the period immediately after inflation or at later stages when the dark energy dominated the matter density. SEPARATE PAPER?] ", "introduction": "%[{\\bf hideki: I have corrected the name ``Zel'dovich''}] Over the last 40 years there has been much interest in how fast a black hole formed in the early universe, when the density is usually radiation-dominated, would grow. As first pointed out by Zel'dovich and Novikov~\\cite{zn1967}, a simple Bondi-type accretion analysis suggests that a primordial black hole (PBH) would not grow much at all if it were much smaller than the cosmological horizon at formation but that it could grow at the same rate as the universe if its initial size were comparable to it. (The term ``cosmological horizon'' should here be interpreted as the Hubble horizon if the PBHs form after an inflationary period but the particle horizon otherwise.) One might expect the latter situation to apply, since a PBH must be bigger than the Jeans length at formation~\\cite{h1971}, so this suggests that any PBH might grow to the horizon mass at the end of the radiation era, which is around $10^{17}M_{\\odot}$. Since there is no evidence for such enormous black holes, for a while it was assumed that no PBHs ever formed. However, the validity of the Zel'dovich-Novikov calculation is questionable when the black hole size is comparable to the horizon size because it neglects the expansion of the universe and is not fully relativistic. Indeed, the conclusion that a PBH could grow at the same rate as the universe in the radiation-dominated era was disproved by Carr and Hawking~\\cite{ch1974}. They demonstrated this by proving that there is no self-similar solution which contains a black hole attached to an exact flat Friedmann background via a sonic point (i.e. in which the black hole forms by purely causal processes). The Zel'dovich-Novikov prediction is therefore definitely misleading in this case. Since the PBH must soon fall well below the horizon size, when their argument should be valid, this suggests that PBHs would not grow much at all. This gave the subject of PBHs a new lease of life and motivated Hawking to consider the quantum effects associated with black holes (since only PBHs could be small enough for these to be significant). Ultimately, this led to his discovery of black hole radiation~\\cite{h1975}, so it is ironic that a consideration of PBH accretion led to the conclusion that they evaporate! Carr and Hawking also claimed that there are self-similar solutions which are {\\it asymptotically} -- rather than {\\it exactly} -- Friedmann at large distances from the black hole. However, these correspond to special acausal initial conditions, in which matter is effectively thrown into the black hole at every distance; they do not contain a sonic point because they are supersonic everywhere. Indeed, such solutions exist in the ``dust'' case, when the cosmological fluid is pressureless~\\cite{ch1974}. They even found solutions in which the whole universe is in some sense inside a black hole; these are now termed ``universal'' black holes. Subsequently, the Carr-Hawking analysis was extended to perfect fluids with equation of state $p=(\\gamma -1)\\rho c^2$ where $p$ is the pressure, $\\rho$ is the mass density and %$1 < \\gamma <2$ so that the pressure is positive; $\\gamma$ is a constant ($4/3$ in the radiation case). This is the most general form for a barotropic equation of state compatible with self-similarity. It was proved that the non-existence of self-similar black holes in an exact Friedmann background applies for all values of $\\gamma$ in the range $1$ to $2$~\\cite{c1976,bh1978b}. Indeed, %it was shown that the only physical self-similar solution which can be attached to an exact external Friedmann solution via a sonic point is Friedmann itself; as the radial coordinate decreases, the other solutions either enter a negative mass regime or encounter another sonic point where the pressure gradient diverges~\\cite{bh1978b}. As in the radiation case, there are still acausal black hole solutions but these are again supersonic everywhere; the asymptotically Friedmann solutions which reach a sonic point %are still physical but they represent density perturbations which grow at the same rate as the universe rather than black holes~\\cite{cy1990}. Later it was realized that none of these positive-pressure self-similar solutions are strictly asymptotically Friedmann after all~\\cite{mkm2002}: there is a solid angle deficit at large distances which might in principle show up in the angular diameter test. It would therefore be more accurate to describe them as asymptotically ``quasi-Friedmann''. Such solutions are not excluded observationally, %and they could still be physical in the inflationary scenario. at least for some parameter range, but they are not physically well-motivated. The attempt to extend the analysis to stiff fluids ($\\gamma =2$) led to some controversy. Lin {\\it et al.}~\\cite{lcf1976} at first claimed that there {\\it is} a self-similar black hole solution in an exact Friedmann background in this case. However, Bicknell and Henriksen~\\cite{bh1978a} showed that this conclusion is invalid because Lin {\\it et al.} had misidentified the point corresponding to the black hole event horizon. Bicknell and Henriksen did manage to construct a numerical self-similar solution containing a black hole but it required the stiff fluid to turn into null dust at some point. Although this might seem rather contrived, Reed and Henriksen~\\cite{rh1980} later found a solution of this kind by generalizing some work of Hacyan~\\cite{hac1979}, involving a self-similar Vaidya model. However, even this model now seems implausible. The only possibility might be universal black hole solutions which are asymptotically quasi-Friedmann. It is also interesting to consider the growth of a black hole when the density of the universe is dominated by a scalar field, as expected in many cosmological contexts. If the scalar field is massless (i.e. if there is no scalar potential), then one might expect the same result to apply as in the stiff fluid analysis, since it is well known that a scalar field is equivalent to a stiff fluid provided the gradient is everywhere timelike~\\cite{m1988}. Indeed, the conclusion that there is no self-similar non-universal black hole solution in an exact or asymptotically Friedmann background dominated by a massless scalar field is supported by both numerical studies~\\cite{hc2005c} and analytical calculations~\\cite{hmc2006}. However, the situation is more complicated if there is a scalar potential (i.e. if the scalar field is massive, the mass being associated with second derivative of the potential) and the discovery that the universe is currently accelerating suggests that this may be the case at the present epoch~\\cite{supernova}. %The similarity assumption then requires that the potential have an exponential form. This has led to a study of black hole accretion in quintessence-dominated universes. Indeed, an argument similar to that advocated by Zel'dovich and Novikov has resurfaced in this context in order to explain the origin of the $10^6$ to $10^9 M_{\\odot}$ black holes thought to reside in galactic nuclei~\\cite{kr1995}. While there are several scenarios for the formation of such supermassive black holes, one possibility is that they originated in the early universe and grew self-similarly %to their present size through accretion of quintessence before cosmological nucleosynthesis~\\cite{bm2002,ch2005}, as well as by purely astrophysical processes later. Since this proposal is motivated by a Bondi-type argument which neglects the cosmological expansion, it is just as questionable as the original Zel'dovich-Novikov one. %This paper exploits the connection between a massless scalar field and a stiff fluid and considers the flaw in the original Lin et al. analysis in more detail. This has led to the search for self-similar black hole solutions in quintessence-dominated universes. In this case, the similarity assumption requires that the potential have an exponential form. Our analysis in Ref.~\\cite{hmc2006} then shows that there is no self-similar solution with a black hole in an exact or asymptotically Friedmann or asymptotically quasi-Friedmann background if the universe is decelerating and no such solution in an exact Friedmann background if it is accelerating. However, this does not prove non-existence in an asymtotically Friedmann or quasi-Friedmann background for the case in which the background is accelerating. % and asymptotically Friedmann or quasi-Friedmann. Kyo {\\it et al.}~\\cite{khm2008} have shown that there is a one-parameter family of self-similar asymptotically Friedmann solutions in this case, although it is unclear whether they can contain black holes. The acceleration of the universe can also be explained if its density is dominated by a perfect fluid with $0 < \\gamma < 2/3$, so this has motivated us to look for self-similar black hole solutions in this case~\\cite{hmc2007,mhc2007}. Such fluids are very different from positive-pressure ones, since there are no sound-waves (the sound-speed $\\sqrt{p/\\rho}$ being imaginary), so one might expect the conclusion about self-similar black hole solutions to be different. We describe such a fluid as ``dark energy'', although this term is sometimes used more generally and may indeed include quintessence itself. % to $2$ {\\color{red}\\bf when it is negligible.}) %{\\color{red}\\bf [THE DEFINITION OF QUINTESSENCE IS NOT RELATED TO %THE VARIABILITY OF GAMMA. LATER I DESCRIBE %THE STATUS ABOUT THE TERMINOLOGY.]} One might regard quintessence as a form of dark energy in which the parameter $\\gamma$, rather than being constant, may vary as the scalar field rolls down its potential. %but it is different from a $p=(\\gamma -1)\\mu$ fluid with fixed $\\gamma$. The value of $\\gamma$ varies However, there is an important physical difference between (constant $\\gamma$) dark energy and quintessence because there are sound-waves in the latter case, the sound-speed being the speed of light at short wavelengths~\\cite{hmc2006}. Since there are no sound-waves for a dark energy fluid, there can be no black hole solutions in an exact Friedmann background. However, as discussed in Refs.~\\cite{hmc2007,mhc2007}, there {\\it are} asymptotically Friedmann solutions containing black holes in this case. Indeed, unlike the positive-pressure case, these are genuinely asymptotically Friedmann rather than asymptotically quasi-Friedmann and the associated inhomogeneities may arise naturally in the inflationary scenario. %(This conclusion may also apply for $\\gamma > 2$ or for $\\gamma <0$, assuming these cases are physically plausible.) % and a positive scalar potential and an accelerating universe~\\cite{khm2008}. [OK? {\\bf hideki: This is correct, see my note}] There is a one-parameter family of such solutions but (in contrast to the implication of the Zel'dovich-Novikov argument) they only exist if the black hole is not too large compared to the particle horizon. These solutions are not analogous to the positive-pressure solutions which Carr and Hawking were originally seeking but they might nevertheless be physical. If one tries to find an asymptotically Friedmann self-similar solution with a black hole which is larger than the upper limit in the $0 < \\gamma < 2/3$ situation, one obtains a cosmological wormhole instead. The transition occurs as the black hole apparent horizon approaches the cosmological apparent horizon, after which both horizons disappear. This is in contrast to the $2/3 < \\gamma < 2$ case, where the two apparent horizons never merge~\\cite{c1976} and one tends to a separate closed universe as the black hole size increases. (However, the separate universe case is not itself self-similar.) %{\\color{red}\\bf[THE FORMER CASE IS FOR SELF-SIMILAR SOLUTIONS BUT THE LATTER CASE IS NOT FOR SELF-SIMILAR SOLUTIONS. SO THE DIRECT COMPARISON DOES NOT MAKE SENSE.]} In the wormhole solution the metric tends to an asymptotically Kantowski-Sachs form as one approaches the wormhole throat, this corresponding to a minimum physical radius, and the solution then connects to another asymptotically Friedmannn or asymptotically quasi-Friedmann universe. This paper provides a comprehensive discussion of all these solutions. For the most part, we will avoid mathematical technicalities, so the number of equations is minimized. Although the paper is intended as a general review of previous work in this area, we believe that bringing all the cases together is illuminating and leads to some original insights. %together it is also contains some original results useful to bring all the different cases together in order to get a proper overview of the subject. In successive sections, we consider positive-pressure fluids, stiff fluids, scalar and quintessence fields, dark-energy fluids and finally more exotic possibilities (phantom fluids with $\\gamma<0$, negative pressure fluids with $2/3<\\gamma<1$ and tachyonic fluids with $\\gamma>2$). %For $2/3<\\gamma<2$, the conclusion is that there are physically realistic self-similar solutions in which a black hole can grow as fast as the Universe but only in models which are asymptotically quasi-Friedmann. %This is not observationally excluded and might arise naturally in inflation. %{\\color{red}\\bf it may not be physically realistic}. {\\color{red}\\bf [CONTRADICTION!]} We conclude that there are certainly self-similar asymptotically Friedmann black hole solutions in the $0<\\gamma<2/3$ case and there may also be in the quintessence case, so it is interesting that these situations may be observationally favoured at the present epoch. In the final section we %estimate the amount of black hole accretion during the dark-energy-dominated phases of the universe and draw some general conclusions. Two appendices clarify the dimensions of various quantities used in our analysis and the connection between relevant energy conditions. It should be stressed that this paper is not intended to be a review of the more general problem of black holes in a cosmological background, although that problem is of great interest in its own right~\\cite{einstein,sultana,nayak,kastor,mou,gibbons}. It is also much more narrowly focussed than the earlier review of self-similar solutions in general relativity by Carr and Coley~\\cite{cc1999}. ", "conclusions": "In this review, we have discussed the possible existence of self-similar solutions containing a black hole or wormhole in an asymptotically Friedmann background whose density is dominated by a perfect fluid with $p=(\\gamma -1) \\rho c^2$ %and $0<\\gamma<2$ or a scalar field. %we have also considered the $\\gamma<0$ and $\\gamma>2$ cases. A simple Bondi-type analysis predicts an accretion rate of the form given by Eq.~(\\ref{eq:growth}) in all cases, though with a different value of the constant $K$, which appears to permit self-similar growth. Ultimately, this is because the Friedmann equation implies that the density scales as $\\rho \\propto t^{-2}$ in a flat Friedmann background, which is precisely the condition for self-similarity. However, the simple analysis is suspect because it neglects the cosmic expansion, so this has motivated a more careful relativistic analysis which allows for the expansion. In the positive-pressure case, the Bicknell-Henriksen solution (in which a stiff fluid turns into null dust) seems to be the only known self-similar black hole solution with a sound-wave which is {\\it exactly} Friedmann at large distances and even this is rather contrived. There is also the Hacyan radiation-dominated solution (in which the region containing the black hole is described by a Vaidya solution) but this is also artificial because all the photons have to become radially directed within some point and it also violates the 2nd law of thermodynamics. Apart from these examples, it appears that there is no self-similar solution containing a black hole in either an exact or asymptotically Friedmann background for any value of $\\gamma$ in the range $1 \\le \\gamma \\le 2$. There are only asymptotically quasi-Friedmann solutions, including ``universal'' black holes without a black hole event horizon or cosmological particle horizon. However, there are self-similar asymptotically Friedmann solutions for $0<\\gamma<2/3$, which suggests that PBHs {\\it can} grow as fast as the universe in the presence of dark energy (at least for a limited period). This conclusion may also apply for a quintessence field, although this has not been rigorously proved. The difference between the positive and negative pressure solutions is important in two respects. First, while the negative-pressure ones are physically well-motivated in the inflationary scenario, because one might expect the associated density perturbations to extend to ``infinity'', the positive-pressure ones are theoretically unmotivated and may also be observationally excluded for some parameter because they have a solid angle deficit at large distances. Second, self-similar black holes only exist if their size as a fraction of the cosmological horizon is not too small in the positive-pressure case but not too large in the negative-pressure case. This means that less fine-tuning is required in the latter case. It is interesting that there is no accretion in the limit $\\gamma\\to 0$, which is consistent with the Schwarzschild-de Sitter solution. The existence of these self-similar black hole solutions suggests that black holes may increase their mass by a considerable factor during any dark-energy or quintessence dominated era, regardless of whether or not they are ``primordial''~\\cite{bm2002} . There are two contexts in which this effect may be important: (1) in the period immediately after any PBHs formed at the end of inflation; (2) in the recent period when whatever dominates the density of the universe causes it to accelerate. Bean and Magueijo~\\cite{bm2002} focussed on the first situation but the second one may also be interesting. For although the accretion factor may not be very large up to now, the black hole mass will continue to grow like cosmic time so as long as the dark energy dominates the cosmological density. In the simplest models, this applies indefinitely, so the black hole can grow arbitrarily large. We are investigating the astrophysical implications of this result -- and especially its implications for the Bean-Magueijo claim that PBHs can grow large enough to provide the supermassive black holes in galactic nuclei -- in a separate paper. %~\\cite{chm2010}. \\if We have also considered the case of a phantom fluid with $\\gamma<0$ and found that there may also be self-similar black hole solutions in this context, although we have been unable to come to a definite conclusion. The astrophysical implications of this result is of special interest since the black hole mass goes to zero. We have not yet considered the $\\gamma>2$ case since it is not clear that this is physically well motivated. The other case which has been neglected is fluids with $2/3 < \\gamma < 1$. However, we do not believe these are fundamentally different from $1< \\gamma < 2$ as regards the accretion problem. \\fi Finally, we note that our analysis restricts the situations in which the similarity hypothesis~\\cite{cc2000c} applies. This hypothesis claims that there are certain circumstances in which spherically symmetric solutions evolve to self-similar form. The present work shows that there are at least some situations in which this does {\\it not} happen. %Since this analysis shows that there is not always a self-similar solution, in the positive-pressure situation. This is because the hypothesis cannot hold if (Even if there were a self-similar solution, one would still need to show that it was stable in order for it to be an attractor.) %he stability is important for the hypothesis because the existence does not mean that it %is an attractor. Tomohiro: Add a reference like this. %{\\bf [I THINK THE ANALYSIS HAS NOTHING TO DO WITH THE GENERAL VALIDITY OF SIMILAITY HYPOTHESIS. WE JUST FIND IT IS NOT ALWAYS TRUE AND ACTUALLY THE HYPOTHESIS ALREADY INCLUDES THAT PHRASE IN ITS STATEMENT.]} On the other hand, the similarity hypothesis {\\it does} appear to hold in spherical gravitational collapse when the pressure is positive but very small $(0 < \\gamma-1 \\ll 1)$~\\cite{hm2001,snajdr2006}. It may also hold in the negative-pressure situation." }, "1003/1003.5634_arXiv.txt": { "abstract": "We present measurements of carbon monoxide emission in the central region of the nearby starburst NGC~6000 taken with the Submillimeter Array. The $J=2-1$ transition of $^{12}$CO, $^{13}$CO, and C$^{18}$O were imaged at a resolution of $\\sim3''\\times2''$ ($450\\times300$\\,pc). We accurately determine the dynamical center of NGC\\,6000 at $\\alpha_{J2000.0}=15^h49^m49\\fs5$ and $\\delta_{J2000.0}=-29\\arcdeg23\\arcmin13''$ which agrees with the peak of molecular emission position. The observed CO dynamics could be explained in the context of the presence of a bar potential affecting the molecular material, likely responsible for the strong nuclear concentration where more than $85\\%$ of the gas is located. We detect a kinematically detached component of dense molecular gas at relatively high velocity which might be fueling the star formation. A total nuclear dynamical mass of $7\\times10^9\\,M_\\odot$ is derived and a total mass of gas of $4.6\\times10^8M_\\odot$, yielding a $M_{gas}/M_{dyn}\\sim6\\%$, similar to other previously studied barred galaxies with central starbursts. We determined the mass of molecular gas with the optically thin isotopologue C$^{18}$O and we estimate a CO-to-H2 conversion factor $X_{CO}=0.4\\times10^{19}\\,\\rm cm^{-2}(K\\,km\\,s^{-1})^{-1}$ in agreement with that determined in other starburst galaxies. ", "introduction": "The relation between the morphology and kinematics of the gas in the inner region of galaxies and the connection to its nuclear activity is not completely understood. Nuclear regions are often obscured by dust at optical wavelengths but can be viewed in continuum infrared dust emission as well as molecular line emission. CO can be used as a tracer of molecular gas, and its rotational transitions are observed at millimeter and sub-millimeter wavelengths. Observations of the morphology and kinematics of the molecular gas content in the central regions of spirals can provide insight into the dynamical processes occurring there \\citep[e.g.,][]{jogee05, perez00}. Studies indicate that instabilities in bars and dissipation of gas clouds remove angular momentum from material orbiting the galactic center, driving gas and dust inward to fuel star formation \\citep[e.g.,][]{pfenniger90,Sakamoto99b,knapen02,jogee05,Sheth05}. NGC\\,6000 is a nearby \\citep[D$\\sim$31.6\\,Mpc,][]{pizzella05} barred spiral starburst galaxy. Despite its proximity and brightness, this galaxy has not been studied in as great detail as others of similar distance and luminosity, due to its southern declination. Hubble Space Telescope observations reveal several bright sources in circumnuclear star-forming rings and a large-scale ($\\gtrsim$~1~kpc) bar \\citep{carollo97, carollo02, fathi03}. However, no such barred structures are identified towards the more irregular nuclear region \\citep{carollo99, hunt04}. Infrared observations towards the nucleus of NGC~6000 have found significant polycyclic aromatic hydrocarbon (PAH) emission \\citep{siebenmorgen04} and determined a dust temperature of $29.2\\pm2.6$~K \\citep{yang07}. NGC\\,6000 appears bright in CO emission \\citep{young95, mauersberger99} as well as \\ion{H}{1} \\citep{koribalski04} as detected with single-dish telescopes. However no detailed morphological and kinematical study of the gas content has been performed yet. A compilation of the observed and derived properties for NGC\\,6000 are presented in Table~\\ref{galaxyparams}. In this paper we present observations of the carbon monoxide emission in the $J=2-1$ transition of the three brighter isotopologues ($\\rm^{12}CO$,$\\rm^{13}CO$,$\\rm C^{18}O$) towards \\object{NGC\\,6000} using the Submillimeter Array \\citep[SMA;][]{ho04}. This is the first high resolution morphological study of the molecular component towards NGC\\,6000. ", "conclusions": "\\label{Discussion} \\subsection{Bar-driven molecular gas fueling the starburst in NGC~6000} The study of the molecular gas content in a sample of 20 galaxies by \\citet{Sakamoto99b} statistically shows that the gas tends to be more concentrated in the central kiloparsec in barred systems. This is the case of NGC\\,6000 where most of the gas is located in the inner half kiloparsec (Sect.~\\ref{sec.COmass}). The large-scale bar revealed in the NIR may also have a fingerprint in the molecular gas kinematics described in this work. We observe that the P.A.=0$^\\circ$ P-V diagram in Fig.~\\ref{fig:pv} (right) shows how the extended molecular component displays a S-shaped feature. This kinematic signature can be understood as the noncircular motions in the context of a barred potential \\citep{binney92}. The inner material could be moving in the barely resolved $x_2$ orbits which, with the resolution of our maps, would look like a molecular disk. On the other hand the external gas would be moving in large elliptical $x_1$ orbits traced as the S-shaped profile in the P-V diagram. Higher resolution imaging of NGC\\,6000 might support the scenario of a barred potential by resolving the inner disk structure into circular $x_2$ orbits as observed by \\citet{meier08} towards Maffei\\,2. Similar signatures of a barred potential are observed in the P-V diagrams of other galaxies such as NGC\\,1530 and NGC\\,4258 \\citep{downes96,cox96} and equivalent structure are found in their innermost regions. However our measured $M_{gas}/M_{dyn}$ ratio of $6\\%$ is lower than the average measured by \\citet{Sakamoto99b} in starbursts and barred galaxies. This measurement is significantly affected by the way the gas mas is calculated Indeed, adopting their same conversion factor ($X=3\\times10^{19}\\,\\rm cm^{-2}(K\\,km\\,s^{-1})^{-1}$ rather than $X=0.4\\times10^{19}\\,\\rm cm^{-2}(K\\,km\\,s^{-1})^{-1}$, see Sect.~\\ref{COtoH2}) the measured ratio would increase up to $\\sim35\\%$. This ratio is only found in barred starbursts in the \\citet{Sakamoto99b} sample. \\subsection{The high velocity asymmetry} The strong asymmetry in the nuclear ring/disk is also a particularly interesting feature. As seen in Fig.~\\ref{fig:pv} a significant part of the molecular gas traced by CO is skewed towards the region at the lowest velocities, corresponding to the south-east disk component. However, the emission at the highest velocities from 2300 to 2400\\,km\\,s$^{-1}$ in the P-V diagrams, although following the same rotation gradient, seems to be significantly detached from the inner rotating structure. This feature at the highest velocities corresponds to the wing observed in the spectrum shown in Fig.~\\ref{fig:specs}. We discard the possibility of self-absoption at high velocities, which would change the systemic velocity derived from the Gaussian fit to the line, as the optically thin $\\rm ^{13}CO$ and $\\rm C^{18}O$ show similar profiles and are centered at the same velocity. It is surprising that $\\rm ^{13}CO$ and even $\\rm C^{18}O$ are detected in this velocity component between 2300 to 2400\\,km\\,s$^{-1}$. Moreover, the low ratio of $\\rm ^{13}CO$ and $\\rm C^{18}O$ with respect to the main isotopologue implies that a significant opacity is affecting even the $\\rm ^{13}CO$ emission in this molecular component. CO being optically thick can provide some constraint on the extent of the emission as compared to our resolution. From the CO brightness temperature measured $T_{\\rm b}\\sim1$\\,K and assuming an excitation temperature $T_{\\rm ex}=30$\\,K \\citep[as derived from the dust temperature,][and assuming CO being thermalized at this temperature]{yang07}, we estimate the region extent to be $<80$\\,pc ($<0.5''$). This is slightly above the $\\sim 50$\\,pc typical size of giant molecular clouds within our Galaxy. This giant molecular cloud within the inner region of NGC\\,6000 must be very dense and more compact than our estimate to explain such CO opacity. It could have originated as a shock between the circumnuclear disk and the infalling molecular gas along the dust lanes. Moreover, emission arises from the same region as the star forming ring which suggest that this dense and compact molecular cloud might be directly related to the fueling of star formation event in this region. \\subsection{CO-to-H$_2$ conversion factor in starburst galaxies} \\label{COtoH2} We can estimate the column density of molecular gas via the conversion factor $X=N(H_2)/CO=3.0\\times10^{20}\\,\\rm cm^{-2}(K\\,km\\,s^{-1})^{-1}$ based on measurements of galactic disk molecular clouds \\citep{solomon87}. This way, and adopting the CO integrated intensity in this work, we derive a H$_2$ column density of $7.5\\times10^{23}\\rm cm^{-2}$ which is almost an order of magnitude larger than the value derived from the optically thin C$^{18}$O line (Sect.~\\ref{sec.COmass}). As already observed towards the starburst NGC~253 \\citep{mauers96,harrison99}, the conversion factor in the starburst environment might be significantly lower than that in the Galactic disk. From our measurements we calculate a conversion factor $X=0.4\\times10^{19}\\,\\rm cm^{-2}(K\\,km\\,s^{-1})^{-1}$ in agreement with that derived for NGC\\,253 \\citep{mauers96}. This result support the evidences found for a lower conversion factor in the central region of galaxies, and in particular in starburst with respect to the Galactic disk \\citep{downes98} or that measured in star forming regions in the Large Magellanic Clouds \\citep{wang09}." }, "1003/1003.4318_arXiv.txt": { "abstract": "We present CARMA observations of the thermal dust emission from the circumstellar disks around the young stars RY~Tau and DG~Tau at wavelengths of 1.3~mm and 2.8~mm. The angular resolution of the maps is as high as 0.15\\arcsec, or 20 AU at the distance of the Taurus cloud, which is a factor of 2 higher than has been achieved to date at these wavelengths. The unprecedented detail of the resulting disk images enables us to address three important questions related to the formation of planets. (1) What is the radial distribution of the circumstellar dust? (2) Does the dust emission show any indication of gaps that might signify the presence of (proto-)planets? (3) Do the dust properties depend on the orbital radius? We find that modeling the disk surface density in terms of either a classical power law or the similarity solution for viscous disk evolution, reproduces the observations well. Both models constrain the surface density between 15 and 50 AU to within 30\\% for a given dust opacity. Outside this range, the densities inferred from the two models differ by almost an order of magnitude. The 1.3 mm image from RY~Tau shows two peaks separated by 0.2\\arcsec\\ with a decline in the dust emission toward the stellar position, which is significant at about 2-4$\\sigma$. For both RY~Tau and DG~Tau, the dust emission at radii larger than 15 AU displays no significant deviation from an unperturbed viscous disk model. In particular, no radial gaps in the dust distribution are detected. Under reasonable assumptions, we exclude the presence of planets more massive than 5 Jupiter masses orbiting either star at distances between about 10 and 60 AU, unless such a planet is so young that there has been insufficient time to open a gap in the disk surface density. The radial variation of the dust opacity slope, $\\beta$, was investigated by comparing the 1.3 mm and 2.8 mm observations. We find mean values of $\\beta$ of 0.5 and 0.7 for DG~Tau and RY~Tau respectively. Variations in $\\beta$ are smaller than $\\Delta\\beta=0.7$ between 20 and 70 AU. These results confirm that the circumstellar dust throughout these disks differs significantly from dust in the interstellar medium. ", "introduction": "Resolved images of circumstellar disks around young stars provide the most direct tool for investigating the formation of planets. At millimeter wavelengths, the thermal dust emission is generally optically thin and measures the radial distribution of circumstellar dust \\citep{bs90}. However, since circumstellar disks in nearby star forming regions typically have radii between 100 and 500 AU, sub-arcsecond angular resolution is required to spatially resolve the dust emission, even in nearby star-forming clouds. Millimeter-wave interferometers are essential for such studies. Since sub-arcsecond observations at millimeter wavelengths require both high sensitivity and high dynamical range, only a small number of bright disks have been observed at resolutions of 0.4\\arcsec-1\\arcsec\\ to date \\citep{b08,gu99,is07,pi05,pi06, pi07,sim00,t03,wi00}. The Combined Array for Research in Millimeter-wave Astronomy (CARMA) and the new extended configuration of the Sub-Millimeter Array are rapidly enabling more extensive high resolution surveys of circumstellar disks, particularly in the Taurus and Ophiuchus star forming regions \\citep[][hereafter Paper I]{an09,hu09,is09}. The highest angular resolution achieved so far by millimeter-wave interferometers is 0.3\\arcsec-0.4\\arcsec, corresponding to spatial scales of 40-50~AU at the distance of Taurus and Ophiuchus. In most cases, the dust density appears to increase smoothly inward down to the orbital radius resolved by the observations, typically $\\sim$25~AU. However, central cavities in the dust distribution are revealed in a number of disks \\citep{an09,hu09}. It remains a matter of debate whether these cavities are caused by dynamical interactions, inside out disk dispersal mechanisms, dust opacity variations, or viscous evolution \\citep[e.g.,][Paper I]{al06,c05,cm07,dd05}. Nevertheless, these observations still lack the angular resolution required to resolve the innermost part of the disk where the density of the circumstellar material is highest and the formation of planets is more probable. Here we describe CARMA observations of the thermal dust emission towards the young stars DG~Tau and RY~Tau at an angular resolution of 0.15\\arcsec\\ at 1.3~mm and 0.3\\arcsec\\ at 2.8~mm. At the distance of Taurus (140~pc), 0.15\\arcsec\\ corresponds to spatial scales of 20~AU, such that emission on orbital scales comparable to Saturn can be resolved. This is more than a factor of two improvement over previous observations of circumstellar disks at these wavelengths. DG~Tau and RY~Tau are classical T Tauri stars of spectral type M0 and K1 respectively \\citep{mu98,kh95}. Stellar ages inferred from stellar evolutionary models are less than 1~Myr (see Paper I for more details and references). The relative youth of both systems is confirmed by the presence of large amounts of gas and dust extending to 0.1~pc and by associated stellar jets and outflows \\citep[see, e.g,][]{mr04,sb08}. From near-infrared to millimeter wavelengths, both objects exhibit strong emission in excess of that from the stellar photospheres. This is attributed to rotating disks with radii of few hundred AU that first absorb and then re-emit radiation from the central stars. \\citep{ks95,t02}. Our earlier CARMA observations of 1.3~mm thermal dust emission from these disks, at a resolution of 0.7\\arcsec, suggested disk masses between 5 and 150\\% of the stellar mass for both sources (see Paper I). These high disk masses and the youth of RY~Tau and DG~Tau make these prime targets to investigate the earliest stages of planet formation. Our new observations of RY~Tau and DG~Tau have a factor of 5 better angular resolution and a factor of 3 better sensitivity than the previous data. This paper investigates three main questions related to the formation of planets in young circumstellar disks. (1) What is the surface density distribution in the observed disks down to an orbital radius of 10 AU? (2) Are there any signatures of planet formation contained in the dust distribution? Finally, (3), do the dust properties vary with orbital radius? A qualitative answer to the first two questions is proposed in Section~\\ref{sec:morp} where we present the observations and discuss the morphology of the dust emission. A quantitative analysis is described in Section~\\ref{sec:mod}, where we compare the observations with theoretical models of disk emission. Implications of these results for disk structure, for the possible presence of planets, and for the radial variation of the dust opacity are considered in Section~\\ref{sec:res}. The conclusions are presented in Section~\\ref{sec:conc}. ", "conclusions": " \\begin{itemize} \\item[(1)] Both the classical power law disk surface density (Hayashi 1981) and the similarity solution for the viscous evolution of a Keplerian disk (Hartmann 1998) fit the observations well. The surface density is well constrained between 15 and 50 AU. In this region, the two models lead to values of $\\Sigma$ that agree within 30\\% for a fixed dust opacity. At smaller and larger radii, the surface density depends on the assumed model and varies by almost an order of magnitude. We have verified that the assumptions on the dust opacity have a small effect on the model fitting and, therefore, on the radial profile of the dust density. However, the total disk mass may vary by almost two order of magnitude for different dust compositions and grain size distributions. \\item[(2)] The dust emission in DG~Tau is mostly radially symmetric. It is characterized by a single, central peak and smoothly decreases up to an angular distance of about 0.5\\arcsec. Theoretical disk models reproduce the observation very well, with randomly distributed residuals between 3 and 6$\\sigma$. No systematic deviation from the similarity solution for the surface density of a viscous disk are observed. By simulating the presence of planets in the disk via the gap in the surface density produced by tidal torques, we find that the observations exclude the presence of planets more massive than Jupiter orbiting between 5 and 40 AU from the central star, unless the planets are very young ($<10^4$ yr) and have not had the time to open a gap in the disk. The observations lack both the angular resolution and sensitivity to investigate the presence of planets less massive than about 0.5~Jupiter masses. For RY~Tau, the dust emission is characterized by two peaks separated by about 28~AU that suggest a decrease in the surface density, or dust opacity, within 14 AU of the central star. We found that the similarity solution for the disk surface density is characterized by a negative value of $\\gamma$, and provides a reasonable explanation for the double peak intensity observed at 1.3 mm. Depletion of millimeter dust grains \\citep{dd05}, decreasing values of the disk viscosity, or the presence of planetesimals, may produce the observed dust morphology. At larger radii, the dust emission shows a very smooth profile with no asymmetries or gaps. The lack of gaps in the disk suggests that any planets between 10 and 50 AU are less massive than about 5 Jupiter masses, or, as for DG Tau, are very young. \\item[(3)] The best-fit models to the 1.3~mm and 2.8~mm data were compared to investigate the radial dependence of the slope opacity $\\beta$, assuming that the dust opacity at millimeter wavelengths is expressed by a power law $k_\\lambda \\propto \\lambda^{-\\beta}$. We can exclude cases in which $\\beta$ varies by more than 0.7 within 70 AU. Nevertheless, between 20 and 70 AU, the disks around DG~Tau and RY~Tau are characterized by values of $\\beta$ that are smaller than that found in the ISM. This implies that the dust has been reprocessed and has grown in size up to a radius of at least 20 $\\mu$m. The investigation of the radial variation of $\\beta$ is still limited by the angular resolution and by the small separation in wavelength between the observations. In the future, ALMA and the EVLA will play crucial roles in the investigation of the radial dependence of the dust properties by increasing the angular resolution and the interval in wavelength. \\end{itemize}" }, "1003/1003.3012_arXiv.txt": { "abstract": "{The covariant entropy bound states that the entropy, $S$, of matter on a light-sheet cannot exceed a quarter of its initial area, $A$, in Planck units. The gravitational entropy of black holes saturates this inequality. The entropy of matter systems, however, falls short of saturating the bound in known examples. This puzzling gap has led to speculation that a much stronger bound, $S\\lesssim A^{3/4}$, may hold true. In this note, we exhibit light-sheets whose entropy exceeds $A^{3/4}$ by arbitrarily large factors. In open FRW universes, such light-sheets contain the entropy visible in the sky; in the limit of early curvature domination, the covariant bound can be saturated but not violated. As a corollary, we find that the maximum observable matter and radiation entropy in universes with positive (negative) cosmological constant is of order $\\Lambda^{-1}$ ($\\Lambda^{-2}$), and not $|\\Lambda|^{-3/4}$ as had hitherto been believed. Our results strengthen the evidence for the covariant entropy bound, while showing that the stronger bound $S\\lesssim A^{3/4}$ is not universally valid. We conjecture that the stronger bound does hold for static, weakly gravitating systems.} \\begin{document} ", "introduction": "\\paragraph{Covariant entropy bound} The covariant entropy bound~\\cite{CEB1} (see Ref.~\\cite{RMP} for a review) establishes a general relation between quantum information and classical geometry: The entropy of matter on a light-sheet $L$ (a non-expanding null hypersurface) orthogonal to a spatial surface $B$ cannot exceed the surface area $A$, measured in Planck units: \\begin{equation} S[L(B)]\\leq \\frac{A(B)}{4}~. \\label{eq-ceb} \\end{equation} This holds for arbitrary spacelike surfaces $B$ of codimension two, open or closed, in any spacetime satisfying Einstein's equation with physically reasonable matter. It implies that the entropy in the past of any event cannot exceed half of the maximum area of the past light-cone (which is finite in cosmological spacetimes, for example). The covariant bound appears to be rather rigid. Several other entropy bounds of the form $S\\leq A/4$ have been formulated, which do not involve light-sheets or do not impose the condition $\\theta\\leq 0$ on the expansion of the geodesics generating the light-sheet. A simple example is the ``spacelike entropy bound'', the claim that the entropy in any volume of space is less than the area of its boundary. One finds that each of these bounds can be violated by arbitrarily large factors~\\cite{FisSus98,KalLin99,RMP} in perfectly physical spacetimes. Meanwhile, no counter-examples to the covariant entropy bound have been found~\\cite{RMP}, and sufficient conditions have been identified that guarantee the bound's validity in a large class of spacetimes~\\cite{FMW,BouFla03}. From the viewpoint of local field theory, the covariant entropy bound is surprising, since one expects entropy to scale with volume. The holographic principle, in its most general form, is the conjecture that this surprising relation between geometry and information must have a fundamental origin, in a theory of quantum gravity~\\cite{CEB2}. This expectation has been borne out in a special case, by the AdS/CFT correspondence~\\cite{Mal97}. The number of binary CFT degrees of freedom necessary to describe an AdS region of (sufficiently large) surface area $A$ is indeed of order $A$~\\cite{SusWit98}. But the covariant entropy bound applies much more broadly. It holds in cosmological and other highly dynamical spacetimes that lack a fundamental quantum gravitational description, such as the interior of black holes. In this general setting, one expects that the bound will constrain how a quantum gravity theory should be formulated and how spacetime should arise from it. Light-sheets on which the entropy is {\\em equal\\/} to the area seem poised to play a distinguished role in the emergence of a classical geometry. Therefore, matter systems that saturate the covariant entropy bound are of special interest. Remarkably, no explicit examples of such systems have been found to date. Of course, the Bekenstein-Hawking entropy of an event horizon is equal to the horizon area, $S_{\\rm BH}=A/4$, and so, in a sense,\\footnote{The black hole horizon {\\em is\\/} a light-sheet, which is crossed by the matter that created or fell into the black hole. In order to say that a black holes saturates the covariant bound, one would like to view the black hole as a kind of matter object whose worldline crosses a different light-sheet of initial area $A_{\\rm hor}$. This is possible if one regards the stretched horizon as a timelike boundary of the black hole~\\cite{SusTho93} and one terminates the light-sheet there.} it saturates the bound. However, it is striking that no material objects, with ordinary, non-Bekenstein-Hawking entropy, are known to saturate the bound. \\paragraph{A stronger bound?} In fact, it would appear that systems made from ordinary matter obey the far stronger bound \\begin{equation} S\\lesssim A^{3/4}~, \\label{eq-a34} \\end{equation} falling short of the holographic bound by a factor of $A^{1/4}$ in Planck units, an enormous factor for macroscopic systems. Consider, for example, a spherical box of radius $R$ filled with radiation at temperature $T$. By increasing $T$, one can increase the entropy, $S\\sim R^3T^3$. However, $T$ is bounded from above by the requirement that the box should not collapse into a black hole: $E\\sim R^3T^4\\lesssim R$. The largest radiation entropy is attained at the threshold of collapse, when $S\\sim R^{3/2}$, saturating the bound (\\ref{eq-a34}) but, in the semiclassical regime $A\\gg 1$, falling far short of the holographic bound (\\ref{eq-ceb}). Another example that supports the stronger bound (\\ref{eq-a34}) obtains in cosmology. Consider the past light-cone of an observer at the time $t_E$ in a radiation-dominated flat FRW universe with vanishing cosmological constant. If one follows the light-cone towards the past, its cross-sectional area initially expands, but then contracts until it vanishes at the Big Bang. The sphere of maximum area, $A_{\\rm AH}\\sim t_E^2$, can be regarded as the origin of {\\em two\\/} light-sheets, one past-directed and one future-directed, which together form the past light-cone. The covariant bound states that the entropy on each light-sheet must be less than $A_{\\rm AH}/4$. The actual entropy on each light-sheet can be estimated by noting that the past light-cone has comoving size comparable to the volume enclosed by the Hubble horizon at the time $t_E$. Since the evolution is adiabatic, the entropy on the past light-cone is the same as the entropy within the Hubble radius at the time $t_E$. The proper energy density of radiation at this time is $\\rho_{\\rm rad}\\sim t_E^{-2}$, and the proper entropy density is $s\\sim \\rho_{\\rm rad}^{3/4}\\sim t_E^{-3/2}$. The proper horizon volume is $\\sim t_E^3$, so the total entropy on the light-sheet is $S\\sim t^{3/2}\\sim A_{\\rm AH}^{3/4}$. Again one finds the bound (\\ref{eq-a34}) approximately saturated, but the holographic bound far from saturated, with a factor $A^{1/4}$ to spare. There is no contradiction here: the holographic bound is just an inequality, and it is not surprising that many systems fall far short of saturating it. In most cases the ratio $S/A$ is even smaller than in the above examples: consider, for example, a sphere surrounding a region with vanishing entropy, such as empty space or a crystal at zero temperature. What is intriguing, however, is that it appears to be hard to exceed the ratio $S/A\\sim A^{-1/4}$ attained by the above two examples. Can we find ordinary matter systems whose entropy comes close to saturating the covariant entropy bound? Or is the holographic bound needlessly lenient? If there truly was a universal bound of the form (\\ref{eq-a34}) for matter, we would be forced to reconsider the significance of the covariant entropy bound and the holographic principle. What importance could we ascribe to an upper bound that is far from saturated in all known examples? Perhaps it is the quantity $A^{3/4}$ that truly governs the information content in quantum gravity? In this paper, we rule out this possibility, and we reaffirm the fundamental stature of the covariant entropy bound. We will find simple examples in which the covariant entropy bound is indeed saturated, and an even larger class of examples where the entropy on a light-sheet exceeds the stronger bound $A^{3/4}$ by arbitrarily large factors. \\paragraph{Relation to previous work} A number of works have considered systems that may violate the naive bound $S\\lesssim A^{3/4}$ or even saturate the holographic bound $S\\leq A/4$. Prior to the covariant bound, Fischler and Susskind~\\cite{FisSus98} proposed a cosmological holographic bound $S\\leq A/4$ on the future light-cone of a point on the Big Bang, where $A$ is the area of the light-cone at some time $t$, and $S$ is the entropy on the portion of the light-cone below $t$. In flat or open universes, this light-cone is an allowed light-sheet off of any of its cross-sectional surfaces, so the following example~\\cite{FisSus98} applies also to the covariant entropy bound. In a universe filled with a maximally stiff fluid ($p=\\rho$), one finds $S/A\\sim \\sigma(t)$, where $\\sigma$ is the comoving entropy density. The area of the light-cone grows with time, so $A^{1/4}\\sigma(t)$ will eventually exceed unity, and the naive bound $S\\lesssim A^{3/4}$ will be violated, assuming that such a fluid can carry entropy and do so adiabatically ($\\sigma=$const). Moreover, assuming that one can achieve $\\sigma\\sim 1$, the fluid will saturate the covariant bound for arbitrarily large area $A$. Neither of these assumptions, however, has yet been justified in a microscopic model of such matter. In our examples below, the entropy is that of ordinary pressureless particles or radiation, and thus can be explicitly computed rather than posited. In Ref.~\\cite{CEB1}, a shell collapsing onto a black hole was considered and shown to {\\em obey\\/} the covariant bound. Under a number of assumptions that erred on the side of larger entropy, this shell was seen to saturate the bound on a light-sheet well inside a black hole. No attempt was made to show that all the assumptions can in fact be satisfied, and it remains unclear whether they can. The constructions given in the present paper, by contrast, are completely explicit. Among them, the collapsing dust ball studied in Sec.~\\ref{sec-ball} is perhaps the example most similar to (yet definitely distinct from) the collapsing shell of Ref.~\\cite{CEB1}. ``Monsters''~\\cite{HsuRee07,HsuRee09} are matter configurations that violate the {\\em spacelike\\/} entropy bound, adding to a multitude of counterexamples~\\cite{RMP}. But monsters do obey the covariant bound, since the light-sheets off of their boundary surface are truncated by black hole singularities. No evidence has been presented that the stronger inequality $S\\lesssim A^{3/4}$ is violated on these light-sheets. The scaling of entropy like $\\Lambda^{-1}$ ($\\Lambda^{-2}$) in open FRW universes with positive (negative) cosmological constant was noted earlier~\\cite{BouLei09} in the context of the causal entropic principle~\\cite{Bou06,BouHar07}. \\paragraph{Outline} In Sec.~\\ref{sec-main}, we will study open FRW cosmologies with zero, positive, or negative cosmological constant. We will identify light-sheets with entropy $S\\gg A^{3/4}$, and we will show that in the limit of early curvature domination, the holographic bound can be saturated: $S\\to A$. In Sec.~\\ref{sec-patch} we will consider the past light-cone of an observer in such universes. We will show that the observer's sky can be filled with so much entropy as to saturate the covariant bound. This shows that the results of Sec.~\\ref{sec-main} correspond naturally to directly observable situations. As a corollary, we will find that the observable entropy in the presence of a cosmological constant can exceed the naive bound $|\\Lambda|^{-3/4}$, and can become as large as $\\Lambda^{-1}$ for $\\Lambda>0$, and $\\Lambda^{-2}$ for $\\Lambda<0$. In Sec.~\\ref{sec-ball}, we will show that light-sheets with $S\\gg A^{3/4}$ exist not only in open universes. We will demonstrate that such light-sheets can be actively produced by setting up a collapsing ball of pressureless matter with a particular velocity distribution. In Sec.~\\ref{sec-cd}, we go further and show that a light-sheet with $S\\gg A^{3/4}$ can be both set up and observed by a single observer. One relevant example arises when a black hole is slowly fed with quanta not much smaller than its own radius. In Sec.~\\ref{sec-conjecture} we will assess the possible role of the bound $S\\lesssim A^{3/4}$. We will briefly explore the conjecture that this bound holds for static, weakly gravitating systems. The appendix will discuss the application of our results to the generalized covariant entropy bound~\\cite{FMW}. ", "conclusions": "" }, "1003/1003.4884_arXiv.txt": { "abstract": "We study the distribution of exoplanets around main sequence (MS) stars and apply our results to the binary model for the formation of extreme horizontal branch (EHB; sdO; sdB; hot subdwarfs) stars. {{{By Binary model we refer both to stellar and substellar companions that enhance the mass loss rate, {{{{where substellar companions stand for both massive planets and brown dwarfs.}}}} }}} We conclude that sdB (EHB) stars are prime targets for planet searches. We reach this conclusion by noticing that the bimodal distribution of planets around stars with respect to the parameter $M_p a^2$, is most prominent for stars in the mass range $1M_\\odot \\la M_{\\rm star} \\la 1.5M_\\odot$; $a$ is the orbital separation, $M_{\\rm star}$ is the stellar mass and $M_p$ the planet mass. This is also the mass range of the progenitors of EHB stars that are formed through the interaction of their progenitors with planets (assuming the EHB formation mechanism is the binary model). In the binary model for the formation of EHB stars interaction with a binary companion or a substellar object (a planet or a brown dwarf), causes the progenitor to lose most of its envelope mass during its red giant branch (RGB) phase. As a result of that the descendant HB star is hot, i.e., an EHB (sdB) star. The bimodal distribution suggests that even if the close-in planet that formed the EHB star did not survive its RGB common envelope evolution, one planet or more might survive at $a \\ga 1 \\AU$. Also, if a planet or more are observed at $a \\ga 1 \\AU$, it is possible that a closer massive planet did survive the common envelope phase, and it is orbiting the EHB with an orbital period of hours to days. ", "introduction": "\\label{sec:intro} Horizontal branch stars (HB) are Helium burning stars that have evolved from main sequence stars (MS) through the red giant branch (RGB). During the RGB phase the star loses a non-negligible amount of mass. The amount of mass lost determines the properties of the descendant HB star; namely, its location on the HR diagram. The distribution of HB stars on the HR diagram, called HB morphology, has become a growing field of research because HB stars can be the main UV radiation source in old population (Dorman et al. 1993; Bertelli et al. 1996), their formation contains some unsolved problems, and HB stars can even act as standard candles (Fusi Pecci et al. 1996b). The formation and evolution of HB stars depend on the mass of the progenitor on the main sequence, initial Helium abundance (D'Antona et al. 2002), metallicity, and what is more relevant to our study, the mass lost during the RGB phase (Fusi Pecci et al. 1996a; Dorman et al. 1995). HB stars with low mass envelopes have small radii and they are hot. They are called extreme HB (EHB) stars {{{{ in photometric classification}}}} (other names are sdO or sdB or hot subdwarfs {{{{ according to spectroscopic classification; }}}} in this work we will use all these terms indistinguishably). To become an EHB star, the RGB progenitor must lose most of its envelope. The reason for some RGB stars to lose so much mass is a major unsolved issue in stellar evolution. The debate is whether a single star (e.g., Yi 2008) can account for the formation of hot subdwarfs, or whether binary evolution is behind the hot subdwarf phenomenon (e.g., Han et al. 2007). Supporting the binary model is the finding that about half of the sdB stars in the field (not in globular clusters) reside in close binaries with periods as short as one day or less (Maxted et al. 2001; Napiwotzki et al. 2004); the companions are either low-mass main sequence stars or white dwarfs {{{{ (WDs; e.g., Han et al. 2003; Geier et al. 2010, and references therein).}}}} Because the components' separation in these systems is much less than the size of the subdwarfs' RGB progenitors, these systems must have experienced a common envelope (CE) phase (e.g., Han et al. 2002, 2003), where the lower mass companion spirals inside the bloated envelope of the RGB star and finally ejects it. {{{{ Most of the formation channels of sdB stars are summarized by Han et al. (2003), although they omit the substellar channel. Stellar binary interaction can result in a stable Roche lobe overflow (RLOF). In that case an sdB star is formed, but the orbital separation stays large. A different scenario discussed in the literature is the merger of two helium WDs, as suggested by Webbink (1984; for more recent papers with more references see Han et al. 2003, Heber 2008 and Nelemans 2010). In this scenario, gravitational wave radiation causes two WDs with a small orbital separation to coalesce and form an sdB star. This scenario is supported by the research of sdB and sdO mass range done by Zhang et al. (2010). A very recent population synthesis of binary stars is reported by Nelemans (2010). His conclusion is that both interaction of RGB stars with substellar companions, and merger of He WDs can contribute to the formation of sdB stars and single He WD. However, the large number of single He WDs suggest that most of them are the descendent of the interaction of RGB stars with substellar companions. The population synthesis of Nelemans (2010) shows that different aspects of the interaction of RGB stars with substellar objects must be studied. Our present paper aim at comparison with known exoplanet properties. }}}} The CE ejection channel provides a reasonable explanation for the extra mass loss required to form sdB stars. But for about half of all analyzed subdwarfs there is no evidence for close stellar companions. {{{Moreover, in globular clusters stellar companions cannot explain the formation of sdB stars (Catelan 2009). }}} A solution which has received a major boost with the recent discovery by Geier et al. (2009), is that substellar objects influence the evolution of the RGB progenitor (Soker 1998a; Nelemans \\& Tauris 1998; Soker \\& Harpaz 2000, 2007; Soker \\& Hershenhorn 2007; Politano et al. 2008; {{{ Villaver \\& Livio 2007, 2009; Carlberg et al. 2009; }}} Bear \\& Soker 2010; {{{ Nordhaus et al. 2010}}}). {{{{It is also possible that the planets were formed along with the sdB star (second generation planets, see Perets 2010) and are a result of the merger of the two WDs (Silvoti 2008). This explanation has the same shortcomings discussed before.}}}} {{{{ In the present paper substellar objects will stand for both massive planets and brown dwarfs. Grouping of brown dwarfs and massive planets (gas giant planets) has its own merit. Lovis \\& Mayor (2007), for example, raise the possibility that massive planets and brown dwarfs are formed in the same process. A key issue here is that it is more easy to detect brown dwarfs, and they are more likely to survive the RGB phase. However, as there are more planets than brown dwarfs, they are likely to play a larger role than brown dwarfs. Also, as the statistics for planets around main sequence stars is much better than that for brown dwarfs, in this paper we deal only with planets. Adding brown dwarfs will further increase the merit of the planet-induced formation of sdB stars. }}}} {{{{ Substellar objects are known to accompany many different stars in different stages of their life. There is a large body of literature and research on relevant substellar companions. Here we mention a few examples. Machalek et al. (2010) study XO-3b, a high mass hot Jupiter planet ($M_p=11.79\\pm 0.59M_J$), on the verge of deuterium burning, and orbiting an F5V parent star. Another example is the detection of substellar companion with a mass of $M_2 \\sin(i)=2.9M_J$ that orbits HD145457 (a K0 giant of $1.9M_\\odot$) with an orbital period of $P=176$d (Sato et al. 2010). As indicated by Schuh et al. (2010), the increasing number of substellar companions to sdB stars may indicate the existence of an undiscovered population. We do not wish to solve the question of how sdBs were formed, as of now it seems that each scenario might be possible under specific conditions (Geier et al. 2009; Soker 1998a; Han et al. 2002, 2003; Lisker et al. 2005; Nelemans 2010). We deal here with the substellar scenario (planet induced) for the formation of sdB stars. }}}} Geier et al. (2009) announced recently the discovery of a close substellar companion to the hot subdwarf (EHB) star HD 149382. The orbital period is very short, 2.391~days, implying that the substellar companion had evolved inside the bloated envelope of the progenitor RGB star (a CE phase). The mass of the companion is $8-23 M_J$, so either it is a planet or a low mass brown dwarf. This discovery supports the prediction of Soker (1998a) that such planets can survive the common envelope (CE) phase, and more relevant to us, that planets can enhance the mass loss rate on the RGB and lead to the formation of EHB. Other planets that orbit EHB at larger separations have been detected (Silvotti et al. 2007; Lee et al. 2009; Qian et al. 2009). Silvotti et al. (2007) announced the detection of a planet with a mass of $3.2 M_{\\rm J}$, an orbital separation of $1.7 \\AU$, and an orbital period of $P=3.2 \\yr$ around the hot subdwarf V391 Pegasi. Serendipitous discoveries of two substellar companions around the eclipsing sdB binary HW~Vir at distances of $3.6 \\AU$ and $5.3 \\AU$ with orbital periods of $3321~$d and $5767~$d (Lee et al. 2009) and one brown dwarf around the similar system HS~0705+6700 with a period of $2610~$d and a separation of $<3.6 \\AU$ (Qian et al. 2009) followed recently. It is quite plausible that closer planets did interact with the RGB progenitor of the sdB star; they are not observed in these systems. In the present paper we examine whether the known exoplanets support such a scenario. We end by noting that all these substellar companions have been detected in the field. It is commonly assumed that planets don't exist in large enough numbers in globular clusters. However, one planet has been detected in the M4 globular cluster (Sigursson et al. 2003; Beer et al. 2004 and references there in), and the role of planets in the formation of EHB in globular clusters, where metallicity is very low, is an open question. In this paper we are aiming at field stars. ", "conclusions": "\\label{sec:summary} Our goal is to use the properties of known exoplanets to better understand the role planets play in the formation of extreme horizontal branch (EHB; sdO; sdB; hot subdwarfs) stars, and the distribution of planets around EHB stars. EHB stars are hot HB stars with a very low mass envelope. The explanation is that their progenitor RGB star has lost most of its envelope on the RGB. The key process, in cases where there is no close stellar companion, is that a planet or a stellar companion caused this enhanced mass loss process. We focus on the role of planets. To lose most of its mass by interaction with a planet on the RGB the star cannot be too massive. On the other hand the minimum mass is determined by evolution time scale. This limits the progenitor mass of field stars to be $M \\ga 1 M_\\odot$. Over all, the relevant mass range for the main sequence (MS) progenitor is $1M_\\odot \\le M_{\\rm Pro} \\le 1.5M_\\odot$. Following Soker \\& Hershenhorn (2007), we examined the distribution of planets according to $M_pa^2$; this is done in Fig. \\ref{fig:All_planets_2}. This figure reproduces the well known double peak distribution. We examined the distribution for three groups of parent star mass. From Figs. \\ref{fig:All_planets_1} and \\ref{fig:All_planets_3} it is evident that the double peak distribution is strong only in the middle mass range $1M_\\odot \\le M_{\\rm Pro} \\le 1.5M_\\odot$ (we note that selection effects might be important for the upper mass range). This middle mass range coincides with that of the progenitors of EHB stars formed by interaction with planets and brown dwarfs. We then examined (Figs. \\ref{fig:Multi_1}, \\ref{fig:Multi_2} and \\ref{fig:Multi_3}) the double-peak distribution for multi-planet systems. We found that the double-peak distribution also holds for these systems in the middle mass range. In the binary model for the formation of EHB stars, if there is no stellar companion to the EHB star, most likely its progenitor interacted with one or more planets or brown dwarfs. Planets close to the progenitor, mainly in the left-peak in our figures, will enter the CE phase at an early stage and will be destroyed (Soker 1998a). Still, they can enhance the mass loss rate and lead to the formation of an EHB star (Soker 1998a). Our results suggest that in many cases there are also planets in the right-peak, that can survive the RGB evolution. The bimodality of planets in multiplanet system suggests that when we observe an ``outer'' planet ($a_p\\geq 1AU$) around an EHB star, another close in planet was probably engulfed during its formation process. Moreover, if we observe a close in planet ($a\\ll 1AU$) or even if we do not observe it, planets around EHB stars are likely to reside in the outer regions at $1 \\leq a_p \\leq 10AU$. We therefore encourage a search for outer planets around EHB that have close planet as well. Moreover, even if there is no close companion (stellar or substellar), there is a high chance of the existence of an outer planet around the EHB star. Furthermore, if an outer planet is found, most likely another planet (or more) went through the CE phase and caused the RGB progenitor to lose most of its envelope. In some cases this closer-in planet might survive the CE phase, and be found around the EHB star. Our general conclusion from this study is that a single EHB (sdBO) star is likely to have an outer planet(s) in an orbital separation of $1 \\leq a_p \\leq 10AU$. We note that red HB stars (these are stars that maintained most of their envelope) might also have planet at large orbital separations. In such cases either the progenitor was too massive ($\\ga 1.5 M_\\odot$) for an inner planet to expel most of the progenitor envelope, or there were no massive close in planets at all. For example, Mercury and Venus will be engulfed when the Sun evolves of the RGB. Earth might also be engulfed. However, these three planets do not contain enough mass to enhance the mass loss rate from the Sun. Therefore, in 6-7~Gyr Jupiter will orbit a red HB star. Comparing Fig. \\ref{fig:Multi_2} and Fig. \\ref{fig:single_2} should be done with caution, taking into account that the sample of this statistics is not large. However, it appears that according to our model the expected number of red HB with outer planets is $\\sim 5$ times as high as that of EHB with outer planets. Once observations will increase the number of main sequence multi-planet systems and the number of planets around HB stars, a population synthesis should be conducted in order to achieve a better estimate. These conclusions hold as well when the inner object is a low mass MS star (mainly M-type). Indeed, HW Vir and HS 0705+6700 are such close binary systems with substellar objects around them. We encourage the search of planets around similar binary systems, e.g., PG 1336-018 (Kilkenny et al. 1993; Drechsel et al. 2001). We strongly suggest to look at EHB (sdO; sdB; hot subdwarfs) stars as prime targets of planet search. We thank an anonymous referee for comments that improved our manuscript. This research was supported by the Asher Fund for Space Research at the Technion, and the Israel Science foundation. E.B. was supported in part by the Center for Absorption in Science, Ministry of Immigrant Absorption, State of Israel." }, "1003/1003.3695_arXiv.txt": { "abstract": "Using the momentum-dependent MDI effective interaction for nucleons, we have studied the transition density and pressure at the boundary between the inner crust and liquid core of hot neutron stars. We find that their values are larger in neutrino-trapped neutron stars than in neutrino-free neutron stars. Furthermore, both are found to decrease with increasing temperature of a neutron star as well as increasing slope parameter of the nuclear symmetry energy, except that the transition pressure in neutrino-trapped neutron stars for the case of small symmetry energy slope parameter first increases and then decreases with increasing temperature. We have also studied the effect of the nuclear symmetry energy on the critical temperature above which the inner crust in a hot neutron star disappears and found that with increasing value of the symmetry energy slope parameter, the critical temperature decreases slightly in neutrino-trapped neutron stars but first decreases and then increases in neutrino-free neutron stars. ", "introduction": "Studying the properties of neutron stars allows us to test our knowledge on the properties of nuclear matter under extreme conditions. Theoretical studies have shown that a neutron star is expected to have a liquid core surrounded by an inner crust~\\cite{Cha08}, which extends outward to the neutron drip-out region. While the neutron drip-out density $\\rho _{\\rm out}$ has been relatively well determined~\\cite{Rus06}, the transition density $\\rho _{t}$ at the inner edge of the crust is still quite uncertain because of our limited knowledge on the nuclear equation of state (EOS), especially the density dependence of the symmetry energy ($E_{\\rm sym}(\\rho)$) of neutron-rich nuclear matter~\\cite{Lat00,Lat07}. Recently, significant progress has been made in constraining the EOS of neutron-rich nuclear matter using terrestrial laboratory experiments (See Ref.~\\cite{LCK08} for a recent review). In particular, from analyses of experimental data on neutron skin thickness, isobaric analogue states, Pygmy dipole resonances, and giant dipole resonances in nuclei as well as on isospin diffusion, isoscaling, and neutron-proton to triton-$^3$He ratio in intermediate-energy nuclear reactions, significant constraints on $E_{\\rm sym}(\\rho)$ have been obtained for the same sub-saturation density region as expected in the inner edge of neutron star crusts. The extracted slope parameter $L=3\\rho_0 (\\partial E_{\\rm sym}(\\rho)/\\partial\\rho)_{\\rho=\\rho_0}$ of the nuclear symmetry energy from these studies has values in the range $30$ MeV $< L < 80$ MeV~\\cite{She10}. With the MDI interaction together with the value $L=86\\pm25$ MeV constrained from an analysis of the isospin-diffusion data~\\cite{Tsa04,Che05a,LiBA05,Tsa09} in heavy-ion collisions using the isospin-dependent Boltzmann-Uehling-Uhlenbeck (IBUU) transport model with the momentum-dependent MDI interaction~\\cite{Das03}, the density and pressure at the inner edge of the crust of cold neutron stars were studied by considering the boundary of the instability region or the spinodal boundary between the liquid core and inner crust of a cold neutron star in both the thermodynamical approach~\\cite{Kub07,Lat07} and the dynamical approach~\\cite{BPS71,BBP71,Pet95a,Pet95b,Oya07}. This leads to the constraints $0.040$ fm$^{-3}<\\rho_t<0.065$ fm$^{-3}$ and $0.01$ MeV/fm$^{3} < P_t < 0.26$ MeV/fm$^{3}$, respectively, for the transition density and pressure. Together with the crustal fraction of the total moment inertia of the Vela pulsar extracted from its glitches~\\cite{Lin99}, a tighter constraint on the mass-radius relation of cold neutron stars was obtained~\\cite{XCLM09}. Because of the initial high temperature and appreciable proton fraction in a newly-formed neutron star immediately after gravitational collapse of a massive star~\\cite{Bur86,Bet90,Bur03}, neutrinos are abundantly produced from the Urca process in its inner core. Although high energy neutrinos can be trapped at densities as low as 10$^{12}$ g/cm$^3$~\\cite{BBLA79}, the stars cools by neutrino emissions. As neutrinos are emitted from this so-called proto-neutron star, which has an initial temperature of $\\sim 10^{11}K$ (about $10$ MeV)~\\cite{Bur88,Hor04}, its temperature drops to $\\sim 10^{10} K$ (about $1$ MeV) and even lower. Afterwards, the neutron star becomes transparent to neutrinos as their mean free path increases with decreasing energy, and the cooling of the neutron star continues to be dominated by neutrino emission for a long time. It is thus of interest to study the transition density and pressure in newly-born hot neutron stars, as this would help to understand the cooling mechanism and structural evolution of neutron stars. In this paper, we extend the study of Ref.~\\cite{XCLM09} to finite temperature and study the dependence of the transition density and pressure of hot neutron stars on the nuclear symmetry energy, particularly its slope at nuclear saturation density. This paper is organized as follows. We first review in Sec.~\\ref{model} the momentum-dependent MDI interaction for nucleons, in Sec.~\\ref{nsmatter} the properties of hot neutron star matter, and in Sec.~\\ref{approaches} the dynamical approach for locating the inner edge of the crust of a hot neutron star. We then show in Sec.~\\ref{results} the results and conclude with a summary in Sec.~\\ref{summary}. ", "conclusions": "\\label{results} In this section, we show the temperature dependence of the transition density and pressure in newly-born hot neutron stars by using the MDI interaction with different values for the symmetry energy parameter $x$ or the slope parameter $L$ of the symmetry energy. \\begin{figure}[h] \\centerline{\\includegraphics[scale=0.9]{fig4.EPS}} \\caption{(Color online) Transition densities $\\rho_t$ ((a) and (c)) and pressure $P_t$ ((b) and (d)) as functions of the slope parameter $L$ of the symmetry energy at different temperatures for both the neutrino-trapped matter ((a) and (b)) and the neutrino-free matter ((c) and (d)).} \\label{rhotL} \\end{figure} The dependence of the transition density and pressure in hot neutron stars on the slope parameter $L$ of the symmetry energy at different temperatures is shown in Fig.~\\ref{rhotL}. It is seen that the transition density $\\rho_t$ generally decreases with increasing value of the slope parameter $L$ of the symmetry energy. As the transition density can be viewed approximately as the beginning of a first-order liquid-gas phase transition, a stiffer symmetry energy, which corresponds to a softer equation of state at subsaturation densities, leads thus to a smaller phase transition density and therefore a lower core-crust transition density. Furthermore, the transition density decreases with increasing temperature, and for the neutrino-free matter this is more pronounced for larger values of $L$. The temperature effect can again be understood by the decreasing phase transition density with increasing temperature. For the neutrino-trapped matter, the $L$-dependence of the transition density is relatively weak. The weak $L$-dependence is mainly due to the fact that the isospin asymmetry is not so large as shown in Fig.~\\ref{rhodelta}. We note that a similar temperature dependence of the transition density has been obtained in studies based on Skyrme interactions and relativistic mean field models~\\cite{Duc08,Par09,Ava09}. For the $L$-dependence of the transition density, results from these models are, however, not so clear as different values for the incompressibility $K_0$ of the symmetric matter at saturation density and $E_{sym}(\\rho_0)$, which also affect the value of the transition density, have been used. For the transition pressure $P_t$, its value in the neutrino-trapped matter decreases only very slightly with increasing $L$ as a result of the weak $L$-dependence of $\\rho_t$. Its temperature dependence shows, however, a complicated behavior of slightly higher and smaller values at higher temperatures for smaller and larger values of $L$, respectively. This is due to the fact that although the transition density $\\rho_t$ decreases with increasing temperature, the contribution from leptons increases with increasing temperature. For the neutrino-free matter, $P_t$ is seen to decrease rapidly with increasing $L$. As to its temperature dependence, $P_t$ in the neutrino-free matter decreases with increasing temperature at larger values of $L$ but shows a weaker temperature dependence for smaller values of $L$. Also, $P_t$ is larger for the neutrino-trapped matter than for the neutrino-free matter. Interestingly, $P_t$ becomes very small and even negative at $T=1$ with larger $L$ for the neutrino-free matter. This is due to the smaller contributions from the leptons and the asymmetric part of the nuclear interactions to the total pressure. Since the pressure at the inner edge of neutron star crust cannot be negative, our finding thus indicates that either the neutrino-free matter in hot neutron stars cannot reach a temperature above $T=1$ MeV or the symmetry energy cannot have a slope parameter larger than $L\\sim 100$ MeV. \\begin{figure}[h] \\centerline{\\includegraphics[scale=0.9]{fig5.EPS}} \\caption{(Color online) Transition densities $\\rho_t$ ((a) and (c)) and pressure $P_t$ ((b) and (d)) as functions of temperature $T$ with $x=0$ and $x=-1$ for both the neutrino-trapped matter ((a) and (b)) and the neutrino-free matter ((c) and (d)). Note that different scales for temperatures are used for the neutrino-trapped matter and the neutrino-free matter.} \\label{rhotT} \\end{figure} The temperature effect on the transition density and pressure for a fixed symmetry energy parameter is demonstrated in Fig.~\\ref{rhotT} for the symmetry energy parameters $x=1$, $x=0$ and $x=-1$. For the neutrino-trapped matter, the temperature effect is similar for all three $x$ values, and the transition density $\\rho_t$ decreases almost linearly with increasing temperature at lower $T$ and decreases quickly at higher $T$. For the neutrino-free matter, although the transition density $\\rho_t$ decreases smoothly with increasing temperature for $x=1$, it decreases slowly (quickly) at lower temperatures with $x=0$ ($x=-1$) while quickly (slowly) at higher temperatures. The temperature effects on the transition density reflect those on the spinodal region and the abundance of particle species in the hot neutron star matter. For the neutrino-trapped matter, the transition pressure $P_t$ is seen to be insensitive to the temperature for $x=-1$ but increases slightly with increasing temperature for $x=0$ and $x=1$ at lower $T$, and it decreases with increasing temperature at higher $T$ for all values of $x$. For the neutrino-free matter, it is insensitive to temperature for $x=1$ but decreases with increasing temperature for $x=0$, while for $x=-1$ it drops to a negative value and becomes positive again as the temperature increases. Our results again show that the behavior of $P_t$ is dominated by that of $\\rho_t$ and the contribution from the leptons, with the former decreasing and the latter increasing with increasing temperature. \\begin{figure}[h] \\centerline{\\includegraphics[scale=0.9]{fig6.EPS}} \\caption{(Color online) Critical temperature $T_c$ as a function of the slope parameters $L$ of the symmetry energy for both the neutrino-trapped matter and the neutrino-free matter. } \\label{rhotTc} \\end{figure} We have seen in Fig.~\\ref{rhonrhop} that as the temperature of the neutron star matter increases, there will be eventually no cross point between the curve of neutron-proton relative abundance and the boundary of the spinodal region, leading to the disappearance of the transition density in hot neutron stars. In such case, the inner crust (nuclear 'pasta' phase) disappears and the liquid core expends directly to the outer crust. To determine the critical temperature $T_c$ at which the transition density $\\rho_t$ disappears is thus useful for understanding the structural evolution of newly-born hot neutron stars. Figure~\\ref{rhotTc} displays the $L$-dependence of the critical temperature for both the neutrino-trapped matter and the neutrino-free matter. One sees that the critical temperature $T_c$ decreases slightly with increasing value of $L$ in the neutrino-trapped matter, but it first decreases and then increases with increasing $L$ in the neutrino-free matter. This complicated behavior is due to the isospin and temperature effects on the spinodal region and the relative neutron-proton abundance as shown in Fig.~\\ref{rhonrhop}. Our results thus indicate that for neutrino-trapped neutron stars of temperatures higher than $12$ MeV or for neutrino-free neutron stars of temperatures higher than $1.5$ MeV, there exists no inner crust if the value of $L$ is $70 \\sim 80$ MeV. As newly-born hot neutron stars cool, the temperature at which the inner crust can form thus depends on the density dependence of the symmetry energy at subsaturation densities. Again, the magnitude of the critical temperature for both neutrino-trapped and neutrino-free matter is similar to those from Skyrme interactions and relativistic mean field models~\\cite{Duc08,Par09}. The above results were obtained using the dynamical approach that includes the effects of both the density gradient terms and the Coulomb interaction. Neglecting these effects, the resulting thermodynamical approach given by Eq.~(\\ref{ther}) gives higher values for both the transition density and pressure. This is especially the case for the neutrino-trapped matter and/or for smaller values of $L$ as more electrons are present in such hot neutron stars and the effect due to the Coulomb interaction becomes more important. Also, there are recently some studies on the transition density in neutron stars using various nucleon-nucleon interactions~\\cite{Sur09,Vid09}. To see the effect of momentum dependence in the nucleon-nucleon interaction on the transition density and pressure, we have also repeated above calculations using the momentum-independent MID interaction~\\cite{Xu08}, which gives the same equation of state for asymmetric nuclear matter but different single-particle mean-field potential in comparison with the momentum-dependent MDI interaction used in the present study, and we find that the momentum-dependent effect on the transition density and pressure in hot neutron stars is small." }, "1003/1003.4626_arXiv.txt": { "abstract": "{It is generally assumed that a large fraction of stars are initially born in clusters. However, a large fraction of these disrupt on short timescales and the stars end up belonging to the field. Understanding this process is of paramount importance if we wish to constrain the star formation histories of external galaxies using star clusters.} {We attempt to understand the relation between field stars and star clusters by simultaneously studying both in a number of nearby galaxies. } {As a pilot study, we present results for the late-type spiral NGC 4395 using HST/ACS and HST/WFPC2 images. Different detection criteria were used to distinguish point sources (star candidates) and extended objects (star cluster candidates). Using a synthetic CMD method, we estimated the star formation history. Using simple stellar population model fitting, we calculated the mass and age of the cluster candidates.} {The field star formation rate appears to have been roughly constant, or to have possibly increased by up to about a factor of two, for ages younger than $\\sim$300 Myr within the fields covered by our data. Our data do not allow us to constrain the star formation histories at older ages. We identify a small number of clusters in both fields. Neither massive ($>10^5$ M$_\\odot$) clusters nor clusters with ages $\\geq1$ Gyr were found in the galaxy and we found few clusters older than 100 Myr.} {Based on our direct comparison of field stars and clusters in NGC~4395, we estimate the ratio of star formation rate in clusters that survive for $10^7$ to $10^8$ years to the total star formation to be $\\Gamma\\sim0.03$. We suggest that this relatively low $\\Gamma$ value is caused by the low star formation rate of NGC~4395.} ", "introduction": "It is commonly assumed that most (if not all) stars are formed in clusters. Clusters, due to dynamical and stellar evolution, dissolve \\citep{spitzer87} and the stars that belong to them become part of the field stellar population of the galaxy. To use clusters as effective tools to constrain the formation and evolution of galaxies, it is necessary to improve our understanding of what fraction of stars end up as members of bound clusters and in the field, respectively. Both field stars and star clusters have been studied extensively in the Milky Way and nearby galaxies. In principle, field stars hold information about star formation histories over the entire Hubble time and can therefore provide important information about galaxy formation and evolution, e.g., \\citet[][]{edvardsson93}. Main sequence stars are present at all ages, and observations of the main sequence turn-off can constrain epochs of stellar formation especially in dwarf galaxies where relatively distinct bursts are often observed \\citep{mateo98}. Other features of the color-magnitude diagram (CMD), such as the sub-giant and horizontal branches, giant and asymptotic branches, red clump stars, and red and blue supergiants can provide information about specific epochs of star formation. However, it is necessary to apply a more sophisticated modeling of the CMD to reconstruct star formation histories, taking into account different effects, e.g., incompleteness, resolution, depth of the observations, extinction and chemical composition. \\citet{tosi91} developed a method that takes these effects into account and attempts to use all the information available in a CMD to reconstruct the field star formation history of a galaxy. The synthetic CMD method creates a synthetic population that is compared to the observed CMD to constrain the star formation history (SFH) of the field stars in a galaxy. Many subsequent studies have refined this method \\citep{dolphin97,harriszaritsky01,coimbra02}. The SFHs of many galaxies in the Local Group and nearby have been studied using the synthetic CMD method: \\citet{brown08} for M31, \\citet{harriszaritsky09} for the LMC, \\citet{harriszaritsky04} for the SMC, \\citet{barker07} for the M33, \\citet{cole07} for Leo A, \\citet{young07} for Phoenix, \\citet{annibali09} for NGC 1705, \\citet{williams09} for M81, \\citet{larsen07} for NGC1313 and \\citet{rejkuba04} for NGC 5128, among many others. On the other hand, the study of cluster systems and disruption processes can provide important insight into the origin of field stars in a galaxy. Clusters disrupt by means of a variety of mechanisms, including ``infant mortality'', stellar evolution, two-body relaxation, and tidal shocks that have been extensively studied by \\citet{BL03,LL03,lamers05,baumgardt09,fall06,elmegreen08,whitmore07}, and \\citet{bastiangieles08} among others. The analysis of star clusters is often based on a comparison between the observed spectral energy distribution and theoretical models, which provides information about the ages and masses of the studied clusters and, in turn, their dissolution. However, the comparison of models with observations also remains affected by many uncertainties, e.g., binarity, completeness effects. Some studies have started to address the relation between field stars and clusters more explicitly. The number of stars that were formed in clusters has been estimated to be $70\\%-90\\%$ in the solar neighbourhood, while $50\\%-95\\%$ of these embedded clusters dissolve in a few Myrs \\citep{LL03,lamersgieles08}. \\citet{gielesbastian08} estimated that only 2-4\\% of the global star formation rate in the SMC happened in bound star clusters. It is currently unknown what fraction of stars are initially born in clusters in SMC, but if this fraction were as large as in the Solar neighbourhood this would imply that there is also a large infant mortality rate in the SMC. \\citet{bastian08} studied the relation between the cluster formation rate and the star formation rate ($\\Gamma=\\frac{CFR}{SFR}$) using archival Hubble images for high star formation rate galaxies and additional galaxies/clusters from literature. He found this fraction to be $\\Gamma\\sim0.08$ and thus concluded that the fraction of stars formed in (bound) clusters represents $8\\%$ of the total star formation. On the other hand, \\citet{gieles09} found that $\\Gamma$ is given by $0.05\\leq\\Gamma\\leq0.18$ in the galaxies M74, M101, and M51. In most of these cases, however, it is difficult to tell whether there is a genuine ``field'' mode of cluster formation, or whether \\emph{all} stars form initially in clusters of which a large fraction dissolves rapidly. We aim to analyze the relation between field stars and star clusters in different environments and address the question of whether or not there is a constant cluster formation ``efficiency''. To this end, we use \\textit{Hubble Space Telescope} (HST) images of a set of five galaxies (NGC 4395, NGC 1313, NGC 45, NGC 5236, and NGC 7793), which are nearby, face-on, spirals that differ in their morphologies, star, and cluster formation histories. The cluster systems of these galaxies were studied by \\citet{mora09} using the same Hubble images analyzed in this work. Mora et al. observed significant variations in the cluster age distributions of these galaxies. The galaxies are sufficiently nearby ($\\sim4$ Mpc) for the brighter field stars to be well resolved in HST images, so that (recent) field star formation histories can be constrained by means of the synthetic CMD method. We can therefore take advantage of the superb spatial resolution of HST images to study field stars and clusters simultaneously within specific regions of these galaxies. As a pilot work, this paper is devoted to presenting and testing all the analysis procedures, such as detection of stars and cluster candidates, completeness tests and photometry, and derivation of field star and cluster age distributions. We discuss our implementation of the synthetic CMD method as an IDL program and carry out tests of this program. To test our procedures we use the galaxy NGC 4395. The methods described in this paper will be used in our study of the rest of the galaxy sample (Silva-Villa et al.\\ 2010, in prep.). This article has the following structure. In Sect. 2, we provide general information about NGC~4395. We present the observations, data reduction, and photometry, where we differentiate the point sources (star candidates) from the extended objects (cluster candidates) in Sects. 3 and 4. Section 5 is devoted to the analysis of the star and cluster properties. Finally, in Sect. 6 we present the discussion and conclusions. \\begin{table*}[!t]\\centering \\newcommand{\\DS}{\\hspace{6\\tabcolsep}} % \\caption{Journal of HST/ACS and HST/WFPC2 observations for both fields (F1 and F2) in NGC 4395.} \\begin{tabular}{c @{\\DS} cccc } Proposal ID & Date & Filter & Total exp. time (s) \\\\ \\hline \\hline \\multicolumn{3}{c}{NGC 4395 (J2000.0) $\\alpha:12^h 26^m 00^s$ $\\delta: +33^o 31^{'} 04^\"$}\\\\ 9774 & 2004 Jun 12 & F435W (B) & 680 \\\\ 9774 & 2004 Jun 12 & F555W (V) & 680 \\\\ 9774 & 2004 Jun12 & F814W (I) & 430 \\\\ \\hline \\multicolumn{3}{c}{NGC 4395 (J2000.0) $\\alpha:12^h 26^m 00^s.50$ $\\delta: +33^o 30^{'} 58^\".3$}\\\\ 9774 & 2004 Jun 12 & F336W (U) & 2400 \\\\ \\hline \\multicolumn{3}{c}{NGC 4395 (J2000.0) $\\alpha:12^h 25^m 45^s.20$ $\\delta: +33^o 34^{'} 28^\"$}\\\\ 9774 & 2004 Jun 18 & F435W (B) & 680 \\\\ 9774 & 2004 Jun 18 & F555W (V) & 680 \\\\ 9774 & 2004 Jun 18 & F814W (I) & 430 \\\\ \\hline \\multicolumn{3}{c}{NGC 4395 (J2000.0) $\\alpha:12^h 25^m 42^s.72$ $\\delta: +33^o 34^{'} 22^\".6$}\\\\ 9774 & 2004 Jun 18 & F336W (U) & 2400 \\\\ \\hline \\hline \\end{tabular} \\label{tab:data} \\end{table*} \\section {NGC 4395} According to the NASA/IPAC Extragalactic Database (NED), NGC~4395 is a late-type spiral classified as type SA(s)m. It harbours the closest and least luminous known example of a Seyfert 1 nucleus \\citep{filippenkosargent89}. Following \\citet{LR99}, we adopt a distance modulus of $(m-M)=28.1 \\ (D\\approx4.2$ Mpc) and an absolute magnitude $M_B = -17.47$, intermediate between the Small and Large Magellanic Cloud. We assume a Galactic foreground extinction for NGC~4395 of $A_B=0.074$ \\citep{schlegel98}. \\begin{figure}[!h] \\centering \\includegraphics[width=\\columnwidth]{ngc4395-1.eps} \\caption{Second Palomar Sky Survey images of NGC 4395 combined using Aladin software. HST/ACS (white lines) and HST/WFPC2 (yellow lines) fields covered by our observations are indicated.} \\label{fig:ngc4395} \\end{figure} The cluster system of NGC~4395 was first studied by \\citet{LR99}, using ground-based imaging. Using multiband ($UBVRIH\\alpha$) photometry, Larsen \\& Richtler identified 2 young clusters in this galaxy, although their observations were limited to objects with $M_V\\leq -8.5$. In their work, NGC 4395 is part of a sample of 21 galaxies. Compared to the remaining sample, NGC 4395 exhibits an exceptionally small number of star clusters and a low specific luminosity \\citep[ratio of cluster- to total galaxy light;][]{LR00}. Using HST images, \\citet{mora09} could detect clusters based on their sizes and found a total of 44 clusters in NGC~4395 to a magnitude limit of $M_B = -3.3$. Compared to the other 4 galaxies in their sample, Mora et al.\\ reached the same conclusions as Larsen \\& Richtler, i.e., NGC~4395 has a small number of star clusters. Mora et al. estimated the ages and masses of the clusters detected, showing that the star cluster system contains no objects with ages older than $10^9$ yr and includes clusters with masses ranging from $200$ to $\\sim3\\times10^4$ M$_{\\odot}$. Until now, no study of resolved field stars has been performed in this galaxy. However, \\citet{LR00} found that NGC 4395 has one of the lowest area-normalised star formation rates of all objects in their sample. \\section {Observation and data reduction} Images of NGC 4395 were taken using the Wide Field Channel on the {\\it Advanced Camera for Surveys} (ACS/WFC) and the {\\it Wide Field Planetary Camera 2} (WFPC2), both on board the {\\it Hubble Space Telescope} (HST). The resolution of the detectors are $0\\farcs05$, $0\\farcs046$, and $0\\farcs1$ per pixel for ACS/WFC, WFPC2/PC, and WFPC2/WF, respectively. At the distance of NGC 4395, $0\\farcs05$ corresponds to a linear scale of $\\sim 1$~pc. Two different fields of the galaxy were observed, covering the two spiral arms (see Fig. \\ref{fig:ngc4395}). The images were taken using the filters F336W ($\\sim $ U) using WFPC2 and F435W ($\\sim $ B), F555W ($\\sim $ V) and F814W ($\\sim $ I) using ACS, for each field. Each exposure was divided into two sub-exposures to eliminate cosmic-ray hits. For the ACS images, these sub-exposures were also dithered to allow removal of cosmetic defects. Table \\ref{tab:data} summarizes the observations. The data were processed with the standard STScI pipeline. The raw ACS images were drizzled using the {\\it multidrizzle} task \\citep{koekemoer02} in the STSDAS package of IRAF\\footnote{IRAF is distributed by the National Optical Astronomical Observatory (NOAO), which is operated by the Association of Universities for Research in Astronomy, Inc, under cooperative agreement with the National Science Foundation}. The default parameters were used, but automatic sky subtraction was disabled. The WFPC2 images were combined and corrected for cosmic rays using the {\\it crrej} task with the default parameters. \\subsection {Object detection} At the distance of NGC~4395, the spatial resolution of the HST images allows us to distinguish field stars (point sources with typical $FWHM\\sim2$ pixels) and star clusters (extended sources with typically $FWHM\\sim3$ pixels or greater) in the images. We analyzed the two separately and applied different detection criteria optimised for each type of object: \\begin {itemize} \\item {\\it Field star candidates (point sources)}\\\\ To detect field stars, we created an averaged image (using the bands B, V, and I) for each field. We ran the {\\it daofind} task in IRAF for the detection of stars using a $4\\sigma$ detection threshold and a background standard deviation of $\\sim0.02$ in units of counts per second. \\item {\\it Star cluster candidates (extended sources)}\\\\ Cluster detection was also performed on the averaged image. The object detection was performed using SExtractor V2.5.0 \\citep{bertinarnouts96}. The parameters used as input for this program were 6 connected pixels, all of them with 10$\\sigma$ over the background, to remove point- and spurious sources as much as possible. From the output file of SExtractor, we kept the coordinates of the objects detected and the FWHM calculated by this program. The coordinates found with SExtractor were then passed to {\\it ishape} in BAOLab \\citep{larsen99}. {\\it Ishape} models star clusters as analytical \\citet{king62} profiles (with concentration parameter $t_{tidal}/r_{core}=30$) and takes the instrument's PSF, created over the average image during the photometry, into account (see Sect. 4 for details on the creation of the PSF). By minimization of a $\\chi^2$-like function, {\\it ishape} calculates the best-fit cluster coordinates, size, i.e., FWHM, signal-to-noise ratio, and the $\\chi^2$ of the best fit. From the output of {\\it ishape} we saved the coordinate of the objects, the FWHM, the $\\chi^2$ of the fit, and the signal-to-noise ratio (calculated within the fitting radius of 4 pixels). For each WFPC2 chip, between five and seven common stars were visually selected and used to convert the coordinate list from ACS to WFPC2 frames using the task {\\it geomap} in IRAF. The transformations had an {\\it rms} of $\\sim 0.15$ pixels. \\end{itemize} \\section {Photometry} The photometry of NGC 4395 was performed following standard aperture and PSF fitting photometry procedures as described in the following. \\subsection {Field stars} Because of the crowding in our fields, we performed PSF fitting photometry to study the field stars. A set of bona fide stars were selected by eye to construct the PSF (for each band a different group of stars were used because the same stars might have different brightnesses in different bands). The PSF photometry was performed with DAOPHOT in IRAF. We selected the PSF stars by measuring the FWHM (using {\\it imexamine}) and selecting point sources smaller than $FWHM\\approx2.2$ pixels. As far as possible, we tried to include isolated stars distributed over the whole image. The raw magnitudes were converted to the Vega magnitude system using the HST zero-points taken from HST webpages\\footnote{http://www.stsci.edu/hst/acs/analysis/zeropoints/\\#tablestart} after applying aperture corrections to a nominal $0\\farcs5$ aperture (see sect. 4.4 for a more detailed description of how aperture corrections were determined). The zero-points used were $ZP_B$ = 25.767, $ZP_V$ =25.727, and $ZP_I$ = 25.520 magnitudes. Figure \\ref{fig:magerrors} shows the errors in our PSF photometry ($1\\sigma$ error) versus magnitude. The errors increase strongly below magnitudes of $\\sim26$ in each band, corresponding to absolute limits of $\\sim-2$ mag at the distance of NGC~4395. \\begin{figure}[!t] \\centering \\includegraphics[trim= 0mm 0mm 0mm 10mm, width=\\columnwidth,height=80mm]{magerrors_lres.eps} \\caption{Magnitude errors for the stars detected using PSF fitting photometry over the two fields for the bands B, V, and I} \\label{fig:magerrors} \\end{figure} A total of $\\sim$30 000 stars were found in each field. A combined Hess diagram for both fields is shown in Fig. \\ref{fig:hess} (a Hess diagram plots the relative frequency of stars at different color-magnitude positions). Various phases of stellar evolution can be recognised in the Hess diagram: \\begin{enumerate} \\item Main sequence and possible blue He-core burning stars at $V-I \\sim0$ and $-2\\le V\\le-8$; \\item Red He core burning stars at $1.2\\le V-I \\le2.5$ and $-2.5\\le V\\le-6.5$; \\item RGB/AGB stars at $1\\le V-I \\le3$ and $-0.5\\le V\\le-2.5$. \\end{enumerate} Overplotted in Fig. \\ref{fig:hess} are theoretical isochrones from the Padova group \\citep{marigo08} for five different ages using metallicity $Z=0.008$ (red lines). The O abundance of H{\\sc ii} regions in NGC~4395 was measured by \\citet{roy96} to be $12 + \\log$O/H = $8.33\\pm0.25$, or about 1/3 solar \\citep{grevessesauval98}, suggesting an overall metallicity similar to that of the LMC. This is consistent with our analysis in Sect. 5, which shows that isochrones of LMC metallicity reproduce our data more closely. The white line indicates the $50\\%$ completeness limit (see Sect. 4.3). \\begin{figure}[!t] \\centering \\includegraphics[trim= 16mm 10mm 0mm 5mm,width=\\columnwidth]{hess.ps} \\caption{Hess diagram for the field stars in both fields. The dashed white line represents the $50\\%$ completeness for the first field. Red lines are Padova 2008 theoretical isochrones for the ages $10^7$,$10^{7.5}$,$10^8$,$10^{8.5}$, and $10^9$ yr using LMC metallicity.} \\label{fig:hess} \\end{figure} \\subsection {Star clusters} Our detection criteria is met by 16 463 objects, i.e., 6 connected pixels with 10$\\sigma$ above the local background level. For the clusters, we carried out aperture photometry. An aperture radius of 6 pixels was used for the ACS images. At the distance of NGC 4395, 1 ACS/WFC pixel corresponds to $\\sim1$ pc, hence the chosen source aperture contains about 2 half-light radii for a typical star cluster. Our sky annulus had an inner radius of 8 pixels and a width of 5 pixels. For the WFPC2 images, apertures covering the same area were used (source aperture = 3 pixel radius, sky annulus = 4 pixel inner radius and 2.5 pixels width). Of the 16 463 objects detected with SExtractor/{\\it ishape}, a total of 4 472 candidates have measured four band photometry. ACS magnitudes were converted to vega magnitude system using the same tables as for the field stars. WFPC2 magnitudes were converted to the Vega magnitude system using the zero-points taken from the webpages\\footnote{www.stsci.edu/instruments/wfpc2/Wfpc2\\_ch52\\#1933986} of HST. Charge transfer efficiency (CTE) corrections were applied following the equations from \\citet{dolphin00}\\footnote{Last update May 12,2008:\\\\ \\url{http://purcell.as.arizona.edu/wfpc2\\_calib/}}. To select star cluster candidates, we used 3 criteria: \\begin{itemize} \\item Size:\\\\ We measured the sizes of the objects. Figure \\ref{fig:fwhms} shows the FWHM distributions for SExtractor and {\\it ishape}. The SExtractor histogram peaks at $\\sim 2.2$ pixels, corresponding to the stellar PSF. The {\\it ishape} histogram peaks at 0 pixels, as {\\it ishape} takes the PSF directly into account. Based on the FWHM distributions, we decided to use a criteria of $FWHM_{SExtractor} \\geq 2.7$ pixels and $FWHM_{ishape} \\geq 0.7$ pixels as a first selection of extended sources. These two limits correspond to a physical cluster half-light radius of $\\sim1$ pc or greater at the distance of NGC~4395. Many of the candidate clusters we detect are low-mass and have irregular profiles often dominated by a few stars, so we chose to rely on both SExtractor and Ishape size measurements to achieve a more robust rejection of unresolved sources. \\begin{figure}[!t] \\centering \\includegraphics[width=\\columnwidth]{fwhms.ps} \\caption{Histograms of the FWHM of the objects detected with SExtractor (top panel) and {\\it ishape} (bottom panel). The vertical dotted lines represent the limits used to select the extended objects ($FWHM_{SExtractor} \\geq 2.7$ and $FWHM_{ishape} \\geq 0.7$ pixels)} \\label{fig:fwhms} \\end{figure} \\item Magnitude:\\\\ \\citet{mora09} ran a completeness analysis of artificial star clusters of different FWHMs in five nearby galaxies, including NGC 4395. The results presented in their work establish a $50\\%$ limit between $25\\leq m_B\\leq26$ magnitudes for objects with $FWHM=[0.1$(point sources)$,1.8]$ (see \\citet{mora07,mora09} and Sect. 4.3 for more details). By performing four band photometry, we set a magnitude cut-off at $m_V \\leq 23$, which represents objects brighter than $M_V\\sim-5 $ at the distance of the galaxy. This limit is brighter than that found by \\citet{mora07} by $\\sim2$ magnitudes. The main difference between the selection criteria of \\cite{mora07} and this work is in the size limits of {\\it ishape}, of half of a pixel (Mora et al. used $FWHM_{ishape} \\geq 0.2$), hence we expect our data not to be significantly affected by incompleteness to our magnitude cutoff. However, even at our limit of $M_V=-5$, there is a risk that a few stars may dominate the light originating in a cluster, making it difficult to differentiate a real cluster from a couple of stars that, by chance, could be in the same line of sight. \\item Color:\\\\ Without taking into account any significant reddening, all clusters, including globulars, will have colors bluer than $V-I\\sim1.5$, e.g., \\citet{forbes97}, \\citet{larsen01}. We therefore make a color cut at $V-I=1.5$. \\end{itemize} \\begin{figure}[!t] \\centering \\includegraphics[width=\\columnwidth]{stamps.eps} \\caption{Example stamps of the clusters that satisfy the criteria stated in this work. Each cluster image has a dimension of $100\\times100$ pixels ($\\sim100\\times100$ pc).} \\label{fig:stamps} \\end{figure} We wish to emphasise that there is probably no unique combination of criteria that will lead to the detection of all bona-fide clusters in the image and at the same time produce no false detections. At low masses and young ages in particular, the light profiles may be dominated by individual bright stars. Out of 4 472 candidate objects, a total of 22 objects fulfill the three criteria stated above and will be considerate in the remain of this study to be star clusters. These 22 clusters were visually inspected in our images to determine whether they resemble a cluster. An example of the clusters detected is presented in Fig. \\ref{fig:stamps}. We investigated the large number of rejected objects and found that they were rejected for many reasons: (1) the area covered by the two detectors differs by a factor of $\\sim2$; (2) many of the objects are too faint in the U band; (3) the magnitude cut-off removes many objects, e.g., from a total of 4 475 detected objects, only 1 351 satisfy the magnitude limit, leading to the loss of $\\sim70\\%$ detections; and (4) a large fraction of the objects have FWHMs smaller than our limits. For the extended objects, a two-color diagram, based on the photometry in all of the available 4 passbands, is presented in Fig. \\ref{fig:twocolor}. Clusters in both fields are depicted with their respective photometric errors and corrected for foreground extinction ($A_B=0.074$). Using GALEV models \\citep{andersfritze03}, Padova isochrones \\citep{bertelli94}, a Salpeter's IMF \\citep{salpeter55}, and LMC (dashed line) or solar (dash-dotted line) metallicities, we overplotted the track that a cluster follows from ages between $6.6\\le Log(\\tau)\\le10.2$ yr (each $\\Delta \\log(\\tau)=0.5$ dex age in log units being indicated). We observe considerable scatter in the two-color diagram. For these relatively low-mass clusters, the discreteness of the initial mass function will be important and might contribute significantly to the scatter \\citep{girardi95,cervignoluridiana06,maizapellaniz09}. A more detailed study of stochastic effects will be presented in a forthcoming paper (Silva-Villa et al.\\ 2010, in prep.) This scatter was previously observed by \\citet{mora09} (see their Fig. 6). Considering this scatter, it is clear that ages derived by a comparing observed and model colors should be treated with some caution. \\begin{figure}[!t] \\centering \\includegraphics[width=\\columnwidth]{2color.ps} \\caption{Two-color diagram for the star clusters detected in both fields, $ \\times $ represents the first field and $ \\triangle $ the second field. The lines represent the theoretical GALEV track for a cluster with LMC (dashed line) or solar (dash-dotted line) metallicity, adopting Padova isochrones, a Salpeter IMF, and ages between $6.6\\leq log(\\tau/{\\rm yr}) \\leq10.1$}. \\label{fig:twocolor} \\end{figure} \\subsection {Completeness test} To determine the completeness limits of our field-star photometry, we generated synthetic images with artificial stars. From 20 to 28 magnitudes, in steps of 0.5 mag, five images per step were created and analyzed with exactly the same parameters used in the original photometry. For each combination of color and magnitude, $528$ artificial stars were added using 4 different regions over the image to cover crowded and uncrowded areas. These test images were created using the task {\\it mksynth} in BAOLab \\citep{larsen99} and using the original PSF images created during the photometry procedures (see Sect. 4.1). The separation between two consecutive stars was 100 pixels (without any sub-pixel variations), avoiding possible overlaping among the stars. The resulting test images were added to the science images using the task {\\it imarith} in IRAF. As an example, a subsection of an image, both test and science, is presented in Fig. \\ref{fig:cmpl}, where the fake stars have magnitudes $m_V=21$. \\begin{figure}[!t] \\centering \\includegraphics[width=\\columnwidth]{cmpl2.eps} \\caption{Subsection of the original image (right) and completeness image (left).} \\label{fig:cmpl} \\end{figure} To performe a realistic completeness analysis it is in principle necessary to sample the three-dimensional ($B,V,I$) color space. However, since different colors are tightly correlated with each other, the problem can be reduced to a two-dimensional one. Figure \\ref{fig:11relcmpl} was generated using Padova 2008 isochrones \\citep{marigo08}, assuming solar metallicity, in the color range from -1 to 2, for $B-V$ and $V-I$ and shows that there is a nearly 1:1 relation between these two colors (the red line in Fig. \\ref{fig:11relcmpl} represents a 1:1 relation, but not an accurate tof the data). We created the images for the completeness test described above using this approximation between the three bands, i.e., if a $B$ image has stars with $m_B=21$ and $B-V=1$, then the $V$ and $I$ images will have stars of $m_V=20$ and $m_I=19$, respectively, allowing us to perform a study of magnitudes and color variations over the images. Based on these tests, we found an average of $50\\%$ completeness limits for the whole color range at $m_B = 26.69$, $m_V = 26.55$, and $m_I = 26.42$ for the first field and $m_B = 26.71$, $m_V = 26.49$, and $m_I = 26.39$ for the second field. Figure \\ref{fig:cmpl-pos1} shows the completeness diagram obtained from this analysis for the first field (for the second field, the figure is similar) and each magnitude. The color-dependent 50\\% completeness limit is shown as a white dashed line in the Hess diagram (see Fig. \\ref{fig:hess}). Since the completeness functions are very similar for the two fields, we can combine the photometry for both fields and use just one set of completeness tests in the following analysis. \\begin{figure}[!h] \\centering \\includegraphics[width=\\columnwidth]{11relcmpl.ps} \\caption{Two color diagram for theoretical values of B, V and I bands from Padova 2008 isochrones adopting LMC metallicity for ages between $10^{6.6}$ to $10^{10}$ yr. The red line represents a $1:1$ relation between the colors B-V and V-I.} \\label{fig:11relcmpl} \\end{figure} We did not performe a completeness test for star clusters but refer to the tests performed by \\citet{mora07,mora09}, who used the same data and very similar cluster detection procedures. These authors created artificial star clusters using different FWHMs from 0.1pixels (stars) to 1.8 pixels and a range of magnitudes from 16 to 26 for three square grids in different positions over the image, trying to cover crowded and non-crowded areas. To create the fake clusters, they assumed a \\citet{king62} profile with $r_{tidal}/r_{core}=30$. Using the {\\it mkcmppsf} task on BAOLab, fake extended objects were added to an empty image and then added to the science image to performe measurements. The selection criteria in \\citet{mora07,mora09} is a $FWHM\\ge(2.7,0.2)$ pixels for SExtractor and {\\it ishape}, respectively. For high background levels, Mora at al. found a shallower detection limit; nevertheless, all the limits correspond a 50\\% completeness limit between $m_B\\approx25$ and $m_B\\approx26$, 2-3 magnitudes fainter than the cut at V=23 that we apply for the selection of cluster candidates. \\begin{figure}[!h] \\centering \\includegraphics[width=\\columnwidth]{cmplp1.ps} \\caption{Completeness diagrams for the bands B,V, and I studied over the first field. Vertical lines represent the $50\\%$ completeness.} \\label{fig:cmpl-pos1} \\end{figure} \\subsection{Aperture corrections} The aperture corrections were determined separately for field stars and star clusters. The aperture corrections for the field stars were derived following standard procedures, while for extended objects we adopted the relations found by \\cite{mora09}. \\begin{itemize} \\item Field stars:\\\\ By ``aperture corrections'', we here mean corrections from the PSF-fitted instrumental magnitudes to aperture photometry for nominal radii of $0\\farcs5$. These were measured using a set of isolated visually selected stars across the images. From $0\\farcs5$ to infinity, we applied the \\citet{sirianni05} values. The corrections obtained are of the order $\\sim 0.1$ mag (see table \\ref{tab:aperturecorrections}). \\item Star clusters:\\\\ \\citet{mora09} estimated a relation between the aperture corrections and the sizes (FWHM) of star clusters using the same data set used in this paper. Photometric parameters in both \\citet[][and ours]{mora09} work are the same, allowing us to assume the relations found in their work and apply these aperture corrections (size-dependent) to our data, following eq. 1 in \\citet{mora09}. This set of equations is also band-dependent, although we used the sizes of the objects measured on an average image. The aperture corrections by \\citet{mora09} correspond to a nominal aperture of $1\\farcs45$. From this nominal aperture to infinity, we adopted the values presented by \\citet{sirianni05}, although these corrections are $\\sim0.03$ magnitudes for the bands B, V, and I (within $1\\farcs5$ about 97\\% of the total energy is encircled). \\end{itemize} \\begin{table}[!t] \\centering \\caption{For point source. B, V, and I aperture corrections to a nominal aperture of $0.\"5$, estimated in this study.} \\begin{tabular} {c c c c} \\hline \\hline & $B_{F435W}$ & $V_{F555W}$ & $I_{F814W}$ \\\\ & [mag] & [mag] & [mag] \\\\ \\hline \\hline {\\it Field 1} & 0.07 & 0.03 & 0.08 \\\\ {\\it Field 2} & 0.05 & 0.08 & 0.06 \\\\ \\hline \\hline \\end{tabular} \\label{tab:aperturecorrections} \\end{table} ", "conclusions": "" }, "1003/1003.0834_arXiv.txt": { "abstract": "{ We present simulations of the 21-cm signal during the epoch of reionization. We focus on properly modeling the absorption regime in the presence of inhomogeneous Wouthuysen-Field effect and X-ray heating. We ran radiative transfer simulations for three bands in the source spectrum (Lyman, UV, and X-ray) to fully account for these processes. We find that the brightness temperature fluctuation of the 21 cm signal has an amplitude greater than 100 mK during the early reionization, up to 10 times greater than the typical amplitude of a few 10 mK obtained during the later emission phase. More importantly, we find that even a rather high contribution from QSO-like sources only damps the absorption regime without erasing it. Heating the IGM with X-ray takes time. Our results show that observations of the early reionization will probably benefit from a higher signal-to-noise value than during later stages. After analyzing the statistical properties of the signal (power spectrum and PDF) we find three diagnostics to constrain the level of X-ray, hence the nature of the first sources. } ", "introduction": "The epoch of reionization (EoR) started with the formation of the first sources of light around $z=15 - 30$. As shown by the Gunn-Peterson effect \\citep{Gunn65} in the spectra of high-Fredshift quasars (QSO) (e.g., \\citealt{Fan06}), the universe was fully reionized by $z \\sim 6$. WMAP 5-year results show that the optical depth for the Thomson scattering of CMB photons traveling through the reionizing universe is $\\tau= 0.084 \\pm 0.016$ \\citep{Koma08}. Together with the Gunn-Peterson data, this strongly favors an extended reionization period between $z>11$ and $z=6$. While other observations, such as the Lyman-$\\alpha$ emitter luminosity function \\citep{Ouch09}, may produce other constraints on the history of reionization in the next few years, the most promising is the observation of the $21$-cm line in the neutral IGM using large radio-interferometers (LOFAR\\footnote{LOw Frequency ARray, \\url{http://www.lofar.org}}, MWA\\footnote{Murchison Widefield Array, \\url{http://mwatelescope.org/}}, GMRT\\footnote{Giant Metrewave Radio Telescope, \\url{http://www.gmrt.ncra.tifr.res.in}}, 21-CMA\\footnote{21 Centimeter Array, \\url{http://21cma.bao.ac.cn/}}, SKA\\footnote{Square Kilometre Array, \\url{http://www.skatelescope.org/}}). The signal will be observed in emission or in absorption against the CMB continuum. Both theoretical modeling \\citep{Mada97,Furl06a} and simulations (e.g., \\citealt{Ciar03c,Gned04,Mell06b,Lidz08,Ichi09,Thom09}) show that the brightness temperature fluctuations of the 21 cm signal have an amplitude of a few $10$ mK in emission, on scales from tens of arcmin down to sub-arcmin. With this amplitude, and ignoring the issue of foreground cleaning residuals, statistical quantities such as the three-dimensional power spectrum should be measurable with LOFAR or MWA with a few 100 hours integration \\citep{Mora04,Furl06a,Lidz08}. In absorption however, the amplitude of the fluctuations may exceed $100$ mK \\citep{Gned04,Sant08,Baek09}, the exact level depending on the relative contribution of the X-ray and UV sources to the process of cosmic reionization. The signal will be seen in absorption during the initial phase of reionization, probably at $z> 10$, when the accumulated amount of emitted X-ray is not yet sufficient to raise the IGM temperature above the CMB temperature. The duration and intensity of this absorption phase, regulated by the spectral energy distribution (SED) of the sources, are crucial. SKA precursors able to probe the relevant frequency range, 70 - 140 MHz, may benefit from a much higher signal-to-noise than during later periods in the EoR. However, if the absorption phase is confined at redshifts above $15$, RFI and the ionosphere will become an problem. Quite clearly, the different types of sources of reionization and different formation histories produce very different properties for the 21-cm signal. It is therefore important for future observations to explore the range of astrophysically plausible scenarios using numerical simulations. To properly model the signal, it is necessary to use $ > 100 {h}^{-1}\\mathrm{Mpc}$ box sizes \\citep{Bark04}. Together with a large box size, it is desirable to resolve halos with masses down to $10^8 M_{\\odot}$ as these contain sources (able to cool below their virial temperature by atomic processes), or even minihalos with masses down to $10^4 M_{\\odot}$ became they act as an efficient photon sink because of their high recombination rate \\citep{Ilie05}. As this work focuses on improving the physical modeling, we restrict ourselves to resolving halos with a mass $10^{10} M_{\\odot}$ or higher. Indeed, simulating the absorption phase correctly, as we do in this work, requires a more extensive and more costly implementation of radiative transfer. We are exploring the direct implication of this improved physical modeling, and will turn to better mass resolution in the near future. There are three bands in the sources SED that influence the level of the 21 cm signal: the Lyman band, the ionizing UV band, and the soft X-ray band. Lyman band photons are necessary to decouple the spin temperature of hydrogen from the CMB temperature through the Wouthuysen-Field effect \\citep{Wout52,Fiel58}, and make the EoR signal visible. UV band photons are of course responsible for the ionization of the IGM, and soft X-rays are able to preheat the neutral gas ahead of the ionizing front, deciding whether the decoupled spin temperature is less (weak preheating) or greater (strong preheating) than the CMB temperature. While a proper modeling should perform the full 3D radiative transfer in all 3 bands, a simpler modeling has often been used in previous works. Indeed, for the usual source SEDs and source formation histories, once the average ionization fraction of the universe is greater than $\\sim 10 \\%$, the background flux of Lyman-$\\alpha$ photons is so high that the hydrogen spin temperature is fully coupled to the kinetic temperature by the Wouthuysen-Field effect \\citep{Baek09}. Thereafter, computing the Lyman band radiative transfer is unnecessary. In the same spirit, it has usually been assumed that the preheating of the IGM by soft X-ray was strong enough to raise the kinetic temperature much higher than the CMB temperature everywhere early in the EoR. However, both assumptions fail during the early reionization: the absorption phase. Even in the later part of reionization the second assumption may fail, depending on the nature of the sources. We will quantify this possibility in this paper. Computing the full radiative transfer in all three bands is necessary to study the absorption regime. Indeed, fluctuations in the local Lyman-$\\alpha$ flux induce fluctuations in the spin temperature (while the Wouthuysen-Field effect is not yet saturated), which, in turn, modify the power spectrum of the $21$ cm signal \\citep{Bark05, Seme07, Chuz07, Naoz08, Baek09}. The same is true for the fluctuations in the local flux of X-ray photons \\citep{Prit07,Sant08}. Let us emphasize however that, in modeling Lyman-$\\alpha$ and X-ray fluctuations, \\citet{Bark05}, \\citet{Naoz08}, \\citet{Prit07} and \\citet{Sant08} all use the semi-analytical approximation that the IGM has a uniform density of absorbers and ignore wing effects in the radiative transfer of Lyman-$\\alpha$ photons. \\citet{Seme07} and \\citet{Chuz07} have shown that these wing effects do exist. Moreover, once reionization is under way, ionized bubbles create sharp fluctuations in the number density of absorbers (not to mention simple matter density fluctuations). In this work, for the first time, we present results based on simulations with full radiative transfer for both Lyman-$\\alpha$ and X-ray photons. What are the possible candidates as sources of reionization? Usually, two categories are considered: ionizing UV sources (Pop II and III stars), and X-ray sources (quasars). When we study 21 cm absorption, however, we must distinguish between Pop II and Pop III stars beyond the large difference in luminosity per unit mass of formed star. Indeed Pop II stars have a three times larger Lyman band to ionizing UV band luminosity ratio than Pop III stars. It means that the 21 cm signal will reach its full power (near saturated Wouthuysen-Field effect) at a lower average ionization fraction for Pop II stars than for Pop III stars. The relevant question is: how long do Pop III stars dominate the source population before Pop II stars take over? The answer to this question, related to the whole process of star formation, feedback and metal enrichment of the IGM, is a difficult one. At this stage, state of the art numerical simulations of the EoR use simple prescriptions in the best case (e.g. \\citealt{Ilie07a}), or simply ignore this issue. The other category of sources are X-ray sources. They may be mini-quasars, X-ray binaries, supernovae \\citep{Oh01, Glov03}, or even more exotic candidates such as dark stars \\citep{schl09}. The exact level of emission from these sources is a matter of speculation. The generally accepted view is that stars dominate over X-ray sources and are sufficient to drive reionization \\citep{Shap87, Giro96, Mada99, Ciar03a}. Recently, \\citet{Volo09} supported the opposite view. While, in their models, X-ray sources are marginally able to complete reionization by $z \\sim 6$, they find a very low contribution from stars. Indeed they rely on \\citet{Gned08} who find, using numerical simulations, a negligible escape fraction for ionizing radiations from galaxies with total mass less than a few $10^{10} M_{\\odot}$, who should actually contribute to $90\\%$ of the ionizing photon production during the EoR \\citep{Chou07}. While the physical modeling in their innovative simulations is quite detailed, this surprizing behavior of the escape fraction definitely needs to be checked at higher resolution and with different codes. For the time being the best simulations can only explore a plausible range of X-ray contributions, and quantify the impact on observables. When the observations become available we would like to be able, using simulation results, to derive tight constraints on the relative level of emission from ionizing UV and X-ray sources. This work, exploring the 21 cm signal for a few different levels of X-ray emission, is a first step toward this goal. The paper is organized as follows. We present the numerical methods in \\S2 and describe our source models in \\S3. In \\S4, we show the results and analyze the differences between the models. We discuss our findings and conclude in \\S5. ", "conclusions": "We modeled the 21-cm signal during the EoR using numerical simulations putting the emphasis on how various types of sources can affect the signal. The numerical methods used in this work are similar to \\citet{Baek09}. The N-body and hydrodynamical simulations have been run with GADGET2 and post-processed with UV continuum radiative transfer further processed with Ly-$\\alpha$ transfer using LICORICE, allowing us to model the signal in absorption. The main difference from the previous work is a more elaborated source model, including X-ray radiative transfer and He chemistry. We have run 7 simulations to investigate the effect of different IMFs, helium, different spectral indexes and the different luminosities of X-rays sources. The reference simulation in this work, S1, using only hydrogen and stellar type sources, reached the end of reionization at $z\\approx 6.5$ and showed a strong absorption signal until the end of reionization. Our top heavy IMF (model S2) produces $\\sim 2.6$ times more ionizing photons than the Salpeter IMF. S2 reached the end of reionization earlier than the others by $\\Delta z \\approx 1$. In addition the different SED changes the ratio of Ly-$\\alpha$ and to ionizing UV photon numbers, and it slows down the saturation of the Ly-$\\alpha$ coupling and the heating by Lyman-$\\alpha$ in the top heavy IMF case. This modifies the statistical properties of 21-cm signal. The simulation with helium, S3 also has a slightly earlier reionization than the others since the number of emitted photons per baryon is higher. Except for the slightly lower kinetic temperature in the bulk of ionized regions due to the higher ionization potential than for hydrogen, the properties of the 21-cm signal from S3 is similar to S1. We chose QSO type sources with a power-law spectrum as X-ray sources in model S4 to S7. The spectral index $\\alpha$ has large observational uncertainty, so we used two different spectral indexes. S4 and S5 have 0.1\\% of the total luminosity in the X-ray band. S5 uses $\\alpha = 0.6$, while other simulations with X-rays use $\\alpha=1.6$. S5 showed very little difference on the gas temperature with respect to S4. S4, S6 and S7 have different luminosities in the X-ray band, keeping the same values for the other simulation parameters. Using a stronger X-ray luminosity indeed increased the gas temperature in the neutral hydrogen. Accordingly the 21-cm signal and its power spectra are modified. We found an increase of a few kelvin for the neutral gas temperature in our fiducial model, S4, in which X-rays account for 0.1\\% of the total emitted energy. The 21-cm signal in S4 was similar to S1, showing the maximum intensity in absorption, $\\sim 200$ mK, at $z \\approx 9$. Stronger X-ray levels increase the gas temperature and reduce the intensity. We found that in S6 and S7, which uses 1\\% and 10\\% of the total luminosity for X-rays, the absolute maximum intensity in absorption decreases to $\\sim 130$ mK and $\\sim 80$ mK. The 21-cm power spectrum of our work is greater by two or three orders of magnitude than in works focusing on the emission regime \\citep{Mell06b,Zahn07,Lidz07b,Mcqu06}. However, the results are in broad agreement with the work of \\citet{Sant08}, who modeled absorption using semi-analytical methods for X-ray and Lyman-$\\alpha$ transfer. We noticed that the 21-cm fluctuation is dominated by Ly-$\\alpha$ fluctuations during the early phase, X-rays later (or the gas temperature), and the ionization fraction at the end. This is visible on the evolution of the 21-cm power spectrum with redshift. The 21-cm PDF of our work was different from other work, since we do not assume that the spin temperature $T_s \\gg T_\\mathrm{CMB}$. \\textit{The first most important conclusion} from our work is that even including a higher than generally expected level of X-ray, the absorption phase of the 21-cm survives. Its intensity and duration are reduced, but the signal is still stronger than in the emission regime. Heating the IGM with X-rays takes time! \\textit{The second important result} is that we found three diagnostics which could be used in the analysis of future observations to constrain the nature of the sources of reionization. $(i)$ The first and maybe the most robust is the evolution with redshift of large scale modes ( $k \\sim 0.1$ h/Mpc) of the powerspectrum. If reionization is overwhelmingly powered by stars, this evolution should have one local minimum (two local maxima) . However, if the energy contribution of QSO is greater than $\\sim 1\\%$, a second local minimum (third maximum) appears. The higher the X-ray level, the broader the third peak. $(ii)$ The second simple diagnostic is the bimodal aspect of the PDF which disappears when the X-ray level rises above $1\\%$ of the total ionizing luminosity. $(iii)$ Last is the redshift evolution of the skewness of the 21-cm signal PDF. While all other models show a single local maximum at a few percent reionization, a very high level of X-rays ($> 10\\%$ of the total ionizing luminosity) produces a second local maximum appear around $50\\%$ reionization. Modeling the sources in the simulation is complex. It involves taking the formation history, IMF, SED, life time, and more into account. Although detailed models are desirable for the credibility of the results, we believe that the effect in the 21-cm signal can be bundled in 3 quantities. The first is the efficiency: how many photons are produced by atoms locked into a star. This parameter must be calibrated to fit observational constraint: end of reionization between redshift $6$ and $7$, and Thomson scattering optical depth in agreement with CMB experiment. The two other quantities which contain most of the information are two box-averaged ratios: the energy emitted in the Lyman band to the energy emitted in the ionizing band ratio and the same ionizing UV to X-ray ratio. In this work we explored values of 0.32 (model S2) and 0.75 (all other models) for the former and $0.001$, $0.01$ and $0.1$ for the latter. Once additional physics is included in the simulation and using a higher resolution to account for all the sources, it will be interesting to explore the value of these quantities systematically. We mentioned in the introduction that the minimum boxsize for reliable predictions of the signal is $100\\,h^{-1}$Mpc. It is important to realize that this value (confirmed by emission regime simulations, e. g., \\citealt{Ilie06b}) is estimated based on the clustering properties of the sources and applies to the topology of the ionization field. It may be underestimated when we study the early absorption regime, when only the highest density peaks contain sources. Their distribution is the most sensitive to possible non-gaussianities in the matter power spectrum. Moreover, they are distant from each other and, consequently, produce large scale fluctuations in the local flux of Lyman-$\\alpha$ and X-ray photons. We intend to extend our investigation to larger box sizes in a future work. A few final words on additional physics not included in our model. Shock heating from the cosmological structure formation is ignored, but it could have the potential to affect the 21-cm signal by increasing the gas temperature above the CMB temperature. However, it is not sure whether shocks are strong enough in the filaments of the neutral regions to affect the 21-cm signal. Mini-halos ($\\sim 10^4-10^8\\,\\,M_{\\odot}$) form very early during the EoR and are dense and warm enough from shock heating during virialization to emit the 21 cm signal, but \\citet{Furl06c} find that the contribution of mini-halos will not dominate, because of the limited resolution of the instrumentation. However, shock heating is worth investigating with coupled radiative hydrodynamic simulations with higher mass resolution. Also worth investigating is the effect of including higher Lyman lines in the radiative transfer." }, "1003/1003.0005_arXiv.txt": { "abstract": "We present redshifts and optical richness properties of 21 galaxy clusters uniformly selected by their Sunyaev-Zel'dovich signature. These clusters, plus an additional, unconfirmed candidate, were detected in a $178\\,\\deg^2$ area surveyed by the South Pole Telescope in 2008. Using $griz$ imaging from the Blanco Cosmology Survey and from pointed Magellan telescope observations, as well as spectroscopy using Magellan facilities, we confirm the existence of clustered red-sequence galaxies, report red-sequence photometric redshifts, present spectroscopic redshifts for a subsample, and derive \\rtwohundred\\ radii and \\mtwohundred\\ masses from optical richness. The clusters span redshifts from 0.15 to greater than 1, with a median redshift of 0.74; three clusters are estimated to be at $z > 1$. Redshifts inferred from mean red-sequence colors exhibit 2\\% RMS scatter in $\\sigma_z/(1+z)$ with respect to the spectroscopic subsample for $z < 1$. We show that \\mtwohundred\\ cluster masses derived from optical richness correlate with masses derived from South Pole Telescope data and agree with previously derived scaling relations to within the uncertainties. Optical and infrared imaging is an efficient means of cluster identification and redshift estimation in large Sunyaev-Zel'dovich surveys, and exploiting the same data for richness measurements, as we have done, will be useful for constraining cluster masses and radii for large samples in cosmological analysis. ", "introduction": "Galaxy clusters are laboratories for both astrophysics and cosmology \\citep{evrard04}. Clusters represent the most massive dark matter halos, and their number density as a function of cosmic time is highly sensitive to dark energy \\citep{wang98,haiman01,holder01b,battye03,molnar04,wang04,lima07}. The mass of these systems is dominated by dark matter, but the primary means of observing clusters---especially large samples of them---are the luminous baryons of the hot intracluster gas and the galaxies themselves. The formation of the halos is well understood, while the precise behavior of the baryons is not as well modeled \\citep[see][for a review]{voit04}. This gap must be closed so that data from large cluster surveys can place precise constraints on cosmological parameters over a wide range of redshifts. Multi-wavelength observations of a cleanly selected, redshift-independent sample of galaxy clusters are a potentially powerful method of achieving this. Searches for galaxy clusters using the Sunyaev-Zel'dovich (SZ) effect \\citep[][]{sunyaev72} promise to provide such a clean, redshift-independent sample. The SZ effect is scattering of cosmic microwave background photons to higher energy by the hot electrons in galaxy clusters \\citep{birkinshaw99}. The SZ surface brightness is independent of redshift but is closely related to cluster mass and so it is expected to be an excellent method for creating approximately mass-limited samples extending over a wide redshift range \\citep[][]{carlstrom02}. The constraints on cosmological parameters from such samples are complementary to geometrical tests using type Ia supernovae and baryon acoustic oscillations \\citep[e.g.,][]{vikhlinin09}. Two SZ surveys, the Atacama Cosmology Telescope \\citep[ACT;][]{fowler07} and the South Pole Telescope \\citep[SPT;][]{carlstrom09} projects, are well positioned to provide large surveys which can be used for growth of structure studies. \\citet[hereafter S09]{staniszewski08} presented the first discovery of previous unknown galaxy clusters using their SZ signature. Cluster redshifts are needed in addition to SZ data to provide the strongest constraints on dark energy. Coordinated optical follow-up observations can provide the needed redshift measurements. The Blanco Cosmology Survey \\citep[BCS;][and \\url{http://cosmology.illinois.edu/BCS/}]{ngeow06}, an NOAO survey program (2005-2008), provided multiband optical observations for the initial follow-up of portions of the first SPT survey fields. These data were used to identify optical counterparts to the \\citetalias{staniszewski08} sample, search for giant arcs, explore possible cluster superpositions, and derive photometric redshifts. Cluster mass can be estimated using several methods: the SZ and X-ray luminosity, which are sensitive to intracluster electrons; the number, luminosity, and velocity dispersion of cluster galaxies; and from gravitational lensing, which is the most direct probe of total cluster mass. \\citet{menanteau09} characterized the galaxy counts and luminosity of the \\citetalias{staniszewski08} cluster sample, and \\citet{mcinnes09} subsequently explored their weak gravitational lensing signals. Using data acquired by the SPT in 2008, \\citet[][hereafter V10]{vanderlinde10} present an additional 17 SZ-detected clusters. Here we describe coordinated optical imaging of the catalog of 21 uniformly selected SZ detections, and new spectroscopic results on 8 of the clusters. Counterparts to a subset have been found in the catalogs of \\citet[][hereafter A89]{abell89} and \\citet[][hereafter SCS-II]{menanteau10}. Seven clusters fell within the BCS footprint. For the remaining 14 clusters, and also for a subset of the BCS sample, we conducted pointed imaging observations and, for 8 clusters, spectroscopic observations, with the Magellan telescopes. The photometry was used to search for overdensities of red-sequence galaxies near the SZ locations, and if present, estimate their redshifts and also characterize their mass via optical red-sequence galaxy counts, or {\\it richness}. We describe in \\S \\ref{sec:data} the observations and data reduction. Section \\ref{sec:analysis} outlines the redshift and richness analysis we used, and \\S \\ref{sec:results} describes the results on redshift (\\S \\ref{sec:redshifts}) and richness (\\S \\ref{sec:richness}). In Section \\ref{sec:discussion} we discuss the results, and conclude with \\S \\ref{sec:conclusion}. Throughout this paper we assume a flat concordance $\\Lambda\\mathrm{CDM}$ universe, with $(\\Omega_{\\Lambda},\\Omega_{\\mathrm{M}},h) = (0.736,0.264,0.71)$ \\citep{dunkley09}. All magnitudes are in the Sloan Digital Sky Survey (SDSS) $griz$ AB system. ", "conclusions": "\\label{sec:conclusion} We have observed clusters from the 2008 South Pole Telescope SZ survey at optical and near-infrared wavelengths. We estimate redshifts and richness with red-sequence techniques, and we obtain spectroscopic redshifts for a subsample of the clusters. Our red-sequence-derived redshifts exhibit $2\\%$ RMS scatter in $\\sigma_z/(1+z)$ in the subsample with spectroscopic overlap, over the redshift range $0.15 \\Sigma_m m, \\label{eq:IM93} \\end{equation} where $\\Sigma_M$ is the surface density of big bodies of mass $M$ and $\\Sigma_m$ that of the small bodies. \\Eq{IM93} can be transformed into a radius, $R_\\mathrm{rg/oli}$, indicating the turnover from runaway growth into oligarchy (see below, \\eqp{Roli}). Many works have adopted \\eq{IM93} as the start of their oligarchic calculations \\citep[\\eg][]{ThommesEtal2003,IdaLin2004,Chambers2006,Chambers2008,FortierEtal2007,BruniniBenvenuto2008,MiguelBrunini2008,MordasiniEtal2009}. In this letter, we will refine the criterion of \\citet{IdaMakino1993} and present a new expression for $R_\\mathrm{rg/oli}$ (\\eqp{R-rg}). In runaway growth the column density spectrum evolves into a power law, $N(m) = (1/m) d\\Sigma/dm \\propto m^{-p}$, where $p\\approx-2.5$ \\citep[][\\fg{fig1}a]{WetherillStewart1993,KokuboIda1996,KokuboIda2000,BarnesEtal2009}. However, for a $p>-3$ index, the stirring power lies at the high-mass end of the population; that is, during runaway growth the stirring power is already moving away from the initial mass ($m_0$). Despite the stirring, the system continues in its runaway (fast) growth mode. The point is that \\eq{IM93} implicitly assumes that stirring occurs fast and outpaces the accretion, which is true for oligarchy but not for runaway growth. In other words, \\eq{IM93} is a necessary condition for oligarchy but not a sufficient one. Instead, we will argue that the condition for the start of oligarchy is met when the stirring timescale in the two component approximation ($T_\\mathrm{vs}$, see below) drops below the accretion timescale that characterizes the runaway regime, $T_\\mathrm{rg}$. In \\se{def} we first introduce key definitions and obtain the stirring timescale $T_\\mathrm{vs}$ and the accretion timescale $T_\\mathrm{ac}$ for a two component system. In \\se{Trg} we present the results of our runaway growth simulations in terms of $T_\\mathrm{rg}$, the accretion timescale during runaway growth, which follows from our simulations. \\Se{rgot} then presents the new transition radius $R_\\mathrm{tr}$ by equating the timescale expressions. We discusses a few implications and summarize in \\se{summ}. ", "conclusions": "\\label{sec:summ} According to our findings, the transition between the runaway growth and oligarchic growth phases is characterized by the following properties: \\begin{itemize} \\item A power-law size distribution of mass index $p\\approx-2.5$, extending from the initial size $R_0$ to the transition size $R_\\mathrm{tr}$ as given by \\eq{R-rg}; \\item The random velocity $v$ of the planetesimal bodies (of size $\\sim$$R_0$), via \\eq{uvhtr}; \\item The timescale, via \\eq{Trg}. Here, $T_\\mathrm{rg}$ must be multiplied by a term $\\sim$$\\log (R_\\mathrm{tr}/R_0)$ to account for the several e-foldings of enjoyed growth. Then we obtain $t=T_\\mathrm{tr} \\sim \\rho R_0/\\Sigma_0 \\Omega$ for the time until the transition. It is remarkable that this short timescale depends on the initial conditions only, a result that is unique to the runaway growth phase. \\end{itemize} We assess the implications of these findings for the broader context of planet formation. First, we do not expect the final timescales of core formation to be much influenced by the new transition radius, since these are set by the much slower oligarchy stage that supersedes runaway growth. Indeed, semi-analytical studies of oligarchic growth required additional mechanisms like gas damping, fragmentation, or migration to produce embryos on reasonable timescales \\citep[\\eg][]{BruniniBenvenuto2008,Chambers2008}. On the other hand, gap formation \\citep{Rafikov2003ii} will increase formation timescales. Recently, \\citet{LevisonEtal2010} investigate core-formation scenarios using $N$-body techniques. The aim of that study was to understand how efficient planetesimal accretion proceeds during the phase where an Earth-size planetary embryo has to grow to a mass of $\\sim$$10M_\\oplus$ to be able to accrete the nebula gas and become a gas giant. They recognize the importance of processes like planetesimal scattering, planetesimal orbital decay due to gas drag (both processes quench the growth), and planetesimal-driven embryo migration (which is conducive to growth). The outcome of these simulations, furthermore, is found to depend on the initial setup of the simulation, \\ie\\ the size distribution of the planetesimals. Given the significance that is attached to the wholesale redistribution of matter, it would also be of interest to assess the importance of these effects for the early oligarchy stage, \\eg\\ to perform $N-$body simulations with embryos radii starting at the transition mass $R_\\mathrm{tr}$. The second implication of our study concerns the Kuiper belt. Our results strongly suggest that the Kuiper Belt is primarily the product of the runaway growth phase. First, assuming that the initial surface density is approximately MMSN or larger ($\\Sigma \\sim 0.1\\ \\mathrm{g\\ cm^{-2}}$), the biggest $\\sim$10$^3$ km bodies (plutinos) can be produced by runaway growth. Second, the observed mass distribution $N(m)$ for the largest Kuiper belt objects (KBOs) obeys a power-law with $p\\approx-2.5$. Recent studies find a power-law size index $q$ (as in $N(R) \\propto R^{-q}$) of $4.8\\pm0.3$ ($p=-2.3\\pm0.1$; \\citealt{FraserKavelaars2009}) and $4.5^{+1.0}_{-0.5}$ ($p=-2.2^{+0.2}_{-0.3}$; \\citealt{FuentesHolman2008}). (The size distribution of the Kuiper belt's smallest bodies is collisionally dominated, though, and $q$ is much lower.) The large $q$, together with the fact that the KBO size distribution is continuous rather than bimodal as in the end stage of \\fg{fig1}a, argue that it has evolved only very little since runaway growth and has not been significantly shaped by oligarchic growth. In order to produce the KBOs in a sufficiently short time span, Kuiper Belt formation scenarios \\citep[\\eg][]{KenyonLuu1998,ChiangEtal2007} assume that the initial belt contained much more mass than the $\\sim0.01\\ M_\\oplus$ ($\\Sigma \\sim 0.001\\ \\mathrm{g\\ cm^{-2}}$) that is present today \\citep{BernsteinEtal2004}. Neptune formation and/or another dynamical shakeup event (as in the Nice model) subsequently depleted 99\\% of its mass \\citep{FordChiang2007,LevisonEtal2008}. Using our results, we can verify these findings. Assuming $a=35$ AU, we find \\begin{eqnarray} \\label{eq:RT-kuipera} R_\\mathrm{tr} &\\sim& 10^3\\ \\mathrm{km} \\left( \\frac{R_0}{\\mathrm{10\\ km}} \\right)^{3/7} \\left( \\frac{\\Sigma}{\\mathrm{0.1\\ g\\ cm^{-2}}} \\right)^{2/7}; \\\\ \\label{eq:RT-kuiperb} T_\\mathrm{tr} &\\sim& \\frac{R_0 \\rho}{\\Omega \\Sigma} \\sim 10^8\\ \\mathrm{yr}\\left( \\frac{R_0}{\\mathrm{10\\ km}} \\right) \\left( \\frac{\\Sigma}{\\mathrm{0.1\\ g\\ cm^{-2}}} \\right)^{-1}, \\end{eqnarray} which readily shows the need for an enhanced surface density over the current one: for $\\Sigma \\sim 0.001\\ \\mathrm{g\\ cm^{-2}}$ either $R_\\mathrm{tr}$ becomes too low or $T_\\mathrm{tr}$ too long. Thus, we conclude that the KBO size distribution as seen today is consistent with a scenario of being a leftover product of the initial runaway growth phase and has since been depleted. These findings are in line with the simulations of \\citet{KenyonBromley2008,KenyonBromley2009arXiv}, where growth also stalls after $\\sim$$10^3$ km bodies have been formed. Further growth is impeded since oligarchic accretion timescales become too long, even at enhanced surface densities." }, "1003/1003.4729.txt": { "abstract": "{ Composite dark matter is a natural setting for implementing inelastic dark matter --- the $\\OO(100\\keV)$ mass splitting arises from spin-spin interactions of constituent fermions. In models where the constituents are charged under an axial $U(1)$ gauge symmetry that also couples to the Standard Model quarks, dark matter scatters inelastically off Standard Model nuclei and can explain the DAMA/LIBRA annual modulation signal. This article describes the early Universe cosmology of a minimal implementation of a composite inelastic dark matter model where the dark matter is a meson composed of a light and a heavy quark. The synthesis of the constituent quarks into dark hadrons results in several qualitatively different configurations of the resulting dark matter composition depending on the relative mass scales in the system. } \\begin{document} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ", "introduction": "A wide variety of hidden sectors weakly coupled to the Standard Model have been considered in the literature, and the possibility that dark matter is charged under hidden sector gauge forces has received considerable recent attention \\cite{Nelson:2008hj,ArkaniHamed:2008qn,Pospelov:2008jd, HiddenValley, Cui:2009xq, Batell:2009vb, Finkbeiner:2009, Chang:2008, Ross:2009, Kaplan:2009ag, Morrissey:2009ur, Kaplan:2009de}. In CiDM, dark matter is a bound state with constituents charged under a hidden sector gauge group of the form $SU(N_c)\\times U(1)_\\td$, where the $U(1)_\\td$ gauge boson kinetically mixes with Standard Model hypercharge gauge boson, and the $SU(N_c)$ group condenses at the GeV scale. Theories in which dark matter is charged under a new GeV-scale gauge group, as in CiDM, predict a variety of multi-lepton signals in $B$-factories and $\\phi$-factories \\cite{Essig:2009nc, Batell:2009yf, :2009pw, :2009qd, Bossi:2009uw}, low-energy upcoming fixed-target experiments \\cite{Bjorken:2009mm, Batell:2009di, Essig:2010xa, Reece:2009un}, and distinctive astrophysical signatures \\cite{Schuster:2009fc, Schuster:2009au, Meade:2009mu}. The Lagrangian of the axial CiDM model in \\cite{Alves:2009nf} is given by %% \\begin{equation} \\LL = \\LL_{\\text{SM}} + \\LL_{\\Psi}+\\LL_{\\text{Gauge}}+ \\LL_{\\text{Higgs}} \\end{equation} where \\begin{eqnarray} \\label{eq:mix} \\nonumber &\\LL_{\\text{Gauge}}=- \\frac{1}{2} \\Tr G_{\\td}^2 - \\frac{1}{4} F_{A_\\td}^2 + \\half\\epsilon F_{A_\\td}^{\\mu\\nu}B_{\\mu\\nu} \\\\ \\nonumber &\\LL_{\\Psi} = \\bar{L} i \\slD L + \\bar{L}^c i \\slD L^c +\\bar{H} i \\slD H + \\bar{H}^c i \\slD H^c\\\\ &\\LL_{\\text{Higgs}} = |D_\\mu \\phi|^2 - \\lambda( |\\phi|^2 - v_\\phi^2)^2 + (y_L \\phi L L^c + y_H \\phi^* H H^c+\\hc\\!), \\end{eqnarray} %% and the gauge charges are %% \\begin{eqnarray} \\label{table charges} \\begin{array}{|c||c|c|c|} \\hline & SU(N_c)& U(1)_A&U(1)_{H-L}\\\\ \\hline\\hline L& \\tableau{1}& 1&-1\\\\ L^c &\\bar{\\tableau{1}}& 1&+1\\\\ H& \\tableau{1} & -1&+1\\\\ H^c & \\bar{\\tableau{1}} &-1&-1\\\\ \\hline \\phi& \\identity& -2&0\\\\ \\hline \\end{array} \\end{eqnarray} %% The fermion sector has an $U(2)_{\\text{left}} \\times U(2)_{\\text{right}}$ chiral flavor symmetry, broken down to a vector-like $U(1)_H\\times U(1)_L$ by Yukawa interactions with a dark Higgs boson, $\\phi$. The vacuum expectation value of $\\phi$, $\\langle \\phi \\rangle = v_\\phi$, causes the Abelian gauge field to acquire a mass $m_{A_\\td} = 2\\sqrt{2} g_\\td v_\\phi$ and the fermions to pair into a light and a heavy Dirac fermion, %% \\begin{eqnarray} \\Psi_L = (L, \\bar{L}^c) \\qquad \\Psi_H = (H, \\bar{H}^c) \\end{eqnarray} %% with masses $m_L = y_L v_\\phi$ and $m_H=y_H v_\\phi$, respectively. In this minimal model, $m_H$ is near the electroweak scale while the other quark has a mass $m_L$ at or beneath the confinement scale, $\\Lambda_\\td\\sim \\OO(100\\MeV - 10\\GeV)$. In analogy to the Standard Model without weak interactions, the $U(1)_H\\times U(1)_L$ flavor symmetry renders the lightest mesons and baryons charged under this symmetry stable. The lightest stable bound state is a meson containing a single $\\Psi_H$ and a $\\bar{\\Psi}_L$, which will be denoted as $\\pid$. This dark meson is the dark matter candidate in \\cite{Alves:2009nf}. Because the constituents are fermions, $\\pid$ is paired with a vector state, $\\rhod$. The spin-spin interactions of the constituents generate a hyperfine splitting %% \\begin{eqnarray} \\label{Eq: HF} \\Delta E_{\\text{hyperfine}}= m_{\\rhod} - m_{\\pid} \\simeq \\left\\{ \\begin{aligned} \\frac{\\Lambda_\\td^2}{ m_H} , \\quad m_L< \\Lambda_\\td \\qquad \\\\ \\frac{\\alth^4m_L^2}{ m_H} , \\quad m_L> \\Lambda_\\td \\qquad \\end{aligned} \\right. \\end{eqnarray} %% where $\\alth$ is the $SU(N_c)$ 't Hooft coupling, and the parameters are chosen such that the splitting is at the scale $\\mathcal{O}(100\\keV)$ suggested by DAMA/LIBRA. The hyperfine structure is described in more detail in Sec. \\ref{FSSpec}. Another key feature of this model is that the dark gauge boson, $A_\\td^\\mu$, couples axially to the dark quarks. The coupling between the $A^\\mu_\\td$ and the dark mesons are constrained by parity and all leading order scattering channels are forbidden but the $\\pid -\\rhod$ transition \\cite{Alves:2009nf}. $A_\\td^\\mu$ mixes with the Standard Model photon and mediates dark meson/baryon scattering off SM nuclei. In particular, $\\pid$ up-scattering can explain the DAMA/LIBRA annual modulation signal. CiDM direct detection phenomenology is discussed in detail in \\cite{Alves:2009nf, Lisanti:2009vy, Lisanti:2009am}. The mass of the dark Higgs, $\\phi$, is radiatively unstable and introduces a second gauge hierarchy problem to the dark matter - Standard Model theory. The solution to the Standard Model's gauge hierarchy problem may also solve this new hierarchy problem. Supersymmetric extensions of these models may solve both hierarchy problems at once, and may introduce new phenomena into the theory. For instance, if the only communication of supersymmetry breaking to the dark sector occurs through kinetic mixing, then the dark sector may be nearly supersymmetric, resulting in nearly supersymmetric bound states \\cite{Rube:2009yc,Herzog:2009fw}. These susy bound states may have different scattering channels and the phenomenology may be different than minimal CiDM \\cite{Behbahani:2010}. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ", "conclusions": "\\label{Sec:Conclusion} In Composite inelastic Dark Matter, the $\\OO(100\\keV)$ mass splitting suggested by DAMA/LIBRA arises from dynamics inside composite states of a strongly interacting sector. Dark hadrons carrying non-zero flavor quantum numbers are stable and hence potential dark matter candidates. In particular, heavy flavor mesons offer a good candidate for implementing iDM. The thermal relic abundance of dark hadrons is too small to account for all the dark matter due to strong self-interactions, meaning that the dark matter abundance in these theories must originate from a dark flavor (or dark baryon) asymmetry. One of the purposes of this work was to investigate the cosmological evolution of the flavor asymmetry and how it determines the mesonic and baryonic abundances at late times. This question was addressed in the context of the minimal model of Sec.~\\ref{intro:summary}, where the iDM candidate is a spin-0 dark meson $\\pid$ that inelastically scatters to a nearly degenerate vector state $\\rhod$. In a large part of the parameter space, $\\pid$ mesons dominate the abundance of dark matter. In other regions, dark baryons dominate the abundance; however, the residual $\\pid$ component interacts sufficiently strongly to give rise to a viable iDM scenario. In all cases, exotic dark matter components arise with novel elastic scattering properties -- nuclear recoil events are suppressed at low-energy by a dark matter form factor. The relative abundance of $\\rhod$ to $\\pid$ is typically $\\sim 10^{-4}$, suggesting that $\\rhod \\rightarrow \\pid$ down scattering off nuclei is a discoverable signal in the near future. Finally, we studied a variety of long lived meson and baryons states and found that BBN constraints on their decays are not severe. \\vspace{0.2in} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\noindent { \\bf Acknowledgements \\vspace{0.05in}} We would like to thank Mariangela Lisanti for helpful feedback as well as Rouven Essig and Natalia Toro for illuminating discussions. DSMA, SRB, PCS and JGW are supported by the US DOE under contract number DE-AC02-76SF00515. JGW is supported by the DOE's Outstanding Junior Investigator Award. \\appendix %%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%" }, "1003/1003.3626_arXiv.txt": { "abstract": "We report the detection of a strong, organized magnetic field in the helium-variable early B-type star HR~7355 using spectropolarimetric data obtained with ESPaDOnS on the 3.6-m Canada-France-Hawaii Telescope within the context of the Magnetism in Massive Stars (MiMeS) Large Program. HR~7355 is both the most rapidly rotating known main-sequence magnetic star and the most rapidly rotating helium-strong star, with $v \\sin i$ = 300 $\\pm$ 15 km~s$^{-1}$ and a rotational period of 0.5214404 $\\pm$ 0.0000006 days. We have modeled our eight longitudinal magnetic field measurements assuming an oblique dipole magnetic field. Constraining the inclination of the rotation axis to be between $38^{\\circ}$ and $86^{\\circ}$, we find the magnetic obliquity angle to be between $30^{\\circ}$ and $85^{\\circ}$, and the polar strength of the magnetic field at the stellar surface to be between 13-17 kG. The photometric light curve constructed from HIPPARCOS archival data and new CTIO measurements shows two minima separated by 0.5 in rotational phase and occurring 0.25 cycles before/after the magnetic extrema. This photometric behavior coupled with previously-reported variable emission of the H$\\alpha$ line (which we confirm) strongly supports the proposal that HR~7355 harbors a structured magnetosphere similar to that in the prototypical helium-strong star, $\\sigma$ Ori E. ", "introduction": "Hot, massive stars are not expected, a priori, to exhibit magnetic fields as they lack the convective envelopes necessary for driving a dynamo. However, Ap/Bp stars (chemically peculiar A and B type stars) are well known to possess organized magnetic fields with surface strengths up to tens of kG (e.g. Borra \\& Landstreet 1979). The hottest of the Bp stars are the so-called helium-strong stars - early B-type stars on the main sequence with enhanced and typically variable helium lines. Some helium-strong stars show additional emission and variability in Balmer lines, variable photometric brightness and color, variable UV resonance lines, and non-thermal radio emission (Pedersen \\& Thomsen 1977; Walborn 1982; Shore \\& Brown 1990; Leone \\& Umana 1993). All of these quantities vary with a single period, interpreted to be the rotational period of the star. The Balmer line (specifically H$\\alpha$) and UV line variability point to the presence of a stellar wind coupled with a magnetic field to form a circumstellar magnetosphere (Shore \\& Brown 1990). Following their discovery of a magnetic field in the B2Vp star $\\sigma$ Ori E (Landstreet \\& Borra 1978), Borra \\& Landstreet (1979) investigated eight helium-strong stars, finding six new magnetic stars. In that paper, the authors established that a defining characteristic of the helium-strong stars is the presence of an strong, organized magnetic field. HR~7355 (HD~182180), the subject of the present paper, is a bright (V=6.02) helium-strong early B-type star with high $v \\sin i$ (320 km~s$^{-1}$, Abt et al 2002; 270 km~s$^{-1}$, Glebocki \\& Stawikowski 2000). Originally classified as B5IV by Morris (1961), Hiltner et al. (1969) re-classified HR~7355 B2Vn, a classification that has been used subsequently. HR~7355 has also been classified as a classical Be star (Abt \\& Cardona 1984) due to observed emission in the H$\\alpha$ line of its spectrum. HIPPARCOS photometry exhibits variability with a period of $\\sim$ 0.26 days (Koen \\& Eyer 2002). Rivinius et al. (2008) reported emission and variability of the H$\\alpha$ profiles, as well as variability of a variety of spectral absorption lines (i.e. He, C, Si). Those authors revisited the HIPPARCOS photometry noting that a period of 0.52 days (twice the period reported by Koen \\& Eyer (2002)) could also explain the photometric variation. Ultimately, based on the helium variability, the variable H$\\alpha$ emission, and the character of the photometric variation, Rivinius et al. (2008) proposed that HR~7355 hosts a structured magnetosphere qualitatively similar to that of $\\sigma$ Ori E. With its very rapid rotation, the establishment of such a magnetosphere around HR~7355 would be of great interest. We have therefore undertaken observations to search for the presence of a magnetic field in this star, and to further explore its spectroscopic and other physical properties. ", "conclusions": "The sinusoidal variation of the longitudinal magnetic field indicates an important dipole component to the stellar magnetic field. Using the characteristics of the sinusoidal fit to the magnetic data and the stellar parameters with their error bars (derived in \\S3), we can estimate the geometry and strength of the dipole: the inclination of the rotational axis with respect to the line of sight to the observer, $i$, the obliquity of the dipole axis relative to the rotational axis, $\\beta$, and the polar strength of the surface dipole field, $B_{d}$, as described by e.g. Wade et al. (1997). We have determined the allowed range of inclinations $38^{\\circ} \\le i \\le 86^{\\circ}$, assuming rigid rotation. Lower inclinations were not feasible as the implied equatorial rotation velocity (i.e. $v \\sin i / \\sin i$) was higher than the critical break-up velocity (v$_{\\rm{crit}}$; Townsend et al. 2004) calculated for the mass and radius considered. Using this range of $i$ and the observed values of the magnetic extrema (Fig. \\ref{fullper}, lower panel), we calculate the obliquity to be $30^{\\circ} \\le \\beta \\le 85^{\\circ}$. With this inclination and obliquity, and the maximum longitudinal field, we compute that the surface dipolar field strength for HR~7355 is 13-17 kG at the pole. The Rigidly Rotating Magnetosphere (RRM) model was proposed by Townsend \\& Owocki (2005) to describe a rapidly rotating star in which the magnetic field overpowers the stellar wind, allowing plasma to become trapped in a magnetosphere that co-rotates along with the star. The geometry of the magnetosphere depends on the obliquity angle and the rotational velocity of the star (See Townsend 2008). The variability of the observables depends on the inclination angle. This model was applied to the prototypical helium-strong star, $\\sigma$ Ori E, by Townsend et al. (2005). Comparing our picture of the observations and geometry of HR~7355 with the observations and model of $\\sigma$ Ori E, the two show similarities in both the variability of observables, but also in the relative timing of each observation related to the geometry. In both stars, the minimum of the photometric light curve corresponds to minimum emission in H$\\alpha$, as well as a null in the longitudinal field curve. Although the photometric light curves both show similar V-band magnitude changes ($\\sim$0.05 mags), the durations of the photometric ``eclipses'' in HR~7355 are slightly longer than those in $\\sigma$ Ori E; this is not surprising given the more rapid rotation of HR~7355. While the phasing of the H$\\alpha$ variations and the photometric minima of HR~7355 are consistent with circumstellar material, the photometric variations may also be due to photospheric spots on the surface of the star. Based on this small initial dataset, the RRM model appears to be capable of explaining the basic observational phenomena presented by HR~7355. We therefore expect it to be useful to confirm and characterize the magnetosphere of HR~7355 once more data are acquired. HR~7355 is both the most rapidly rotating helium-strong star and the most rapidly rotating magnetic star discovered thus far with a 0.5214404 $\\pm$ 0.0000006 day period and a $v \\sin i$ of 300 km~s$^{-1}$, surpassing the Bp star CU Vir ($P_{\\rm{rot}}$ = 0.52070308 d, $v \\sin i$ = 160 km~s$^{-1}$, Kuschnig et al. 1999). HR~7355 has a strong magnetic field ($B_{d} \\sim$ 13-17~kG) that when coupled with its wind produces metal line variability and a magnetosphere that co-rotates with the star, resulting in variability of H$\\alpha$ and, photometric brightness. HR~7355 is rotating near its critical velocity, making it an extreme laboratory to study the effects of a magnetosphere under these conditions, and perhaps even to provide a link between Bp and Be stars. However, the high rotation of this star also creates challenges for spectral line analysis. At rotational velocities near critical, the star becomes oblate and gravity darkening effects become important. When viewed from the equator, the star can appear reddened and the temperature and surface gravity are lowered (Townsend et al. 2004). Detailed line-profile modeling is required to determine accurate parameters. Although historically HR~7355 is classified as B2Vn, the temperature derived from its spectrum suggests a later type. In addition to continuing to monitor the magnetic field, line profile and photometric variations of HR~7355, we intend to undertake, in a future study, a more physically realistic line profile analysis, taking into account rotational deformation and gravity darkening, to provide a consistent solution for stellar parameters and evolution." }, "1003/1003.4230_arXiv.txt": { "abstract": "{We present a new technique to fit color-magnitude diagrams of open clusters based on the Cross-Entropy global optimization algorithm. The method uses theoretical isochrones available in the literature and maximizes a weighted likelihood function based on distances measured in the color-magnitude space. The weights are obtained through a non parametric technique that takes into account the star distance to the observed center of the cluster, observed magnitude uncertainties, the stellar density profile of the cluster among others. The parameters determined simultaneously are distance, reddening, age and metallicity. The method takes binary fraction into account and uses a Monte-Carlo approach to obtain uncertainties on the determined parameters for the cluster by running the fitting algorithm many times with a re-sampled data set through a bootstrapping procedure. We present results for 9 well studied open clusters, based on 15 distinct data sets, and show that the results are consistent with previous studies. The method is shown to be reliable and free of the subjectivity of most previous visual isochrone fitting techniques.} ", "introduction": "Galactic open clusters are a key class of objects used in a wide range of investigations due to their wide span of ages and distances as well as the precision to which these parameters can be determined through their color-magnitude diagrams (CMD). With the publication of the Hipparcos Catalog (ESA 1997) and its derivatives the Tycho and Tycho2 \\citep{ESA1997,Hog2000} as well as individual efforts using modern ground based instrumentation there has been increased interest in studies involving open clusters in the Galaxy. Individually the results obtained from the study of open clusters can provide important constraints for theoretical models of stellar formation and evolution. Comparison of observed CMDs to model isochrones can provide important information on the effect of overshooting \\citet{VandenBerg2004} and chemical abundances \\cite{Meynet1993, Kassis1997}. Open clusters can also be used in the study of variable stars, especially Cepheids (\\cite{Kang2008,Majaess2009}), in the search for connections between stellar magnetic fields and their evolution and in the search for extra-terrestrial planets, among many others. Our group has focused efforts in the investigation of open clusters and their use in the understanding of the Galactic spiral structure (\\cite{Dias2005,lepine2008}). The results are based on the catalog of open clusters published in \\cite{dias2002} witch is now in version 2.10 and, since 2002, can be accessed on line at http://www.astro.iag.usp.br/$\\sim$wilton. The results presented in our catalog are compiled from data published from many authors, using different instruments, techniques, calibrations and criteria which results in a heterogeneous sample. However, in \\cite{PN06}, the authors show that our data has the same statistical significance as the data they use to define the standard parameters of the chosen open clusters when considering the calculated errors. To derive the fundamental parameters of open clusters the main sequence ``fitting'' in most works up to today has been done mostly with subjective visual fitting. This is mainly due to the fact that the isochrones have no simple parametric form so that a usual least square technique can be applied. The available isochrones in the literature are usually in the form of tabulated points for a set of fundamental parameters such as age and metallicity. The traditional ``fitting'' method utilized is to first determine the reddening by adjusting a Zero Age Main Sequence (ZAMS) to the observed color-color diagram (usually $(B-V)$ vs. $(U-B)$ of the cluster and then, keeping this value fixed, adjusting the distance and age using the observed CMD and tabulated isochrones. Both of these steps are usually performed by eye as mentioned, leading to inevitable subjectivity in the fitting of an isochrone to a given observed CMD. The subjectivity in the determination of the reddening is specially problematic since it affects the subsequent determination of the distance and age of the cluster. It is also important to point out that in some cases the reddening is not determined through the use of U filter photometry which can further compromise the results obtained. The lack of homogenization of the data and especially the subjectivity of the methods used to obtain fundamental parameters of clusters from their CMD indicate the great need for a method that circumvents at least the subjectivity. The subject of automatic fitting and non subjective criteria to choose the best fit was addressed in some papers such the recently published work by \\cite{Naylor2006}, where the authors propose a maximum likelihood method for fitting two dimensional model distributions to stellar data in a CMD. In this work the authors also discuss the most important attempts at performing isochrone fits using different methods and we refer the reader to their very complete discussion and references therein. The authors also present the main problems involved in this kind of study, as the data precision (see also \\cite{Narbutis2007}), the importance of non resolved binaries (see \\cite{Schlesinger1975, Fernie1961} and in more modern observational studies \\cite{Montgomery1993, vonHippel1998}). In this work we present a new technique to fit models to open cluster photometric data using a weighted likelihood criterion to define the goodness of fit and a global optimization algorithm known as Cross-Entropy to find the best fitting isochrone. We have successfully applied the optimization algorithm in the study of jet precession in Active Galactic Nuclei (\\cite{CMA09}) where we demonstrate the robustness of the method. In the present work we adapt the method to study open clusters and show that it can find the best parameters and eliminate the subjectivity (author analysis dependence) of the main sequence fitting process by using well defined control parameters and a weighted likelihood. In the next section we introduce the optimization technique and how it was adapted to the problem of ishocrone fitting. In Sec 3. we define the likelihood function used and in Sec. 4 how the weight function is obtained. In Sec. 5 we demonstrate the validity of the method by applying it to synthetic clusters and in Sec. 6 we apply the method to the data of 9 clusters carefully chosen. In Sec. 7 we discuss the results and in Sec. 8 we give our final conclusions. \\section[]{Cross entropy global optimization} \\subsection{The optimization method} The Cross Entropy technique (CE) was first introduced by \\cite{rubi97}, with the objective of estimating probabilities of rare events in complex stochastic networks, having been modified later by \\cite{rubi99} to deal with continuous multi-extremal and discrete combinatorial optimization problems. Its theoretical asymptotic convergence has been demonstrated by \\cite{marg04}, while \\cite{kro06} studied the efficiency of the CE method in solving continuous multi-extremal optimization problems. Some examples of robustness of the CE method in several situations are listed in \\cite{deb05}. The CE procedure uses concepts of importance sampling, which is a variance reduction technique, but removing the need for a priory knowledge of the reference parameters of the parent distribution. The CE procedure provides a simple adaptive way of estimating the optimal reference parameters. Basically, the CE method involves an iterative procedure where in each iteration the following is done: \\begin{enumerate}[i] \\item Random generation of the initial parameter sample, respecting pre-defined criteria; \\item Selection of the best candidates based on some mathematical criterion; \\item Random generation of updated parameter samples from the previous best candidates to be evaluated in the next iteration; \\item Optimization process repeats steps (ii) and (iii) until a pre-specified stopping criterion is fulfilled. \\end{enumerate} The CE algorithm is based on a population of solutions in a similar manner to the well known genetic algorithm and in this sense it is also a type of evolutive algorithm where some fraction of the population is selected in each iteration based on a given selection criteria. However, a detailed discussion of evolutive algorithms and their similarities and comparative performance is beyond the scope of this work. In the work of \\cite{deb05}, many standard benchmark optimization problems, such as the Travelling Salesman, are studied using the CE method and its efficiency is also discussed. A great advantage of the CE method over the genetic algorithm for example is its simplicity to code. For problems with many free parameters there is no need to deal with genes and their definitions, crossing, mutation rates and other details. Here we have implemented the CE method to find the best fitting model isochrone to open cluster data as discussed in the following sections. \\subsection[]{Cross entropy isochrone fitting} Let us suppose that we wish to study a set of $N_{d}$ observational data in terms of an analytical model characterized by $N_{p}$ parameters $p_1, p_2, ..., p_{N_{p}}$. The main goal of the CE continuous multi-extremal optimization method is to find a set of parameters $p^*_i$ ($i=1, ..., N_{p}$) for which the model provides the best description of the data (\\cite{rubi99,kro06}). It is performed generating randomly $N$ independent sets of model parameters ${\\bf X}=({\\bf x}_1,{\\bf x}_2,...,{\\bf x}_N)$, where ${\\bf x}_i=(p_{1_i},p_{2_i},...,p_{N_{{p}i}})$, and minimizing the objective function $S({\\bf X})$ used to transmit the quality of the fit during the run process. If the convergence to the exact solution is achieved then $S\\rightarrow 0$, which means ${\\bf x}\\rightarrow{\\bf x}^*=(p^*_1,p^*_2,...,p^*_N)$. In order to find the optimal solution from CE optimization, we start by defining the parameter range in which the algorithm will search for the best candidates: $\\xi^{min}_i\\leq p_i \\leq \\xi^{max}_i$. Introducing $\\bar{\\xi}_i(0)=(\\xi^{min}_i+\\xi^{max}_i)/2$ and $\\sigma_i(0)=(\\xi^{max}_i-\\xi^{min}_i)/2$, we can compute ${\\bf X}(0)$ from: \\begin{equation} X_{ij}(0)=\\bar{\\xi}_i(0)+\\sigma_i(0) G_{i,j}, \\end{equation} where $G_{i,j}$ is an $N_{p}\\times N$ matrix with random numbers generated from a zero-mean normal distribution with standard deviation of unity. The next step is to calculate $S$ for each component of ${\\bf X}(0)$, ordering them from the lowest to the highest value of $S$. Then the first $N_{elite}$ set of parameters is selected, i.e. the $N_{elite}$-samples with lower $S$-values, which will be labeled hereafter as the elite sample array ${\\bf X}^{elite}$. Having determined ${\\bf X}^{elite}$ at $k$th iteration, the mean and standard deviation of the elite sample are calculated, $\\bar{{\\bf x}}^{elite}_i(k)$ and ${\\bf \\sigma}^{elite}_i(k)$ respectively, using: \\begin{equation} \\bar{{\\bf x}}^{elite}_i(k)=\\frac{1}{N_{elite}}\\sum\\limits_{j=1}^{N_{elite}}X^{elite}_{ji}, \\end{equation} \\begin{equation} {\\bf \\sigma}^{elite}_i(k)=\\sqrt{\\frac{1}{\\left(N_{elite}-1\\right)}\\sum\\limits_{j=1}^{N_{elite}}\\left[X^{elite}_{ji}-\\bar{{\\bf x}}^{elite}_j(k)\\right]^2}. \\end{equation} In order to prevent convergence to a sub-optimal solution due to the intrinsic rapid convergence of the CE method, \\cite{kro06} suggested the implementation of a fixed smoothing scheme for $\\bar{{\\bf x}}^{elite,s}_i(k)$ and ${\\bf \\sigma}^{elite,s}_i(k)$: \\begin{equation} \\bar{{\\bf x}}^{elite,s}_i(k)=\\alpha^\\prime\\bar{{\\bf % x}}^{elite}_i(k)+\\left(1-\\alpha^\\prime\\right)\\bar{{\\bf x}}^{elite}_i(k-1), \\end{equation} \\begin{equation} {\\bf \\sigma}^{elite,s}_i(k)=\\alpha_{d}(k){\\bf % \\sigma}^{elite}_i(k)+\\left[1-\\alpha_{d}(k)\\right]{\\bf \\sigma}^{elite}_i(k-1), \\end{equation} where $\\alpha^\\prime$ is a smoothing constant parameter ($0<\\alpha^\\prime< 1$) and $\\alpha_{d}(k)$ is a dynamic smoothing parameter at $k$th iteration: \\begin{equation} \\alpha_{d}(k)=\\alpha-\\alpha\\left(1-k^{-1}\\right)^q, \\end{equation} with $0<\\alpha< 1$ and $q$ being an integer typically between 5 and 10 (\\cite{kro06}). As mentioned before, such parametrization prevents the algorithm from finding a non-global minimum solution since it guarantees polynomial speed of convergence instead of exponential. The array ${\\bf X}$ at $k$th iteration is determined analogously to equation (1): \\begin{equation} X_{ij}(k)=\\bar{{\\bf x}}^{elite,s}_i(k)+{\\bf \\sigma}^{elite,s}_i(k) G_{i,j}, \\end{equation} The optimization stops when either the mean value of ${\\bf \\sigma}^{elite,s}_i$ is smaller than a pre-defined value or the maximum number of iterations $k_{max}$ is reached. \\section[]{Defining the likelihood} To implement an objective function we use a weighted likelihood function similar to the one proposed by \\cite{FJ82} and Hernandez \\& David Valls-Gabaud (2008) to fit theoretical isochrones to star cluster data. In the first work the authors derive what they called the Near Point Estimator based on a Maximum Likelihood analysis of cluster data where the fitting statistic measures the overall coincidence of model isochrones and observed data, assuming uniform star distribution along the isochrone. The Near Point Estimator is derived in \\cite{FJ82} and we refer the reader to this work for the details. The authors show that the probability of a given measurement to be a cluster star related to a given theoretical isochrone is given by $ln(Prob) \\propto d_{min}^2$, where $d_{min}^2$ is the minimum distance from the observed point to the model isochrone, being valid for any number of distinct measurements. In the latter work the authors use a similar method but take the Bayesian approach to solving the problem by setting a likelihood function very similar to the Near Point Estimator of the previous work. Both works relate the probability of a given star of belonging to a given model isochrone to distances in the CMD. In this work we adopt a similar path to define our likelihood function which will then be maximized by the CE method described previously. The major difference is that we include a weighting factor, discussed in detail below, in a semi-Bayesian approach, which defines our prior knowledge based on the observed data-set. Unlike the work of \\cite{FJ82} we do not assume uniform distribution of stars along the isochrone. To accomplish this we define the isochrone points by sampling from an initial mass function (IMF), randomly generating a number of stars in the mass range of the original isochrone. Because we sample from a given IMF, we are also able to directly account for binaries. We have done so assuming a binary fraction of 100\\% with companions drawn randomly from the same IMF as the cluster stars. With these randomly generated stars we obtain a synthetic cluster which can be compared to the observed data set through a given metric. Because this procedure populates the CMD in the correct manner through the IMF there is no need to introduce approximations to account for weights along the isochrone, therefore a statistic which sums over the distances in magnitude space of a given observed star to the generated points will give the observed point probability of belonging to that specific model isochrone, which is calculated by: \\begin{eqnarray} P(V,BV,UB|I_N)_l= \\sum_{m} \\frac{1}{\\sigma_{V_l}\\sigma_{BV_l}\\sigma_{UB_l}} \\times \\\\ EXP\\left [-\\frac{1}{2} \\left ( \\frac{V_l-{I_N}_{,V_m}} {\\sigma_{V_l}} \\right )^2 \\right ] \\times \\nonumber\\\\ EXP\\left [-\\frac{1}{2}\\left ( \\frac{BV_l-{I_N}_{,BV_m}}{\\sigma_{BV_l}} \\right )^2 \\right ] \\times \\nonumber\\\\ EXP\\left [-\\frac{1}{2}\\left ( \\frac{UB_{l}-{I_N}_{,UB_m}}{\\sigma_{UB_l}} \\right )^2 \\right ] \\nonumber \\end{eqnarray} where $I_N$ is a tabulated Isochrone function defined by $m$ points, $V_l$ is the observed V magnitude, $BV_l$ and $UB_l$ the color indexes, $\\sigma_{V_l}$ is the error on the $V$ magnitude of star $l$, $\\sigma_{BV_l}$ and $\\sigma_{UB_l}$ are the errors on the $BV$ and $UB$ colors of star $l$. Because the $I_N$ isochrone is a discretely tabulated function, the continuous optimization algorithm described previously was adapted to find the nearest parameters thus introducing a grid resolution error in the values obtained. The grid used is comprised of isochrones of ages from $log(age)=6.6$ to $log(age)=10.15$ with a step of $log(age)=0.05$. The uncertainty resulting from the grid resolution has been incorporated in the final quoted errors. The weighted likelihood function is then given in the usual manner by: \\begin{equation} \\mathcal L = \\prod_{l}^{N_d}P(V,BV,UB|I_N)_l \\times W_l \\end{equation} where $W_l$ is the weight for a given star as determined from the data using the non-parametric technique described in the following section. The likelihood above is used to define the objective function $S({\\bf X})$ of the optimization algorithm as follows: \\begin{equation} S({\\bf X}) = -log(\\mathcal L({\\bf X})) \\end{equation} where ${\\bf X}$ is the vector of parameters that define a given isochrone ${\\bf I_N}$ and the optimization is then done with respect to N. \\section[]{Determining the weight function} Before determining the weights of each observed star we must first deal with the fact that the data is not free of contamination from stars that do not belong to the cluster itself. The contamination is usually from field stars that can be at different distances and have different reddening values than that of the cluster. Therefore we must introduce schemes to filter out contaminating stars as well as to determine which stars are more likely to belong to the cluster in a way that can be easily reproduced given simple and clear parameters. \\subsection[]{Magnitude cut-off} The first step in the decontamination is to inspect the magnitude cut-off of the observations. In many cases the observations are clearly not complete down to the faintest magnitude as can be easily seen in the histogram of the observed $V$ magnitude. With the $V$ magnitude histogram, obtained with a bin size of 0.5, we determine in which magnitude it peaks and then reject stars that have magnitudes higher than this threshold. This procedure takes care of a good part of the contamination from faint stars, that are in fact intrinsically brighter and further away than the cluster itself and thus show up as faint stars in the CMD. However, this procedure is not sufficient to remove contamination from stars that are within the completeness limit of the observation. The stars that are within the completeness limit of the observations are usually the main source of contamination in the fainter (lower mass) regions of the CMD where larger magnitude errors also contribute to the confusion. Typically when a field is very crowded and this contamination type is large, it can be identified as a triangular region where the stars concentrate with high density and no clear clustering around a main sequence. To treat these situations when there is high density of field stars and no detectable clustering around an isochrone main sequence we introduced a magnitude cut-off that can be defined by the user. Since the turn-off point and the red giant region of the open clusters are the two most important features that constrain the determination of fundamental parameters, this cut-off does not affect negatively their final value in most situations. \\subsection[]{Cluster density profile} Having eliminated the most obvious types of contamination we proceed to estimating the weight of a given cluster star using non-parametric techniques. The main assumption made here is that cluster stars are concentrated in a limited region in the observed field and also in a specific region of the CMD assuming they formed according to a single isochrone. To determine the region in the observed field where the cluster stars are located we use the position of the cluster center, usually provided with the observational data set obtained from the literature. In some cases where this information is not provided it can be determined by obtaining a two dimensional histogram of star positions and determining the location of the peak, provided sufficient number of stars. With the position of each star measured relative to the cluster center we obtain their radial distance (measured in pixels). We then use the radial distances to calculate the number of stars at a given radial distance using a histogram. The bin size of the histogram is calculated with $bin=0.05\\times MAX({\\bf r})$, where ${\\bf r}$ is the vector of all star radial distances. For a given cluster $h$ bins will determined and the density of stars as a function of radial distance is then estimated by $\\rho({\\bf r}_h)=N_h/4 \\pi {\\bf r}_h^2$, where $N_h$ is the number of stars in the $h$th bin. Typically the density profile of a cluster falls off as the radius increases. Integrating this density profile we can define a cluster radius where we find a given percentage of the total number of stars in the field. The user defined percentage value, which we call star fraction or $F_{star}$, is then used to define two regions, Cluster and Field, using the following integral: \\begin{equation} \\int_0^{R_{cluster}} \\rho(r)/N_{star}~dr = F_{star} \\end{equation} where $R_{cluster}$ is the cluster radius for the given $F_{star}$ fraction. \\subsection[]{Photometric uncertainty} To determine the weight of each star we also need the photometric error $\\sigma_{phot}$ of the observed data. The error is defined as a percentage and errors for each star in magnitude and color are calculated using this factor. For the color errors we assume that the photometric errors in each filter are independent of each other and therefore calculate the error with the usual propagation formula. Since few works in the literature include a full study of the errors involved in obtaining magnitudes we adopted values that were consistent with the ones obtained by \\cite{M01}, where the author presents error values for stars observed multiple times with the same instrumentation. We also used the results of the filtering described below to guide the final error value adopted, aiming for the most efficient elimination of contamination. \\subsection[]{Non-parametric weight function} With the Cluster and Field regions as well as the photometric errors defined, the weight for each star is then estimated by comparing the characteristics of the stars in an area around the $lth$ star in the CMD defined by a box with dimensions $3\\sigma_{V_{l}}$ by $3\\sigma_{(B-V)_{l}}$. We then calculate the average and standard deviation of $V$ and $(B-V)$ for the stars that fall within this box and belong to the Cluster region defined earlier. The assumption made here is that this statistic provides the most likely position for a cluster star in the CMD region defined by the error box of the $lth$ observed star. This is clearly not the case when there is no detectable concentration of cluster stars relative to field stars within the error box. This situation is more likely to happen in clusters with heavy field contamination and large magnitude errors as in the lower magnitude regions of the CMD. However, as discussed earlier the magnitude cut-off can resolve this issue by eliminating these stars from the sample to be fitted. We then calculate the weight for the $l$th star with the expression: \\begin{eqnarray} W_{l} = \\frac{1}{\\sigma_{V_{l}}\\sigma_{BV_{l}}\\sigma_{UB_{l}}} \\times EXP^{\\frac {-(V_l-\\overline{V_c})^2} {2\\sigma^2_{V_c}} } \\times \\nonumber\\\\ EXP^{\\frac{-(BV_l-\\overline{BV_c})^2} {2\\sigma^2_{BV_c}} }\\times \\nonumber\\\\ EXP^{\\frac{-(UB_l-\\overline{UB_c})^2} {2\\sigma^2_{UB_c}} }\\times \\nonumber\\\\ EXP^{\\frac{-r_l^2} {2 \\left ( \\frac{R_{cluster}} {3} \\right ) ^2} } \\end{eqnarray} where $V_{l}$, $BV_{l}$ and $UB_{l}$ are the observed V magnitude and the color indexes of the $lth$ star, $\\overline{V_{c}}$, $\\overline{BV_{c}}$ and $\\overline{UB_{c}}$ are the average V magnitude and the average color indexes of the stars that fall within the $3\\sigma$ error box and belong to the Cluster region as defined earlier. Note that according to this procedure stars that fall outside of the Cluster region are automatically given $W_{l}=0$. In Fig. 1 we show the results of the decontamination and weighting process for the cluster \\object{NGC~2477} and data set of \\cite{Kassis1997}. The black dots in the left graph are the selected stars after decontamination, the open circles are stars that fall outside the defined cluster radius (which are then eliminated from the sample to be fitted) and light dots are stars for which no statistic was available due to low numbers. The right graph shows open circles with sizes scaled to their weights, with larger sizes meaning larger weights. It is clear from these graphs that the decontamination scheme we have adopted is not perfect and in regions where field stars have large densities in the CMD non-cluster stars are likely to survive the process. However, as we can see in the right graph, the weighting procedure does a good job of assigning low values to these fields stars. In any case, as discussed before, we have introduced a magnitude cut-off to eliminate these regions altogether from the fitting process if necessary. \\subsection[]{Implementing the Cross entropy algorithm} In our problem $I_N$ is a tabulated function taken from a grid of models calculated by Padova database of stellar evolutionary tracks and isochrones \\cite{Girardi2000, Marigo2008}. \\begin{figure*} \\begin{center} \\includegraphics[scale=0.5]{fig-selected} \\includegraphics[scale=0.5]{fig-pesos} \\end{center} \\caption{Result of the decontamination (left) and weighting process (right) for the cluster \\object{NGC~2477}. The black dots in the left graph are the selected stars after decontamination, the open circles are stars that fall outside the defined cluster radius and light dots are stars for which no statistic was available due to low numbers. The right graph shows open circles with sizes scaled to their weights, with larger sizes meaning larger weights.} \\end{figure*} The tabulated isochrones are defined by 2 parameters, namely, age and metallicity and to compare to observed data we also need distance and extinction, which we consider constant. These are the parameters we wish to optimize thus fitting the tabulated isochrone to the observed data. The parameter ranges for generating ${\\bf X}$ are pre-defined by the user and should be representative of the problem being optimized. In general, the CE algorithm is very forgiving of large parameter spaces, being very efficient in quickly zoning in on optimal regions. In the isochrone fitting done in this work we defined the parameter space as follows: \\begin{enumerate} \\item {\\bf Age}: from $log(age)=6.60$ to $log(age)=10.15$ encompassing the full range of theoretical isochrones; \\item {\\bf distance}: from 1 to 10000 parsecs \\item ${\\bf E(B-V)}$: from 0.0 to 3.0 \\item {\\bf Z}: fixed (used literature values) \\end{enumerate} The algorithm has been written to allow the optimization of the metallicity as well and work is under way to fully implement this feature. However, since we are mainly concerned with benchmarking the method and comparing results to the literature values provided by \\cite{PN06} we have kept it fixed to values obtained from the literature. \\begin{table*} \\caption{Results for synthetic clusters studied by the fitting method.} \\centering \\begin{tabular}{lcccccc} \\hline\\hline {Cluster} & {$N_{stars}$} & {Contamination (\\%)} & {$3\\sigma_{phot}$(\\%)} & {$E(B-V)$(mag)} & {Distance (pc)} & {$log(Age)$ (yr)} \\\\ \\hline SC~01 & 432 & 0\\% & 1.0 & $0.40 \\pm 0.01$ & $2112 \\pm 51$ & $8.65 \\pm 0.05$ \\\\ SC~02 & 480 & 20\\% & 1.0 & $0.40 \\pm 0.01$ & $2062 \\pm 43$ & $8.71 \\pm 0.05$ \\\\ SC~03 & 444 & 50\\% & 1.0 & $0.38 \\pm 0.01$ & $2073 \\pm 30$ & $8.70 \\pm 0.07$ \\\\ SC~04 & 65 & 0\\% & 1.0 & $0.38 \\pm 0.02$ & $2008 \\pm 94$ & $8.70 \\pm 0.07$ \\\\ SC~05 & 113 & 20\\% & 1.0 & $0.40 \\pm 0.03$ & $2060 \\pm 60$ & $8.75 \\pm 0.06$ \\\\ SC~06 & 61 & 50\\% & 1.0 & $0.40 \\pm 0.02$ & $2102 \\pm 58$ & $8.70 \\pm 0.07$ \\\\ \\hline \\end{tabular} \\tablefoot{Synthetic clusters were generated with parameters log(age)=8.70 yr, distance=2100 pc, E(B-V)=0.40, Z=0.019) and $N_{stars}$, including the given contamination fraction and photometric accuracy of $3\\sigma_{phot}$.} \\label{tab1} \\end{table*} The filtered sample is then fed through the optimization algorithm described previously which then minimizes the objective function ${\\bf S}$ thus maximizing the likelihood function to find the best fitting values for age, distance and reddening for a given metallicity. An important point to be made is that although in the algorithm all parameters can be fit simultaneously, we opted to take the usual procedure of determining the $E(B-V)$ parameter first, using only the colour-colour diagram and the ZAMS, and then performing the fit for the other parameters in the CMD. However, to ensure that the fit has some liberty in accommodating other possibilities in the second stage of the fitting we allow $E(B-V)$ to vary in a range of $10\\%$ of the value determined in the first step. An advantage of this fitting procedure is that it allows for determination of parameter errors through Monte-Carlo techniques. To accomplish this we perform the fit for each data set $N_{Run}$ times, each time re-sampling from the original data set with replacement to perform a bootstrap procedure as well as generating new isochrone points from the adopted IMF as described previously. For each run we also replace the stars chosen in the new bootstrap sample with ones obtained by randomly generating values of $(U,B,V)_l$ drawn from a normal distribution centered at the original data value and with $\\sigma=\\sigma_{phot}$. The final uncertainties in each parameter are obtained by calculating the standard deviation of the $N_{Run}$ fit values. ", "conclusions": "As pointed out by \\cite{PN06}, studying phenomena using open cluster physical parameters, is highly dependent on their precision. The authors also show that major discrepancies still exist even in well studied clusters. In our work we provide a new technique for the determination of open cluster physical parameters that is not dependent on the user and is reproducible within the statistical uncertainties given well defined conditions. Our method, based on the Cross-Entropy optimization algorithm, was tailored to the fitting of theoretical isochrones as the ones of \\cite{Girardi2000, Marigo2008} used in this work. The procedure is simple and allows for the use of any tabulated theoretical isochrones and thus provides also an unbiased means of comparing fits using different theoretical models given the same constraints and fitting procedure. In this work we have concentrated in the validation of the method limiting ourselves to fitting synthetic clusters and well studied open clusters with the tabulated isochrones of Padova. The results show that the method is capable of recovering the original parameters with good accuracy even in cases where we included considerable non-uniform field contamination demonstrating its robustness. The validation using these synthetic clusters is by no means complete, however, we explored the most typical situations present in open cluster data. The results using the observed data available in the literature show that the parameters determined through our technique are consistent with the results obtained by other authors and especially those given by \\cite{PN06} as shown in Table 3. In all cases where we encountered significant discrepancies these could be explained by data quality, level of contamination, which stars were selected by our filtering technique or some combination of these. In any case, the fact that all the steps related to our fitting procedure are quantifiable, allows us to perform objective comparisons of different parameters for a given data set, removing the subjectivity of which stars are selected. The filtering and weighting technique defined by a precise set of conditions is central to our method. As shown in the results for Trumpler 1, in cases where the cluster has low sampling, i.e. low number of stars, it becomes difficult to determine the weight for cluster stars. It is likely that better statistical tools may improve the efficiency of this step. In general, the filtering technique performed well in eliminating most of the contamination and assigning weights in the observed clusters, especially in the well sampled cases. The final results show that there is good agreement in general with the results adopted as standards in the literature, but also indicates that some issues still remain unresolved. Perhaps the most important of these issues is related to the metallicity. As mentioned before, we kept the metallicity values fixed in our fits and attempted to use the ones provided by the observers, except in cases where we believed more reliable values were available. Many of the clusters show results that could be clearly improved by changing the metallicity, as for example the case of \\object{NGC 7044}. Given the considerations above, it is also important to point out that the values we used for comparison taken from the proposed standard list of \\cite{PN06} do not take metallicities into account. The different author determinations for the parameters that were averaged possibly were derived using different metallicities. Another aspect that plays a major role in the final results is the binary fraction. This is a characteristic of the clusters that is not easily accounted for even in visual fits and so it is difficult to evaluate how much this is relevant in each individual case. We have accounted for this effect by assuming a 100\\% binary fraction and drawing companions from the same IMF used to generate the fitting points. While this is clearly not the correct binary fraction for all clusters, the effect of adopting that value is relatively small, as shown in Fig. 5, due to the fact that effectively only binary systems with similar mass will show a significant difference in magnitude on the CMD. Given all the consideration above, we show that our method is reliable and robust and although the results presented in this work are consistent with literature values all clusters we believe that there is room for improvement in the accepted parameter values for other clusters. The possibility of re-evaluating previous results with more quantifiable means is important as it removes the subjectivity inherent in most open cluster studies up to today." }, "1003/1003.2416_arXiv.txt": { "abstract": "It has recently been proposed that the large-scale bias of dark matter halos depends sensitively on primordial non-Gaussianity of the local form. In this paper we point out that the strong scale dependence of the non-Gaussian halo bias imprints a distinct signature on the covariance of cluster counts. We find that using the full covariance of cluster counts results in improvements on constraints on the non-Gaussian parameter $\\fnl$ of three (one) orders of magnitude relative to cluster counts (counts + clustering variance) constraints alone. We forecast $\\fnl$ constraints for the upcoming Dark Energy Survey in the presence of uncertainties in the mass-observable relation, halo bias, and photometric redshifts. We find that the DES can yield constraints on non-Gaussianity of $\\sigma(\\fnl)\\sim 1$-$5$ even for relatively conservative assumptions regarding systematics. Excess of correlations of cluster counts on scales of hundreds of megaparsecs would represent a smoking gun signature of primordial non-Gaussianity of the local type. ", "introduction": "\\label{sec:intro} Primordial non-Gaussianity provides cosmology one of the precious few connections between primordial physics and the present-day universe. Standard inflationary theory with a single-field, slowly rolling scalar field, predicts that the spatial distribution of structures in the universe today is very nearly Gaussian random (e.g.\\ \\cite{maldacena,Acquaviva:2002ud,Creminelli:2003iq,Lyth_Rodriguez,Seery_Lidsey}; for an excellent recent review, see \\cite{Chen_AA}). Departures from Gaussianity, barring contamination from systematic errors or late-time non-Gaussianity due to secondary processes, can therefore be interpreted as violation of this ``vanilla'' inflationary assumption. Constraining or detecting primordial non-Gaussianity is therefore an important and basic test of the cosmological model. Constraints on primordial non-Gaussianity have been traditionally obtained from observations of the cosmic microwave background, as nonzero non-Gaussianity generates a non-zero three-point correlation function (or its Fourier transform, the bispectrum) of density fluctuations \\cite{Falk_Ran_Sre,Luo_Schramm,Gangui_etal,Wang_Kam,Komatsu_Spergel,Bartolo_AA,Liguori_AA}. Increasingly sophisticated algorithms have been developed to constrain non-Gaussianity, \\cite{Babich:2005en,Babich_shape,Creminelli_estimators,Smith_Zaldarriaga,Fergusson_Shellard} and, to the extent that it can be measured, Gaussianity has so far been confirmed \\cite{wmap3,Creminelli_wmap,SenSmiZal09}. For example, the most recent constraints from the Wilkinson Microwave Anisotropy Probe (WMAP) indicate $\\fnl\\approx 32\\pm 21$ ($1\\sigma$; \\cite{wmap7}), where the exact constraints depend somewhat on the choice of the statistical estimator applied to the data, the CMB map used, and details of the foreground subtraction. Here $\\fnl$ is the parameter describing non-Gaussianity in the widely studied ``local'' model, where the non-Gaussian potential $\\Phi_{\\rm NG}$ is defined by \\begin{equation} \\Phi_{\\rm NG}(x)=\\Phi_{\\rm G}(x) + \\fnl (\\Phi_{\\rm G}^2(x)-\\langle\\Phi_{\\rm G}^2\\rangle), \\label{eq:fnl_local} \\end{equation} and where $\\Phi_{\\rm G}$ is the Gaussian potential. Corresponding constraints can be obtained on other classes of non-Gaussian models. For example, for ``equilateral'' models where most power comes from equilateral triangle configurations, % $\\fnleq=26\\pm 140$ ($1\\sigma$; \\cite{wmap7}). The CMB is not the only cosmological probe to be sensitive to the presence of primordial non-Gaussianity. It has been known for a relatively long time that the abundance of dark matter halos \\cite{Lucchin:1987yv,Robinson:1999se,Benson:2001hc,MVJ,verde01,sco04,Komatsu:2003fd} (or voids \\cite{Kam_voids,songvoids}) is sensitive to the presence of primordial non-Gaussianity. This dependence is easy to understand: halos populate the high tail of the probability density distribution of structures in the universe, and the shape of this distribution is sensitive to departures from Gaussianity. However, while the halo abundance is rather powerful in constraining models that are non-Gaussian in the density (rather than the potential) \\cite{Verde_tests}, for the popular models of the local type (cf. Eq.~(\\ref{eq:fnl_local})) the abundance is much less constraining than the CMB anisotropy and not competitive with the CMB constraints (e.g.\\ \\cite{Verde_CMBLSS,Sefusatti}). Some of us \\cite{Dalal} have recently shown that the clustering of dark matter halos is very sensitive to primordial non-Gaussianity of the local type. This exciting development paves way to using the large-scale structure to probe primordial non-Gaussianity nearly three orders of magnitude more accurately than using the abundance of halos. \\citet{Dalal} found, analytically and numerically, that the bias of dark matter halos acquires strong scale dependence \\begin{equation} b(k)=b_0 + \\fnl(b_0-1)\\delta_c\\, \\frac{3\\Omega_mH_0^2}{a\\,g(a) T(k)c^2 k^2}, \\label{eq:bias} \\end{equation} where $b_0$ is the usual Gaussian bias (on large scales, where it is constant), $\\delta_c\\approx 1.686$ is the collapse threshold, $a$ is the scale factor, $\\Omega_M$ is the matter fraction relative to critical, $H_0$ is the Hubble constant, $k$ is the wavenumber, $T(k)$ is the transfer function, and $g(a)$ is the growth suppression factor\\footnote{The usual linear growth $D(a)$, normalized to be equal to $a$ in the matter-dominated epoch, is related to the suppression factor $g(a)$ via $D(a)=ag(a)$.}. This result has been confirmed by other researchers using a variety of methods, including the peak-background split \\cite{Afshordi_Tolley,MV,Slosar_etal}, perturbation theory \\cite{McDonald,Taruya08,GP}, and numerical (N-body) simulations \\cite{Grossi,Desjacques_Seljak_Iliev,PPH}. Astrophysical measurements of the scale dependence of the large-scale bias, using galaxy and quasar clustering as well as the cross-correlation between the galaxy density and CMB anisotropy, have recently been used to impose constraints on $\\fnl$ % already comparable to those from the cosmic microwave background (CMB) anisotropy \\cite{Slosar_etal, Afshordi_Tolley}, giving $\\fnl=28\\pm 23$ ($1\\sigma$), with some dependence on the assumptions made in the analysis \\cite{Slosar_etal}. In the future, constraints on $\\fnl$ are expected to be of order a few \\cite{Dalal,Carbone,Carbone2}. The sensitivity of the large-scale bias to other models of primordial non-Gaussianity has not been investigated yet (though see preliminary analyses in \\cite{Desjacques_gnl,MV09}). {\\it Clustering} of galaxy clusters, in particular, can very strongly constrain primordial non-Gaussianity. Clusters have an advantage of being large, relatively simple objects that are easy to find using either optical or X-ray light, or else from their Sunyaev-Zeldovich signature. Clusters already provide interesting constraints on dark energy \\cite{Mantz07,Vikhlinin} and they hold promise for precision measurements of cosmological and dark energy parameters (e.g.\\ \\cite{clusters_PC}). Since clusters are massive and hence significantly biased objects, their counts (via the mass function) and clustering (via the mass function and bias) are both sensitive to primordial non-Gaussianity. Recently, \\citet{Oguri09} has argued that the variance of cluster counts (i.e.\\ scatter measured in each cell individually), in combination with the cluster counts, leads to interesting improvements on $\\fnl$ constraints relative to the counts-only case. In this paper we point out that including the {\\it covariance} of cluster counts in angle and redshift leads to very significant further improvements in the cluster constraints on local primordial non-Gaussianity. The principal reason for the improvement is simple: covariance is determined by the cluster power spectrum, which is proportional to the halo bias squared. At large scales, the non-Gaussian contribution to the halo bias dominates (cf.\\ Eq.~(\\ref{eq:bias})), and this results in a strong $\\fnl$ signal in the covariance. Furthermore, we explore the sensitivity of the constraints to various assumptions about statistical and systematic errors in modeling the cluster mass-observable relation, as well as the presence of other cosmological parameters. We find that the bulk of the information about local non-Gaussianity comes from the far-separation covariances of cluster counts-in-cells. This paper is organized as follows. In Sec.~\\ref{sec:method}, we describe the methodology that we use to obtain constraints from both counts and clustering of galaxy clusters. In Sec.~\\ref{sec:calc} we describe our fiducial assumptions about the cosmological model and data as well as solutions to various challenges in calculating the constraints. In Sec.~\\ref{sec:results} we describe the forecasted constraints on $\\fnl$ from the Dark Energy Survey. We discuss our results in Sec.~\\ref{sec:disc}, and conclude in Sec.~\\ref{sec:concl}. ", "conclusions": "\\label{sec:concl} In this paper we studied how well primordial non-Gaussianity of the local type can be probed with galaxy clusters. We took into account cluster number counts, as well as the full covariance of cluster counts-in-cells. We allowed generous uncertainties in the knowledge of the cluster mass-observable relation, the photometric redshifts, and the Gaussian halo bias (we did not consider systematics due to uncertainties in angular selection, which may be important.) As we discuss at length in Sec.~\\ref{sec:calc}, the Fisher matrix calculation is computationally challenging, and we resorted to a number of conservative approximations, the most important of which is using very large pixels. Since angular selection issues are expected to be most significant at small angular scales, our pixel choices partly justify neglecting angular uncertainties. We found that most information on primordial non-Gaussianity comes from the previously neglected covariance of counts. The covariance links cluster overdensities across large distances, and thus benefits the constraints on primordial non-Gaussianity of the local type. The reason is easy to understand: the non-Gaussian parameter $\\fnl$ enters through the term proportional to $k^{-2}$ in the bias, and correlates cluster counts in bins separated by hundreds of megaparsecs. Other cosmological parameters do not lead to these far-separation correlations in cluster counts (see the right panel of Fig.~\\ref{fig:Sij}). Correlations of cluster counts across vast spatial distances of hundreds of megaparsecs therefore represent a smoking-gun signature of primordial non-Gaussianity of the local type. The combination of counts and clustering is particularly effective at breaking degeneracies of $\\fnl$ with cosmological and nuisance parameters, since the two statistical probes complement each other very well. While our full set of 23 freely varying nuisance parameters can degrade $\\fnl$ constraints by factors of a few, even modest prior uncertainties on some of them break degeneracies and restore the accuracy in $\\fnl$. For example, the bias in each photo-z bin needs to be known to 0.01 to keep $\\fnl$ constraints within $15\\%$ of their values for the case of perfectly known photo-z's. We investigated the sensitivity of our results to the choice of fiducial value of $\\fnl$ and found that the uncertainty in $\\fnl$ at $\\fnl \\neq 0$ is smaller than that for $\\fnl=0$. In other words, a non-zero small value of $\\fnl$ may even be more sensitively differentiated from the $\\fnl=0$ case than indicated in our Tables; the reason for this is explained in Sec.~\\ref{sec:disc_fiducial}. Our forecasts indicate very strong constraints on primordial non-Gaussianity, which is perhaps surprising. However, closer inspection reveals a number of effects that help clusters achieve these numbers; we discuss these in Sec.~\\ref{sec:clus_gal}. In particular, we use the pixel-based estimator, which is well suited for extracting signal from very large scales. Previous error forecasts of non-Gaussianity from galaxy clustering used the suboptimal FKP estimator; dark-energy studies that did use the pixel-based estimator only considered variance of cluster counts. To achieve the full potential of forecasted constraints discussed here, a few more issues need to be carefully studied. Particularly important are theoretical uncertainties in linking dark matter halos to observed clusters of galaxies, and observational systematics across large angular scales. While constraints on primordial non-Gaussianity have improved two orders of magnitude between COBE \\cite{Komatsu_thesis} and WMAP \\cite{wmap7}, another one or even two orders of magnitude improvement may be possible with upcoming surveys of large-scale structure, especially if they include both dark matter halo counts and their clustering covariance." }, "1003/1003.2550_arXiv.txt": { "abstract": "{} {Distances to Bok globules and small dark nebulae are important for a variety of reasons. We provide new distance estimates to several small clouds, some of them known to harbor YSO and molecular outflows, and thus being of particular interest. } { We use a procedure based on extinctions determined from the $(H-K)$ vs. $(J-H)$ diagram, and stellar distances based on a $Hipparcos$ calibration of the main sequence locus: $M_J[(J-K)_0]$. The cloud confinement on the sky is determined from contours of the average $(H-K)$ color formed in reseaus. Along the sight line stars affected by the clouds extinction may be extracted from the variation of the number density of atomic hydrogen $n_H\\sim A_{V,\\star}/D_\\star$ to provide the cloud distance and its uncertainty.} { According to our estimates, the group of three globules CB~24, CB~25 and CB~26 is located at 407$\\pm27$ pc, farther than the previous estimates. CB~245 and CB~246 are found at 272$\\pm20$ pc, suggesting that the current distance to these clouds is underestimated. Toward CB~244 we detect a layer at 149$\\pm16$ pc and the cloud at 352$\\pm18$, in good agreement with previous studies. CB~52 and CB~54, though to be at 1500 pc, are found at 421$\\pm28$ pc and slightly beyond 1000 pc, respectively. It seems that the most distant Bok globule known, CB~3, is located at about 1400 pc, also significantly closer than currently accepted. } {} ", "introduction": "Distances to small dark clouds are important for several reasons. First, they are necessary to obtain luminosities of the young stellar objects or protostars embedded in the clouds (e.g., Yun \\& Clemens 1990). In addition, distances are needed for calculating the masses and densities of the clouds (Clemens, Yun, \\& Heyer 1991). While volume densities for cloud cores can be obtained via millimeter-wavelength spectral line studies (e.g., Kane, Clemens, \\& Myers 1994), the core mass determinations require distance knowledge. Finally, accurate information about the properties of small dark clouds is needed in order to test models of these star-forming regions. For starless Bonor-Ebert spheres, e.g. Barnard 68 (Bergin, et al. 2006) the cloud distance is essential to assess its response to gravity. With the currently adopted distance of 125 pc, Barnard~68 seems to be just on the verge of instability. Given core temperature and density the distance estimate decides the stability issue. Being the most simple and regular molecular clouds in our Milky Way, Bok globules are now considered ideal laboratories for the study of the formation of low-mass stars (Vallee et al. 2000). Distances to small molecular clouds are also important in the context of studying the structure of the Milky Way star-forming fields. All Bok globules are too small and too opaque to easily apply star counts or photometric methods as distance estimators. Since most of the known globules should not be located beyond 1 kpc (Bok \\& Cordwell-McCarthy 1974), their radial velocities are dominated by peculiar motions rather than by the systematic rotational velocity field of the Galaxy. Thus, the kinematic method of distance determination can not be reliably applied. Another approach is to assume their association with larger molecular cloud complexes. Such approach has been recently used by Launhardt \\& Henning (1997), increasing the number of globules with known distances. However, in many cases this procedure is quite uncertain. There is also a significant number of globules which could not be associated with any known large molecular cloud structure and for them simply the average distance of 500 pc has been adopted (see Launhardt \\& Henning 1997). Due to these difficulties, at present only very few globules have reliably determined distances. Despite the difficulties, at present the method of photometric distance determination seems to be most reliable. Optical or infrared photometry of moderately obscured stars located at the peripheries of the globules allows us to derive individual stellar distances and thus estimate the distance to the clouds. For example Franco (1988) used \\uvbyb photometry to obtain an accurate distance to the dark cloud L1569. Recently Piehl, Briley, \\& Kaltcheva (2010) obtained \\uvbyb photometry of stars at the peripheries of CB~3, CB~52, CB~54 and CB~246 and provided new distance estimates to some of these clouds. A broadband BVI photometry of reddened M dwarfs located in front of and behind CB~24 has been used by Peterson \\& Clemens (1998) to establish a method bracketing the cloud's distance. Snell (1981) studied reddened background stars to obtain an upper limit to the distance of nine Bok globules. Maheswar \\& Bhatt (2006) obtained distances to another nine dark globules using a method based on optical and near-infrared photometry of stars projected towards the field containing the globules. Knude (2010) developed a statistical method based on the 2MASS catalog for obtaining distances to molecular clouds with a distance uncertainty of less than 10 pc, but for the small features discussed presently the uncertainty changes to a few times 10 pc. The method has been developed to be applicable to larger clouds, but it also provides rather reliable distances for small isolated globules. In this paper we demonstrate the application of this method to small-scale fields in direction of several Bok globules and obtain new estimates of their distances. ", "conclusions": "Utilizing the 2MASS catalog and applying a statistical method based on extinctions determined from the $(H-K)$ vs. $(J-H)$ diagrams, and stellar distances from a $Hipparcos$ calibration of the main sequence, we obtain new distances to several small molecular clouds. We find that the group of three globules CB~24, CB~25 and CB~26 is located at 407$\\pm27$ pc, a distance larger than the previous estimates. The collection of LDN clouds found $\\sim0.5$\\d ~below the three CB clouds is beyond 225 pc and possibly has a distribution in depth. CB~245 and CB~246 are both at 272$\\pm20$ pc. A layer at 149$\\pm16$ pc is detected in front of CB~244, which is located at 352$\\pm18$ pc. CB~52 and CB~54 are found at 421$\\pm28$ pc and slightly beyond 1000 pc, respectively. We estimate a distance about 1400 pc to CB~3, ruling out its connection to the Perseus Arm. The way we estimate the distance to a globule from 2MASS-$Hipparcos$ was develloped for more extended molecular features. It is thus encouraging that this method may possibly also be of some use to much smaller interstellar clouds when they are within $\\sim$500 pc and located in dense stellar fields. The size of the formal distance error from the curve fitting has increased but is still on the $\\la$10\\% level." }, "1003/1003.3999_arXiv.txt": { "abstract": "{The matter power spectrum as derived from large scale structure (LSS) surveys contains two important and distinct pieces of information: an overall smooth shape and the imprint of baryon acoustic oscillations (BAO). We investigate the separate impact of these two types of information on cosmological parameter estimation for current data, and show that for the simplest cosmological models, the broad-band shape information currently contained in the SDSS DR7 halo power spectrum (HPS) is by far superseded by geometric information derived from the baryonic features. An immediate corollary is that contrary to popular beliefs, the upper limit on the neutrino mass $m_\\nu$ presently derived from LSS combined with cosmic microwave background (CMB) data does not in fact arise from the possible small-scale power suppression due to neutrino free-streaming, if we limit the model framework to minimal $\\Lambda$CDM+$m_\\nu$. However, in more complicated models, such as those extended with extra light degrees of freedom and a dark energy equation of state parameter $w$ differing from $-1$, shape information becomes crucial for the resolution of parameter degeneracies. This conclusion will remain true even when data from the Planck spacecraft are combined with SDSS DR7 data. In the course of our analysis, we update both the BAO likelihood function by including an exact numerical calculation of the time of decoupling, as well as the HPS likelihood, by introducing a new dewiggling procedure that generalises the previous approach to models with an arbitrary sound horizon at decoupling. These changes allow a consistent application of the BAO and HPS data sets to a much wider class of models, including the ones considered in this work. All the cases considered here are compatible with the conservative 95\\%-bounds $\\sum m_{\\nu} < 1.16\\ {\\rm eV}$, $N_{\\rm eff}= 4.8 \\pm 2.0$.} \\begin{document} ", "introduction": "\\label{sec:introduction} Our best source of information about cosmological parameters at present is the precision measurement of the cosmic microwave background (CMB) anisotropies by the Wilkinson Microwave Anisotropy Probe (WMAP)~\\cite{Komatsu:2010fb}. However, except in the simplest models, the CMB on its own does not provide very tight constraints on specific model parameters because of parameter degeneracies. One very well known example is the bound on neutrino masses, which in the simplest vanilla+$m_\\nu$ model can be {\\it significantly} improved by adding information extracted from surveys of the large scale structure (LSS) distribution. Moreover, if the model space is extended, degeneracies with other parameters such as the dark energy equation of state parameter quickly deteriorate the neutrino mass bound from CMB data alone. In such cases it is necessary to appeal to other cosmological probes (e.g., LSS, Type Ia supernov\\ae) in order to alleviate the degeneracies. In many recent analyses (e.g., \\cite{Komatsu:2010fb}), the only information from LSS surveys employed in the parameter estimation pipeline is the length scale associated with the baryon acoustic oscillation (BAO) peak in the two-point correlation function. We call this geometric information, since a known and measured length scale (a ``standard ruler'') allows for the determination of the angular diameter distance to the object of interest simply via geometric effects. The common view is that the BAO length scale is a more robust observable than the broad-band shape of the power spectrum which may suffer from ill-understood nonlinear effects, such as nonlinear clustering or redshift- and scale-dependent galaxy/halo bias.% \\footnote{The turning point of the matter power spectrum in $k$ space corresponds to the comoving Hubble radius at the time of matter-radiation equality and in principle also constitutes a geometric measure. However, since this length scale has yet to be measured, we prefer to consider the turning point as part of the broad-band shape of the power spectrum.} Indeed, two recent studies of cosmological parameter constraints from combining CMB information with either the BAO scale from SDSS alone~\\cite{Percival:2009xn}, or including {\\it both} the broad-band shape of the SDSS DR7 halo power spectrum and BAO information~\\cite{Reid:2009xm} find very similar parameter estimates and uncertainties for the simplest vanilla model. Somewhat surprisingly, the same conclusion holds also when the vanilla model is extended with a finite neutrino mass which should in principle be very sensitive to the broad-band shape. Thus circumstantial evidence seems to suggest that in the simplest cosmological models, geometric information from the BAO wiggles supersedes the information contained in the broad-band shape of the matter power spectrum. However, as we shall demonstrate, for more complicated models the additional information contained in the broad-band shape of the matter power spectrum can make a very substantial difference to parameter inference . For example, in models where the number of relativistic degrees of freedom $N_{\\rm eff}$ and the dark energy equation of state parameter $w$ are added as free parameters, the difference in the parameter uncertainties between including and excluding the power spectrum shape information can be a factor of two or more (see also Ref.~\\cite{Biswas:2009ej} for a related discussion in the context of CMB data and dark energy models). The purpose of the present work is to clarify the roles of ``geometric'' and ``shape'' information extractable from the current generation of LSS surveys, and to stress that in extended models geometric information alone does not optimally exploit the available data. This will remain true even when CMB data from the Planck spacecraft become available. The paper is organised as follows: In section~\\ref{sec:analysis} we describe the specific cosmological parameter space used, the data sets, as well as the analysis method. We present our results in section~\\ref{sec:results}, including a forecast for Planck. Our conclusions can be found in section~\\ref{sec:conclusions}. Appendices~\\ref{sec:appendix_hps} and \\ref{sec:appendix_bao} contain details of our methodology. ", "conclusions": "\\label{sec:conclusions} We have studied in detail how various subsets of the power spectrum information from a large scale structure survey can be used to constrain cosmological parameters. At present an often used approach is to restrict the LSS information to the geometric distance information contained in the BAO peak. For the minimal $\\Lambda$CDM model (with neutrino mass included) this does indeed provide exactly the same constraint as the use of the full power spectrum data, and is far superior to using a smoothed no-wiggle power spectrum which contains only shape information. This indicates that the neutrino mass is currently more strongly constrained by its effect on the background evolution~\\cite{Lesgourgues:2006nd}, and the contribution of present LSS data consists mainly in alleviating the geometrical degeneracy with $h$ and $\\Omega_{\\rm m}$ \\cite{Komatsu:2008hk,Thomas:2009ae}, rather than constraining the possible small-scale power suppression in the large scale matter power spectrum due to free-streaming. However, this simple picture changes when more complex cosmological models are studied.\\footnote{We should point out a small caveat here: the usefulness of the LSS shape information depends not only on the cosmological model under consideration, but also on the combination of data sets used in the analysis. For example, Reid et al.~\\cite{Reid:2009xm} find that in a vanilla+$N_{\\rm eff}$ model, WMAP5+HPS yields much better constraints on $N_{\\rm eff}$ than WMAP5+BAO. However, in this model the LSS shape information loses its usefulness as soon as one adds the HST constraint, which breaks the ($H_0$-$N_{\\rm eff}$)-degeneracy more efficiently -- leading to conclusions similar to those found in our subsection~\\ref{subsec:v+f+n}.} As an example we have tested a model with a variable number of neutrino species, and a dark energy equation of state, $w$, different from -1. In this model some parameters are still as well constrained by BAO alone as by the full power spectrum. This is true for example for the number of neutrino species, which is mainly probed by CMB data, and the dark matter density which is highly sensitive to the position of the BAO peak. However, for other parameters such as $\\sum m_\\nu$, $w$ and $n_{\\rm S}$, there is additional information in the shape of the power spectrum which is crucial for constraining these parameters. For example the upper 95\\% bound on neutrino mass goes from $1.47\\ {\\rm eV}$ to $1.16\\ {\\rm eV}$ when BAO information is replaced with the full halo power spectrum. However, even in this model, the entire shape information of the HPS is contained in the data points at wavenumbers smaller than $0.12\\ h {\\rm Mpc}^{-1}$, the higher-$k$ information being diluted due to uncertainties in nonlinear modelling. In other words, due to our ignorance of the processes governing the power spectrum at smaller scales, we basically lose almost half of the available data points (the half that is less subject to sample variance at that!). Clearly, a better understanding of the mildly nonlinear physics at these scales would be highly desirable. Needless to say, our conclusions do not alter the fact that in the future, when better data from galaxy, cluster, weak lensing or 21cm surveys will be available, the best way to probe the neutrino mass will be through the information contained in the shape and the scale-dependent growth factor of the large scale structure power spectrum~\\cite{Song:2004tg,Lesgourgues:2004ps,Wang:2005vr,Lesgourgues:2006nd,Hannestad:2007cp,Vikhlinin:2008ym,Pritchard:2009zz}. In this paper we have also presented a new method for separating the geometric BAO information from the shape information in the no-wiggle spectrum for more complex models than previously studied. The method is based on removing the oscillating part of the power spectrum by use of a fast sine transform and then removing the BAO peak by smoothing the resulting ``correlation function''. Finally, the smoothed function is transformed back to provide the no-wiggle power spectrum. The method has been demonstrated to be extremely fast and robust to even radical changes in the cosmological model, making it easy and safe to implement in \\texttt{CosmoMC}. Along the same lines we have also implemented a version of the SDSS BAO likelihood code which allows for models in which the sound horizon is modified compared to the $\\Lambda$CDM model with $N_{\\rm eff}=3$. In addition to constraints using current data we have also performed an estimate of how the SDSS measurements can be used to improve the Planck constraints on some parameters in extended models. Most parameters will be so well determined by Planck that little can be gained from adding the SDSS data in any form. However, with $\\sum m_\\nu$ and $w$ the situation is different. With these parameters the SDSS data can lead to very significant improvements in sensitivity. Furthermore, we have also shown that even for Planck the SDSS halo power spectrum contains important information beyond what is in the geometric BAO data - for both $\\sum m_\\nu$ and $w$ the shape information can improve the sensitivity by 30-40\\%. As shown in Ref.~\\cite{Rassat:2008ja}, in future large scale structure surveys the relative impact of the shape information is expected to increase, so extracting the full power spectrum information and at the same time improving the theoretical modeling of small-scale perturbations remains a crucial goal." }, "1003/1003.1214_arXiv.txt": { "abstract": "{Many sources listed in the 4$^{\\rm th}$ IBIS/ISGRI survey are still unidentified, i.e. lacking an X-ray counterpart or simply not studied at lower energies ($<$ 10 keV). The cross-correlation between the list of IBIS sources in the 4$^{\\rm th}$ catalogue and the \\emph{Swift}/XRT data archive is of key importance to search for the X-ray counterparts; in fact, the positional accuracy of few arcseconds obtained with XRT allows us to perform more efficient and reliable follow-up observations at other wavelengths (optical, UV, radio). In this work, we present the results of the XRT observations for four new gamma-ray sources: IGR J12123--5802, IGR J1248.2--5828, IGR J13107--5626 and IGR J14080--3023. For IGR J12123--5802 we find a likely counterpart, but further information are needed to classified this object, IGR J1248.2--5828 is found to be a Seyfert 1.9, for IGR J13107--5626 we suggest a possible AGN nature, while IGR J14080--3023 is classified as a Seyfert 1.5 galaxy.} \\FullConference{The Extreme sky: Sampling the Universe above 10 keV - extremesky2009,\\\\ October 13-17, 2009\\\\ Otranto (Lecce) Italy} \\begin{document} \\begin{table*} \\begin{center} \\footnotesize \\caption{XRT position and classification of the counterpart of the IBIS sources.} \\label{Tab1} \\begin{tabular}{lccccc} \\hline \\hline Source & R.A. & Dec & Error & Counterpart & Type \\\\ & (J2000)& (J2000) & (arcsec) & & \\\\ \\hline \\hline IGR J12123--5802 & $12^{\\rm h}12^{\\rm m}25^{\\rm s}.97$ & $-58^\\circ00^{\\prime}23^{\\prime \\prime}.1$ & 3.7 & 2MASS J12122623--5800204 & unidentified \\\\ \\hline IGR J1248.2--5828 & $12^{\\rm h}47^{\\rm m}57^{\\rm s}.82$ & $-58^\\circ29^{\\prime}59^{\\prime \\prime}.1$ & 4.0 & 2MASX J12475784--5829599 & Seyfert 1.9 \\\\ \\hline IGR J13107--5626 & $13^{\\rm h}10^{\\rm m}37^{\\rm s}.27$ & $-56^\\circ26^{\\prime}56^{\\prime \\prime}.7$ & 4.4 & 2MASX J13103701--5626551 & AGN candidate \\\\ \\hline IGR J14080--3023 & $14^{\\rm h}08^{\\rm m}06^{\\rm s}.57$ & $-30^\\circ23^{\\prime}52^{\\prime \\prime}.6$ & 3.6 & 2MASX J14080674--3023537 & Seyfert 1 \\\\ \\hline \\hline \\end{tabular} \\end{center} \\end{table*} ", "introduction": " ", "conclusions": "In this work, we show how follow-up observations in X-rays are of key importance to search for counterparts of high energy emitters. The cross-correlation between the 4$^{\\rm th}$ IBIS catalogue and the \\emph{Swift}/XRT data archive allowed us to pinpoint unambiguously the counterpart of four IBIS sources. IGR J1248.2--5828 is classified as a Seyfert 1.9 galaxy, while the X-ray properties of IGR J14080--3023 are compatible with its classification as a Seyfert 1.5 galaxy; based on its X-ray characteristics, IGR J13107--5626 is likely an AGN, while for IGR J12123--5802 the data available so far do not allow us to assess its nature and only optical measurements are needed to firmly establish the nature of this new gamma-ray source." }, "1003/1003.2217_arXiv.txt": { "abstract": "We present photometric and spectroscopic observations of SN~2007if, an overluminous ($M_V = -20.4$), red ($B-V = 0.16$ at $B$-band maximum), slow-rising ($t_\\mathrm{rise} = 24$~days) type Ia supernova (SN Ia) in a very faint ($M_g = -14.10$) host galaxy. A spectrum at 5~days past $B$-band maximum light is a direct match to the super-Chandrasekhar-mass candidate SN~Ia 2003fg, showing \\ion{Si}{2} and \\ion{C}{2} at $\\sim 9000$~km~s$^{-1}$. A high signal-to-noise co-addition of the SN spectral time series reveals no \\ion{Na}{1}~D absorption, suggesting negligible reddening in the host galaxy, and the late-time color evolution has the same slope as the Lira relation for normal SNe~Ia. The ejecta appear to be well mixed, with no strong maximum in $I$-band and a diversity of iron-peak lines appearing in near-maximum-light spectra. SN~2007if also displays a plateau in the \\ion{Si}{2} velocity extending as late as $+10$~days, which we interpret as evidence for an overdense shell in the SN~ejecta. We calculate the bolometric light curve of the SN and use it and the \\ion{Si}{2} velocity evolution to constrain the mass of the shell and the underlying SN ejecta, and demonstrate that SN~2007if is strongly inconsistent with a Chandrasekhar-mass scenario. Within the context of a ``tamped detonation'' model appropriate for double-degenerate mergers, and assuming no host extinction, we estimate the total mass of the system to be $2.4 \\pm 0.2$ \\Msol, with $1.6 \\pm 0.1$ \\Msol\\ of \\nickel\\ and with 0.3--0.5 \\Msol\\ in the form of an envelope of unburned carbon/oxygen. Our modeling demonstrates that the kinematics of shell entrainment provide a more efficient mechanism than incomplete nuclear burning for producing the low velocities typical of super-Chandrasekhar-mass SNe~Ia. ", "introduction": "Type Ia supernovae (SNe~Ia) are of vital importance as luminosity distance indicators for measuring the expansion history of the Universe \\citep{riess98,scp99}. % They have a small dispersion ($\\sim 0.35$ mag) in intrinsic peak luminosity, which can be further reduced to 0.16--0.18 mag by applying a well-known correction dependent on the width, or decay rate, of the light curve \\citep{phillips99,salt2,mlcs2k2}. Searches for other luminosity correlates with which to derive even more accurate luminosity distances from SNe~Ia are underway, with some methods delivering core Hubble residual dispersions as low as 0.12 mag \\citep{sjb09,wang09,csp09}. SNe~Ia are generally understood to result from the thermonuclear explosion of at least one carbon/oxygen white dwarf. However, the underlying distribution of SN~Ia progenitor systems and explosions mechanisms, and the relative rates of possible different physical subclasses of SN~Ia events, remain poorly constrained, with potential consequences for the average luminosity of events. Any systematic effect which may influence the luminosity of different SN~Ia subpopulations at the level of a few percent has become a cause for concern for next-generation experiments, particularly redshift-dependent effects \\citep{kim04,linder04}. A better understanding of the progenitor systems will place these corrections on a much firmer conceptual footing and place limits on hitherto uncontrolled astrophysical systematics in SN~Ia luminosity distance measurements. Rare SN~Ia events displaying extreme characteristics or evidence of unusual physics can often point the way toward other, less extreme instances of similar physics which may be lurking in the otherwise undifferentiated sample of ``normal'' events. For example, one commonly-invoked rationale for the uniformity in pre-correction SN~Ia luminosities is that they start with the same amount of fuel, and are triggered by the same physics: the SN~Ia progenitor explodes when its mass nears the Chandrasekhar limit, $\\mch = 1.4\\ \\Msol$, its mass is completely unbound and converted mostly to heavier elements, especially \\nickel, the decay of which powers the SN~Ia light curve. The \\emph{single-degenerate} scenario \\citep[SD;][]{wi73} ensures that the white dwarf slowly approaches \\mch\\ via accretion from a non-degenerate companion. In contrast, the \\emph{double-degenerate} scenario \\citep[DD;][]{it84}, in which two white dwarfs in a binary system merge and explode, provides a way for SN~Ia progenitors to exceed \\mch\\ and to give rise to more luminous events. There are also some arguments that single-degenerate, differentially rotating white dwarfs with mass exceeding \\mch\\ significantly can exist \\citep{yl05}, although the inclusion of baroclinic and magnetohydrodynamic instabilities appears to preclude this \\cite[e.g.,][]{piro08}. There may therefore be a population of SNe~Ia with a distribution of masses greater than \\mch, with different explosion physics that interferes with luminosity standardization. The relative rate of such events among SNe~Ia in general may also depend on redshift, and unless they can be identified or their luminosities accurately calibrated, they need not be common to produce significant biases in reconstructions of the dark energy equation of state. There are at least four documented examples of overluminous SN~Ia explosions with progenitor mass potentially exceeding \\mch. The first known example of the class was SN~2003fg \\citep[\\champagne;][]{howell06}; SN~2006gz \\citep{hicken07}, SN~2007if \\citep{cbet2007if}, and SN~2009dc \\citep{tanaka09,yamanaka09,silverman10} were discovered later as events spectroscopically similar to SN 2003fg. The main evidence cited for a very massive progenitor in each of these cases was the extremely high luminosity of each of these events, and by inference unusually large \\nickel\\ synthesis. SNe~2003fg, and later SNe 2006gz and 2009dc, were noted for being overluminous ($M_V \\sim -20$), with unusually wide light curves (``stretch'' $s > 1.1$) and \\ion{C}{2} lines, evidence for unburned carbon, in early or near-maximum spectra. In SN~2003fg, narrow, low-velocity ($8000$--$9000$~km~s$^{-1}$) \\ion{Si}{2} lines near maximum light have been interpreted as evidence for a high gravitational binding energy, further supporting the hypothesis of a very massive progenitor; low velocities were also found in SN~2009dc. Ejecta velocities inferred from spectra of SN~2006gz were closer to those of normal SNe~Ia. SN~2007if was discovered by the Texas Supernova Search \\citep{cbet2007if} in unfiltered ROTSE-IIIb images taken on 2007~August~16.3 UT. It was found independently as SNF20070825-001 by the Nearby Supernova Factory \\citep[SNfactory;][]{snf}, in an image taken in a red (RG610) filter with the QUEST-II camera \\citep{baltay07} on the Palomar Observatory Oschin 1.2~m Schmidt telescope (``Palomar-QUEST'') on 2007~August~25.4. No host galaxy was visible in the discovery images; nor in any available sky survey images, including the Sloan Digital Sky Survey (SDSS), POSS and USNO, making the redshift determination and interpretation of early-phase spectra uncertain. An optical spectrum we obtained with the SuperNova Integral Field Spectrograph \\citep[SNIFS;][]{snifs} on the University of Hawaii 2.2~m telescope on 2007~August~26.5 UT revealed a blue continuum not obviously like a type~Ia supernova. Spectroscopy taken at the Hobby-Eberly Telescope on 2007~August~29 also failed to identify the nature of the event \\citep{cbet2007if}. However, a later spectrum we obtained with the Double Spectrograph on the Hale 5-m telescope at Palomar, on 2007~September~6.5 UT, identified SN~2007if as a SN~Ia, apparently well before peak. A further SNIFS spectrum taken on September~10.5 UT turned out to be an unambiguous match to a published spectrum of SN~2003fg \\citep{howell06}. Cross-correlation of the September~10.5 SNIFS spectrum with the SNLS spectrum of SN~2003fg suggested a redshift of $0.070 \\pm 0.005$. The faintness of the host, coupled with the unusually large luminosity of the supernova, presented a challenge for the detection of line emission from the host as late as a full year after explosion. More recently, however, an optical spectrum of the host galaxy was obtained on 2009~August~24.5 with the Low Resolution Imaging Spectrograph (LRIS) at the Keck Telescope on Mauna Kea, showing H$\\alpha$ and \\ion{O}{3}~$\\lambda 3727$ at a heliocentric redshift of $0.07416 \\pm 0.00082$. This new redshift measurement allows more accurate determination of the SN luminosity and the ejecta velocity scale, which in turn enables a measurement of the total mass in the explosion. In Sections 2 and 3 we present our detailed photometric and spectroscopic observations of SN~2007if and its host galaxy. In Section 4 we present the bolometric light curve of the SN and an estimate of the mass of \\nickel\\ synthesized in the explosion, which we find to nominally exceed \\mch. In Section 5 we argue that the red color of the SN, its unusually long (24-day) rise time, and the existence of a plateau in the inferred photospheric velocity are best explained by the existence of an overdense shell in the ejecta, probably caused by the entrainment of an unburned carbon-oxygen envelope. We also estimate the total mass ejected in the SN, and the fraction of that mass residing in the shell and envelope. Since SN~Ia mass estimates are often sensitive to the assumed kinetic energy of the explosion, we consider in Section 6 the importance of shell structure on arguments associated with mass estimates in the literature, and ask whether shell structure may be more common in super-Chandrasekhar-mass SN~Ia candidates than previously believed. We summarize and conclude in Section 7. ", "conclusions": "" }, "1003/1003.2484_arXiv.txt": { "abstract": "We have modeled the displacement of luminous X-ray binaries from star clusters in star-burst galaxies with an evolutionary population synthesis code developed by \\citet{Hurley00,Hurley02}. In agreement with \\citet {kaaret04}, we find significant spatial offset of X-ray sources from their parent clusters, and the apparent X-ray luminosity vs. displacement correlation can be roughly reconstructed. The correlation is not sensitive to the fundamental properties of the clusters (e.g., initial mass functions of the binary stars) and the kick velocity imparted to the newborn compact stars, except the common envelope parameter $\\alpha_{\\rm CE}$. We present the distributions of the main parameters of the current X-ray binaries, which may be used to constrain the models for the formation and evolution of X-ray binaries with future optical observations. ", "introduction": "X-ray binaries (XRBs) contribute a significant fraction of the X-ray radiation of normal galaxies \\citep{fabb89}. They are binary systems containing an accreting neutron star (NS) or black hole (BH) and a normal companion star. Based on the masses $M_{\\rm op}$ of the optical companions, XRBs are conventionally divided into high-mass X-ray binaries (HMXBs) and low-mass X-ray binaries (LMXBs) \\citep[e.g.][]{verbunt94}. In HMXBs, the massive ($M_{\\rm op}\\ga 10M_{\\odot}$) companions generally have strong stellar winds (with mass loss rate $\\sim 10^{-8}-10^{-5} M_{\\odot}$yr$^{-1}$), part of which can be captured by the compact star; while LMXBs, in which $M_{\\rm op}\\la 1.5 M_{\\odot}$, experience mass transfer through Roche-lobe overflow (RLOF) of the companion, at a rate of $\\sim 10^{-10}-10^{-8} M_{\\odot}$ yr$^{-1}$. Between them are intermediate-mass X-ray binaries (IMXBs), in which the companion stars' masses are in the range $\\sim 2 - 10 M_{\\odot}$ \\citep{heuvel75}. Mass transfer in these binaries also occurs through RLOF, but on much faster, (sub)thermal timescale of $\\sim 10^4-10^5$ yr. XRBs in galactic disks are thought to have evolved from primordial binaries, in which a high-mass primary star ($M\\ga 10M_{\\odot}$) formed the compact star and the secondary star as its companion. The formation and evolution of XRBs are often accompanied with mass transfer and loss of mass and orbital angular momentum \\citep[][for a review]{tauris06}. If the primary star evolves to be a (super)giant and fills its Roche lobe (RL), its mass is transferred to the secondary via RLOF. If the mass ratio of the primary and secondary stars is sufficiently high or the primary star has a deep convective envelope, the mass transfer process occurs on a dynamical timescale and is highly unstable, so that a common-envelope (CE) enshrouding binary results. The secondary star is captured by the expansion of the giant star and is forced to move through the giant's envelope. The resulting frictional drag will cause its orbit to shrink rapidly while, at the same time, ejecting the envelope before the naked core of the giant star explodes to form the NS/BH. The binary, if it survives the supernova (SN), will evolve to be an XRB when the secondary begins to transfer mass to the compact star. Note that during the formation and evolution processes of XRBs there may be several instances of mass transfer and CE phases. For example, HMXBs may end up in a CE phase, as the NS (or BH) is engulfed by the extended envelope of its companion. If the system survives after the CE, an XRB with a Helium companion may be produced. The formation of LMXBs in globular clusters often invokes dynamical process such as tidal capture of a low-mass main sequence (MS) star by a NS \\citep{bailyn87} and exchange encounters between an NS and a primordial binary \\citep{davies98}. Young star clusters and X-ray sources from them have many interesting aspects on modern astrophysics. Recent data shows that compact young massive clusters contain a rich population of massive stars, evidently following a standard Salpeter-like upper initial mass function \\citep[IMF;][]{massey98}. However, there are also indications that some compact young massive clusters have either a flatter than normal upper mass function or a cut-off at low mass \\citep{sternberg98,smith01,mcCrady03}. Additionally, observational studies on the mass ratio of binaries have reached widely varying, even disparate conclusions \\citep[reviewed in][]{abt83,larson01}. Recent studies find that there are two populations of secondaries in clusters, which leads to a bimodal distribution of mass ratios \\citep[for more, see \\S1.2 in][and references therein]{kobulnicky07}. For example, \\citet{lucy06}, using data from the Ninth Catalogue of Spectroscopic Binary Orbits \\citep{pourbaix04}, reassessed the \\citet{hogeveen90} study and concluded that the data support an excess of $q\\simeq 1$ twin systems. These cases all reveal that potentially more massive binaries will be produced in these young clusters, which may present different observational properties when they turn on X-rays. So investigations on XRBs in young star clusters may help explore, besides the formation and evolution of XRBs, the fundamental properties of the clusters as well as recent star formation processes in galaxies, including stellar population in galaxies \\citep{liu07}, massive star formation and evolution \\citep{Kaper07}. Using observations from {\\it Chandra\\/} and NICMOS on board {\\it Hubble Space Telescope\\/} ({\\it HST\\/}), \\citet{kaaret04} examined the spatial offsets between X-ray point sources and star clusters in three star-burst galaxies. They found that (1) the X-ray sources are preferentially located near the star clusters, indicating that the X-ray sources are young objects associated with current star formation, and (2) brighter X-ray sources preferentially occur closer to clusters. The displacements of the X-ray sources observed in these starburst galaxies are likely due to the motion of the X-ray sources caused by the SN explosions and/or dynamical interactions with other stars and binaries in the clusters. The absence of bright X-ray sources with large displacements suggests that there is some correlation between the luminosity of an X-ray source and its motion when the X-ray luminosity $L_{\\rm X}>10^{38}$ ergs$^{-1}$. They proposed that these high luminosity sources may be RLOF BH-XRBs with intermediate mass companions, if emitting isotropically and running away at a speed of $v\\sim$ 10 kms$^{-1}$. If $v\\sim$ 50 kms$^{-1}$ instead, the short lifetime ($\\sim$ 4 Myr) of these luminous sources needs very massive companion stars or alternatively the luminosity of the sources decreases with age, which can be examined by an evolutionary population synthesis (EPS) calculation. They also pointed out that the correlation appears inconsistent with the highly beamed X-ray emission model \\citep{king01,kording02,kaaret03} because the delay time between the formation of a BH and the onset of the thermal time-scale mass transfer is long enough for the X-ray source to move to $\\sim 1$ kpc from its point of origin. The spatial offset between the X-ray point sources and the star clusters, especially the X-ray luminosity versus displacement correlation is determined by the velocity of the binary after the birth of the NS/BH, the time passed since the SN, and the mass transfer process. Spatial distribution of X-ray sources in galactic environment has been investigated both observationally and theoretically. For example, \\citet{paradijs95} and \\citet{white96} have investigated the spatial distribution of NS and BH LMXBs in our Galaxy, from which they suggested that the compact objects had received a kick during the SN explosion. \\citet{paczynski90} also studied the spatial distribution of Galactic NSs in the scope of SN kicks. Recently \\citet{zuo08} modeled the spatial distribution of Galactic XRBs, incorporating the kinematic evolution of kicked binary systems. \\citet {kiel09} calculated the scale-heights as well as the radial and space velocity distributions of pulsars, considering their kinematic evolution within the Galactic potential. However, for XRBs born in star clusters, besides mass ejection from the binary \\citep{nelemans99,heuvel00} and the kick velocity imparted on the newborn NS caused by the SN explosion \\citep{lyne94}, an ejection speed via dynamical interactions in clusters \\citep{phinney91,kulkarni93,sigurdsson93} should also be considered. In our calculation we only consider the former two mechanisms on the motion of XRBs \\citep[i.e., primordial binaries;][]{webbink83,webbink92} though the third one can also influence the binary population, but not significantly in our situation, as explained below. Dynamical interactions include tidal capture, physical collisions and exchange encounters. The formation rate of NS-XRBs produced through tidal capture process using Eq.~(3) from \\citet{verbunt87} can be estimated as: \\begin{equation} R=n_{\\rm ns}nv_{\\rm rel}\\sigma\\simeq 6\\times10^{-11}\\frac{n_{\\rm ns}}{10^2\\rm pc^{-3}}\\frac{n}{10^4\\rm pc^{-3}} \\frac{M_1+M_2}{M_{\\odot}}\\frac{d}{R_{\\odot}}\\frac{10\\rm kms^{-1}}{v_{\\rm rel}}\\rm yr^{-1}\\rm pc^{-3} \\end{equation} where $n_{\\rm ns}$ and $n$ are the number densities of NSs and other stars, respectively, $v_{\\rm rel}$ is the relative velocity between the stars at infinity, $\\sigma\\simeq\\pi d[2G(M_1+M_2)/v_{\\rm rel}^2]$ \\citep[i.e., Eq.~(2) in][]{verbunt87} is the cross section of the two passengers with small relative velocities, $M_1$ the NS mass, $M_2$ the mass of the companion star, and $d$ is the distance of closest approach of the two passengers. Here we adopt $(M_1+M_2)\\sim 2M_{\\odot}$ and $d_{\\rm max}\\sim$ 10$R_{\\odot}$ which is estimated roughly from Eq.~(1) in \\citet[][]{verbunt87}. Given the typical values of young massive clusters as $v_{\\rm rel} \\sim 10$ kms$^{-1}$, the core radius $r_{\\rm c} \\sim 1$ pc, the central density $\\rho_0 \\sim 10^4 M_{\\odot}$ pc$^{-3}$ \\citep{mcCrady03}, we can get $n\\sim 2\\times 10^4$ pc$^{-3}$, $n_{\\rm ns}\\sim 0.5f\\times 10^2$ pc$^{-3}$ assuming an IMF of \\citet{Kroupa}, following the approach (i.e., \\S 3.1) of \\citet[][]{verbunt87}, here $f\\sim 0.2$ is the fraction of NSs remaining in the cluster. Then we can estimate the predicted number of XRBs in one cluster is $N_{\\rm X}\\simeq\\frac{4}{3}\\pi r_{\\rm c}^3R\\times T\\sim 10^{-2}$ with an X-ray lifetime $T= 10$ Myr. So the expected number of XRBs produced through this channel in the \\citet{kaaret04} sample is very small ($\\ll 1$). Exchange encounters can also occur on time-scale of a few Myr \\citep{portegies99}, although the formation rate of XRBs through exchange collisions barely competes with two-body encounters \\citep{hut83}. However, during the exchange encounters a lower mass binary star tend to be replaced by a more massive participant. Hence the effect of exchange encounters is to modify the mass-ratio distribution of the binaries, making equal-mass binaries more likely. We examined this effect in our models (i.e., models M3 and M7) and found no significant changes. In the present work, we investigated the kinematic consequences of XRBs from star clusters in star-burst galaxies from a theoretical point of view. We used an EPS code to calculate the expected cumulative distribution of XRB displacements from their parent clusters. Following the approach of \\citet{zuo08}, we calculated the spatial offset distribution of XRBs with luminosities $>10^{36}$ ergs$^{-1}$. We mainly examined several parameters, such as IMF, the secondary star mass function, common envelop efficiency and kick velocity, which may affect the formation, evolution and motion of massive XRBs significantly. The objective of this study is to use the apparent X-ray luminosity versus displacement correlation to constrain the model parameters and the fundamental properties of clusters in star-burst galaxies. We also aim to explore why such correlation exists, which may help understand the nature of the sources and may be testified by future observations. This paper is organized as follows. In \\S 2 we describe the population synthesis method and the input physics for XRBs in our model. The calculated results are presented in \\S 3. Our discussion and conclusions are in \\S 4. ", "conclusions": "We have used an EPS code to calculate the spatial offset distribution of XRBs from their parent star cluster in star-burst galaxies. We used the apparent X-ray luminosity versus displacement correlation to constrain models of XRBs. Our study shows that the correlation can be roughly reproduced with all models considered, but significant differences exist when changing the common envelope parameter $\\alpha_{\\rm CE}$. In the $\\alpha_{\\rm CE}=1.0$ cases (models M1-M4), the $L_{\\rm X}$ vs. $R$ relation is constructed by both high-luminosity ($L_X>10^{38}$ ergs$^{-1}$), small-offset ($1010^{38}$ ergs$^{-1}$) sources in the $10$ 80$\\%$) NS-XRBs with HeMS companions (from $\\sim$ 50$\\%$ in model M4 to $\\sim$ 85$\\%$ in model M7), transferring mass through RLOF or wind-capture. The orbital periods are $\\sim 1-20$ hr for models M1 and M3-M8, and can reach $\\sim 10^3$ hr for model M2. The companion masses are $\\sim 1-4 M_{\\odot}$ for models M1 and M4, $\\sim 1-20 M_{\\odot}$ for models M5-M8 and can reach $\\sim 60 M_{\\odot}$ for models M2 and M3. (3) The XRBs in the $30010 M_{\\odot}$), the value of $\\lambda$ may be low as 0.1. So incorporating $\\lambda$ as a function of stellar radius will help model the $L_{\\rm X}$ vs. $R$ relation more precisely, and put more realistic constraints on the model parameters, though it is beyond the scope of this paper." }, "1003/1003.0481_arXiv.txt": { "abstract": "This article presents the first computation of the complete bispectrum of the cosmic microwave background temperature anisotropies arising from the evolution of all cosmic fluids up to second order, including neutrinos. Gravitational couplings, electron density fluctuations and the second order Boltzmann equation are fully taken into account. Comparison to limiting cases that appeared previously in the literature are provided. These are regimes for which analytical insights can be given. The final results are expressed in terms of equivalent $\\fNL$ for different configurations. It is found that for moments up to $\\ell_{\\rm max}=2000$, the signal generated by non-linear effects is equivalent to $\\fNL\\simeq 5$ for both local-type and equilateral-type primordial non-Gaussianity. ", "introduction": "The Cosmic Microwave Background (CMB) anisotropies are now observed with a high precision and have become a key observation of modern cosmology. They are in particular very precious to constrain the theories of the primordial Universe~\\cite{WMAP7}. So far the temperature anisotropies have been found to have statistical properties that are compatible with Gaussian statistics~\\cite{WMAP7}. The CMB data can therefore entirely be captured in its power spectrum and the latter has been used to set constraints on the cosmological parameters and on the shape of the inflationary potential. There is hope however that future observations such as Planck~\\cite{PLANCK}, that will provide better data on the statistical properties of the temperature and polarization fields, open a new window on the physics of the early Universe with the use of higher-order statistical properties of the CMB sky such as its bispectrum. In the analysis of those data, one should keep in mind though that the properties of the CMB temperature and polarization anisotropies depend both on the properties of the initial conditions and on their evolution. As long as measurements are restricted to second-order statistics such as the angular correlation function or the power spectra, a linear perturbation theory suffices, for the required precision, to relate the (2-dimensional) angular power spectrum to the (3-dimensional) initial power spectrum of the metric perturbations at the end of inflation. It has thus been understood that the characteristic features observed in the CMB temperature spectrum originate from the developments of acoustic oscillations encoded in the linear transfer~\\cite{Hu1994b,Hu1995} function while its overall amplitude and its scale dependence are fixed by the initial power spectrum, the shape of which agrees with the predictions from an inflationary era~\\cite{Mukhanov1992,Linde2007}. At this level of description, all aspects are now fully understood and is part of textbooks~\\cite{2003moco.book.....D,2005pfc..book.....M,FrancisLivre,PeterUzanTrans,2008cmb..book.....D,2008cosm.book.....W} and since the metric and matter perturbations are linearized, any model that predicts Gaussian initial conditions, as standard single field inflation, is expected to produce Gaussian statistical properties for the CMB temperature field. In general bispectra arise whenever non-linear mode couplings are at play during the cosmological evolution and general relativity being in essence a non-linear field theory, deviations from Gaussianity are expected to be ubiquitous, arising either from the inflationary era (and thus called primordial non-Gaussianity) or from the post-inflationary evolution. Gravity mediated couplings are however generally small and in full agreement with the current data that clearly favors only mild non-Gaussianities if any~\\cite{WMAP7}. In particular, it is now widely accepted that standard single field inflation cannot produce significant non-Gaussianities since its amplitude is mostly dictated by gravity induced couplings~\\cite{Maldacena2003}. On the other hand, significant deviations from Gaussianity can arise from non-standard kinetic terms, for which models based on the Dirac-Born-Infeld action are typical models \\cite{2004PhRvD..70l3505A} or in the context of of multiple-field inflation specially when non-gravity type couplings are at play~\\cite{Bernardeau2002,Komatsu2002,Bartolo2002,Bernardeau2003,Bernardeau2006,Rigopoulos2006,2008PThPh.120..159S,2009PhRvL.103g1301B}. It is therefore generally admitted that non-Gaussianity searches can open a window on the details of the inflationary mechanism at work in the early Universe (see Refs.~\\cite{Bartolo2004,2009astro2010S.158K,Bartolo2010revue} for general reviews). The level with which primordial non-Gaussianity could actually be detected is however still largely debated, mainly because this source of non-Gaussianity is in ``competition'' with the couplings induced by non-linear effects throughout the whole recombination and photon propagation processes. This has motivated a series of general studies aiming at characterizing the bispectrum~\\cite{Bartolo2004b,Pyne1996,Mollerach1997,Goldberg1999,Komatsu2001,Babich2004,Bartolo2004a,Komatsu2005,Liguori2006,Bartolo2008,Khatri2008,Senatore2008b,Nitta2009,Hanson2009,Boubekeur2009,Mangilli2009,Liguori2010} and even the trispectrum~\\cite{Komatsu2002,Okamoto2002,Kogo2006} to be expected from the observation made by CMB experiments. While the identification of the mode couplings (at the quantum level) during the inflationary phase has been set on secure grounds~\\cite{Maldacena2003,Bernardeau2004a,Weinberg2005,Weinberg2006,2008PhRvD..78f3534W}, the evolution of the cosmological perturbations during the post-inflationary era is still largely unexplored, the primordial statistical properties being often related to the observed statistical properties only through a linear transfer function. However, in order to relate the angular bispectrum of the observed CMB to the spatial bispectrum of the metric perturbations at the end of inflation, one needs to derive transfer functions up to second order that incorporate all type of couplings. The work we present in this article is the end result of a task which was initiated in Ref.~\\cite{PUB2008}, that was followed by some partial reanalysis~\\cite{Bartolo2008,Senatore2008a,Khatri:2008kb}, In this article we described what we thought is the main mechanism at play for the generation of the bispectra at small angular scales and for which we could give some physical insights. The present article extends this analysis using the theoretical developments of the non-linear perturbation theory described in Refs.~\\cite{Nakamura2007,Pitrou2008,Pitrou2008letter}, and in particular concerning the second order Boltzmann equation that we need in order to describe the evolution of radiation and neutrinos. Its goal is three-fold: (i) to support the predictions and analytic understanding obtained on small scales in Ref.~\\cite{PUB2008} by a full numerical integration of the second order Boltzmann equation, {\\em without neglecting any term}, (ii) to compute the expected bispectrum on intermediate scales, and (iii) to revisit the large-scale behaviour of the bispectrum which has been presented in the previous literature~\\cite{Bartolo2004a,Boubekeur2009,Bartolo2005}. We emphasize that the amplitude of this non-Gaussianity is completely fixed once the amplitude and power spectrum of the first order scalar perturbations are constrained so that this offers a definite prediction of the minimal amount of non-Gaussianity expected in any CMB observation, that is provided Einstein theory remains a good description of gravity. Let us mention that our present analysis refines our previous descriptions by classifying the evolution effects, linear or not, in two categories. First, the \\emph{primary} or \\emph{early} effects arise from the evolution of the perturbations in the radiation dominated era after horizon crossing, during recombination between protons and electrons and until the potentials are constant in their linear evolution. Then \\emph{secondary} or \\emph{late} effects arise at late time when the potential starts to evolve due the late time acceleration of the Universe or due to reionisation effects (see Ref.~\\cite{2008RPPh...71f6902A} and references therein). This separation is unambiguous for the linear evolution and is therefore pertaining to the power spectrum computation but also to the calculation of the bispectrum that mix linear and second order source terms. This is actually clear from Fig. \\ref{FigSources1} below, and it is summarized in table~\\ref{Tableclassement}. \\begin{table} \\begin{tabular}{|l|l|l|} \\hline Effect & {\\bf Linear evolution} & {\\bf Non-linear evolution} \\\\ \\hline Primordial & Primordial power spectrum $P(k)$ & Primordial $\\fNL$ \\\\ \\hline Primary (early) & Sachs-Wolfe and Doppler effects & This article \\\\ \\hline Secondary (late) & Integrated Sachs-Wolfe effect, & Lensing-ISW correlation\\\\ &reionization& \\\\ \\hline \\end{tabular} \\caption{Classification of the linear and non-linear effects. \\label{Tableclassement}} \\end{table} This present work focuses on the non-linear evolution of the field during the early period and the resulting bispectra it induces. In particular we compute the shape, amplitude and bispectrum of the temperature anisotropies\\footnote{Note that these results actually depend on the actual definition of the temperature one uses (see the discussion of \\S~\\ref{subsec30}). This is due to the spectral distortion that second order effects necessarily induce as described in Ref.~\\cite{Pitrou2009ysky}.} due to these effects and eventually compare them to secondary sources, and specifically the ISW-lensing effects, and to primordial coupling effects (through standard parameterizations). In this work the non-linear evolution of the fluids, including neutrinos, is therefore treated exactly up to second-order in the linear perturbations until reionization is complete. That includes of course the use of the second order Boltzmann equation. Adiabatic initial conditions are also assumed. All the results described here can actually be reproduced from a {\\it Mathematica} code which is freely available -- with its documentation -- on the webpage~\\cite{CMBquick}. The article is organized as follows. Section~2 describes the main concepts of the non-linear cosmological perturbation theory while all equations are gathered in~\\ref{app_equations}. Then, in Section~3 we determine the initial conditions for both the metric variables and cosmic fluids (baryon, cold dark matter, photons, neutrinos). In Section~4 we describe, after a careful definition of the temperature in \\S~\\ref{subsec30}, the numerical integration based on the flat sky approximation presented in Section~3. The numerical results concerning the bispectrum are discussed in Section~5 while in Section~6 we provide an analytical understanding of these results in various limiting cases. ", "conclusions": "This article presents a complete investigation of the imprint of the non-linear dynamics on the CMB bispectrum. The calculations were carried with all the matter fields of the standard cosmological $\\Lambda$CDM model included with three families of massless neutrinos. The numerical calculations make use of a full numerical integration of the coupled system of the second order Boltzmann and Einstein equations for both the photons, with their polarization, and the neutrinos. Furthermore, line of sight integrations include first order effects in the recombination history. The initial conditions correspond to adiabatic initial conditions with a vanishing intrinsic primordial non-Gaussianity (in the context of standard single-field inflation, that implies that a contribution of the order of the slow-roll parameters has been neglected). The numerical integrations were done with the cosmological parameters of the best fit model derived from the WMAP data~\\cite{WMAP5}. This article is focused on the bispectrum of the temperature anisotropies. We have been forced however to define the temperature we use, namely in this article the bolometric temperature, since second order effects are bound to induce spectral distortions. This effect has been described in details in Ref.~\\cite{Pitrou2009ysky}, but we do not expect though that a change of definition for the temperature, like the occupation number temperature, would change significantly the conclusions we have reached. This work also demonstrates that the second order Boltzmann equation can be exactly integrated numerically and be used to produce bispectra. The resulting shape and amplitude of those bispectra is the result of intricate phenomena. It is possible though to obtain theoretical insights into peculiar cases, at small or large scale for instance. Those results confirm in particular that, at small scales, the major mechanism at play is the impact of the gravitational coupling of the dark matter potential during the matter dominated era as it had been put forward in Ref.~\\cite{PUB2008}. It is obviously difficult to grasp those results in details. In order to be able to compare the amplitude of those effects to primordial couplings, we have defined and computed equivalent $\\fNL$ parameters, $\\fNLeq$. They are defined in such a way that it is the signal a statistical indicator designed to measure primordial $\\fNL$ would get from the amplitude of the temperature bispectrum. We have found that for both primordial non-Gaussianity of local or equilateral types we have $\\fNLeq\\simeq5$ for $\\ell$ above 500. When compared to secondary effects, namely ISW-lensing couplings, primary effects are found to be of comparable amplitude. The former are however more efficient in producing a $\\fNLeq$ for the local type ; the situation is reverse for the equilateral type for which the signal comes predominantly from the primary effects. Evaluations of signal to noise ratio however show that the non-Gaussianity induced by the primary second order effects in the temperature field alone can only be marginally detected by the \\emph{Planck} mission. Secondary effects are more likely to be detected. This is however the first ever explicit and complete computation of these effects and even though the concordant model does not offer a good chance of detection, it might be a good way to put constraints on alternative cosmological models. We remind that the numerical tools used in this article are freely available and can be downloaded at~\\cite{CMBquick} so that many other bispectrum configurations and transfer functions can be investigated at will. The code makes uses of a flat sky approximation. It is accurate enough for our purpose and this approximation does not interfere with the resolution of the Boltzmann equation. The code should also be complemented with secondary non-linear effects, that have been investigated in Refs.~\\cite{Hanson2009,Mangilli2009,Spergel1999}. \\ack C.P. is supported by STFC and would like to thank Institut d'Astrophysique de Paris for its kind hospitality during part of this project. \\appendix" }, "1003/1003.2398_arXiv.txt": { "abstract": "We present optical spectroscopy of the microquasar SS\\,433 covering a significant fraction of a precessional cycle of its jet axis. The components of the prominent stationary H$\\alpha$ and H$\\beta$ lines are mainly identified as arising from three emitting regions: (i) a super-Eddington accretion disc wind, in the form of a broad component accounting for most of the mass loss from the system, (ii) a circumbinary disc of material that we presume is being excreted through the binary's L2 point, and (iii) the accretion disc itself as two remarkably persistent components. The accretion disc components move with a Keplerian velocity of $\\gtapprox 600$~\\kms\\ in the outer region of the disc. A direct result of this decomposition is the determination of the accretion disc size, whose outer radius attains $\\sim$8~\\Rsun\\ in the case of Keplerian orbits around a black hole mass of 10~\\Msun. We determine an upper limit for the accretion disc inner to outer radius ratio in SS\\,433, $R_{\\rm in}/R_{\\rm out} \\sim 0.2$, independent of the mass of the compact object. The Balmer decrements, H$\\alpha/$H$\\beta$, are extracted from the appropriate stationary emission lines for each component of the system. The physical parameters of the gaseous components are derived. The circumbinary ring decrement seems to be quite constant throughout precessional phase, implying a constant electron density of $\\log N_{\\rm e}({\\rm cm}^{-3})\\simeq 11.5$ for the circumbinary disc. The accretion disc wind shows a larger change in its decrements exhibiting a clear dependence on precessional phase, implying a sinusoid variation in its electron density $\\log N_{\\rm e}({\\rm cm}^{-3})$ along our line-of-sight between 10 and 13. This dependence of density on direction suggests that the accretion disc wind is polloidal in nature. ", "introduction": "\\label{sec:intro} Microquasars are X-ray binaries which undergo a wide range of physical processes including accretion onto a compact object (black hole or neutron star) and the launch of relativistic jets. SS\\,433, one of the most studied microquasars, became famous as the first known source of relativistic jets in the Galaxy, and it is the only system, X-ray binary or active galactic nucleus (AGN), for which atomic emission lines have so far been associated with the jets, hence implying a baryonic content \\citep[i.e., $e^- p^+$, ][]{mil79,cra81,fen00}. The optical spectrum of SS\\,433 is characterised by the presence of numerous broad emission lines with complex profiles, on top of a bright continuum. As at near-infrared and X-ray wavelengths, these emission lines can be divided into two groups: lines that are referred to as \\textit{stationary}, albeit highly variable in strength and profile, and lines that are \\textit{moving}. The latter are thought to originate in the two oppositely-directed relativistic jets moving with a speed $v \\sim 0.26~c$. The system SS\\,433 shows four main periodicities: the binary's orbital motion, with a period of about 13.08~days \\citep{cra81}; the jet axis precession and nutation, with periods of about 162 and 6~days \\citet{kat82}, respectively, and a recently discovered 550~day precession of the radio ruff \\citep{doo09}. A configuration where two jets emerge in opposite directions from the central object \\citep{fab79}, in which the jet axis undergoes a precession cycle every 162~days was proposed to describe the motion of the emission lines \\citep{mil79}, and this is referred to as the \\textit{kinematical model}. The stationary optical and near-infrared spectrum of SS\\,433 is dominated by hydrogen and {He\\,\\sc i} emission lines \\citep{mar84} and at least 15 per-cent of the flux in the hydrogen lines is contributed by the accretion disc itself \\citep{seba09}. The emission-line spectra of accretion discs, and their evolution with orbital and precessional phases, comprise most of the information we can obtain about the temperature and density variations as well as velocity gradients within the disc \\citep{ski00}. The Balmer decrements of the stationary lines in SS\\,433 have not previously been studied due to the high interstellar extinction towards the object, which makes the detection of the H$\\beta$ line rather difficult \\citep{pan97}. The decrements are highly dependent on physical parameters of the gas such as its temperature, optical depth and also the nature of the source of radiation \\citep{dra80}. The hard radiation field around compact objects and the expected high electron densities for SS\\,433 \\citep[$N_{\\rm e} \\ge 10^{13}$~cm$^{-3}$, ][]{pan97} require that we use an adequate treatment for such environments in order to study the emission-line gas. At high densities, and in the presence of heavy elements, excitation by collisional processes become a relevant factor \\citep{fer88}. \\citet{dra80} performed theoretical calculations of the emission-line spectrum from a slab of hydrogen at moderate to high densities ($10^8 < N_{\\rm e} < 10^{15}~$cm$^{-3}$) over a wide range of physical parameters, including values close to those observed in SS\\,433's gas \\citep{pan97}. We compare our estimates with Drake \\& Ulrich's findings in Section~\\ref{sec:results}. The existence of dust mixed with the emitting gas would have effects on the Balmer decrement that are by no means negligible \\citep{ost06}. SS\\,433's optical spectrum reveals evidence of severe dust extinction, such as the remarkably red continuum and the presence of prominent interstellar absorption lines in the form of diffuse interstellar bands \\citep{mar84}. In our optical data described in subsequent sections, and in agreement with previous observations \\citep[e.g.,][among others]{mur80,mar84,gies02}, we have detected the diffuse interstellar bands at 4430, 5778 and 5780~\\AA\\ and also the interstellar lines Ca~H\\,$\\lambda3968$, Ca~K\\,$\\lambda3934$ as well as Na~D\\,$\\lambda5890$. \\citet{mur80} reported an interstellar absorption (i.e., Galactic extinction) of $A_V \\sim 8$~mag toward SS\\,433, obtained from infrared measurements. This result has been corroborated by \\citet{gies02} by fitting the spectral energy distribution, and the currently most accepted value is $A_V = 7.8$~mag. Fig.~\\ref{fig:dust} shows the map of Galactic extinction in the direction of the SS\\,433/W50 complex. We constructed this extinction image by converting the colour excess map of \\citet{sch98} into an estimate of the reddening, assuming a selective extinction, $R_V \\equiv A_V/E(B-V)$, equal to the average value for the Galaxy \\citep[i.e., $R_V=3.1$,][]{car89}, where $E(B-V)$ is the colour excess. It is clear from the image that the $\\sim$7.8~mag of extinction is consistent with the interstellar absorption gradient. The extinction estimate from this map towards SS\\,433's position is $A_V = 7.813$~mag. Although the value for $E(B-V)$ is quite accurate, the choice of $R_V$ is not. The intrinsic error introduced by $R_V$ implies that the accuracy of the extinction estimate is not better than to a tenth of a magnitude. Therefore, we corrected our spectra (whose reduction is described in Section~\\ref{sec:obs}) for $A_V = 7.8$~mag of interstellar absorption using the reddening law of \\citet{car89}, before calculating the Balmer decrements. In this paper we decompose the profiles of the stationary emission lines, H$\\alpha$ and H$\\beta$, with a number of Gaussian components (Section~\\ref{sec:decomposition}). Each model component is identified with its corresponding emitting origin. In Section~\\ref{sec:results} we analyse the physical conditions of each component present in SS\\,433 and its behaviour as a function of orbital and precessional phases, via the Balmer decrement. In Section~\\ref{sec:size} we present our calculations for the accretion disc radii and compare our findings with discs from other objects that share some common features with SS\\,433. \\begin{figure} \\centering\\includegraphics[angle=270, width=\\columnwidth]{figs/dust_overlay.eps} \\caption{Colour-scale image shows the Galactic extinction ($A_V$) towards the W50/SS\\,433 system from the IRAS/COBE all-sky survey \\citep{sch98}. Colour scale corresponds to visual extinction in magnitudes and it is centred on SS\\,433's position (given by the red circle). Grey contours on the colour image are 3, 3.6, 5.4, 8.4, 12.6 and 18~mag. The radio continuum emission at 1465~MHz is shown as black contours \\citep[data from ][]{dub98}. Radio contours are 12, 14 and 20~mJy~beam$^{-1}$. } \\label{fig:dust} \\end{figure} ", "conclusions": "We have presented optical spectroscopy of the microquasar SS\\,433 covering a significant fraction of the precessional cycle of its accretion disc. The components of the prominent stationary H$\\alpha$ and H$\\beta$ lines have been identified as arising from three emitting regions: an accretion disc wind which is super-Eddington \\citep{seba09}, in the form of a broad component accounting for most of the mass loss in the system, a circumbinary disc of material probably excreted through the binary's L2 point, and the accretion disc itself as two persistent components, having an outer region Keplerian velocity of $\\gtapprox 600$~\\kms. A direct result of this decomposition using our UKIRT data published in \\citet{seba09} was the determination of the accretion disc size, whose outer radius attains 8~\\Rsun, for an assumed black hole mass of 10~\\Msun. With the data presented in this paper we determined the accretion disc inner to outer radius ratio in SS\\,433, $R_{\\rm in}/R_{\\rm out}$ to be $\\sim$0.2, independent of the mass of the compact object. The Balmer decrements, H$\\alpha/$H$\\beta$, were extracted from the stationary emission lines for each component of the system. The decrement of the circumbinary ring seems to be quite constant throughout precessional phase, implying a fairly constant electron density of $\\log N_{\\rm e}\\simeq 11.5$ for the circumbinary disc. The accretion disc wind shows larger changes in its decrements as a function of precessional phase, implying variations in its density, $N_{\\rm e}$, between $10^{10}$ and $10^{13}$~cm$^{-3}$. Thus, the physical parameters of the gaseous components imply rather dense environments emitting the Balmer lines." }, "1003/1003.1744_arXiv.txt": { "abstract": "We discuss the potential of using the \\HeI\\ $584~$\\AA\\ forest to detect and study \\HeII\\ reionization. Significant $584~$\\AA\\ absorption is expected from intergalactic \\HeII\\ regions, whereas there should be no detectable absorption from low density gas in \\HeIII\\ regions. Unlike \\HeII\\ Ly$\\alpha$ absorption (the subject of much recent study), the difficulty with using this transition to study \\HeII\\ reionization is not saturation but rather that the absorption is weak. The Gunn-Peterson optical depth for this transition is $\\tau \\sim 0.1 \\,x_{\\rm HeII} \\,\\Delta^2 \\, [(1+z)/5]^{9/2}$, where $x_{\\rm HeII}$ is the fraction of helium in \\HeII\\ and $\\Delta$ is the density in units of the cosmic mean. In addition, \\HeI\\ $584~$\\AA\\ absorption is contaminated by lower redshift \\HI\\ Ly$\\alpha$ absorption with a comparable flux decrement. We estimate the requirements for a definitive detection of redshifted \\HeI\\ absorption from low density gas ($\\Delta \\approx 1$), which would indicate that \\HeII\\ reionization was occurring. We find that this objective can be accomplished (using coeval \\HI\\ Ly$\\alpha$ absorption to mask dense regions and in cross correlation) with a spectral resolution of $10^4$ and a signal-to-noise ratio per resolution element of $\\sim 10$. Such specifications may be achievable on a few known $z\\sim 3.5$ quasar sightlines with the Cosmic Origins Spectrograph on the Hubble Space Telescope. We also discuss how \\HeI\\ absorption can be used to measure the hardness of the ionizing background above $13.6$~eV. ", "introduction": "Starlight produced by the first galaxies is the leading candidate for ionizing the hydrogen as well as singly ionizing the helium at $z\\sim 6$. It takes a harder source of radiation to doubly ionize the helium, so the reionization of this species is likely deferred until $z\\sim 3$ when quasars produce a sufficient hard UV background \\citep{madau99, furlanetto08, mcquinn09}. However, the helium could have been doubly ionized at nearly the same cosmic time that hydrogen was reionized if more exotic sources ionized the hydrogen, such as the first generation of metal-free stars \\citep{bromm01, venkatesan03, tumlinson04} or miniquasars \\citep{madau99, volonteri09}. In a third potential scenario, early sources doubly ionized the helium and then shut off. Afterward, the \\HeII\\ recombined such that quasars could again reionize it at $z\\sim 3$ \\citep{wyithe03, venkatesan03}. If \\HeII\\ reionization were completing at $z \\sim 3$, an epoch for which there are numerous observations of the intergalactic medium (IGM), it should be an easier task to definitively detect this process compared to detecting $z \\gtrsim 6$ reionization processes. Furthermore, if \\HeII\\ reionization were ending at $z \\sim 3$, it should have significantly affected the temperature of the intergalactic gas and the ultraviolet radiation background. These motivations, along with recent additions to the Hubble Space Telescope (HST), have inspired a significant effort of late to understand the signatures and the detection prospects of \\HeII\\ reionization \\citep{furlanetto08, mcquinn09, lidz09, bolton09, mcquinn09b, dixon09, furlanetto09, syphers09, syphers09b, mcquinn3He}. Three separate observations of the $z\\sim 3$ IGM suggest that \\HeII\\ reionization was ending around this redshift: First, several studies have measured the temperature of the intergalactic gas from the widths of the narrowest lines in the \\HI\\ Ly$\\alpha$ forest, and the majority of these studies have found evidence for an increase in the IGM temperature of $\\sim 10^4$~K between $z \\approx 4$ and $z \\approx 3$, before a decline to lower redshift \\citep{schaye00, ricotti00, lidz09}. These trends have been attributed to the heating from \\HeII\\ reionization. Second, observations of \\HeII\\ Ly$\\alpha$ absorption from gas at $2.8 < z < 3.3$ show tens of comoving Mpc (cMpc) regions with no detected transmission \\citep{reimers97, heap00}, which may indicate that \\HeII\\ reionization was not complete. Thirdly, \\citet{songaila98} and \\citet{agafonova07} detected evolution in the column density ratios of certain highly ionized metals at $z \\approx 3$, which they argued was due to a hardening in the ionizing background around $50~$eV and, thus, the end of \\HeII\\ reionization. However, the interpretations of all of these indications for \\HeII\\ reionization are controversial. Temperature measurements of the IGM are difficult, and not all measurements detected the aforementioned trends. It is often argued that \\HeII\\ Ly$\\alpha$ absorption saturates at \\HeII\\ fractions that are too small ($\\sim 10^{-3}$ at the mean density) to study \\HeII\\ reionization (although, see \\citealt{mcquinn09b}). Lastly, inferences from metal lines require significant modeling, and studies have reached different conclusions regarding the degree of their evolution at $z\\sim 3$ \\citep{boksenberg03}. This paper discusses intergalactic absorption by the \\HeIallowedn\\ transitions of \\HeI\\ as an unsaturated observable of \\HeII\\ reionization. We primarily focus on the strongest and longest wavelength of these absorption lines, the \\HeI\\ \\HeIallowed, $584~$\\AA\\ line. For a given optical depth in the \\HI\\ Ly$\\alpha$ forest, the amount of absorption in the \\HeI\\ forest can be directly estimated if both the fraction of helium that is \\HeII\\ and the ratio of the \\HI\\ and \\HeI\\ photoionization rates are known. In addition, because the \\HeI\\ ionization edge is relatively close to that of hydrogen, the photoionization rate that \\HeI\\ experiences is similar to this for hydrogen. Thus, the most important determinant of the amount of \\HeI\\ $584$~\\AA\\ absorption is the \\HeII\\ fraction. Other studies have discussed the \\HeI\\ forest, but did not focus on its usefulness as a probe of helium reionization. \\citet{tripp90} was the first to discuss intergalactic absorption by the \\HeI\\ $584~$\\AA, and this study attempted unsuccessfully to detect this absorption in the spectrum of a $z=1.7$ quasar. \\citet{reimers93} targeted \\HeI\\ absorption from $z\\approx 2$ \\HI\\ Lyman-limit systems, and reported the first (and at present only) detection of intergalactic \\HeI\\ $584~$\\AA\\ absorption. Finally, \\citet{miralda92} showed that the \\HeI\\ forest could be a useful probe of the hardness of the ultraviolet background, and they focused on using this absorption to rule out a particular nonstandard model for the dark matter. \\citet{santos03} followed up on this idea, arguing that the $z\\sim 5$ \\HeI\\ forest could be a useful diagnostic of the hardness of the ultraviolet background after hydrogen reionization. As with the \\HeII\\ Ly$\\alpha$ transition at $304~$\\AA, the \\HeI\\ $584~$\\AA\\ transition falls blueward of the hydrogen Lyman limit and, therefore, is subject to continuum absorption by hydrogen, in this case from \\HI\\ systems with $z > z' \\equiv 3.2 \\, (1 +z_{\\rm HeI})/5 -1$. This continuum absorption may even be worse in terms of obscuring the \\HeI\\ $584~$\\AA\\ forest compared to the \\HeII\\ Ly$\\alpha$ forest because this spectral region is more affected by higher redshift \\HI\\ continuum absorbers. It is unlikely that there are more than a handful of quasar sightlines at $z>4$ with sufficient near ultraviolet (NUV) flux for detection with the present generation of instruments (and almost certainly not at $z>6$, during \\HeI\\ reionization; \\citealt{santos03}). Therefore, our focus is on applications of \\HeI\\ absorption at $z\\lesssim 4$. We have performed a cursory search for candidate \\HeI\\ sightlines among the relatively few published HeII Lya forest spectra that extend redward to $584 \\, (1+z_{\\rm QSO})~$\\AA. Of note, QSO~OQ~172 ($z=3.54$) has $F_\\lambda \\approx 2 \\times 10^{-16}~$erg~s$^{-1}$~cm$^{-2}$~\\AA$^{-1}$ at $584 \\, (1+z) ~ $\\AA\\ \\citep{lyons95}, and QSO~0055-269 ($z=3.67$) has $F_\\lambda \\approx 1 \\times 10^{-16}$ (Gabor Worseck, priv. com.). In fact, we were surprised to find that most \\HeII\\ sightlines in our search had significant flux in the relevant band. Many of the existing \\HeII\\ sightlines had been selected by their far ultraviolet flux. NUV selection would be more optimal for identifying candidate \\HeI\\ forest sightlines. For example, HS~1140~+3508 ($z=3.15$) is obscured in the far ultraviolet by a Lyman-limit system, but has a NUV flux of $F_\\lambda \\approx 2 \\times 10^{-16}$ (Gabor Worseck, priv. com.). Another difficulty with the \\HeI\\ forest is that line absorption from foreground systems can contaminate the \\HeI\\ forest, the most important of which is \\HI\\ Ly$\\alpha$ absorption. \\HI\\ Ly$\\alpha$ absorption from a system with a redshift of $z_{\\rm HI, 1216} = 2.4 \\, (1 +z_{\\rm HeI, 584})/5 -1$ falls directly on top of the \\HeI\\ absorption from a redshift of $z_{\\rm HeI, 584}$. Fortunately, the Ly$\\alpha$ forest is quite thin at relevant $z_{\\rm HI, 1216}$ (with a flux decrement of several percent), and we find that both this contaminant's mean transmission and variance tend to be comparable to that of the \\HeI\\ forest. Also, the \\HeI\\ forest correlates strongly with the coeval \\HI\\ Ly$\\alpha$ forest, which allows it to be extracted in cross correlation despite this contamination. This study is timely because the HST reservicing mission installed the Cosmic Origins Spectrograph (COS) in May 2009. COS is capable of measurements of HeI $584~$\\AA\\ absorption at redshifts relevant to \\HeII\\ reionization ($2.8 \\lesssim z < 4.5$), and ground-based spectrographs can cover higher redshifts ($z \\gtrsim 4.3$, corresponding to $\\gtrsim 3100$~\\AA). COS is able to achieve higher signal-to-noise ratios than previous instruments in the ultraviolet. A $60~$hr exposure with COS would achieve a signal-to-noise ratio of $10$ at ${\\it R} \\approx 30,000$ for a flux of $F_\\lambda = 2 \\times 10^{-16}~$erg~s$^{-1}$~cm$^{-2}$~\\AA$^{-1}$ (the flux of QSO~OQ~172 and HS~1140~+3508) at $2600~$\\AA\\ (\\HeI\\ absorption from $z=3.5$). A $40~$hr exposure with COS would achieve a signal-to-noise ratio of $10$ for this flux at ${\\it R} \\approx 20,000$ at $2300~$\\AA\\ ($z=2.9$).\\footnote{http://etc.stsci.edu/webetc/. As of mid-January 2010, the dark current for the COS NUV gratings is $5$ times higher than specifications (and than what is assumed by this ETC), and, thus, a significantly longer observation is required unless this issue is resolved.} This paper is organized as follows. Section~\\ref{sec:forest} discusses the physics of the \\HeI\\ forest. Section~\\ref{sec:calculations} contrasts simulated spectra for this forest under different assumptions regarding the ionization state of the helium. Section~\\ref{sec:tests} quantifies the spectral quality an observation must achieve to verify whether \\HeII\\ reionization was occurring from the \\HeI\\ forest. Appendix~\\ref{sec:hardness} outlines how to measure the hardness of the ionizing background with \\HeI\\ absorption, and Appendix~\\ref{ap:SNRcc} derives formulae for the significance with which \\HeI\\ absorption can be detected in cross correlation with the coeval Ly$\\alpha$ forest. This paper assumes a flat $\\Lambda$CDM cosmology with $h =0.7$, $\\Omega_b = 0.046$, $\\Omega_m = 0.28$, $\\sigma_8 = 0.82$, $n_s = 1$, and $Y_{\\rm He} = 0.24$, consistent with recent measurements \\citep{komatsu08}. However, the simulation used to calculate the Ly$\\alpha$ forest spectra at $z < 1.5$, the D5 simulation in \\citet{springel03}, assumes a slightly different cosmology, with the most notable differences being $\\Omega_b = 0.04$ and $\\sigma_8 = 0.9$. The photoionization and recombination rates used in this study are from \\citet{hui97}. ", "conclusions": "This paper discussed the usefulness of the \\HeI\\ $584~$\\AA\\ forest to study \\HeII\\ reionization at $3 \\lesssim z \\lesssim 4.5$. The optical depth of the \\HeI\\ $584~$\\AA\\ line is proportional to the optical depth in \\HI\\ Ly$\\alpha$ by the factor $0.025 \\, x_{\\rm HeII} \\, \\Gamma_{\\rm HI}/\\Gamma_{\\rm HeI}$ (ignoring differences in the amount of thermal broadening between \\HI\\ and \\HeI). The factor $x_{\\rm HeII}$ should have been spatially variable, but the ratio $\\Gamma_{\\rm HI}/\\Gamma_{\\rm HeI}$ should have been essentially spatially independent and roughly equal to unity at relevant redshifts. Therefore, this absorption can be used to study \\HeII\\ reionization through its dependence on $x_{\\rm HeII}$. Our best method at present to probe \\HeII\\ reionization, \\HeII\\ Ly$\\alpha$ absorption, saturates at neutral fractions of a part in a thousand at the cosmic mean density. In contrast, \\HeI\\ $584~$\\AA\\ absorption is unsaturated in all except the densest regions, even for $x_{\\rm HeII} = 1$. This absorption provides a complementary window into the ionization state of intergalactic helium at $z\\sim 3$ that can definitively test whether an opaque region in the \\HeII\\ Ly$\\alpha$ forest was due to a large-scale \\HeII\\ region. Even for realistic amounts of instrumental noise and foreground \\HI\\ Ly$\\alpha$ absorption, we showed that the coeval \\HI\\ Ly$\\alpha$ forest absorption can be used to construct a matched filter that can detect the \\HeI\\ absorption at high significance from a single quasar sightline. A detection of \\HeI\\ absorption from $\\Delta_b \\sim 1$ gas at $z \\sim 3$ would definitively indicate that \\HeII\\ reionization was occurring. If a significant fraction of the intergalactic helium was in \\HeII, we found that \\HeI\\ $584~$\\AA\\ forest absorption from low density gas could be identified in a quasar spectrum with SNR~$\\sim10$ and ${\\cal R} \\sim 10^4$ in an interval of $\\Delta z = 0.2$. These specifications may be achievable with the COS instrument on the HST for a few known $z>3$ targets. \\\\ We thank Claude-Andr{\\'e} Faucher-Gigu{\\`e}re, Wayne Hu, Gabor Worseck, and especially J. Xavier Prochaska for useful discussions. We also thank Claude-Andr{\\'e} Faucher-Gigu{\\`e}re, Lars Hernquist, Adam Lidz, and Volker Springel for providing the simulations used in this work. E. S. acknowledges support by NSF Physics Frontier Center grant PHY-0114422 to the Kavli Institute of Cosmological Physics. \\\\" }, "1003/1003.4211_arXiv.txt": { "abstract": "{} {Intrinsic alignments constitute the major astrophysical systematic for cosmological weak lensing surveys. We present a purely geometrical method with which one can study gravitational shear-intrinsic ellipticity correlations directly in weak lensing data.} {Linear combinations of second-order cosmic shear measures are constructed such that the intrinsic alignment signal is boosted while suppressing the contribution by gravitational lensing. We then assess the performance of a specific parametrisation of the weights entering these linear combinations for three representative survey models. Moreover a relation between this boosting technique and the intrinsic alignment removal via nulling is derived.} {For future all-sky weak lensing surveys with photometric redshift information the boosting technique yields statistical errors on model parameters of intrinsic alignments whose order of magnitude is compatible with current constraints determined from indirect measurements. Parameter biases due to a residual cosmic shear signal are negligible in case of quasi-spectroscopic redshifts and remain sub-dominant for typical values of the photometric redshift scatter. We find good agreement between the performance of the intrinsic alignment removal based on the boosting technique and standard nulling methods, both reducing the cumulative signal-to-noise by about a factor of 6, which possibly indicates a fundamental limit in the separation of lensing and intrinsic alignment signals.} {} ", "introduction": "\\label{sec:intro} Weak gravitational lensing of the large-scale structure is going to be one of the major cosmological probes contributing to reveal the properties of dark matter and dark energy in the near future \\citep{albrecht06,peacock06}. Within the past decade the method has evolved from its first detections \\citep{bacon00,kaiser00,vwaer00,wittman00} to maturity, nowadays yielding statistical constraints which are compatible to other probes (for recent measurements see e.g. \\citealp{benjamin07}; \\citealp{fu07}; \\citealp{schrabback09}; for a recent review see \\citealp{munshi08}). Planned surveys measuring weak lensing on cosmological scales, or cosmic shear in short, include Pan-STARRS\\footnote{\\texttt{http://pan-starrs.ifa.hawaii.edu}}, KIDS\\footnote{\\texttt{http://www.astro-wise.org/projects/KIDS}}, DES\\footnote{\\texttt{https://www.darkenergysurvey.org}}, LSST\\footnote{\\texttt{http://www.lsst.org}}, and Euclid\\footnote{\\texttt{http://sci.esa.int/science-e/www/area/\\\\index.cfm?fareaid=102}}. The increasingly large statistical power of these surveys demands a more and more thorough treatment of systematic errors. The major astrophysical contamination to cosmic shear is constituted by the intrinsic alignment of galaxies. To infer cosmic shear information from the correlation of galaxy ellipticities, it is usually assumed that the intrinsic shapes of galaxy images are purely random, so that only the desired correlations of gravitational shear (GG in the following) remain. However, due to interactions with the surrounding matter structure, galaxy shapes can intrinsically align, causing correlations between the intrinsic ellipticities of galaxies (II hereafter). Moreover matter can influence the shape of a close-by galaxy via tidal forces and at the same time contribute to the lensing signal of a background galaxy, thereby producing gravitational shear-intrinsic ellipticity correlations (GI hereafter). Intrinsic alignments have been subject to extensive studies, both analytical and using simulations \\citep{croft00,heavens00,lee00,pen00,catelan01,crittenden01,jing02,mackey02,bosch02,hirata04,heymans06,bridle07a,semboloni08,okumura08,okumura09,brainerd09}. Results vary widely, but are mostly consistent with a contamination of the order $10\\,\\%$ by both II and GI signals for future weak lensing surveys, which can lead to serious biases on cosmological parameters if left untreated \\citep[e.g.][]{bridle07}. Intrinsic alignments depend intricately on the formation and evolution of galaxies within their dark matter environment, so that models cannot be expected to develop far beyond the current crude level in the near future. For the most recent advancement in intrinsic alignment modelling see \\citet{schneiderm09}. Using uncertain models of limited accuracy for assessing systematics in statistical analyses is risky \\citep{kitching08}. Therefore observational data which can put limits on the possible range of intrinsic alignment signals are highly warranted. It should be noted that in principle intrinsic alignments constitute an interesting cosmological signal worth investigating, shedding light onto the interaction between galaxies, their haloes, and the large-scale structure. Both II and GI correlations have been subject to investigations in several data sets \\citep{brown02,heymans04,mandelbaum06,hirata07,brainerd09,mandelbaum09}, results ranging from null to significant detections, depending strongly on the type and colour of galaxies considered. However, none of these observations were direct measurements of intrinsic alignments for the galaxy populations and redshifts which are most interesting for cosmic shear because in those cases the shear signal clearly dominates the correlations of galaxy ellipticities. While the II signal is observed at small redshifts where cosmic shear is negligible, the GI term is usually inferred from cross-correlations between galaxy number densities and ellipticities in samples with spectroscopic redshifts. The latter approach requires the assumption of a simple form of the galaxy bias, which is of limited accuracy and inapplicable on small scales. If one wishes to analyse larger galaxy samples for which only photometric redshift information is available, further signals such as galaxy-galaxy lensing contribute and need to be modelled carefully (see \\citealp{bernstein08,joachimi10} for an overview on the types of signals contributing to correlations between galaxy number density and ellipticity). The II signal is less of a concern because, in order to intrinsically align, a pair of galaxies has to have interacted physically, and hence to be both close on the sky and in redshift. This fact can be used to remove II correlations \\citep{king02,king03,heymans03,takada04b}, partly in a fully model-independent way with only marginal loss of statistical power if precise redshift information is available. The GI signal is not restricted to physically close pairs of galaxies, but it can also be eliminated in a purely geometrical way via nulling techniques \\citep{joachimi08b,joachimi09}. However, a considerable loss of cosmological information is inherent to nulling, and hence, it is still desirable to have a reliable model of GI correlations at one's disposal to be used with other methods controlling this systematic \\citep{king05,bridle07,bernstein08,zhang08,joachimi10}. In the following we will develop a model-independent technique to extract the GI signal from a cosmic shear data set, thereby allowing for direct measurements of GI correlations on the most relevant galaxy samples. This \\lq GI boosting\\rq\\ approach can be regarded as complementary to nulling both in its purpose and in its implementation. Analogous to the nulling technique, we will construct linear combinations of second-order cosmic shear measures, making only use of the well-known characteristic redshift dependence of the GI and GG terms. This paper is organised as follows. In Sect.$\\,$\\ref{sec:method} we present the principle of GI boosting and derive general conditions, which are used in Sect.$\\,$\\ref{sec:weights} to explicitly construct weight functions for the boosting transformation of the cosmic shear signal. Section \\ref{sec:modelling} details the modelling which we apply in Sect.$\\,$\\ref{sec:performance} to assess the performance of the boosting technique. In Sect.$\\,$\\ref{sec:nulling} we construct a method to remove GI correlations based on the GI boosting technique and investigate the relation between the new approach and the standard nulling method of \\citet{joachimi08b,joachimi09}, before we summarise and conclude in Sect.$\\,$\\ref{sec:conclusions}. ", "conclusions": "\\label{sec:conclusions} In this paper we presented a method which extracts shear-ellipticity correlations (the GI signal) from a tomographic cosmic-shear data set. The approach relies neither on models of intrinsic alignments nor on knowledge of the cosmological parameters that characterise the cosmic shear (GG) signal, making only use of the typical and well-understood redshift dependencies of both the GI and GG term. We derived constraints which a linear transformation of second-order cosmic shear measures has to fulfil in order to boost the GI signal and simultaneously suppress the lensing contribution. We studied in depth a particular parametrisation of the weights entering this transformation and analysed the performance of the resulting GI boosting technique for three representative survey models. Applying the GI boosting to future all-sky cosmic shear surveys, it should be possible to isolate the GI signal with subdominant biases due to a residual GG term, and with constraints that are comparable to current results from indirect measurements of shear-ellipticity correlations \\citep{mandelbaum09}. If one restricts the analysis to galaxies with photometric redshift information of good quality, i.e. a redshift scatter of not more than $\\sigma_{\\rm ph}(1+z)$ with $\\sigma_{\\rm ph}=0.03$, one can achieve $1\\,\\sigma$-errors on the GI signal amplitude $A$ in the parametrisation of (\\ref{eq:GImodel}) of better than 0.2 when varying only the amplitude, and a marginalised error of approximately 0.7 when fitting an additional redshift dependence. Using all galaxies from a survey fulfilling $\\sigma_{\\rm ph} \\leq 0.05$, the statistical constraints degrade only marginally but the parameter bias due to the residual GG contribution can attain more significant values of up to $b_{\\rm sys}/\\sigma_{\\rm stat} \\lesssim 2/3$. We also considered a survey with high-quality photometric or spectroscopic redshifts. However, the expected low number density of galaxies of $n_{\\rm g}=1\\,{\\rm arcmin}^{-2}$, even for future surveys, does not permit us to place competitive constraints on intrinsic alignment models. In this case of highly accurate redshift information the residual bias on parameters is negligible. Although we have modelled scatter in photometric redshifts for our investigations, we did not consider other effects affecting the accuracy of redshift information, such as an error in the median of the galaxy redshift distributions or catastrophic failures in the determination of photometric redshifts. As several studies of intrinsic alignment removal techniques have demonstrated \\citep[e.g.][]{bridle07,joachimi09,joachimi10}, the ability to separate the GI from the GG signal depends vitally on these parameters characterising the accuracy of and knowledge about redshifts. The same can be expected for the GI boosting technique, possibly to an even larger extent since in this case one attempts to suppress the originally strongest contribution to ellipticity correlations, the GG signal. Hence, we hypothesise that the requirements of future ambitious weak lensing surveys, like a negligible fraction of catastrophic failures and an error in the mean of each redshift distribution of not more than $0.002(1+z)$ \\citep{laureijs09}, are both necessary and sufficient for a success of GI boosting. We leave a detailed assessment of the requirements on the quality of redshift information to future work. Moreover, we did not yet include intrinsic ellipticity correlations (II) into our considerations. Since the II signal is generated by physically close pairs of galaxies, it has a redshift dependence that is clearly distinct from the GI and GG terms, and can thus be removed relatively easily \\citep{king02,king03,heymans03,takada04b}. In tomographic cosmic shear data it mainly affects auto-correlations and cross-correlations of adjacent photometric redshift bins with significant overlap of their corresponding distributions of true redshifts. One of the aforementioned II removal techniques could precede the GI boosting, causing an increased shape noise contribution in particular in the auto-correlations due to the reduced number of available galaxy pairs. Alternatively, the downweighting of the II signal could also be readily incorporated into the boosting technique by introducing the additional condition $\\partial^2 G^{(i)}/ \\partial \\chi^2 |_{\\chi_i} = 0$, implying $B^{(i)}(\\chi_i)=0$ and therefore a downweighting of auto-correlations as well as cross-correlations of adjacent redshift distributions, see (\\ref{eq:defpi}). Our findings still have the potential for significant improvement because we have only considered one specific parametrisation of the weight function that governs the boosting transformation. While this choice is intuitive and allows analytical progress, a more versatile approach could be to assume the weight function $B^{(i)}(\\chi)$ as piecewise linear, with nodes placed at the median redshift of every galaxy redshift sample. The constraints on GI boosting and GG suppression could then be directly imposed on the discretised version of the boosting transformation, thereby fixing a subset of the values of $B_Q^{(i)}(\\chi)$ at its nodes. The remaining freedom in the weight function could for instance be used to maximise the signal-to-noise of the expected transformed GI signal. We also constructed a method of GI removal, directly based on a slightly modified version of the GI boosting technique. In principle, we showed that if one is able to isolate the GI signal via boosting, one can simply subtract a rescaled version of the GI term from the original cosmic shear measures to eliminate the intrinsic alignment systematic. We find that the residual contamination of the cosmic shear signal by GI correlations is indeed small, and that the cumulative signal-to-noise of the thus treated cosmic shear signal decreases by about a factor of 6. This value is remarkably close to the result for the standard GI nulling technique as introduced by \\citet{joachimi08b,joachimi09}, in spite of the differing approaches. The underlying reason for this agreement may be due to a fundamental limit in the ability to separate GI and GG signals relying only on the dependence on redshift, which is worth to be addressed in future investigations. Of course, such a limit would also imply a maximum accuracy with which parameters of intrinsic alignments can be constrained via GI boosting. Like the method devised in this work, the standard nulling technique is also a purely geometrical method. Hence, a combined application of GI boosting and nulling to a cosmic shear data set would still be based on a minimum of assumptions about the actual forms of signals or the values of model parameters. For instance one could use an initial analysis based on nulling to yield robust estimates of the cosmic shear signal and the corresponding cosmological model. This could then be used to construct weights for the GI boosting transformation such that even in the case of standard photometric redshift quality (which we assumed to be $\\sigma_{\\rm ph}=0.05$ in this paper) the bias due to the residual GG signal would be negligible, thereby enabling an equally robust estimate of the GI signal. Ultimately, the cosmic shear analysis, the treatment of intrinsic alignments, and the inclusion of additional information from galaxy number density correlations \\citep[as in][]{mandelbaum06,hirata07,mandelbaum09} will all be efficiently combined into a simultaneous analysis of the form presented in \\citet{bernstein08} and \\citet{joachimi10}, provided one can summon the computational power. Yet the model-independent, direct, and robust boosting technique, as well as nulling and the combination of the two, will prove useful e.g. to provide reliable priors on the large set of parameters entering the integrative approaches and in addition serve as a valuable consistency check in cosmic shear analyses." }, "1003/1003.2437_arXiv.txt": { "abstract": "We present observations of newly discovered 24 \\mum\\ circumstellar structures detected with the Multiband Imaging Photometer for {\\em Spitzer} (MIPS) around three evolved stars in the Cygnus-X star forming region. One of the objects, BD+43 3710, has a bipolar nebula, possibly due to an outflow or a torus of material. A second, HBHA 4202-22, a Wolf-Rayet candidate, shows a circular shell of 24 \\mum\\ emission suggestive of either a limb-brightened shell or disk seen face-on. No diffuse emission was detected around either of these two objects in the {\\em Spitzer} 3.6-8 \\mum\\ Infrared Array Camera (IRAC) bands. The third object is the luminous blue variable candidate G79.29+0.46. We resolved the previously known inner ring in all four IRAC bands. The 24 \\mum\\ emission from the inner ring extends $\\sim1\\farcm2$ beyond the shorter wavelength emission, well beyond what can be attributed to the difference in resolutions between MIPS and IRAC. Additionally, we have discovered an outer ring of 24 \\mum\\ emission, possibly due to an earlier episode of mass loss. For the two shell stars, we present the results of radiative transfer models, constraining the stellar and dust shell parameters. The shells are composed of amorphous carbon grains, plus polycyclic aromatic hydrocarbons in the case of G79.29+0.46. Both G79.29+0.46 and HBHA 4202-22 lie behind the main Cygnus-X cloud. Although G79.29+0.46 may simply be on the far side of the cloud, HBHA 4202-22 is unrelated to the Cygnus-X star formation region. ", "introduction": "} Cygnus-X is arguably the biggest and brightest massive star formation region within 2 kpc of the Sun (d$\\sim$1.7 kpc, Schneider et al. 2006, 2007). It contains several hundred HII regions, numerous OB associations with over a thousand OB stars, and thousands of low mass stars. We have performed an unbiased survey of $\\sim$24 sq degrees in Cygnus-X with the Infrared Array Camera (IRAC, Fazio et al. 2004) and the Multiband Imaging Photometer for \\spi\\ (MIPS, Rieke et al. 2004) on the {\\em Spitzer Space Telescope} (Werner et al. 2004). The primary objectives of this \\spi\\ Legacy project were to study the evolution of high mass protostars, clustering of both high and low mass stars in the complex, as well as several other goals described by Hora et al. (2008b). Follow-up spectroscopy with \\spi's Infrared Spectrograph (IRS; Houck et al. 2004) was obtained for several dozen promising sources in the MIPS and IRAC survey regions. In addition to the numerous young stars that are the primary focus of the Legacy project, there are dozens of evolved objects in the survey region, from Miras and Cepheids to Wolf-Rayet and carbon stars to planetary nebulae and supernova remnants. After the initial discovery of a bipolar nebula at 24 \\mum\\ around a carbon star during the data processing quality control effort, a systematic search for additional circumstellar nebulae was made. SIMBAD was used to compile positional lists of evolved objects that lie within the bounds of the 24 \\mum\\ observations. The 24 \\mum\\ images were then inspected for the presence of point sources and associated nebulosity, if any. More than 80\\% (146) of the 191 known evolved objects are detected at 24 \\mum\\ as point sources (Table 1). These types of objects can be bright in the infrared for a variety of reasons such as emission from dust or forbidden lines. This, of course, is an incomplete survey of the evolved stars as it does not include the faint population of asymptotic giant branch (AGB) stars detected with \\spi\\ (also known as contaminants in the young stellar object studies). Be that as it may, only three evolved stars in the Cygnus-X region were distinguished from the vast majority by exhibiting extended 24 \\mum\\ emission. The emission around two of these stars, which are listed in Table 2, was first reported by Kraemer et al. (2009). Here, we discuss the morphological characteristics of the circumstellar structures as well as the spectral properties of the dust and the central exciting objects. ", "conclusions": "} \\subsection{BD+43 3710} BD+43 3710, located in the extreme northeast corner of the Cygnus-X survey region, is a little known carbon star candidate classified as spectral type R: by Nassau \\& Blanco (1954). It shows a remarkable bipolar structure at 24 \\micron\\ that is completely absent at the shorter wavelengths (it is also outside the 70 \\micron\\ MIPS coverage). Figure \\ref{fig.bd_3c} shows a three-color image of BD+43 3710 combining the 3.6 \\mum, 8 \\mum, and 24 \\mum\\ data while Figure \\ref{fig.bd_5} shows the individual images from all four IRAC bands plus the MIPS 24 \\mum\\ image. The 24 \\mum\\ emission is roughly $2\\farcm7\\times0\\farcm9$ or $\\sim1.34\\times0.45$ pc, assuming the Cygnus-X distance of 1.7 kpc. The central star is in the Bonner Durchmusterung (Argelander 1903), Tycho (H\\o g et al. 2000), 2MASS (Cutri et al. 2003), and {\\em Midcourse Space Experiment} (\\msx; Egan et al. 2003) point source catalogs. However, no references discussing its physical properties could be found in the literature. The visual magnitude did change from a reported 9.2 mag in the Bonner Durchmusterung (1903) to 10.1 mag in the early 1990s as measured by Tycho. Other than indicating variability in this object, not an unusual feature in a carbon star, such a sparse sampling of the temporal baseline (albeit long) does not allow us to say anything meaningful about the pulsation properties of the star. Figure \\ref{fig.sed} shows the spectral energy distribution for BD+43 3710, as well as the other two sources. The smooth curves are Planck curve fits to the available photometry. A single temperature graybody does not well represent the emission from any of these sources, although the BVR data are probably suppressed by an unknown amount of extinction. As mentioned above, we obtained low-resolution IRS data on BD+43 3710 and one position in each lobe. Figure \\ref{fig.irs_star} (left) shows the locations where the spectra were taken. The right side shows the resulting spectrum for the central star. The spectrum is largely featureless. In comparison to the naked stars in the ISOSWS atlas of Sloan et al. (2003), it is most similar to HD 19557, a carbon star of type R noted by Goebel et al. (1983) as having particularly weak carbon features in visible and near-IR bands. In particular, the inflection in the IRS spectrum at 9 \\micron\\ is well-matched in the SWS spectrum of HD 19557, although it is not apparent in most of the naked carbon stars in the SWS atlas (class 1.NC of Kraemer et al. 2002) or in the other dust-free SWS spectral classes (1.N, 1.NO, 1.NE, and 1.NM). The IRS spectrum, however, does not show the turnover at $\\sim$5.5 \\micron\\ expected from C$_3$, CO, and CN absorption features (Goebel et al. 1978) typically observed in carbon stars (e.g. Aoki et al. 1998, Zijlstra et al. 2006). Several hydrogen recombination lines are detected in emission, as annotated in the figure. The strengths of the Pf $\\alpha$, Hu $\\alpha$, and HI 8-7 do not change between the raw data and the background-subtracted data, unlike the typical nebular lines such as the [Ne II] feature at 12.8 \\micron, which nearly vanishes after background subtraction. Thus, the hydrogen recombination lines are likely associated with BD+43 3710 and its nebula, and are not simply in the foreground or background cloud. Although not typical for carbon stars, Balmer lines have been detected in emission in the carbon star UV Aur A (Herbig 2009). Thus, we find that the carbon star spectral classification for BD+43 3710 is supported by the IRS spectrum but not definitively confirmed. The spectra from the two lobes, shown in Figure \\ref{fig.irs_ew}, are quite similar to each other, with a cool dust spectrum that rises with wavelength past the end of the IRS band and a few broad features and low-excitation fine structure lines. The SED of the western lobe rises a bit more steeply with wavelength than that of the eastern lobe, possibly indicative of slightly warmer dust. Neither spectrum is consistent with a single temperature graybody, but the majority of the dust must be cooler than $\\sim$100 K since the SED peaks beyond 35 \\mum. A 108 K graybody is the ``best'' fit to the LL data for both lobes and is shown in the figure for comparison. As can be seen from the spectra, unlike in a small number of sources where the circumstellar emission in the MIPS 24 \\mum\\ band has been found to arise from a high excitation [O IV] line (e.g. Morris et al. 2004 (a WN star), Morris et al. 2006 (a SNR candidate), Billot et al. 2009 (sources of unknown nature, possibly SNR or PNe candidates), Flagey et al. in preparation), or the $\\sim26-30$ \\mum\\ MgS feature seen in some post-asymptotic giant branch stars and planetary nebulae (e.g. Forrest et al. 1981; Goebel \\& Moseley 1985; Hony et al. 2002; Bernard-Salas et al. 2009), the nebular emission around BD+43 3710 arises from small dust grains. They are probably carbon-rich grains, dominated by amorphous carbon, given the lack of either a silicate emission feature around 9.7 \\mum\\ or silicon carbide around 11 \\mum. The broad features present in the spectra at 6.3, 7-8, 11.3, and 16-18 \\mum, typically attributed to polycyclic aromatic hydrocarbons (PAHs), are not likely to be from the BD+43 3710 nebula since there is no corresponding nebulosity in the IRAC 5.8 or 8.0 \\mum\\ image that can be readily associated with the 24 \\mum\\ lobes (Fig. \\ref{fig.bd_5}). The PAHs and low energy fine structure lines seen in Figure \\ref{fig.irs_ew} are extremely common in active star forming regions such as Cygnus-X, so those features likely arise from the fore/background Cygnus-X cloud. \\subsection{HBHA 4202-22} HBHA 4202-22 is located on the western edge of the survey region. Kohoutek \\& Wehmeyer (1999) report H$\\alpha$ emission from the star, citing ``Dolidze (1971) No. 3'' which probably corresponds to Dolidze (1971). This may be where the WR candidacy given in SIMBAD comes from although we could not confirm this due to the unavailability of the Russian circular, and it is otherwise unnoted in the literature. Figure \\ref{fig.hb_3c} shows a three-color image from the 3.6, 5.8, and 24 \\mum\\ data, and the individual images are in Figure \\ref{fig.hb_all}. We cannot say definitively there is no emission from PAHs, as we do not have the 8 \\mum\\ data which would contain the (typically) strongest feature. However, there is little extended emission in the 5.8 \\mum\\ data that would have contained the 6.2 \\mum\\ feature if it were present at any significant strength. Figure \\ref{fig.hbslice} shows a slice through the ring at position angle 30$\\arcdeg$, which corresponds to the narrowest point in the ring. The sharper, southwest edge of the ring is $\\sim1\\arcmin$ from the central star, or about 1.9 pc, assuming a distance of 6.5 kpc (see modeling discussion below). The ring is more extended to the northeast, $\\sim1\\farcm2$, or about 2.3 pc, and the edge is less well-defined. Two possible explanations come to mind. First, the surrounding medium could be less dense to the northeast, allowing the shell to expand more easily in that direction. However, the 5.8 \\mum\\ data indicate that if anything, there is more material to the east-northeast than toward the southwest. A second possibility is that the star is moving toward the southwest and the material on the ``front'' side of the shell is being compressed. Indeed, the proper motion of the HBHA 4202-22 is $\\mu_{\\alpha}=-4.5$ mas yr$^{-1}$, $\\mu_{\\delta}=-2.7$ mas yr$^{-1}$ (Zacharias et al. 2004), which corresponds to a position angle of 31\\arcdeg, nicely consistent with the compression direction. The distance to this object combined with the proper motion suggest that it may be a runaway star. We used a modified version of the radiative transfer code of Egan, Leung, \\& Spagna (1988) to model the dust shells around HBHA 4202-22. We modeled the central star as an A0 supergiant with T$_*$=10,000 K, L$_*$=$10^5$ \\lsun\\ at a distance of 6.5 kpc, i.e., well behind the Cygnus-X region. While these parameters are formal inputs to the model, they, too, were varied in order to get the best fit to the observed data. As noted above, the shell around HBHA 4202-22 is compressed toward the southwest compared to the northeast. Therefore, we fit two models to the two portions. There is also a hint of a ring at $\\sim35$\\arcsec\\ in the 5.8 \\micron\\ data, so an inner shell is also included in the model. Table \\ref{tab.hbhac} gives the derived parameters using amorphous carbon, with a particle size of 0.135 \\micron, an opacity of $\\tau_{8.035~\\micron}=7.5\\times10^{-6}$ ($\\tau_{0.55~\\micron}= 1.77\\times10^{-2}$), and grain constants from Mathis \\& Whiffen (1988). Figures \\ref{fig.hbhamodel} and \\ref{fig.hbhaarcs} show the results of the model compared to the data for the primary shells (compressed and not) at 24 \\micron. Models using silicate-rich dust were also considered, which result in cooler dust grains. The match to the data was not as good, though, so only the carbon-rich results are presented in Table \\ref{tab.hbhac}. \\subsection{G79.29+0.46} Figure \\ref{fig.g79_3c} shows the three-color image of G79.29+0.46 from the 3.6, 8, and 24 \\mum\\ data. It was first found to have an $\\sim$4\\arcmin\\ circumstellar ring in the radio by Higgs et al. (1993), who also suggested that the central object is probably a luminous blue variable (LBV). Subsequently, the same ring was detected in the infrared at a number of wavelengths from 8 to 60 \\mum\\ with the {\\em Infrared Astronomical Satellite} ({\\em IRAS}; Waters et al. 1996), the {\\em Infrared Space Observatory} ({\\em ISO}; Wendker et al. 1998), and {\\em MSX} (Egan et al. 2002, Clark et al. 2003). This inner ring is detected with all four IRAC bands as well as the 24 \\mum\\ MIPS band. Figure \\ref{fig.g79_5} shows the individual images from all five \\spi\\ bands. The apparent break in the south of the ring, particularly in the IRAC bands, is actually due to an infrared dark cloud (IRDC) that extends south-southwest in the figures. This suggests that as with HBHA 4202-22, G79.29+0.46 may be more distant than the 1.7 kpc assumed for the Cygnus-X complex. Interestingly, though, the CO observations of Rizzo et al. (2008) also show a break in the CO emission in roughly the same position as the IRDC. Since IRDCs are not known to absorb molecular line emission in the submillimeter, this suggests that the two structures might actually be interacting in some fashion. Additionally, the infrared ring appears to be somewhat flattened in the southeast, also consistent with possible interaction. Also visible in the 24 \\mum\\ data in Figures \\ref{fig.g79_3c} and \\ref{fig.g79_5} is a newly discovered outer ring $\\sim7\\arcmin$ across. This large ring probably represents an earlier episode of mass loss from G79.29+0.46 compared to the previously known inner ring. Additionally, the 24 \\mum\\ emission from the inner ring is much more extended than that from the shorter wavelengths. This kind of structure, where the IRAC emission ring is interior to a larger round emission structure at 24 \\mum, is occasionally seen in the MIPSGAL catalog of 24 \\micron\\ rings and disks (Mizuno et al. 2010). Usually, the 8 \\micron\\ emission appears to be co-spatial with the 24 \\micron\\ emission, if present, but most often it is absent entirely (as may be the case with HBHA 4202-22). As with HBHA 4202-22, we modeled the dust shells for G79.29+0.46. The stellar parameters are $T_*=18,000$ K, $L_*=4\\times10^5$ \\lsun, $R_*=4.567\\times10^{12}$ cm, with a distance of $d=9.25\\times10^{21}$ cm, i.e. slightly behind Cygnus-X at $\\sim$3 kpc. Table \\ref{tab.g79} lists the derived parameters for the inner and outer shells. A mix of amorphous carbon grains and PAHs (Li \\& Draine 2001) was used in the model, as the dust alone could not produce the flux detected in the IRAC bands and the strong PAH emission seen in the IRS observations of Umana et al. (in preparation)\\footnote{Morris et al. (2008) present an infrared spectrum of G79.29+0.46 that included IRS data as well as ISOSWS and near-infrared spectra. Presumably the IRS spectrum is from the central star as it does not show the strong PAHs that are detected by Umana et al. toward the shell.}. The amorphous carbon grains dominate the mass and opacity, although the (much) smaller PAHs are more numerous. While the inner shell is modeled as a single shell with a density law of $r^{-3.5}$, the 24 \\micron\\ flux profile (Fig. \\ref{fig.g79_rad}) suggests that a more complex model with a different density law for the interior of the shell might give better results." }, "1003/1003.2565_arXiv.txt": { "abstract": " ", "introduction": "The importance of studying faint outskirts of galaxies for our understanding of the galaxy formation and evolution has become increasingly apparent in recent years. Due to their long dynamical timescales, galactic outer regions have retained the fossil record from the epoch of galaxy assembly in the form of spatial distribution, kinematics, ages and metallicities of their stars \\citep{bullockjohnston05,freemanbh02}. In addition, secular evolution of spirals leaves the most conspicuous clues in the low density regions in galactic outskirts \\citep{roskar08a,roskar08b,schoenrich08,sanchezblazquez09,martinezserrano09}, providing an opportunity for testing scenarios of galaxy evolution using observations of outer disks of spirals. In this paper we look into recent advances in observational studies and numerical simulations of spiral disks, and the new insights they have provided into the processes of disk galaxy formation and evolution. ", "conclusions": "Thanks to a confluence of exciting new results from both observational and theoretical work, outer spiral disks have in recent years become a fast-growing research area. While more deep stellar photometry is necessary in order to test predictions of a growing body of numerical simulations of disk formation and evolution, our current understanding can be summarized as follows: (i) Origin of the diversity in outer disk structure remains a puzzle. While scenarios have been proposed to explain sub- and super-exponential profiles, the full picture which self-consistently explains existence of all three types of light profiles is lacking. (ii) Radial migrations potentially play a significant role in the evolution of disk galaxies; this could have profound consequences on how we model galaxy evolution, in particular chemical evolution in spirals. (iii) Star counts have been shown to be a superior method for probing faint outer disks in individual galaxies compared to traditional surface photometry. (iv) Age behavior in disks is very challenging to determine from resolved stellar photometry. However, a small number of studies seems to indicate that in case of galaxies with sub-exponential profiles, the minimum stellar age is observed at the break radius, in agreement with simulations of radial mixing in spirals. (v) Metallicity gradient which flattens in the outermost regions seems to be a general feature of spiral disks. However, some spirals experience single-slope negative gradient with no flattening." }, "1003/1003.5942_arXiv.txt": { "abstract": "We have analyzed photometry from space- and ground-based cameras to identify {\\it all} bright red giant branch (RGB), horizontal branch (HB), and asymptotic giant branch (AGB) stars within 10\\arcmin ~ of the center of the globular cluster M13. We identify a modest (7\\%) population of HB stars redder than the primary peak (including RR Lyrae variables at the blue end of the instability strip) that is somewhat more concentrated to the cluster core than the rest of the evolved stars. We find support for the idea that they are noticeably evolved and in the late stages of depleting helium in their cores. This resolves a disagreement between distance moduli derived from the tip of the red giant branch and from stars in or near the RR Lyrae instability strip. We identified disagreements between HB model sets on whether stars with $\\teff \\la 10000$ K (near the ``knee'' of the horizontal branch in optical CMDs) should evolve redward or blueward, and the differences may depend on the inclusion of diffusion in the stellar interior. The sharp cut at the red end of M13's HB provides strong evidence that stars from the dominant HB group must still be undergoing blue loops, which implies that diffusion is being inhibited. We argue that M13's HB is a somewhat pathological case --- the dominant HB population occurs very near the ``knee'' in optical CMDs, and evolved stars exclusively appear redward of that peak, leading to the incorrect appearance of a continuation of the unevolved HB. We identify two stars as ``blue hook'' star candidates --- the faintest stars in optical bands that remain significantly subluminous in the shortest ultraviolet wavelength photometry available. M13 also has a distinct group of stars previously identified with the ``second $U$ jump''. Based on far UV photometry, we find that these stars have genuinely high temperatures (probably 26000 K $\\la \\teff \\la 31000$ K), and are not produced by a jump in brightness at lower temperature ($\\teff \\approx 22000$ K) as previously suggested. These stars are brighter than other stars of similar color (either redder or bluer), and may be examples of ``early hot flashers'' that ignite core helium fusion shortly after leaving the red giant branch. We used ultraviolet photometry to identify hot post-HB stars, and based on their numbers (relative to canonical AGB stars) we estimate the position on the HB where the morphology of the post-HB tracks change to $I \\sim 17.3$, between the two peaks in the HB distribution. Concerning the possibility of helium enrichment in M13, we revisited the helium-sensitive $R$ ratio, applying a new method for correcting star counts for larger lifetimes of hot horizontal branch stars. We find that M13's $R$ ratio is in agreement with theoretical values for primordial helium abundance $Y_P = 0.245$ and inconsistent with a helium enhancement $\\Delta Y = 0.04$. The brightness of the horizontal branch (both in comparison to the end of the canonical HB and to the tip of the red giant branch) also appears to rule out the idea that the envelopes of the reddest HB stars have been significantly enriched in helium. The absolute colors of the turnoffs of M3 and M13 may potentially be used to look for differences in their mean helium abundances, but there are inconsistencies in current datasets between colors using different filters that prevent a solid conclusion. The numbers of stars on the lower red giant branch and in the red giant bump agree very well with recent theoretical models, although there are slight indications of a deficit of red giant stars above the bump. There is not convincing evidence that a large fraction of stars leave the RGB before undergoing a core helium flash. ", "introduction": "M13 was one of the first globular clusters identified as having unusually blue horizontal branch (HB) stars, and it remains one of the prototypes of the ``long blue tail'' with stars approaching the main sequence for helium stars. As \\citet{smith} notes, differences between stars on both giant branches and stars on the horizontal branch can be discerned from data in papers as early as \\citet{barn09,barn14}. Along with the nearby, massive, and little reddened cluster M3, M13 forms half of the best known ``second parameter'' pair. The chemical composition of the hydrogen envelopes of HB stars affects their observable properties via opacity and mean molecular weight. The ``first parameter'' is heavy element content, where higher metallicity produces higher envelope opacity and generally redder stars. Because M3 and M13 have nearly the same iron abundances ($\\langle$[Fe/H]$\\rangle = -1.53$ for M13 versus $\\langle$[Fe/H]$\\rangle = -1.45$ for M3; \\citealt{sned}) but M3 has a much redder HB including a huge number of RR Lyrae variable stars, a second parameter is needed. In addition to the color shift between the HBs of M3 and M13, the HB stars in M13 show a bimodal distribution that is not present in M3. On its own, this fact implies that the HB stars were produced by at least two different populations of cluster stars or involve different methods of producing HB stars. M3 and M13 share a number of similarities beyond iron abundance, and we tabulate some of their characteristics in Table \\ref{m13m3}. Many theories have been proposed to explain the HB differences, and two of the main goals of this paper are to 1) assemble a large and complete set of photometric data for M13, and 2) use the photometric data to examine questions bearing on the production of M13's horizontal branch stars. Because of the complexity of the HB, it is very doubtful that one explanation can cover all of its aspects. Before we describe our results, we briefly summarize the main hypotheses we will examining, and the primary reasons they are viable. We emphasize that they are not mutually exclusive. \\underline{The $\\Delta t$ Hypothesis.} Early models showed that as the mass of the hydrogen-rich envelope of an HB star is decreased, the surface temperature increases with relatively little change in luminosity. In this hypothesis, age differences between populations of stars lead to differences in mass between stars leaving the main sequence and between stars reaching the HB. While it is natural to expect that clusters in the Milky Way were born at different times, age differences are hard to prove except for clusters that are much younger than the average. \\citet{rey} compared M13 with M3 and found that turnoff-to-giant branch color differences and changes in HB morphology were consistent with an age difference $\\Delta t = 1.7 \\pm 0.7$ Gyr (with M13 older). \\citet{catrev} finds that the age differences implied by differences in the cluster CMDs near the turnoff can explain the HB morphology as long as M3 is younger than about 12 Gyr and the EHB stars in M13 are presumed to arise from a process that is unrelated to the age. However, there are aspects of M13's population that cannot be explained in this hypothesis. For example, neither of the studies above could reproduce the bluest HB stars in M13 in synthetic HB simulations using the same chemical composition and dispersion in stellar mass used for M3. \\underline{The $\\Delta Y$ Hypothesis.} Variations in helium content ($Y$) result in differences in position on the HB, largely because greater helium abundance allows lower mass stars to leave the main sequence at the present day \\citep{dant08b}. \\citet{jb} proposed that a helium abundance difference $\\Delta Y \\sim 0.05$ (with M13 the more helium rich) could be responsible for many of the unusual features of the color-magnitude diagram (CMD), including an interesting difference in the slopes of the subgiant branch. \\citet{cda} also examined data for M3 and M13, finding the luminosity of the red giant bump and RR Lyrae stars relative to the cluster turnoff are consistent with an enhancement $\\Delta Y \\sim 0.04$. This picture has been taken very seriously with the discovery of multiple stellar populations in some clusters. For example, NGC 2808 was found to have at least three identifiable main sequences \\citep{pio07}, while $\\omega$ Cen has a blue main sequence \\citep{bed04} that appears to be helium enriched \\citep{pio05}. $\\omega$ Cen and NGC 2808 are among the most massive clusters known in the Milky Way, which may enable them to retain gas that has been processed and released by a first generation of stars. While the helium abundance is very difficult to measure except in limited circumstances, spectroscopic observations of other heavy element species lead to the belief that helium was probably enriched in some clusters. Stars in M13 are well-known to have star-to-star abundance differences in O and Na that can be traced from the giant branch \\citep{sbcn,yong,jkp,sbh,sned} to the main sequence turnoff \\citep{cm,briley}. O depletion and Na enrichment can only be accomplished in hydrogen-fusion regions where significant production of helium is accomplished \\citep{dandd,lang}, and star-to-star variations on the main sequence require that they must have been present in the gas forming the stars. The $\\Delta Y$ hypothesis is attractive for M13 because it may explain the blueward shift of the main body of HB stars compared to M3's population, and the bimodality of the HB (as an additional population within the cluster). \\citet{dant08b} conducted a fit to M13's HB using helium-enriched models, and found that a fit required 70\\% of the population to be enriched to $0.27 < Y < 0.35$ (a mean $\\Delta Y \\approx 0.04$), and the remaining 30\\% to be enriched to $Y \\sim 0.38$. There are some difficulties with this picture though. In the \\citeauthor{dant08b} models of M13, they still needed to assume a rather large (but constant) total mass loss on the RGB ($0.18 \\msun$). Because M13 has few HB stars in the instability strip or redward (where M3's HB is heavily populated), {\\it virtually all} of M13's HB stars must also be more helium rich than M3's. This is in striking contrast to more massive clusters that show strongly bimodal HB star distributions in which the redder HB population is interpreted as a first generation of stars formed from primordial material, while subsequent generations have varying degrees of enrichment and are bluer. The massive clusters that are inferred to have such large spreads in $Y$ generally also have multiple main sequence or subgiant branch populations, whereas M13 has shown no sign of multiple populations to date. There is also not a clear bimodality in the spectroscopic abundances and a definite connection has never been made between the abundances and HB morphology \\citep[although see][for indications that the maximum $\\teff$ extent of the HB is correlated with the extent of observed Na-O anticorrelations]{carrb}. \\citet{rey} attribute some of \\citeauthor{jb}'s conclusions to slight missteps in the implementation of their relative age comparison. \\underline{The $\\dot{M}$ Hypothesis.} From early models, it was recognized that a significant amount of mass loss is needed, probably on the red giant branch (RGB), to produce the colors of the majority HB populations in most clusters. Dispersion in colors was then taken to mean that there are star-to-star differences in mass loss, although a mechanism to produce these differences has not been identified. Independent of the majority of HB stars, there is a population that seems to {\\it require} strong mass loss: the ``blue hook'' stars \\citep{cast93,dcruzbhk,cast06,millb}. In ultraviolet color-magnitude diagrams of some of the most massive clusters \\citep{dcruzomega,brown}, stars are found fainter and redder than the zero-age HB (ZAHB) at its blue end, meaning that they must have almost no hydrogen envelope. If a star loses virtually all of its hydrogen envelope before reaching the tip of the RGB, it can leave the RGB without igniting core helium fusion. As the star contracts onto the He white dwarf cooling curve, a late He flash can be ignited that drives a convection zone that reaches hydrogen rich layers \\citep{brown}. \\underline{The Evolution Hypothesis.} As relatively cool HB stars convert He into carbon and oxygen, they are expected to eventually evolve brightward and redward toward the asymptotic giant branch (AGB). Depending on the distribution of stars on the HB, evolving HB stars could be mistaken for fainter, more slowly evolving HB stars, thereby misrepresenting the brightness of the HB. Clusters with large blue HB populations (like M13) are most susceptible to this effect because evolutionary tracks may nearly parallel the ZAHB at the blue end of the HB distribution where the relative number of unevolved stars drops rapidly. Because the HB is a frequently used standard candle in astronomy, it is worth studying the degree to which this affects stellar populations. Evolutionary effects have been inferred from the pulsation properties of RR Lyrae stars. For example, \\citet{jurc} used magnitudes and periods to identify RR Lyraes in different stages of their HB evolution. \\citet{caccm3} identified mean lines in the period-amplitude diagram for different subsets of M3 variables, and labeled them as regular or ``well-evolved''. They showed that in at least some other clusters, the majority of variables could be identified with one group or the other. For the purposes of this paper we focused on post main-sequence evolution. Using datasets from telescopes and instruments having a wide range in spatial resolution and field size, we attempted to completely survey the evolved stars from the center far into the outskirts of the cluster. We discuss the observational material and the analysis of the photometry in \\S 2. In \\S 3, we describe the steps used to identify the evolutionary status of the evolved cluster stars. We examine the red giant branch and horizontal branch populations in greater detail in \\S 4 and \\S 5, respectively. In \\S \\ref{ratios} we look at population ratios and their relationship to the evolution timescales for stars in different stages. Finally in \\S \\ref{disc} we discuss the body of evidence involving second parameter effects (cluster to cluster variations) and intracluster differences between stars. ", "conclusions": "To our minds, some of the most important questions regarding M13 remain in dispute. One question that we have reopened here is whether the reddest of the blue HB stars in M13 are significantly evolved (and whether they are therefore good representatives of the brightness of the horizontal branch). The weight of the observational and theoretical evidence leans toward the idea that they are significantly evolved, and that the red edge of the primary HB population is a decent indicator of the true HB level. The distribution of stars on the horizontal branch in M13 is complex, and the most notable questions regard 1) how the color of the primary peak in M13 could have been shifted so far relative to M3's when the gross composition of clusters appear nearly identical, and 2) how a large fraction of M13's RGB stars become blue stars near the end of the canonical HB. Our examination of the luminosity function shows little sign that a large fraction of stars leave the bright RGB before having a core flash, in agreement with the massive cluster NGC 2419 \\citep{sh} but not with NGC 2808 \\citep{sm}. We do not find any clear evidence of helium enrichment among the stars of the dominant (redder) HB population, and in fact, the helium abundance indicator $R$ and the relative brightness of the HB argue against significant enrichment. The HB and RGB stars (and different subsets of these) do not show significant signs of radial segregation within the cluster. The small color difference between the main sequence turnoff and the giant branch of M13 (in comparison to M3) remains unexplained, but careful examination of the absolute colors of both clusters would provide a new test. Our thorough search of M13's HB population has revealed second $U$ jump and blue hook stars that imply that many of these stars have very low-mass hydrogen-rich envelopes. Far UV observations show that many of the stars in the second $U$ jump are more luminous than stars with similar colors at the end of the HB. The reason is unclear, however. Spectroscopic data on similar stars in NGC 6752 \\citep{moni} indicate that the excess brightness is not related to enhanced atmospheric helium abundance, so further study is required. Spectroscopic measurements may help to clarify our understanding of extreme HB stars in a number of ways. We particularly encourage studies of: stars near the red end of the HB in M13 where helium abundances can be accurately determined \\citep{villa}; stars in the extreme HB to look for signs of unusual Mg abundances (a species that appears to be minimally affected by diffusion) that could connect them to giant stars; O, Na, and Mg for stars at the red giant tip of other clusters to determine whether they are super O-poor; relatively unevolved turnoff and subgiant stars in M13 and NGC 2808 to search for large O depletions ([O/Fe]$ < -0.4$) and check whether this is the result of external pollution or not." }, "1003/1003.2423_arXiv.txt": { "abstract": "We present and discuss \\emph{Spitzer} and near-infrared H$_{2}$ observations of a new bi-polar protostellar outflow in the Rosette Molecular Cloud. The outflow is seen in all four IRAC bands and partially as diffuse emission in the MIPS 24 $\\mu$m band. An embedded MIPS 24 $\\mu$m source bisects the outflow and appears to be the driving source. This source is coincident with a dark patch seen in absorption in the 8 $\\mu$m IRAC image. \\emph{Spitzer} IRAC color analysis of the shocked emission was performed from which thermal and column density maps of the outflow were constructed. Narrow-band near-infrared (NIR) images of the flow reveal H$_2$ emission features coincident with the high temperature regions of the outflow. This outflow has now been given the designation MHO 1321 due to the detection of NIR H$_2$ features. We use these data and maps to probe the physical conditions and structure of the flow. ", "introduction": "Outflows and jets from young stellar objects (YSOs) accompany the early stages of star formation. Outflows can manifest themselves as jets and knots of shocked material visible at optical and near-infrared wavelengths and also molecular emission observable at longer wavelengths. The outflowing material plays a role in removing the excess angular momentum from the YSOs allowing them to evolve into stars. Outflows are able to trace the history of mass loss and accretion of their driving sources. Studying the structure and properties of these flows may provide clues to understanding the connection between jets and the associated wide angle molecular flows \\citep{rei2001}. Additionally, this outflowing material interacts with its surroundings and may affect its environment, possibly regulating further star formation and cluster evolution. The energy and momentum inputted by outflows may disrupt the surrounding ambient gas, contribute to the turbulence in the cloud, and affect chemical processes \\citep{bal2007}. \\citet{yba2009} developed a technique to study the thermal structure of shocked H$_2$ gas using color analysis of observations from the \\emph{Spitzer} InfraRed Array Camera (IRAC). Given the vast amount of \\emph{Spitzer} data available, this technique can be used to survey large regions and simultaneously find and analyze shocked emission. The IRAC color analysis enables the construction of temperature maps of the shocked gas which may in turn be used to probe the interaction of outflow with its surroundings. These maps may also be used to compare the properties of outflow with those of simulations allowing a better understanding of the physics involved and estimating the energy and momentum inputted by outflows into their environment. The Rosette Molecular Cloud (RMC) is a star forming region located at a distance of 1.6 kpc. Near-infrared imaging studies have revealed nine embedded clusters across the cloud \\citep{phe1997,rom2008}. Outflow activity in the cloud has been revealed through the [\\ion{S}{2}] narrowband imaging survey of \\citet{yba2004} and the $^{12}$CO survey of \\citet{den2009}. In an analysis of the \\emph{Spitzer} IRAC images of the Rosette Molecular Cloud, we have discovered a structure with the morphology of a bipolar outflow that is visible in the images from all four IRAC bands. This structure can be seen in the images published by \\citet{pou2008} although it is not discussed in their paper. In this study we analyze the outflow using near infrared (NIR) narrowband imaging of the flow to confirm the presence of shocked gas inferred from analysis of the the IRAC images. We improve the IRAC color analysis of \\citet{yba2009} and use it to create temperature and column density maps of the outflow. Using both the NIR and IRAC data, we probe the physical conditions and structure of the outflow. \\begin{figure*} \\plotone{f1lr.eps} \\caption{ Spitzer IRAC images of the outflow. The origin is set at ($\\alpha$,$\\delta$)(J2000) = ($06^{\\rm{h}}35^{\\rm{m}}25\\fs 0$, $+03\\arcdeg56\\arcmin21\\arcsec$) } \\end{figure*} ", "conclusions": "We present the discovery of a new bi-polar outflow in the Rosette Molecular Cloud and use NIR narrowband and Spitzer imaging data to study the flow. We show that IRAC color analysis can be used to interpret the interaction of an outflow with its surrounding environment. Using our calculations of the IRAC space of non-dissociative shocked gas we fit analytic forms to the color-temperature and column density-temperature relationships. We verify that IRAC color analysis can reveal regions of shocked gas and find that the NIR H$_{2}$ knots correspond to regions of high temperature and or column density determined through color analysis. We find diffuse MIPS 24 $\\mu$m emission, most likely from [\\ion{Fe}{2}] lines, to be coincident with regions of high temperature thus confirming the validity of using the non-dissociative shock IRAC color space The NIR line ratios combined with the temperature estimates allow for the determination of extinction along the line of sight which is used to create a column density map of the shocked H$_{2}$ gas. We deduce that the asymmetry in the outflow is due to interactions with the dense material to the west of the outflow source causing deflection and possibly deceleration of the outflowing material." }, "1003/1003.5035_arXiv.txt": { "abstract": "{} {The objective of this study was to examine the kinematics of coronal mass ejections (CMEs) using EUV and coronagraph images, and to make a quantitative comparison with a number of theoretical models. One particular aim was to investigate the acceleration profile of CMEs in the low corona. } {We selected two CME events for this study, which occurred on 2006 December 17 (CME06) and 2007 December 31 (CME07). CME06 was observed using the EIT and LASCO instruments on-board SOHO, while CME07 was observed using the SECCHI imaging suite on STEREO. The first step of the analysis was to track the motion of each CME front and derive its velocity and acceleration. We then compared the observational kinematics, along with the information of the associated X-ray emissions from GOES and RHESSI, with the kinematics proposed by three CME models (catastrophe, breakout and toroidal instability). } {We found that CME06 lasted over eight hours while CME07 released its energy in less than three hours. After the eruption, both CMEs were briefly slowed down before being accelerated again. The peak accelerations during the re-acceleration phase coincided with the peak soft X-ray emissions for both CMEs. Their values were $\\sim$60~m~s$^{-2}$ for CME06 and $\\sim$600~m~s$^{-2}$ for CME07. CME07 reached a maximum speed of over 1000~km~s$^{-1}$ before being slowed down to propagate away at a constant, final speed of $\\sim$700~km~s$^{-1}$. CME06 did not reach a constant speed but was moving at a small acceleration by the end of the observation. Our comparison with the theories suggested that CME06 can be best described by a hybrid of the catastrophe model and breakout model while the characteristics of CME07 were most consistent with the breakout model. Based on the catastrophe model, we deduced that the reconnection rate in the current sheet for CME06 was intermediate, the onset of its eruption occurred at a height of $\\sim$200~Mm, and the Alfv\\'en speed and the magnetic field strength at this height were approximately 130--250~km~s$^{-1}$ and 7~Gauss, respectively.} {} ", "introduction": "Coronal mass ejections (CMEs) are the ejections of large amount of mass and magnetic flux from the Sun to interplanetary space. The energy released during the process is of the order of $10^{32} - 10^{33}$ erg. Statistical studies have reported variations in the observed properties and projected kinematics of CMEs \\citep[e.g.,][]{CB2004A&A,Zhang_etal2004ApJ} Nevertheless, CMEs often rise with a small initial speed of a few kilometers per second, followed by a rapid upward expansion reaching several hundred or several thousand kilometers per second, and, eventually, propagate through the interplanetary space at the ambient solar-wind speed. Several models have been proposed to explain the driving mechanism and observed properties of the CMEs. A valid model must be able to produce the observed kinematics, dynamics and properties of CMEs. In Sec.~\\ref{sec:model}, we give a review of three representative models: breakout model (BO) \\citep{Breakout1999ApJ}, catastrophe model (CA) \\citep[see, e.g.,][]{VanTend_Kuperus1978SoPh,FI1991ApJ,FP1995ApJ,Lin_etal1998ApJ}, and toroidal instability model (TI) \\citep[see, e.g.,][]{Chen_1989ApJ,KT2006PhRvL}, which have been investigated in this study. Although the profile of the CME kinematics predicted by each model is qualitatively consistent with the observed three-phase profile (i.e., slow rise, eruption and slow down), the exact profile of each model is different. Several simulations also showed that a same model can produce different kinematic profiles by simply varying the initial conditions and magnetic environment \\citep{Lynch_etal2008ApJ,Schrijver_etal2008ApJ,Chen_1989ApJ}. In addition to the discrepancies among theoretical predictions, observational studies that were not based on any model, also reported different functional forms that best-fitted their respective data. For instance, \\citet{Sheeley_etal1999} fitted their data by $r(t)=r_0 + 2 r_0 \\ln \\cosh \\left[ v_a(t + \\Delta t_0)/2 r_a \\right]$; \\citet{Gallagher_etal2003ApJ} formulated a double exponential function $a(t)=\\left[a^{-1}_r \\exp(-t/\\tau_r) + a^{-1}_d \\exp(t/\\tau_d)\\right]^{-1}$ to describe the fast rising and decaying acceleration profile; and \\citet{AMN_2002GeoRL} demonstrated that their fast CME was best fitted by the polynomial, $h(t)= h_0 + v_0t + c t^m$, with $m \\approx 3.7$, which is consistent with the results by \\citet{Schrijver_etal2008ApJ}. Although these studies were able to derive how the velocity and acceleration evolve through the duration of CMEs, they did not provide a physical explanation and driving mechanism behind such kinematic profiles. Besides, while most of these studies chose polynomials or exponential functions or a simple combination of the two, there is no guarantee that the chosen function is the only one that can fit the observational profile. Studies to verify a specific model by comparing observations with models have also been carried out. For example, by selecting only those that can be characterized as flux-rope type CMEs, \\citet{Krall_etal2001ApJ} showed that the kinematic curves predicted by the TI model matched the kinematics in the higher corona but not the lower coronal region. An extensive examination of a fast CME by \\citet{MK2003ApJ} based on the data from various instruments showed supportive evidence for the BO model. However, \\citet{Bong_etal2006ApJ} reported that while the magnetic-field configuration and X-ray loop field connectivity in their data indicated a breakout process, the acceleration profile they obtained contradicted the BO model. The studies to verify the CA model \\citep[e.g.,][]{AMN_2002GeoRL,Schrijver_etal2008ApJ}, however, often deduced their conclusions by comparing the early-stage expression of CA model with the acceleration phase and/or the late stage of the observed CME kinematic curves. As already been demonstrated in all observational studies, the kinematics of CMEs can change greatly through the duration of the event. Hence, it is unsurprising that the functional form for the early stage does not match with the profiles of the later stages. The objectives of this paper are first to examine the features and kinematics of CME events, and then to determine the possible mechanism behind the observed phenomena by comparing the observations with various theories. The theory most consistent with the observation can thus provide information on the magnetic environment where the eruption occurs. We have chosen two CME events for this study. We first derived their kinematic profiles and examined the associated X-ray emissions and flares, and then compared the results with those derived from the three different models. The kinematic profiles from the models have been mainly obtained from simulations. The only stage of which an analytical expression of the kinematics is available is the early stage of the eruption. Hence, we first focused on a qualitative comparison of the characteristics between the observations and different models, and then carried out an additional quantitative examination by fitting the observational data with a number of model expressions. The objectives of the quantitative examination were to verify the assumptions employed in the models to derive these expressions, and/or to obtain an empirical expression for the observed kinematics. The values of the fitting parameters may be used to infer the magnetic environment and the initial conditions at the onset of the eruption. The rest of the paper is organized as follows: the three models are reviewed in Sec.~\\ref{sec:model}. The two selected CME events and how their motions were traced are described in Sec.~\\ref{sec:observation}. The qualitative and quantitative comparisons with the models are explained and discussed in Secs.~\\ref{sec:exam_qual} and \\ref{sec:exam_quant}. We conclude our study in Sec.~\\ref{sec:conclusion}. ", "conclusions": "\\label{sec:conclusion} One of the main goals of this study was to investigate the possible driving mechanisms behind coronal mass ejections. We have analyzed, both qualitatively and quantitatively, two CME events, and compared them with three CME models for this purpose. The first CME was observed on 2006 December 17 by EIT and LASCO, and the associated X-ray emissions were recorded by GOES and RHESSI. This CME lasted for over eight hours, and the SXR emission exhibited two rising phases. A comparison of the SXR emission profile and the kinematics of the CME showed that the CME was launched earlier than the rising of the SXR emissions, reached an acceleration peak during the second rising phase of the SXR emission. The CME was first decelerated after the eruption, and then re-accelerated to another acceleration peak of $\\sim$60~m~s$^{-2}$, at which time the SXR emissions reached the maximum and a flare was detected. The CME propagated away at a small, constant acceleration at the end. Comparing the above information with the three CME models discussed in this paper, CA, TI and BO, we propose that this CME can be best described by a combination of CA and BO models. That is, the CME was initiated by an ideal MHD process (CA), as hinted by the absence of SXR emissions during the initial stage. The rising and expanding of the flux rope against the ambient magnetic fields then led to reconnections with these fields (BO). As it continued to rise, a current sheet was formed underneath, which exerted a dragging force to slow down the rising motion, and, lastly, the current sheet was cut off by reconnections and the CME was re-accelerated. Our quantitative examinations also revealed that the CA-type expressions gave the best fit to the EIT data. From the fitting results, we deduced that the critical height at which the CME erupted was approximately 200~Mm and the Alfv\\'en speed at this height was around 130--250~km~s$^{-1}$, which leads to a magnetic field of about 7 Gauss. They are of reasonable order of magnitude for the low corona. Based on the model results by \\citet{LinForbes2000JGR}, we can also infer from the occurrence of the deceleration of this CME that the reconnection rate in the current sheet may correspond to a Alfv\\'en Mach number between 0.005 and 0.041. The second CME was observed on 2007 December 31. This CME was formed in a magnetic arcade system. Our analysis indicated that the initial eruption stage might have been obscured by the arcade system. After this un-detected eruption, the acceleration of the CME dropped first and then rose again to a peak value of $\\sim$600~m~s$^{-2}$. The CME remained at this peak acceleration until its velocity reached a peak value of $\\sim$1000~km~s$^{-1}$. The CME then decelerated and eventually propagated away at a constant speed of $\\sim$700~km~s$^{-1}$. A comparison between the kinematics of the CME and the SXR and HXR emission information from GOES, MESSENGER SAX, and RHESSI \\citep[cf., Fig.~\\ref{fig:kin_cme071231};][]{Raftery_etal2010apj,Krucker2009SPD,Krucker_etal2009ApJ} showed that the SXR emission began to rise before the emergence of the CME and peaked around the same time as the peak of the CME velocity, and that the HXR emission due to particle accelerations occurred at the time of the CME emergence and at the moment when the acceleration of the CME began to rise again. These observed features can be well described by a BO model as follows: the SXR emission before the CME emerged from the magnetic arcades and the high velocity and acceleration at the beginning of its emergence can be explained as that the breakout reconnections were happening as the CME loop was rising through and breaking the overlying arcade. The reconnections accelerated the CME, and caused the X-ray emissions. By the time the CME broke through the arcades, a trailing current sheet may have been formed to act as a drag force, which was reflected in the decrease of the acceleration. However, we note that some overlying features, such as streamers, could also slow down the CME. The subsequent rising of the acceleration (after 01:00~UT) and the detection of a HXR signal indicated that this drag force was being reduced by the reconnections in the current sheet. After the current sheet was completely cut off, the ambient coronal magnetic fields moved in to reform and settle to a lower potential state. The last stage of the CME indicated that it was being slowed down by the solar wind and/or interstellar medium. To be in agreement with a CA model, which proposes an ideal MHD process to initiate the eruption, a CME should have a beginning rising phase during which there was no X-ray emission. This is inconsistent with the observations of this CME, which detected flares and X-ray emissions from the very beginning of the eruption. However, since the initial rising stage of the CME was occulted, we cannot completely rule out the possibility that an ideal MHD process had taken place. The missing of this ideal MHD process in our data was also reflected in our unsuccessful fitting using CA-type expressions. Both events exhibited a re-acceleration phase, during which the maximum acceleration coincided with the peak SXR emission. If higher SXR emissions indicate greater reconnection rates, as explained in Sec.~\\ref{sec:mdl_CA}, such coincidence can be explained as that the faster reconnections, which lead to faster reduction of the tethering force, increased the acceleration of the CME. In conclusion, we have investigated two CME events. Our results showed that the first event, CME06, can be best described by a combination of the CA and BO models while the other event, CME07, can be well explained by the BO model alone. However, it is also possible that CME07 may have been initiated by an ideal MHD process. We conclude that TI model is the least likely driving mechanism for either of the events mainly because our observationally derived kinematic profiles do not match the profiles and expressions from the TI model. Specifically, the decreasing of acceleration in the early stage of our CME events is not seen in TI simulations to-date \\citep{KT2006PhRvL,TK2007AN328.743T,Schrijver_etal2008ApJ}. Since our data is a two-dimensional projection of the actual motion, we acknowledge the possibility that this phase be a projection effect. Hence, to improve the accuracy of our analysis, we plan to utilize stereoscopic data to obtain an actual, three-dimensional kinematics of CMEs in our future work. At last, to improve the current CME models, our study suggests that the CA model include the reconnections/interactions between the CME and its ambient magnetic fields, BO model consider the possibility of an ideal MHD process in the early stage, and that TI model provide possible mechanisms for a drag force during the eruptive phase." }, "1003/1003.1420_arXiv.txt": { "abstract": "Modern population synthesis models estimate that 50\\% of the restframe $K$-band light is produced by TP-AGB stars during the first Gyr of a stellar population, with a substantial fraction continuing to be produced by the TP-AGB over a Hubble time. Between 0.2 and 1.5 Gyr, intermediate mass stars evolve into TP-AGB C stars which, due to significant amounts of circumstellar dust, emit half their energy in the mid-IR. We combine these results using published mid-IR colors of Galactic TP-AGB M and C stars to construct simple models for exploring the contribution of the TP-AGB to 24$\\mu$m data as a function of stellar population age. We compare these empirical models with an ensemble of galaxies in the CDFS from $z=0$ to $z=2$, and with high quality imaging in M81. Within the uncertainties, the TP-AGB appears responsible for a substantial fraction of the mid-IR luminosities of galaxies from $z=0$ to $z=2$, the maximum redshift to which we can test our hypothesis, while, at the same time, our models reproduce much of the detailed structure observed in mid-IR imaging of M81. The mid-IR is a good diagnostic of star formation over timescales of $\\sim 1.5$ Gyr, but this implies that on-going star formation rates at $z=1$ may be overestimated by factors of $\\sim 1.5-6$, depending on the nature of star formation events. Our results, if confirmed through subsequent work, have strong implications for the star formation rate density of the universe and the growth of stellar mass over time. ", "introduction": "\\label{sec:introduction} The infrared provides a critical window into obscured star formation in our Galaxy and in nearby galaxies \\cite[e.g.][and many since]{helou1988}, allowing us to peer into a range of star forming environments at the present epoch \\cite[e.g.][]{roussel2001,calzetti2007}. Space-based IR observations have improved and extended our views to greater distances and large lookback times \\citep[e.g.][]{salim2009}, making these wavelengths crucial for studies of galaxy structure \\citep[see, e.g.,][]{regan2004} and evolution \\citep[e.g.][]{papovich2007}. The $24\\mu $m data, in particular, provided the first of what were to be unbiased histories of the star formation rate density of the universe over time \\citep{lefloch2005}. \\begin{figure*} \\centerline{ \\includegraphics[scale=0.37]{figure1.png}} \\caption{The mid-IR contributions of the TP-AGB. (a) Simplified schematic for the fraction of $K$-band light produced by the TP-AGB for an SSP. This approximation allows us, in tandem with observed near- to mid-IR colors of TP-AGB stars, allows us to estimate mid-IR luminosities. The blue region specifies the time over which stars can become C rich and emit half their energy in the mid-IR. Outside of this age range, we use the ensemble color of the Galactic TP-AGB M stars. (b) The correlation between $L_{24/(1+z)}$ and $L_K$ for galaxies in CDFS at $z\\sim 1$ \\citep[courtesy of][]{wuyts2008}, where those with $24\\mu$m flux densities significant at a level less than $2\\sigma$ are shown in black, less than $5\\sigma$ in green, and the rest in orange. Non-detections are shown as red points along the bottom of the figure. The maximum and minimum mid-IR luminosities of galaxies are defined by the ensemble colors of TP-AGB stars, shown by the blue and red lines, for young and old stellar populations, respectively. (c) Galaxies in CDFS at $z\\sim 1.8$. Lavender circles mark passively evolving galaxies at $z\\sim 2$ in the HUDF \\citep{maraston2006,maraston2007}. (d) Galaxies in CDFS and the HUDF where the blue and red circles denote the colors of the TP-AGB C and M ensembles in each of the four bands. The solid line traces the MIPS $24\\mu$m bandpass where $k$-corrections have been computed using the \\cite{dale2002} templates. \\label{fig:agbf} } \\end{figure*} Such analyses utilize template SEDs in order to estimate IR bolometric corrections \\citep[e.g.][]{chary2001,dale2002,rieke2009}. These, in turn, are used with empirical calibrations of on-going SFRs derived from nearby galaxies under the assumption that the mid-IR flux arises from the reprocessed light from young, luminous stars. However, these circumstances are not well understood, partly because of the different physical mechanisms and timescales probed by the IR compared to recombination lines or UV emission \\citep[e.g.][]{kennicutt1998,salim2009}. Here we recognize that when intermediate mass stars join the Thermally-Pulsating Asymptotic Giant Branch (TP-AGB), they do not uniformly have the colors of lower-mass TP-AGB M stars but, as C stars, are particularly luminous in the mid-IR, with most of their energy emitted between 20$\\mu m$ and 45$\\mu $m \\citep[e.g.][]{guandalini2006}. Knowing the importance of the TP-AGB at red and near-IR wavelengths \\citep{maraston2005,bruzual2009,conroy2009}, and, using near- to mid-IR colors of Galactic TP-AGB populations, we empirically calibrate the contributions of such stars to the integrated mid-IR luminosities of stellar populations. Using simple models, we show that the mid-IR luminosities of galaxies are specifically in the sensitivity of the mid-IR to the amount of stellar mass formed in the previous 1.5 Gyr, naturally complimenting the optical and near-UV. ", "conclusions": "\\label{sec:summary} The contribution of the TP-AGB to the $K$-band has been combined with with the mean IR colors of Galactic TP-AGB C and M stars in order to estimate the contributions of both young and old stellar populations to mid-IR observations of galaxies. Without tuning, we find that the resulting mid-IR luminosities of the TP-AGB can reproduce the MIPS $24\\mu$m fluxes for galaxies back to at least $z=2$ in a manner consistent with restframe optical colors. We have also tested the validity of the model on local scales in the galaxy M81 and find reasonable agreement. The origins of correlations between optical colors and mid-IR luminosities seen by others, such as \\cite{salim2009}, can now be understood. \\begin{figure*} \\centerline{ \\includegraphics[scale=0.36]{figure4.png}} \\caption{(a) M81 at $24\\mu$m. (b) The model $24\\mu$m image of M81, computed using extinction-corrected $u-g$ and $K_s$-band images and the model locus in Fig. \\ref{fig:M81}(c). (c) The flux enclosed within concentric apertures is shown for both $24\\mu$m and $8\\mu$m. \\label{fig:M81b} } \\end{figure*} With careful modeling of SEDs from the UV through the mid-IR, more detailed histories of star formation should be possible. Unfortunately, stellar spectral libraries and theoretical modeling are neither sufficient for verifying nor reducing the uncertainties our models \\citep{conroy2009}. This is largely due to the great difficulty in modeling post-main-sequence evolution, including the envelopes of TP-AGB stars, though the UV may provide further constraints \\citep{buzzoni2008}. We are optimistic that improved characterization of the mid-IR colors of the TP-AGB can be incorporated into SED fitting, though our calculations have uncertainties perhaps on the order of a factor of two due to uncertainties in the ensemble colors of the TP-AGB populations at different ages. With refinement we anticipate that incorporating the mid-IR into multiwavelength analysis of SEDs will provide the strongest constraints on the star formation histories of galaxies. There is little doubt that star formation and the growth of stellar mass was occurring more rapidly in the distant universe than today, but the nature of that growth has remained largely unknown. Earlier results \\citep[e.g.][]{lefloch2005} had implied that $\\sim 1/3$ of the stellar mass at the present epoch was formed after $z=1$ --- a result that appears to be at odds with the evolution in the stellar mass function to $z=1$ \\citep[e.g.][]{cirasuolo2007}. But the model presented here implies that the mid-IR provides the total mass in stars formed in windows stretching back 1.5 Gyr in cosmic time. As a result, such observations must be used with care when constraining the star formation rate density of the universe at $z < 2$, or when considering whether variations in the initial mass function are warranted by the data \\citep[e.g.][]{dave2008,wilkins2008}. The detection of galaxies in the mid-infrared over most of a Hubble time has helped change our view of galaxy assembly, and the determination of star formation rates associated with that assembly has remained a difficult task \\citep[see][]{chary2001,calzetti2007,salim2009,rieke2009}. Perhaps the most important implication of this {\\it Letter\\/} is that modeling the TP-AGB has allowed us to derive the relationship between mid-IR luminosities and star formation rates from ``first principles'' for the first time. With such models, it should now be possible to more accurately constrain the detailed history of star formation in the universe back to early times." }, "1003/1003.4519_arXiv.txt": { "abstract": "We present a study of the vertical magnetic field of the Milky Way towards the Galactic poles, determined from observations of Faraday rotation toward more than 1000 polarized extragalactic radio sources at Galactic latitudes $|b| \\ge 77^\\circ$, using the Westerbork Radio Synthesis Telescope and the Australia Telescope Compact Array. We find median rotation measures (RMs) of $ 0.0 \\pm 0.5$~rad~m$^{-2}$ and $+6.3\\pm0.7$~rad~m$^{-2}$ toward the north and south Galactic poles, respectively, demonstrating that there is no coherent vertical magnetic field in the Milky Way at the Sun's position. If this is a global property of the Milky Way's magnetism, then the lack of symmetry across the disk rules out pure dipole or quadrupole geometries for the Galactic magnetic field. The angular fluctuations in RM seen in our data show no preferred scale within the range $\\approx0\\fdg1$ to $\\approx25^\\circ$. The observed standard deviation in RM of $\\sim9$~rad~m$^{-2}$ then implies an upper limit of $\\sim1$~$\\mu$G on the strength of the random magnetic field in the warm ionized medium at high Galactic latitudes. ", "introduction": "\\label{section:introduction} Large scale coherent magnetic fields are observed in our Milky Way and in external galaxies \\citep{beck2008b,beck2008a}; these fields play crucial roles in many astrophysical processes in the interstellar medium (ISM) --- they help to exert pressure to balance ordinary matter against gravity and trigger star formation, they are also responsible for the confinement of cosmic rays, and they can regulate and trace a large scale galactic wind. Therefore, to better understand galaxy evolution, it is necessary to investigate the structure, origin and evolution of galactic magnetic fields. The primordial theory and the dynamo theory are the two possible explanations for the existence of a galactic scale magnetic field. These two theories make certain predictions on the symmetry of the large scale magnetic field with respect to the rotation axis and the mid-plane of the galaxy \\citep[e.g.,][]{beck1996}. The characterization of the overall magnetic field geometry of a galaxy allows one to distinguish between various primordial and dynamo models. Studies of the large scale magnetic field geometry in the Milky Way have been mainly focused on the disk field symmetry with respect to the rotation axis \\citep[e.g.,][]{han2006,brown2007} even though sight lines towards high Galactic latitude may be less turbulent and have less tangled fields than are seen in the Galactic plane and hence global field patterns can be identified more easily. The strength and the symmetry of the vertical, azimuthal and radial components of the Galactic magnetic field across the plane can provide us with unique insights on the mechanisms which maintain the Galactic magnetic field. For example, a reversal across the Galactic plane, in the azimuthal component of the large-scale magnetic field but not in its vertical component would indicate a field of dipolar structure, which could result from weak differential rotation (near solid body rotation) \\citep{ferriere2005}; or a substantial primordial field \\citep{zweibel1997}. On the other hand, if there is a reversal in the vertical but not the azimuthal component across the mid-plane, it would indicate a field of quadrupolar structure resulting from dynamo action due to differential rotation in the Galactic disk \\citep{zweibel1997}. The dynamo theory predicts a weak vertical field compared to the horizontal field on galactic scales. If the measured large-scale vertical field strength is substantially larger, it could imply a primordial component to the Galactic magnetic field \\citep{ruzmaikin1988}. The ratio of the vertical to horizontal field strength also regulates the confinement of cosmic rays to the Galactic disk, which can help us better understand the disk-halo interaction in the Milky Way. Some information on the large scale structure of the Galactic magnetic field can come from optical starlight polarization and radio synchrotron polarization \\citep[see][for a summary]{beck1996,beck2008b,beck2008a}. A series of papers: \\cite{berdyugin2000,berdyugin2001a,berdyugin2001b,berdyugin2004} found that towards the south Galactic pole, the polarization of stars traces a field orientation along $\\ell$ of 80$^\\circ$ which is parallel to the local spiral arm field, while the same field direction could not be traced towards the North Galactic pole. Most of our knowledge of the geometry of the large-scale Galactic magnetic field to date comes from the Faraday rotation measure of distant extragalactic radio sources (EGSs) and pulsars. Faraday rotation is a birefringence effect when linearly polarized light travels through a magnetized media. The plane of the polarization rotates through an angle $\\Delta$$\\psi$ (in radians) given by \\begin{equation} \\Delta\\psi= {\\rm RM} \\lambda^{2}, \\end{equation} where $\\lambda$ is the wavelength of the radiation measured in meters and RM is the rotation measure, defined as the integral of the line of sight magnetic field $B_{\\parallel}$ (in $\\mu$G) weighted by the thermal electron density $n_{e}(l)$ (in cm$^{-3}$) over a line-of-sight line element $d$$\\it{l}$ (in pc). \\begin{equation} {\\rm RM} =0.812 \\int ^{observer}_{source} {n_{e}(l){B_{\\parallel} (l)}} dl~~~\\rm{rad~m^{-2}}. \\label{eq:rmdef} \\end{equation} The sign of the RM gives the direction of the line of sight component of the average field: a negative RM represents a magnetic field whose line of sight component is directed away from us. With independent knowledge of the thermal electron density, one can determine the average line-of-sight magnetic field strength. Faraday rotation is complementary to other measurement techniques since RMs provide the direction of the magnetic field, while most other techniques (apart from Zeeman splitting) provide the field orientation, but not its direction. High latitude magnetic fields in the Milky Way have been investigated using RMs of EGSs by \\cite{morris1964}, \\cite{andreasyan1980}, \\cite{andreassian1988} and \\cite{han1994}. They all found a strong antisymmetric RM pattern between the northern and southern Galactic hemispheres. These authors attribute the pattern to a horizontal magnetic field that reverses direction above and below the Galactic plane. \\cite{han1994} and \\cite{han1999} estimated the local vertical magnetic field strength to be $\\sim$ 0.3 $\\mu$G, pointing from the south to the north Galactic pole assuming a priori a dipolar field geometry. A wavelet analysis on an all sky RM catalogue presented by \\cite{frick2001} found that the RM distribution at the largest scales is shifted to negative latitudes, indicative of a stronger magnetic field in the southern hemisphere and hence the possible existence of an antisymmetric\\footnotemark[1]\\footnotetext[1]{In this paper, the terms symmetric and antisymmetric are used to describe the vertical and horizontal components of the Galactic magnetic field instead of the total magnetic field vector itself.} halo magnetic field\\footnotemark[2]\\footnotetext[2]{Throughout the paper, we refer to halo fields as non-disk fields that are in regions with sufficient diffuse interstellar electrons to produce Faraday rotation}. Unfortunately, local distortions such as that of a Parker instability loop cannot be ruled out. A dipolar halo field is also suggested by \\cite{sun2008}, who demonstrated that an oppositely directed torus-like halo field above and below the Galactic plane and a symmetric disk field provide a reasonable fit to the latitude extension of the Canadian Galactic Plane Survey (CGPS) RM measurements, but it is unclear if this model can predict RMs that are consistent with observations for all ranges of $\\ell$ and $b$. \\cite{jansson2009} found that the halo field might not be anti-symmetric globally since a halo field that reverses across the plane only towards the inner Galaxy provides a better fit to the WMAP 22GHz all sky polarization map and an all-sky catalogue of $\\sim$ 1,400 RMs of EGSs. Recently, \\cite{taylor2009} used EGS RMs calculated from the NRAO VLA Sky Survey (NVSS) catalogue, a 1.4GHz continuum survey of the entire sky north of DEC=$-$40$^\\circ$, to derive a vertical field strength of 0.30 $\\pm$ 0.03 $\\mu$G for $b$$<$0$^\\circ$ pointing from the south pole towards the north pole and a field strength of 0.14 $\\pm$ 0.02 $\\mu$G for $b$$>$0$^\\circ$ pointing from the north towards the south pole. This new result is incompatible with the simple dipolar field model concluded in earlier studies. The current knowledge of the parity of the Galactic magnetic field is based on the sparse all-sky RM measurements and their potential unreliable values because they were derived using polarization angles at only a few and sometimes widely separated wavelengths. Moreover, most conclusions were drawn using low-latitude RMs which are likely to be contaminated by turbulence and tangled fields in the Galactic disk. In this paper, we present an accurate rotation measure survey of more than 1,000 polarized extragalactic sources previously identified in the NVSS catalogue towards both the north and the south Galactic poles at a sampling density of approximately 1 source/deg$^2$. Our goal is to measure the magnetic field structure at high latitude in great details and to study the symmetry of the vertical Galactic magnetic field across the disk plane. In \\S~\\ref{section:observations}, we describe data acquisition and reduction procedures. The RM extraction procedure is outlined in \\S~\\ref{section:rmcomputation}. The results are presented in \\S~\\ref{section:results}. We discuss different possible origins of the measured RM pattern in \\S~\\ref{section:discussion} and derive the corresponding magnetic field properties, including both the coherent (\\S~\\ref{section:largescalefield}) and the random magnetic field strength (\\S~\\ref{section:random}). ", "conclusions": "In this section, we explore possible origins of the observed RM pattern towards the Galactic caps: an overall RM consistent with zero towards the north and a positive median RM towards the south. We compare predictions from different models with the observed RM pattern. We first investigate local sources/ events that might give rise to the observed RMs: the local interstellar medium, the local bubble and a Parker instability loop\\footnotemark[11]\\footnotetext[11]{We have excluded the NPS two shell model \\citep{wolleben2005} from consideration as we have already discarded RMs of EGSs whose sight lines intercept the NPS (see \\S~\\ref{subsection:NPS})}. We then estimate the magnetic field strength towards the Galactic poles from the observed RMs by assuming that the observed Faraday rotation occurs in diffuse ionized gas in the Galaxy. Finally, we consider the likelihood that the observed RM pattern has been generated by global events in the Galaxy, such as a large scale Galactic wind, large scale Galactic dynamos or a relic field. \\subsection{Local origins} \\label{subsection:local} \\subsubsection{Local Interstellar medium} \\label{subsection:lism} The immediate medium surrounding the sun could potentially produce the observed RM pattern at high Galactic latitude. Even though the sun resides in a low density cavity called the Local Bubble (LB) (see more in \\S~\\ref{subsection:localbubble} ), there are warm partially ionized clouds surrounding the Sun \\citep{linsky2007}. These clouds are magnetized and have free electrons, and therefore are capable of rotating the polarization plane of incident radiation. \\cite{spangler2009} estimated the upper limit of $|$RM$|$ produced by the LISM to be 0.32$-$1.1 rad m$^{-2}$ by assuming that typical clouds have an electron density of 0.12 cm$^{-3}$, a volume filling factor of 5.5$\\%$ - 19$\\%$ \\citep{redfield2008a,redfield2008b} and a magnetic field strength of 4 $\\mu$G. Different studies have yield different estimates for the local magnetic field strength. \\cite{snowden1998} suggested that a magnetic field of up to 7$\\mu$G is required in the LB to counter balance the enormous thermal pressure ($p/k$$\\sim$ 15,000 cm$^{-3}$ K ) exerted by the enclosing hot X-ray gas, but the recent discovery of X-ray emission associated with charge exchange between solar wind ions and heliospheric plasma has greatly alleviated the need of non-thermal pressure support of the LB and has lowered the required magnetic field strength to $\\sim$ 2.8 $\\mu$G \\citep{welsh2009a}. Other works \\citep{opher2007,wood2007} find a local field strength of $\\sim$ 2 $\\mu$G. Since various estimations of the local magnetic field strength are lower than the 4 $\\mu$G adopted by \\cite{spangler2009}, the LISM $|$RM$|$ can be a factor of two smaller than the estimates of \\cite{spangler2009}. We note that the estimated LISM $|$RM$|$ contribution of $\\sim$ 1 rad m$^{-2}$ is an upper limit as $|$RM$|$ would be much smaller if the magnetic field reverses within the clouds or from cloud to cloud. As the typical RM measurement error in our experiment is a few rad m$^{-2}$, contribution of RM from the LISM is likely to be negligible. \\subsubsection{Effects of the local bubble wall} \\label{subsection:localbubble} In this section, we consider the possibility that the Faraday rotation originates in the wall of the Local Bubble. As mentioned in \\S~\\ref{subsection:lism}, the Sun is situated in a low-density cavity thought to be created by star formation and subsequent supernova explosions that occurred in the past 25-60 Myrs \\citep{frisch2007}. Such processes would sweep up materials and magnetic fields in the solar neighborhood into a dense shell and might produce measurable RM for sight lines through it. To test if the LB wall is responsible for producing the observed RM towards the Galactic caps, we have adopted the LB wall model constructed by \\cite{cl2002,cl2003} with slight modifications: we have modeled the LB wall as a slanted cylinder of constant radius 0.085 kpc that extends to 0.2 kpc both above and below the Galactic plane. Unlike the original model in \\cite{cl2002,cl2003}, our cylindrical wall has open ends since the recent expansion of NaI measurements made by \\cite{lallement2003} and \\cite{welsh2009b} find no continuous neutral LB boundary at high Galactic latitudes. We have computed the projection of the modeled LB boundary towards the Galactic poles and plotted it in Fig~\\ref{fig:rm_dis} as a green dotted line. None of the sight lines towards the north Galactic cap intercepts the local bubble wall, while 267 out of 341 sight lines intercept the modeled wall towards the south Galactic cap. Towards the south Galactic pole, the median RM of sight lines that penetrate the LB wall is $+$6.1 $\\pm$ 0.7 rad m$^{-2}$ whereas those do not penetrate the LB wall have a median of $+$ 7.8 $\\pm$ 1.8 rad m$^{-2}$. Since the inferred RM through the LB wall is consistent with zero ($-$2 $\\pm$ 2 rad m$^{-2}$), we conclude that the local bubble wall is unlikely to be the major contributor to the observed RM towards the caps if the adopted wall model is realistic. The 3$\\sigma$ upper limit of $|$RM$|$ $\\sim$ 7 rad m$^{-2}$ through the local bubble wall can be used to infer the magnetic field strength in the LB wall. Since \\cite{bhat1998} found that the scintillation measures of 20 nearby pulsars were well modeled by a scattering structure with local electron density enhancement of a factor of 10 which roughly coincides with the neutral LB wall, we assume that the electron density wall can be well traced by the neutral wall. If energetic events have swept up magnetic fields and electrons of initial density 0.025 cm$^{-3}$ \\citep{bhat1998} into a cylindrical shell of radius $R$ $\\sim$ 85 pc \\citep{cl2002,cl2003} centered at the Sun, the thickness of the shell in pc ($\\delta R$) is related to the electron density enhancement ($x$$=$10) in the shell and its radius $R$ by conserving the mass of electrons before and after the formation of the shell by $\\delta R = \\frac {R}{2x}$. The thickness of the cylindrical LB shell is estimated to be $\\sim$ 4 pc. If the RM produced by the wall is 7 rad m$^{-2}$, it implies a magnetic field strength of $\\sim$ 9 $\\mu$G . This prediction from a simple theoretical model of the LB is in rough agreement with the magnetic field strength of 8 $\\mu$G estimated by \\cite{andersson2006} \\citep[or earlier][]{leroy1999}, who have applied the \\cite{chandrasekharfermi1953} method on starlight polarization measurements towards stars at distances from 40$-$200 pc in the direction of $\\ell$=300$^\\circ$, $b$=0$^\\circ$. \\subsubsection{Small Scale outflows: a Parker instability loop} \\cite{parker1979} demonstrated that a system of horizontal magnetic field and cosmic rays in a vertical gravitational field can be unstable with respect to the bending of magnetic field lines in the vertical plane. For example, energetic stellar events can produce such waviness in a Galactic disk \\citep{kronberg1994}. It is possible that the entire surface of the Galactic disk is packed with Parker loops with height $\\sim$ 1 kpc and width of 0.1$-$1 kpc \\citep{parker1992}. If these magnetic loops thread the warm ionized gas in the Galaxy, it is plausible that they contribute towards the observed RM patterns at high Galactic latitude. \\cite{frick2001} performed a wavelet analysis on the all sky RM distribution and found that the RM structure at the largest scale has been shifted to a negative Galactic latitude of $-$15$^\\circ$, even after omitting sight lines that intercept Loop I. While one can interpret this shift as due to a stronger large-scale field in the southern Galactic hemisphere, possibly due to a separate halo dynamo with opposite parity to the disk field, it can also be due to a Parker instability loop. The authors suggest a scenario that a Parker instability loop with the Sun located near its top ($\\sim$ 50 pc above the galactic plane), a horizontal extent of $\\sim$ 400 pc and a magnetic field strength enhancement of 0.5 $\\mu$G can shift the symmetry axis of the RM structure to $b$=$-$15$^\\circ$. While the physical parameters of the inferred loop is consistent with that proposed by \\cite{parker1992}, it has difficulties explaining the observed RM pattern towards the Galactic caps on its own. Since the sun is located near the top of this loop, when looking towards a cone of radius 13$^\\circ$ around the Galactic poles, the extra $|$RM$|$ produced by the loop is small $\\sim$ 1-2 rad m$^{-2}$, comparable to individual RM measurement errors. Furthermore, the median RM towards the south Galactic cap would be zero if the loop is sufficiently symmetric as half of the south cap region should have the exact opposite RMs to the other half. This is inconsistent with what has been observed: the RM towards the entire south Galactic cap is positive (Fig~\\ref{fig:smooth_rm}). If the Sun is not located exactly at the top of this loop, then one would expect median RMs of different signs towards each cap because the magnetic field lines should be continuous. This is again inconsistent with the observations because the average RM towards the north cap is consistent with zero. We therefore conclude that the Parker instability loop proposed by \\cite{frick2001} alone cannot reproduce our high latitude RM measurements. We note that the argument above is based on the assumption that Parker instability loops exist in the warm ionized gas (typical WIM density $\\sim$ 0.1 cm$^{-3}$) in the Galaxy. If these instabilities are associated with the hot phase of the ISM where the typical magnetic field strength is $\\sim$ 0.1 $\\mu$G \\citep{beck1996} and the typical density is very low 10$^{-3}$ - 10$^{-4}$ cm$^{-3}$ \\citep{sembach2003}, the expected RM from such loops would be much smaller than the observed $|$RM$|$ of a few rad m$^{-2}$. In this case, such loops could not produce the RMs seen towards the caps. \\subsection{Large scale magnetic field in the halo of the Milky Way} \\label{section:largescalefield} In this section, we consider the possibility that the observed RM pattern originates on the Galactic scale. We first estimate the implied magnetic field strength and direction from the observed RMs assuming that the observed Faraday rotation occurs in diffuse ionized gas. In reality, it is mostly in the WIM that the Faraday rotation takes place because even with a higher filling factor, the hot halo electrons ($\\sim$ 10$^5$ $-$ 10$^6$ K) are of very low densities ($\\le$ 10$^{-3}$ - 10$^{-4}$ cm$^{-3}$). Therefore, it is reasonable to assume that the hot halo electrons have negligible contributions to the observed RMs. We then consider the possibilities that the observed magnetic field is the result of either a Galactic wind, a large scale Galactic dynamo or a primordial field. \\subsubsection{Vertical Magnetic Field} \\label{subsection:vertical} After the removal of extreme RMs and anomalous RM regions from the data set following the method described in \\S~\\ref{section:results}, one can infer the properties of the Galactic magnetic field from the remaining EGSs RMs assuming that the WIM is where most of the Faraday rotation takes place. The observed RM is the integral of the projection of the Galactic magnetic field along the line of sight weighted by thermal electron density. At the highest latitude, RMs measure mostly the vertical component of the Galactic magnetic field, as the projection of the horizontal component along the line of sight is very small. In a right handed coordinate system centered at the location of the Sun where the positive $z$ axis points from the south Galactic pole to the north Galactic pole and the positive $x$ axis points towards the Galactic center, the Galactic magnetic field is of the form \\begin{equation} \\vec{B} = B_H \\cos (l_0) \\hat{x} +B_H \\sin(l_0) \\hat{y}+B_z \\hat{z}, \\end{equation} where $B_z$ is the coherent component perpendicular to the Galactic disk (defined such that a positive $B_z$ implies a field pointing from the south to the north Galactic pole), and $B_H$ is the coherent component parallel to the Galactic disk directed along Galactic longitude $l_0$. We consider $B_z$, $l_0$ and $B_H$ to be constant within the 13$^\\circ$ radius cone that we are probing around the Galactic poles. If we use the definition of RM (Eq~\\ref{eq:rmdef}) and assume that the volume averaged thermal electron density $n_e$ in the Milky Way is an exponential disk of mid-plane density $n_{e,0}$ and scale height $H_0$ : \\begin{equation} \\label{eq:ne} n_e(z) = n_{e,0} e^{-|z|/H_0} , \\end{equation} then the RM towards Galactic coordinates ($\\ell$,$b$) can be expressed as \\begin{equation} \\label{eq:rmmodel} {\\rm RM} = 0.812 n_{e,0} H_{0} (a B_z - B_H \\cos(\\ell-\\ell_0)/\\tan|b|), \\end{equation} where $a$ =+1 for $b$$<$$-$77$^\\circ$ and $a$= $-1$ for $b$$>$+77$^\\circ$. If we take the average RM, $ \\langle {\\rm RM} \\rangle$, along many sight lines with full coverage in $\\ell$ towards the caps, the contribution from the horizontal field vanishes. The vertical magnetic field strength is related to the integrated thermal electron column density towards the Galactic poles, ${\\rm DM_\\perp}$=$n_{e,0} H_0$ by \\begin{equation} \\label{eq:bz} B_z = \\frac{ \\langle {\\rm RM} \\rangle } {0.812 a \\rm {DM_\\perp}}. \\end{equation} We note that while there remain controversies on the exact values of the scale height $H_0$ and the mid-plane electron density $n_{e,0}$, the total column density of thermal electrons is well constrained to be $\\sim$ 25 pc cm$^{-3}$ using pulsar DMs at high $z$ \\cite[see for example,][]{cl2003,gaensler2008}. From Eq~\\ref{eq:bz}, we found a vertical magnetic field towards the north Galactic pole to be consistent with zero (0.00 $\\mu$G $\\pm$ 0.02 $\\mu$G), with a 3$\\sigma$ upper limit on the vertical magnetic field strength of 0.07 $\\mu$G. On the other hand, the vertical magnetic field towards the south Galactic pole is found to be $+$0.31 $\\mu$G $\\pm$ 0.03 $\\mu$G. We note that the reduced $\\chi^2$s of a model with the derived vertical magnetic field strengths towards the Galactic poles exceed unity ($\\chi_r^2$ $\\sim$ 7.9 towards the north and $\\chi_r^2$ $\\sim$ 10.6 towards the south). This is because systematic RM scatter is introduced by intrinsic RMs of EGSs and small-scale Milky Way foreground electron density and magnetic field fluctuations. \\cite{wu2009} have found, using isothermal magnetohydrodynamic turbulence simulation, a relation between the distribution of normalized RM and the line of sight magnetic field strength. One can estimate the vertical magnetic field strength towards the poles using Equation (3) of \\cite{wu2009} and the median and standard deviation of the RM data set towards the north and the south caps reported in \\S~\\ref{subsection:rmdist}. The estimated vertical magnetic field strength towards the north and south Galactic pole are roughly 0 $\\mu$G and 0.46 $\\mu$G respectively, which agree in general with our results. However, as \\cite{wu2009} have pointed out, the relation is valid only for a Mach number of unity and thus it is unclear if it holds in the diffuse ionized gas at high Galactic latitude. There are two pulsars with measured RMs towards the Galactic caps: PSRs J0134-2937 and B1237+25 \\citep{taylor1993,han1999}. Their coordinates, DMs and RMs are listed in Table~\\ref{table:pulsar}. One can estimate the vertical magnetic field between the Sun and the pulsar using DM$_{\\rm pulsar}$ and RM$_{\\rm pulsar}$ \\begin{equation} \\label{eq:pulsarbz} B_z = \\frac { {\\rm RM_{pulsar}} \\sin{|b|}} {0.812 a {\\rm DM_{pulsar}}}, \\end{equation} where $b$ is the Galactic latitude of the pulsar. The vertical magnetic field derived from B1237+25 is $+$0.044 $\\pm$ 0.008 $\\mu$G, while that derived from J0134-2937 is +0.7 $\\pm$ 0.1 $\\mu$G. The vertical magnetic field strength derived using the pulsars roughly agree with results obtained using RMs of EGSs. Since PSR J0134-2937 has a much higher DM than PSR B1237+25, it probes through a longer path length in the WIM than the northern pulsar. \\cite{taylor2009} reported a vertical field of $-$0.14 $\\mu$G $\\pm$ 0.02 $\\mu$G towards the north Galactic pole and +0.30 $\\mu$G $\\pm$ 0.03 $\\mu$G towards the south Galactic pole. While our measured vertical field towards the south Galactic cap is consistent with \\cite{taylor2009} within errors, our estimation of the vertical field towards the north Galactic cap disagrees with \\cite{taylor2009}. This is likely due to the fact that \\cite{taylor2009} have averaged over a larger region around the north Galactic pole without discarding outliers and anomalous RM regions around the pole before performing the fit. Different individual RMs derived from the NVSS catalogue and our WSRT observations due to multi-RM component sources and different ionospheric conditions (\\S~\\ref{subsection:rmcompare}) might also contribute to this discrepancy. \\subsubsection{Horizontal Field} \\label{subsection:horizontal} In \\S~\\ref{subsection:vertical}, we used the fact that the median RM of EGSs is zero to reach the conclusion that there is no vertical field towards the north Galactic pole. However, it is still possible that there exists a horizontal magnetic field at high positive Galactic latitude. We can test this by fitting the north Galactic cap RMs to a model with only a horizontal component\\footnotemark[12]\\footnotetext[12]{This is obtained by setting $B_z$ = 0 $\\mu$G in Eq~\\ref{eq:rmmodel}.} \\begin{equation} \\label{eq:rmmodelhorizontal} {\\rm RM} = -0.812 {\\rm DM_\\perp} (B_H \\cos(\\ell-\\ell_0)/\\tan|b|), \\end{equation} minimizing the $\\chi^2$ between the observed RMs and the modeled RMs predicted by the above equation. The best fit parameters are $B_H$=$0.6_{-0.4}^{+0.8}$ $\\mu$G and $\\ell_0$= $153^{\\circ{+53^\\circ}}_{~-101^\\circ} $ with a reduced $\\chi^2$ of 7.7. Comparing this model to a model with no magnetic fields using the F-test, the significance of a horizontal field of strength 0.6 $\\mu$G is only at 2.4 $\\sigma$ level. This is not surprising as no obvious sinusoidal variation of RM as a function of $\\ell$ could be seen in the smoothed RM map in the top panel of Fig~\\ref{fig:smooth_rm}. Both the derived horizontal field strength and its direction differ from the best fit values towards positive mid-latitude (0.39 $\\mu$G at $\\ell_0$= 281$^\\circ$) obtained by \\cite{taylor2009}. As illustrated in Figure 7 of \\cite{taylor2009} , the best fit parameters obtained when fitting to high positive latitude RMs are different. The authors attributed this to a potentially more complicated halo magnetic field. We can also test if there is a horizontal magnetic field component in addition to the detected vertical component towards the south Galactic pole by performing a least square fit to a Galactic magnetic field model with both a vertical and a horizontal component (Eq~\\ref{eq:rmmodel}). The reduced $\\chi^2$ of the best fit to such a model is 10.6. The F-test suggests the existence of an additional horizontal field to the measured vertical field is significant at only 1.7 $\\sigma$ level. This is expected because no sinusoidal RM variations in $\\ell$ could be seen in the bottom panel of Fig~\\ref{fig:smooth_rm}. We conclude that there is little evidence for a horizontal field towards the Galactic poles. Instead of attempting to fit for the horizontal field, we can subtract its contribution from the measured RMs using the best fit values obtained by \\cite{taylor2009} to check if it changes our estimation of the vertical magnetic fields in \\S~\\ref{subsection:vertical}. This is justified only if the halo magnetic field at high latitude is the same as that at mid-latitude. At $|$$b$$|$=77$^\\circ$, the horizontal component contributes a maximum $|$RM$|$ of $\\sim$ 1.9 rad m$^{-2}$ towards the north Galactic pole and 4.3 rad m$^{-2}$ towards the south Galactic pole. We have subtracted the contribution of the \\cite{taylor2009} horizontal halo magnetic field from our measured RMs and found that it does not alter the vertical magnetic field estimates presented in \\S~\\ref{subsection:vertical} within the errors. \\subsection{Galactic Wind} X-shaped polarization pattern observed in halos of nearby edge-on galaxies imply large vertical magnetic fields increasing with height above the galactic disk \\citep{beck2008b}. A kinematic disk dynamo, which generates dominant toroidal magnetic fields, cannot alone explain the existence of these fields. As some of these galaxies exhibit evidence of cosmic ray driven winds, an alternative explanation of the large vertical field is wind advection that transports magnetic fields from the disk into the halo, distorting the expected field structure from dynamo actions \\citep[see for example, ][]{heesen2009}. Recent studies by \\cite{everett2008,everett2009} have successfully reproduced the Milky Way's diffuse soft X-ray emission and synchrotron emission using a 1D thermally and cosmic ray driven wind model for the Galaxy. In their cosmic ray driven wind model, the wind is launched within a Galactocentric radius of $\\sim$ 4.5 kpc, but flares to larger Galactocentric radii above and below the plane, and hence there is no wind launched at the location of the Sun. Similarly, \\cite{breitschwerdt1991,breitschwerdt1993} have argued that a wind launched from the solar neighborhood would lead to too much cosmic ray escape and hence inconsistent with the inferred residence time of cosmic rays in the Galaxy. \\cite{everett2008,everett2009} have adopted a flared-cylinder wind geometry where the wind stays well confined within a cylinder of constant cross-section up to a height z$_{break}$ above/below the Galactic plane. Beyond z$_{break}$, the cross-sectional area of the wind increases as a power law of z. At some z, this wind is directly above/below the location of the Sun. However, the best fit value of z$_{break}$ is found to be $\\sim$ 4 kpc \\citep{everett2009}. At this height above/below the Galactic plane, the density of the exponential WIM is too low to produce any observable Faraday rotation. Therefore, we conclude that the observed RM towards the Galactic caps is unlikely to be due to a large scale Galactic wind. \\subsection{Mean Field Dynamo Theory} \\label{subsection:dynamo} The existence of galactic-scale coherent magnetic fields in the Milky Way disk and in other normal spiral galaxies with significant differential rotation can be explained by the standard mean field $\\alpha$-$\\omega$ dynamo \\citep{beck1996}, although the theory and its application to galaxies has been questioned on theoretical grounds \\citep[see for example][for a recent summary]{cattaneo2009}. On a time scale of a few Gyrs, this process amplifies and orders the field by turbulence rising into the halo, transforming an azimuthal field into a poloidal one (the $\\alpha$-effect) and by differential rotation in the underlying disk, transforming the radial component of the poloidal field back into an azimuthal one (the $\\omega$-effect) \\citep{shukurov2004}. In this section, we examine if dynamo theory can predict the strength and geometry of a vertical field that is consistent with the values reported in \\S~\\ref{subsection:vertical}. The $\\alpha$-$\\omega$ dynamo predicts an azimuthal field that dominates over the vertical field by more than a factor of 10 because the $\\omega$-effect operates more efficiently than the $\\alpha$-effect \\citep{ferriere2005}. We can crudely estimate the expected local vertical magnetic field strength by evoking $\\nabla$ $\\cdot$ $\\vec{B}$ = 0 using a local horizontal field of $\\sim$ 2 $\\mu$G at a pitch angle of 15$^\\circ$ \\citep{beck1996}, and a Galactic disk of total thickness 2$h$ and diameter 2$R$ (where $h$ $\\sim$ 2 kpc and $R$ $\\sim$ 15 kpc). The estimated vertical magnetic field strength at the location of the Sun is $\\sim$ B$_{radial}$ ($h$/$R$) $\\sim$ 0.07 $\\mu$G. Another estimation of this ratio follows from dynamo theory: B$_z$ $\\sim$ B$_{radial}$ $\\sqrt{h/R}$ \\citep{ruzmaikin1988}. The expected vertical magnetic field strength is $\\sim$ 0.2 $\\mu$G, which is similar to the vertical magnetic field strength we obtained for the south Galactic cap. The dynamo mechanism is thus capable of producing the observed vertical magnetic field strength. The large scale magnetic field configuration of a galaxy can be classified by axial symmetry with respect to the rotation axis as well as by vertical symmetry with respect to the galactic mid-plane. For a strongly differentially rotating galaxy, a classical mean-field dynamo favors even field symmetry (a quadrupolar field), in which the toroidal component is symmetric across the mid-plane while the vertical component reverses direction \\citep{shukurov2004}. For a weakly differentially rotating galaxy, or a halo, odd field symmetry (a dipolar field) is preferred, in which the toroidal component reverses direction across mid-plane while the vertical component does not \\citep{ferriere2005}. However, numerical simulations carried out by \\cite{ferriere2000} found comparable growth rates of the odd and the even symmetry modes, suggesting that the present Galactic magnetic field might be of mixed parity rather than of pure even or odd parity. As mentioned in \\S~\\ref{section:introduction}, studies of the Milky Way's field parity have yielded diverging results, possibly because they focus on different Galactic latitude ranges and because it is difficult to distinguish local magneto-ionic effects from a genuine large scale field. The most convincing piece of work is the wavelet analysis of all sky RMs conducted by \\cite{frick2001}. They found that the local magnetic field has an even parity -- the horizontal field component does not reverse direction across the Galactic plane. If the Milky Way's large scale field is indeed quadrupolar in nature, then one expects the vertical magnetic field to reverse direction across the mid-plane. On the other hand, \\cite{han1994} and \\cite{han1997,han1999} concluded from an all sky smoothed RM map that the RM distribution towards the inner Galaxy is anti-symmetric across the Galactic plane. The authors attribute this to a large scale dipolar field. If this is the case, then one expects the vertical magnetic field to have the same direction above and below the Galactic plane. In \\S~\\ref{subsection:vertical}, we found that the vertical field is consistent with zero for $b$$>$ +77$^\\circ$ and that $B_z$ =$+$0.31 $\\pm$ 0.03 $\\mu$G for $b$$<$-77$^\\circ$. This is inconsistent with either a quadrupole or a dipole large scale field configuration. Even though a simple disk dynamo can produce a vertical field strength comparable to that being observed towards the south cap, it cannot explain the observed vertical field geometry. \\cite{sokoloff1990} demonstrated that the mean field dynamo can operate in the halo of a rotating galaxy since the mean helicity (the $\\alpha$-effect) of gaseous halos are non-zero due to galactic fountains/Parker instabilities. The dominant mode is an axisymmetric field with odd parity with respect to the Galactic disk; this is of the opposite parity from the mode excited by the dynamo action in the Galactic disk. The asymmetric RM pattern that we see towards the north/south Galactic hemisphere seems to support this idea. In particular, the shift of RM structure on the largest scales to $b$=$-$15$^\\circ$ found by \\cite{frick2001} can be explained by a dipolar field in the halo of the Galaxy such that the vertical magnetic field from the disk and the halo add up in the southern hemisphere. \\cite{sun2008} were able to obtain a reasonable fit to the latitude extension of CGPS measurements using an asymmetric halo field plus a symmetric disk field. More recently, \\cite{taylor2009} found that mid-latitude NVSS RMs can be well fitted with a $\\sim$ 0.4 $\\mu$G toroidal halo field that reverses direction across the mid-plane. If indeed an anti-symmetric halo field and a symmetric disk field co-exist in the Milky Way, then on the side of the Galactic disk (in this case $b$$>$0$^\\circ$), the vertical component of the halo field and that of the disk field would partly cancel out, as they are oppositely directed. These two vertical components would add up on the other side of the disk ($b$$<$0$^\\circ$). Depending on the relative vertical field strength between the halo/disk field and their extension above/below the disk, it is not impossible that they can cancel each other out exactly towards the northern hemisphere resulting in a net RM that is consistent with zero. However, this theory has been challenged by numerical work. \\cite{brandenburg1992} showed that the halo dynamo requires more than a Hubble time to reach the steady state configuration and thus the field that one observes at this moment might merely be a transient field of mixed parity. \\cite{moss2008} solved the mean field dynamo equations for a system with a disk and a halo dynamo. These authors were not able to produce any co-existing system of a dipole-like halo field and a quadrupole-like disk field; instead, one always dominates over the other, though \\cite{moss2008} acknowledged that using a larger turbulence diffusivity ratio between the halo and the disk might mitigate the problem. We conclude that the observed vertical magnetic field geometry towards the Galactic poles is not consistent with predictions from a pure disk dynamo. A separate halo dynamo of odd parity could potentially account for the observed vertical magnetic fields at high latitude but until now no numerical simulation has successfully produced a co-existing disk and halo fields of opposite parity in a steady state. \\subsection{Primordial origin} The competing theory to the dynamo origin of galactic magnetic fields is that of primordial origin. The primordial field theory suggests the following: if the IGM field is frozen into the gas, then the total magnetic field would be enhanced by a few orders of magnitude as gas clouds collapse to form a protogalaxy. This relic field would then be modified by the differential rotation of the galaxy, producing the present day galactic magnetic fields \\citep[see for example,][]{beck1996,howard1997}. The component of the seed field parallel to the galactic disk can be removed diffusively and through large scale flow, but the component of the magnetic field parallel to the rotation axis of the galaxy is trapped \\citep{ruzmaikin1988}. Since galactic rotation is symmetric with respect to the plane, one expects the azimuthal component of such a field to reverse its direction across the mid-plane while the vertical component preserves its direction, thus resulting in a dipolar type field \\citep{ruzmaikin1988,beck1996}. The strength of a vertical galactic magnetic field of primordial origin depends highly on the initial orientation of the intergalactic seed field with respect to the rotation axis. The collapsed gas that forms the proto-galaxy could potentially increase its strength by two orders of magnitude \\citep{kulsrud2008}. This small field is preserved until the present day since the total vertical magnetic flux is conserved. The concept of a primordial field may be overly simplistic because galactic disks probably build up over time through mergers and infall. The addition of new material can add new magnetic flux, but whatever flux is added is subject to the constraints described above. The observed vertical field geometry is not consistent with that from a pure dipole field of primordial origin: the vertical field from the south Galactic pole is directed toward us, while the vertical field toward the north Galactic pole is consistent with zero. Therefore, we cannot attribute the observed vertical magnetic field to a primordial field alone. \\label{section:conclusions} In this paper, we have presented an RM survey with the ATCA and the WSRT of polarized extragalactic radio sources towards the north and the south Galactic poles at $|b|$ $\\ge$ 77$^\\circ$. Using rotation measure synthesis, we have obtained 813 reliable RMs towards the Galactic poles. No preferred RM fluctuation scale was apparent from the flat RM structure functions towards the Galactic caps. After discarding outliers and anomalous RM regions in \\S~\\ref{section:results}, we obtain a median RM of 0.0 $\\pm$ 0.5 rad m$^{-2}$ towards the north Galactic pole; and a median RM of +6.3 $\\pm$ 0.7 rad m$^{-2}$ towards the south Galactic pole. In \\S~\\ref{subsection:local}, we have ruled out the possibility that local sources/events such as the LISM, the LB and a Parker's instability loop produce the observed RM pattern. In \\S~\\ref{subsection:vertical} and~\\S~\\ref{subsection:horizontal}, we have derived the halo magnetic field properties from the observed RMs assuming that they are produced by the diffuse interstellar free electrons. We found no evidence for vertical and horizontal magnetic field towards the north Galactic pole. On the other hand, a vertical field of strength +0.31 $\\pm$ 0.03 $\\mu$G was detected at $>$9 $\\sigma$ towards the south Galactic pole, but there is no evidence for an additional horizontal component. Although a dynamo or a primordial field can explain the derived vertical magnetic field strength towards the Galactic poles, a pure dipole/ quadrupole field cannot explain the geometry of the observed vertical field across the mid-plane. One possible explanation of the derived magnetic field properties is that proposed by \\cite{sokoloff1990}, in which a disk and a halo dynamo of different parities are simultaneously at work in the Galaxy. This could potentially lead to part cancellation of RM produced by the vertical magnetic field in the northern Galactic hemisphere, which is compatible with our RM measurements. However, until now, no numerical simulation has been able to produce a co-existing system of a dipole-like halo field and a quadrupole-like disk field. Numerical works that explore larger parameter space (especially turbulent diffusivity) is needed to test this hypothesis. Finally, we have estimated the random magnetic field strength in the halo of the Milky Way by constructing a plane-parallel cell model of the WIM and the standard deviation of RMs to derive a random magnetic field strength of $\\sim$ 1 $\\mu$G in the Galactic halo, which is smaller than the random field in the mid-plane of the Galaxy, but in equipartition with the lower turbulent energy density inferred for the halo. Exploration of cosmic magnetism is one of the key sciences of the next generation radio telescopes -- the Square Kilometre Array (SKA) and its prototypes such as the Australian Square Kilometre Array Pathfinder (ASKAP), which are capable of providing accurate RMs of EGSs densely sampled over the entire sky \\citep{johnston2007,johnston2008}. For example, one of the approved Survey Science project of ASKAP -- the Polarization Sky Survey of the Universe's Magnetism (POSSUM) aims to perform RM synthesis and obtain a grid of RMs over a large fraction of the sky. Similar projects will provide a more detailed picture on the magnetic field and turbulence properties at high Galactic latitude. Also, pulsar searches at high Galactic latitudes will allow one to probe the vertical magnetic field as a function of height above/ below the Galactic disk by simultaneously using pulsar DM and RM, which can further constrain the structure of the Galactic halo magnetic field. \\textbf{Acknowledgements} We thank Eve Meyer and Gemma Anderson for helping to carry out the ATCA observations; Observer's friend Ger de Bruyn and Gyula Jozsa to help with the preparation of the WSRT observations. Robert Braun for helping with the WSRT data calibration with AIPS, Douglas Finkbeiner for useful discussions, and Justin Kasper for detailed discussion on the ionospheric rotation measure correction. This work was supported in part by an Australian Research Council Federation Fellowship (FF0561298) awarded to B. M. G.. The Wisconsin H-Alpha Mapper is funded by the National Science Foundation. The Southern H-Alpha Sky Survey Atlas (SHASSA) is supported by the National Science Foundation. The Westerbork Synthesis Radio Telescope is operated by the ASTRON (Netherlands Institute for Radio Astronomy) with support from the Netherlands Foundation for Scientific Research (NWO). The Australia Telescope Compact Array is part of the Australian Telescope, which is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO. {\\it Facilities:} ATCA WSRT \\clearpage" }, "1003/1003.3049_arXiv.txt": { "abstract": "We present a series of three-dimensional hydrodynamical simulations of central AGN driven jets in a dynamic, cosmologically evolved galaxy cluster. Extending previous work, we study jet powers ranging from $L_{\\rm jet}=10^{44}$~erg/s to $L_{\\rm jet}=10^{46}$~erg/s and in duration from $30$~Myr to $200$~Myr. We find that large-scale motions of cluster gas disrupt the AGN jets, causing energy to be distributed throughout the centre of the cluster, rather than confined to a narrow angle around the jet axis. Disruption of the jet also leads to the appearance of multiple disconnected X-ray bubbles from a long-duration AGN with a constant luminosity. This implies that observations of multiple bubbles in a cluster are not necessarily an expression of the AGN duty cycle. We find that the ``sphere of influence'' of the AGN, the radial scale within which the cluster is strongly affected by the jet, scales as $R \\propto L_{jet}^{1/3}$. Increasing the duration of AGN activity does not increase the radius affected by the AGN significantly, but does change the magnitude of the AGN's effects. How an AGN delivers energy to a cluster will determine where that energy is deposited: a high luminosity is needed to heat material outside the core of the cluster, while a low-luminosity, long-duration AGN is more efficient at heating the inner few tens of kpc. ", "introduction": "The X-ray emitting gas in the centre of many galaxy clusters has a cooling time less than the Hubble time. Supported by large central concentrations in the surface brightness profile, early models of cluster evolution posited that cooling gas is funneled onto the central galaxy in the form of a cooling flow, ultimately resulting in a cooling catastrophe as more and more gas condenses out of the hot phase (Fabian 1994). However, high-resolution spectra of clusters show a lack of cool gas below a temperature floor of about one third of the virial temperature ($\\sim2$~keV) and a lack of star formation at the levels expected for the early estimates of cooling rates in excess of 100 solar masses per year in powerful cooling flows (e.g., Peterson et al. 2001; Fabian et al. 2001; Kaastra et al. 2001; Tamura et al. 2001). The most natural interpretation is that some source of heat is preventing the gas from cooling. The heat source must respond to the cooling rate of the cluster gas in such a way that heating and cooling are balanced, on average, in a self-regulating feedback cycle. Many different mechanisms for heating in clusters have been proposed, including AGN outflows, thermal conduction, dynamical friction, gravitational heating, cosmic rays and preheating, each either alone or in combination (see Conroy \\& Ostriker 2008; Brighenti \\& Mathews 2002 for reviews). However, recent studies have found evidence for AGN activity in the central galaxy of nearly all cooling-core clusters (Burns 1990; Mittal et al. 2009). The now widely accepted picture is that radio-loud AGN inflate bubbles of under-dense, relativistic plasma which displaces the local ICM. The plasma bubbles appear as dark cavities in X-ray observations (see, for example, B\\^irzan et al. 2004; Dunn \\& Fabian 2004; Rafferty et al. 2006). While the presence of cavities is ubiquitous, the longer term evolution of the plasma, and the direct effects on the gas are still uncertain and a topic of ongoing work (e.g. Reynolds et al. 2002; Ruszkowski et al. 2004b; Zanni et al. 2005; Gaibler et al. 2009). Although it is possible that AGN outflows are not solely or primarily responsible for cluster feedback, understanding how AGN interact with the cluster environment is still necessary to understand how clusters evolve. Interactions with the cluster environment can have an effect on other types of cluster heating. Recent work by Parrish et al. (2010) and Ruszkowski \\& Oh (2010) has found that turbulence in the cluster above a threshold value can suppress the heat-flux-driven buoyancy instability which would other wise inhibit thermal conduction. Turbulence can, therefore, allow efficient thermal conduction to occur, so it is important to understand the velocity structure of the cluster and how AGN alter that structure. \\subsection{AGN simulations} The details of how an AGN delivers energy to the cluster are still largely unknown. Simulations of AGN jets or hot, under-dense bubble in idealized, spherically symmetric clusters have found that energy is confined to a narrow angle around the jet axis, rather than being spread throughout the cluster (e.g. Churazov et al. 2001; Reynolds et al. 2001; Saxton et al. 2001; Br\\\"uggen et al. 2002; Br\\\"uggen \\& Kaiser 2002; Reynolds et al. 2002; Ruszkowski et al. 2004a; Dalla Vecchia et al. 2004; Omma et al. 2004; Omma \\& Binney 2004; O'Neil et al. 2009; O'Neil \\& Jones 2010). Turbulence produced by RT instabilities has been shown to help in distributing energy and in producing a self-regulated AGN jet (Scannapieco \\& Br\\\"uggen 2008, Br\\\"uggen \\& Scannapieco 2009). However, a simulation carried out by Heinz et al. (2006) found that in a realistic, cosmologically evolved cluster the motion of cluster gas effectively distributed the effects of the AGN over a wide angle. There is also observational evidence for the misalignment of radio jets and X-ray cavities in Abell 4059, presumably due to bulk motions in the ICM (Heinz et al. 2002). MHD simulations of clusters with AGN and turbulence from star formation (Falceta-Gon\\c{c}laves et al. 2010b) were able to distribute AGN energy isotropically and produce a filamentary structure similar to the observed in Perseus. An alternate means of distributing energy over a wide angle is to change the properties of the AGN rather than the cluster. Two-dimensional simulations in spherically symmetric clusters of slow, wide opening angle jets (Sternberg et al. 2007) or narrow, rapidly precessing jets with a wide precession angle (Sternberg \\& Soker, 2008) have been able to effectively distribute energy in the cluster core. Three-dimensional simulations by Falceta-Gon\\c{c}laves et al. (2010a) of slowly precessing jets were also able to produce an approximately isotropic energy distribution. These simulations were also able to create multiple bubbles if a large precession angle ($\\approx 60\\degr$) was used. In this paper, we directly investigate the role of cluster weather and jet power, as well as the injection history, on the evolution of radio lobes and X-ray cavities in the ICM. \\subsection{Jet intermittency} One critical question concerning the effect of jet activity is what the jet duty cycle is, i.e., what fraction of the time a black hole at the centre of a cool core cluster is actively driving a jet. The average jet power is often used to estimate the amount of AGN heating. However, the heating from a short, powerful period of AGN activity could be different from a long-lived, low-luminosity AGN that injects the same amount of total power. Observations of multiple radially segregated sets of cavities in several clusters have been used to argue that jet activity is intermittent. Examples include Perseus (B\\\"{o}hringer et al. 1993; Fabian et al. 2000, 2003, 2006), Hydra A (Wise et al. 2007), Virgo (Forman et al. 2007), and Abell 262 (Clarke et al. 2009). The duty-cycle inferred from these observations has been taken as evidence for self-regulation between cooling of the central cluster gas and AGN activity on timescales of tens of millions of years. Similar arguments have been made on the basis of sound and shock waves observed in deep Chandra observations of nearby clusters (Virgo, Perseus), which associate an AGN outburst with every ripple. As we will argue, the underlying assumption that surface brightness features can be associated one-to-one with activity in the central engine is likely overly optimistic. Sternberg \\& Soker (2009) have found that multiple sound waves can be excited by a single episode of bubble formation. Falceta-Gon\\c{c}laves et al. (2010a) were able to create multiple bubbles in simulations of a precessing AGN jet with a large precession angle. Cluster weather and the dynamic nature of the evolution of radio lobes can introduce features very similar to those observed in nearby clusters even in the case of continuously powered jets with no actual modulation of the jet power. In this paper, we present the results of hydrodynamical simulations of AGN jets in a realistic galaxy cluster with a variety of AGN properties, focusing on how the interaction of the AGN with a dynamic ICM affects the evolution of the morphology and energy distribution of the AGN. The paper is arranged as follows: Section 2 describes the setup and technical details of our simulations, Section 3 presents the results and discusses the implications for cluster heating and observations, Section 4 presents a analytic toy model to estimate the characteristic timescale for an individual bubble form and break off from the jet, and Section 5 presents a summary of our results and conclusions. ", "conclusions": "We have carried out a series of high-resolution hydrodynamic simulations of AGN jets in the centres of galaxy clusters. Simulations are carried out in a realistic, dynamic cluster and cover a range of AGN luminosities and durations. We find that the interaction of AGN jets with the motion of cluster gas is critical for determining the evolution of both the cluster and radio lobes. In particular, we find that: \\begin{itemize} \\item Multiple bubbles can be formed from a single period of AGN activity with a constant luminosity. As the AGN develops, bubbles can be broken off from the jet by the motion of gas within the cluster, leading to many generation of bubble formation. These bubbles can be pushed away from the jet axis by cluster weather. Therefore, observations of multiple X-ray or radio bubbles in a cluster does not necessarily give you any information about the duty cycle of the AGN or about the past alignment of the AGN jet. \\item A toy model balancing pressure in an expanding cocoon against ram pressure due to circular motion in the cluster core provides a reasonable estimate for the timescale of bubble breakoff. \\item Energy from the AGN is distributed over all angles by large-scale motions in the cluster. Jet material is distributed throughout the centre of the cluster and any information about the original orientation of the jet is lost on a time scale of about 100 Myr. This is not the case in a hydrostatic cluster, where the absence of large-scale flows allows the jets to propagate without being deflected. \\item A jet of a given luminosity will create a low-density, high-entropy cavity that expands to a fixed radius and then stops, limited by the interaction with large-scale flows in the cluster. The radius reached scales approximately as $R \\propto L_{jet}^{1/3}$. The exact radius reached is likely to depend on the detailed velocity and density structure within the cluster, but it does indicate that only high-luminosity AGN will be able to directly heat gas at 100 kpc or more from the cluster centre. \\item How an AGN delivers its energy determines where that energy ends up in the cluster. Lower luminosity AGN that are active for long periods are more efficient at heating the inner few 10's of kpc, while high luminosity AGN are necessary to deliver energy to large radii. In our simulations, an AGN with a luminosity of $10^{45}$~erg/s that was active for 90 Myr was as effective at removing mass and increasing entropy in the inner 30 kpc of the cluster as a $10^{46}$~erg/s AGN active for 30 Myr, despite emitting only 30\\% as much total power. However, the effects of the $10^{45}$~erg/s AGN were limited to half the radius of the $10^{46}$~erg/s AGN. \\item In a hydrostatic cluster, AGN evolution is quite different. A long-duration AGN inflates two large bubbles rather than many smaller ones. Jet material remains concentrated near the jet axis and the radius reached by the jet material continues to increase with time beyond the value we find in the dynamic case. \\end{itemize} The relationship between jet power and the radius of the jets ``sphere of influence'' has consequences for the impact jets have on clusters. The strong effect that the motion of cluster gas has on the AGN development means that exactly how and where an AGN deposits energy will be strongly affected by the inflation history of the jet and the dynamical state of the cluster. Low-power AGN are more efficient at heating the central cluster, but the heat is confined, implying that an additional heat source may be needed farther from the centre of the cluster. One possible heat source is conduction of heat from warm gas in the outer cluster to cooling gas toward the centre. Recent work by Parrish et al. (2010) and Ruszkowski \\& Oh (2010) has shown that turbulence can suppress the heat-flux-driven buoyancy instability and allow efficient thermal conduction to occur. AGN activity can act as a source of turbulence, potentially providing a switch that allows conduction to occur when an AGN is active. Heinz et al. (2010), using some of the same simulations presented here, found an increase in the turbulent velocity dispersion of cluster gas due to AGN activity of up to several $100$~km~s$^{-1}$. Central heating in cooling flows (where the entropy has to be injected in a relatively small volume) is more likely to result from continuously operating lower power jet than episodic powerful outbursts. This is, in a sense, numerical confirmation of the effervescent model, where continuous low-level activity generated multiple generations of bubbles (Ruszkowski \\& Begelman 2002). However, in our case, the effect of cluster weather is primarily responsible for breaking off bubbles, not necessarily the bubble's buoyant escape. The initial structure of a cluster will play a large role in determining the specific morphology of an AGN outflow. Although we only use one realistic cluster setup for our simulations, it is clear that the assumption of a simple one-to-one correspondence between observations of X-ray cavities or waves and periods of activity from an intermittent AGN is not valid. Multiple bubbles could be formed by intermittent AGN activity, but it is possible to produce multiple generations of X-ray bubbles from a single, continuously active AGN. The simulations presented here do not include magnetic fields. The presence of magnetic fields could stabilize the bubbles created by the AGN, allowing them to rise farther away from the cluster centre before becoming disrupted (Ruszkowski et al. 2007). Observations of an increase in X-ray cavity size with distance from the cluster centre have been interpreted as favoring a current-dominated MHD jet model (Diehl et al. 2008). However, Br\\\"uggen et al. (2009) argued that pure hydro simulations could produce similar observations. Even if the bubble evolution is not dominated by magnetic fields, the presence of weak magnetic fields will affect the mixing of jet and cluster material and can change where energy is deposited. Future simulations including realistic magnetic fields will assess the effect of magnetic fields on cluster evolution." }, "1003/1003.3878_arXiv.txt": { "abstract": "We introduce a powerful semi-numeric modeling tool, 21cmFAST, designed to efficiently simulate the cosmological 21-cm signal. Our code generates 3D realizations of evolved density, ionization, peculiar velocity, and spin temperature fields, which it then combines to compute the 21-cm brightness temperature. Although the physical processes are treated with approximate methods, we compare our results to a state-of-the-art large-scale hydrodynamic simulation, and find good agreement on scales pertinent to the upcoming observations ($\\gsim$ 1 Mpc). The power spectra from 21cmFAST agree with those generated from the numerical simulation to within 10s of percent, down to the Nyquist frequency. We show results from a 1 Gpc simulation which tracks the cosmic 21-cm signal down from $z=250$, highlighting the various interesting epochs. Depending on the desired resolution, 21cmFAST can compute a redshift realization on a single processor in just a few minutes. Our code is fast, efficient, customizable and publicly available, making it a useful tool for 21-cm parameter studies. ", "introduction": "\\label{sec:intro} Through challenging observational efforts, the high-redshift frontier has been incrementally pushed back in recent years. Glimpses of the $z\\sim$ 6--8 Universe were provided by quasars (e.g. \\citealt{Fan06}), candidate Lyman break galaxies (e.g. \\citealt{Bouwens08, McLure09, Bouwens09, Ouchi09}), Lyman alpha emitters (LAEs; e.g. \\citealt{Shimasaku06, Kashikawa06}), and GRBs (e.g. \\citealt{Cusumano06, Greiner09, Salvaterra09}). Unfortunately, these precious observations currently provide only a limited set of relatively bright objects. Luckily, we will soon be inundated with observations probing this and even earlier epochs. These observations should include infrared spectra from the {\\it James Webb Space Telescope} (JWST), the Thirty Meter Telescope (TMT), the Giant Magellan Telescope (GMT), the European Extremely Large Telescope (E-ELT), wide-field LAE surveys from the Subaru HyperSupremeCam, as well as the E-mode CMB polarization power spectrum measured by the {\\it Planck} satellite. Some of the most important information will come in the form of the redshifted 21-cm line of neutral hydrogen. Several interferometers will attempt to observe the cosmological 21-cm signal, including the Mileura Wide Field Array (MWA; \\citealt{Bowman05})\\footnote{http://web.haystack.mit.edu/arrays/MWA/}, the Low Frequency Array (LOFAR)\\footnote{http://www.lofar.org}, the Giant Metrewave Radio Telescope (GMRT; \\citealt{Pen08}), the Precision Array to Probe the Epoch of Reionization (PAPER; \\citealt{Parsons09}), and eventually the Square Kilometer Array (SKA)\\footnote{http://www.skatelescope.org/}. The first generation of these instruments, most notably LOFAR and MWA, are not only scheduled to become operational within a year, but should also yield insight into the 3D distribution of HI, provided the systematics can be overcome (see the recent reviews of \\citealt{FOB06, MW09}). However, interpreting this data will be quite challenging and no-doubt controversial, as foreshadowed by the confusion surrounding the scant, currently-available observations. There are two main challenges to overcome: (1) an extremely large parameter space, due to our poor understanding of the high-redshift Universe; (2) an enormous dynamic range (i.e. range of relevant scales). Theoretically, the dawn of the first astrophysical objects and reionization could be modeled from first principles using numerical simulations, which include the complex interplay of many physical processes. In practice however, simulating these epochs requires enormous simulation boxes. Gigaparsec scales are necessary to statistically model ionized regions and absorption systems or create accurate mock spectra from the very rare high-redshift quasars. However, the simulations also require high enough resolution to resolve the underlying sources and sinks of ionizing photons and the complex small-scale feedback mechanisms which regulate them. Thus one is forced to make compromises: deciding which physical processes can be ignored, and how the others can be parameterized and efficiently folded into large-scale models. Furthermore, large-scale simulations are computationally costly (even when they sacrifice completeness for speed by ignoring hydrodynamic processes) and thus are inefficient in large parameter studies. On the other hand, analytic models, while more approximate, are fast and can provide physical insight into the import of various processes. However, analytical models are hard-pressed to go beyond the linear regime, and beyond making fairly simple predictions such as the mean 21-cm signal \\citep{Furlanetto06}, the probability density function (PDF; \\citealt{FZH04_21cmtop}) and power spectrum \\citep{PF07, Barkana09}. The 21-cm tomographic signal should be rich in information, accommodating many additional, higher-order statistical probes, such as the bi-spectrum (Pritchard et al., in preparation), the difference PDF \\citep{BL08}, etc. In this paper, we follow a path of compromise, attempting to preserve the most useful elements of both analytic and numeric approaches. We introduce a self-consistent, semi-numerical\\footnote{By ``semi-numerical'' we mean using more approximate physics than numerical simulations, but capable of independently generating 3D realizations.} simulation, specifically optimized to predict the high-redshift 21-cm signal. Through a combination of the excursion-set formalism and perturbation theory, our code can generate full 3D realizations of the density, ionization, velocity, spin temperature, and ultimately 21-cm brightness temperature fields. Although the physical processes are treated with approximate methods, our results agree well with a state-of-the-art hydrodynamic simulation of reionization. However, unlike numerical simulations, realizations are computationally cheap and can be generated in a matter of minutes on a single processor, with modest memory requirements. Most importantly, our code is publicly available at http://www.astro.princeton.edu/$\\sim$mesinger/Sim.html. We name our simulation {\\it 21cmFAST}. Semi-Numerical approaches have already proved invaluable in reionization studies \\citep{Zahn05, MF07, GW08, Alvarez09, CHR09, Thomas09}. Indeed, 21cmFAST is a more specialized version of our previous code, DexM (\\citealt{MF07}; hereafter MF07). The difference between the two is that 21cmFAST bypasses the halo finding algorithm, resulting in a faster code with a larger dynamic range and more modest memory requirements. In this work, we also introduce some new additions to our code, mainly to compute the spin temperature. In \\S \\ref{sec:post_heat}, we compare predictions from 21cmFAST with those from hydrodynamic simulations of the various physical components comprising the 21-cm signal in the post heating regime. Density, ionization, peculiar velocity gradient, and full 21-cm brightness temperature fields are explored in \\S \\ref{sec:den}, \\S \\ref{sec:ion}, \\S \\ref{sec:dvdr}, \\S \\ref{sec:21cm_cmp}, respectively. In \\S \\ref{sec:heating}, we introduce our method for computing the spin temperature fields, with results from the complete calculation (including the spin temperature) presented in \\S \\ref{sec:results}. Finally in \\S \\ref{sec:conc}, we summarize our findings. Unless stated otherwise, we quote all quantities in comoving units. We adopt the background cosmological parameters ($\\Omega_\\Lambda$, $\\Omega_{\\rm M}$, $\\Omega_b$, $n$, $\\sigma_8$, $H_0$) = (0.72, 0.28, 0.046, 0.96, 0.82, 70 km s$^{-1}$ Mpc$^{-1}$), matching the five--year results of the {\\it WMAP} satellite \\citep{Komatsu09}. ", "conclusions": "\\label{sec:conc} We introduce a powerful new semi-numeric modeling tool, {\\it 21cmFAST}, designed to efficiently simulate the cosmological 21-cm signal. Our approach uses perturbation theory, excursion set formalism, and analytic prescriptions to generate evolved density, ionization, peculiar velocity, and spin temperature fields, which it then combines to compute the 21-cm brightness temperature. This code is based on the semi-numerical simulation, DexM, (MF07). However, here we bypasses the halo finder and operate directly on the evolved density field, thereby increasing the speed and decreasing memory requirements. In the post-heating regime, 21cmFAST can generate a realization in a few minutes on a single processor, compared to many days on $>$ 1000-node supercomputing cluster required to generate the same resolution boxes using state-of-the-art numerical simulations. 21cmFAST realizations in the pre-heating regime require $\\sim$ one day of computation time. Conversely, RT simulations of the pre-heating regime which resolve most sources currently do not exist, as they are too computationally expensive. Our code is publicly available at http://www.astro.princeton.edu/$\\sim$mesinger/Sim.html We compare maps, PDFs and power spectra from 21cmFAST, with corresponding ones from the hydrodynamic numerical simulations of \\citet{TCL08}, generated from the same initial conditions. We find good agreement with the numerical simulation on scales pertinent to the upcoming observations ($\\gsim$ 1 Mpc). The power spectra from 21cmFAST agree with those generated from the numerical simulation to within 10s of percent down to the Nyquist frequency. We find evidence that non-linear peculiar velocity effects enhance the 21-cm power spectrum, beyond the expected geometric, linear value. This enhancement quickly diminishes during the onset of reionization, remaining only on small scales at $\\avenf \\lsim 0.7$. Interestingly, we also find that the large-scale power is {\\it decreased} as a result of peculiar velocities in the advanced stages of reionization. The reason for this is due to the ``inside-out'' nature of reionization on large-scales: the remaining HI regions are preferentially underdense, in which peculiar velocities decrease the 21-cm optical depth and brightness temperature. Our code can also simulate the pre-reionization regime, including the astrophysical processes of X-ray heating and the WF effect. We show results from a 1 Gpc simulation which tracks the cosmic 21-cm signal down from $z=250$, highlighting the various interesting epochs. There are several large 21-cm interferometers scheduled to become operational soon. Interpreting their upcoming data will be difficult since we know very little about the astrophysical processes at high redshifts. Furthermore, there is an enormous range of scales involved, making numerical simulations too slow for efficient parameter exploration. 21cmFAST is not. \\vskip+0.5in We thank Hy Trac for permission to use the initial conditions and the output from his radiative transfer hydrodynamic simulation. We thank Dave Spiegel for his considerable Matlab experience, without which making 21-cm movies would have been considerably more difficult. Support for this work was provided by NASA through Hubble Fellowship grant HST-HF-51245.01-A to AM, awarded by the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., for NASA, under contract NAS 5-26555. SRF was partially supported by the David and Lucile Packard Foundation and by NASA through the LUNAR program. The LUNAR consortium (http://lunar.colorado.edu), headquartered at the University of Colorado, is funded by the NASA Lunar Science Institute (via Cooperative Agreement NNA09DB30A) to investigate concepts for astrophysical observatories on the Moon. Support was also partially provided by NASA grants NNG06GI09G and NNX08AH31G to RC." }, "1003/1003.4669_arXiv.txt": { "abstract": "Tidal interactions between neighboring objects span across the whole admissible range of lengths in nature: from, say, atoms to clusters of galaxies i.e. from micro to macrocosms. According to current cosmological theories, galaxies are embedded within massive non-baryonic dark matter (DM) halos, which affects their formation and evolution. It is therefore highly rewarding to understand the role of tidal interaction between the dark and luminous matter in galaxies. The current investigation is devoted to Early-Type Galaxies (ETGs), looking in particular at the possibility of establishing whether the tidal interaction of the DM halo with the luminous baryonic component may be at the origin of the so-called ``tilt'' of the Fundamental Plane (FP). The extension of the tensor virial theorem to two-component matter distributions implies the calculation of the self potential energy due to a selected subsystem, and the tidal potential energy induced by the other one. The additional assumption of homeoidally striated density profiles allows analytical expressions of the results for some cases of astrophysical interest. The current investigation raises from the fact that the profile of the (self + tidal) potential energy of the inner component shows maxima and minima, suggesting the possible existence of preferential scales for the virialized structure, \\ie\\ a viable explanation of the so called \"tilt\" of the FP. It is found that configurations related to the maxima do not suffice, by themselves, to interpret the FP tilt, and some other relation has to be looked for. ", "introduction": " ", "conclusions": "" }, "1003/1003.1750_arXiv.txt": { "abstract": "% ", "introduction": " ", "conclusions": "" }, "1003/1003.0340_arXiv.txt": { "abstract": "Several authors have shown that precise measurements of transit time variations of exoplanets can be sensitive to other planetary bodies, such as exo-moons. In addition, the transit timing variations of the exoplanets closest to their host stars can provide tests of tidal dissipation theory. These studies, however, have not considered the effect of the host star. There is a large body of observational evidence that eclipse times of binary stars can vary dramatically due to variations in the quadrupole moment of the stars driven by stellar activity. In this paper we investigate and estimate the likely impact such variations have on the transit times of exoplanets. We find in several cases that such variations should be detectable. In particular, the estimated period changes for WASP-18b are of the same order as those expected for tidal dissipation, even for relatively low values of the tidal dissipation parameter. The transit time variations caused by the Applegate mechanism are also of the correct magnitude and occur on timescales such that they may be confused with variations caused by light-time travel effects due to the presence of a Jupiter-like second planet. Finally, we suggest that transiting exoplanet systems may provide a clean route (compared to binaries) to constraining the type of dynamo operating in the host star. ", "introduction": "\\label{sec:intro} Since the discovery of the first exoplanet around a solar-like star by \\cite{mayor95}, the exoplanet field has bloomed with over 400 currently known. While the majority of exoplanets have been discovered through radial velocity measurements of the Doppler wobble effect, it is the systems that exhibit transits that are most highly prized. These planets are crucial for determining exoplanet bulk densities (via radii and mass measurements) as well as their atmospheric properties (via infrared measurements of their day/night variations and transmission spectroscopy). Transiting exoplanets, however, also offer the opportunity to detect other planets within the system since an additional planet may alter the period of the observed transits. This can occur in two ways. In the first case the gravitational influence of the perturbing body can can alter the orbital period of the transiting exoplanet directly. This effect is particularly strong for planets in mean motion resonances and can even allow Earth-massed objects to be detected, while `exo-moons' orbiting the transiting planet itself also induce a similar effect (e.g. \\citealt{simon07}). In the second case, a perturbing mass in a wider orbit can cause the transiting planet / star system to wobble around the barycentre, again altering the observed transit times but by changing the light travel-time. Although searching for transit timing variations (hereafter, TTVs) can potentially uncover the existence of Earth-mass objects (see \\citealt{gibson10}; \\citealt{rabus09}; \\citealt{bean09}; \\citealt{gibson09}; \\citealt{millerricci08}, and references therein for recent observational studies), there are other effects that can lead to TTVs. These include the precession of orbits due to general relativistic effects (\\citealt{pal08}), tidal dissipation, torques due to the spin-induced quadrupole moment of the host star (\\citealt{miralda02}), perturbations of transit times due to star spots, as well as reorientation of the planetary orbit with respect to the Earth as a result of proper motion (\\citealt{rafikov09}). However, there is a wealth of observations of many different eclipsing binary stars (e.g. \\citealt{hall80}; \\citealt{glownia86}; \\citealt{bond88}; \\citealt{warner88}; \\citealt*{baptista92}; \\citealt{echevarria93}; \\citealt{wolf93}; \\citealt*{baptista00}; \\citealt{baptista02}; \\citealt{baptista03}; \\citealt{borges08}) that have shown quasi-periodic variations in eclipse times over timescales of years to decades that are comparable to, or larger than, the effects being searched for amongst transiting exoplanets. The favoured explanation for these observed variations in the orbital periods of eclipsing binary stars is known as the Applegate effect (\\citealt{applegate92}). This mechanism invokes magnetic activity cycles in the low-mass components of such binaries to redistribute angular momentum within the interior of the star, thereby changing the stellar quadrupole moment which leads to changes in the orbital period of the components. Later, \\cite*{lanza98} proposed that the Applegate mechanism could also be driven by effectively converting rotational kinetic energy and magnetic energy back and forth. Regardless of the details of the exact physical mechanism at work, the Applegate effect should also operate in most exoplanet systems since the host stars are (by selection) low-mass stars with a convective outer layer which should exhibit some form of dynamo activity. It will therefore be important to know the magnitude of the Applegate effect for exoplanet systems when interpreting any TTVs. In this paper we briefly review the Applegate mechanism in the next section, before applying the analysis of \\cite{applegate92} to estimate the effects on known transiting exoplanet systems. Finally, we look at the implications that the Applegate effect has for TTV work in detecting additional planets as well as for the measurement of the level of tidal dissipation in very-hot Jupiter's. ", "conclusions": "We have estimated the likely transit timing variations induced by changes in the quadrupole moment of the host star in transiting exoplanet systems driven by the Applegate effect. Depending on the length of the activity cycle, TTVs of several minutes are plausible for a number of the currently known transiting exoplanets. While the timescales and sizes of Applegate driven TTVs are of the wrong magnitude to be confused with TTVs driven by additional planets in mean motion resonances, there appears to be much scope for confusion with TTVs caused by light travel time effects caused by massive, Jupiter-like planets on wide orbits. The magnitude of TTVs driven by the Applegate mechanism also grow as the star-planet separation decreases (assuming all other factors are equal). Indeed, for the shortest period transiting exoplanets, such as WASP-18b, the Applegate mechanism could potentially be mistaken for orbital period changes due to tidal dissipation. Indeed, the orbital decay due to tidal dissipation could even be temporarily reversed since the orbital period variations due to the Applegate effect can take on either sign. We therefore urge caution when interpreting TTVs, especially those that appear to be occurring on timescales of years to decades. In all cases, the clear signature that the Applegate effect is at work is that the TTVs are quasi-periodic. Only once the strict periodicity of any TTV has been ascertained should investigators be confident in their final interpretation. Finally, an alternative mechanism for driving quasi-periodic stellar quadrupole variations (and hence orbital period variations) was put forward by \\cite{lanza98}. This rests on the principle that changes in the azimuthal magnetic field intensity can change the stellar quadrupole moment by altering the effective centrifugal acceleration, the stellar dynamo effectively interchanges magnetic energy and rotational kinetic energy. The main feature of note in this prescription is that relative changes in the angular velocity required to drive orbital period changes are a factor of 2 smaller than that required by \\cite{applegate92}. While it is difficult to assess what TTVs may be expected from the \\cite{lanza98} work, it is quite possible that non-periodic TTVs with magnitudes exceeding those outlined in this paper could be observed. This latter point should also be taken in light of the fact that we have been very conservative (in comparison to \\citealt{applegate92}) in the energy budget we have assumed to be available for driving the stellar quadrupole variations. Indeed, as outlined earlier, quasi-periodic variations larger than those that the \\cite{applegate92} prescription could sustain may already have been observed in a few binary systems. Studies of the long-term trends in the transit times of short-period exoplanets could provide crucial evidence for settling may of the points above, presenting valuable insights into the working of stellar dynamos across a broad range of fundamental stellar properties." }, "1003/1003.5753_arXiv.txt": { "abstract": "{The Cosmic Ray Energetics And Mass (CREAM) balloon experiment had two successful flights in 2004/05 and 2005/06. It was designed to perform energy measurements from a few GeV up to 1000 TeV, taking advantage of different detection techniques. The first instrument, CREAM-1, combined a transition radiation detector with a calorimeter to provide independent energy measurements of cosmic-ray nuclei. Each detector was calibrated with particle beams in a limited range of energies. In order to assess the absolute energy scale of the instrument and to investigate the systematic effects of each technique, a cross-calibration was performed by comparing the two independent energy estimates on selected samples of oxygen and carbon nuclei.} \\begin{document} ", "introduction": "CREAM is a balloon-borne experiment designed to perform direct measurements of the energy spectra and elemental composition of cosmic rays (CR) up to the PeV scale. Two instruments, launched from McMurdo in 2004 and 2005, flew over Antarctica for 42 and 28 days, respectively. Both instruments achieved single-element discrimination by means of multiple measurements of the particle charge provided by a pixelated silicon charge detector (SCD), a segmented timing-based particle-charge detector (TCD) and a Cherenkov detector (CD). The particle energy was measured by a thin ionization calorimeter (CAL) preceded by a graphite target. During the first flight, the payload was equipped with a Transition Radiation Detector (TRD), thus allowing redundant energy measurements. A detailed description of the instrument can be found elsewhere \\cite{ref1}. In this paper, we present an analysis, based on the first flight data, that shows how it is possible to cross-calibrate the TRD and the calorimeter to assess the absolute scale of energy measurements in CREAM. ", "conclusions": "A preliminary analysis of the data from the first flight of CREAM confirmed the possibility to cross-calibrate the energy measurements of TRD and calorimeter." }, "1003/1003.5109_arXiv.txt": { "abstract": "{Recent work on several $\\beta$ Cephei stars has succeeded in constraining both their interior rotation profile and their convective core overshoot. In particular, a recent study focusing on $\\theta$ Ophiuchi has shown that a convective core overshoot parameter of $\\alpha_{ov}$ = 0.44 is required to model the observed pulsation frequencies, significantly higher than for other stars of this type.} {We investigate the effects of rotation and overshoot in early type main sequence pulsators, such as $\\beta$ Cephei stars, and attempt to use the low order pulsation frequencies to constrain these parameters. This will be applied to a few test models and the $\\beta$ Cephei star $\\theta$ Ophiuchi. } {We use the 2D stellar evolution code ROTORC and the 2D linear adiabatic pulsation code NRO to calculate pulsation frequencies for 9.5 \\msun\\ models evolved to an age of 15.6 Myr. We calculate low order p-modes ($\\ell \\leq 2$) for models with a range of rotation rates and convective core overshoot parameters. These low order modes are the same range of modes observed in $\\theta$ Ophiuchi. } {Using these models, we find that the convective core overshoot has a larger effect on the pulsation frequencies than the rotation, except in the most rapidly rotating models considered. When the differences in radii are accounted for by scaling the frequencies by $\\sqrt(GM/R(40^{\\circ})^3)$, the effects of rotation diminish, but are not entirely accounted for. Thus, this scaling emphasizes the differences produced by changing the convective core overshoot. We find that increasing the convective core overshoot decreases the large separation, while producing a slight increase in the small separations. We created a model frequency grid which spanned several rotation rates and convective core overshoot values. We used this grid to define a modified $\\chi^2$ statistic in order to determine the best fitting parameters from a set of observed frequencies. Using this statistic, we are able to recover the rotation velocity and convective core overshoot for a few test models. We have also performed a ``hare and hound\" exercise to see how well 1D models can recover these parameters. Finally, we discuss the case of the $\\beta$ Cephei star $\\theta$ Oph. Using the observed frequencies and a fixed mass and metallicity, we find a lower overshoot than previously determined, with $\\alpha_{ov}$ = 0.28 $\\pm$ 0.05. Our determination of the rotation rate agrees well with both previous work and observations, around 30 \\kms.} {} ", "introduction": "Recently, great progress in the asteroseismology of $\\beta$ Cephei stars has been made thanks to extensive observational campaigns, which have allowed constraints to be placed on the interior properties of some of these stars. For example, \\citet{aerts03} and \\citet{aerts04v836} have compiled and analyzed 21 years of photometry for V836 Cen, identifying six frequencies and their degree and order. Subsequent modeling has placed constraints on the mass, age, metallicity, convective core overshooting ($\\alpha_{ov}$) and internal rotation profile for this star \\citep{dupret04}. They find strong evidence for $\\alpha_{ov}$ = 0.1 $\\pm$ 0.05 in the absence of rotational mixing. Although the rotation rate for this star is quite slow, around 2 \\kms, the observed frequencies can not be matched with a uniformly rotating model. A second $\\beta$ Cephei star, $\\nu$ Eridani, has also been studied extensively with both photometric and spectroscopic campaigns \\citep{handler04,aerts04nueri}. Nine modes were detected for this star, including the radial mode and two $\\ell$ = 1 triplets \\citep{deridder04}. Modeling of $\\nu$ Eridani has also shown that non-uniformly rotating models are required \\citep{pamyatnykh04, ausseloos04}. As for V836 Cen, this star cannot be uniformly rotating to match the observed frequencies, and some convective core overshooting may be required. There is also some indication that the interior chemical composition is not homogeneous, with Fe overabundant in the driving zone. Several other $\\beta$ Cephei stars have also been successfully modeled. Using observations by the MOST satellite, \\citet{aerts06} were able to place constraints on the physical parameters of $\\delta$ Ceti, including constraining the convective core overshoot ($\\alpha_{ov}$ = 0.2 $\\pm$ 0.05). Seismic modeling of $\\beta$ CMa has constrained the core overshoot to $\\alpha_{ov}$ = 0.2 $\\pm$ 0.05, as well as placing constraints on the mass and age of the star \\citep{mazumdar06}. Most recently, \\citet{briquet07} have successfully modeled the $\\beta$ Cephei star $\\theta$ Ophiuchi using rigid rotation. This star has been the subject of both photometric and spectroscopic campaigns \\citep{handler05,briquet05}, with 7 frequencies identified. These frequencies are thought to be the radial fundamental, one $\\ell$ = 1 triplet, and three components of an $\\ell$ = 2 quintuplet based on both spectroscopic and photometric mode identification. Recent independent modeling agrees with the $\\ell$ identifications, although in some cases they assign a different $m$ to the modes \\citep{dd09}. Spectroscopic observations were also used to determine the metallicity of $\\theta$ Ophiuchi, with a best value of Z = 0.0114 using the new Asplund mixture \\citep{briquet07}. Based on these observations, the best fitting model for $\\theta$ Oph has been found to have a mass of about 8.2 \\msun, X$_c$ = 0.38, T$_{eff}$ = 22 260, log L/L$_{\\odot}$ = 3.85 and a rotational velocity of about 30 \\kms\\ \\citep{briquet07}. The convective core overshoot determined for $\\theta$ Ophiuchi is surprisingly large, $\\alpha_{ov}$ = 0.44 $\\pm$ 0.07. This result is more than double that found for similar $\\beta$ Cephei stars, which have $\\alpha_{ov}$ around 0.1-0.2. It is possible that the unusually high overshoot is a result of 1D modeling, which does not take into account rotational effects on the evolution. Although $\\theta$ Oph is relatively slowly rotating for a B star, its rotation velocity , $vsini$ $\\sim$ 30 \\kms, is significantly higher than the stars discussed above, which have rotation velocities around 2 \\kms\\ (V836 Cen) and 6 \\kms\\ ($\\nu$ Eri). Rotation and convective core overshoot may be complimentary effects, in which case including rotation could reduce the amount of convective core overshoot required to match the observed frequencies. The $vsini$ of $\\theta$ Oph is about 30 \\kms and the star is thought to be viewed nearly equator on \\citep{briquet05}. Although this is not particularly rapid rotation for a star of this type, the rotation should produce some effect on the structure and frequencies. We have chosen to model $\\theta$ Oph using models which are uniformly rotating on the ZAMS. In this work, we use 2D stellar evolution and linear adiabatic pulsation codes to determine the effects of rotation and overshoot on pulsation frequencies. By including a 2D treatment of rotation, we investigate whether the overshoot of $\\theta$ Oph could be reduced while still matching the observed frequencies. This paper is organized as follows. In section \\ref{method}, we discuss the stellar evolution and pulsation calculations. In section \\ref{structure} we recall the effects of convective overshoot on the structure of the star, and in section \\ref{overshoot} we discuss the effects of rotation and overshoot on the observed frequencies in these models. Using the resulting variation we attempt to determine the rotation rate and overshoot using the observed frequencies in section \\ref{idit}. As asteroseismic modeling is more commonly done using 1D models, we perform a ``hare and hound\" exercise to determine how different the results from the two methods can be. This exercise is discussed in section \\ref{HH}. Finally, we constrain the rotation and overshoot of $\\theta$ Oph in section \\ref{thetaoph}. Our results are summarized in section \\ref{conclusions}. ", "conclusions": "\\label{conclusions} We have found that both rotation and overshoot do have an effect on stellar frequencies, although this is predominately through the effects of changing stellar radius. Increasing convective core overshoot increases the unscaled frequencies, while rotation causes the unscaled frequencies to decrease, although certain choices of scaling can minimize this effect for rotation. Although the frequencies themselves may change, as expected we find that the overshoot has no effect on the mode splitting. We also investigated the effect of overshooting on the large and small separations. We find that increasing convective core overshoot decreases the large separations, but increases the small separation. Unfortunately, both of these effects are similar to the effects of rotation noted by \\citet{me09}, and may not be useful for disentangling the effects of rotation and overshoot. We have attempted to use the changes in frequency to constrain the best fitting convective core overshoot. Using individual modes, we were able to easily constrain the convective core overshoot, although not the rotation. As should be expected, different modes gave different results, so we developed a modified $\\chi^2$ statistic, simultaneously fitting all known frequencies for models of known mass and age (9.5 \\msun, 15.6 Myr). We found this was most effective if the $\\ell_o$s of all frequencies were known. When the $\\ell_o$s are included, we are able to correctly determine both rotation and overshoot for models in our grid as well as for three test models with slightly different $v_{eq}$ and $\\alpha_{ov}$. Although the uncertainties on our results are large as a result of the coarse grid spacing used, in all cases the best solutions surrounded the true parameters of the test models. We also conducted a hare and hound exercise using frequencies calculated using 2D stellar models and pulsation calculations. We found that using 1D stellar models, we were able to find a model that reproduced the frequencies for reasonably close values of age, mass, rotation and convective core overshoot. The 1D models were able to find a velocity quite close to the true rotation rate, although the determination of core overshoot was not as accurate. For a model with true parameters (108, 0.25), the 1D models returned a best fit of (125, 0.1), while the 2D models found a best fit of (145, 0.28). The mass and age returned by the 1D models were also quite close to the true values. Finally, we applied these methods to $\\theta$ Oph, a rotating $\\beta$ Cephei star with seven observed frequencies. Using models at a fixed mass and metallicity, we find an overshoot slightly lower than that determined by \\citet{briquet07}, with $\\alpha$ = 0.28 $\\pm$ 0.1, as expected for our higher metallicity models. The rotation velocity for this model agrees well with the observed value, around 30 \\kms. We also find that a good match to $\\theta$ Oph requires a more massive star than determined by \\citet{briquet07}, around 9.5 \\msun\\ vs 8.2 \\msun\\ in their models. This difference in mass could be a a consequence of including rotation in our models, which tends to make models appear less massive, but may also be a result of the higher metallicity used in our models. Decreasing the convective core overshoot from 0.44 to 0.28 H$_p$ brings $\\theta$ Oph closer to the range of overshoots found in other $\\beta$ Cephei stars, around 0.1-0.2 H$_{p}$, but is still high. As discussed above, (Section \\ref{HH}), the fact that the 1D calculation returns a lower overshoot than the 2D models while the 1D calculations find a higher overshoot for $\\theta$ Oph is probably a result of the 1D treatment of rotation. At slow rotation velocities ($<$ 100 \\kms), like $\\theta$ Oph, the 1D method works well, while for more rapidly rotating models, as in the hare and hound exercise performed here, the 1D treatment of rotation results in an underestimate of the convective core overshoot." }, "1003/1003.0030_arXiv.txt": { "abstract": "We report the discovery of two strongly-lensed $z\\sim 3$ Lyman Break Galaxies (LBGs) discovered as $u$-band dropouts as part of the SDSS Giant Arcs Survey (SGAS). The first, SGAS J122651.3+215220 at $z=2.9233$ is lensed by one of several sub-clusters, SDSS J1226+2152, in a complex massive cluster at $z=0.43$. Its ($g,r,i$) magnitudes are ($21.14,20.60,20.51$) which translate to surface brightnesses, $\\mu_{g,r,i}$, of ($23.78,23.11,22.81$). The second, SGAS J152745.1+065219, is an LBG at $z=2.7593$ lensed by the foreground SDSS J1527+0652 at $z=0.39$, with $(g,r,z)$=($20.90,20.52,20.58$) and $\\mu_{g,r,z}$=($25.15,24.52,24.12$). Moderate resolution spectroscopy confirms the redshifts suggested by photometric breaks, and shows both absorption and emission features typical of LBGs. Lens mass models derived from combined imaging and spectroscopy reveal that SGAS J122651.3+215220 is a highly magnified source ($M \\simeq 40$), while SGAS J152745.1+065219 is magnified by no more than $M\\simeq 15$. Compared to LBG survey results \\citep{steidel03}, the luminosities and lensing-corrected magnitudes suggest that SGAS J122651.3+215220 is among the faintest $\\simeq 20\\%$ of LBGs in that sample. SGAS J152745.1+065219, on the other hand, appears to be more representative of the average LBG, similar to the ``Cosmic Eye''. \\subjectheadings{galaxies: formation, galaxies: high-redshift, gravitational lensing, cosmology: early universe} ", "introduction": "In the sequence of gravitational collapse, heating, stellar ignition and death that govern the evolution of baryons in the earliest overdensities, Lyman Break Galaxies (LBGs) serve as high-redshift way points on the path to the $z \\sim 0$ galaxies we observe today \\citep[e.g.,][]{adelberger98,steidel03}. In star-forming galaxies at $z \\sim 3$, the Lyman continuum break at 912~\\AA~ resides in blue optical bands ($u,g$), while the continuum itself can be detected at redder wavelengths ($g,r$). These features have motivated the construction of photometric surveys that rely on this ``dropout'' technique \\citep[][and references therein]{steidel95} to select hundreds of likely LBGs at $z \\sim 3$ \\citep{steidel96a,adelberger03,steidel03}. However, typical samples of LBGs consist of tens of objects with fluxes too faint to permit detailed spectroscopic observations \\citep[e.g.,][]{nesvadba06} of individual systems. \\citet{shapley03} addressed this shortcoming by creating a high S/N composite spectra from low S/N spectra of $\\simeq 1000$ LBGs to infer the properties of the average LBG. Alternatively, since the discovery of MS1512-cB58 at $z=2.7$ \\citep[cB58;][]{yee96}, strongly-lensed LBGs magnified tens of times have offered a high S/N window \\citep{teplitz00,pettini02} into the conditions of individual star-forming galaxies when the Universe was less than 2 Gyrs old. For example, the presence of ultraviolet absorption lines in cB58 \\citep{pettini00,pettini02} reveal a chemically diverse ISM, which suggests that most of the metal enrichment occurred within the previous $\\sim 300$ Myr, and that the energetic star-formation drives a bulk outflow of the ISM that exceeds the star-formation rate. Studies of cB58 with \\emph{Spitzer} \\citep[e.g][]{siana08} have recently brought the IR observations into the picture, suggesting that the UV-inferred star-formation rate is a factor of $\\sim 3-5$ lower than that measured in the IR. Following the discovery of cB58, three more strongly-lensed LBGs have been added: 1E0657-56 at $z=3.24$ \\citep{mehlert01}, the 8 o'clock arc \\citep{allam07} at $z=2.73$, and the Cosmic Eye at $z=3.07$ \\citep{smail07}. Further spectroscopy \\citep{finkelstein09}, space-based IR \\citep{siana09}, and millimeter \\citep{coppin07} observations of the latter two systems have begun to fill out our understanding of $z=3$ LBGs but still leave many unanswered questions. In this Letter, we report the discovery of two more strongly-lensed LBGs, and include a description of their basic properties. Given the burgeoning rate of discovery of large samples of lensed sources \\citep[e.g.,][]{gladders05,cabanac07,hennawi08} we refrain from assigning nicknames to the sources discussed here; instead, we introduce designations of the form SGAS JXXXXXX+XXXXXX to denote lensed sources and SDSS J????+???? to denote the cluster lenses. Where necessary, we assume a flat $(\\Omega_m,\\Omega_\\Lambda)=(0.3,0.7)$ cosmology with $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. \\newpage ", "conclusions": "While we require additional deep imaging to better constrain the lens models for SGAS J122651.3+215220 and SGAS J152745.1+065219, the existing MagE spectra, other spectroscopy, and additional IR and UV imaging form a basis for a series of forthcoming papers that investigate the intrinsic properties of the LBGs in this study (Rigby et al., 2010). These discoveries nearly double the number of known lensed LBGs. The ambitious follow-up program built into SGAS will continue to grow the $z \\sim 3$ lensed-LBG sample as well as other lensed source populations (e.g., Ly-$\\alpha$ emitters, Bayliss et al., 2010). The systematic nature of this program offers the possibility of building large samples of strongly-lensed high redshift galaxies whose high fluxes enable both the acquisition of high S/N spectra and a look at the faint end of the luminosity function." }, "1003/1003.2205.txt": { "abstract": "% context heading (optional) {Star formation theories are currently divergent regarding the fundamental physical processes that dominate the substellar regime. Observations of nearby young open clusters allow the brown dwarf (BD) population to be characterised down to the planetary mass regime, which ultimately must be accommodated by a successful theory.} % aims heading (mandatory) {We hope to uncover the low-mass population of the $\\rho$ Ophiuchi molecular cloud and investigate the properties of the newly found brown dwarfs.} % methods heading (mandatory) {We use near-IR deep images (reaching completeness limits of approximately 20.5~mag in \\emph{J}, and 18.9~mag in \\emph{H} and \\emph{K$_{s}$}) taken with the Wide Field IR Camera (WIRCam) at the Canada France Hawaii Telescope (CFHT) to identify candidate members of $\\rho$ Oph in the substellar regime. A spectroscopic follow-up of a small sample of the candidates allows us to assess their spectral type, and subsequently their temperature and membership.} % results heading (mandatory) {We select 110 candidate members of the $\\rho$ Ophiuchi molecular cloud, from which 80 have not previously been associated with the cloud. We observed a small sample of these and spectroscopically confirm six new brown dwarfs with spectral types ranging from M6.5 to M8.25.} % conclusions heading (optional), leave it empty if necessary {} ", "introduction": "\\label{introduction} The determination of the initial mass function (IMF) across the entire stellar and substellar mass spectrum is a fundamental constraint for star formation theories \\citep[see, for example,][and references therein]{Bonnell2007}. Although there are general accepted views on the way star formation occurs and young stellar objects (YSOs) evolve to the main sequence \\citep{Shu1987,Larson1973}, the existing theories have not yet converged to an agreed paradigm that can explain the wide range of existing observational properties of YSOs. In particular, since their discovery, hundreds of brown dwarfs (BD) with masses down to the planetary regime have been uncovered in star-forming regions and the solar neighbourhood, with a ratio of the number of BDs to stars of approximately 1/5 \\citep[see, for example,][ and references therein]{Luhman2007c}, implying that a successful star and planet formation theory must account for them. Different theories for the formation of BDs are currently debated, according to which they could either form by gravitational fragmentation and collapse of molecular cores \\citep{Padoan2007,Hennebelle2008}, from early ejection from stellar embryos \\citep{Reipurth2001,Whitworth2005}, or from fragmentation of massive circumstellar discs \\citep{Stamatellos2009}. The extension of the IMF to the brown dwarf and planetary mass regime and the search for the end of the mass function is therefore crucial to determine the dominant formation process of substellar objects and its relation with the surrounding environment \\citep{Moraux2007,Andersen2008,Luhman2007}. Brown dwarfs are brighter when they are young \\citep{Chabrier2000} and their detection down to a few Jupiter masses can be attained with the current technology by studying them in young star-forming regions \\citep{Lucas2000,ZapateroOsorio2002,Weights2009,Burgess2009,Marsh2009}. For that reason, one of the prime goals of modern observations is to achieve completeness at the lower mass end, i.e., the brown dwarf and planetary mass regime, for different environments across several young star-forming regions \\citep[][among many others]{Bihain2009,Bouy2009b,Bouy2009a,Lodieu2009,Luhman2009,Scholz2009}. The main motivation of our survey of the $\\rho$~Ophiuchi molecular cloud is to uncover the low-mass population of the cluster down to the planetary regime. Despite being one of the youngest ($\\sim$1~Myr) and closest star-forming regions \\citep[120 to 145~pc,][]{Lombardi2008,Mamajek2008}, the high visual extinction in the cloud's core, with A$_V$ up to 50-100 mag \\citep{Wilking1983}, make it one of the most challenging environments to study low-mass YSOs. The main studies previously conducted in $\\rho$~Oph have been summarised in a recent review \\citep{Wilking2008}, which includes a census with the $\\sim$300 stellar members that have been associated with the cloud up to now, from which only 15 are estimated to have masses in the substellar regime. \\citet{Marsh2009} reported the discovery a young brown dwarf with an estimated mass of $\\sim$2~$-$3~Jupiter masses in $\\rho$~Oph, although we here question its membership to the cloud (see Sect. ~\\ref{comp:surveys}). We conducted a deep near-IR (\\emph{J}, \\emph{H}, and \\emph{K$_{s}$}) photometric survey centred approximately on the cloud's core and covering $\\sim$1~deg$^{2}$, which we use to identify candidate members in the substellar mass regime. Near-IR surveys are particularly suitable to study this star-forming region because most of its population is visibly obscured. Previous near-IR studies of this cluster have been done from the ground down to a sensitivity limit of \\emph{K}~$<$~13-14~mag for a larger area of the cloud \\citep{Greene1992,Strom1995,Barsony1997}, and of \\emph{K}~$<$~15.5~mag for a smaller region \\citep[200~arcmin$^2$,][]{Comeron1993}. Deeper observations were done from space with a small coverage of 72~arcmin$^2$ and a sensitivity of \\emph{H}~$<$~21.5~mag \\citep{Allen2002}. The WIRCam near-IR survey presented takes advantage of a new generation of wide-field imagers on 4 meter-class telescopes, to reach completeness limits of approximately 20.5 in \\emph{J}, and 18.9 in \\emph{H} and \\emph{K$_{s}$} over the entire degree-size area of the sky occupied by the $\\rho$~Ophiuchi central cloud. This work complements the previous surveys both in the area it covers and in sensitivity. Of comparable characteristics is the near-IR survey recently conducted by \\citet{Alvesdeoliveira2008}, which uses a different technique, near-IR variability, to select candidate members. Our selection method allows BDs with masses down to a few Jupiter masses (according to evolutionary models) to be detected through $\\sim$20 magnitudes of extinction. Extensive use of archive data at optical and IR wavelengths is made to further characterise the candidate members. In a pilot study, a spectroscopic follow-up of a subsample of these candidates has confirmed six new brown dwarfs. In Sects.~\\ref{data} and \\ref{archdata}, the observations and reductions for new and archive data are described. Section~\\ref{select:cmd} explains the methods used to select candidate members of $\\rho$~Oph and the results, and in Sect.~\\ref{discussion_phot} we discuss their properties. Section~\\ref{spec} describes the numerical fitting procedure used to analyse the data from the spectroscopic follow-up and the spectral classification. These results are then discussed through Sect. ~\\ref{properties}. Conclusions are given in Sect. ~\\ref{conclusion}. %______OBSERVATIONS_________________________________________ ", "conclusions": "\\label{conclusion} We identify 110 substellar candidate members of $\\rho$~Ophiuchi from a deep, near-IR photometric survey, from which 80 were not previously associated with the cloud. By extensive use of archive multi-wavelength data, we find evidence of mid-IR excess for 27\\% of the candidates and a variability behaviour consistent with that of YSOs for 15\\%, further supporting the membership of these candidates. We started a spectroscopic follow-up of the substellar candidate members, and present the first results for 16 sources. We identify six new members of $\\rho$~Ophiuchi with spectral types ranging from $\\sim$M6.5 to $\\sim$M8.25, and classify them as new confirmed brown dwarfs according to the evolutionary models of \\citet{Baraffe1998}. We confirm the spectral type derived by \\citet{Cushing2000} for a previously known very low-mass star close to the substellar limit, and based on the SED constructed from optical to mid-IR photometry, we report the discovery of a candidate edge-on disc around this star. We cannot derive accurate spectral types for five sources which have extremely red spectra. Two of these show water absorption features and are classified with spectral types M5 and M6. However, since they lack a \\emph{J}-band spectra and given the poor fit they remain as candidate members. The remaining three sources could be T~Tauri star members of the cluster, because they show strong mid-IR excess and one of them is emitting in X-rays. We found signatures of outflow activity in two of the sources studied spectroscopically where H$_{2}$~1~-~0 S(0) emission (2.12~$\\mu$m) was detected. Four sources out of the 16 were found to be contaminant field dwarfs." }, "1003/1003.3669.txt": { "abstract": "We performed a spectroscopic galaxy survey, complete to $m_{F814W}$$\\leq$20.3 ($L_B>0.15L_B^{\\star}$ at $z=0.3$), within 100$\\times$100$''$ of the quasar Q1127--145 ($z_{em}=1.18$). The VLT/UVES quasar spectrum contains three $z_{abs}$$<$0.33 {\\MgII} absorption systems. We obtained eight new galaxy redshifts, adding to the four previously known, and galaxy star formation rates (SFRs) and metallicities were computed where possible. A strong {\\MgII} system [$W_r(2796)=1.8$~\\AA], which is a known damped Ly$\\alpha$ absorber (DLA), had three previously identified galaxies; we found two additional galaxies associated with this system. These five galaxies form a group with diverse properties, such as a luminosity range of $0.04\\leq L_B\\leq0.63 L_B^{\\star}$, an impact parameter range of $17\\leq D \\leq 241$~kpc and velocity dispersion of $\\sigma=115$~\\kms. The DLA group galaxy redshifts span beyond the 350~{\\kms} velocity spread of the metallic absorption lines of the DLA itself. The two brightest group galaxies have SFRs of $\\sim$few $M_{\\odot}$~yr$^{-1}$ and should not have strong winds. We have sufficient spectroscopic information to directly compare three of the five group galaxies' (emission-line) metallicities with the DLA (absorption) metallicity: the DLA metallicity is 1/10th solar, substantially lower than the three galaxies' which range between less than 1/2 solar to solar metallicity. HST/WFPC--2 imaging shows perturbed morphologies for the three brightest group galaxies, with tidal tails extending $\\sim$25 kpc. We favor a scenario where the DLA absorption originates from tidal debris in the group environment. Another absorber exhibits weak {\\MgII} absorption [$W_r(2796)=0.03$~\\AA] and had a previously identified galaxy at a similar redshift. We have identified a second galaxy associated with this system. Both galaxies have solar metallicities and unperturbed morphologies in the HST/WFPC--2 image. The SFR of one galaxy is much lower than expected for strong outflows. Finally, we have also identified five galaxies at large impact parameters with no associated {\\MgII} absorption [$W_r(2796) \\lesssim 5.7$~m{\\AA}, 3~$\\sigma$] in the spectrum of Q1127--145. ", "introduction": "Absorption lines detected in the spectra of background quasars and gamma ray bursts remain one of the best probes of intervening multiphase gas throughout the Universe. Pioneering work of \\citet{bergeron88} and \\citet{bb91} led to the first galaxies identified in close proximity to a quasar sight-line and at the same redshift as metal-enriched absorption traced by the {\\MgIIdblt} doublet. Since then, there has been numerous studies of {\\MgII} absorption line systems aimed at interpreting the properties of galaxy halos at a variety of redshifts \\citep[e.g.,][]{lebrun93,sdp94,csv96,archiveII,steidel02,ellison03,bouche06,zibetti07,kacprzak08,chen08,barton09,rubin09a,pollack09,menard09}. $\\hbox{{\\rm Mg}\\kern 0.1em{\\sc ii}}$~absorption~lines are ideal for studying a large dynamic range of structures and environments in and around galaxies since they trace low ionization metal-enriched gas with neutral hydrogen column densities of $10^{16} \\lesssim \\hbox{N(\\HI)} \\lesssim 10^{22}$~{\\cmsq}\\citep{archiveI,weakII}. This large density range allows for detections of {\\MgII} in absorption out to $\\sim120$~kpc from the host galaxy \\citep{zibetti07,chen08,kacprzak08}. %A significant theoretical effort %performed in order to interpret and understand these absorption %systems using semi--analytical models and galaxy simulations Significant theoretical efforts have employed semi-analytical models and single halo galaxy simulations to interpret and understand absorption systems \\citep[e.g.,][]{mo96,burkert00,lin00,maller04,chelouche08,chen08,tinker08,kaufmann09}. These models and simulations have helped constrain halo sizes, covering fractions, gas kinematics, physical gas conditions, etc. However, the majority of these studies modeled galaxies as isolated systems/halos and lack the important dynamic influences of the cosmic structure and local environments, which also may contribute to a significant fraction of the detected {\\MgII} absorption \\citep{kacprzak10}. %the {\\MgII} cross section and equivalent with of {\\MgII} %absorption systems in galaxy halos Halo gas masses and cross sections are suggested to increase due to tidal streams produced by the interactions/minor mergers and/or increased star formation-induced winds caused by gas rich minor mergers \\citep{york86,rubin09a}. {\\MgII} gas outflowing at $\\sim500$~{\\kms} has been detected in winds of galaxies at $z\\sim1$ \\citep[e.g.,][]{tremonti07,weiner09,rubin09b}. Winds have also been suggested to be responsible for high equivalent width {\\MgII} absorbers \\citep{bouche06}. Evidence of galaxy interactions producing {\\MgII} absorption was discussed by \\citet{kacprzak07} who reported a suggestive correlation between the {\\MgII} rest equivalent width, $W_r(2796)$, and the galaxy morphological asymmetries normalized by impact parameter. They suggest that perturbations from minor galaxy mergers may be responsible for producing low equivalent width systems [$W_r(2796)<1.5$~\\AA]. These results are consistent with low redshift {\\HI} surveys where galaxies having a perturbed/warped disk, from a previous or ongoing minor merger, also have more extended {\\HI} disks/halos \\citep{puche92,swaters97,rand00,fraternali02,chynoweth08,sancisi08}. The aforementioned results suggest that galaxy environments may play a role in the metal enrichment of galaxy halos. \\citet{lopez08} performed the first statistical environmental study of absorption systems associated with 442 x-ray selected galaxy clusters ($z=$0.3--0.9) out to transverse distances of 2~h$^{-1}$~Mpc. It was determined that galaxy clusters produce a factor of 15 over-abundance of strong equivalent width systems [$W_r(2796) > 2$~\\AA] compared to field galaxies. This over-abundance is higher in the centers of clusters than in the outer parts and also increases with cluster mass. In contrast, the $dN/dz$ of weak {\\MgII} systems [$W_r(2796) < 0.3$~\\AA] in clusters is consistent with those derived from environmentally unbiased samples. \\citet{lopez08} argue that the detected over-abundance of strong systems is a result of the over-density of galaxies in a cluster region. The lack of an over-abundance of weak systems may imply they were destroyed by the cluster environment. \\citet{padilla09} modeled these results and found environmental evidence of truncated {\\MgII} halo sizes as a function of cluster radii. Their models require a median {\\MgII} halo size of $r < 10$~h$^{-1}$~kpc, compared to $35-85$~h$^{-1}$~kpc for field galaxies, in order to reproduce the observed absorption-line statistics of \\citet{lopez08}. These x-ray selected clusters represent more extreme environments than in galaxy groups where the majority of galaxies reside. In x-ray clusters, ram-pressure stripping is an important environmental effect since the intracluster medium and galaxy velocities are much higher than in galaxy groups where ram-pressure stripping is negligible \\citep{mulchaey98}. However, in groups the galaxy velocities are smaller and the interactions and mergers more frequent, resulting in increased gas covering fractions of the cool intragroup gas \\citep{zabludoff98}. A significant fraction of strong {\\MgII} absorption systems are damped Lyman alpha systems (DLA) \\citep{rao03}. Since the discovery of DLAs \\citep{wolfe86} their host galaxy properties and environments have remained largely unknown. Only a small fraction of DLA hosts have been identified and appear to be isolated galaxies in close proximity of the quasar LOS \\citep{lacy03,moller02,rao03,chun06}. Models support an array of origins of the DLA absorbing gas from rotating thick disks \\citep[e.g.,][]{prochaska97,prochaska98,prochaska02}, gas rich dwarf galaxies \\citep{matteucci97}, irregular protogalactic clumps \\citep{haehnelt98}, and tidal gas or processes such as superwinds and outflows \\citep{zwaan08}. In this paper, we perform a spectroscopic survey of the galaxies in the Q1127$-$145 quasar field. A VLT/UVES quasar spectrum shows that there are three absorption systems in this field, one of which is DLA system which has three previously identified galaxies at a similar redshift. However, this field contains many unidentified bright galaxies within 50$''$ of the quasar line of sight. We perform a spectroscopic survey, to a limiting magnitude of $m_{F814W}\\leq 20.3$, in an attempt to obtain spectroscopic redshifts for the remaining galaxies within the field. In \\S~\\ref{sec:data} we describe our sample and analysis. In \\S~\\ref{sec:results} we present the results of our redshift survey. We discuss morphologies of the galaxies and we also compute galaxy star formation rates (SFRs) and emission line metallicities when possible. We compare galaxy metallicities to the absorption line metallicity derived for the DLA. In \\S~\\ref{sec:dis}, we discuss the possible origins of the {\\MgII} absorption and our concluding remarks are in \\S~\\ref{sec:conclusion}. Throughout we adopt an H$_{\\rm 0}=70$~\\kms Mpc$^{-1}$, $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$ cosmology. ", "conclusions": "\\label{sec:conclusion} We have performed a spectroscopic galaxy survey to limiting magnitude of $m_{F814W}\\leq 20.3$ ($L_B>0.15L_*$ at $z=0.3$) within 100$\\times$100$''$ of the quasar Q1127$-$145. This field has a large number of bright galaxies near the quasar line of sight and has three {\\MgII} absorption systems detected in the quasar spectrum, including one DLA. Here we have obtained spectroscopic redshifts for eight galaxies in this field, adding to the four previously identified \\citep{bb91,gb97,kacprzak10}. Our main results can be summarized as follows: \\begin{enumerate} \\item We have identified two galaxies (G6 and G14) associated with the DLA at $z=0.313$, which, in addition to the three known galaxies, form a group of at least five galaxies. The group has a luminosity range of $0.04\\leq L_B\\leq0.63 L_B^{\\star}$ and an impact parameter range of $17.4\\leq D \\leq 240.8$~kpc. The group velocity dispersion is $\\sigma=115$~{\\kms} having a full velocity range of $\\sim 350$~\\kms. The group redshift is offset $80$~{\\kms} blueward of the {\\MgII} absorption redshift. The galaxy redshift distribution spans the entire range of the absorption velocities. Furthermore, the rotation curves of G2 and G4 alone cover the entire range of absorption velocities. Star formation rates of two of the brightest galaxy members are too low to drive strong winds, reducing the likelihood that winds are responsible for the absorbing gas. Metal enriched winds are also unlikely since the DLA metallicity is 1/10th solar, whereas three of the five galaxies have metallicities range between less than 1/2 solar to solar. Although stellar metallicity gradients in the literature are consistent with our findings, it is has yet to be demonstrated that these gradients can be extrapolated to 50~kpc. The favored scenario for the origin of the absorption is from tidal debris. The deep WFPC--2 F814W imaging shows the perturbed morphologies for three galaxies and optical tidal tails extending $\\sim 25$~kpc away from the disks. These features suggest merger/harassment events, consistent with the more frequent galaxy harassment/merging expected in the group environment we have identified. \\item We have identified a galaxy (G3), in addition to previously identified G5 \\citep{kacprzak10}, associated with the $z=0.328$ weak {\\MgII} absorption system, $W_r(2796)=0.029$~\\AA. There is no evidence of recent interactions since both galaxies have unperturbed morphologies and they are separated by 140~kpc. Even armed with the star-formation rate and rotation velocities of G5 and the metallicities of both galaxies, it remains difficult to determine which galaxy hosts the absorber. We can only conclude that this weak absorption system can arise in a variety of cosmic structures in either or both halos of the galaxy pair. \\item We have identified five galaxies (G7, G9, G11, G12, and G15) with $0.21\\leq z\\leq0.33$ that are not associated with any detectable {\\MgII} absorption (3$\\sigma$ detection limits of $4.8-5.7$~m{\\AA}). These galaxies appear to be normal star-forming spiral disks. All non-absorbing galaxies have impact parameters $D>118$~kpc. This is consistent with previous results on {\\MgII} halo sizes, which suggest we should not expect to detect absorption beyond impact parameters of $\\sim 120$~kpc. \\end{enumerate} The DLA-galaxy group at $z=0.313$ is quite different from the standard examples in literature of DLA-plus-(apparently) isolated galaxy \\citep[e.g.,][]{lacy03,moller02,rao03,chun06}. The group of galaxies associated with the $z=0.313$ DLA suggests that interactions, which are common in groups of galaxies, might be responsible for at least some DLA absorption systems as well. This may explain why searches for host galaxies of DLAs and strong {\\MgII} systems have a low success rate of 30--40\\% using small field of view IFUs \\citep[e.g.,][]{bouche07}. It is likely that we need to survey further out from the quasar line of sight if there are many other cases where tidal debris produces the absorption. It is also interesting to note that if this galaxy group was at a slightly higher redshift, we would not be able to detect the 0.04$L_B^{\\star}$ galaxy that is closest to the quasar line of sight, which could even be the DLA host. Given the low redshift of the DLA and even using the deep {\\it HST} imaging, star formation rates, and metallicities, it is difficult to understand this complex system and determine the origins of the absorbing gas. We emphasize that we should take caution in concluding the origins of absorbing gas drawn from studies of individual DLAs at higher redshifts. %conclusions reached about systems at $z\\sim2$ are therefore likely to %be biased, depending of the frequency group systems, towards only %certain conclusions about the physical nature/origin of the absorbers. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%" }, "1003/1003.3854_arXiv.txt": { "abstract": "{} {We provided accurate estimates of distances, radii and iron abundances for four metal-rich Cepheids, namely V340~Ara, UZ~Sct, AV~Sgr and VY~Sgr. The main aim of this investigation is to constrain their pulsation properties and their location across the Galactic inner disk.} {We adopted new accurate NIR (J,H,K) light curves and new radial velocity measurements for the target Cepheids to determinate their distances and radii using the Baade-Wesselink technique. In particular, we adopted the most recent calibration of the IR surface brightness relation and of the projection factor. Moreover, we also provided accurate measurements of the iron abundance of the target Cepheids.} {Current distance estimates agree within one $\\sigma$ with similar distances based either on empirical or on theoretical NIR Period-Luminosity relations. However, the uncertainties of the Baade-Wesselink distances are on average a factor of 3-4 smaller when compared with errors affecting other distance determinations. Mean Baade-Wesselink radii also agree at one $\\sigma$ level with Cepheid radii based either on empirical or on theoretical Period-Radius relations. Iron abundances are, within one $\\sigma$, similar to the iron contents provided by Andrievsky and collaborators, thus confirming the super metal-rich nature of the target Cepheids. We also found that the luminosity amplitudes of classical Cepheids, at odds with RR Lyrae stars, do not show a clear correlation with the metal-content. This circumstantial evidence appears to be the consequence of the Hertzsprung progression together with the dependence of the topology of the instability strip on metallicity, evolutionary effects and binaries. } {} ", "introduction": "Classical Cepheids are used both as standard candles and tracers of young stellar populations \\citep{mac09, mio09}. They are bright and variable objects and thanks to the Hubble Space Telescope they have been identified and accurately measured in Local Group (d $\\lesssim$ 1 Mpc) and in Local Volume (d $\\lesssim$ 10 Mpc) galaxies \\citep{free01, tamm03, bo08}. They obey to a Period-Luminosity (PL) relation and are the most popular primary distance indicators \\citep{feast99, macr06, fo07, groe07, diben08, groe08,ker08, kan09a, mar09, sco09}. In spite of these outstanding observational and theoretical efforts \\citep[][and references therein]{marc09} the universality of optical and near-infrared (NIR) PL relations still lacks a firm empirical validation \\citep{ben07, ro08, sand09}. This thorny problem remains even if the zero-point and the slope of the PL relation can be estimated with a variety of independent methods. Ideally the calibration and the validation of the PL relation should be rooted on distances measured with a geometrical method such as the trigonometric parallaxes. This approach was recently adopted by \\cite{ben07} who provided parallaxes with a mean accuracy of 8\\% for a sample of nine Galactic Cepheids using the {\\em Fine Guidance Sensor} (FGS) on board of the Hubble Space Telescope (HST). A new revision of HIPPARCOS parallaxes for Galactic Cepheids (244 objects) has been recently provided by \\cite{vanl07} and by \\cite{vanletal07}. The accuracy of the new measurements is on average a factor of two better than the old ones. However, the Cepheids with the most accurate HIPPARCOS parallaxes are Polaris ($\\alpha$ UMa) and the prototype $\\delta$ Cep. The accuracy of the former one is 1.6\\% \\citep{vanletal07}, while the latter is similar to the accuracy of the FGS@HST parallax, namely 5.2\\% versus 4.1\\% \\citep{mer05}. The Baade-Wesselink method (BW, \\cite{ba26, wess46}) provides an independent empirical approach to measure Cepheid absolute distances and it can be applied to variable stars. This method relies on two observables: the radial velocity ($v_r$) and the variation of the angular diameter ($\\theta$). The latter parameter was historically substituted by the variation in color along the pulsation cycle. However, direct measurements of the Cepheid angular diameter have been recently provided by \\cite{ker04b} using the Very Large Telescope Interferometer (VLTI). In particular, the use of VLT INterferometer Commissioning Instrument (VINCI) gave the unique opportunity to provide angular diameter measurements along the pulsation cycle for seven Cepheids. These pioneer measurements encouraged a detailed comparison between theory and observations concerning the limb darkening and the atmosphere of variable stars \\citep{mar03,nar07}, but the number of Cepheids for which these measurements are available is still very limited. \\begin{table*}\\label{mag} \\caption{Intrinsic parameters, mean NIR magnitudes and mean radial velocities for the target Cepheids.} \\begin{center} \\begin{tabular}{l c c c c c c c c c c} \\hline\\hline NAME &$\\alpha$(J2000)$^a$&$\\delta$(J2000)$^a$&$\\log$ P&E(B-V)$^b$&$^c\\pm \\sigma(J)$&$^c\\pm \\sigma(H)$&$^c\\pm\\sigma(K)$&$N_s^d$&$^e$&$\\Delta v_r^e$\\\\ \\hline V340~Ara& 16 45 19&-51 20 33&1.32&0.574&7.382$\\pm$0.011&6.809$\\pm$0.007&6.619$\\pm$0.008&25+2&-80.8$\\pm$1.2&59.1 \\\\ UZ~Sct & 18 31 22&-12 55 00&1.17&1.071&7.502$\\pm$0.049&6.818$\\pm$0.042&6.564$\\pm$0.045&25+2& 40.2$\\pm$0.5&49.4 \\\\ AV~Sgr & 18 04 49&-22 43 00&1.19&1.267&6.909$\\pm$0.040&6.081$\\pm$0.035&5.758$\\pm$0.033&25+1& 19.3$\\pm$1.8&59.3 \\\\ VY~Sgr & 18 12 05&-20 42 00&1.13&1.283&7.174$\\pm$0.069&6.375$\\pm$0.046&6.068$\\pm$0.040&25+2& 16.0$\\pm$1.8&59.2 \\\\ \\hline \\multicolumn{11}{l}{$^a$ Cepheid coordinates: units of right ascension are hours, minutes, and seconds; units of declination are degrees, arcminutes and arcseconds.} \\\\ \\multicolumn{11}{l}{$^b$ Reddening according to \\cite{fer95}.} \\\\ \\multicolumn{11}{l}{$^c$ Mean NIR magnitudes.} \\\\ \\multicolumn{11}{l}{$^d$ Number of spectra collected with FEROS at the 2.2m MPG/ESO telescope.} \\\\ \\multicolumn{11}{l}{$^e$ Mean radial velocity and velocity amplitude (km s$^{-1}$).} \\\\ \\end{tabular} \\end{center} \\end{table*} To overcome the difficulties in the interferometric measurement of the angular diameter several variants of the BW method were suggested in the literature. Among them the methods based on the Surface-Brightness (SB) relations link variations in color with variations in angular diameters. This method can be applied to Cepheids for which accurate radial velocities and multi-wavelength light curves are available. \\cite{sto04}, \\cite{ker04a} and \\cite{groe07} (hereafter G07) derived such relations, using an optical-NIR ($V$-$K$) color (IRSB), which gives the highest precision in the derived quantities. However, the most relevant limit of the currently adopted BW methods is the value of the projection factor ($p$-factor). This parameter links radial velocity changes to radius changes and still lacks firm theoretical and empirical constraints \\citep{nar09}. It can be empirically estimated using Cepheids for which accurate interferometric angular diameter measurements, radial velocity curves and trigonometric parallaxes are available. This approach was applied to $\\delta$ Cep by \\cite{mer05} using the new optical interferometric measurements obtained with the CHARA Array, the FGS@HST trigonometric parallax by \\cite{ben07} and the radial velocities available in the literature. The $p$-factor they found --$p$=$1.27\\pm0.06$-- agrees quite well with theoretical predictions by \\cite{nar04b}. More recently, G07 found that the use of a constant $p$-factor ($p$=$1.27\\pm0.05$) for six Galactic Cepheids, with interferometrically measured angular diameter variations and known distances, agrees quite well with HST parallaxes. Moreover, he found that a strong period dependence of the $p$-factor ($p \\sim -0.15\\cdot \\log P$, \\cite{gie05}), could also be ruled out. However, a moderate period dependence ($p \\sim -0.03 \\cdot \\log P$) as suggested either by \\cite{gie93, gie97,gie98, bar03, sto04}, or more recently by \\cite{nar09} ($p \\sim -0.08 \\cdot \\log P$) is still consistent with currently available data (G07). In this investigation we plan to apply the BW method using the most recent calibration of the IRSB relation (Groenewegen 2010, hereinafter G10, in preparation), to estimate the distance of four metal-rich Galactic Cepheids. In particular, we plan to use new radial velocity measurements, new accurate NIR ($J$,$H$,$K$) light curves and $V$-band light curves available in the literature \\citep{ber92}. Moreover, we discuss in \\S 4 the iron abundance of the four target Cepheids using \\ion{Fe}{i} and \\ion{Fe}{ii} lines. \\S 5 deals with the pulsation amplitude of metal-rich Cepheids, while in \\S 6 we summarize current findings and outline future developments of this project. \\begin{figure} \\begin{center} \\includegraphics[height=10.5truecm, width=7.25cm,angle=0]{curveJHK.ps} \\vspace{0.8cm} \\caption{From left to right $J$,$H$,$K$-band light curves for the four metal-rich selected Cepheids. In each panel are also plotted the intrinsic scatter ($\\sigma$) of the fit with a cubic spline (red line) and the luminosity amplitude.}\\label{LCJHK} \\end{center} \\end{figure} ", "conclusions": "We provided accurate BW distances and radii for four metal-rich Cepheids, namely V340~Ara, UZ~Sct, AV~Sgr and VY~Sgr. Current distance estimates, taken at face value, agree quite well with similar estimates based either on empirical \\citep{fo07} or on theoretical NIR PL relations. However, the uncertainties affecting the BW distances, by summing in quadrature the errors in the fit and the errors estimated with the Monte Carlo simulations, are on average a factor of 3-4 smaller than for distances based on predicted and empirical NIR PL relation. The same outcome applies to the mean Cepheid radii, but the uncertainties on the BW radii, by summing in quadrature errors on the fit and errors from Monte Carlo simulations, are on average a factor of two larger than for radii based either on the empirical or on the theoretical PR relation provided by G07 and by \\cite{petr03}, respectively. We also collected high-resolution, high signal-to-noise ratio spectra to measure the iron abundances of the target Cepheids. Special attention was paid to provide accurate estimates of intrinsic parameters (effective temperature, surface gravity, microturbulent velocity) directly from observed spectra. We performed detailed measurements of iron abundances using large samples of \\ion{Fe}{i} and \\ion{Fe}{ii} lines. Current abundances indicate that selected Cepheids are super metal-rich and agree, within 1$\\sigma$, with iron abundances provided by \\cite{and02a} using a similar approach. We adopted a sample of 259 Galactic Cepheids for which are available either spectroscopic iron measurements or metallicity estimates based on the Walraven metallicity index \\citep{ped08}. We found that classical Cepheids do not seem to show in the Bailey diagram (luminosity amplitude vs pulsation period), in contrast with low-mass helium burning RR Lyrae stars, a clear correlation between luminosity amplitude and metallicity. The lack of such a correlation might be the consequence of the Hertzsprung progression. We also found that a good fraction of metal-rich ([Fe/H]$\\ge$0.13 dex) Cepheids are located among long-period ($\\log P \\ge$ 1.0) variables. However, for the moment this can only be considered as circumstantial evidence, since the current sample is probably affected by selection bias. Metal-rich Cepheids are located in the inner disk and are typically affected by large extinctions. The selected Cepheids have Galactocentric distances smaller than 6.5 kpc and their reddening ranges from 0.6 to 1.3 mag. Detailed analysis concerning the pulsation properties of metal-rich Cepheids best awaits more complete samples. However, the game is worth the candle, since classical Cepheids are excellent tracers of young stellar populations. Their pulsation properties and their radial distribution across the inner Galactic disk and the bar can provide robust constraints on the bar-driven formation scenario \\citep{vanloo03, deba04, zoc06} on short (10-100 Myr) timescales." }, "1003/1003.0362.txt": { "abstract": "{}{}{}{}{} % 5 {} token are mandatory \\abstract % context heading (optional), leave it empty if necessary {Overall spherically symmetric, geometrically thin gas and dust shells have been found around a handful of asymptotic giant branch (AGB) carbon stars. Their dynamical ages lie in the range of 10$^3$ to 10$^4$ years. A tentative explanation for their existence is that they have formed as a consequence of mass-loss-rate modulations during a He-shell flash.} % aims heading (mandatory) {The detached shells carry information on their formation process, as well as on the small-scale structure of the circumstellar medium around AGB stars due to the absence of significant line-of-sight confusion.} % methods heading (mandatory) {The youngest detached shells, those around the carbon stars R~Scl and U~Cam, are studied here in great detail in scattered stellar light with the Advanced Survey Camera on the Hubble Space Telescope. Quantitative results are derived assuming optically thin dust scattering.} % results heading (mandatory) {The detached dust shells around R~Scl and U~Cam are found to be consistent with an overall spherical symmetry. They have radii of 19$\\farcs$2 (corresponding to a linear size of 8$\\times$10$^{16}$\\,cm) and 7$\\farcs$7 (5$\\times$10$^{16}$\\,cm), widths of 1$\\farcs$2 (5$\\times$10$^{15}$\\,cm) and 0$\\farcs$6 (4$\\times$10$^{15}$\\,cm), and dust masses of 3$\\times$10$^{-6}$ and 3$\\times$10$^{-7}$\\,$M_\\odot$, respectively. The dynamical ages of the R~Scl and U~Cam shells are estimated to be 1700 and 700\\,yr, respectively, and the shell widths correspond to time scales of 100 and 50\\,yr, respectively. Small-scale structure in the form of less than arcsec-sized clumps is clearly seen in the images of the R~Scl shell. Average clump dust masses are estimated to be about 2$\\times$10$^{-9}$\\,$M_\\odot$. Comparisons with CO line interferometer data show that the dust and gas shells coincide spatially, within the errors ($\\le$\\,1$\\arcsec$ for U~Cam and $\\approx$\\,2$\\arcsec$ for R~Scl).} % conclusions heading (optional), leave it empty if necessary {The results are consistent with the interpretation of geometrically thin gas and dust shells formed by a mass-loss eruption during a He-shell flash, and where interaction with a previous wind plays a role as well. The mass loss responsible for the shells must have been remarkably isotropic, and, if wind interaction plays a role, this also applies to the mass loss prior to the eruption. Clumpy structure is present in the R~Scl shell, possibly as a consequence of the mass loss itself, but more likely as a consequence of instabilities in the expanding shell.} ", "introduction": "Mass loss from the surface is an important characteristic of stellar evolution on the asymptotic giant branch (AGB). It is a common property of most M-type, S-type, and all C-type AGB stars that has been well established for decades \\citep{olofetal93a, ramsetal09}. Yet, many of its details, such as the mechanism behind it, and its evolution with time and e.g. dependence on stellar mass, are essentially unknown and remain the major obstacle for understanding stellar evolution on (and beyond) the AGB in detail as well as the contribution that AGB stars make to the galactic chemical evolution \\citep{habi96, will00, schrsedl01}. It is particularly important to understand the temporal evolution and the dependence on direction of the stellar mass loss. The former determines to a large extent how the star evolves, while the latter has a profound effect on the circumstellar evolution beyond the AGB, e.g., the formation of planetary nebulae. Geometrically thin detached shells \\citep{schoetal05b} are a phenomenon that bears on both issues. Some carbon stars, of which less than ten are known, show this phenomenon. It has been suggested that these shells are the result of strong mass-loss modulations during a thermal pulse \\citep{olofetal90,schretal98,wachetal02} and that they are additionally affected by an interaction with the surrounding (relic) circumstellar envelope (CSE) \\citep{stefscho00, schoetal05b}. Recently, \\citet{mattetal07} presented models where the response to structure variations during a thermal pulse of the dynamical atmosphere and the expanding gas and dust are studied in some considerable detail. Geometrically thin shells appear under certain circumstances as an effect of a mass-loss eruption and a subsequent interaction with a previous slower wind. Remarkably, these shells imply that at least during this phase, the mass loss is very close to isotropic \\citep{olofetal96, olofetal00, gonzetal03a, maeretal10}. The study of thin, detached shells also has a bearing on the small-scale structure of the circumstellar medium, since the line-of-sight confusion is limited by the thinness of the shells \\citep{olofetal00}. Most of the information on these detached shells stems from CO radio line observations. The first detections were made by \\citet{olofetal88}, and to this date there are seven known carbon stars with this type of shell \\citep{schoetal05b}. \\citet{lindetal99} and \\citet{olofetal00} used the IRAM PdB mm-wave interferometer to show that the shells are geometrically thin (width/radius $<$\\,0.1) and remarkably spherically symmetric. \\citet{schoetal05b} found evidence of effects of interacting winds, i.e., the shells are affected by their progress in a previous slower stellar wind. \\citet{gonzetal01} showed for the first time that these shells could be imaged in stellar light scattered in the circumstellar medium, and \\citet{gonzetal03a} and \\citet{maeretal10}, following up on this using polarimetric imaging, showed that both dust and atoms act as scattering agents. No M-type or S-type AGB star was found with similar geomtrically thin shells, despite extensive searches \\citep[see e.g.,][]{kersolof99, ramsetal09}. Detached dust shells around AGB and post-AGB objects have been seen as well, but with much coarser resolution observations \\citep{wateetal94, izumetal96, izumetal97, hashetal98, specetal00}. We note here that some detached shells around AGB stars may have a different origin. Neutral hydrogen 21\\,cm observations of extended CSEs show that some AGB stars are surrounded by large detached shells, whose presence is most likely due to interactions between circumstellar winds of different epochs or an interaction with the surrounding interstellar medium \\citep{libeetal07}. Similar conclusions are drawn based on Spitzer observations \\citep{wareetal06}. In the CO mm observations by \\citet{olofetal96, olofetal00} it was found that the detached shells showed a clumpy structure. The optical observations by \\citet{gonzetal01} verified this finding optically to some extent, but the angular resolution was severely limited by seeing. Moreover, scattering in the terrestrial atmosphere limited the studies of the shells close to the star. In order to increase the resolution and study the circumstellar envelope closer to the star, we decided to use the Hubble Space Telescope (HST). We present broadband filter images of the circumstellar environments of the carbon stars R~Sculptoris (R~Scl) and U~Camelopardalis (U~Cam) obtained with the Advanced Camera for Surveys (ACS) on the Hubble Space Telescope (HST). % ---------------------------------------------------------------- ", "conclusions": "\\subsection{Characteristics of the detached shells} At the distances of R~Scl and U~Cam the angular radii of the shells correspond to linear radii of 8$\\times$10$^{16}$\\,cm and 5$\\times$10$^{16}$\\,cm, respectively. \\citet{schoetal05b} estimate that the expansion velocities for the gas shells of R~Scl and U~Cam are 15.5\\,km\\,s$^{-1}$ and 23.0\\,km\\,s$^{-1}$, respectively, and, since there appears to be no separation between the gas and the dust shells, the corresponding dynamical ages of the shells are consequently about 1700 yr (R~Scl) and 700 yr (U~Cam). Thus, these are the youngest detached shells observed in detail so far. The average shell widths convert into linear widths of 5$\\times$10$^{15}$\\,cm and 4$\\times$10$^{15}$\\,cm for R~Scl and U~Cam, respectively, and the corresponding time scales are 100 yr (R~Scl) and 50 yr (U~Cam). The shell width/radius ratios lie in the range 0.05\\,$-$\\,0.1. The estimated dust and gas shell masses are 3$\\times$10$^{-6}$\\,$M_{\\odot}$ and 2.5$\\times$10$^{-3}$\\,$M_{\\odot}$, respectively, for R~Scl and 3$\\times$10$^{-7}$\\,$M_{\\odot}$ and 1$\\times$10$^{-3}$\\,$M_{\\odot}$, respectively, for U~Cam. To this we can add the recent result by \\citet{maeretal10} for the carbon star U~Ant, where the shell has a dynamical age of 2700 yr, a width corresponding to 140 yr, and the dust and gas masses are estimated to be 5$\\times$10$^{-5}$\\,$M_{\\odot}$ and 2.5$\\times$10$^{-3}$\\,$M_{\\odot}$. Our new dust mass estimates agree with the finding by \\citet{schoetal05b} that the shell masses increase with the age of the shell, which is possibly an effect of interaction between the shell and a previous slower stellar wind. The dust-to-gas ratio also increases with the age of the shell (3$\\times$10$^{-4}$, 1$\\times$10$^{-3}$, and 2$\\times$10$^{-2}$ for U~Cam, R~Scl, and U~Ant, respectively). This is curious, because one expects the dust to gradually leave the gas shell as it expands \\citep{maeretal10}, at least if the wind is smooth. Clumpiness, possibly arising as the shell expands, and/or wind-wind interaction may be an explanation here as this will affect the gas-grain drift and may lead to accumulation of material. \\subsection{Origin of the detached shells} It has been argued by various authors that the geometrically thin, detached shells found around carbon stars are the effects of mass-loss rate modulations during thermal pulses \\citep{olofetal90, olofetal96, schretal98, gonzetal03a, schoetal05b, mattetal07, maeretal10}. In particular, the detailed models of \\citet{mattetal07} show that geometrically thin shells appear under certain circumstances as an effect of a mass-loss eruption and a subsequent interaction with a previous slower wind. The results on R~Scl and U~Cam presented here are consistent with such a scenario. They add further support to the presence of highly isotropic mass loss during the shell formation event, and, if interaction plays a major role, this must apply also to the mass loss prior to the eruption. If this interpretation is correct, both stars must be in the aftermath of a very recent thermal pulse ($<$\\,2000 years ago). The shell formation time scale is as short as $<$\\,100\\,yr. This is an estimate that is likely not affected by interaction with a previous slow wind because of the relative youth of the shells. The amount of mass ejected is relatively small, $\\approx$\\,10$^{-3}$\\,$M_\\odot$. An estimate of the mass-loss rate during the formation of the shells is obtained using the shell gas mass estimates of \\citet{schoetal05b} and the time scales estimated here. The results are 3$\\times$10$^{-5}$ and 2$\\times$10$^{-5}$\\,$M_{\\odot}$\\,yr$^{-1}$ for R~Scl and U~Cam, respectively. This ignores any effects of e.g sweeping up of material. An estimate of the mass-loss rate prior to the shell ejection is difficult to give. We limit ourselves here to an estimate of the density contrast between the shell and the surrounding medium. In the light of optically thin scattering, the density contrast between the shell and the medium outside the shell is at least a factor of 30 in the case of R~Scl (for which the data has the highest S/N). This should be a rough estimate of the mass-loss rate contrast as well. The mass-loss rate following the ejection of the shell is difficult to estimate from our data, because it depends critically on the psf subtraction and the properties of the grain scattering (in particular its angular dependence). \\citet{schoetal05b} estimated, based on CO radio line modelling, that the present mass-loss rates are 3$\\times$10$^{-7}$ and 2$\\times$10$^{-7}$\\,$M_{\\odot}$\\,yr$^{-1}$ for R~Scl and U~Cam, respectively. In conclusion, the mass-loss rate modulations are substantial, of the order of two magnitudes. \\begin{figure} \\centering \\includegraphics[width=8.5cm]{simim814.eps} \\caption{A simulated image of the detached shell around R~Scl seen in dust-scattered light in the f814 filter (see the text for details).} \\label{f:simimage} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=7cm]{simraaps_sharp.eps} \\includegraphics[width=7cm]{simfftap_sharp.eps} \\caption{Same as Fig.~\\ref{f:rsclraaps}, but with the data input from the sharpened version of the simulated image in Fig.~\\ref{f:simimage} (see text for details).} \\label{f:simraaps} \\end{figure} \\subsection{Origin of clumps in the detached shells} The small-scale structure of the circumstellar medium is of interest for various reasons, e.g., it will affect the radiative transfer and hence the estimates (like those of mass-loss rates and molecular abundances) obtained from modelling molecular line emission, it will affect the circumstellar chemistry (incl. the photodissociation), and it will affect the efficiency of the mass loss. The observational evidence for clumpiness in AGB CSEs is scattered and not very conclusive. High angular resolution near-IR observations certainly suggest a very inhomogeneous medium close to the star, where even proper motions of clumps have been measured \\citep{weigetal02}. Far-infrared and radio observations of the circumstellar medium normally do not have the resolution to identify any small-scale structure. However, observations of in particular SiO and H$_2$O maser emission suggest the presence of clumps also in the CSEs. The very long baseline interferometry (VLBI) observations of SiO masers towards AGB stars suggest spot sizes as small as a few 10$^{12}$\\,cm in regions close to the star \\citep{coloetal92}. Indeed, for reasonable mass-loss rates the gas within $10^{14}$\\,cm of the star must be substantially clumped to provide the densities required for the SiO masers (about 10$^{13}$\\,cm$^{-3}$). Also H$_2$O masers show spot sizes as small as a few 10$^{12}$\\,cm \\citep{imaietal97}. The characteristic OH maser spot size appears to be a few $10^{14}$\\,cm at a radius of a few $10^{15}$\\,cm, i.e., further out in the envelope \\citep{chapetal94}. Note though that maser observations are difficult to interpret in terms of a detailed density structure, and they tend to enhance any density contrasts present. Larger-scale patchiness in CO radio line brightness maps of detached shells have been interpreted as a consequence of a large number of small clumps \\citep{bergetal93, olofetal96, olofetal00}. On the other hand, high resolution optical observations of planetary nebulae (PNe) show conclusive evidence of small-scale structure, in particular, the cometary globules seen in the Helix nebula. These are about 1$\\arcsec$ in size and are clearly affected by the strong radiation field from the central white dwarf \\citep{meabetal96, meabetal98, odelhand96}. \\citet{huggetal92, huggetal02} detected CO radio line emission from these clumps and estimate their masses to be about 10$^{-5}$\\,$M_{\\odot}$. It is likely that similar structures are present also in other PNe \\citep{wils50, zans55, specetal03}; the Helix happens to be one of the most nearby PNe. \\citet{huggmaur02} studied the small-scale structure in two well-known CSEs, those around the carbon star IRC+10216 and the young PN NGC 7027, with the aim to identify at which point the small-scale structure arises. They found no evidence of clumpiness in any of the two envelopes. We note in passing that clumps have also been observed in expanding shells around stars in other circumstances. Examples are O associations \\citep{frie54}, HII regions \\citep{reipetal97}, and in the environments of young stellar objects \\citep{rudowelc88, ketoho89}. The combination of optically thin scattering due to dust and the narrow widths of the detached shells is excellent for studying the small-scale structure. A potential drawback is that the mass-dominant component, the gas, may not be distributed in the same way as the dust \\citep[see e.g.,][]{maeretal10}. The R~Scl sharpened image lends itself to a quantitative analysis of the clump properties. It suggests clump sizes of about 0$\\farcs$9 in diameter, corresponding to linear clump diameters of 4$\\times$10$^{15}$\\,cm. There is an estimated total of about 2000 clumps in the shell. The average clump dust mass is consequently about 2$\\times$10$^{-9}$\\,$M_{\\odot}$. The estimated dust-to-gas ratio suggests a clump gas mass of about 2$\\times$10$^{-6}$\\,$M_{\\odot}$. This is somewhat lower than the mass estimates of the Helix nebula cometary globules \\citep{huggetal02}, but the uncertainties in both estimates are substantial. A likely explanation for the creation of detached shells around carbon stars is an increased mass-loss during the He-shell flash, with a subsequent interaction of a faster wind running into a previous, slower wind \\citep{stefscho00, schoetal05b, mattetal07}. The question arises whether the clumpy structure is already present in the ejected gas [as argued e.g. by \\citet{dysoetal89}], or whether it emerges in the expanding gas as a consequence of instabilities [as argued e.g. by \\citet{capr73} for clumps in PNe]. It is difficult to judge whether the clumps reflect inhomogeneities already present early, e.g. in the wind-accelerating zone at a distance of a few stellar radii from the star, or even in the stellar photosphere or below, as convective cells. We can identify two problems in this context. The clumps ejected from the star in this scenario are likely to have a distribution of properties (mass, size, etc.). If so, one would expect clumps with different properties to reach somewhat different terminal velocities, and once they start to interact with the previous slower wind (which is presumably also clumpy), one would expect clumps of different sizes or masses to accrete different amounts of material from the slow wind, and thus aquire different speeds for the second time. Thus it is hard to explain why the clumps seem to be located at equal distances from the star in a thin shell. On the other hand, as a result of the faster wind colliding with a slower wind, instabilities of various types may occur, e.g., Rayleigh-Taylor, Kelvin-Helmholtz, thermal, and gravitational instabilities. It is beyond the scope of this paper to analyse this in detail, but we note that numerical models indicate that structures roughly corresponding to the observed sizes and densities may indeed occur due to Rayleigh-Taylor instabilities \\citep{myasetal00}. Once the clump characteristics are determined and their formation mechanism understood, the clump properties can possibly be used as a diagnostic of e.g. stellar evolution." }, "1003/1003.2313_arXiv.txt": { "abstract": "The Higgs portal of the Standard Model provides the opportunity for coupling to a very light scalar field $\\phi$ via the super-renormalizable operator $\\phi(H^\\dagger H)$. This allows for the existence of a very light scalar dark matter that has coherent interaction with the Standard Model particles and yet has its mass protected against radiative corrections. We analyze ensuing constraints from the fifth-force measurements, along with the cosmological requirements. We find that the detectable level of the fifth-force can be achieved in models with low inflationary scales, and certain amount of fine-tuning in the initial deviation of $\\phi$ from its minimum. ", "introduction": "About 95\\% of the energy budget of the Universe consists of \"dark\" -- and unknown -- components. This is a strong motivation for considering and studying hidden sectors beyond the Standard Model (SM). Gravitational effects of dark matter cannot reveal the mass of its constitutents, and indeed a wide variety of mass ranges, from the inverse galactic size to the super-Planckian scales, is conceivable. While many models that possess stable particles with masses comparable to the SM energy scales have been a subject of incessant theoretical and experimental activity, models with light sub-eV mass scale dark matter received far less attention. Below the eV mass scale the dark matter would have to be of integer spin, and be produced non-thermally. The only chance of detecting such dark matter non-gravitationally would occur if such particles are converted into electromagnetic radiation in the external fields or they modify the interaction stength of SM particles. But if light dark matter interacts with the SM, then immediately its lightness comes to question as the quantum loops with SM particle may easily destabilize the mass scale. A prominent particle in this category is the QCD axion \\cite{QCD} that interacts with the SM currents derivatively, $j_\\mu\\partial_\\mu a $, and has its tiny mass generated by the non-perturbative QCD effects protected at any loop level. Because of the pseudoscalar nature of $a$ and its derivative couplings, it does not generate a long-range attractive force. A very natural question to ask is whether SM allows for couplings to other types of sub-eV dark matter fields that lead to additional observable effects. For a recent review of the light sector phenomenology see, {\\em e.g.} \\cite{Ringwald}. Real scalar field $\\phi$ and the vector field $V_\\mu$ provide such opportunities with their couplings to the SM fields via the so-called Higgs and vector portals: \\begin{eqnarray} \\label{portals} (A \\phi + \\lambda \\phi^2) H^\\dagger H & & \\qquad {\\rm Higgs~portal} \\\\ \\nonumber J_\\mu V_\\mu;~~ \\partial_\\mu J_\\mu = 0 & & \\qquad {\\rm Vector~portal}, \\end{eqnarray} where $H$ is the Higgs doublet, $A$ and $\\lambda$ are parameters and $J_\\mu$ is some locally conserved SM current, such as hypercharge of baryon current. If there is some initial value for $\\phi$ or $V_\\mu$ fields with respect to their zero energy configurations, one can source part/all of the Universe's energy density from the coherent oscillations around the minimum. The perils of low mass scale stabilization are immediately apparent in Eq. (\\ref{portals}). Indeed, any loops of the SM fields would tend to induce the correction to the mass of $\\phi$ field $\\sim \\lambda \\Lambda_{UV}^2$, where $ \\Lambda_{UV}$ is the highest energy scale in the problem serving as the ultra-violet cutoff. Therefore, $\\lambda$ should be taken to incredibly small values, making this portal irrelevant for the phenomenology of sub-eV dark matter. In contrast, the vector portals and the super-renormalizable Higgs portal, $A\\phi H^\\dagger H$, allow to avoid problems with technical naturallness. In the latter case loop corrections scale only as $A^2 \\log\\Lambda_{UV}$, while the quadratic divergences affect only the term linear in $\\phi$, which can typically be absorbed in an overall field shift. In this paper we examine generic consequences of this coupling for the sub-eV scalar dark matter, leaving vector dark matter to future studies. ", "conclusions": "The model we considered in this work is very similar to the linearized version of the Brans-Dicke (BD) theory when the scalar field is supplied with the mass term. Indeed, the transformation from the Jordan to the Einstein frame puts the BD scalar in front of any dimensionful parameter. Therefore, the $A$ parameter from the model considered here can be identified with $A \\sim m_h^2 / (\\omega^{1/2} M_{P})$, where $\\omega$ is the BD parameter. It is very important to keep in mind, however, one crucial difference. In the BD theory, the $\\phi$-field also couples to all massive states that may exist beyond the SM states, and therefore, even at the electroweak scale one should expect the extension of Eq. (\\ref{potential}) by additional higher-dimensional operators. Such terms alter the couplings of BD scalar to matter, and make couplings to gauge bosons, {\\em e.g.} $g_{\\phi\\gamma\\gamma}$, different from the values in the model considered here. Moreover, the BD theory requires explicit UV completion, while the model with coupling via the super-renormalizable portal assumes that higher-dimensional operators are absent from the beginning and generated only via the SM loops with the $\\varphi$-independent UV cutoff. The key feature of the model considered here is its technical naturalness. It allows to have a relatively light scalar dark matter that generate medium-range attractive force without extra fine tuning of the parameters in the Lagrangian. A detectable level of the fifth force, would have to be combined with inflationary scenarios with low $r$ and face with the potential fine-tuning problem in initial condition of the scalar field value. One of the most interesting (albeit fine-tuned) scenarios that can have particle physics implications not considered in this paper is the $\\phi$-dependence of the electroweak phase transition. If $m_\\varphi$ is taken comparable to the Hubble rate at $T=100$ GeV, $A>m_\\varphi$ can lead to $|A\\phi_*| \\sim 10^4$ GeV$^2$, thus altering the properties of the electorweak sector close to the phase tansition point. This way, one could change the order of the phase transition, and make it first order if the effective Higgs mass is pushed below 50 GeV. The model considered here falls into the class of the \"super-cool\" dark matter models, such as axion dark matter. Another example, worth of investigation is the vector dark matter. There, the coupling of vector fields to the SM and the mass of the vector fields do not have to follow the strength$\\times$range=const constraint of the scalar case. This could open more room for the fifth-force mediated by the vector-like sub-eV dark matter. We would like to thank N. Afshordi, N. Barnaby, J. Bond, C. Burgess, A. Erickcek and A. Nicolis for useful discussions. Research at the Perimeter Institute supported in part by the Government of Canada through NSERC and by the Province of Ontario through MEDT." }, "1003/1003.0911_arXiv.txt": { "abstract": "In the last decade or so, there have been numerous searches for hot subdwarfs in close binaries. There has been little to no attention paid to wide binaries however. The advantages of understanding these systems can be many. The stars can be assumed to be coeval, which means they have common properties. The distance and metallicity, for example, are both unknown for the subdwarf component, but may be determinable for the secondary, allowing other properties of the subdwarf to be estimated. With this in mind, we have started a search for common proper motion pairs containing a hot subdwarf component. We have uncovered several promising candidate systems, which are presented here. ", "introduction": "The mass of a star is of fundamental importance to our understanding of stellar evolution. While there are several robust methods for determining the masses of stars on the Main Sequence, for more evolved non-degenerate stars the situation is much more difficult and many assumptions must be made. Stars at the extreme blue end of the horizontal branch - the hot subdwarfs - fall into this category. Hot subdwarfs (sdB, sdO) are core helium-burning stars that are found at the blue end of the horizontal branch. They are found in all Galactic stellar populations and are sufficiently numerous to account for the UV-upturn of early-type galaxies. While it is generally accepted they will evolve directly into white dwarfs and not return to the giant phase, exactly how the stars arrived at this point in the HR diagram is still an open question. There are several possible formation channels: binary evolution through mass transfer and common envelope ejection \\citep[e.g.][]{Han2003}; delayed core helium flashes in post-Red Giant Branch (RGB) evolution, \\citep[e.g.][]{Lanz2004}; the merger of two helium-core white dwarfs \\citep[e.g. ][]{SJ2000}; and non-core helium burning post-RGB evolution \\citep[e.g.][]{CC1993}. The first three channels give rise to a mass distribution peaked near 0.48\\,$M_\\odot$, while the latter channel produces helium-core white dwarfs with masses closer to $\\sim$0.3\\,$M_\\odot$. There are only a handful of useful parallax measurements of hot subdwarfs because they were -- in general -- too faint to be observed with Hipparcos \\citep[see][]{Heber2002}. This in turn means they have very few reliable mass determinations. \\citet{Heber1992} noted \\emph{``Hot subluminous stars in binary systems could provide an important tool for checking the evolutionary scenarios ... since they possibly allow stellar masses to be determined. Visual binaries and eclipsing spectroscopic binaries are of utmost importance in this respect.''} Indeed, about half of the sdBs reside in close binaries with white dwarf or late-type Main Sequence companions. There are several eclipsing spectroscopic binaries containing the latter combination, however these all also show a reflection effect and therefore rely on our incomplete knowledge of the impacts of stellar irradiation \\citep[for example, see][]{Drechsel2001,Vuckovic2007}. There are some tremendous advantages of studying a hot subdwarf in a resolved binary system. Firstly, the stars \\emph{do not} interact, and most likely have \\emph{never} interacted. This means one does not have to worry about the physics of binary interactions. Secondly, the individual components can be studied independently, without the problems and assumptions surrounding composite spectra. Lastly, Main Sequence and Red Giant Branch stars are \\emph{well understood}, at least relative to hot subdwarfs. We have therefore searched for common proper motion pairs and visual binaries containing a hot subdwarf component. This is the focus of the present work, where we present our most promising candidates and some very preliminary results. The reader should note that there are no definite or clear conclusions in the work presented here; this is an exploratory pilot-type study, which is in part intended to inspire others to explore alternative methods to determining the fundamental parameters of these enigmatic objects. ", "conclusions": "" }, "1003/1003.0250_arXiv.txt": { "abstract": "Gravitational lensing is generally treated in the geometric optics limit; however, when the wavelength of the radiation approaches or exceeds the Schwarzschild radius of the lens, diffraction becomes important. Although the magnification generated by diffractive gravitational lensing is well understood, the astrometric signatures of diffractive microlensing are first derived in this paper along with a simple closed-form bound for the astrometric shift. This simple bound yields the maximal shifts for substellar lenses in solar neighbourhood observed at 20~GHz, accessible to high sensitivity, high angular resolution radio telescopes such as the proposed Square Kilometre Array (SKA). ", "introduction": "Gravitational microlensing is a powerful tool to probe the constituents of the solar neighbourhood, the Galaxy and beyond \\citep[e.g.][]{wambsganss06:_gravit_lensin}. In particular \\citet{2005ApJ...635..711G} have propose astrometric microlensing as a technique to detect sub-stellar objects in the solar neighbourhood, and \\citet{2009arXiv0910.3922H,2010arXiv1002.3007H} argued that diffraction could provide important constraints on lensing objects in the Kuiper belt and beyond. The combination of diffraction and astrometric lensing offers a new dimension to microlensing surveys. Several authors have examined gravitational lensing including the effects of diffraction \\citep[e.g.][]{1978RaF....21...87I,1980Ap&SS..71..171E,1981Ap&SS..78..199B,1986ApJ...307...30D,1995ApJ...442...67U,2004A&A...423..787T}. However, the focus has almost entirely been on the magnification of the image. An exception is the work of \\citet{Labeyrie:1994p1906} that examines the possibility of using a planetary mass lens as a telescope. This letter will examine the astrometry of diffractive lensing; that is how does lensing affect the centroid of the light distribution including the effects of diffraction. As diffraction can amplify the magnification of a gravitational lens, so too does it increase the motion of the image. Measuring the motion of the image can provide constraints on the lens, source and their relative motion. The commissioning of the Square Kilometre Array (SKA) over the next decade will offer an unprecedented view of the radio sky. \\citet{2000pras.conf..213K} outlines some prospects for using the SKA to understand strongly lensed quasars and especially the small-scale structure of the lensing object. This letter also examines primarily the lensing of quasars but focuses on nearby lensing objects with the hopes to provide constraints on the number of small bodies in the solar neighbourhood. Such constraints are difficult to obtain otherwise. The letter is divided into a calculation (\\S~\\ref{sec:calculations}) of the astrometric signature of lensing both in the diffractive and geometric optics regimes, a description of the results (\\S~\\ref{sec:results}) and an evaluation of the prospects of observing this effect (\\S~\\ref{sec:conclusions}). ", "conclusions": "\\label{sec:conclusions} The continuous monitoring of compact, distant radio sources may provide new way to probe the constituents of our solar neighbourhood, in particular freely, floating sub-stellar objects. The astrometric signatures of diffractive microlensing can provide an estimate of the mass, distance and proper motion of the lensing object, possibly allowing follow-up observations of the lens itself. Astrometric lensing even without diffraction effects can provide this information as well \\citep{wambsganss06:_gravit_lensin}; however, diffraction typically amplifies the astrometric signature and radio observations often offer much higher angular resolution on the order of ten milliarcseconds versus several hundred milliarcseconds in the optical. This letter has used the specifications of the SKA as a benchmark. Clearly the high angular resolution and high frequency offered by the SKA are helpful for the detection of astrometric lensing in the radio; however, the high sensitivity of the SKA may not strictly be necessary if one focuses on bright radio sources. Perhaps, a purpose-built very-large baseline array of phased dipoles could achieve the needed angular resolution (and possibly even a finer resolution than the SKA) with a sufficient sensitivity to continuously determine the centroids the brightest radio sources to the needed accuracy to detect low-mass objects in the solar neighbourhood. Furthermore, such a monitoring campaign could yield new insights on quasar physics as well as other ancillary results. The low expected optical depth for these events of about $2\\times 10^{-9}$ would required the monitoring of 100,000 radio sources to achieve even the modest event rate of once per decade. These sources could be quasars or bulge giants, although the effect should be more pronounced with the high brightness-temperture quasars." }, "1003/1003.2580_arXiv.txt": { "abstract": "V-type asteroids in the inner Main Belt ($a<2.5$ AU) and the HED meteorites are thought to be genetically related to one another as collisional fragments from the surface of the large basaltic asteroid 4 Vesta. We investigate this relationship by comparing the near-infrared ($0.7-2.5~\\mu m$) spectra of 39 V-type asteroids to laboratory spectra of HED meteorites. The central wavelengths and areas spanned by the 1 and 2 $\\mu m$ pyroxene-olivine absorption bands that are characteristic of planetary basalts are measured for both the asteroidal and meteoritic data. The band centers are shown to be well correlated, however the ratio of areas spanned by the 1 and 2 $\\mu m$ absorption bands are much larger for the asteroids than for the meteorites. We argue that this offset in band area ratio is consistent with our currently limited understanding of the effects of space weathering, however we can not rule out the possibility that this offset is due to compositional differences. Several other possible causes of this offset are discussed. Amongst these inner Main Belt asteroids we do not find evidence for non-Vestoid mineralogies. Instead, these asteroids seem to represent a continuum of compositions, consistent with an origin from a single differentiated parent body. In addition, our analysis shows that V-type asteroids with low inclinations ($i<6^\\circ$) tend to have band centers slightly shifted towards long wavelengths. This may imply that more than one collision on Vesta's surface was responsible for producing the observed population of inner belt V-type asteroids. Finally, we offer several predictions that can be tested when the Dawn spacecraft enters into orbit around Vesta in the summer of 2011. ", "introduction": "} The basaltic howardite, eucrite and diogenite meteorites (HEDs) and the large Main Belt asteroid 4 Vesta have traditionally been linked due to their spectroscopic similarity and the lack of any other large asteroid with the characteristic spectral signature of magmatic basalts \\citep[][]{1970Sci...168.1445M,1977GeCoA..41.1271C}. The presence of a collisional family associated with Vesta supports this link. Vesta-family members, often referred to as the Vestoids, are spectroscopically classified as V-types and are dynamically linked to Vesta. The term non-Vestoid refers to any V-type asteroid that originated on a parent body other than Vesta. The Vestoids extend from the $\\nu_6$ secular resonance at the inner edge of the Main Belt to the 3:1 mean motion resonance with Jupiter at 2.5 AU \\citep[Fig. \\ref{fig.inner},][]{1993Sci...260..186B}. These resonances act as a dynamical escape hatch from the Main Belt and can transport fragments removed from the surface of Vesta (or one of the Vestoids) to the Earth as HED meteorites \\citep{1997Sci...277..197G}. Resolved images of Vesta reveal a large crater ($\\sim460$ km in diameter) on its south pole \\citep{1997Sci...277.1492T}, supporting a scenario of collisional formation for the Vesta family. Hydrocode simulations of the collision that formed the Vesta family \\citep{1997M&PS...32..965A} suggest that km-size fragments would have been removed with ejection velocities ($\\Delta v$) of no greater than approximately 0.6 km/s. Simplified versions of Gauss's equations can be used to quantify the distance from Vesta in orbital element space corresponding to this ejection velocity \\citep{1996Icar..124..156Z}.\\footnote{The following calculations assume that the true anomaly and argument of perihelion at the time of formation of the Vesta family were equal to those calculated by \\citet{1996Icar..124..156Z}. However, these authors did not consider the effects of orbital migration due to the Yarkovsky force \\citep{2006AREPS..34..157B}. Therefore, the calculations presented here are approximations accurate to $\\sim20\\%$.} The maximum range of semi-major axes for collisionally-produced fragments can be estimated by assuming that a hypothetical Vestoid was ejected with a 0.6 km/s velocity vector aligned exclusively in a direction tangential to its orbit \\citep[Equation 4 in][]{1996Icar..124..156Z}. This calculated range is $2.23-2.49$ AU, centered on the semi-major axis of Vesta (2.36 AU). Similar calculations can be done for both eccentricity and inclination, producing ranges of $0.06-0.13$ and $4.8-7.9^\\circ$ respectively. This region of orbital element space is enclosed by the ellipse in Figure \\ref{fig.inner}. Figure \\ref{fig.inner} shows numerous V-type asteroids in the inner Main Belt with values of $\\Delta v$ much larger than 0.6 km/s, some with values in excess of 2 km/s. When these objects were first discovered \\citep[e.g.][]{1993Sci...260..186B,2001M&PS...36..761B,2002Icar..159..178F,2004Icar..172..179L,2006A&A...459..969A} it was unclear how they could have reached such orbits. However, recent progress in the use of numerical integrators has helped to clarify this issue. \\citet{2005A&A...441..819C} showed that three-body and weak secular resonances could lead to the migration of some Vestoids to orbits with $\\Delta v>0.6$ km/s. \\citet{2008Icar..193...85N} showed that a combination of these resonances and the Yarkovsky effect could disperse the orbits of Vestoids to nearly the full extent of the inner Main Belt. However, these authors found that the observed number of V-type asteroids at low inclination ($i<6^\\circ$) was too large to be explained by their model of Vestoid migration. Three possibilities exist to explain this over-abundance of low-$i$ V-types. First, these objects could be fragments of basaltic crust from a non-Vestoid differentiated parent body. In this case these objects could be spectroscopically distinct from the Vestoids, as is the case for non-Vestoid V-types in the outer Main Belt \\citep[e.g. ][]{2000Sci...288.2033L,2008ApJ...682L..57M}. Second, they may be from Vesta, but were removed from the surface before the Late Heavy Bombardment (LHB), before the primary family forming collision, and were scattered to their current orbits as mean motion and secular resonances swept through the Main Belt during the LHB \\citep{1997AJ....114..396G}. In this case these objects would represent an older population of Vestoids, removed from a different region on Vesta's surface and thus might be spectroscopically distinct. Third, these objects may have been ejected from the Vesta parent body at the time of family formation and have since migrated to their current orbits by some unexplored dynamical mechanism. In this case these objects should not appear spectroscopically different from other V-type asteroids in the inner Main Belt. A similar line of reasoning motivated \\citet{1998AMR....11..163H} to investigate the visible-wavelength spectral features (namely the slope and 1 $\\mu m$ band depth) of 20 V-type asteroids in the inner Main Belt as a function of their orbital elements. These authors found that V-type asteroids with large values of $\\Delta v$ tended to have steeper spectral slopes than V-types with smaller $\\Delta v$. However, the largest value of $\\Delta v$ considered by these authors was 0.65 km/s, very close to the expected ejection velocity of Vestoid fragments. Furthermore, all of the objects that were studied have since been incorporated into the Vesta dynamical family as detection completeness has increased in the last decade. Thus it is surprising that this spectroscopic trend was observed as a function of orbital parameters for objects that plausibly originated at the same time from the same parent body. \\citet{2004Icar..171..120D} attempted a similar investigation into the spectroscopic diversity of Vestoids at NIR wavelengths. This study produced unexpected results: the Band II to Band I area ratios and the Band I and Band II centers for the majority of the V-types in their sample did not agree with those of the HEDs (Band I and II refer to the 1 and 2 $\\mu m$ absorption bands common to basaltic material, see \\S\\ref{sec.tools} for definitions of these parameters). Although it has been suggested that band area ratios are sensitive to variations in grain size, temperature and space weathering \\citep{2002LPI....33.2023U}, band centers should be less sensitive to these effects and thus comparable between genetically related populations (i.e. the HEDs and Vestoids). In light of recent advances in dynamical simulations \\citep{2005A&A...441..819C,2008Icar..193...85N} and improvements in NIR spectroscopic instrumentation \\citep{2003PASP..115..362R}, we revisit the issue of the diversity of basaltic asteroids by measuring the NIR spectral properties of 39 inner Main Belt V-type asteroids. The goals of this study are threefold: (1) address the reported spectro-dynamical correlation amongst V-type asteroids in the inner Main Belt \\citep{1998AMR....11..163H} by extending our analysis out to NIR wavelengths and by including V-types across a wider range of orbital element space; (2) address the findings of \\citet{2004Icar..171..120D} to look for spectroscopic differences between inner belt V-type asteroids and HED meteorites; (3) determine if any of the V-type asteroids in the inner Main Belt have spectroscopic properties suggestive of a non-Vestoid mineralogy. It is important to note that we do not attempt to extract detailed mineralogical information for individual asteroids. Instead we characterize our relatively large data set using band analysis techniques \\citep[e.g.][]{1986JGR....9111641C} with the intent of making statistically significant statements about the gross spectral properties of V-type asteroids relative to those of the HED meteorites. ", "conclusions": "We have observed and analyzed the spectra of 39 V-type asteroids. Comparison of their band parameters to those of HED meteorites from the RELAB database reveals a close correlation between band centers. We do not find the wide range of band centers that was reported by \\citet{2004Icar..171..120D}. We suspect that this difference is due to the lower S/N of the \\citet{2004Icar..171..120D} data set. We were able to confirm an offset in BAR between the HEDs and V-type asteroids \\citep{2005M&PS...40..445D} and argue that this offset is consistent with our initial understanding of space weathering effects on Vesta-like mineralogies However, further work is necessary to understand whether a combination of grain size, composition and/or weathering could be responsible for this offset. Several other possible causes were discussed and found to be unlikely. We were unable to reproduce at NIR wavelengths the spectro-dynamical association found by \\citet{1998AMR....11..163H}, namely inner belt V-types do not show any correlation between slope across their 1 $\\mu m$ bands and distance from Vesta in orbital element space. We note that a search for correlation between other spectroscopic characteristics (band centers, slopes, depths and areas) and dynamical properties (semi-major axis, inclination, eccentricity, $\\Delta v$) did not reveal any significant results. We did not find any new evidence to suggest the presence of V-type asteroids with non-Vesta mineralogies [e.g.~1459 Magnya and 21238 (1995 WV7)]. Instead, the band parameters of these objects seem to represent a continuum of compositions that are consistent with an origin from a single parent body, most likely 4 Vesta. Only asteroid 2579 Spartacus is found to have a band area ratio that is significantly offset from the general trend represented by the other targets. This has been noted before \\citep{2001M&PS...36..761B} and could be due to its origin on another parent body, however it could also be a large fragment that originated from greater depths (relative to the other targets) within the Vesta parent body. The lack of additional spectroscopic outliers amongst the V-type asteroids in the inner Main Belt implies that they are of a common origin. However, this does not preclude the possibility that inner belt V-types include basaltic crustal fragments from multiple differentiated parent bodies that are indistinguishable with band analysis techniques. We have reported that V-type asteroids with low inclinations ($i<6^\\circ$) in the inner Main Belt tend to have band centers shifted to longer wavelengths. This is compelling in light of the dynamical results of \\citet{2008Icar..193...85N}, however additional data should be obtained to confirm or refute the significance of this finding. In particular, a focused study on low-inclination V-type asteroids may provide further insight on whether these objects are spectroscopically distinct. This study has resulted in several predictions that can be tested by the Dawn spacecraft when it enters into orbit around Vesta in the summer of 2011. First, we have parameterized the dependence between temperature and BAR (Equations \\ref{eqn.temp} and \\ref{eqn.heliofit}). Spectroscopic observations with VIR-MS of regions at various temperatures on Vesta's surface can be used check this dependence. We have also suggested that fresher, less weathered surfaces on Vesta (e.g. impact craters) should have smaller BARs than the surrounding terrain. And finally, if the low inclination V-type asteroids in our study are predominantly eucritic in composition and were removed from one of the minor impact craters in Vesta's northern hemisphere \\citep{1997Sci...277.1492T}, then we expect that these craters will be devoid of significant quantities of diogenite-like material." }, "1003/1003.5475_arXiv.txt": { "abstract": "{To understand the formation and evolution of galaxies, it is important to have a full comprehension of the role played by the metallicity, star formation rate (SFR), morphology, and color. The interplay of these parameters at different redshifts will substantially affect the evolution of galaxies and, as a consequence, the evolution of them will provide important clues and constraints on the galaxy evolution models. In this work we focus on the evolution of the SFR, metallicity of the gas, and morphology of galaxies at low redshift in search of signs of evolution.} % {To analyze the S2N2 (log({H$\\alpha$}/[{S\\,\\textsc{ii}}]) vs. log({H$\\alpha$}/[{N\\,\\textsc{ii}}])) diagram as a possible segregator of star--forming, composite, and AGN galaxies, to study the evolution of the Baldwin, Phillips $\\&$ Terlevich (1981) diagrams, as well as the evolution of the SFR, metallicity, and morphology, through the mass--metallicity, luminosity--metallicity, SFR--stellar mass, and SFR--metallicity relationships of star--forming galaxies from SDSS--DR5 (Sloan Digital Sky Survey--Data Release 5), using redshift intervals in bins of 0.1 from $\\sim$0 to 0.4.} {We used data processed with the STARLIGHT spectral synthesis code, correcting the fluxes for dust extinction, and estimating metallicities using the $R_{23}$ method. We use the S2N2 diagnostic diagram as a tool to classify star--forming, composite, and AGN galaxies. We analyzed the evolution of the three principal BPT diagrams, estimating the SFR and specific SFR (SSFR) for our samples of galaxies, studying the luminosity and mass-metallicity relations, and analyzing the morphology of our sample of galaxies through the $g-r$ color, concentration index, and SSFR.} {We found that the S2N2 is a reliable diagram to classify star--forming, composite, and AGNs galaxies. We demonstrate that the three principal BPT diagrams show an evolution toward higher values of [{O\\,\\textsc{iii}}] $\\lambda$5007/{H$\\beta$} due to a metallicity decrement. We found an evolution in the mass--metallicity relation of $\\sim$ 0.2 dex for the redshift range $0.3 < z < 0.4$ compared to our local one. From the analysis of the evolution of the SFR and SSFR as a function of the stellar mass and metallicity, we discovered a group of galaxies with higher SFR and SSFR at all redshift samples, whose morphology is consistent with those of late--type galaxies. Finally, the comparison of our local ($0.04$ -0.2. Following with the objective of segregate SF from composite and AGNs galaxies, in this work we study the S2N2 diagram as a reliable segregator of galaxies. This log({H$\\alpha$}/[{S\\,\\textsc{ii}}]) vs. log({H$\\alpha$}/[{N\\,\\textsc{ii}}]) diagram was introduced by Sabbadin et al. (1977) as a useful tool to separate galactic planetary nebula (PNe), {H\\,\\textsc{ii}} regions, and supernova remnants (SNRs). This diagram was later applied to Herbig-Haro objects (Cant\\'o 1981), Galactic PNe (Garc\\'{\\i}a-Lario et al. 1991, Riesgo $\\&$ L\\'opez 2005), and extragalactic PNe (Magrini et al. 2003). The S2N2 diagram has been used also as a metallicity and ionization parameter indicator for extragalactic {H\\,\\textsc{ii}} regions by Viironen et al. (2007). The S2N2 diagram has been also applied to galaxies by some authors. For example, Moustakas $\\&$ Kennicutt (2006) studied whether there was a difference between integrated spectra of galaxies and the spectra of individual {H\\,\\textsc{ii}} regions. Dopita et al. (2006) used the S2N2 diagram, among others, for abundance diagnostics using photoionization models. Nevertheless, the [{S\\,\\textsc{ii}}] flux shows always deficiences when generated by photoionization models (e.g. Levesque et al. 2010). Also, Lamareille et al. (2009) and P\\'erez-Montero et al. (2009) used the S2N2 diagram as a segregator of SF from Seyfert 2 galaxies, but using different ratios: log([{N\\,\\textsc{ii}}]/{H$\\alpha$}) vs log([{S\\,\\textsc{ii}}]/{H$\\alpha$}). However, in their division Lamareille et al. (2009) do not distinguish between SF and composite galaxies, also, they used equivalent widths instead of emission line fluxes, which could affect the results (Kobulnicky $\\&$ Kewley 2004). The formation and evolution of galaxies at different cosmological epochs are driven mainly by two linked processes: the star formation history and the metal enrichment. Thus, from an observational point of view, the star formation rate (SFR), the metallicity and the stellar mass of the galaxies at different epochs will give us important clues on the evolution of galaxies. The first quantitative SFRs were derived from evolutionary synthesis models of galaxy colors (Tinsley 1968, 1972, Searle et al. 1973), confirming the trends in SFRs and star formation histories along the Hubble sequence, and giving the first predictions of the evolution of the SFR with cosmic lookback time. The development of more precise direct SFR diagnostics includes the integrated emission--line fluxes (Cohen 1976, Kennicutt 1983), near-ultraviolet continuum fluxes (Donas $\\&$ Deharveng 1984), and infrared continuum fluxes (Harper $\\&$ Low 1973, Rieke $\\&$ Lebofsky 1978, Telesco $\\&$ Harper 1980); see Kennicutt (1998) for a review. The hydrogen Balmer line {H$\\alpha$} is currently the most reliable tracer of star formation, since in {H\\,\\textsc{ii}} regions and star-forming galaxies, the Balmer emission-line luminosity scales directly with the total ionizing flux of the embedded stars. A widely known calibration of the {H$\\alpha$} line as SFR tracer is the one devised by Kennicutt (1998). However, it is important to take into account corrections for stellar absorption and reddening to obtain SFRs in agreement with the ones derived using other wavelengths (e.g. Rosa-Gonz\\'alez et al. 2002, Charlot et al. 2002, Dopita et al. 2002). In parallel, other diagnostics have been developed using the oxygen doublet [{O\\,\\textsc{ii}}] $\\lambda$3726, 3729 for the redshift range $z \\sim 0.4-1.5$ (e.g. Gallagher et al. 1989, Kennicutt 1998, Rosa-Gonz\\'alez et al. 2002, Kewley et al. 2004). Moreover, this diagnostic is usefull when the {H$\\alpha$} line is not easily observable at higher redshifts ($z \\gtrsim 0.4$ in the optical). However, the [{O\\,\\textsc{ii}}] doublet presents problems in reddening and abundance dependence (Jansen, Franx $\\&$ Fabricant 2001, Charlot et al. 2002). Alternatively, it is possible to estimate the SFR from the soft X-ray luminosity, which is comparable to that determined from the {H$\\alpha$} luminosity (Rosa Gonz\\'alez et al. 2009, Rovilos et al. 2009). A strong dependence of the SFR and the stellar mass and its evolution with redshift has been found, with the bulk of star formation occurring first in massive galaxies, and later in less massive systems (e.g. Guzm\\'an et al. 1997, Brinchmann $\\&$ Ellis 2000, Juneau et al. 2005, Bauer et al. 2005, Bell et al. 2005, P\\'erez-Gonzalez et al. 2005, Feulner et al. 2005, Papovich et al. 2006, Caputi et al. 2006, Reddy et al 2006, Erb et al. 2006, Noeske et al. 2007a, Buat et al. 2008). In the local universe, several studies have illustrated a relationship between the SFR and stellar mass, identifying two populations: galaxies on a star-forming sequence, and $``$quenched\" galaxies, with little or no detectable star formation (Brinchmann et al. 2004, Salim et al. 2005, Lee 2006). At higher redshift, Noeske et al. (2007a) showed the existence of a $``$main sequence\" (MS) for SF galaxies in the SFR--stellar mass relation over the redshift range $0.2 < z < 1.1$. From the galaxies considered in this study, it was shown that the slope of the MS remains constant to $z>1$, while the MS as a whole moves to higher SFR as $z$ increases. Metallicity is another important property of galaxies, and its study is crucial for a deep understanding of galaxy formation and evolution, since it is related to the whole past history of the galaxy. Metallicity is a tracer of the fraction of baryonic mass already converted into stars and is sensitive to the metal losses due to stellar winds, supernovae and active nuclei feedbacks. A detailed description of the different metallicity methods and calibrations are given in Lara-L\\'opez et al. (2009a,b). Stellar mass and metallicity are strongly correlated in SF galaxies, with massive galaxies showing higher metallicities than less massive galaxies. This relationship provides crucial insight into galaxy formation and evolution. The mass-metallicity ($M-Z$) relation was first observed by Lequeux et al. (1979), has been intensively studied (Skillman et al. 1989; Brodie $\\&$ Huchra 1991; Zaritsky et al. 1994; Richer $\\&$ McCall 1995; Garnett et al. 1997; Pilyugin $\\&$ Ferrini 2000, among others), and it is well established by the work of Tremonti et al. (2004, hereafter T04) for the local universe (z $\\sim$ 0.1) using SDSS data. The study of the redshift evolution of the $M-Z$ relation has provided us with crucial information on the cosmic evolution of star formation. Regarding the evolution of the $M-Z$ relation for SF galaxies at z $<$ 1, Savaglio et al. (2005), have investigated the mass--metallicity relations using galaxies at 0.4 $<$ z $<$ 1, finding that metallicity is lower at higher redshift by $\\sim$ 0.15 dex. Moreover, Maier at al. (2005), Hammer et al. (2005), and Liang et al. (2006) found that emission line galaxies were poorer in metals at z $\\sim$ 0.7 than present--day spirals. A study of Lamareille et al. (2009) focused on the evolution of the $M-Z$ relation up to z $\\sim$ 0.9, suggesting that the $M-Z$ relation is flatter at higher redshifts. However, Carollo $\\&$ Lilly (2001), from emission--line ratios of 15 galaxies in a range of 0.5 $<$ z $<$ 1, found that their metallicities appear to be remarkably similar to those of local galaxies selected with the same criteria. Also, Lilly et al. (2003), from a sample of 66 SF galaxies with 0.47 $<$ z $<$ 0.92, claim a smaller variation in metallicity of $\\sim$ 0.08 dex compared with the metallicity observed locally, showing only modest evolutionary effects (for more details about the $M-Z$ relation, see Lara-L\\'opez et al. 2009b). In a recent study, Calura et al. (2009) have demonstrated the importance on the morphology of galaxies when deriving the $M-Z$ relation since, at any redshift, elliptical galaxies present the highest stellar masses and the highest metallicities, whereas the irregulars are the least massive galaxies, characterised by the lowest O abundances. In this paper, we consistently approach several topics, starting with the introduction of the S2N2 as a reliable diagram to classify galaxies, the analysis of the metallicity evolution of galaxies in the three BPT diagrams, and, for a better understanding of the processes involved in the observed evolution of galaxies at low redshift, we studied the mass, metallicity and SFR relations, such as the $M-Z$, metallicity-SFR and mass-SFR relations. We also point out that the morphology of galaxies play an important role when deriving conclusions, since late--type galaxies will result in lower metallicity estimates and higher SFRs than early--type (Calura et al. 2009). This paper is structured as follows, in Sect. 2 we detail the data used for this study, the dust extinction correction and the metallicity estimates for our sample of galaxies, in Sect. 3 we introduce the S2N2 as a reliable diagram to segregate SF, composite, and AGNs galaxies. In Sect. 4 we analyzed the evolution of the BPT diagrams. In Sect. 5 we investigate the evolution of the mass-metallicity and luminosity-metallicity relations, whereas in Sect. 6 we discuss the relations between the SFR and SSFR with stellar mass and metallicity, as well as the morphology of our galaxies using colors, concentration index, and SSFRs. Finally, conclusions are given in Sect. 7. ", "conclusions": "We analyzed a sample of emission line galaxies selected in four redshift intervals from $\\sim$0 to 0.4 in bins of 0.1, taking into account the magnitude completeness of every redshift interval. In this paper we introduced the S2N2 diagram as a star-forming, composite, and AGNs galaxy classificator, we estimated metallicities using the R$_{23}$ method and analyzed the evolutive effects of galaxies from the three BPT diagrams. Additionally, we studied the evolution of the $M-Z$ and $L-Z$ relations, and analyzed the evolution and implications of the galaxy morphology in the SFR--mass and metallicity relations. From these analysis we conclude the following: \\begin{itemize} \\item Using the Kew01 photoionization grids, and the Kauf03 and Kew01 SF, and starburst limit respectively, in the [{N\\,\\textsc{ii}}] /{H$\\alpha$} vs [{O\\,\\textsc{iii}}] $\\lambda$5007/{H$\\beta$} diagram, we have demonstrated that the S2N2 is a well--behaved diagnostic diagram efficiently classifying star-forming, composite, and AGNs galaxies. \\item We analyzed the galaxy evolution using the three main BPT diagrams: [{N\\,\\textsc{ii}}] /{H$\\alpha$}, [{S\\,\\textsc{ii}}] /{H$\\alpha$}, and [{O\\,\\textsc{i}}] $\\lambda$6300/{H$\\alpha$} vs [{O\\,\\textsc{iii}}] $\\lambda$5007/{H$\\beta$} in our four redshift bins, observing an evolution toward higher values of the [{O\\,\\textsc{iii}}] $\\lambda$5007/{H$\\beta$} ratio. This evolution is a consequence of the metallicity evolution as redshift increases, reflected in the three BPT diagrams, because the ratio [{O\\,\\textsc{iii}}] $\\lambda$5007/{H$\\beta$} is a good metallicity indicator. As a result, a metallicity decrement will be reflected in higher values of this ratio. \\item We analyzed the evolution of the $M-Z$ and $L-Z$ relations, observing that at higher redshift values, both relations evolve towards lower values of metallicity. We discovered that the flat zone of the $M-Z$ relation reported by Tremonti et al. (2004) for galaxies with log(M$_{star}$/M$_{\\odot}$) $\\gtrsim$ 10.5, is mainly constituted by galaxies at $z > 0.1$ (samples at $z_1$, $z_2$ and $z_3$). Galaxies at $z_0$ redshift could be fitted with a linear function. Our $M-Z$ relation at redshift $z_3$ is $\\sim$0.2 dex lower than our local one. \\item Our fit to the $M-Z$ relation for sample $z_3$ is in agreement with the one of Erb et al, (2006a) at $z\\sim2.2$. We attribute this similarity to the galaxy morphology in the different samples, since our $z_3$ sample is conformed by late--type galaxies, while the sample of Erb et al. is composed by a mix of early and late--type galaxies. According to Calura et al. (2009), the $M-Z$ relation of late--type galaxies will have systematically lower metallicities than a $M-Z$ relation conformed by a mix of early and late--type galaxies. \\item We analyzed the evolution of the mass-to-light ratio, observing lower $M/L$ ratios as redshift increase. For a small range of absolute magnitudes in the $z_3$ sample, we have a wide range of mass, making it possible to generate the $M-Z$ relation, but difficult to generate the $L-Z$ relation. \\item The decrement in metallicity observed in previous papers for galaxies at redshift $z_3$ (Lara-L\\'opez et al. 2009a,b) is also observed, even though in this study we are not restricting our galaxy luminosities as in our previous studies. \\item We estimated the SFR and SSFR for our sample of galaxies and analyzed its relation with 12+log(O/H) and log(M$_{star}$/M$_{\\odot}$), confirming the existence of a main sequence reported by Noeske et al. (2007) in the log(SFR) vs. log(M$_{star}$/M$_{\\odot}$) plot. Consistently, we found that higher SFRs and SSFRs increase with redshift. \\item We analyzed the morphology of our galaxies through the $g-r$ color, the concentration index R$_{90}$/R$_{50}$, and the SSFR, concluding that the best method to determine the morphology was combining both, a color of $g-r < 0.6$, and a log(SSFR)$>$10 for selecting late--type galaxies. \\item Our $z_3$ sample of galaxies is mainly formed by late--type galaxies, a fact that helped us to classify morphological types at lower redshift. The fact that at higher redshift the fraction of late--type galaxies is larger, was confirmed by using mock galaxy catalogues from Millennium simulations. \\item We found at the higher redshift, a population with higher SFR and SSFR than the galaxies in the $z_0$ sample. After classifying late and early--type galaxies in the $z_0$ sample, we realized that the observed $tail$ showing higher SFR and SSFR is formed by late--type galaxies, demonstrating the connection of the galaxy morphology with the SFR in a new fashion. \\end{itemize} Our work provide a useful tool for classifying galaxies with the S2N2 diagram, and demonstrating how galaxies evolve on the BPT diagrams as a consequence of metallicity evolution. We also analyzed the mass, metallicity and SFR relations, noting that galaxies in the redshift sample $z_3$ have lower values of metallicity, higher SFRs, and morphology indicators associated to late--types. In this study we pointed out the importance of the morphology of galaxies when deriving conclusions. Since a sample conformed by late--type galaxies will show lower values of metallicity than ones formed by a mix of morphological types." }, "1003/1003.3535_arXiv.txt": { "abstract": "We develop an analytical model to follow the cosmological evolution of magnetic fields in disk galaxies. Our assumption is that fields are amplified from a small seed field via magnetohydrodynamical (MHD) turbulence. We further assume that this process is fast compared to other relevant timescales, and occurs principally in the cold disk gas. We follow the turbulent energy density using the Shabala \\& Alexander (2009) galaxy formation and evolution model. Three processes are important to the turbulent energy budget: infall of cool gas onto the disk and supernova feedback increase the turbulence; while star formation removes gas and hence turbulent energy from the cold gas. Finally, we assume that field energy is continuously transferred from the incoherent random field into an ordered field by differential galactic rotation. Model predictions are compared with observations of local late type galaxies by Fitt \\& Alexander (1993) and Shabala \\etal\\/ (2008). The model reproduces observed magnetic field strengths and luminosities in low and intermediate-mass galaxies. These quantities are overpredicted in the most massive hosts, suggesting that inclusion of gas ejection by powerful AGNs is necessary in order to quench gas cooling and reconcile the predicted and observed magnetic field strengths. ", "introduction": "\\label{sec:introduction} Radio synchrotron emission of high energy electrons in the interstellar medium (ISM) indicates the presence of magnetic fields in galaxies. Rotation measures (RM) of background polarized sources indicate two varieties of field: a random field, which is not coherent on scales larger than the turbulence of the ISM; and a spiral ordered field which exhibits large-scale coherence (e.g. Stepanov \\etal\\/ 2008). For a typical galaxy these fields have strengths of a few $\\mu$G. In a galaxy such as M\\,51, the coherent magnetic field is observed to be associated with the optical spiral arms \\cite{PatrikeevEA06}. Such fields are important in star formation and the physics of cosmic rays, and could also have an effect on galaxy evolution, yet, despite their importance, questions about their origin, evolution and structure remain largely unsolved. The Square Kilometre Array (SKA) will help us answer questions such as these. The SKA will observe polarized synchrotron emission \\cite{GaenslerEA04}, and the ``All-Sky SKA Rotation Measure Survey'' will expand RM data sets by five orders of magnitude, mapping the magnetic fields of galaxies in unprecedented detail \\cite{StepanovEA08,Gaensler06}. In particular, the SKA will provide data on the evolution of galactic magnetic fields to high redshift \\cite{Gaensler06}, making a theoretical model for this process very valuable. Here we present such a model within the context of hierarchical structure formation. Our model is based on the observed properties of galactic magnetic fields. Observations and simulations suggest that the random field is generated by turbulence in the ISM, which is modeled as a single-phase magnetohydrodynamic (MHD) fluid, within which magnetic field lines are frozen (e.g. Cho \\etal\\/ 2009). Simulations have shown that in such a turbulent MHD fluid, the random magnetic field energy and turbulent fluid kinetic energy are approximately equal after a few turnover times \\cite{ChoEA09}, so by writing an equation for the turbulent energy in the ISM the random magnetic field energy can also be evaluated. The large-scale ordered field is produced by the differential rotation of the galaxy, which winds the random field into a spiral - this is the basic operation of a dynamo. A simple model for this process is postulated. The energy sources which contribute or remove turbulence from the ISM are supernovae, star formation and accretion of gas from the hot gas halo, and these parameters can be determined from well-established semi-analytic galaxy models, in which galaxies form from gas condensing at the centres of hierarchically merging haloes \\cite{WhiteRees78,LaceySilk91,CrotonEA06}. The formation of magnetic fields in galaxies is intimately linked to galaxy formation, since the same physical processes are at work in both cases. Recently, Arshakian \\etal\\/ (2009) have presented a qualitatively similar model for the evolution of magnetic fields in late-type galaxies. These authors considered three processes. At high redshift, the initial seed field of $\\sim 10^{-18}$~Gauss is amplified via the Biermann battery mechanism to generate what they refer to as the regular field. This field is rapidly (on eddy turnover timescale, $\\sim 10^7$~years for the Milky Way) amplified by virialization turbulence. The mean-field galactic dynamo mechanism then amplifies this regular seed field, with a typical $e$-folding time of $\\sim 10^8$ years, until it reaches equilibrium with turbulent dissipation. By contrast, given the short (compared to cosmological) timescales involved, in our model we simply assume instantaneous equality between the magnetic field energy and turbulent kinetic energy to estimate what we refer to as the random field. This field is then ordered by the mean-scale dynamo on much longer timescales. While our treatment of large-scale dynamo generation is simplified compared to that of Arshakian \\etal\\/, it encapsulates the relevant physics and has the advantage of being simple enough to be included in a cosmological framework. In that sense this work is highly complementary to the Arshakian \\etal\\/ (2009) study. We present here the first (to our knowledge) attempt to include magnetic fields in a self-consistent galaxy formation and evolution model. A number of galaxy properties are predicted, and we compare these with available data. This paper primarily focuses on the radio properties of local late-type galaxies. The paper is organised as follows. In Section~\\ref{sec:model} we outline the adopted galaxy formation model, and develop models for both the random and ordered magnetic fields. Section~\\ref{sec:parameterSpace} investigates the salient features of our model and constrains various parameters. Predictions of our model are compared with available observational data in Sections~\\ref{sec:BfieldVsStellarMass} and \\ref{sec:localRLF}. We summarise our findings in Section~\\ref{sec:conclusions}. Throughout the paper, we adopt a flat cosmology of $\\Omega_{\\rm M}=0.3$, $\\Omega_\\Lambda=0.7$, $h=0.7$ and $\\sigma_8 = 0.9$, consistent with the 2dFGRS \\cite{CollessEA01} and WMAP \\cite{SeljakEA05} results. ", "conclusions": "\\label{sec:conclusions} We have developed an analytical model to follow the cosmological evolution of magnetic fields in disk galaxies. This is done by assuming equality between the magnetic field energy density, and the turbulent energy density in cold disk gas. The gas cooling and star formation histories are followed using the Shabala \\& Alexander (2009) galaxy evolution model, which includes a physically-motivated prescription for AGN feedback, together with supernovae and reionization feedback. The model successfully reproduces the observed stellar mass functions and shutdown of star formation in massive galaxies out to redshifts of $1.5$. Three mechanisms alter the turbulent energy budget in the cold gas. Star formation removes turbulence associated with gas parcels that collapse to form the stars. The turbulence is increased by gravitational infall of cool gas onto the disk; and also by high-mass stars driving shocks through the ISM as they end their lives in supernovae explosions. Two types of magnetic fields are identified: the random field, which is incoherent on scales comparable with the turbulent eddy scale; and the ordered field, generated from the random field by differential rotation of the galaxy. Two local samples are used to test the models. The model reproduces magnetic field strengths and radio luminosities well across a wide range of low and intermediate-mass galaxies. However, the radio luminosities and magnetic field strenths are overpredicted significantly in high mass galaxies due to an overprediction in the amount of gas cooling. Inclusion of outward gas transport by powerful radio sources is required to reconcile the model with observations." }, "1003/1003.4315_arXiv.txt": { "abstract": "Motivated by the recent discovery of massive planets on wide orbits, we present a mechanism for the formation of such planets via disk fragmentation in the embedded phase of star formation. In this phase, the forming disk intensively accretes matter from the natal cloud core and undergoes several fragmentation episodes. However, most fragments are either destroyed or driven into the innermost regions (and probably onto the star) due to angular momentum exchange with spiral arms, leading to multiple FU-Ori-like bursts and disk expansion. Fragments that are sufficiently massive and form in the late embedded phase (when the disk conditions are less extreme) may open a gap and evolve into giant planets on typical orbits of several tens to several hundreds of AU. For this mechanism to work, the natal cloud core must have sufficient mass and angular momentum to trigger the burst mode and also form extended disks of the order of several hundreds of AU. When mass loading from the natal cloud core diminishes and the main fragmentation phase ends, such extended disks undergo a transient episode of contraction and density increase, during which they may give birth to a last and survivable set of giant planets on wide and relatively stable orbits. ", "introduction": "The likelihood of giant planet formation via direct gravitational instability of circumstellar disks around solar-type stars has been the subject of intense research in the past years. Despite much effort in this field, increasingly sophisticated numerical hydrodynamics simulations and analytical considerations continue to yield conflicting results. On one hand, some studies indicate that giant planets can form in massive disks, particularly in their outer parts where conditions for disk fragmentation are less extreme and the competing core-accretion model is less viable \\citep[e.g.,][]{Johnson03,Stamatellos07,Mayer07,Boss08,Dodson09,Nero09}. On the other hand, many studies show that gravitational fragmentation is unlikely, particularly in the inner few tens of AU due to insufficient disk cooling and strong stellar/envelope irradiation \\citep[e.g.,][]{Matzner05,Rafikov05,Boley06,Boley07,Rafikov07,Stamatellos08,Cai08} In spite of a great deal of sophistication, the aforementioned studies miss one important aspect---circumstellar disks are {\\it not} isolated in the early embedded phase of star formation (hereafter, EPSF). In this stage, they are subject to intense mass loading from a natal cloud core, which can significantly alter the disk's ability to fragment. A self-consistent handling of this process in numerical simulations is not easy and requires a much larger spatial scale (than just that of the disk). A spatial resolution of less than 1~AU is usually needed for planetary-mass fragments to form and {\\it survive}. Global numerical hydrodynamics simulations of the gravitational collapse and fragmentation of molecular clouds have demonstrated that forming stars are indeed surrounded by accreting gravitationally unstable disks \\citep[e.g.,][]{Bate03,Krumholz07}. Yet, these simulations resolve only massive disks, which, if fragmented, produce brown dwarfs or low-mass stellar companions rather than giant planets. Moreover, such numerical simulations are very computationally intensive and are unable to explore a wide parameter space and long evolution times. On the other hand, semi-analytic models and simplified numerical simulations of the gravitational collapse of dense cloud cores can explore a wide range of initial conditions and can give us a valuable insight into the required conditions for disk fragmentation. Using the thin-disk approximation, we were able to self-consistently follow the process of cloud core collapse and star/disk formation for at least several Myr after the formation of a central stellar object \\citep{VB05,VB06,VB09a}. These studies have shown that circumstellar disks may be gravitationally unstable and susceptible to fragmentation if the rate of gas deposition onto the disk from the cloud core is greater than that from the disk onto the star, disk viscosity is not too high, and the natal cloud cores are characterized by sufficiently large rotation rates. More sophisticated numerical hydrodynamics simulations, though with an approximate treatment of gas infall onto the disk, and semi-analytic studies have confirmed the susceptibility of non-isolated disks to fragmentation, particularly at large radii \\citep[e.g.,][]{Kratter08,Rice09,Boley09a,Boley09b,Clarke09,Rafikov09}. The feasibility of disk fragmentation and giant planet formation is only one part of the problem. The other part is the likelihood of survival of giant planets formed via disk fragmentation. Rapid radial migration due to gravitational interaction of a giant planet with a natal gas disk \\citep{Goldreich80} has traditionally been one of the stumbling blocks for the theory of giant planet formation and many mechanisms have been proposed to stop this migration in the late evolution phase \\citep[see e.g.,][]{Thommes06,Crida07,Ida08,Matsumura09}. In the early EPSF, this problem may be even more severe due to the fact that disks are more massive and profoundly non-axisymmetric. Indeed, our previous numerical studies have shown that fragments forming in the EPSF are quickly driven into the inner regions and probably onto the protostar due to exchange of angular momentum with spiral arms \\citep{VB05,VB06}. We have speculated that only those fragments that form in the late EPSF, when gravitational instability starts to gradually decline with time, may have a chance to survive. In this paper, we present confirmation that the fragments formed in the EPSF can survive through this extreme phase and form giant protoplanets (hereafter, GPPs) on large, relatively stable orbits. This finding is made possible by the employment of expanded computational resources, by improvements in the numerical model, and by the use of a wider parameter space in comparison to previous works. ", "conclusions": "We have studied the long-term evolution of disks that are formed by the self-consistent collapse of prestellar cores. Our model yields gas giant formation starting from initial conditions of the early stages of star formation. The initial cores are more gravitationally unstable and have greater angular momenta than similar models studied in the past, and a large number of models have been run with relatively high resolution. An early burst mode of evolution is characterized by the formation of clumps which are then driven into the inner disk. However, in a small subset of models, massive fragments are formed on wide orbits and settle into stable orbits of radius $\\gtrsim 50$ AU. An interesting feature is that minidisks around these fragments can be {\\it counterrotating} with respect to the disk. Sometimes, the final orbit can be much larger or alternatively there can be a slowly continuing inward migration. We believe that our results can explain the purported observations of giant planets on wide orbits. By extrapolation, they may also represent the first stages of the eventual formation of a low mass brown dwarf companion." }, "1003/1003.1706_arXiv.txt": { "abstract": "We use $\\sim$88 arcmin$^2$ of deep ($\\gtrsim$26.5 mag at $5\\sigma$) NICMOS data over the two GOODS fields and the HDF South to conduct a search for bright $z\\gtrsim7$ galaxy candidates. This search takes advantage of an efficient preselection over 58 arcmin$^2$ of NICMOS $H_{160}$-band data where only plausible $z\\gtrsim7$ candidates are followed up with NICMOS $J_{110}$-band observations. $\\sim$248 arcmin$^2$ of deep ground-based near-infrared data ($\\gtrsim25.5$ mag, $5\\sigma$) is also considered in the search. In total, we report 15 $z_{850}$-dropout candidates over this area -- 7 of which are new to these search fields. Two possible $z\\sim9$ $J_{110}$-dropout candidates are also found, but seem unlikely to correspond to $z\\sim9$ galaxies (given the estimated contamination levels). The present $z\\sim9$ search is used to set upper limits on the prevalence of such sources. Rigorous testing is undertaken to establish the level of contamination of our selections by photometric scatter, low mass stars, supernovae (SNe), and spurious sources. The estimated contamination rate of our $z\\sim7$ selection is $\\sim$24\\%. Through careful simulations, the effective volume available to our $z\\gtrsim7$ selections is estimated and used to establish constraints on the volume density of luminous ($L_{z=3}^*$, or $\\sim$$-$21 mag) galaxies from these searches. We find that the volume density of luminous star-forming galaxies at $z\\sim7$ is 13$_{-5}^{+8}$$\\times$ lower than at $z\\sim4$ and $>$25$\\times$ lower ($1\\sigma$) at $z\\sim9$ than at $z\\sim4$. This is the most stringent constraint yet available on the volume density of $\\gtrsim L_{z=3}^{*}$ galaxies at $z\\sim9$. The present wide-area, multi-field search limits cosmic variance to $\\lesssim$20\\%. The evolution we find at the bright end of the $UV$ LF is similar to that found from recent Subaru Suprime-Cam, HAWK-I or ERS WFC3/IR searches. The present paper also includes a complete summary of our final $z\\sim7$ $z_{850}$-dropout sample (18 candidates) identified from all NICMOS observations to date (over the two GOODS fields, the HUDF, galaxy clusters). ", "introduction": "The recent WFC3/IR camera on the Hubble Space Telescope (Kimble et al.\\ 2006) has completely revolutionized our ability to search for galaxies at $z\\gtrsim7$ due to its extraordinary imaging capabilities in the near-infrared -- allowing for large areas to be surveyed to great depths. Already some 40 credible $z$$\\sim$7-8 galaxy candidates have been identified in the first hundred orbits of observations (Oesch et al.\\ 2010a; Bouwens et al.\\ 2010a,b; McLure et al.\\ 2010; Bunker et al.\\ 2010; Yan et al.\\ 2010; Finkelstein et al.\\ 2010; Wilkins et al.\\ 2010a,b). This compares with $\\sim$15 credible candidates reported thus far from deep, wide-area ground-based observations (Ouchi et al.\\ 2009; Castellano et al.\\ 2010; Hickey et al.\\ 2010) and $\\sim$12 thusfar with NICMOS (e.g., Bouwens et al.\\ 2008, 2009a; Bradley et al.\\ 2008; Oesch et al.\\ 2009; Zheng et al.\\ 2009; $\\sim2$ from Richard et al.\\ 2008). Whereas $\\sim$100 orbits of NICMOS observations were required to obtain 1 $z\\sim7$ credible galaxy candidate (see e.g. \\S4 of Bouwens et al.\\ 2009a), only $\\sim$2.5 orbits of WFC3/IR observations are required to find a similar $z\\sim7$ candidate. Despite these significant advances in our observational capabilities with WFC3/IR to reach deep and identify large numbers of faint $z\\gtrsim7$ galaxies, a full characterization of the galaxy population at $z\\sim7$ requires that we identify large numbers of galaxies at \\textit{both} high and low luminosities. All but $\\sim$6 galaxies in early selections of $z$$\\sim$7-8 galaxies from early WFC3/IR observations over the ERS/HUDF09 fields have magnitudes faintward of 26.5 mag (e.g., Oesch et al.\\ 2010a; Wilkins et al.\\ 2010a; Bouwens et al.\\ 2010c). As such, it is somewhat challenging to characterize the properties of relatively luminous galaxies at $z\\gtrsim6.5$, and some expansion of the number of sources known brightward of 26.5 mag would be beneficial. Such samples are particularly valuable over fields such as GOODS (Giavalisco et al.\\ 2004) where other valuable multiwavelength data exist like deep IRAC (Dickinson \\& GOODS Team 2004) or Chandra coverage (Brandt et al.\\ 2001; Rosati et al.\\ 2002). \\begin{figure*} \\epsscale{1.15} \\plotone{f1.eps} \\caption{\\scriptsize{Deep, wide-area near-IR data available to search for $z\\gtrsim7$ galaxies over the CDF-South GOODS (\\textit{left}) and HDF-North GOODS (\\textit{right}) fields (see also Table~\\ref{tab:obsdata}). Most of the ultra-deep NICMOS and deep NICMOS observations were already presented in Bouwens et al.\\ (2008) and included here on this figure (red and dark orange regions corresponding to regions with $5\\sigma$ depths of $\\gtrsim28$ mag and $\\gtrsim26.5$ mag, respectively). What is new in the current analysis is the moderately deep (26.9 mag at $5\\sigma$), wide-area ($\\sim$58 arcmin$^2$: $>$60 NIC3 pointings) NICMOS $H_{160}$-band data (light orange squares: Conselice et al.\\ 2010). These data can be used to search for $z\\gtrsim7$ galaxy candidates by identifying those sources that are not detected in the ACS optical bands, but have very red $z_{850}-H_{160}$ colors and possess reasonably blue $H-5.8\\mu m$ and $3.6\\mu m-5.8\\mu m$ colors. Unfortunately, these criteria (while very demanding) are not sufficient to place strong enough constraints on the redshift of the candidates, and so we also obtained deep (27.0 mag at $5\\sigma$) NICMOS $J_{110}$-band imaging (\\textit{magenta squares}) over the best $z\\gtrsim7$ candidates (Table~\\ref{tab:candlist}: see \\S3.2) with GO program 11144 (PI: Bouwens). Also shown in dark orange -- and annotated with ``Y'' or ``H'' (for the GO11192 H. Yan et al.\\ [2010, in prep] or Henry et al.\\ 2009 fields, respectively) -- are several additional NICMOS fields we used to search for $z\\gtrsim7$ galaxies not considered by Bouwens et al.\\ (2008). NICMOS fields previously considered by Bouwens et al.\\ (2008) in searches for $z\\gtrsim7$ galaxies are indicated as follows: ``D'' denotes the HDF-North Dickinson field (Dickinson 1998), ``T'' denotes the HDF-North Thompson field (Thompson et al.\\ 1999), ``U'' denotes the HUDF Thompson field (Thompson et al.\\ 2005), ``S'' denotes the HUDF05 NICMOS field over the HUDF (Oesch et al.\\ 2009), ``1'' denotes the first set of NICMOS parallels to the HUDF (NICP12: Oesch et al.\\ 2009), ``2'' denotes the second set of NICMOS parallels to the HUDF (NICP34: Oesch et al.\\ 2009), and ``G'' denotes the GOODS Parallel NICMOS fields. The blue and dark blue regions correspond to regions with deep and very deep optical ACS coverage, respectively ($5\\sigma$ depths of $\\gtrsim28$ and $\\gtrsim29$ mag). The position of the ground-based ISAAC+MOIRCS search areas is not shown here to minimize confusion (but is presented in Figure 1 of Bouwens et al.\\ 2008). Also not included on this figure are the NICMOS search fields (Zirm et al.\\ 2007) over the HDF South (Williams et al.\\ 2000). The position of the Early Release Science WFC3/IR observations with the CDF South (not used here for $z\\sim7$ LF constraints) is indicated by the light shaded red region. A link to our NICMOS reductions over the GOODS fields is provided at http://firstgalaxies.org//astronomers-area.} \\label{fig:obsdata}} \\end{figure*} Fortunately, for the selection of luminous $z$$\\sim$7-8 galaxies, some $\\sim$88 arcmin$^2$ of deep ($>$26.5 mag, $5\\sigma$), wide-area NICMOS observations exist over the two GOODS fields and the HDF South. While $\\sim$23 arcmin$^2$ of those observations have already been used to identify $z$$\\sim$7-8 galaxies (Bouwens et al.\\ 2008; Oesch et al.\\ 2009), $\\gtrsim$60 arcmin$^2$ of those data have yet to be used. Most of these data are associated with the GOODS NICMOS survey (Conselice et al.\\ 2010) or are NICMOS parallels associated with other HST programs. In total, these NICMOS observations cover 2$\\times$ as much area as available in the WFC3/IR observations that made up the Early Release Science Program (PI O'Connell: GO 11359). Meanwhile, these observations cover comparable area to that available in the wide-area HAWK-I $Y$-band observations over the CDF-South (Castellano et al.\\ 2010; Hickey et al.\\ 2010) and less area than the Subaru Suprime-Cam (Ouchi et al.\\ 2009) observations over the HDF-North, but are deeper on average (by $\\sim$0.2 and $\\sim$0.7 mag, respectively). In the present work, we use these wide-area NICMOS observations to identify a small sample of luminous star-forming galaxies at $z\\sim7$, adding significantly to that known, and performing a search at $z\\sim9$. Some of the present $z\\sim7$ sample has already been used by Gonzalez et al.\\ (2010) and Labb{\\'e} et al.\\ (2010b) to do stellar population modelling of luminous $z\\sim7$ galaxies and to extend measures of the stellar mass density and specific star formation rate to $z\\sim7$. Here we describe the selection of those $z\\gtrsim7$ candidates in detail, discuss the properties and layout of the NICMOS observations, and summarize the properties of the sample. We also estimate the contamination rates and the selection volume for this sample. Finally, we use this search to derive a constraint on the volume density of $L_{z=3}^*$ (or $\\sim$$-21$ mag) galaxies at $z\\sim7$. We will also incorporate results from the $\\sim$248 arcmin$^2$ Bouwens et al.\\ 2008 search for $z\\sim7$ galaxies in deep ground-based data over GOODS. The structure of this paper is as follows. In \\S2, we summarize the observational data. In \\S3, we describe our search results for $z\\gtrsim7$ galaxies. In \\S4, we use the observational search results to derive a constraint on the bright end of the $z\\sim7$ and $z\\sim9$ $UV$ LFs. Finally, we provide a summary (\\S5). Throughout this work, we often denote luminosities in terms of the characteristic luminosity $L_{z=3}^{*}$ at $z\\sim3$ (Steidel et al.\\ 1999), i.e., $M_{UV,AB}=-21.07$. Where necessary, we assume $\\Omega_0 = 0.3$, $\\Omega_{\\Lambda} = 0.7$, $H_0 = 70\\,\\textrm{km/s/Mpc}$. Although these parameters are slightly different from those determined from the WMAP seven-year results (Komatsu et al.\\ 2010), they allow for convenient comparison with other recent results expressed in a similar manner. We express all magnitudes in the AB system (Oke \\& Gunn 1983). \\begin{deluxetable*}{ccccccc} \\tablecolumns{7} \\tablecaption{ACS + NICMOS + Ground-Based Imaging data used for our $z$ and $J$ dropout searches.\\tablenotemark{a}\\label{tab:obsdata}} \\tablehead{ \\colhead{} & \\colhead{} & \\multicolumn{4}{c}{5$\\sigma$ Depth\\tablenotemark{b}} & \\colhead{} \\\\ \\colhead{Name} & \\colhead{Area} & \\colhead{$z_{850}$} & \\colhead{$J_{110}$} & \\colhead{$H_{160}$} & \\colhead{$K_{s}$} & \\colhead{Ref\\tablenotemark{c}}} \\startdata \\multicolumn{7}{c}{New NICMOS Fields} \\\\ GOODS NICMOS Survey & 44.6 & 27.5 & ---\\tablenotemark{d} & 26.9 & $\\sim$25\\tablenotemark{h} & [1] \\\\ Teplitz Parallels\\tablenotemark{e} & 9.3 & 27.5 & ---\\tablenotemark{d} & 26.9 & $\\sim$25\\tablenotemark{h} & [2] \\\\ HDF-South & 4.3 & 26.4 & 26.3\\tablenotemark{f} & 26.7 & 25.7 & [3] \\\\ Yan Survey & 3.1\\tablenotemark{g} & 27.5 & 26.9 & 26.7 & --- & [4] \\\\ Henry Fields & 3.4\\tablenotemark{g} & 27.5 & 26.9 & 26.7 & --- & [5] \\\\ \\multicolumn{7}{c}{NICMOS Fields Already Considered by Bouwens et al.\\ (2008)} \\\\ HDF-North Dickinson & 4.0 & 27.8 & 27.0 & 27.0 & 25.6 & [6,7] \\\\ HDF-North Thompson & 0.8 & 27.8 & 28.0 & 28.1 & 25.6 & [8,7] \\\\ HUDF Thompson & 5.8 & 29.0 & 27.6 & 27.4 & 26.0 & [9,10] \\\\ HUDF Stiavelli & 0.7 & 29.0 & 28.1 & 27.9 & 26.0 & [10,11] \\\\ HUDF-NICPAR1 & 1.3 & 28.6 & 28.6 & 28.4 & -- & [11,12] \\\\ HUDF-NICPAR2 & 1.3 & 28.6 & 28.6 & 28.4 & -- & [11,12] \\\\ GOODS Parallels & 9.3 & 27.5 & 27.0 & 26.9 & $\\sim25$\\tablenotemark{h} & [12] \\\\ \\multicolumn{7}{c}{Fields with Ground-Based Data Already Considered in Bouwens et al.\\ (2008)} \\\\ ISAAC v2.0 (CDFS) & 136 & 27.5 & $\\sim25.4$\\tablenotemark{h} & $\\sim24.8$\\tablenotemark{h} & $\\sim25$\\tablenotemark{h} & [13] \\\\ MOIRCS GTO-2 (HDFN) & 28 & 27.5 & 25.6 & -- & 25.6 & [7] \\\\ MOIRCS GTO-1,3,4 (HDFN) & 84 & 27.5 & 24.2 & -- & 24.4 & [7] \\\\ \\enddata \\tablenotetext{a}{The layout of these search fields is illustrated in Figure~\\ref{fig:obsdata}.} \\tablenotetext{b}{$5\\sigma$ depths for ACS and NICMOS data given for a $0.6''$-diameter aperture and for a $\\sim1.0''$-diameter aperture for the ground-based $K_s$-band data. No correction has been made for the nominal light outside these apertures (for example, for a point source, the correction is typically $\\sim$0.2 mag) to keep the present estimates as empirical as possible. This is in contrast to the convention that we use in some previous work (e.g., Bouwens et al.\\ 2008) where such corrections have been made.} \\tablenotetext{c}{References: [1] Conselice et al.\\ 2010, [2] Siana et al.\\ (2007), [3] Labb{\\'e} et al.\\ (2003), Zirm et al.\\ (2007) [4] H. Yan et al.\\ (2010, in prep) [5] Henry et al.\\ (2009), [6] Dickinson 1998, [7] Kajisawa et al.\\ 2006, Ouchi et al.\\ 2007, [8] Thompson et al. 1999, [9] Thompson et al. (2005), [10] Labb{\\'e} et al.\\ (2006), [11] Oesch et al.\\ 2007, [12] Bouwens \\& Illingworth (2006), Riess et al.\\ (2007), Siana et al.\\ (2007), [13] Retzlaff et al.\\ 2010, Mannucci et al.\\ 2007, Stanway et al.\\ (2008).} \\tablenotetext{d}{NICMOS $J_{110}$-band observations, with $5\\sigma$ depths of 27.0 mag ($0.6''$-diameter aperature) were acquired in those fields with $z\\gtrsim7$ candidates.} \\tablenotetext{e}{40-orbit NICMOS $H_{160}$-band observations taken in parallel with ACS SBC far-UV observations of the HUDF (Siana et al.\\ 2007; GO10403: PI Teplitz)} \\tablenotetext{f}{The $J$-band observations here are from the deep FIRES observations over the HDF-South with ISAAC (Labb{\\'e} et al. 2003).} \\tablenotetext{g}{Not including overlap with the GOODS NICMOS Survey (Conselice et al.\\ 2010)} \\tablenotetext{h}{The depth of the near-IR data over the CDF-South varies by $\\sim$0.2-0.4 mag depending upon the observational conditions in which the ISAAC data were taken (Mannucci et al.\\ 2006; Stanway et al.\\ 2008; Retzlaff et al.\\ 2010).} \\end{deluxetable*} ", "conclusions": "We have taken advantage of $\\sim$88 arcmin$^2$ of deep, wide-area NICMOS data to search for $z\\gtrsim7$ galaxies within the HUDF, the two GOODS fields, and the HDF South. This search incorporates $\\sim$65 arcmin$^2$ of wide-area NICMOS data not previously used to identify $z\\gtrsim7$ galaxies.\\footnote{A small fraction of these 65 arcmin$^2$ (4 arcmin$^2$) had previously been used by Henry et al.\\ (2009) for a $z\\sim7$ $z_{850}$-dropout search.} $\\sim$248 arcmin$^2$ of deep ground-based data ($\\gtrsim25.5$ mag, $5\\sigma$) previously used by Bouwens et al.\\ (2008) is also considered. In total, we find $\\sim$7 plausible $z\\sim7$ $z_{850}$-dropout candidates in the new NICMOS observations (six of which are being reported for the first time) and $\\sim$2 possible (but probably unlikely) $z\\sim9$ $J_{110}$-dropout candidates. These candidates significantly add to the number of luminous $z\\gtrsim7$ candidates known within the GOODS fields and improve our constraints on the volume density of luminous galaxies at $z\\sim7$. These candidates have recently been used to model the stellar populations of bright $z\\sim7$ galaxies (Gonzalez et al.\\ 2010; Labb{\\'e} et al.\\ 2010b). When taken together with the NICMOS data already considered in Bouwens et al.\\ (2008), we have identified 15 $z\\sim7$ $z_{850}$-dropout candidates in total from $\\sim$88 arcmin$^2$ of deep NICMOS data. After running detailed simulations to estimate the selection volumes and contamination rates, we use our new $z\\sim7$ samples (plus ground-based search area) to update the Bouwens et al.\\ (2008) determination of the $UV$ LF at $z\\sim7$ and to strengthen our constraints on the LF at $z\\sim9$. We find that the bright end of the UV LF at $z\\sim7$ is 13$_{-5}^{+8}$$\\times$ lower at $z\\sim7$ than at $z\\sim4$. At $z\\sim9$, the $UV$ LF is a factor of $>$25$\\times$ lower ($1\\sigma$) than at $z\\sim4$, assuming that at most one of the $J_{110}$-dropout candidates identified here is at $z\\sim9$." }, "1003/1003.5125_arXiv.txt": { "abstract": " ", "introduction": "The postulation of the existence of additional spacelike dimensions in nature, that can be as large as a few micrometers \\cite{ADD} or even infinite in size \\cite{RS}, has led to the idea of a higher-dimensional gravitational theory with a fundamental energy scale $M_*$ much smaller than the traditional Planck scale $M_P$. If this can be realised with $M_*$ close to the TeV scale, present of future experiments may accelerate particles at energies beyond this new gravity scale. This will unavoidably lead to the occurence of strong gravity effects in particle collisions and the production of heavy final states, including miniature black holes \\cite{creation}. The lifetime of these black holes is expected to be very short as they instantaneously decay via the emission of Hawking radiation \\cite{Hawking} (for detailed reviews of their properties, see \\cite{Kanti, reviews}). Since these black holes will be created and decay in front of our detectors, it is anticipated that the emission of Hawking radiation will be the main obervable signature of their creation and, at the same time, a manifestation of the existence of additional spacelike dimensions in nature. As a result, the study of the emission of radiation by higher-dimensional black holes has been the subject of an intensive research activity over the last years. This includes the emission from both spherically-symmetric \\cite{KMR, HK1, Barrau, Jung, BGK, Naylor, Park, Cardoso, CEKT1, Dai} and rotating \\cite{FS-rot, IOP, Nomura, HK2, IOP2, Jung-super, Jung-rot, DHKW, CKW, CDKW, CEKT2, CEKT3, CEKT4, CDKW2, KKKPZ, CDKW3} black holes in the form of zero and non-zero spin fields. In order to simplify the analysis, the emitted fields are assumed to be minimally-coupled to gravity but otherwise free as well as massless. Nevertheless, in the context of the four-dimensional analysis \\cite{Page} it was found that for certain particles and mass of the black hole, the particle mass can significantly (up to 50\\%) suppress the emission rate. Recently, a set of works \\cite{Sampaio} has addressed the question of the role of the mass of the emitted field (as well as that of the charge) for emission on the brane by a higher-dimensional black hole. Here, we extend this analysis by considering the case of a higher-dimensional black hole with a non-vanishing angular momentum emitting massive scalar fields. We perform a comprehensive study of the absorption probability and energy emission rate for a range of values of the mass of the emitted field, number of extra dimensions, and angular momentum of the black hole. By integrating over the entire frequency range, we compute the total emissivities and obtain the suppression factors in each case. We also consider the cases of both bulk and brane emission, and pose the additional question of whether the presence of the mass of the emitted field can affect the bulk-over-brane energy ratio and threaten the dominance of the brane channel. The outline of this paper is as follows: In section 2, we study the emission of massive scalar fields by a higher-dimensional, simply rotating black hole in the bulk; we compute the value of the absorption probability both analytically and numerically and compare the two sets of results; finally, we derive the exact energy emission spectra and discuss their behaviour. In section 3, we turn to the brane and perform the same tasks. The total emissivities for bulk and brane emission are derived in section 4, and the bulk-over-brane ratio is computed for a large number of values of the parameters of the theory. We close with our conclusions in section 5. ", "conclusions": "In this work, we have moved towards the direction of considering the emission of realistic particle states by a higher-dimensional, simply rotating black hole. We have studied the emission of massive scalar fields both in the bulk and on the brane, and investigated the role that the mass of the field plays in the corresponding energy spectra profiles and in the bulk-over-brane energy ratio. The emission of Hawking radiation in the bulk in the form of massive scalar fields was studied first. The radial part of the field equation was first solved analytically, and an expression for the absorption probability was found that helped us investigate low-energy aspects of the emission. Next, by using numerical analysis, the exact value of the absorption probability was determined and its dependence on the mass of the emitted field, in conjunction with the number of extra dimensions and angular-momentum of the black hole, was studied. As expected, the presence of the mass term caused the suppression of the absorption probability as additional energy is required for the emission of a massive field. Our numerical and analytical results were directly compared, and found to be in excellent agreement in the low and intermediate energy regimes for scalar fields with a mass smaller than (0.5-1) TeV. The exact numerical value of the absorption probability was subsequently used to derive the differential emission rate per unit time and unit frequency in the bulk. Particular care was taken so that a large enough number of scalar modes ($N_{bu} \\simeq 5500$) was summed up in our computation of the energy spectra. The mass term caused the suppression of the energy spectra in the low and intermediate-energy regimes, compared to the massless case: for low values of $n$ and $a_*$ and $m_\\Phi=0.8$, the suppression is of the order of 50\\%, while it becomes smaller in magnitude as either $n$ or $a_*$ increases. The same task was performed for the emission of massive scalar fields on the brane. The value of the absorption probability was again found both analytically and numerically, and it was shown that the two sets of results are in very good agreement, in the lowest part of the spectrum, up to masses of order (250-500) GeV. The exact profile of the energy spectra on the brane was found next in terms of the parameters ($m_\\Phi, n, a_*$), with the mass term causing again a significant suppression in their value. The suppression was larger than the one in the bulk decreasing the value of the energy emission rate to approximately 40\\% of that in the massless case, for low values of $n$ and $a_*$ and for $m_\\Phi=0.8$. As in the case of bulk emission, a considerable number of modes ($N_{br} \\simeq 1700$) was summed up in our calculation so that the computed spectra are as close as possible to the real ones. The role of the mass of the emitted field in the bulk-over-brane energy ratio was also investigated. The total energy emissivities of bulk and brane emission were derived and directly compared. In agreement with previous analyses \\cite{EHM, HK1, CDKW2} -- that we have generalised by considering a larger range of parameters of both $n$ and $a_*$ -- we found that the bulk channel remains sub-dominant to the brane one; nevertheless, the bulk-over-brane ratio takes a considerable value especially for a large number of extra dimensions and a slowly rotating black hole. We further found that the presence of the mass of the emitted field increases the percentage of energy which is spent by the black hole in the bulk. For a small number of extra dimensions and a low value of the angular-momentum of the black hole, the enhancement of the bulk channel over the brane one can reach the value of 33\\% if $m_\\Phi=0.8$. In conclusion, in this work we have performed a comprehensive study of the emission of massive scalar fields by a higher-dimensional, simply rotating black hole both in the bulk and on the brane. We have studied the dependence of the absorption probabilities and energy emission rates on all parameters of the theory, and compared analytic and numerical methods for the computation of their value. We have confirmed the importance of the emission of a higher-dimensional black hole both in the bulk and on the brane, and demostrated that properties of the emitted field, such as its mass which was up to now largely ignored, can play a significant role in the bulk-over-brane energy balance. \\bigskip {\\bf Acknowledgments.} We are deeply grateful to Athanasios Dedes for useful discussions regarding our numerical code. We also acknowledge participation in the RTN Universenet (MRTN-CT-2006035863-1 and MRTN-CT-2004-503369)." }, "1003/1003.5852_arXiv.txt": { "abstract": "USNO-B1.0 and 2MASS are the most widely used full-sky surveys. However, 2MASS has no proper motions at all, and USNO-B1.0 published only relative, not absolute (i.e. on ICRS) proper motions. We performed a new determination of mean positions and proper motions on the ICRS system by combining USNO-B1.0 and 2MASS astrometry. This catalog is called PPMXL\\footnote{VO-access to the catalog is possible via http://vo.uni-hd.de/ppmxl} , and it aims to be complete from the brightest stars down to about $V \\approx 20$ full-sky. PPMXL contains about 900 million objects, some 410 million with 2MASS photometry, and is the largest collection of ICRS proper motions at present. As representative for the ICRS we chose PPMX. The recently released UCAC3 could not be used because we found plate-dependent distortions in its proper motion system north of -20$^\\circ$ declination. UCAC3 served as an intermediate system for $\\delta \\leq -20^\\circ$. The resulting typical individual mean errors of the proper motions range from 4 mas/y to more than 10 mas/y depending on observational history. The mean errors of positions at epoch 2000.0 are 80 to 120 mas, if 2MASS astrometry could be used, 150 to 300 mas else. We also give correction tables to convert USNO-B1.0 observations of e.g. minor planets to the ICRS system. ", "introduction": "According to IAU Resolution B2 of the XXIIIrd General Assembly (1997), the Hipparcos catalog \\citep{1997yCat.1239....0E} is the primary realisation of the International Celestial Reference System (ICRS) at optical wavelengths. Its first and basic extension to higher star densities and fainter limiting magnitudes is Tycho-2 \\citep{2000A&A...355L..27H}, based on observations of the Tycho experiment onboard the ESA-Hipparcos satellite. The early epoch observations of Tycho-2 were taken from new reductions \\citep{1998AJ....115.1212U} of the observations made for the Astrographic Catalog and 143 other ground-based astrometric catalogs. Tycho-2 contains about 2.5 million stars and is 90 percent complete down to $V$ = 11.5. \\citet{2008A&A...488..401R} published the PPMX catalog of positions and proper motions of 18 million stars with limiting magnitude around 15 in a red band. The typical accuracy of the proper motions is about 2 mas/y for 4.5 million stars with first epoch in the Astrographic catalog and about 10 mas/y for all other stars. Very recently, the UCAC3 catalog \\citep{2009yCat.1315....0Z} was released. UCAC3 is based on a new full-sky astrometric survey made in the years 1998 to 2004. The catalog contains some 100 million stars down to $r_U$ = 16 mag. The largest catalog in the optical regime is USNO-B1.0 \\citep{2003AJ....125..984M} with more than one billion objects. However, USNO-B1.0 is not in the system of ICRS; it contains relative, not absolute proper motions \\citep[see][]{2003AJ....125..984M}. A comparison of USNO-B1.0 and PPMX performed in the present work yielded systematic differences (in areas of square degrees) of up to 15 mas/y in proper motion and up to 0.6 arcseconds in positions at epoch 2000.0. The Two Micron All Sky Survey \\citep{2006AJ....131.1163S}, 2MASS, is a complete Sky-Survey in the J, H and K$_s$ bands performed in the years from 1997 to 2001. The Point Source Catalog of about 471 million entries is also a source of accurate astrometric positions, but contains no proper motions. For kinematical studies in the Milky Way a catalog of proper motions in the ICRS system and with a well-defined completeness limit is indispensable. PPMX fulfills this requirement, but with only 18 million stars it is rather small. USNO-B1.0 with inertial proper motions would be a big step forward. When SDSS observations became available, inertial proper motions have been constructed from a combination of SDSS with USNO-B1.0 using SDSS galaxies as reference \\citep{2004AJ....127.3034M,2008AJ....136..895M,2004ApJS..152..103G}. This is, of course, restricted to the SDSS part of the sky. The fact that USNO-B1.0 is not on ICRS creates a problem for minor planet observers. Right ascensions and declinations based on USNO-B1.0, when combined with older epoch observations on ICRS may cause biases in orbit determinations of minor planets. This is of great importance for fly-by manoeuvers of interplanetary spacecraft \\citep{2009DDA....40.1704C}. Combining USNO-B1.0 with 2MASS is such an obvious idea that it is already mentioned by \\citet{2003AJ....125..984M}. In the following we first give a schematic overview of the entire procedure(section \\ref{over}), then we describe in detail the initial steps to {\\em coarsely} put USNO-B1.0 to the ICRS (section \\ref{prel}). This is followed by a description the combination with 2MASS observations (section \\ref{lsqadj}), the details of the construction of the system of positions and proper motions on ICRS (section \\ref{system}), and we close with an overview of the properties of the catalog (section \\ref{finalcat}). Our approach can be considered an affordable effort to put USNO-B1.0 onto ICRS, and improve the individual proper motions by inclusion of 2MASS. A sophisticated re-reduction of all the material might deliver superior results provided that a better reference catalog were available before Gaia. ", "conclusions": "" }, "1003/1003.5846_arXiv.txt": { "abstract": "We have studied the implications of high sensitivity polarization measurements of objects from the WMAP point source catalogue made using the VLA at 8.4, 22 and 43~GHz. The fractional polarization of sources is almost independent of frequency with a median of $\\approx 2$ per cent and an average, for detected sources, of $\\approx 3.5$ per cent. These values are also independent of the total intensity over the narrow range of intensity we sample. Using a contemporaneous sample of 105 sources detected at all 3 VLA frequencies, we have investigated the spectral behaviour as a function of frequency by means of a 2-colour diagram. Most sources have power-law spectra in total intensity, as expected. On the other hand they appear to be almost randomly distributed in the polarized intensity 2-colour diagram. This is compatible with the polarized spectra being much less smooth than those in intensity and we speculate on the physical origins of this. We have performed an analysis of the correlations between the fractional polarization and spectral indices including computation of the principal components. We find that there is little correlation between the fractional polarization and the intensity spectral indices. This is also the case when we include polarization measurements at 1.4~GHz from the NVSS. In addition we compute 45 rotation measures from polarization position angles which are compatible with a $\\lambda^2$ law. We use our results to predict the level of point source confusion noise that contaminates CMB polarization measurements aimed at detecting primordial gravitational waves from inflation. We conclude that some level of source subtraction will be necessary to detect $r\\sim 0.1$ below 100~GHz and at all frequencies to detect $r\\sim 0.01$. We present estimates of the level of contamination expected and the number of sources which need to be subtracted as a function of the imposed cut flux density and frequency. ", "introduction": "Synchrotron radiation is the primary emission mechanism for sources detected in the radio band. Hence, radio sources are likely to be linearly polarized with a fractional polarization of typically a few percent, depending on the level of order of the magnetic field in the emission producing regions. Large-scale polarization catalogues exist for sources selected at low frequencies such as 1.4~GHz (NVSS - Condon et al. 1998) with $\\approx 2\\times 10^5$ detections. At higher frequencies Jackson et al. (2007) have looked at flat spectrum sources in the CLASS survey (Myers et al., 2003; Browne et al., 2003) and have $\\approx 5000$ detections of 8.4~GHz polarizations. In recent times more information has become available at even higher frequencies. Ricci et al. (2004) have detected significant polarization in 170 sources at 18.5~GHz, while Massardi et al. (2007) present the results from follow-up of the AT20G survey (Sadler et al., 2006). They have detected polarization at 22~GHz in 218 sources with total intensity greater than 0.5Jy at 22~GHz with additional information at 4.8 and 8.4~GHz. \\changed{Murphy et al. (2010) present the full AT20G survey containing 5890 sources stronger than 40mJy at 20GHz. Most sources have near-simultaneous 5 and 8~GHz measurements and 1559 have detections of polarised emission at one or more frequencies.} In addition, Agudo et al (2009) have presented measurements of 146 sources at 86~GHz made using the IRAM telescope and Lopez-Caniego et al. (2009) have created a catalogue of polarized sources detected from the WMAP maps between 22 and 90~GHz. Our main motivation for studying the properties of polarized radio sources is to assess their potential contamination of Cosmic Microwave Background (CMB) polarization data. Measurements have been made of the so-called E-mode polarization which is produced by scalar density fluctuations (Kovac et al. 2002, Readhead et al. 2004, Montroy et al. 2005, Page et al. 2007, Brown et al. 2009, Chiang et al. 2009). Future instruments are aimed at detecting the B-mode signature of primordial gravitational waves produced by inflation (see, for example, Baumann et al. 2009) which is often quantified in terms of the scalar-to-tensor ratio, $r$. These will require not only high signal-to-noise, but also exquisite control of systematics and separation of astrophysical foregrounds. To date the only work on this issue (de Zotti et al 1999, Mesa et al 2002, Tucci et al., 2004) has relied on extrapolation of polarized source counts from frequencies which are factors of $\\approx 20$ lower than where the CMB observations will take place. We will focus on the implications of the measurements we have obtained for a sample of 203 radio sources extracted from the WMAP 22~GHz catalogue which was complete to $\\approx 1{\\rm Jy}$. Most of these sources were also detected in the NVSS (Condon et al., 1998) at 1.4~GHz and 71 were detected at 86~GHz using the IRAM telescope by Agudo et al. (2009); we will make use of this information later. We studied sources in the region with declinations greater than $-34^{\\circ}$ and all were observed at 22 and 43~GHz using the Very Large Array (VLA). A subset of 134 were also observed at 8.4~GHz. Observations were missing for 3 sources due to misidentifications and a further 7 sources were deemed inappropriate to include in the present statistical analysis for a variety of reasons such as the source being very extended and being heavily resolved in our VLA observations. Polarized emission was detected for 123, 169 and 167 sources at 8.4, 22, 43~GHz respectively and 105 were detected at all 3 frequencies. No polarization bias correcttion was applied since all detections are greater than $5\\sigma$ and many are substantially stronger. This ``contemporaneous sample'' provides the first direct information at frequencies relevant to CMB measurements allowing more accurate extrapolations to be made. The details of the observations and the catalogue are presented in Jackson et al. (2010). ", "conclusions": "We have presented some statistical properties of polarized radio sources from our observations of the WMAP point source catalogue at 8.4, 22 and 43~GHz. The main results are: \\begin{itemize} \\item the median level of fractional polarization is around $2-2.5$ per cent and there is no evidence for any dependence on frequency in the range 8.4 to 43~GHz; \\item we have argued that our data are compatible with the spectra in polarized intensity being less smooth than those in total intensity; \\item there is no evidence for a strong correlation between a single intensity spectral index and the fractional polarization, although most of our sources are flat spectrum; \\item we were able to compute rotation measures for 45 out of the 105 sources which were detected contemporaneously at all three frequencies; \\item there is statistical evidence for large rotation measures with a significant number of sources having rotation measures $>1000\\,{\\rm rad}\\,{\\rm m}^{-2}$. \\end{itemize} We then used our results, in conjunction with the dZ05 model, to model the polarized source counts in order to compute the level of confusion noise expected in CMB polarization experiments designed to detect primordial gravitational waves from inflation. We find that \\begin{itemize} \\item the source confusion due to jet-powered radio sources is likely to be the dominant contribution to the total up to relatively high frequencies - in our estimates they even dominate at 220~GHz but we have also pointed out that this could be an artefact of our assumption of a single spectral index for each source; \\item source subtraction will be important at frequencies below 100~GHz if one is to detect $r\\sim 0.1$; \\item if one is to detect $r\\sim 0.01$ some source mitigation strategy will be required at all frequencies. \\end{itemize} In addition, there appears to be good agreement with our modelling and that in the PSM at 30~GHz for flux densities $\\approx 1{\\rm Jy}$. There is one final point we should make. Many of the sources, particularly those with high flux densities, are expected to be significantly variable. This will manifest itself not only in variable total and polarized flux densities, but also in the polarization position angle. This could make it difficult to subtract the effects of source confusion from the CMB polarization signal. If the observations of a particular CMB field are performed over a timescale which is shorter than the timescale of variability, this will require the high resolution observations performed to assist source subtraction to take place contemporaneously in order to make an accurate subtraction. Conversely, if the timescale for observations is longer than the variability, for example, if observations on a particular region are built up a series of short integrations over many days, then the variability of the source polarization could easily average out to a significantly lower observed polarization when the individual integrations are stacked. This effect will make accurate source subtraction difficult since it is likely to be impractical to monitor the level of variability for a substantial number of sources." }, "1003/1003.2657_arXiv.txt": { "abstract": " ", "introduction": "Clouds occur at various geographical locations on Titan with distinct morphologies, which suggests that different mechanisms and meteorological conditions underly cloud formation. The diversity of clouds provides insight into Titan's climate dynamics and exotic methane-based meteorological cycle. Since the first spectroscopic indications of clouds \\citep{Toon1988,Griffith1998, Griffith2000}, they have been imaged near the South pole; in the region of maximum solar insolation at the time \\citep{Brown2002,Roe2002b}. The polar storm clouds seen by Cassini/ISS are clusters of individual small-scale convective clouds \\citep{Porco2005}. These southern polar clouds occasionally develop into massive storms and their frequency changes seasonally \\citep{Schaller2006a, Schaller2006b, Turtle2009}. There are signs of changes on the surface after large storms, which may be the accumulation of precipitation into lakes at the south pole \\citep{Turtle2009}. Clouds that are preferentially aligned near 40$^{\\circ}$S latitude were first observed by \\citet{Roe2005a}. These ``southern temperate'' or ``southern mid-latitude'' clouds evolve over hourly timescales \\citep{Griffith2005,Porco2005}. Clouds composed of ethane have been discovered near the North polar tropopause that likely form from Hadley cell subsidence there \\citep{Griffith2006}. A thin cloud of solid methane was inferred from the relative humidity profile near the Huygens landing site, 110$^{\\circ}$W, 10$^{\\circ}$S \\citep{Tokano2006a}, which seems to form a stratiform layer over the globe \\citep{Adamkovics2007}. The characteristics of clouds in the tropics is consistent with a dry climate there now, unlike the polar atmosphere \\citep{Griffith2009}. Streaky clouds have been catalogued during the Cassini mission that may be closely related to the northern polar lakes \\citep{Brown2009a}. The south pole has low-altitude fog, which indicates the release of methane into the atmosphere from the surface \\citep{Brown2009b}. Both seasonal changes in meteorological conditions and geographical differences between the poles could lead to the distinct cloud morphologies that are observed at various locations. Regions of upwelling in general circulation models (GCM) of Titan \\citep{Tokano2005a,Rannou2006, Mitchell2006,Mitchell2008} are consistent with the location of both southern midlatitude and southern polar clouds. The GCMs each predict that as the large scale motion of the atmosphere changes with seasons, so will the frequency and location of these convective clouds. There are differences among these models for when and where such changes will occur. The observed change in cloud top altitude is also consistent with convective cloud formation \\citep{Griffith2006}, as is the scarcity of clouds, which coincides with the relatively small amount of convective cells compared to Earth \\citep{Lorenz2003}. \\citet{Barth2007} used a dynamic cloud model to reproduce the cloud-top height, horizontal extent, and short lifetimes of individual mid-latitude (or temperate) clouds observed by \\citet{Griffith2006} with Cassini/VIMS. Such cloud models could be used to study the specific triggering mechanism for convection, potentially disentangling the role of synoptic-scale effects from those of global circulation. Measurements of the evolution of clouds on larger spatial scale and longer timescales provide observational constraints for models of cloud formation. The observations we present were made during the course of ground-based observational support of the Cassini flybys and have captured the daily evolution of a massive cloud system on Titan. We monitor the evolution of a southern mid-latitude cloud system over a broad range of timescales and suggest that meso-scale meteorology may have an impact on cloud system formation. ", "conclusions": "After the first observations of mid-latitude clouds, \\citet{Roe2005a} described the possible link between circulation and the latitude of cloud formation, noting that the seasonal change in circulation should lead to changes in the locations of clouds. Dynamical models of circulation were used to support this hypothesis \\citep{Rannou2006} and to quantify predictions of the seasonal changes in convection and precipitation \\citep{Mitchell2006}. In Jan 2007, more than four years after southern summer solstice (SSS) on Titan, the mean daily isolation is greatest near 25$^{\\circ}$S, and yet the formation of clouds is still observed at essentially the same latitude as when the south polar clouds were first imaged. While this is not inconsistent with the range of latitudes where convection is predicted to occur \\citep{Mitchell2006}, the preferred formation of clouds near 40$^{\\circ}$ \\citep{Roe2005b,Turtle2009} and the modest seasonal change in this preferred latitude, is smaller than expected \\citep{Brown2009c,Rodriguez2009}. If the border of the tropical vs. polar climate regions is $\\sim$60$^{\\circ}$ \\citep{Griffith2008}, then clouds at $\\sim$40$^{\\circ}$S should be in a region that is well-approximated by the relative humidity profile measured by Huygens \\citep{Niemann2005, Fulchignoni2005}. But if the distribution of moisture is not meridionally uniform --- perhaps because of the availability of energy to evaporate methane \\citep{Griffith2008} or because of the seasonal cycles of precipitation and a large reservoir of methane in the soil \\citep{Mitchell2008} --- then methane availability may be an additional factor that confines cloud formation to a particular latitude band. Entrainment of methane from moist polar air, or sources of moisture near the surface \\citep{Roe2005b}, would both result in localized enhancements of methane. While there is no overall longitudinal preference for the mid-latitude clouds \\citep{Brown2009c,Turtle2009}, the localization of cloud formation may suggest a non-negligible interaction between the surface and the atmosphere at this time. \\citet{Barth2007} describe how nearby cells compete for and quickly exhaust CAPE within one convective cycle. This contrasts our observations (Figure \\ref{fVIMSobs}), where multiple cells are observed near each other directly prior to a large storm outburst on 29 Jan 2007 (Figure \\ref{fobs}). If large scale circulation is resupplying the necessary CAPE \\citep{Barth2007}, then the evolution of this cloud system may be constrained by the amount of energy needed, and therefore the mid-latitude circulation. \\citet{Schaller2009} hypothesize an atmospheric teleconnection of tropical and south polar cloud formation, occurring over roughly two weeks, mediated by Rossby waves. In our observations, a region of southern temperate cloud formation persisted for at least four consecutive days, remaining stationary for three days, and did not correspond with cloud formation near the south pole. Tropical clouds may require a fundamentally different formation mechanism than temperate clouds. We have presented observations of the spatial localization of a southern mid-latitude cloud system on Titan, with retrievals of the cloud-top altitude and opacity using both ground-based (VLT/SINFONI) and Cassini/VIMS data. The unresolved, possibly stratiform, cloud system is composed of multiple cloud forming cells that are likely driven by convection, with moist air that may be supplied from higher latitude regions, where insolation is sufficient to cause evaporation from the surface. The mechanisms for southern mid-latitude cloud formation may be related to the mesoscale dynamics of large convective systems on Earth." }, "1003/1003.5900_arXiv.txt": { "abstract": "Using the Very Long Base Array, we observed the young stellar object EC 95 in the Serpens cloud core at eight epochs from December 2007 to December 2009. Two sources are detected in our field, and are shown to form a tight binary system. The primary (EC 95a) is a 4--5 $M_\\odot$ proto-Herbig AeBe object (arguably the youngest such object known), whereas the secondary (EC 95b) is most likely a low-mass T Tauri star. Interestingly, both sources are non-thermal emitters. While T Tauri stars are expected to power a corona because they are convective while they go down the Hayashi track, intermediate-mass stars approach the main sequence on radiative tracks. Thus, they are not expected to have strong superficial magnetic fields, and should not be magnetically active. We review several mechanisms that could produce the non-thermal emission of EC 95a, and argue that the observed properties of EC 95a might be most readily interpreted if it possessed a corona powered by a rotation-driven convective layer. Using our observations, we show that the trigonometric parallax of EC 95 is $\\pi$ = 2.41 $\\pm$ 0.02 mas, corresponding to a distance of 414.9$^{+4.4}_{-4.3}$ pc. We argue that this implies a distance to the Serpens core of 415 $\\pm$ 5 pc, and a mean distance to the Serpens cloud of 415 $\\pm$ 25 pc. This value is significantly larger than previous estimates ($d$ $\\sim$ 260 pc) based on measurements of the extinction suffered by stars in the direction of Serpens. A possible explanation for this discrepancy is that these previous observations picked out foreground dust clouds associated with the Aquila Rift system rather than Serpens itself. ", "introduction": "An accurate knowledge of the physical properties of young stellar objects (like their mass, age and luminosity) is important to constrain theoretical pre--main sequence evolutionary models. The determination of these properties, however, depends critically on the availability of accurate distances. Unfortunately, since distances to regions of star formation are often uncertain by more than 20 or 30\\%, errors on the luminosity and age of young stars are typically about 70\\%. Significant progress has been possible in recent years thanks to Very Long Baseline Interferometry (VLBI) observations, particularly with the Very Long Base Array (VLBA --Loinard et al.\\ 2005, 2007, 2008; Torres et al.\\ 2007, 2009; Menten et al.\\ 2007; Xu et al.\\ 2006). Owing to the very accurate astrometry delivered by such instruments, trigonometric parallaxes (and therefore distances) can be measured very precisely if multi-epoch observations spread over a few years are obtained. VLBI instruments are only sensitive to high surface brightness emission, and can only detect objects where non-thermal processes are at work. Such non-thermal sources must, therefore, be identified in the regions of interest before their distance can be measured using multi-epoch VLBI observations. Fortunately, many young stars are magnetically active, and do exhibit detectable levels of non-thermal radio emission.\\footnote{Theoretically, only {\\it low-mass} young stars are expected to be convective, and to have strong superficial magnetic fields. However, the young 6 $M_\\odot$ B4V star S1 in Ophiuchus (Loinard et al.\\ 2008) is known to exhibit non-thermal radio emission, although theoretical arguments suggest it should be radiative. We will come back to this point in Sect.\\ 4.2.} This type of emission is typically characterized by strong variability, some level of circular polarization, and a negative spectral index. Also, for magnetically active stars, there is a good correlation between X-ray and non-thermal radio emission (Benz \\& G\\\"udel 2004) so young stars with detectable levels of non-thermal radio emission are associated with bright X-ray sources. In this work, we will focus on the star-forming region associated with the Serpens molecular cloud (Strom et al.\\ 1974; see Eiroa 1992 and Eiroa et al.\\ 2008 for two recent reviews). More specifically, we will concentrate on the SVS~4 region (Strom et al.\\ 1976), an infrared cluster of at least 11 pre-main sequence sources deeply embedded within the Serpens core, and one of the densest young stellar sub-clusters known, with a stellar mass density of $\\sim\\, 10^5\\, {\\rm M}_\\odot$ {\\rm pc}$^{-3}$ (Eiroa \\& Casali 1989). In the direction of SVS 4, Preibisch (1998) detected a bright X-ray source ($L_X$ $\\sim$ 4 $\\times$10$^{31}$ erg s$^{-1}$; Preibisch 2003a) now known to be associated with the infrared object EC 95 (Preibisch 2003a, Eiroa \\& Casali 1992). Using IR spectroscopy, Preibisch (1999) showed that EC 95 is $\\sim$ K2 star, and a comparison with theoretical tracks indicates that it is a very young ($\\sim$ 10$^5$ yr), intermediate mass star, presumably a precursor of a $\\sim$ 4 {\\rm M$_{\\odot}$} Herbig AeBe star. As a consequence, Preibisch (1999) argued that EC 95 is a {\\em proto Herbig AeBe star}. \\begin{landscape} \\begin{table*}[!htbp] \\small \\begin{center} \\caption{Measured Source Positions and Flux Densities} \\begin{tabular}{lccccccccc}\\hline\\hline Mean UT date& &Julian Day & $\\alpha$(J2000.0) & $\\sigma_{\\alpha}$ & $\\delta$(J2000.0) & $\\sigma_\\delta$ & $f_\\nu$ $\\pm$ $\\sigma_{f_\\nu}$&$\\sigma$&T$_b$\\\\ (yyyy.mm.dd&hh:mm)& & $18^{\\rm h}29^{{\\rm m}}$& & $+1^{\\circ}12'$& & (mJy)& (mJy beam$^{-1}$)& ($10^7$ K)\\\\ \\hline \\noindent EC 95a& & & & & & & & &\\\\ 2008.11.29&21:34&2454800.40&57\\rlap.{$^{\\rm s}$}8918723 &0\\rlap.{$^{\\rm s}$}00000168&46\\rlap.{''}110068&0\\rlap.{''}000078&1.24 $\\pm$ 0.08&0.08&1.71\\\\ 2009.02.27&15:42&2454890.15&57\\rlap.{$^{\\rm s}$}8921768 &0\\rlap.{$^{\\rm s}$}00000049&46\\rlap.{''}106950&0\\rlap.{''}000017&4.01 $\\pm$ 0.15&0.08&5.55\\\\ 2009.12.05&21:12&2455171.38&57\\rlap.{$^{\\rm s}$}8922183 &0\\rlap.{$^{\\rm s}$}00000070&46\\rlap.{''}095333&0\\rlap.{''}000032&2.77 $\\pm$ 0.16&0.09&4.37\\\\ \\noindent EC 95b & & & & & & & & &\\\\ 2007.12.22&19:38&2454457.32&57\\rlap.{$^{\\rm s}$}8909548 &0\\rlap.{$^{\\rm s}$}00000094&46\\rlap.{''}108014&0\\rlap.{''}000032&1.79 $\\pm$ 0.10&0.05&2.14\\\\ 2008.06.29&07:39&2454646.82&57\\rlap.{$^{\\rm s}$}8909523 &0\\rlap.{$^{\\rm s}$}00000370&46\\rlap.{''}107340&0\\rlap.{''}000145&0.47 $\\pm$ 0.16&0.09&0.61\\\\ 2008.09.15&02:33&2454724.61&57\\rlap.{$^{\\rm s}$}8908014 &0\\rlap.{$^{\\rm s}$}00000087&46\\rlap.{''}105880&0\\rlap.{''}000029&4.51$\\pm$ 0.23&0.11&4.43\\\\ 2008.11.29&21:34&2454800.40&57\\rlap.{$^{\\rm s}$}8908963 &0\\rlap.{$^{\\rm s}$}00000168&46\\rlap.{''}104204&0\\rlap.{''}000078&1.03 $\\pm$ 0.08&0.08&1.45\\\\ 2009.02.27&15:42&2454890.15&57\\rlap.{$^{\\rm s}$}8911287 &0\\rlap.{$^{\\rm s}$}00000575&46\\rlap.{''}103781&0\\rlap.{''}000148&0.41 $\\pm$ 0.16&0.08&0.45\\\\ 2009.06.03&09:22&2454985.89&57\\rlap.{$^{\\rm s}$}8910820 &0\\rlap.{$^{\\rm s}$}00000408&46\\rlap.{''}104144&0\\rlap.{''}000096&1.57 $\\pm$ 0.24&0.11&1.96\\\\ 2009.08.31&03:35&2455074.65&57\\rlap.{$^{\\rm s}$}8909107 &0\\rlap.{$^{\\rm s}$}00000085&46\\rlap.{''}103210&0\\rlap.{''}000030&3.11 $\\pm$ 0.14&0.07&4.19\\\\ \\hline\\hline \\label{tab:pa} \\end{tabular} \\end{center} \\end{table*} \\end{landscape} Rodr\\'{\\i}guez et al.\\ (1980) and Eiroa et al.\\ (1995) detected EC 95 in VLA observations as an unresolved and very strong, radio source (one of the strongest in the survey of Eiroa et al.\\ 2005). This radio emission is strongly variable, and has a negative spectral index, suggesting a non-thermal (gyrosynchrotron) origin (Smith 1999). In order to constrain further the origin of that radio emission, Forbrich et al. (2007) obtained combined high sensitivity Very Large Array (VLA) and High Sensitivity Array (HSA; a heterogeneous VLBI instrument that includes the VLBA, the VLA, and the Effelsberg and Green Bank single-dish telescopes) observations. The source is detected with the VLA as a compact $\\sim$ 0.5 mJy source, but not on the long baselines of the HSA (at a 3$\\sigma$ upper limit of about 50 $\\mu$Jy). They argue that the emission detected at the VLA traces free-free radiation from an ionized wind that might absorb the emission from the underlying active magnetosphere. In the present article, however, we will describe new multi-epoch VLBA observations of EC 95, in which compact non-thermal emission {\\em is} detected on very long baselines. We will use these data to estimate the trigonometric parallax of that source, and of the Serpens core as a whole. As discussed by Eiroa et al.\\ (2008), the distance to Serpens has been a matter of some controversy over the years, with estimates ranging from 210 {\\rm pc} to 700 {\\rm pc} (see Sect.\\ \\ref{sect:distance}). \\begin{figure}[!t] \\begin{center} \\includegraphics[width=0.4\\linewidth,angle=-90]{fig1.eps} \\end{center} \\caption{Relative position of the astronomical target, as well as of the main and secondary calibrators.} \\end{figure} ", "conclusions": "In this paper, we have presented multi-epoch VLBA observations of EC 95, a young stellar object located in the SVS 4 sub-cluster of the Serpens molecular core. Our data demonstrate that EC 95 is in fact a tight binary system with a separation of about 15 mas. The primary (EC 95a) appears to be a 4--5 $M_\\odot$ intermediate-mass Herbig AeBe protostar, whereas the secondary (EC 95b) is most likely a low-mass T Tauri. At radio wavelengths, the secondary is on average brighter than the primary. It is, therefore, also likely to dominate the fairly bright X-ray emission from the system. The primary, on the other hand, contributes most of the infrared and of the bolometric luminosity of EC 95. This might naturally explain why the extinction to EC 95 based on infrared observations was systematically found to be much larger than the extinction based on X-ray data. Interestingly, both members of EC 95 appear to be non-thermal radio emitters. While a low-mass T Tauri star such as EC 95b is expected to generate that type of emission, an intermediate-mass such as EC 95a is not. We discussed several mechanisms that could explain the presence of non-thermal emission on EC 95a, and argue that the observed properties of EC 95a might be most readily explained if it possessed a corona powered by a thin, rotation-driven convective layer. The trigonometric parallax of EC 95 appears to be $\\pi$ = 2.41 $\\pm$ 0.02 mas, corresponding to a distance of 414.9$^{+4.4}_{-4.3}$ pc. We argue that this implies a distance to the Serpens core of 415 $\\pm$ 5 pc, and to the Serpens molecular cloud of 415 $\\pm$ 25 pc. This is significantly larger than previous distance estimates based on measurements of the extinction suffered by stars in the direction of Serpens ($d$ $\\sim$ 260 pc). A possible explanation for this discrepancy is that these measurements correspond to the distance to clouds associated with the Aquila Rift located in the foreground of the Serpens cloud. Since our observations resolve EC 95 into a binary, future radio monitoring of that system will allow us to measure the dynamical mass of that system. To our knowledge, this will be the first time that a dynamical mass is obtained for a young Herbig AeBe system, and it will enable us to uniquely constrain theoretical pre-main sequence evolutionary for intermediate-mass stars. The orbital period of the system appears to be 10--20 years, so observations in the next two decades will be required to obtain an accurate mass estimate. To Serpens cloud is about 20 pc across on the plane of the sky, so it is likely to be also about 20 pc deep. Observations similar to those presented here for a sample of young stars distributed across the cloud would help determine the structure of the region, and would allow us to further constrain the mean distance to this important region of star-formation." }, "1003/1003.3717_arXiv.txt": { "abstract": "We present optical and near infrared observations of $Swift$ GRB 080325 classified as a \\lq\\lq Dark GRB\\rq\\rq. Near-infrared observations with Subaru/MOIRCS provided a clear detection of afterglow in $K_{s}$ band, although no optical counterpart was reported. The flux ratio of rest-wavelength optical to X-ray bands of the afterglow indicates that the dust extinction along the line of sight to the afterglow is $A_{V}$ = $2.7 - 10$ mag. This large extinction is probably the major reason for optical faintness of GRB 080325. The $J - K_{s}$ color of the host galaxy, ($J - K_{s}$ = 1.3 in AB magnitude), is significantly redder than those for typical GRB hosts previously identified. In addition to $J$ and $K_{s}$ bands, optical images in $B$, $R_{c}$, $i$', and $z$' bands with Subaru/Suprime-Cam were obtained at about one year after the burst, and a photometric redshift of the host is estimated to be $z_{photo} =$ 1.9. The host luminosity is comparable to $L^{*}$ at $z \\sim 2$ in contrast to the sub-$L^{*}$ property of typical GRB hosts at lower redshifts. The best-fit stellar population synthesis model for the host shows that a large dust extinction ($A_{V} = 0.8$ mag) attributes to the red nature of the host and that the host galaxy is massive ($M_{*} = 7.0 \\times 10^{10} M_{\\odot}$) which is one of the most massive GRB hosts previously identified. By assuming that the mass-metallicity relation for star-forming galaxies at $z \\sim 2$ is applicable for the GRB host, this large stellar mass suggests the high metallicity environment around GRB 080325, consistent with inferred large extinction. ", "introduction": "The origin of a \\lq\\lq dark\\rq\\rq gamma-ray burst (GRB) remains one of the most serious mysteries of long GRB (hereafter just referred as GRBs) phenomena. A dark GRB is characterized by the faintness of its optical afterglow. While the X-ray afterglow is currently well-explored since the launch of $Swift$, an optical and/or near-infrared detection is reported in only about half of cases, suggesting the significant fraction of dark events. This optically dark nature may be attributed to the (i) large dust extinction around a GRB, (ii) attenuation by the neutral hydrogen in the host and/or intergalactic medium due to a high redshift event, (iii) intrinsically low luminous event of a small GRB fluence, and (iv) low-density medium surrounding a GRB progenitor. However, the actual nature of dark GRBs has not yet been revealed well. Recently, the metallicity environment of dark GRBs is getting a lot more attention. The low metallicity has been theoretically suggested to be required to produce the rapid rotating progenitors and relativistic explosions associated with GRBs \\citep{2006ARA&A..44..507W,2006A&A...460..199Y}. As for observational studies, \\cite{2006Natur.441..463F} showed that GRBs preferentially occur in small faint irregular galaxies. This suggests the low-metallicity environment around GRBs because mass (or luminosity)-metallicity relation exists among galaxies. In fact, the low-metallicity nature of GRB hosts based on spectroscopic observations are reported by \\cite{2006AcA....56..333S}. \\cite{2008AJ....135.1136M} compared the chemical abundances at the sites of 12 nearby ($z < 0.14$) type Ic supernovae (SN Ic) that showed broad lines, but had no observed gamma-ray burst (GRB), with the chemical abundances in five nearby ($z < 0.25$) galaxies at the sites of GRBs. They showed that there is a critical metallicity below which GRBs occur in contrast to type Ic SNe, which occur in more metal rich environment. In this context, \\cite{2009ApJ...702..377K} predicted the critical stellar mass of GRB hosts, which is a stellar-mass counterpart of the critical metallicity, as a function of redshift, assuming an evolution of the mass-metallicity relation by \\cite{2005ApJ...635..260S} and the presence of the critical metallicity \\citep{2008AJ....135.1136M}. They showed that most GRB hosts collected by \\cite{2009ApJ...691..182S} have stellar masses smaller than or comparable to the critical mass. There are, however, exceptionally massive hosts at $z \\gtrsim 0.4$, i.e., possibly exceptionally high metallicity: hosts of GRB 020127 \\citep[$z \\sim 2$:][]{2007ApJ...660..504B}, GRB 020819, and GRB 051022 \\citep[$z =$ 0.4 and 0.8, respectively:][]{2009ApJ...691..182S}. For GRB 020819, \\cite{2010ApJ...712L..26L} reported metallicity larger than the critical metallicity at the GRB position as well as in the host. In addition, for GRB 051022 metallicity of the host is slightly larger than the critical value above mentioned \\citep{2009AIPC.1133..269G}. Both GRBs are classified as dark\\footnotemark[1]. \\footnotetext[1]{It is not clear whether GRB 020127 is classified as \\lq\\lq dark\\rq\\rq or not because of no strong constraint on optical-to-X-ray spectral index of the afterglow \\citep[$\\beta_{\\rm OX} < 1.24$:][]{2004ApJ...617L..21J}.} This suggests that high metallicity environment around a GRB may be somehow related to its \\lq\\lq dark\\rq\\rq nature. Therefore, investigating the properties of massive dark GRB hosts is a key to reveal the origin of dark GRBs. The $Swift$ Burst Alert Telescope (BAT) detected the long-duration GRB 080325 (T$_{90} = 128.4 \\pm 34.2$ s) on 2008 March 25 at 04:10:32 UT with an initial localization of 3$\\arcmin$ radius \\citep{2008GCN..7512....1V}. The $Swift$ X-Ray Telescope (XRT) began observing the field at 151.9 s after the BAT trigger, and found a bright, uncatalogued X-ray source \\citep{2008GCN..7512....1V}. The enhanced XRT position was reported by \\cite{2008GCN..7513....1O} by using 199 s of overlapping XRT Photon Counting mode with an uncertainty of 2$\\arcsec$.6 radius. After the trigger, several attempts were made to identify the afterglow in optical and near-infrared wavelength \\citep{2008GCN..7516....1D,2008GCN..7518....1B, 2008GCN..7520....1C,2008GCN..7522....1K, 2008GCN..7529....1I,2008GCN..7563....1M}. However, no detection of the afterglow or host galaxy has been reported in optical and near-infrared wavelength, suggesting that the GRB 080325 is a dark GRB. In this paper we report the successful photometric identification of the afterglow of GRB 080325 in near-infrared wavelength and its host galaxy at $z \\sim 2$, and present a case study of properties of the afterglow and its host galaxy of a dark GRB. This paper is organized as follows: in \\S \\ref{NEAR-INFRARED AND OPTICAL IMAGING WITH SUBARU} the near-infrared and optical observations and data analysis are presented. In \\S \\ref{GRB 080325 AS A DARK GRB} we show the faintness of the optical and near-infrared afterglow of GRB 080325 compared with its X-ray afterglow, and dust extinction along the line of sight to GRB 080325. Results of spectral energy distribution fitting analysis of the GRB 080325 host and discussion are presented in \\S \\ref{MASSIVE RED HOST GALAXY AT} and \\S \\ref{COMPARISON WITH OTHER GRB AFTERGLOWS AND HOSTS}, respectively. Throughout this paper we use the AB magnitude system and the cosmological parameters of $H_{0}$ = 71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{M}$ = 0.27, and $\\Omega_{\\Lambda}$ = 0.73. ", "conclusions": "\\label{SUMMARY} We successfully detected an afterglow of GRB 080325 in $K_{s}$ band at 8.7 hours after the burst and its host galaxy in $K_{s}$ and $J$ band with Subaru/MOIRCS although no optical counterpart was reported. GRB 080325 is classified as a \\lq\\lq Dark GRB\\rq\\rq based on optical-to-X-ray spectral index of $\\beta_{\\rm OX} < 0.2-0.5$. The flux ratio of rest-wavelength optical to X-ray bands of the afterglow shows a very large dust extinction along the line of sight to the afterglow ($A_{V}$ = $2.7 - 10$ mag). To reveal host properties, we obtained optical images in $B$, $R_{c}$, $i$', and $z$' bands with Subaru/Suprime-Cam at about one year after the burst, and clearly detected the host galaxy in these bands. The SED fitting analysis for the host galaxy was performed assuming the Calzetti et al. extinction law and various star formation histories including instantaneous, exponentially declining ($e^{-\\frac{t}{\\tau}}$ with $\\tau$ = 1 Myr, 10 Myr, 100 Myr,1 Gyr, and 10 Gyr), and constant star formation rate. The best-fit stellar population synthesis model ($\\tau$ = 1 Gyr) indicates that the host is at $z_{photo} =$ 1.9. The dust extinction for the entire host ($A_{V} = 0.8$ mag) is larger than those of typical GRB hosts. Although GRB 080325 is located at the outskirt of the host galaxy, the dust extinction along the line of sight to GRB 080325 is larger than that for the entire host. Therefore, a major reason for the optical faintness of GRB 080325 relative to X-ray brightness is likely attributed to the local dusty environment around GRB 080325 rather than extinction by dust distributed over the entire host. We found that the host is very luminous (comparable to $L^{*}$ at $z \\sim 2.0$) and very massive ($M_{*} = 7.0 \\times 10^{10} M_{\\odot}$) in contrast to the faint and less massive properties of GRB hosts at lower redshifts. Considering the mass-metallicity relation for star-forming galaxies at $z \\sim 2$ and the stellar mass of the GRB 080325 host, the high metallicity environment around GRB 080325 is suggested. This possible high metallicity environment favors the binary merger scenario of the GRB phenomenon rather than the single star explosion, although the future spectroscopic observation at the GRB position is essential." }, "1003/1003.6132_arXiv.txt": { "abstract": "We use large-scale simulations to investigate the morphology of reionization during the final, overlap phase. Our method uses an efficient three-dimensional smoothing technique which takes into account the finite mean free path due to absorption systems, $\\mfp$, by only smoothing over scales $R_s<\\mfp$. The large dynamic range of our calculations is necessary to resolve the neutral patches left at the end of reionization within a representative volume; we find that simulation volumes exceeding several hundred Mpc on a side are necessary in order to properly model reionization when the neutral fraction is $\\simeq 0.01-0.3$. Our results indicate a strong dependence of percolation morphology on a large and uncertain region of model parameter space. The single most important parameter is the mean free path to absorption systems, which serve as opaque barriers to ionizing radiation. If these absorption systems were as abundant as some realistic estimates indicate, the spatial structure of the overlap phase is considerably more complex than previously predicted. In view of the lack of constraints on the mean free path at the highest redshifts, current theories that do not include absorption by Lyman-limit systems, and in particular three-dimensional simulations, may underestimate the abundance of neutral clouds at the end of reionization. This affects predictions for the 21 cm signal associated with reionization, interpretation of absorption features in quasar spectra at $z \\sim ~5-6$, the connection between reionization and the local universe, and constraints on the patchiness and duration of reionization from temperature fluctuations measured in the cosmic microwave background arising from the kinetic Sunyaev-Zel'dovich effect. ", "introduction": "On scales relevant to the large-scale structure of reionization -- hundreds of Mpc -- numerical calculations cannot self-consistently track all the sources and sinks of ionizing radiation. Resolving the IGM on the small scales relevant to QSO absorption systems optically-thick at the Lyman edge, also known as Lyman limit systems \\citep[e.g.,][]{weymann/etal:1981, tytler:1982, mo/miralda-escude:1996}, is particularly difficult. These absorption systems were likely to have impeded the progress of reionization \\citep[e.g.,][]{miralda-escude/etal:2000, gnedin/fan:2006}. What are likely values for the comoving absorption system mean free path at the end of reionization, $\\mfp$? \\citet{prochaska/etal:2009} determined the abundance of Lyman-limit absorption systems, $dN/dz = 1.9 [(1+z)/4.7]^{5.2}$, over the redshift range $3.6 ", "introduction": "When a potential exo-Earth is detected, the first thing we want to know is, ``Is this an Earth?'' The question has three parts: first ``Is the mass of the planet roughly the same as our Earth?''; second, ``Is the orbit of the planet in the habitable zone of the star?''; and third, ``What does the spectral characterization of the planet tell us?'' Direct imaging measures only the motion of the planet. Because the albedo, the phase, and surface scattering properties of the planet are unknown, planet mass cannot be determined from images. And while spectroscopic information may indicate whether the planet's atmosphere is qualitatively similar to that expected for a Neptune or a terrestrial planet, the planet's mass can only be quantitatively determined by astrometry -- measuring the motion of the star around the planet-star center of gravity. Measuring the orbit of a planet can be done by either astrometric or direct imaging missions. But because of the limitation of the IWA an exo-Earth is generally observable over a smaller fraction of its orbit, for a direct imaging mission. Figure~\\ref{fig:Figure1} shows the orbit of a planet (blue dot) around another star (represented by the yellow dot). The large blue circular region is the IWA and the yellow arcs are the parts of the orbit when the planet is observable. With a coronagraph whose IWA is only slightly smaller than the maximum star-planet separation, orbital parameters such as inclination can't be measured. If a planet is non-self-luminous and its surface is a Lambertian scatterer, then for a circular orbit, the planet's apparent brightness can vary by a factor of three from the ``full moon'' phase to the ``half moon'' phase (see Equation~\\ref{eq:Lambertian}). In multiple planet systems, two images with one planet in each image leaves open the possibility that they are two separate planets. If the IWA is substantially smaller, say half of the maximum star-planet separation, the planet becomes observable over most of its orbit. In this case it will be possible to look for seasonal variations in brightness. Such variations may occur because the surface is a non-Lambertian scatterer. Another source of seasonal variation may be a change in albedo. In the winter time, an ocean's surface may be covered with ice, which has a much higher albedo than liquid water. Seasonal changes in brightness are a double edged sword; they reveal important information about the surface of the planet, but their presence also complicates the identification of which dot is which planet and the determination of the planet's orbit around the star. The key to measuring the orbit of a planet is to have many images of the planet at different times of its year. But some types of coronagraphs are seriously limited in their ability to take images at many epochs. Recent studies have shown that an astrometric mission will have a completeness for detection of $\\geq95\\%$ for star-planet systems within 10 pc of the Sun \\cite{tra10}. On the other hand, the completeness for detection for a direct imaging mission, in the absence of data from an astrometric mission, has been shown to be $\\leq35\\%$ \\cite{sav10}. Given proof-of-existence, masses, and orbits from an astrometric mission, a direct imaging mission is free to plan and optimize its resources for the function for which it is uniquely-equipped, the spectral characterization of planetary atmospheres. With this information in hand, the exoplanet's resum\\'{e} --- existence, mass, orbit, and atmospheric spectra --- is complete. Only then can we answer the question, ``Is this an Earth?'' In this paper, we examine and quantify the effect of confusion on imaging surveys. We investigate the synergy between astrometry and direct imaging in measuring the orbit of a planet, and we examine how an astrometric orbit helps a direct imaging mission determine that a `pale blue dot' is an exo-Earth and not a false alarm. We consider only the case of circular orbits. We estimate the number of observations needed for orbit determination with imaging alone and with a combination of imaging and astrometry. Using simulations based on reasonable assumptions, we show that knowledge from a prior astrometric survey can greatly increase the science yield of an imaging mission. In sections 2 and 3 we consider two different situations. In section 2 the planet is first discovered by an astrometric mission and the role of imaging is to confirm the discovery and improve on its orbit determination. In section 3 we consider the measurement of an orbit from direct imaging alone. In section 4 we discuss how to realistically model planet populations. We simulate an example external occulter mission, O3, and show that the number of exo-Earths with confirmed orbits is significantly different with and without a prior astrometric mission. In section 5 we extend the discussion to internal coronagraphs and compute the time required to get orbits for a medium to large internal coronagraph mission. ", "conclusions": "In this paper we have simulated the planet-finding capability of the Occulting Ozone Observatory (O3), with and without a prior astrometric mission. O3 is an imaging mission which uses a 1.1 m telescope plus an external starshade occulter to measure orbits of exo-Earths. In our analysis, we used the same target list as in \\cite{kas10}. We simulated realistic planet distributions, which allow multiple-planet systems. This study is the first that we are aware of to account for the confusion that can arise among exo-Earths, super-Earths, and Neptune-class planets for an imaging mission. In the near future, RV and microlensing data will further constrain the Neptune and super-Earth populations, and Kepler data will reveal the statistics of exo-Earths, allowing us to improve the realism of our planet population models. In this work, we have assumed circular planet orbits, and we have assumed that 4 images are sufficient to fit an orbit. In future work, we will include eccentric orbits, and we will model the detection process by actually fitting the planet orbits. A planet's orbit can be measured by space astrometry or by multiple direct images, or a combination of the two. Prior astrometry helps in two important ways. If an exo-Earth has been discovered and its orbit measured astrometrically, the orbital phase extrapolated to 5 years in the future would have a sizable error bar. However, as we have shown in this paper, given an orbital solution from a prior astrometric mission, the orbital phase can be pinned down very accurately by a single follow-up image. Without this knowledge, a coronagraph would need a much larger number of images to discover the planet and measure the orbit (see Figure ~\\ref{fig:Figure5}). Secondly, knowing beforehand which stars do not have exo-Earths would save the many hours of coronagraph search time at these stars that would be necessary to make such a determination. This is potentially a large fraction of mission time, which could become available for spectroscopy of the exo-Earths, or for more orbit measurements. With information from prior astrometry, the coronagraph gains efficiency; it could survey more targets for exo-Earths in the same mission time. The impact of astrometry is most pronounced if the follow-on direct detection mission is an external occulter. Small occulters might be limited to $\\sim60$ pointings, large occulters to $\\sim120$ pointings. A small occulter (e.g. the O3 mission) as a follow-on to an astrometric mission could detect $\\sim6$ exo-Earths, if $\\eta_{\\oplus} = $10\\%. By itself, this small external occulter would detect just $\\sim1.5$ exo-Earths, with a sizeable probability that zero are detected. Small internal coronagraphs that only have $\\sim10$ potential targets would spend a large fraction of their mission visiting each of these targets 10 to 15 times each. For internal coronagraphs, there is no propellant limitation and if one builds a large enough telescope, there will be sufficient integration time to survey 60 to 100 nearby stars. When the target list extends to 100 stars, a survey of all 100 stars with 4 images of the planet at a contrast of 4$\\times10^{-11}$ would take 40 years. But knowledge of which stars have Earths and which do not could reduce the integration time by roughly a factor of 20. The factor of 20 corresponds to slightly more than a factor of 2 in the diameter of the coronagraphic telescope. When it comes time to build a large internal coronagraph, knowing where the exo-Earths are will make it possible to design a mission just large enough to do what is scientifically compelling, rather than a mission that is much more expensive, to insure against bad luck." }, "1003/1003.2594_arXiv.txt": { "abstract": "We present high-resolution 3D smoothed particle hydrodynamics simulations of the formation and evolution of protostellar discs in a turbulent molecular cloud. Using a piecewise polytropic equation of state, we perform two sets of simulations. In both cases we find that isolated systems undergo a fundamentally different evolution than members of binary or multiple systems. When formed, isolated systems must accrete mass and increase their specific angular momentum, leading to the formation of massive, extended discs, which undergo strong gravitational instabilities and are susceptible to disc fragmentation. Fragments with initial masses of $\\mcla$, $\\mclb$ and $\\altmcl$ are produced in our simulations. In binaries and small clusters, we observe that due to competition for material from the parent core, members do not accrete significant amounts of high specific angular momentum gas relative to isolated systems. We find that discs in multiple systems are strongly self-gravitating but that they are stable against fragmentation due to disc truncation and mass profile steeping by tides, accretion of high specific angular momentum gas by other members, and angular momentum being redirected into members' orbits. In general, we expect disc fragmentation to be less likely in clusters and to be more a feature of isolated systems. ", "introduction": "The overall paradigm under which star formation occurs remains under debate. One hypothesis is that stars are thought to form as a collection of low mass fragments in collapsing clumps of gas, and then undergo {\\em competitive accretion} as the fragments try to accrete gas from their common reservoir (\\citealt{bonnellbateclarke01}; \\citealt{bonnellbate06}). The main competing hypothesis is {\\em gravitational collapse}, where massive star-forming clumps collapse and form multiple cores. Each star forms from the gas that is available in its own core, with limited accretion from other material in the parent clump \\citep{krumholzmckeeklein05}. Regardless of the star formation paradigm, in the case of solar-type stars, the end result is most likely membership in a binary or multiple system (\\citealt{duquennoymayor91}; \\citealt{eggenbergerudrymayor04}). Scenarios for the formation of such systems abound. Amongst the simplest, most idealised scenarios are the fission of a bar-unstable core (\\citealt{durisengingoldtohlineboss86}; \\citealt{burkertbatebodenheimer97}) and the fragmentation of centrally condensed, rotating, magnetised cores \\citep{boss97}. More complicated hypotheses appeal to the chaos of the cloud environment, such as core-core collisions \\citep{turneretalcap95}, protostellar encounters \\citep{shenetal10}, dynamical capture in unstable multiple systems \\citep{batebonnellbrommbin02}, and accretion-triggered fragmentation (\\citealt{bonnell94}; \\citealt{whitworthetalrot95}; \\citealt{hennebelleetalcomp04}; \\citealt{offnerklein08}; \\citealt{kratter_2010}). Star formation scenarios involving the rapid collapse of a protostellar core following the loss of support against gravity require the formation of a massive accretion disc around the central object (e.g. \\citealt{vorobyovbasu07} and \\citealt{walch09}). Such discs undergo a short-lived $\\sim 0.1\\rm\\,Myr$ stage where the disc is massive relative to its host ($0.1 < M_d/M_{\\star} < 1$) and where gravitational instabilities operate to transport mass through the disc onto the central protostar \\citep{vorobyovbasu07}. Massive, accreting, and extended protostellar discs have been shown to be susceptible to fragmentation, which could be responsible for a range of phenomena such as FU Orionis events and early dust processing if the clumps are disrupted (\\citealt{boleyetal10}; \\citealt{nayakshin2010}), and the formation of substellar companions otherwise (\\citealt{vorobyovbasu07} and \\citealt{boleytwomodes09}). Given the problem that both the standard core accretion planet formation timescale and protostellar disc lifetimes are typically a few Myr \\citep{haischladalada01}, the idea of creating giant planets in a few orbital times via gravitational instabilities was revived \\citep{bosssci97} and has since been the subject of sustained interest (\\citealt{boss02,boss08}; \\citealt{mayeretal04,mayeretal07}; \\citealt{pickettetal00}; \\citealt{pickett07}; \\citealt{boleyetal06}; \\citealt{boleytwomodes09}) . Analytical works constraining inner disc fragmentation \\citep{rafikov05,rafikov07}, and the short cooling times required to form long-lived clumps within $\\sim 10\\rm\\,AU$ in simulations of protoplanetary discs, along with observations of massive planets on wide orbits (e.g. \\citealt{maroisetal08,maroisetal2010}), have lead to a shift in focus to outer disc ($> 40\\rm\\,AU$) fragmentation (e.g. \\citealt{stamatellosfrag07}; \\citealt{boleytwomodes09}; \\citealt{dodson2009}; \\citealt{vorobyovbasu2010}). Whether the extended disc models used in these studies are similar to discs formed from collapsed molecular cloud cores remains to be seen. In light of the likely connection between outer disc fragmentation and early protostellar systems, we examine in this paper the formation of discs in detail using 3D SPH simulations of core collapse in turbulent molecular clouds. In particular we compare the early evolution in mass, surface density, specific angular momentum, and disc stability between several different systems under near-identical conditions. In one set of simulations we compare an isolated and a binary system, and in the other we compare an isolated system and a small cluster. Previous studies of gravitational instabilities in binary systems have yielded mixed results, with some finding that the perturbing companion hinders fragmentation through disc truncation and tidal heating (\\citealt{nelson00}; \\citealt{mwqs05}) and alternatively promotes fragmentation, also through tidal perturbations \\citep{boss06}. However, the evolved systems considered in these studies may not be as susceptible to fragmentation as their protobinary counterparts, which we investigate here, due to the enhanced importance of gravitational instabilities during protostellar disc formation. The paper is laid out as follows: in Section 2 we discuss the simulations, initial conditions and the clump identification procedure. In Section 3 we present the results and analysis. Further discussion is provided in Section 4, and our conclusions are laid out in Section 5. ", "conclusions": "The primary motivation of this paper was to compare the evolution of an isolated and binary/multiple systems during the early stages of prestellar development within the same environment. Of particular interest was the potential for fragmentation of such early, accreting and self-gravitating discs. The following are our findings: \\begin{itemize} \\item The initial collapse occurs within filaments that tend to collimate infalling gas onto the central prestellar objects, whose specific angular momenta tend to remain aligned to that of the environment (and perpendicular to the embedding filament) on scales of $\\sim 1000\\rm\\,AU$. \\item During formation, isolated systems must accrete mass and increase their specific angular momentum, leading to the formation of massive, extended discs, which undergo strong gravitational instabilities and are susceptible to fragmentation \\item In systems starting out as binaries and multiples, an effective specific angular momentum ceiling exists, limiting the maximum $j$ of the systems but not their mass, making disc fragmentation unlikely. The ceiling is the product of tidal interactions which strip cold gas from members, and concentrate their mass profiles, and the redirection of angular momentum into orbits. \\end{itemize}" }, "1003/1003.0905_arXiv.txt": { "abstract": "We calculate a signature of cosmic strings in the polarization of the cosmic microwave background (CMB). We find that ionization in the wakes behind moving strings gives rise to extra polarization in a set of rectangular patches in the sky whose length distribution is scale-invariant. The length of an individual patch is set by the co-moving Hubble radius at the time the string is perturbing the CMB. The polarization signal is largest for string wakes produced at the earliest post-recombination time, and for an alignment in which the photons cross the wake close to the time the wake is created. The maximal amplitude of the polarization relative to the temperature quadrupole is set by the overdensity of free electrons inside a wake which depends on the ionization fraction $f$ inside the wake. The signal can be as high as $0.06 {\\rm \\mu K}$ in degree scale polarization for a string at high redshift (near recombination) and a string tension $\\mu$ given by $G \\mu = 10^{-7}$. ", "introduction": "In recent years there has been renewed interest in the possibility that cosmic strings contribute to the power spectrum of curvature fluctuations which give rise to the large scale structure and cosmic microwave background (CMB) anisotropies which we see today. One of the reasons is that many inflationary models constructed in the context of supergravity models lead to the formation of gauge theory cosmic strings at the end of the inflationary phase \\cite{Rachel}. Secondly, in a large class of brane inflation models the formation of cosmic superstrings \\cite{Witten} at the end of inflation is generic \\cite{CS-BI}, and in some cases (see \\cite{Pol1}) these strings are stable (see also \\cite{recentCS} for reviews on fundamental cosmic strings). Cosmic superstrings are also a possible remnant of an early Hagedorn phase of string gas cosmology \\cite{SGrev} In models which admit stable strings or superstrings, a scaling solution of such strings inevitably \\cite{Kibble} results as a consequence of cosmological dynamics (see e.g. \\cite{CSrevs} for reviews on cosmic strings and structure formation). In a scaling solution, the network of strings looks statistically the same at any time $t$ if lengths are scaled to the Hubble radius at that time. The distribution of strings is dominated by a network of ``infinite\" strings \\footnote{Any string with a mean curvature radius comparable or greater than the Hubble radius or which extends beyond the Hubble radius is called ``infinite\" or ``long\".} with mean curvature radius and separation being of the order of the Hubble radius. The scaling solution of infinite strings is maintained by the production of string loops due to the interaction of long strings. This leads to a distribution of string loops with a well defined spectrum (see e.g. \\cite{TuBr}) for all radii $R$ smaller than a cutoff radius set by the Hubble length. Whereas the scaling distribution of the infinite strings is reasonably well known as a result of detailed numerical simulations of cosmic string evolution (see \\cite{CSsimuls} for some references), there is still substantial uncertainty concerning the distribution of string loops. It is, however, quite clear that the long strings dominate the energy density of strings. The distinctive signals of strings which we will focus on are due to the long strings. Cosmic strings give rise to distinctive signatures in both the CMB and in the large-scale structure. These signatures are a consequence of the specific geometry of space produced by strings. As studied initially in \\cite{Vil}, space perpendicular to a long straight string is locally flat but globally looks like a cone whose tip coincides with the location of the string (the smoothing out of the cone as a consequence of the internal structure of the cosmic string was worked out in \\cite{Ruth}). The deficit angle is given by \\be \\label{deficit} \\alpha \\, = \\, 8 \\pi G \\mu \\, . \\ee where $\\mu$ is the string tension and $G$ is Newton's constant. Hence, a cosmic string moving with velocity $v$ in the plane perpendicular to its tangent vector will lead to line discontinuities in the CMB temperature of photons passing on different sides of the string. The magnitude of the temperature jump is \\cite{KS} \\be \\label{KSsig} {{\\delta T} \\over T} \\, = \\, 8 \\pi \\gamma(v) v G \\mu \\, , \\ee where $\\gamma$ is the relativistic gamma factor associated with the velocity $v$. As a consequence of the deficit angle (\\ref{deficit}), a moving string will generate a cosmic string wake, a wedge-shaped region behind the string (from the point of view of its velocity), a region with twice the background density \\cite{wake} (see Figure 1). Causality (see e.g. \\cite{Traschen}) limits the depth of the distortion of space due to a cosmic string. The details were worked out in \\cite{Joao} where it was shown that the deficit angle goes to zero quite rapidly a distance $t$ from the string. Hence, the depth of the string wake is given by the same length. String wakes lead to distinctive signatures in the topology of the large-scale structure, signatures which were explored e.g. in \\cite{topology}. \\begin{figure} \\includegraphics[height=5cm]{wakewedge.eps} \\caption{Sketch of the mechanism by which a wake behind a moving string is generated. Consider a string perpendicular to the plane of the graph moving straight downward. From the point of view of the frame in which the string is at rest, matter is moving upwards, as indicated with the arrows in the left panel. From the point of view of an observer sitting behind the string (relative to the string motion) matter flowing past the string receives a velocity kick towards the plane determined by the direction of the string and the velocity vector (right panel). This velocity kick towards the plane leads to a wedge-shaped region behind the string with twice the background density (the shaded region in the right panel).} \\label{fig:1} \\end{figure} Wakes formed at arbitrarily early times are non-linear density perturbations. For wakes formed by strings present at times $t_i > t_{rec}$, where $t_{rec}$ is the time of recombination, the baryonic matter inside the wake undergoes shocks \\cite{Rees} (see e.g. \\cite{Sornborger} for a detailed study). The shocks, in turn, can ionize the gas - although as it turns out the residual ionization from decoupling is larger. Photons passing through these ionized regions on their way from the last scattering surface to the observer can thus be polarized - and it is this polarization signature which we aim to study here. The tightest constraints on the contribution of scaling strings to structure formation (and thus the tightest upper bound on the tension $\\mu$ of the strings) comes from the analysis of the angular power spectrum of CMB anisotropies. As discussed in \\cite{Perivo,Albrecht,Turok1}, the angular power spectrum does not have the acoustic ringing which inflation-seeded perturbations generate. The reason is that the string network is continuously seeding the growing mode of the curvature fluctuation variable on super-Hubble scales. Hence, the fluctuations are ``incoherent\" and ``active\" as opposed to ``coherent\" and ``passive\" as in the case of inflation-generated fluctuations. The contribution of cosmic strings to the primordial power spectrum of cosmological perturbations is thus bounded from above, thus leading to an upper bound \\cite{Wyman1,Fraisse,Slosar1,Bevis,Battye} on the string tension of about $G \\mu < 3 \\times 10^{-7}$. Past work on CMB temperature maps has shown that signatures of cosmic strings are easier to identify in position space than in Fourier space \\cite{KS, Moessner, Lo, Smoot} and recent studies show that high angular resolution surveys such as the South Pole Telescope project \\cite{SPT} have the potential of improving the limits on the string tension by an order of magnitude \\cite{Amsel, Stewart, Rebecca}. To date there has been little work on CMB polarization due to strings. Most of the existing work focuses on the angular power spectra of the polarization. Based on a formalism \\cite{Turok1} (see also \\cite{Andy}) to include cosmic defects as source terms in the Boltzmann equations used in CMB codes, the power spectra of temperature and temperature polarization maps were worked out \\cite{Turok2} in the case of models with global defects such as global cosmic strings. Since in cosmic defect models vector and tensor modes are as important as the scalar metric fluctuations \\cite{Turok1,Durrer1}, a significant B-mode polarization is induced. In fact, in the case of global strings with a tension close to the upper bound mentioned above, whereas the amplitude of the temperature and E-mode polarization power spectra are so small as to make the string signal invisible compared to the signal from the scale-invariant spectrum of adiabatic fluctuations (e.g. produced in inflation), the contribution of strings dominates the amplitude of the B-mode polarization power spectrum. In the case of local strings, these conclusions were confirmed in the more recent analyses of \\cite{Wyman1,Slosar2,Battye,Bevis,Wyman2}. The maximal amplitude of the B-mode polarization power spectrum for strings with $G \\mu = 3 \\times 10^{-7}$ was shown to be taken on at angular harmonic values of $l \\sim 500$ and to be of the order $0.3 \\mu K^2$ \\cite{Wyman2}. However, the analyses of \\cite{Wyman1,Slosar2,Battye,Bevis,Wyman2} do not take into account the effects of the gravitational accretion onto cosmic string wakes. In related work, the conversion of E-mode to B-mode polarization via the gravitational lensing induced by cosmic strings was studied in \\cite{Francis} \\footnote{While this paper was being finalized for submission, a preprint appeared \\cite{Durrer2} computing the local B-mode polarization power spectrum from cosmic strings.}. Similarly to what was found in the analysis of CMB temperature maps from cosmic strings, we expect that a position-space analysis will be more powerful at revealing the key non-Gaussian signatures of strings in CMB polarization. Hence, in this work we derive the position space signature of a cosmic string wake in CMB polarization maps. In the following section we briefly review cosmic string wakes. In Section 3 we then analyze the polarization signature of wakes, and we conclude with some discussion. ", "conclusions": "In this Letter we have discussed a position space signal of a cosmic string wake in CMB polarization maps. In the same way that the line discontinuities in the CMB temperature maps predicted by the Kaiser-Stebbins (KS) effect yield a promising way to constrain/detect cosmic strings in the CMB (see e.g. \\cite{Moessner,Lo,Smoot,Amsel,Stewart,Rebecca}, we believe that the signal discussed in this paper will play a similar role once CMB polarization maps become available. We have shown that a single wake will produce a rectangular patch in the sky of dimensions given by equation (\\ref{size2}), average magnitude given by equation (\\ref{result}) and amplitude increasing linearly in one direction across the patch. For a value of the string tension $G \\mu = 3 \\times 10^{-7}$ (the current upper limit), the amplitude of the signal is within the range of planned polarization experiments for wakes produced sufficiently close to the surface of recombination. These wakes are also the most numerous ones. The brightest edge in the polarization map corresponds to the beginning location of the string, not the final location. Since the KS discontinuity in the CMB temperature map will occur along the line corresponding to the final position, the polarization signal discussed here provides a cross-check on a possible string interpretation of a KS signal. A scaling distribution of strings will yield a distribution of patches in the sky, the most numerous ones and the ones with the largest polarization amplitude being set by wakes laid down at times close to the time of recombination which are crossed by the CMB photons at similarly early times." }, "1003/1003.2307_arXiv.txt": { "abstract": "The fidelity of radio astronomical images is generally assessed by practical experience, i.e. using rules of thumb, although some aspects and cases have been treated rigorously. In this paper we present a mathematical framework capable of describing the fundamental limits of radio astronomical imaging problems. Although the data model assumes a single snapshot observation, i.e. variations in time and frequency are not considered, this framework is sufficiently general to allow extension to synthesis observations. Using tools from statistical signal processing and linear algebra, we discuss the tractability of the imaging and deconvolution problem, the redistribution of noise in the map by the imaging and deconvolution process, the covariance of the image values due to propagation of calibration errors and thermal noise and the upper limit on the number of sources tractable by self calibration. The combination of covariance of the image values and the number of tractable sources determines the effective noise floor achievable in the imaging process. The effective noise provides a better figure of merit than dynamic range since it includes the spatial variations of the noise. Our results provide handles for improving the imaging performance by design of the array. ", "introduction": "The radio astronomical community is currently building or developing a number of new instruments such as the low frequency array (LOFAR) \\cite{Bregman2005-1}, the square kilometer array (SKA) \\cite{Hall2004-1} and the Mileura wide field array (MWA) \\cite{Lonsdale2005-1}. Imaging and self calibration of these radio telescopes will be computationally demanding tasks due to the large number of array elements. Much research is therefore focused on finding clever short-cuts to reduce the amount of processing required, such as $w$-projection \\cite{Cornwell2005-1} or facet imaging \\cite{Perley1994-1} and different variants of CLEAN \\cite{Cornwell1999-1}. The validity and quality of these methods is generally assessed by practical experience. Attempts to do a rigorous analysis are done for some aspects and cases \\cite{Schwarz1978-1, Tan1986-1, Kulkarni1989-1, Wijnholds2006-1}, but most of the time rules of thumb are used. This paper presents the first comprehensive mathematical framework capable of describing the fundamental limits of radio astronomical imaging problems. The data model used in this paper applies to snapshot observations, i.e. variations in time and frequency are not considered. However, using a multi-measurement data model such as those in \\cite{Leshem2000-1, Veen2004-1, Tol2007-1} it is straightforward to extend the data model to synthesis observation and still apply the framework described herein. The resolution of the final image (or {\\em map}) is normally determined by the size and configuration of the array and the spatial taper function. Under the assumption that the sky is mainly empty, i.e.\\ that the image contains only a few sources, maps with higher resolution than predicted by the array configuration (superresolution) can be made using CLEAN. The Maximum Entropy Method (MEM) \\cite{Narayan1986-1} imposes a similar constraint by aiming for a solution that is as featureless as possible. In the array processing literature, superresolution is achieved by high resolution direction of arrival (DOA) estimation techniques such as MUSIC \\cite{Schmidt1986-1} and weighted subspace fitting \\cite{Viberg1991-1, Viberg1991-2}. In all these approaches the goal is to disentangle the spatial response of the array and the source structure, a process called deconvolution. In section \\ref{sec:imaging} we formulate imaging as an estimation problem, an approach called model based imaging, and obtain an analytic expression for its least squares solution that allows us to formulate the deconvolution problem as a matrix inversion problem. This provides a powerful tool to assess the tractability of the deconvolution problem and to demonstrate the impact on the array configuration on the deconvolution problem and the redistribution of noise in the imaging and deconvolution process. The dynamic range of an image is generally defined as the power ratio between the strongest and the weakest meaningful features in the map. In practice, the limitations of an instrument are more conveniently described by the achievable noise floor in an imaging observation since the dynamic range strongly depends on the strength of the strongest source within the field-of-view and because the noise varies over the map. This noise floor is a combination of calibration errors, thermal noise and confusion noise. In this paper the term ``effective noise'' refers to the net result of these constituents in the image plane. In section \\ref{sec:effnoise} analytical expressions are derived that describe the components of the effective noise in terms of the covariance of the image values, a concept which we will refer to as image covariance. The consequences of these expressions are illustrated with a few examples in section \\ref{sec:implications}. These examples suggest that the contribution of propagated calibration errors to the image covariance is considerably smaller than the contribution of thermal noise even if the calibration is done on data with similar SNR. They also indicate that self calibration causes higher covariance between source power estimates than pure imaging does. \\emph{Notation}: Overbar $\\overline{(\\cdot)}$ denotes complex conjugation. The transpose operator is denoted by $^T$, the complex conjugate (Hermitian) transpose by $^H$ and the Moore-Penrose pseudo-inverse by $^\\dagger$. The expectation operator is denoted by $\\expect \\{ \\cdot \\}$, $\\odot$ is the element-wise matrix multiplication (Hadamard product), $(\\cdot)^{\\odot n}$ is used to denote the element-wise matrix exponent with exponent $n$, $\\otimes$ denotes the Kronecker product and $\\circ$ is used to denote the Khatri-Rao or column-wise Kronecker product of two matrices. $\\diag(\\cdot)$ converts a vector to a diagonal matrix with the elements of the vector placed on the main diagonal, $\\vec(\\cdot)$ converts a matrix to a vector by stacking the columns of the matrix and $\\vecdiag(\\cdot)$ converts the main diagonal of its argument to a column vector. $\\mathrm{circulant}(\\cdot)$ creates a square circulant matrix by circularly shifting the entries of its vector argument to form its columns. $\\circledast$ will be used to denote circular convolution of two vectors, i.e., for vectors of length $n$, $(\\mathbf{x} \\circledast \\mathbf{y})_j = \\sum_{i=0}^{n-1} x_i y_{j-i\\,\\mathrm{mod}\\, n}$. For matrices and vectors of compatible dimensions, we will frequently use the following properties: \\begin{eqnarray} \\vec(\\mathbf{A} \\mathbf{B} \\mathbf{C}) &=& (\\mathbf{C}^T \\otimes \\mathbf{A}) \\vec(\\mathbf{B}) \\label{eq:prop1} \\\\ \\vec(\\mathbf{A} \\diag(\\mathbf{b}) \\mathbf{C}) &=& (\\mathbf{C}^T \\circ \\mathbf{A}) \\mathbf{b} \\label{eq:prop2} \\\\ (\\mathbf{A} \\circ \\mathbf{B})^H (\\mathbf{C} \\circ \\mathbf{D}) &=& \\mathbf{A}^H \\mathbf{C} \\odot \\mathbf{B}^H \\mathbf{D} \\label{eq:prop3} \\\\ (\\mathbf{A} \\otimes \\mathbf{B}) (\\mathbf{C} \\circ \\mathbf{D}) &=& \\mathbf{A} \\mathbf{C} \\circ \\mathbf{B} \\mathbf{D} \\label{eq:prop4} \\end{eqnarray} ", "conclusions": "In this paper we presented an analytic solution for snapshot imaging including deconvolution based on a data model (measurement equation) for the antenna signal covariance matrix or visibilities. The presented comprehensive framework is sufficiently flexible to enable extension of this analysis to synthesis observations, since the data model for a synthesis observation has the same form \\cite{Leshem2000-1, Veen2004-1, Tol2007-1}. This framework allowed us to make the first complete rigorous assessment of the effective noise floor, which is the combined effect of propagated calibration errors, thermal noise and source confusion, in the image in terms of the covariance of the image values. Our simulations for a 2D uniform rectangular array indicate that the effect of propagated calibration errors is strongly concentrated at the source locations but is considerably smaller than the thermal noise at other image points. The results also suggest that if the instrument is sufficiently stable, a separate calibration step is to be preferred over a self calibrated image since it allows better source separation in the imaging process. The effects of deconvolution can be described by a deconvolution matrix that describes the amount of linear independence (orthogonality) of the spatial signature vectors weighted by the actual gains of the receiving elements. A diagonal deconvolution matrix not only ensures the best possible spatial separation between the sources, but also ensures a homogeneous noise distribution over the map. This poses the question whether this matrix can be diagonalized by array design or by applying appropriate weights to the array elements. Since this factor is related to the array beam pattern, the latter is equivalent to finding weights that suppress the side lobe patterns at least in the direction of other sources, which suggest that techniques like Robust Capon Beamforming should provide the requested weighting \\cite{Tol2005-1}. The condition number of the deconvolution matrix can be used to assess the quality of the solution to the deconvolution problem. Compared to a redundant array (ULA, URA), an array without redundant element spacings provides much better possibilities to approach the maximum number of solvable image points for a fixed number of antenna elements, thereby allowing the system to reach the theoretical self calibration confusion limit." }, "1003/1003.1132_arXiv.txt": { "abstract": "We present a new $N$-body model for the tidal disruption of the Sagittarius (Sgr) dwarf that is capable of simultaneously satisfying the majority of angular position, distance, and radial velocity constraints imposed by current wide-field surveys of its dynamically young ($\\lesssim 3$ Gyr) tidal debris streams. In particular, this model resolves the conflicting angular position and radial velocity constraints on the Sgr leading tidal stream that have been highlighted in recent years. While the model does not reproduce the apparent bifurcation observed in the leading debris stream, recent observational data suggest that this bifurcation may represent a constraint on the internal properties of the Sgr dwarf rather than the details of its orbit. The key element in the success of this model is the introduction of a non-axisymmetric component to the Galactic gravitational potential which can be described in terms of a triaxial dark matter halo whose minor/major axis ratio $(c/a)_{\\Phi} = 0.72$ and intermediate/major axis ratio $(b/a)_{\\Phi} = 0.99$ at radii $20 < r < 60$ kpc. The minor/intermediate/major axes of this halo lie along the directions $(l, b) = (7^{\\circ}, 0^{\\circ})$, $(0^{\\circ}, 90^{\\circ})$, and $(97^{\\circ}, 0^{\\circ})$ respectively, corresponding to a nearly-oblate ellipsoid whose minor axis is contained within the Galactic disk plane. This particular disk/halo orientation is difficult to reconcile within the general context of galactic dynamics (and CDM models in particular), suggesting either that the orientation may have evolved significantly with time or that inclusion of other non-axisymmetric components (such as the gravitational influence of the Magellanic Clouds) in the model may obviate the need for triaxiality in the dark matter halo. The apparent proper motion of Sgr in this model is estimated to be $(\\mu_l \\textrm{cos}\\, b, \\mu_b) = (-2.16, 1.73)$ mas yr$^{-1}$, corresponding to a Galactocentric space velocity $(U,V,W$) = (230, -35, 195) \\kms. Based on the velocity dispersion in the stellar tidal streams, we estimate that Sgr has a current bound mass $M_{\\rm Sgr} = 2.5^{+1.3}_{-1.0} \\times 10^{8} M_{\\odot}$. We demonstrate that with simple assumptions about the star formation history of Sgr, tidal stripping models naturally give rise to gradients in the metallicity distribution function (MDF) along the stellar debris streams similar to those observed in recent studies. These models predict a strong evolution in the MDF of the model Sgr dwarf with time, indicating that the chemical abundances of stars in Sgr at the present day may be significantly different than the abundances of those already contributed to the Galactic stellar halo. We conclude by using the new $N$-body model to reevaluate previous claims of the association of miscellaneous halo substructure with the Sgr dwarf. ", "introduction": "With the advent of progressively deeper photometric surveys, it has become apparent that galactic haloes are threaded with the phase-mixed detritus of multiple generations of dwarf satellites that have been destroyed by the inexorable tides of their host's gravitational potential. Surveys have indicated the presence of such streams in both the Milky Way (e.g., Majewski et al. 1996; Helmi et al. 1999; Newberg et al. 2002; Rocha-Pinto et al. 2004; Belokurov et al. 2006; Duffau et al. 2006; Grillmair \\& Dionatos 2006; Starkenburg et al. 2009) and other nearby galaxies (e.g., Shang et al. 1998; Ibata et al. 2001a; McConnachie et al. 2003; Kalirai et al. 2006; Mart{\\'{\\i}}nez-Delgado et al. 2008, 2009), possibly constituting the primary source of the Galactic stellar halo (e.g., Searle \\& Zinn 1978; Majewski 1993, 2004; Majewski et al. 1996; Bullock et al. 2001b; Bullock \\& Johnston 2005; Font et al. 2006; Bell et al. 2008). Perhaps the most prominent such stream is that from the Sagittarius dwarf spheroidal galaxy (hereafter Sgr dSph), whose lengthy tidal streams wrap entirely around the Milky Way. The first observations of Sgr were presented by Ibata et al. (1994), since which time significant observational effort has been invested in detecting and characterizing the stellar streams emanating from the dwarf. Observations of high-latitude carbon stars (Totten \\& Irwin 1998) and various small-field surveys (e.g., Dinescu et al. 2002, and references therein) fueled a host of early efforts to model the Sgr --- Milky Way system by Johnston et al. (1995, 1999), Velazquez \\& White (1995), Edelsohn \\& Elmegreen (1997), Ibata et al. (1997), G{\\'o}mez-Flechoso et al. (1999), Helmi \\& White (2001), Mart{\\'{\\i}}nez-Delgado et al. (2004). In the past seven years however, our understanding of the scope and significance of the streams has been revolutionized by the deep, wide-field views provided by the Two Micron All-Sky Survey (2MASS) and Sloan Digital Sky Survey (SDSS). In particular, the 2MASS survey showed a large population of young, relatively metal-rich Sgr M-giants wrapping $360^{\\circ}$ or more across the sky (Majewski et al. 2003). SDSS observations have shown (Belokurov et al. 2006) that the debris streamer leading Sgr along its orbit continues to be well-defined through the North Galactic Pole as it passes over the solar neighborhood towards the Galactic anticenter. Follow-up spectroscopy (Majewski et al. 2004; Law et al. 2004; Yanny et al. 2009) has confirmed that the streams are similarly well-defined in radial velocity, and indicate a significant evolution in the metallicity distribution function (MDF) of the component stellar populations with increasing separation from the Sgr dwarf core along the tidal streams (e.g., Bellazini et al. 2006; Chou et al. 2007, 2009; Monaco et al. 2007; Starkenburg et al. 2009; Carlin et al. {\\it in prep.}). In the inner Galaxy, dynamical tracers such as tidal streams can provide strong constraints on the mass distribution of the baryonic components of the Milky Way (e.g., Font et al. 2001; Ibata et al. 2003; Koposov et al. 2009). At the larger radii $r \\sim 15-60$ kpc traced by the orbit of the Sgr dwarf however the gravitational potential is dominated by dark matter, and observations of the luminous tidal streams can be used to constrain the shape, orientation, and mass of the Milky Way's dark halo (e.g., Ibata et al. 2001b; Helmi 2004; Johnston et al. 2005; Law et al. 2005 [hereafter LJM05]; Fellhauer et al. 2006; Mart{\\'{\\i}}nez-Delgado et al. 2007). The breadth and quality of the recent observational data is such that it has highlighted a `halo conundrum' (variously discussed by LJM05; Fellhauer et al. 2006; Mart{\\'{\\i}}nez-Delgado et al. 2007; Newberg et al. 2007; and Yanny et al. 2009): In a static axisymmetric Milky Way dark halo no one model has been capable of simultaneously reproducing both the angular position and distances/radial velocities of tidal debris in the Sgr leading arm. Depending on which criterion a given study chose to focus, claims have been made in favor of an oblate halo (e.g., Johnston et al. 2005; Mart{\\'{\\i}}nez-Delgado et al. 2007), an approximately spherical halo (e.g., Ibata et al. 2001b; Fellhauer et al. 2006), and a prolate halo (e.g., Helmi 2004). As we demonstrated using massless test-particle orbits in Law, Majewski, \\& Johnston (2009; hereafter LMJ09) however, it is possible to resolve this conundrum by adopting a triaxial halo in which the minor axis is approximately aligned with the line of sight to the Galactic Center. In this contribution we extend the results of LMJ09 by performing detailed $N$-body simulations of the tidal disruption of the Sgr dwarf within a range of Galactic halo parameterizations that were favored by test-particle orbits. In \\S \\ref{basicapproach.sec} we outline our basic numerical approach, describing the generalized form of the Milky Way --- Sgr system along with an overview of the observational constraints on the Sgr tidal streams. We next describe our solution to the problem of constraining the basic orbital parameters of the Sgr stream (for a given Milky Way model) in \\S \\ref{sgrprop.sec}, using these various simulations to discriminate between different realizations of the underlying Galactic potential in \\S \\ref{mwprop.sec}. The final parameters describing our best-fit model of the Milky Way --- Sgr system are summarized in \\S \\ref{overview.sec}. In \\S \\ref{halodisc.sec} we discuss the best-fit model for the dark matter halo of the Milky Way that was derived in \\S \\ref{mwprop.sec}, and present possible theoretical challenges to, alternatives for, and methods of testing the triaxial halo model. We demonstrate in \\S \\ref{feh.sec} that by adopting simple assumptions for the star formation history of Sgr it is possible to reproduce the trends in the MDF observed along the stellar tidal streams. We compare the best-fitting model of the Sgr streams to other observations of halo substructure that have previously been postulated to be associated with the Sgr stream in \\S \\ref{otherobs.sec}, and conclude in \\S \\ref{summary.sec} with a summary of the successes and failures of our new model of the Milky Way --- Sgr system and an overview of the key issues that need to be addressed in future work. Throughout this paper we adopt the heliocentric Sgr coordinate system ($\\Lambda_{\\odot}, B_{\\odot}$) defined by Majewski et al. (2003) in which the longitudinal coordinate $\\Lambda_{\\odot} = 0^{\\circ}$ in the direction of Sgr and increases along trailing tidal debris, and $B_{\\odot}$ is positive towards the orbital pole $(l,b)_{\\rm pole} \\approx (274^{\\circ}, -14^{\\circ})$. ", "conclusions": "\\label{halodisc.sec} Cold dark matter (CDM) models of galaxy formation generally predict that dark haloes should be mildly triaxial (e.g., Bullock 2002; Jing \\& Suto 2002; Bailin \\& Steinmetz 2005; Allgood et al. 2006; and references therein), predominantly prolate systems ($T_{\\Phi} = \\frac{1 - (b^2/a^2)_{\\Phi}}{1 - (c^2/a^2)_{\\Phi}} \\gtrsim 0.66$; Allgood et al. 2006) with characteristic central axis ratios\\footnote{We adopt a convention in which axial ratios such as $c/a$ are given with a subscript of $\\Phi$ ($\\rho$) to indicate that they are measured with respect to the potential (density) contours.} $(c/a)_{\\Phi} \\sim 0.72$, $(b/a)_{\\Phi} \\sim 0.78$ (Hayashi et al. 2007). While there is a tendency for haloes to be more spherical at larger radii (Allgood et al. 2006; Hayashi et al. 2007), the halo of Milky Way-like systems at a distance of $\\sim 50$ kpc is expected to be prolate with $T_{\\Phi} \\approx 0.88$ (based on the Via Lactea simulation of Diemand et al. 2007; Kuhlen et al. 2007). Observationally, typical haloes appear to be rounder than predicted by collisionless CDM simulations, as indicated by gravitational lensing mass models (e.g., Hoekstra et al. 2004; Mandelbaum et al. 2006; Evans \\& Bridle 2009), polar ring galaxies (e.g., Iodice et al. 2003), and X-ray isophotal studies (e.g., Buote et al. 2002; although c.f. Diehl \\& Statler 2007; Flores et al. 2007). This possible mismatch is perhaps unsurprising because the growth of baryonic structures within these haloes tends to make them considerably rounder (e.g., Debattista et al. 2008; Valluri et al. 2009), inflating the principal axis ratios by as much as $\\sim 0.2 - 0.4$ within the virial radius (Kazantzidis et al. 2004). However, it is difficult to constrain the actual triaxiality of individual haloes because galaxies are projected onto the two-dimensional plane of the sky and observations are relatively insensitive to the mass distribution along the line of sight (although c.f. Corless et al. 2009). It is therefore noteworthy that for the only galaxy where we have a fully three-dimensional view from the interior of the mass distribution (i.e., the Milky Way) it appears (\\S \\ref{combinedconst.sec}) that the best-fitting axial ratio at radii $20 < r < 60$ kpc describes an almost perfectly {\\it oblate} ellipsoid with minor/major axis ratio $(c/a)_{\\Phi} = 0.72$, intermediate/major axis ratio $(b/a)_{\\Phi} = 0.99$, triaxiality $T_{\\Phi} = 0.06$, and the minor axis located in the plane of the Galactic disk in the direction $(l, b) = (7^{\\circ}, 0^{\\circ})$. In this section we discuss the physical interpretation of these results, along with possible effects that could influence our derivation of the halo shape. \\subsection{Physical Validity of the Triaxial Halo Model} \\subsubsection{Tracing the Underlying Mass Distribution} \\label{massmodel.sec} In \\S \\ref{model.sec} we introduced flattening directly into the contours of the gravitational potential because this was analytically tractable and consistent with the method of previous studies (in particular LJM05). It is well-known however that as the ellipticity of the potential becomes more extreme, the corresponding mass distribution required to generate it may become unphysical (i.e., have regions of negative mass). In order for any model of the Galactic halo to be believed, it must correspond to a physically realizable density distribution. In Figure \\ref{density.fig} we show the corresponding contours of the mass distribution underlying our best-fit model for the Milky Way (as calculated by taking the Laplacian of the gravitational potential). As expected, the mass distribution is less spherical than the isopotential surface; the contours of the dark matter halo are approximately elliptical at small radii (i.e., within $\\sim 40$ kpc) with $(c/a)_{\\rho} = 0.44$, $(b/a)_{\\rho} = 0.97$, and $T_{\\rho} = 0.07$. Although some pinching of the dark matter density contours is apparent at larger radii the mass distribution is everywhere positive. Figure \\ref{density.fig} also demonstrates that while addition of the baryonic mass component significantly changes the net density distribution at small radii and near the Galactic disk plane, the density distribution at larger radii probed by the orbit of Sgr is dominated by the dark halo component. \\begin{figure*} \\plotone{f10.eps} \\caption{Density maps viewed along different axes for the best-fit Galactic halo model in an $(X', Y', Z)$ coordinate system aligned with the principal axes of the halo. Density maps are shown for the dark halo alone, and for the sum of all mass components (i.e., logarithmic halo, Miyamoto-Nagai disk, and Hernquist spheroid). The greyscale/red contours are logarithmic and with the same scale in all panels. Note that while the axisymmetric density profile of the disk dominates in the disk plane ($Z=0$), at only 1 kpc above the plane the halo contribution produces noticeably elliptical contours in the $X'-Y'$ plane at radii $r \\gtrsim 20$ kpc. The dashed green circles indicate the mean apocenter/pericenter of the Sgr orbit.} \\label{density.fig} \\end{figure*} We additionally confirmed the validity of our results by constructing a triaxial NFW (Navarro et al. 1996) halo model in which the axial flattenings were introduced in a more physically motivated manner directly into the mass distribution and where the resulting accelerations were computed using the approximate form of the corresponding potential given by Lee \\& Suto (2003). As expected, there are slight differences in the path of satellites orbiting in this NFW halo and in our best-fit logarithmic potential, but these differences are minor and likely due to the fact that we did not fine tune the NFW model. \\subsubsection{Orientation of Principal Axes} While the mass distribution underlying our best-fit model of the Galactic halo is physically reasonable, the {\\it orientation} of the halo (i.e., strongly non-axisymmetric in the plane of the Galactic disk, with the angular momentum vector of the disk aligned with the intermediate axis of the halo) is unexpected because orbits about intermediate axes tend to be unstable (e.g., Adams et al. 2007). This is not a concern for stellar orbits in the Galactic disk because the gravitational potential at small radii $r \\lesssim 20$ kpc is dominated by the (nearly) axisymmetric baryonic mass distribution rather than the triaxial dark halo. Indeed, preliminary results of orbit integration in the gravitational potential derived here indicate that loop orbits are permitted, suggesting that a self-consistent disk may still be supported (R. Johnson and K.V. Johnston, private communication, 2010). However, this situation begs the question of why the Galactic disk formed in its present orientation unless the principal axes of the dark halo either twist substantially with radius (in contrast to the general prediction of Hayashi et al. 2007) or evolve with time. While we cannot rule out the possibility that rotation of the dark halo has contributed in some regard to our results, we neglect such effects because typical haloes are expected to rotate extremely slowly (e.g., Bullock et al. 2001a; Herbert-Fort et al. 2007, 2008; Bett et al. 2009). Typically, simulations of disk galaxy formation suggest (e.g., Debattista et al. 2008) that the minor axes of the dark halo and the baryonic disk should be approximately aligned (to within an average of $\\sim 30^{\\circ}$; see discussion by Warren et al. 1992; van den Bosch 2002; Bett et al. 2009) because both are formed from accreted matter with similar specific angular momenta. If a disk were to form in an elliptical halo potential such as that proposed in \\S \\ref{combinedconst.sec}, simulations suggest that the disk too should become elongated and non-axisymmetric (see, e.g., Gerhard \\& Vietri 1986; Debattista et al. 2008). It will therefore be intriguing to investigate whether the relative orientation of the host dark matter halo and baryonic disk proposed here can be strongly motivated within the current CDM paradigm. It is worth asking as well whether such an inclined oblate halo can explain the observed distribution of Galactic satellites, which simulations suggest (Zentner et al. 2005) should be distributed along the major axis of the host halo (i.e., have a pole aligned with the minor axis). The Fornax-Leo-Sculptor great circle has a pole $(l,b) = (135^{\\circ}, -2.9^{\\circ})$ (Lynden-Bell 1982; see also Kunkel \\& Demers 1976; Majewski 1994; Fusi Pecci et al. 1995), and a recent update (Metz et al. 2007, 2009) including the faint new dwarfs discovered in the SDSS gives a revised satellite pole $(l,b)=(159.7, -6.8)$. The updated pole of Metz et al. (2009) lies within $\\sim 30^{\\circ}$ of the minor axis we derive for the Galactic halo, but it is difficult to determine whether this similarity is simply an unrelated coincidence.\\footnote{If the probability distribution of the satellite pole were uniform across the entire sky, the chance for it to be coincidentally aligned with $(l,b) = (7^{\\circ}, 0^{\\circ})$ to within $30^{\\circ}$ is $\\sim 13$\\%.} \\subsection{Alternatives to the Triaxial Halo Model} In LJM05, we conducted an extensive parameter space search in an effort to find a combination of Milky Way + Sgr parameters capable of simultaneously matching both the observed angular position and distance/radial velocity trends of leading tidal debris. In particular, we explored the effect of varying the proper motion of Sgr ($v_{\\rm tan}$), the distance to Sgr ($D_{\\rm Sgr}$), the distance from the Sun to the Galactic Center ($R_{\\odot}$), the total mass of the Galactic disk (via the scaling parameter $\\alpha$ in Equation \\ref{diskeqn}), the dark halo scale length ($d$), the dark halo flattening perpendicular to the disk ($q_z$), and the total mass scale of the Milky Way as parameterized by the rotation speed of the Local Standard of Rest ($v_{\\rm LSR}$). Although the Galactic disk contributes significantly to the overall flattening of our adopted Galactic potential at small radii (e.g., Figure \\ref{density.fig}), it has only a minor effect at the range of radii ($r \\sim 20 - 60$ kpc) probed by the orbit of Sgr. As demonstrated by Figure \\ref{rotcurve.fig}, the dark halo in our model is the dominant contributor to the total gravitational potential at these radii. Indeed, if we decrease the mass of the Galactic disk by 50\\% (by setting $\\alpha = 0.5$ in Equation \\ref{diskeqn}) and increase the total halo mass correspondingly (so that the Local Standard of Rest remains constant at $v_{\\rm LSR} = 220$ \\kms), we can repeat the parameter space search described in \\S \\ref{mwprop.sec}. Scaling the proper motion $v_{\\rm tan}$ of Sgr to again best-fit the trailing-arm velocity trend, we find optimal values for $q_1, q_z, \\phi$ almost identical to those derived previously (although the overall $\\chi^2$ of the best-fit model is not as good as for the case $\\alpha = 1.0$). We are therefore confident that the details of the Galactic disk model do not significantly bias our results. \\begin{figure*} \\plottwo{f11a.eps}{f11b.eps} \\caption{Left-hand panel: Contribution of the dark halo, stellar disk, and stellar spheroid to our model Galactic rotation curve as a function of radius in the Galactic disk plane (for purposes of illustration the halo is taken to be axisymmetric with $q_1 = q_2 = 1.0$). The vertical dotted lines denote the mean pericenter/apocenter of the Sgr orbit. Right-hand panel: Acceleration of the Sgr dwarf in the gravitational potential of the Milky Way as a function of time over its assumed 8 Gyr orbital history. The dark halo, stellar disk, stellar spheroid, and total contributions to the rotation curve and acceleration of Sgr are indicated by dashed, long-dashed, dot-dashed, and solid lines respectively in both panels.} \\label{rotcurve.fig} \\end{figure*} Similarly, the details of the triaxial Galactic bar (which we did not include in our model) and stellar halo are unlikely to change our results. Neither are sufficiently massive to have a great effect on the gravitational potential at distances $r \\gtrsim 20$ kpc, and neither have orientations akin to that which we derive for the dark halo. The major axis of the bar is thought to lie within $\\sim 15^{\\circ}-20^{\\circ}$ of the $X$-axis (e.g., Blitz \\& Spergel 1991; Nakada et al. 1991; Morris \\& Serabyn 1996; Babusiaux \\& Gilmore 2005), similar to the minor axis of the dark halo. While the stellar halo is significantly triaxial ($[c/a]_{\\rho} \\approx 0.65$, $[b/a]_{\\rho} = 0.75$, $T_{\\rho} = 0.76$; Newberg \\& Yanny 2006; see also Larsen \\& Humphreys 1996; Xu et al. 2006) with a major axis lying approximately in the Galactic plane ($\\sim 50-70^{\\circ}$ from the $X$ axis) the minor axis is not apparently contained within the $X-Y$ plane but located $\\sim 13^{\\circ}$ from the $Z$ axis. Throughout the entirety of the LJM05 parameter space search, the only parameter that was found to have a significant effect on the angular position/distance/radial velocity trend of the Sgr leading arm was the flattening of the dark halo $q_z$. Although no single choice of $q_z$ can simultaneously reproduce all observational constraints, we have demonstrated (see also LJM09) that a fully triaxial model in which the short axis is approximately aligned with the Galactic $X$ axis is capable of doing so. Such a triaxial halo is therefore obviously attractive in the sense that {\\it only} $N$-body models in such a halo succeed in modeling the Sgr -- Milky Way system where previous efforts (e.g., Helmi et al. 2004; LJM05; Fellhauer et al. 2006; Mart{\\'{\\i}}nez-Delgado et al. 2007) have failed. Despite the extensive exploration of parameter space undertaken by ourselves (this paper; LJM05; LMJ09) and other groups (e.g., Helmi et al. 2004; Fellhauer et al. 2006; Mart{\\'{\\i}}nez-Delgado et al. 2007) however, the space explored to date is far from exhaustive and it is possible that the triaxial halo model may prove to be simply a numerical crutch that mimics the effect of some as-yet unidentified alternative. In the following sections, we briefly discuss a few possibilities that may obviate the need for triaxiality to explain the extant data. \\subsubsection{Alternative Formulations for the Gravitational Potential} While we have focused on a logarithmic formulation of the Galactic gravitational potential, other forms could also reasonably be adopted. In LJM05 we explored axisymmetric NFW (Navarro et al. 1996) models and were unable to resolve the halo conundrum. As demonstrated in \\S \\ref{massmodel.sec} however, fully triaxial NFW models can give rise to orbits for Sgr that match the observed angular precession and distance/radial velocity trends of leading tidal debris. We find that orbits within logarithmic and NFW haloes are sufficiently similar that there is no reason to differentiate between the two for our present purposes. Fellhauer et al. (2006) claim that a mass distribution of the form specified by Dehnen \\& Binney (1998) can produce some effects similar to prolateness in a logarithmic halo. However, using such a model these authors were nevertheless unable to reproduce the distance trend of leading tidal debris as seen in the SDSS (Belokurov et al. 2006) and concluded that their Set D of proper motions for Sgr (i.e., those derived by LJM05) in a logarithmic halo provided the best overall fit to the observational data then available. We therefore find no reason to prefer a Dehnen \\& Binney (1998) model over the triaxial logarithmic halo model presented here. \\subsubsection{Evolution in the Orbit} Although our model dSph has been orbiting in the Galactic potential for $\\sim$ 8 Gyr, we have not considered the influence of any evolution in either the orbit of Sgr or the underlying gravitational potential of the Milky Way. We expect there was actually evolution in both the depth, and the {\\it shape} of the potential (e.g., Kuhlen et al. 2007). While we expect the streams to be independent of past evolution in the potential of the Milky Way (since tidal streams respond adiabatically to changes in their host potential; see discussion by Pe{\\~n}arrubia et al. 2006) some evolution has certainly occurred in the orbit of Sgr since it fell into the Milky Way. Dynamical friction and other such evolutionary effects therefore reflect an uncertainty in our models, although we note that (as discussed by LJM05) such effects alone are unlikely to be able to resolve the halo conundrum. Further investigation of these effects must (among other things) properly account for the relative distribution and mass-loss history of both the dark and baryonic matter in Sgr, and is therefore beyond the scope of this contribution. \\subsubsection{Dwarf --- Tail Gravitational Interaction} As discussed by Choi et al. (2007), the gravitational influence of the bound satellite on debris in the tidal tails can shift the trajectories of actual tidal debris away from the orbital path traced by the bound core of the dSph. For sufficiently massive satellites (greater than $\\sim 0.1$\\% the virial mass of the host), this effect can be appreciable and degenerate with that caused by flattening (or triaxiality) in the underlying gravitational potential of the Milky Way. However, we note that despite claims to the contrary (Choi et al. 2007) our simulations fully account for this effect; the self-gravity of the bound satellite is applied to all $N$-body particles at all times regardless of whether or not they are in the Sgr core or in the tidal tails. \\subsubsection{Gravitational Influence of the Magellanic Clouds} The Large Magellanic Cloud (LMC) is $\\sim 50$ kpc distant in the direction $(l,b) = (280.5^{\\circ}, -32.8^{\\circ}$) (van der Marel et al. 2002), placing it within $\\sim 0.5$ kpc of the Galactic $Y-Z$ plane. Based on great-circle fits to the Magellanic stream, the pole of the system appears to be $(l_{\\rm p},b_{\\rm p}) = (188.5^{\\circ}, -7.5^{\\circ}$) (Nidever et al. 2008). That is, the pole of the Magellanic stream is aligned with the short axis of the Galactic dark halo as derived in this contribution to within $\\sim 1^{\\circ}$ in Galactic longitude.\\footnote{They differ by $7.5^{\\circ}$ in Galactic latitude, but the symmetry axes of the dark halo were fixed to lie at $b =0^{\\circ}$.} This suggests two possibilities: (1) The Magellanic Clouds may have fallen into the Milky Way along the plane of the long and intermediate axes of the dark matter halo (i.e., the alignment is a {\\it consequence} of the dark matter distribution), or (2) the Magellanic Clouds may exert a significant gravitational acceleration on the Sgr dwarf, with the mass of the Clouds at large radii in the $Y-Z$ plane mimicking the effect of an extended dark matter distribution in this plane (i.e., the alignment is the {\\it cause} of the apparent dark matter distribution). A precise description of the second scenario would require a comprehensive model for the orbital and mass-loss history of the Magellanic Clouds, but it is possible to understand the magnitude of their gravitational influence on the path of Sgr tidal debris by integrating massless test-particles along the Sgr orbit in a Milky Way with a spherical dark halo and simplified model of the LMC. We first test the effect of introducing a fixed point-mass into the gravitational potential of the Milky Way at the present location of the LMC, whose mass is allowed to vary from $1$ - 10\\% of the total mass of the Milky Way within 50 kpc (i.e., $M_{\\rm LMC} = 6 \\times 10^9 M_{\\odot} - 6\\times 10^{10} M_{\\odot}$). As illustrated in Figure \\ref{LMCeffect.fig}, a fixed LMC with a mass 10\\% that of the Milky Way can produce a significant deviation on the order of $20^{\\circ}$ in the orbital path of leading Sgr tidal debris. Of course the LMC is not fixed in space, but describes an orbit about the Milky Way confined approximately to the $Y-Z$ plane. We therefore develop a second, slightly more realistic model in which the LMC is described in a time-integrated sense by a ring of mass at a radius of 50 kpc in the $Y-Z$ plane. The gravitational acceleration at an arbitrary point in space resulting from such a configuration is described using an analytical approximation to the numerically evaluated elliptical integrals (see derivation for the electrostatic case by Zypman 2005). This second model produces even more noticeable deviations from the original orbital path, with appreciable shifts in both the leading arm angular position and radial velocities expected for models in the range of masses $1 - 10$\\% that of the Milky Way. \\begin{figure*} \\epsscale{0.8} \\plotone{f12.eps} \\caption{\\scriptsize Effect of the LMC on the angular path, distance, and line of sight velocity of the leading/trailing Sgr orbit (left- and right-hand panels respectively). Observational data (open symbols) are as in Figure \\ref{ModelSummary.fig}. The cyan line represents the path of the Sgr dwarf in a simple spherical halo model, green solid/dashed lines the path in a spherical halo with a point-mass LMC that is 1\\%/10\\% of the Milky Way mass, and red solid/dashed lines the path in a spherical halo with a ring-model LMC that is 1\\%/10\\% of the Milky Way mass. It is difficult to see the solid green line in places since the cyan line lies almost on top of it. Note the strong deviations of the models incorporating the LMC from the predictions of a simple spherical halo model, especially for larger values of the LMC mass. The black line indicates the path of the dwarf in the triaxial halo model for comparison.} \\label{LMCeffect.fig} \\end{figure*} The detailed effect of the Magellanic Clouds on the orbit of Sgr is extremely uncertain. Both our point-mass and ring-mass models are obviously oversimplifications, neglect the SMC entirely, and assume that the gravitational influence of the LMC has been fixed over the entire interaction history of the Sgr dwarf ($\\sim 8$ Gyr in the current model). Since the Clouds are currently thought to be near the pericenter of their orbit, for the majority of the lifetime of Sgr they likely lay at greater distances where their gravitational force would be smaller. Indeed, some recent lines of investigation suggest that they may have only recently fallen into the gravitational potential of the Milky Way for the first time (e.g., Besla et al. 2007), drastically limiting the time during which they could have torqued the orbit of the Sgr dwarf. Of the various orbital paths shown in Figure \\ref{LMCeffect.fig}, only the triaxial halo model actually succeeds in reproducing all of the observational data. However, given the myriad uncertainties in the Milky Way - LMC - Sgr system the important conclusion to draw is that gravitational peturbations from the Magellanic Clouds {\\it may} be sufficient to torque the leading arm of the Sgr stream significantly, especially if the effect is enhanced by the response of the Milky Way dark matter halo itself to the passage of the Clouds (see, e.g., Weinberg 1998). \\subsubsection{Alternative gravity} Modified Newtonian gravity (MOND; see Milgrom 1983; Bekenstein 2004) on Galactic scales has been proposed as an alternative theory to explain the flat rotation curves typical of galaxies without recourse to an extended halo of dark matter. Read \\& Moore (2005) demonstrated that MOND was an equally viable alternative to the oblate dark matter halo model presented by LJM05, capable of fitting the angular position of the leading tidal debris but not the radial velocities. Now that a dark matter model has been able to successfully match {\\it both} observational constraints, it is worthwhile to investigate whether a similarly successful MOND model can be constructed. Since the key to resolving the halo conundrum appears to have been the inclusion of a strongly non-axisymmetric component to the gravitational potential, it is not immediately obvious how such a potential could be produced in a MOND model by the predominantly axisymmetric distribution of baryonic matter typically adopted for the Milky Way (although c.f. Wu et al. 2008; Widrow 2008). \\subsection{Testing the Model} Fortunately, the model presented here for the shape, size, and orientation of the Galactic gravitational potential is easily testable; it must be capable of explaining the orbital paths of other Galactic satellites that similarly provide constraints on the mass distribution in the Milky Way. Typically, previous studies of tidal debris systems in the Galactic halo have only explored the effects of changing the axial scalelength perpendicular to the Galactic disk (i.e., $q_z$), and it is unclear whether satisfactory solutions for other satellites can be found in such a strongly non-axisymmetric potential as that suggested here. In particular, can realistic orbits be found to match such systems as the Monoceros tidal stream (e.g., Pe{\\~n}arrubia et al. 2005), the Magellanic Clouds\\footnote{While simulations generally favor haloes which are extended along the Galactic $Z$ axis (similar to those in the triaxial model), preliminary results indicate that the Magellanic Clouds are relatively insensitive to the distribution of mass along the Galactic $Y$ axis and can neither confirm nor deny the triaxial model presented here (G. Besla, priv. comm.).} (e.g., Besla et al. 2007; R{\\u u}{\\v z}i{\\v c}ka et al. 2007), the ``orphan stream'' (Belokurov et al. 2007), the GD-1 stream (Koposov et al. 2009), the newfound Cetus Polar Stream (Newberg et al. 2009), and tidally disrupting Galactic globular clusters such as Pal 5 (Odenkirchen et al. 2009)? Similarly, is a triaxial Galactic halo consistent with observations of the Galactic HI disk (e.g., Olling \\& Merrifield 2000; Kalberla et al. 2005; and references therein)? It is intriguing that Saha et al. (2009) discuss evidence to suggest that the dark matter halo must be similarly non-axisymmetric in the Galactic $XY$ plane in order to explain the asymmetric flaring of the disk (although c.f. Narayan et al. 2005). Future observations of hypervelocity stars at large distances may also provide key diagnostics (e.g., Gnedin et al. 2005; Yu \\& Madau 2007; Perets et al. 2009). \\label{otherobs.sec} The model presented in this paper has the attractive quality that it provides a reasonable match to the major dynamical characteristics of Sgr tidal debris measured from wide-field surveys. It is worthwhile also investigating the degree of correspondence between the $N$-body model and various narrow-field observations of subgiant stars, K-giants, and RRL that have been postulated to trace the Sgr tidal stream. In Figure \\ref{OtherObs.fig} we plot the $N$-body model of the Sgr stream overlaid by a selection of detections from recent narrow-field surveys (this comparison is not exhaustive, see also discussion by Majewski et al. 2003). We discuss each data set in detail below. Note that we do not consider possible associations with globular clusters or faint dwarf satellites; these are treated separately in a forthcoming contribution (Law et al. {\\it in prep.}). \\begin{figure*} \\epsscale{0.8} \\plotone{f19.eps} \\caption{\\scriptsize Model (colored points) of the Sgr stream overlaid with various detections of potential Sgr tidal debris. We include 2 Gyr more of tidal debris in this plot (green points) than in previous plots (e.g., Figure \\ref{ModelSummary.fig}) to indicate the possible path of debris if Sgr has been orbiting the Milky Way for more than $\\sim 5$ Gyr. Overlaid on the $N$-body model is observational data from Vivas et al. (2005; open boxes), the SDSS `Stripe-82' (Cole et al. 2008; Watkins et al. 2009; filled circle), Duffau et al. (2006; open circle), `VOD' observations from Vivas et al. (2008; filled boxes), K-giant stars from Starkenburg et al. (2009; skeletal triangles) and Kundu et al. (2002; filled triangles), RR-Lyrae stars from Prior et al. (2009; open triangles), and red clump stars from Majewski et al. (1999; crosses). Note that at $\\Lambda_{\\odot} \\approx 260^{\\circ} - 300^{\\circ}$ it is difficult to distinguish the open circle of Duffau et al. (2006; because it lies atop the open triangles of Prior et al. 2009 in $\\Lambda_{\\odot}$ vs distance space), and the skeletal triangles of Starkenburg et al. (2009; because they lie amidst the open boxes of Vivas et al. 2005).} \\label{OtherObs.fig} \\end{figure*} \\subsection{SDSS Stripe 82} Cole et al. (2008) used a maximum likelihood method to characterize the spatial properties of F-turnoff stars in stripe 82 of the SDSS. They observed tidal debris located at $(\\alpha, \\delta, R) = (31.37^{\\circ} \\pm 0.26^{\\circ}, 0.0^{\\circ}, 29.22 \\pm 0.20 \\, {\\rm kpc})$, corresponding to Sgr longitudinal coordinates $(\\Lambda_{\\odot}, B_{\\odot}) = (105.2^{\\circ}, -1.15^{\\circ})$. As illustrated by Figure \\ref{OtherObs.fig} (filled circles) the T1 wrap of the $N$-body model is an extremely good match to these F-turnoff stars in both angular position and distance (matching to within $0.7^{\\circ}$ in $B_{\\odot}$, and 1 kpc in distance). Cole et al. (2008) measured the width of the stream perpendicular to the line of sight at this $\\Lambda_{\\odot}$ to have a FWHM $ = 6.74 \\pm 0.06$ kpc. The $N$-body stream is slightly fatter with FWHM $ = 7.9$ kpc, but given the single component mass-follows-light model employed for Sgr the agreement with the apparent width of a specific stellar population is extremely good. The radial velocity of a sample of RRL in Stripe 82 has recently been measured by Watkins et al. (2009), who detect a component with $v_{\\rm GSR} = -130$ \\kms, which is matched by the mean radial velocity of the T1 stream at this position to within 5 \\kms. Watkins et al. (2009) also measure the metallicity of these RRL, finding that $\\langle \\feh \\rangle = -1.41 \\pm 0.19$. This value is appreciably more metal-poor than either the M-giant sample or the mean of the $N$-body model at this longitude ($\\feh \\sim -0.7$). This difference in measured metallicity emphasizes the strong sensitivity of derived metallicities of a stellar stream to the biases imposed by the type of stellar tracer used (e.g., metal-poor populations do not generally form M giants while metal-rich populations do not generally form RRL). Thus, until an unbiased assessment of the MDF along the Sgr stream is obtained, models such as that presented in \\S \\ref{feh.sec}, which was constrained by a single type of metallicity-biased population tracer, must be considered as merely a demonstration of how MDF-variations can originate. \\subsection{RRL and K-giants at $\\Lambda_{\\odot} \\sim 270^{\\circ} - 300^{\\circ}$} Vivas et al. (2005) published VLT spectroscopy of 16 RRL from the QUEST survey in the range $\\Lambda_{\\odot} \\sim 270^{\\circ} - 300^{\\circ}$. All but 1-2 of these stars, (which have highly discrepant velocities, at $v_{\\rm GSR} = -106$ and $-197$ \\kms) are in agreement with the angular position, distance and radial velocity trend of the L1 arm of the $N$-body model (see Figure \\ref{OtherObs.fig}, open boxes). Starkenburg et al. (2009) have also detected an overdensity of 5 K-giant stars at similar orbital longitude that are well-matched to the distance and radial velocity of the $N$-body model stream (Figure \\ref{OtherObs.fig}, skeletal triangles), although they appear to be located towards the edge of the stream in $B_{\\odot}$.\\footnote{Note that the star added to the Starkenburg et al. (2009) group using the `4dist' $\\leq 0.08$ criterion is not as well-matched to the $N$-body stream as the other 5 stars, with the former discrepant by $\\sim 60$ \\kms.} Both samples of stars are well-reproduced by the $N$-body model, although we note that the Vivas et al. (2005) and Starkenburg et al. (2009) samples have mean metallicity $\\langle \\feh \\rangle = -1.77$ and $-1.68$ respectively. These metallicities are entirely consistent with each other, but significantly more metal-poor than the M-giant sample, which has $\\langle \\feh \\rangle \\sim -0.9$. As was the case for SDSS Stripe 82, this reinforces the conclusion that stellar populations significantly more metal-poor than the M-giant subsample exist at a range of orbital longitudes throughout the Sgr tidal streams. As expected on the basis of dynamical age, the RRL of Vivas et al. (2005) which correspond to tidal debris $\\sim 1.5 - 3$ Gyr old (i.e., magenta points in Figure \\ref{OtherObs.fig}) are slightly more metal-poor than the RRL in SDSS Stripe 82 which correspond to $\\sim 0 - 1.5$ Gyr old debris. We note that an additional sample of metal-poor K-giant stars in a similar longitude range was also pointed out by Kundu et al. (2002; filled triangles in Figure \\ref{OtherObs.fig}). While the L2/T2 wraps of the $N$-body stream match the radial velocity trend of these stars, the angular positions are discrepant by up to $20^{\\circ}$ for L2 and $30^{\\circ}$ for T2. At a typical distance $\\sim 5$ kpc the Kundu et al. (2002) stars are also much closer than expected for the T2 arm ($d \\sim 50$ kpc), but may be consistent with the L2 arm ($d \\sim 10$ kpc). If these stars genuinely belong to the Sgr stream, it is most likely that they belong to the L2 wrap, and may suggest that the angular coordinates and/or distance of this wrap in the $N$-body model require adjustment. \\subsection{Virgo: The VSS and the VOD} Recent years have witnessed a wealth of substructure discovered in the direction of Virgo, which we follow Mart{\\'{\\i}}nez-Delgado et al. (2007) in classifying into the Virgo Stellar Stream (VSS) and the Virgo Overdensity (VOD). The VSS (as traced by QUEST RRL; Duffau et al. 2006) is a relatively well-defined feature that lies at an angular position $(\\Lambda_{\\odot},B_{\\odot}) = (265^{\\circ}, 14^{\\circ})$, i.e., projected on the outer edge of the L1/L2 streams and significantly off the T1/T2 streams in sky position. At a distance of 19 kpc, it is roughly coincident with the L2/T1 stream, but has a radial velocity ($v_{\\rm GSR} = 99.8$ \\kms; Duffau et al. 2006) that is discrepant with these wraps by 220/94 \\kms respectively. Curiously, the proper motion of the VSS ($[\\mu_{\\alpha} \\textrm{cos} \\delta, \\mu_{\\delta}] = [-4.85 \\pm 0.85, 0.28 \\pm 0.85]$; Casetti-Dinescu et al. 2009) is reproduced by the T1 stream to within $1\\sigma$ ($[\\mu_{\\alpha} \\textrm{cos} \\delta, \\mu_{\\delta}] = [-4.18 \\pm 0.41, 0.04 \\pm 0.25]$), but given the angular position and radial velocity mismatch we conclude that the VSS is unlikely to be associated with the Sgr stream. In contrast, the VOD (Juri{\\'c} et al. 2008) is a much less well-defined, diffuse clump of stars around $(l,b) = (300^{\\circ}, 65^{\\circ})$ that covers more than $\\sim 1000$ deg$^2$ on the sky. The possible association of the VOD with Sgr has been addressed in detail by Mart{\\'{\\i}}nez-Delgado et al. (2007), who noted that it was roughly coincident with tidal debris from models of the Sgr dwarf (LJM05) disrupting in an axisymmetric, oblate Galactic halo. In the triaxial halo model derived in this contribution, we find that the model Sgr stream arcs significantly over the solar neighborhood, and tidal debris in the L1 tidal stream at the angular coordinates of the VOD lies much too far away ($\\sim 46$ kpc for the Sgr stream versus $\\sim 5-17$ kpc for the VOD) to be associated (see Figure \\ref{OtherObs.fig}, filled squares). Similarly, both the T1 and T2 wraps lie too far away to produce the VOD, in addition to lying in the opposite angular direction $B_{\\odot}$. Interestingly, this angular position and distance {\\it does} correspond closely to the predicted location of the secondary L2 wrap of leading Sgr tidal debris, suggesting that the VOD may be evidence for old tidal debris torn from Sgr $\\gtrsim 3$ Gyr ago. If such an identification is correct, we should expect the radial velocity of the VOD to be $v_{\\rm GSR} = -131 \\pm 22$ \\kms. Recent observational work on QUEST RRL in the direction of the VOD by Vivas et al. (2008) derived radial velocity signatures at $v_{\\rm GSR} = +215, -49, -171$ \\kms. The first (and by far most significant) two of these velocity structures are strongly discrepant with the expected signal by $\\sim 345$ and $\\sim 80$ \\kms respectively. The third structure is possibly consistent with the velocity signature expected of the L2 Sgr stream, differing by only $\\sim 40$ \\kms in this part of the simulated debris where the $N$-body stream is ill-constrained. Only 7\\% of the QUEST RRL (i.e., 3/43) are contained in this $v_{\\rm GSR} = -171$ \\kms velocity structure however, and we therefore concur with the conclusion of Vivas et al. (2008) that while Sgr may make some minor contribution to the VOD it is not the dominant origin of the stellar overdensity. \\subsection{SEKBO RRL} An additional sample of RRL stars from the SEKBO (Moody et al. 2003) survey has recently been studied by Keller et al. (2008), with radial velocities presented by Prior et al. (2009). The Prior et al. (2009) subsample focuses on two groups of RRL which they group into the ``VSS region'' at $\\alpha = 12.4 - 14$ hr (i.e., $\\Lambda_{\\odot} \\approx 260^{\\circ}-290^{\\circ}$) and the ``Sgr region'' at $\\alpha = 20 - 21.5$ hr (i.e., $\\Lambda_{\\odot} \\approx 10^{\\circ}-40^{\\circ}$). These two groups of stars are plotted against our $N$-body model in Figure \\ref{OtherObs.fig} (open triangles). Considering first the angular coordinate plot $\\Lambda_{\\odot}$ vs. $B_{\\odot}$, we note that neither of these groups of stars are well-centered on the predicted path of the $N$-body tidal stream, lying at the extreme outer edge of the L1/L2 streams and fully outside the angular path of the T1/T2 streams. This is consistent with Figure 15 of Keller et al. (2008): The plane of the SEKBO survey intersects the Sgr plane at an angle and best overlaps at $\\alpha \\approx 18$h (i.e., behind the Galactic Center) and $\\alpha \\approx 5$h. In the regions traced by the Prior et al. (2009) subsample we confirm that the RRL tend to lie $\\sim 10^{\\circ} - 15^{\\circ}$ away from the bulk of the Sgr stream. It is perhaps unsurprising then that the distances (typically 16-21 kpc) and radial velocities of the RRL do not obviously match any particular branch of the Sgr stream. While Prior et al. (2009) discussed a possible match between some RRL and trailing Sgr tidal debris in the oblate-halo model of LJM05, no such match obviously exists for the updated model presented here. We note that the distance ($\\sim 20$ kpc) and radial velocity range ($v_{\\rm GSR} \\approx -200$ - $+200$ \\kms) of the $\\Lambda_{\\odot} \\approx 10^{\\circ}-40^{\\circ}$ sample of Prior et al. (2009) is similar to that found for a sample of red clump stars by Majewski et al. (1999; crosses in Figure \\ref{OtherObs.fig}). These red clump stars lie near the angular position of the L1 stream, and LJM05 discussed the possibility that their large spread in radial velocities may indicate the range of velocities occupied by the Sgr stream at this longitude. Given the similarities in distance and radial velocity of the Prior et al. (2009) and Majewski et al. (1999) samples, it is at least likely that they trace the same Galactic substructure. If this substructure is confirmed to be the Sgr stream, it will suggest that the L1 stream at $\\Lambda_{\\odot} \\sim 25^{\\circ}$ has a greater angular width, smaller distance, and larger spread in radial velocities than predicted by the current $N$-body model. \\subsection{Subgiant Populations in the Leading Tidal Stream} Keller (2009) has recently described a population of subgiant stars in the range $\\alpha \\approx 120^{\\circ} - 180^{\\circ}$ from the SDSS that may correspond to the L1 stream of Sgr. While this paper gives insufficient information to include these stars in Figure \\ref{OtherObs.fig}, we note that the $\\Lambda_{\\odot}$ vs distance trend illustrated in Figure 7 of Keller (2009) is reproduced by the revised $N$-body model to within $\\sim 1$ kpc, in contrast to previous $N$-body models of the Sgr stream (e.g., LJM05). Further characterization of this subgiant population may therefore provide even greater insight into the nature of the Sgr leading tidal stream." }, "1003/1003.1768_arXiv.txt": { "abstract": "We present the first results from a $1.1\\,\\rm{mm}$ confusion-limited map of the Great Observatories Origins Deep Survey-South (GOODS-S) taken with the AzTEC camera on the Atacama Submillimeter Telescope Experiment. We imaged a $270\\,\\rm{arcmin}^2$ field to a $1\\sigma$ depth of $0.48-0.73\\,\\rm{mJy/beam}$, making this one of the deepest blank-field surveys at mm-wavelengths ever achieved. Although by traditional standards our GOODS-S map is extremely confused due to a sea of faint underlying sources, we demonstrate through simulations that our source identification and number counts analyses are robust, and the techniques discussed in this paper are relevant for other deeply confused surveys. We find a total of 41 dusty starburst galaxies with signal to noise ratios $S/N\\ge3.5$ within this uniformly covered region, where only two are expected to be false detections, and an additional seven robust source candidates located in the noisier ($1\\sigma\\approx1\\,\\rm{mJy/beam}$) outer region of the map. We derive the $1.1\\,\\rm{mm}$ number counts from this field using two different methods: a fluctuation or ``$P(d)$'' analysis and a semi-Bayesian technique, and find that both methods give consistent results. Our data are well-fit by a Schechter function model with $(S^{\\prime},N_{3\\rm{mJy}},\\alpha)=(1.30^{+0.19}_{-0.25}\\,\\rm{mJy},160^{+27}_{-28}\\,\\rm{mJy}^{-1}\\rm{deg}^{-2},-2.0)$. Given the depth of this survey, we put the first tight constraints on the $1.1\\,\\rm{mm}$ number counts at $S_{1.1\\rm{mm}}=0.5\\,\\rm{mJy}$, and we find evidence that the faint-end of the number counts at $S_{850\\mu\\rm{m}}\\lesssim2.0\\,\\rm{mJy}$ from various SCUBA surveys towards lensing clusters are biased high. In contrast to the 870\\,\\micron~survey of this field with the LABOCA camera, we find no apparent under-density of sources compared to previous surveys at 1.1\\,mm; the estimates of the number counts of SMGs at flux densities $\\mathbf{>1}$\\,mJy determined here are consistent with those measured from the AzTEC/SHADES survey. Additionally, we find a significant number of SMGs not identified in the LABOCA catalogue. We find that in contrast to observations at $\\lambda\\le500$\\,\\micron, MIPS 24\\,\\micron~sources do not resolve the total energy density in the cosmic infrared background at $1.1\\,\\rm{mm}$, demonstrating that a population of $z\\gtrsim3$ dust-obscured galaxies that are unaccounted for at these shorter wavelengths potentially contribute to a large fraction {\\bf ($\\sim2/3$)} of the infrared background at $1.1\\,\\rm{mm}$. ", "introduction": "\\label{sec:int} Galaxies selected at submillimetre (sub-mm) and millimetre (mm) wavelengths (hereafter SMGs) comprise a population of dust-obscured starburst or active galactic nuclei (AGN) host galaxies at high redshift \\citep[$z\\gtrsim1$; see review by][]{blain02}. With far-infrared (FIR) luminosities $L_{\\rm{FIR}}\\gtrsim10^{12}\\,\\rm{L}_\\odot$, these systems appear to be scaled-up analogs to the ultra-luminous infrared galaxies (ULIRGs) observed in the local Universe \\citep{sanders96}. Their FIR to mm spectral energy distributions (SEDs) are characterised by thermal dust emission with temperatures of $T_d\\sim35-40$\\,K \\citep{chapman05, kovacs06, pope06, coppin08}, peaking in the FIR at rest-frame $\\lambda\\sim100$\\,\\micron. Due to the steep rise with frequency of the SED on the Rayleigh-Jeans tail \\citep[$S_{\\nu}\\propto\\nu^{3-4}$;][]{dunne00, dunne01}, the FIR peak is increasingly redshifted into the sub-mm/mm observing bands with increasing distance, resulting in a strong negative k-correction that roughly cancels the effects of cosmological dimming with redshift for observations at $\\lambda\\gtrsim500$\\,\\micron. This makes SMGs of a given bolometric luminosity equally detectable between $110^4$~cm$^{-3}$, while CO can be thermally excited for densities $\\simlt10^3$~cm$^{-3}$. Based on the HCN/$L_{\\rm IR}$ relation, the bimodal trends of $L_{\\rm IR}$ to $M_{\\rm gas}$ in disks versus starbursts can be interpreted in terms of a similarly bimodal behavior for the dense gas fractions, with roughly 10 times higher fraction of dense gas in the starbursts compared to disks at fixed $L_{\\rm IR}$ (but also with about twice higher dense gas fractions in BzKs versus local spirals). All the observations might thus be explained by a genuine increase of SFR efficiency in some galaxy classes, probably due to the concentration of the gas at high volume densities. Major mergers or other kinds of instabilities appear to be the most natural explanation for the increased efficiencies in the starbursts (although not all mergers will necessarily produce this effect, di Matteo et al.\\ 2008). This is not a new scenario. A higher SF efficiency in merger-driven starbursts than in rotating disks has often been implemented in semianalytical models of galaxy formation \\citep[e.g.,][]{gui98,som01}, although this has been usually interpreted as the natural outcome of a single SF law with an exponent $>1$, as long as mergers make gas lose angular momentum and concentrate in the galaxy center. However, in such implementations the occurrence of very high SF rates in gas-rich disks is neglected; they thus have a bimodality in surface density, not in the SF law. Our analysis suggests that the original K98 calibration can account for the properties of disk galaxies at low and high redshifts but would underestimate the SF efficiency of starbursts by a factor of 10. A SF law with a higher exponent of 1.7 (Bouche et al.\\ 2007) would in turn overestimate SF efficiency of gas-rich disks by a similar amount. An implementation of such a double SF law would surely influence predictions from semianalytical models of galaxy formation. The difference in $\\alpha_{\\rm CO}$ for disks and starbursts helps to {\\it hide} what we are interpreting here as large differences in the SF efficiency, expressed in terms of $L_{\\rm IR}/M_{\\rm gas}$, by reducing the observed differences in $L_{\\rm IR}/L'_{\\rm CO}$. There seems to be a conspiracy at work such that the particular physical conditions that lead to high $L_{\\rm IR}/M_{\\rm gas}$ in starbursts also determine variations in $\\alpha_{\\rm CO}$ that obscure observationally the differences in SF efficiency. We emphasize thus that the distinction of starburst (or merging vs non merging) systems is important for interpreting CO observations, although this may be difficult on a case by case basis. We caution that blindly applying the same conversion factor to all high-z observations can lead to confusion. Also, care must be taken that SFR/IR-luminosities are accurately derived. The estimates for BzK galaxies here are based on the cross-comparison of 3 independent SFR indicators that agree within each other very well, and overall on the global assessment by Daddi et al.\\ (2007ab) on SFR measuraments of near-IR selected galaxies at $z\\sim2$. On the other hand, purely radio selected or mid-IR selected populations are likely to produce a mixed bag of merging systems and disk galaxies and can be affected by AGNs." }, "1003/1003.3840_arXiv.txt": { "abstract": "The properties of a massive star prior to its final explosion are imprinted in the circumstellar medium (CSM) created by its wind and termination shock. We perform a detailed, comprehensive calculation of the time-variable and angle-dependent transmission spectra of an average-luminosity Gamma-Ray Burst (GRB) which explodes in the CSM structure produced by the collapse of a 20 $M_{\\sun}$, rapidly rotating, $Z=0.001 $ progenitor star. We study both the case in which metals are initially in the gaseous phase, as well as the situation in which they are heavily depleted into dust. We find that high-velocity lines from low-ionization states of silicon, carbon, and iron are initially present in the spectrum only if the metals are heavily depleted into dust prior to the GRB explosion. However, such lines disappear on timescales of a fraction of a second for a burst observed on-axis, and of a few seconds for a burst seen at high-latitude, making their observation virtually impossible. Rest-frame lines produced in the termination shock are instead clearly visible in all conditions. We conclude that time-resolved, early-time spectroscopy is not a promising way in which the properties of the GRB progenitor wind can be routinely studied. Previous detections of high velocity features in GRB UV spectra must have been due either due to a superposition of a physically unrelated absorber or to a progenitor star with very unusual properties. ", "introduction": "For a few hours after their onset, the afterglows of Gamma-Ray Bursts (GRBs) are the brightest sources in the far Universe. Their high luminosity, together with their powerlaw, featureless spectrum, make them ideal sources to probe their surrounding environment through the absorption lines imprinted in their spectra. Theoretical studies have suggested that long Gamma Ray Bursts (GRBs) are produced by the collapse of rapidly-rotating, chemically homogeneous, massive stars (e.g. Macfadyen \\& Woosley 1999). The association between GRBs and massive stars, which has been observationally supported \\citep{Stanek_etal03, Hjorth_etal03}, makes absorption studies potentially useful as a new way to probe star-forming regions at intermediate and high redshifts, and/or the last hundreds of years of the progenitor evolution. Absorption lines imprinted by the material ejected by the star prior to its explosion allow one to probe the velocity structure and metal content of the ejecta. High-resolution spectroscopy of GRB afterglows has been performed in a number of cases (e.g. \\citet{Moller_etal02, Matheson_etal03, Mirabal_etal03, Schaefer_etal03, Starling_etal05, Fiore_etal05, D'Elia_etal07, Chen_etal08}; Prochanska et al. (2008a), (2008b); Fox et al. (2008); Thoene et al. (2008); see also Whalen et al. 2008 for an extended discussion on absorbers in GRBs), yielding constraints on the nature of the absorbing medium. An especially well studied burst was GRB~021004, whose high-resolution spectroscopy revealed a complex velocity structure of the absorbing material, with velocities up to $\\ga 3000$ km/s. The interpretation of these features has been controversial. Initial studies claimed that the high velocity lines were a direct proof of the association of GRBs with WR stars. \\citet{Starling_etal05} argued that the lines must be produced in a fossil stellar wind with hydrogen enrichment from a companion. \\citet{Mirabal_etal03} and \\citet{Schaefer_etal03} interpreted the lines as the result of shells of material which are present around the progenitor. More recent studies have however cast doubts on the initial interpretation. \\citet{Lazzati_etal06} performed a detailed time dependent analysis, taking into account the burst flash ionization. Even though they still considered the wind of the WR progenitor star as the best absorber candidate, they pointed out that such an interpretation would require a termination shock at a distance of at least 100~pc. Such a large radius of the termination shock is somewhat at odds with current wind models and could be accounted for only if the progenitor were an extremely massive star evolving in a fairly low-density environment. \\citet{Chen_etal07} studied a sample of 5 GRBs with high velocity features (among which GRB~021004). Using fine-structure transitions and the presence of low-ionization species, they argued that the location of the absorber is very likely at a large distance from the burster ($>1$~kpc), favoring an intervening halo, physically unrelated to the GRB, as the location of the absorbing material. Similarly, \\citet{Vreeswijk_etal07} modeled fine structure lines present in the spectrum of GRB 060418 with a UV pumping model and concluded that the absorbers are at a distance of 1.7 $\\pm$ 2 kpc. These results are consistent with a GRB completely ionizing the local CSM, and with high-velocity absorbers elsewhere in the host galaxy or intervening intergalactic space (see also Prochanska et al. (2008a,b) and D'Elia et al. (2009a) for similar findings in the case of other GRBs they studied). Most of these studies do not take into account the possible shielding effects of an unusually dense CSM, as well as any effects due to dust. Additionally, when photoionization calculations are performed, most assume a steady state solution, whereas the radiation fields of GRBs are highly temporally variable, with initial burst timescales of a few to tens of seconds, and the afterglow varying over minutes to hours. Disentangling the physical origin of the spectral absorption features seen in GRB afterglows is clearly of great importance for a full understanding of the GRB phenomenon. The aim of this paper is a detailed, systematic study of the time-dependent transmittance spectrum of a GRB exploding in a dusty circumburst medium (CSM) shaped by its progenitor star. Recent studies have suggested that long GRBs are produced by the collapse of rapidly-rotating, chemically homogeneous, massive stars \\citep{YoonLanger05, WoosleyHeger06, Yoon_etal06, Cantiello_etal07}. This kind of evolution is believed to take place at low metallicity, where stellar winds are weak and angular momentum loss by winds is marginal. The wind of numerous types of GRB and SN progenitors has been simulated by \\citet{vanMarle_etal08}, including one model which recreates an anisotropic wind close to the star. This theoretical scenario has found support in recent observations. \\citet{Campana_etal08} performed a detailed study of the X-ray spectrum of GRB~060218, discovering a larger than normal N/O ratio in the surrounding of the burst progenitor. They concluded that only a progenitor star characterized by a fast stellar rotation and sub-solar initial metallicity could produce such a metal enrichment. In this paper we perform a detailed and comprehensive study of the effects of the GRB and afterglow radiation on the CSM structure produced by the stellar wind of a rapidly-rotating, low-metallicity, massive progenitor star \\citep{vanMarle_etal08}. We aim at predicting the time-variable, angle-of-sight dependent transmittance spectra of a GRB afterglow which resulted from the collapse of such a star. We study two cases, that of a CSM whose metals are initially in the gaseous phase (i.e. no dust), and that of a CSM whose metals are heavily depleted into dust prior to the GRB explosion. Even though the formation and survival of dust in the wind of Wolf-Rayet stars may seem surprising, dust is observed in WR winds \\citep{Allen_etal72,Williams_etal90}. The gas temperature and the exposure to the star's UV radiation make the wind a very harsh environment for dust particles. If the dust were to share the temperature of the gas phase, the grains would quickly sublimate. One key ingredient that allows the dust to survive is its capability of radiatively cooling to temperatures lower than the gas in which it is embedded (e.g. \\citep{Lazzati08}). Even so, there are still many unknown riddles in the way in which dust particles are born and survive in WR winds, but their presence has been clearly detected (e.g. Crowther 2003). Although it is presently unknown whether the particular stellar evolution model that we consider here is truly representative of critically rotating stars in general, we believe that the ability to compare theoretical absorption features with observational data can be potentially useful to the study of both GRBs and massive star evolution. ", "conclusions": "In this paper we have presented a comprehensive calculation of the spectra of a typical GRB which explodes in the CSM created by a massive, chemically homogeneous, low-metallicity rotating star. We have explored both a dusty CSM as well as a dust-free one. These stars are strong candidates for the typical progenitors of long GRBs; therefore, the models presented here, in combination with an average GRB luminosity and spectrum, are expected to yield the typical expectations for the spectra during both the prompt and the afterglow phase of the GRB. We emphasize, however, that an accurate prediction can be obtained only from a tailored time-dependent simulation for each particular burst, and hence our results should rather be seen as an example of the overall evolution, whose details will then depend on the spectrum and luminosity of the specific burst. We have performed the calculation of the transmittance for different lines of sights with respect to the GRB jet axis (assumed to coincide with the rotation axis of the star), thus exploring the latitudinal variation of the spectra, which include the $\\theta$-dependence of both the CSM profiles as well as of the GRB luminosity. The highest velocity components of the CSM, reaching velocities above 5000 km/s, exist predominantly at small viewing angles (with respect to the rotation axis of the progenitor). This is where the burst is most luminous, and thus ionization timescales are very fast. We do see the presence of several ions at these velocities, but this region becomes completely ionized in a fraction of a second, and an actual observation of the high velocity features seems hopeless, at least for the star considered in this study. Lines of sight at higher latitude are characterized by smaller maximum velocities, with the maximum being around 2800 km/s in the case of a line of sight at $80^\\circ$. {In our model, even at these latitudes, ions} with lines around 1000 km/s cannot survive the entire burst without being photoionized. From a purely theoretical prospective, the time variability of the high velocity lines, and in particular, the transition from a nearly featureless spectrum to high column density lines back to a featureless or weak lined system like the one observed in our study, is a clear signature that the lines are associated with the CSM, and furthermore, that they are due to the destruction of dust in the CSM created by the progenitor star. Practically, however, observational detection of such variability with current instruments is not feasible. Observationally, high-velocity lines have been reported in GRB~021004 by \\citet{Fiore_etal05} after about 16 hr from the trigger. From the analysis presented in this paper, we find that these features, if associated with the free-streaming wind of a progenitor star like the one discussed here, could have survived till these late times only if the burst had a very low luminosity. Since this is not supported by observations (the luminosity of GRB~021004 was rather high), the only way to associate these lines to a wind is by requiring a very special event, in which the progenitor star would possess an unusually strong wind, and it were to explode in a very low density medium. Alternatively, the observed late-time, high-velocity lines could be simply due to intervening absorbers (Chen et al. 2007). {In summary, during} the initial moments of the GRB explosion, the optical absorption spectra present a very rich phenomenology due to the different speeds at which the wind components move away from the progenitor side \\citep{vanMarle_etal08}. If we were able to observe these early phases, we could derive a wealth of information on the properties of the progenitor stars of GRBs. More generally, this would be of great importance to the study of the final stages of stellar evolution. Observationally, however, it is extremely challenging to catch variability on such small timescales, shorter even than the time it usually takes to localize the burst. Therefore we conclude that, if the massive stars progenitors of GRBs are indeed of the kind studied here (as suggested by theoretical investigations), then the high-velocity lines produced in their winds are practically invisible, even for very low-luminosity bursts (such as the $\\theta=80^\\circ$ case that we considered). The only observational signatures of the GRB progenitor stars are imprinted in non-variable absorption lines produced in the termination shock of the wind. Previous detections of high-velocity lines must therefore be due to intervening absorbers, unless the GRB progenitor star had extreme properties well outside of the range considered here." }, "1003/1003.1845_arXiv.txt": { "abstract": "{} {A wide observational campaign was carried out in 2004-2009 aimed to complete the ground-based investigation of Lutetia prior to the Rosetta fly-by in July 2010.} {We have obtained BVRI photometric and V-band polarimetric measurements over a wide range of phase angles, and visible and infrared spectra in the 0.4-2.4 $\\mu$m range. We analyzed them together with previously published data to retrieve information on Lutetia's surface properties.} {Values of lightcurve amplitudes, absolute magnitude, opposition effect, phase coefficient and BVRI colors of Lutetia surface seen at near pole-on aspect have been determined. We defined more precisely parameters of polarization phase curve and showed their distinct deviation from any other moderate-albedo asteroid. An indication of possible variations both in polarization and spectral data across the asteroid surface was found. To explain features found by different techniques we propose that (i) Lutetia has a non-convex shape, probably due to the presence of a large crater, and heterogeneous surface properties probably related to surface morphology; (ii) at least part of the surface is covered by a fine-grained regolith with particle size less than 20 $\\mu$m; (iii) the closest meteorite analogues of Lutetia's surface composition are particular types of carbonaceous chondrites or Lutetia has specific surface composition not representative among studied meteorites.} {} ", "introduction": "Asteroid 21 Lutetia has been extensively observed by different techniques for more than 30 years. Initially the interest in this object was connected with its classification as an M-type asteroid with a possible metallic composition (see Bowell et al. 1978). Since 2004 when Lutetia was selected as a target of the Rosetta mission, the volume of observational data on this asteroid is rapidly growing (see Barucci and Fulchignoni 2009 for a review). On the basis of photometric data obtained in 1962-1998 Torppa et al. (2003) determined the pole coordinates $\\lambda$$_{p}$ =39$^o$ (220$^o$), $\\beta$$_{p}$=3$^o$ and the sidereal rotation period P$_{sid}$=8.165455 h. The shape was found to have some irregular features with rough global dimensions a/b = 1.2 and b/c = 1.2. Recently Drummond et al. (2009) gave new estimations of these parameters including in analysis adaptive optics images of Lutetia at the Keck telescope: $\\lambda$ $_{p}$=49$^o$, $\\beta$$_{p}$=-8$^o$ and a shape of 132x101x76 km with formal uncertainties of 1 km for the equatorial dimensions, and 31 km for the shortest axis. On the basis of spectral and polarimetric observations, three types of meteorites are generally taken into consideration as possible analogues: iron meteorites (Bowell et al. 1978, Dollfus et al. 1979), enstatite chondrites (Chapman et al. 1975, Vernazza et al. 2009) and some types of carbonaceous chondrites, mainly CO3 or CV3 (Belskaya \\& Lagerkvist 1996, Birlan et al. 2004, Barucci et al. 2008, Lazzarin et al. 2009). The main problem in spectral data interpretation is the featureless spectrum of Lutetia. A presence of few minor features in the visible range was reported and interpreted as indicative of aqueous alteration material consistent with carbonaceous chondrites composition (see Lazzarin et al. 2009 and references therein). A 3 $\\mu$m feature associated with hydrated minerals was found by Rivkin et al. (2000). In the emissivity spectra a narrow 10 $\\mu$m emission feature was found (Feierberg et al. 1983, Barucci et al. 2008). It was interpreted as indicative the presence of fine silicate dust (Feierberg et al. 1983). According to Barucci et al. (2008) the emissivity spectrum is similar to that of the CO3 and CV3 carbonaceous chondrites with a grain size less than 20 $\\mu$m. To constrain surface composition knowledge of Lutetia's albedo is considered to be very important. However up to now the diversity in albedo estimations by different techniques is quite large, spanning from 0.1 (Zellner et al. 1976) to 0.22 (Tedesco et al. 2002). Polarimetric method of albedo determination gives contradictory results (Zellner et al. 1976, Gil-Hutton 2007). An accuracy of radiometric albedo strongly depends on adopted absolute magnitude which is not well-determined for Lutetia due to lack of observations at small phase angles. In this paper we present new photometric, polarimetric and spectral observations of Lutetia carried out in 2004-2009. They were aimed to determine absolute magnitude and albedo, and to put additional constrains on surface properties. Their analysis together with previously published data is given. We think that such analysis is important not only for deriving physical characteristics of this particular asteroid but first of all for checking the efficiency of remote techniques in the study of atmosphereless bodies. ", "conclusions": "On the basis of a detailed analysis of new photometric, polarimetric and spectral data on the asteroid 21 Lutetia, together with observational data available in literature, we can draw some conclusions which can be checked during Rosetta fly-by: 1. Lutetia has a non-convex shape, probably due to the presence of a large crater, and heterogeneous surface properties probably due to variations of texture and/or mineralogy related to surface morphology. 2. At least part of Lutetia's surface is covered by regolith composed of particles having a mean grain size less than 20 $\\mu$m. 3. The closest meteorite analogues of Lutetia's surface composition are particular types of carbonaceous chondrites {CO, CV, CH). It is also possible that Lutetia has specific surface composition not representative among studied meteorites or has a mixed mineralogy, e.g. due to surface contamination. Flyby observations of Lutetia by the Rosetta spacecraft in July 2010 will provide ground truth for Earth-based remote sensing." }, "1003/1003.0673_arXiv.txt": { "abstract": "We search a sample of photometric luminous red galaxies (LRGs) measured by the Sloan Digital Sky Survey (SDSS) for a quadrupolar anisotropy in the primordial power spectrum, in which $P(\\veck)$ is an isotropic power spectrum $\\bar P(k)$ multiplied by a quadrupolar modulation pattern. We first place limits on the 5 coefficients of a general quadrupole anisotropy. We also consider axisymmetric quadrupoles of the form $P(\\veck) = \\bar P(k)\\{ 1 + g_*[(\\hatk\\cdot\\hatn)^2-\\frac13]\\}$ where $\\hatn$ is the axis of the anisotropy. When we force the symmetry axis $\\hatn$ to be in the direction $(l,b)=(94^\\circ,26^\\circ)$ identified in the recent Groeneboom {\\em et~al.} analysis of the cosmic microwave background, we find $g_*=0.006\\pm0.036$ ($1\\sigma$). With uniform priors on $\\hatn$ and $g_*$ we find that $-0.410$} \\\\ Y_{L0}&\\mbox{if $M=0$} \\\\ \\frac{(-1)^M}{{\\rm i}\\sqrt{2}}(Y_{LM}^*-Y_{LM})&\\mbox{if $M<0$}\\,; \\end{array} \\right. \\end{eqnarray} these are easily seen to obey the usual orthonormality rules, but have the advantage of making the $g_{LM}$ coefficients real. The expressions for $L=2$ are given in Appendix \\ref{A:realhar}. The purpose of this paper is to measure or constrain the anisotropy using large scale structure data. Given the recent debate over the detection of quadrupolar anomalies in WMAP, and the evidence that the signal is contaminated by systematic effects \\cite{G09, Bennett:2010jb}, it is worth using other datasets as well to constrain models with anisotropic power. In this paper we will assume for simplicity that $g_{LM}$ is scale-invariant. This is both a simplifying assumption, but is also a good first approximation in at least some classes of modified inflationary models \\cite{Ackerman:2007nb}. We will also focus only on the quadrupole anisotropy $g_{2M}$; this is the phenomenologically simplest type of anisotropy allowed, and also emerges from anisotropic inflation models in the limit of very weak anisotropy \\cite{Ackerman:2007nb}. Galaxy surveys probe matter fluctuations because on large scales, the galaxy density is related to the matter density in accordance with a linear bias model: \\begin{eqnarray} \\label{E:bias} \\VEV{\\delta_g(\\veck)\\delta_g^*(\\veck')} = (2\\pi)^3\\delta_D(\\veck-\\veck')b_g^2P(\\veck)\\, , \\end{eqnarray} where $\\delta_g(\\veck)$ is the Fourier amplitude of the fractional galaxy density perturbation, and $b_g$ is the linear galaxy bias. The galaxy survey probe has been used to estimate $P(\\veck)$ by stacking the measured angular matter power spectra $C_l$ in eight photometric redshift slices ranging from $z=0.2 \\mbox{ to } 0.6$ \\cite{Padmanabhan:2006ku}. By performing a similar analysis and including the anisotropy parameters $g_{2M}$ in the power spectrum, we can use galaxy surveys to estimate quadrupole anisotropy while assuming fiducial values for the other cosmological parameters. The plan of our paper is as follows: In Section \\ref{S:sample} we describe the SDSS data used and why we choose LRGs to trace the galaxy distribution. Section \\ref{S:form} calculates the angular correlations statistical anisotropy produces in galaxy surveys and constructs estimators of the $g_{2M}$s and other systematic power spectrum variations. We present estimates of these parameters in Section \\ref{S:result}, and in Section \\ref{S:conclude} we present our conclusions. Wherever not explicitly mentioned, we assume a flat $\\Lambda$CDM cosmology with $\\Omega_M = 0.3$, $\\Omega_b = 0.05$, $h = 0.7$, $n_s = 1.0$, and $\\sigma_8 = 0.9$. Since ours is a search for anisotropy, small changes in the cosmology will result only in changes in the calibration of the $g_{2M}$ estimator; they do not alter the null hypothesis. ", "conclusions": "\\label{S:conclude} We have conducted a search for statistical anisotropy in the galaxy distribution. Statistical anisotropy can manifest from the direction-dependent primordial power spectrum shown in Eq.~\\ref{E:powerspectrum} with the magnitude of the anisotropy parametrized by $g_{LM}$. This phenomenon causes the angular galaxy power spectrum $C_{g,l}$ to be generalized by $D_{g,ll'}^{LM}$, which includes $g_{LM}$. We used estimators formulated by Padmanabhan {\\em et al.}~\\cite{Padmanabhan:2006ku} and a sample of LRGs from SDSS to search for evidence of quadrupolar anisotropy parametrized by $g_{2M}$. We found $g_{2M}$ for all $M$ to be within $2\\sigma$ of zero. Using our estimates of $g_{2M}$ and assuming a symmetry axis in the direction $(l,b)=(94^\\circ,26^\\circ)$, we calculated the anisotropy amplitude $g_*=0.006\\pm0.036\\,(1\\sigma)$. This confirms that the previously identified anisotropy in the WMAP maps (already believed to be a systematic effect) is not of primordial origin. When marginalizing over the symmetry axis direction and assuming a uniform prior for $g_*$, we constrain $-0.410.5$~\\Ms \\citep{Kroupa2001}.}, stars in the mass range 25--120~\\Ms{} account for only 10\\% of the mass of the whole first stellar generation \\citep[see][]{DecressinCharbnnel2007} and their slow winds account for 2.5\\% of the mass only. If pollution is due to AGB stars, a similar constraint arises: the wind released by stars between 5 and 6.5~\\Ms{}, for which nucleosynthesis agrees with the observations according to \\citet{VenturaDantona2008a}, represents less than 3\\%. After taking into account the possible dilution of these slow winds with the pristine gas present in the ISM to explain the observed Li abundance variation \\citep{DecressinCharbnnel2007}, we find that the mass available to form the second generation low-mass stars compared to the first generation of low-mass stars is only about 10\\%. This is in sharp contrast with observations, which show that more than half and up to 85\\% of the stars in GCs are second generation stars \\citep{PrantzosCharbonnel2006,CarrettaBragaglia2008}, i.e., which show anticorrelations in light elements. Thus a rather extreme reduction of the first generation stars relative to the second generation stars is needed to reproduce the observations. However massive binaries have recently been proposed as polluters of the proto-GC by \\citet{deMinkPols2009}. In this case the mass-budget is more favourable as more slow winds are ejected and more second generation stars are formed. This could help to reduce the fraction between first and second generation stars found in the present paper therefore supporting a high-mass star pollution scenario. One possible way to reconcile the pollution scenario with the observations is to consider a top-heavy initial mass function (IMF) of first generation stars. In the case of pollution by fast-rotating massive stars, an IMF slope as flat as 1.55 (compared to the canonical value of 2.3) is required to reproduce the high number of stars with abundance anomalies in the cluster NGC~6752 \\citep{DecressinCharbnnel2007}, whereas the AGB scenario requires an even flatter IMF slope \\citep[see][]{PrantzosCharbonnel2006}. A second way to reconcile the pollution scenario with observations is to consider that first generation stars are preferentially lost from the cluster during its evolution so that an initially relatively small population of second generation stars can become the dominant population after several Gyr. To allow this preferential loss of first generation stars requires the GCs to be initially mass-segregated (i.e., that more massive stars occupy the central part of the clusters). In this case the matter released in the disks of massive stars is more concentrated in the cluster centre, and second generation stars are born in the centre while first generation stars are present throughout the cluster. The viability of a self-enrichment scenario by fast-rotating massive stars has been recently explored by \\citet[Paper I]{DecressinBaumgardt2008}.\\defcitealias{DecressinBaumgardt2008}{Paper I} They have shown that first generation low-mass stars are preferentially lost from the cluster, which is assumed to be initially in dynamical equilibrium and mass-segregated, before two-body relaxation induces a spread of second generation stars and a full mixing of the cluster. \\citet{DErcoleVesperini2008} find similar results with the AGB scenario. Afterwards, the evolution is smoother and the variation of the fraction of second generation stars takes longer. Any radial difference between first and second generation stars is erased after 10-12~Gyr of evolution as the cluster relaxation time (a few Gyr) is much shorter than the age of the clusters.\\footnote{The only exception is the GC $\\omega$~Cen, for which the relaxation time at the center is comparable to its age. Indeed in this cluster stars on the blue main sequence (i.e., He-rich) are found more centrally concentrated than red main sequence stars \\citep{VillanovaPiotto2007,BelliniPiotto2009}.} In \\citetalias{DecressinBaumgardt2008} we show that even if the relaxation-driven evaporation increases the fraction of second generation (which harbour abundance anomalies) to about 25\\%, this ratio remains too low to fully explain the observations (between 50--85\\%, \\citealp{CarrettaBragaglia2009}). The increase of the fraction of second generation stars mainly occurs in the early times and points towards the high sensitivity of the fraction of second generation stars on cluster dynamics. In this paper, we aim to quantify the increase of the fraction of second generation stars to the total number of low-mass stars by another dynamical mechanism not taken into account in the above studies, namely the effect of primordial gas expulsion (i.e., the fast ejection of the remaining gas left by star formation after the onset of supernovae). Gas expulsion can strongly modify the total binding energy of the cluster and can lead to an efficient loss of first generation stars from the cluster. We emphasise that we discuss generic properties of gas expulsion models based on the simplified assumption that a cluster contains only two stellar generations with the same [Fe/H]. Multiple populations with different [Fe/H] would require other physical mechanisms, whereby notably gas accretion form the surrounding inter stellar medium \\citep{PflammAltenburgKroupa2009} may play a role, and recycling of SN ejecta also \\citep{TenorioTagleWunsch2007}. In \\S~2 we present the N-body models used in this study. Then our results are discussed in \\S~3. In \\S~4 we present a complete scenario for the evolution of GCs and our conclusions are in \\S~5. ", "conclusions": "In this paper we have studied the influence of primordial gas expulsion by supernovae during the early dynamical evolution of globular clusters. in the context of clusters with two chemically and dynamically distinct stellar populations. In particular we investigate if this dynamical process can explain the high number of observed second generation stars which harbour abundance anomalies in light elements. We deduce the following: \\begin{itemize} \\item If the two populations have a different radial extent with second generation stars more concentrated, primordial gas expulsion is able to expel most of the first generation stars while most second generation stars can be retained. \\item For a given fractional mass loss by the cluster, the fraction of second generation stars is nearly independent of the gas expulsion parameters (see \\S~3.2). \\item The final observed fraction of second generation stars can constrain the initial properties of GCs as this fraction is highest for clusters with SFE around 0.3-0.33, with concentrated clusters relative to the tidal field, and with a fast timescale for gas expulsion relative to the crossing time. \\item We infer proto-GC cloud masses of several $10^6$~\\Ms{} and up to $9\\times 10^6$~\\Ms{} for clusters which show a large fraction of chemically different second generation stars like NGC~6752. Their initial half-mass radii are in the range of $\\sim1$--3~pc (4--5~pc for the most massive cases). \\item It is possible to reproduce the fraction of second generation stars in present-day GCs through cluster dynamical processes by combining gas expulsion and tidal stripping during long-term evolution of initially mass-segregated clusters. \\item The primordial gas expulsion process can also be at the origin of the observational trend observed by \\citet{Carretta2006}, who shows that clusters with large orbital period and with high orbital inclinations relative to the Galactic plane produce more extended O-Na and Mg-Al anticorrelations. \\end{itemize}" }, "1003/1003.1169_arXiv.txt": { "abstract": "We investigate high-magnification events caused by planets in wide binary stellar systems under the strong finite-source effect, where the planet orbits one of the companions. From this investigation, we find that the pattern of central perturbations in triple lens systems commonly appears as a combination of individual characteristic patterns of planetary and binary lens systems in a certain range where the sizes of the caustics induced by a planet and a binary companion are comparable, and the range changes with the mass ratio of the planet to the planet-hosting star. The inside and outside edge regions of a circle with a radius corresponding to that of the source star and its center located at the center of the caustic, show the binary-lensing pattern, while the outside region of the circle shows the planetary-lensing pattern. Specially, we find that because of this central perturbation pattern, the characteristic feature of high-magnification events caused by the triple lens systems appears in the residual from the single-lensing light curve despite the strong finite-source effect, and it is discriminated from those of the planetary- and binary-lensing events and thus can be used for the identification of the existence of both planet and binary companion. This characteristic feature is a simultaneous appearance of two features. First, double negative-spike and single positive-spike features caused by the binary companion appear together in the residual, where the double negative spike occurs at both moments when the source enters and exits the caustic center and the single positive spike occurs at the moment just before the source enters into or just after the source exits from the caustic center. Second, the magnification excess before or after the single positive-spike feature is positive due to the planet, and the positive excess has a remarkable increasing or decreasing pattern depending on the source trajectory. ", "introduction": "The microlensing signal of a planet is a short-duration perturbation on the smooth standard light curve of the primary-induced lensing event that occurred on a background source star. To detect extrasolar planets using microlensing, survey and follow-up observations are now being carried out toward the Galactic bulge field. The survey observations (OGLE: Udalski 2003, MOA: Bond et al. 2002) monitor a large area of sky and alert ongoing events by analyzing data in real time, while the follow-up observations ($\\mu$FUN: Dong et al. 2006, PLANET: Albrow et al. 2001, RoboNet: Burgdorf et al. 2007) intensively monitor the alerted events. For a planetary-lensing event, the perturbation of the lensing light curve is induced by two sets of disconnected caustics that are composed of the central and planetary caustics. The perturbation induced by the central caustic always occurs close to the peak of the lensing light curve, while the perturbation induced by the planetary caustic can occur at any part of the light curve. The lensing event caused by the central caustic becomes a high-magnification event with central perturbation and this event is very sensitive for the detection of a planet \\citep{griest98}. The current follow-up observations thus focus on high-magnification events. However, the lensing event by a very close or a very wide binary can also produce central perturbation of the high-magnification event. Fortunately, \\citet{han08a} found that high-magnification events with a double-peak structure caused by a planet and a binary companion can be immediately distinguished by the shape of the interpeak region of the light curves. The size of the source star becomes important in high-magnification events because the source star passes very close to the central caustic. In planetary-lensing events where the source diameter is considerably larger than the central caustic and thus the finite-source effect is strong, the central perturbation is greatly buried and the resulting light curve appears like that of a single-lensing event induced by the primary star. Thus, planetary- and binary-lensing events with a double-peak structure affected by the strong finite-source effect cannot be distinguished any more by the shape of the interpeak region of the light curves. High-magnification events with strong finite-source effects have been recently reported by \\citet{subo09} and \\citet{janczak10}. They reported that, on casual inspection, these events appear as single-lensing events with pronounced finite-source effects, but a more detailed analysis reveals they are planetary-lensing events with a buried signature of the planet. \\citet{hankim09} investigated the planetary-lensing signals of high-magnification events under strong finite-source effects. They found that the characteristic features of planetary-lensing events commonly appear in the residuals and thus they can be used for the diagnosis of the existence of the planet. However, they noted that the characteristic residual features can be also produced by very close or very wide binary-lensing events and thus the existence of these features does not necessarily confirm the existence of the planet. We expect that many more high-magnification events with buried signatures of the planet and binary companion will be detected by future observations with high cadence monitoring including the ground-based observations (e.g., Korea Microlensing Telescope Network (KMTNet) Project: B.-G. Park 2009, private communication) and space-based observations (e.g, {\\it Microlensing Planet Finder} ({\\it MPF}): Bennett et al. 2004). Therefore, it is very important to distinguish planetary- and binary-lensing events with the buried signature. \\citet{han09} found that some planetary- and binary-lensing events with the buried signature can be distinguished by the difference of the characteristic features of the individual residual patterns. For triple-lensing events caused by planets in binary stellar systems, where the sizes of the caustics induced by a planet and a binary companion are similar to each other, the signatures of both planet and binary companion can show up in the lensing light curves, and thus it is possible to detect planets in binary stellar systems \\citep{lee08}. However, if the events are strongly affected by finite-source effects, both signatures are greatly washed out and thus do not obviously appear in the light curves. After all, triple-lensing events are very difficult to be identified as events caused by planets in binary systems. Fortunately, since current follow-up observations have a high photometric precision ($\\sim 1 \\%$) at the peaks of high-magnification events and it is expected that future observations would have more high photometric precision, it could be feasible for the identification of the planetary signatures in binary stellar systems affected by the strong finite-source effect. In this paper, we investigate high-magnification events caused by planets in binary stellar systems to find out whether the signatures of both planet and binary companion can be identified despite the strong finite-source effect. The paper is organized as follows. In Section 2, we briefly describe the central caustics caused by a planet and a binary companion. In Section 3, we investigate the central perturbation patterns of planets in binary systems under the strong finite-source effect. We also compare the residual patterns of the events caused by the triple lens systems with those of planetary- and binary-lensing events. We summarize the results and conclude in Section 4. ", "conclusions": "We have investigated high-magnification events caused by planets in binary stellar systems under the strong finite-source effect, especially for wide-separation binaries with low mass planets. From this study, we found that the central perturbation patterns of the triple lens systems commonly appear as the combination of the characteristic perturbation patterns of planetary and binary lens systems in a certain range that changes as the planet/primary mass ratio. The inside and outside edge regions of a circle with a radius corresponding to that of the source and its center located at the caustic center show the binary-lensing pattern, while the outside region of the circle shows the planetary-lensing pattern. In particular, we found that due to this central perturbation pattern, the characteristic feature of the high-magnification events caused by the triple lens systems appears in the residual from single-lensing light curve, and it is distinguished from those of the planetary- and binary-lensing events and thus can be used for the diagnosis of the identification of the existence of both planet and binary companion. This characteristic feature is a simultaneous appearance of the following two features. First, double negative-spike and single positive-spike features caused by the binary companion appear together in the residual, where the double negative spike occurs at both moments when the source enters and exits the caustic center, and the single positive spike occurs at the moment just before the source enters into or just after the source exits from the caustic center. Second, the excess before or after the single positive spike in the residual is positive due to the planet, and the positive excess has a remarkably increasing or decreasing pattern depending on the source trajectory. We thank I. A. Bond for making helpful comments. \\begin{figure}[t] \\epsscale{1.0} \\plotone{fig1.eps} \\caption{\\label{fig:one} Magnification excess maps of planets in binary stellar systems with various planet/primary mass ratios and size ratios of the caustics induced by a planet and a binary companion. The coordinates ($\\hat{\\xi},\\hat{\\eta}$) represent the axes that are parallel with and normal to the binary axis and are centered at the effective position of the primary star. Here the notation with the hat represents the length scale normalized by the Einstein radius of the primary, $\\thetaeone$. In this map, the companion/primary mass ratio and projected primary-planet separation are $(\\qb, \\spp) = (0.5, 1.23)$, the position angle of the planet measured from the binary axis is $50^\\circ$, $\\sbb$ is the projected separation between the binary components in units of the Einstein radius corresponding to the total mass of the lens system, $\\thetae$, $\\qp$ is the planet/primary mass ratio, and $\\delxip$ and $\\delxib$ are the sizes of the caustics induced by the planet and binary companion, respectively. The color changes into darker scales when the excess is $|\\epsilon| = 2\\%,\\ 4\\%,\\ 8\\%,\\ 16\\%,\\ 32\\%$, and $64\\%$, respectively. The dashed circle represents the moment when the source star is located at the center of the caustic. The straight lines with arrows represent the source trajectories. }\\end{figure} \\begin{figure}[t] \\epsscale{1.0} \\plotone{fig2.eps} \\caption{\\label{fig:two} Light curves for the source trajectories presented in Figure 1. In the upper panel, solid and dashed curves represent the light curves of the triple- and single-lensing events, respectively. The lower panel shows the residuals from the single-lensing light curve. In the lower panel, the horizontal line indicates the magnification excess of $|\\epsilon|=0.0$. The arrows in the left bottom panel represent the moments when the source enters and exits the dashed circle located at the caustic center in Figure 1. }\\end{figure} \\begin{figure}[t] \\epsscale{1.0} \\plotone{fig3.eps} \\caption{\\label{fig:three} Magnification excess maps of planet-binary, stellar binary, and planetary lens systems together with the residuals resulting from the source trajectories presented in the individual maps. Solid, dashed, and dashed-dot curves represent the residuals of the planet-binary (triple), binary-, and planetary-lensing events, respectively. Shaded regions represent the perturbations occurring over the outside edge of the dashed circle. }\\end{figure}" }, "1003/1003.3446_arXiv.txt": { "abstract": "We present HST/WFC3 grism spectroscopy of the brightest galaxy at $z>1.5$ in the GOODS-South WFC3 Early Release Science grism pointing, covering the wavelength range $0.9\\,\\mu$m -- 1.7\\,$\\mu$m. The spectrum is of remarkable quality and shows the redshifted Balmer lines H$\\beta$, H$\\gamma$, and H$\\delta$ in absorption at $z=1.902 \\pm 0.002$, correcting previous erroneous redshift measurements from the rest-frame UV. The average rest-frame equivalent width of the Balmer lines is 8\\,\\AA\\,$\\pm 1$\\,\\AA, which can be produced by a post-starburst stellar population with a luminosity-weighted age of $\\approx 0.5$\\,Gyr. The $M/L$ ratio inferred from the spectrum implies a stellar mass of $(4\\pm 1) \\times 10^{11}$\\,\\msun. We determine the morphology of the galaxy from a deep WFC3 $H_{160}$ image. Similar to other massive galaxies at $z\\sim 2$ the galaxy is compact, with an effective radius of $2.1\\pm 0.3$\\,kpc. Although most of the light is in a compact core, the galaxy has two red, smooth spiral arms that appear to be tidally-induced. The spatially-resolved spectroscopy demonstrates that the center of the galaxy is quiescent and the surrounding disk is forming stars, as it shows H$\\beta$ in emission. The galaxy is interacting with a companion at a projected distance of 18\\,kpc, which also shows prominent tidal features. The companion has a slightly redder spectrum than the primary galaxy but is a factor of $\\sim 10$ fainter and may have a lower metallicity. It is tempting to interpret these observations as ``smoking gun'' evidence for the growth of compact, quiescent high redshift galaxies through minor mergers, which has been proposed by several recent observational and theoretical studies. Interestingly both objects host luminous AGNs, as indicated by their X-ray luminosities, which implies that these mergers can be accompanied by significant black hole growth. This study illustrates the power of moderate dispersion, low background near-IR spectroscopy at HST resolution, which is now available with the WFC3 grism. ", "introduction": "\\noindent The formation history of massive galaxies is not well understood. Present-day galaxies with stellar masses $\\gtrsim 3\\times 10^{11}$\\,\\msun\\ are typically giant elliptical galaxies in the centers of galaxy groups. These galaxies have old stellar populations and follow tight scaling relations between their velocity dispersions, sizes, surface brightnesses, line strengths, and other parameters (e.g., {Djorgovski} \\& {Davis} 1987; {Thomas} {et~al.} 2005). At redshifts $z\\sim 2$ massive galaxies form a more complex population. A fraction of the population is forming stars at a high rate, as determined from their brightness in the rest-frame UV or IR, emission lines such as H$\\alpha$, and other indicators (e.g., {Steidel} {et~al.} 1996; {Blain} {et~al.} 2002; {Rubin} {et~al.} 2004; {Papovich} {et~al.} 2006, and many other studies). However, others have no clear indications of ongoing star formation and have spectral energy distributions (SEDs) characterized by strong Balmer- or 4000\\,\\AA\\ breaks (e.g., {Daddi} {et~al.} 2005; {Kriek} {et~al.} 2006). The existence of these ``quiescent'' galaxies at this early epoch is in itself remarkable, and provides constraints on the accretion and thermodynamics of gas in massive halos at $z>2$ (e.g., Kere{\\v s} et al.\\ 2005; {Dekel} \\& {Birnboim} 2006). What is perhaps even more surprising is that these galaxies are structurally very different from early-type galaxies in the nearby Universe: their effective radii are typically 1--2\\,kpc, much smaller than nearby giant ellipticals (e.g., {Daddi} {et~al.} 2005; {Trujillo} {et~al.} 2006; {van Dokkum} {et~al.} 2008; {Cimatti} {et~al.} 2008). Several explanations have been offered for the dramatic size difference between local massive galaxies and quiescent galaxies at high redshift. The simplest is that observers underestimated the sizes and/or overestimated the masses. Although subtle errors are almost certainy present in the interpretation of the data, recent studies suggest that it is difficult to change the sizes and the masses by more than a factor of 1.5, unless the IMF is altered (e.g., Muzzin et al.\\ 2009, Cassata et al.\\ 2010, Szomoru et al.\\ 2010). Other explanations include extreme mass loss due to a quasar-driven wind ({Fan} {et~al.} 2008), strong radial age gradients leading to large differences between mass-weighted and luminosity-weighted ages ({Hopkins} {et~al.} 2009; {La Barbera} \\& {de Carvalho} 2009), star formation due to gas accretion ({Franx} {et~al.} 2008), and selection effects (e.g., {van Dokkum} {et~al.} 2008; {van der Wel} {et~al.} 2009). Perhaps the most plausible mechanism for bringing the compact $z\\sim 2$ galaxies onto the local mass-size relation is (minor) merging (e.g., {Bezanson} {et~al.} 2009; {Naab}, {Johansson}, \\& {Ostriker} 2009; {van Dokkum} {et~al.} 2010; Carrasco, Conselice, \\& Trujillo 2010). Numerical simulations predict that such mergers are frequent ({Guo} \\& {White} 2008; {Naab} {et~al.} 2009); furthermore, they may lead to stronger size growth than mass growth ({Bezanson} {et~al.} 2009). From an analysis of mass evolution at fixed number density, {van Dokkum} {et~al.} (2010) infer that massive galaxies have doubled their mass since $z=2$, and suggest that $\\sim 80$\\,\\% of this mass growth can be attributed to mergers. \\setcounter{figure}{1} \\noindent \\begin{figure*}[t] \\epsfxsize=17cm \\epsffile[86 528 563 688]{imspec.eps} \\caption{\\small HST ACS and WFC3 images of \\gala\\ and its companion \\galb. The ACS color image was created from the $B_{435}$, $V_{606}$, and $z_{850}$ bands, and the WFC3 image from the $Y_{098}$, $J_{125}$, and $H_{160}$ bands. \\gala\\ has a compact core and spiral arms, which may be the result of an interaction with \\galb. Red circles are the locations of X-ray sources in the Luo et al.\\ (2008) catalog, with the size of the circles indicating the uncertainties in the positions. Both galaxies host an AGN. The SEDs of the two galaxies (from Wuyts et al.\\ 2008) are shown in the right-most panel. The galaxies are both red and have broadly similar SEDs. \\vspace{0.0cm} \\label{imspec.plot}} \\end{figure*} Although qualitatively consistent with observations and theory the minor merger scenario currently has little direct evidence to support it. It is also not clear whether properties other than sizes and masses are easily explained in this context; one of the open questions is why present-day elliptical galaxies are so red and homogeneous if half of their mass was accreted from the general field at relatively recent times. Ideally we would identify and study the infalling population directly at high redshift, but so far this has been hampered by the limitations of ground-based spectroscopy and ground- and space-based near-IR imaging. In this {\\em Letter}, we use the exquisite WFC3 grism on the {\\em Hubble Space Telescope} (HST), in combination with WFC3 imaging, to study the environment of a quiescent compact galaxy at $z=1.9$. As we show below, the observations presented here provide the first direct evidence for minor mergers as a mechanism for the growth of compact galaxies at high redshift. We use $H_0=70$\\,\\kms\\,Mpc$^{-1}$, $\\Omega_m=0.3$, and $\\Omega_{\\Lambda}=0.7$. Magnitudes are on the AB system. ", "conclusions": "\\noindent The WFC3 grism and imaging data of \\gala\\ may provide ``smoking gun'' evidence for minor mergers as an important growth mechanism of massive galaxies: \\gala\\ is a massive, compact galaxy at $z\\sim 2$ which is interacting with a $\\sim 10\\times$ less massive companion. The quiescent spectrum of the primary galaxy is qualitatively consistent with the spectra of other compact high redshift galaxies and with the old stellar ages of present-day early-type galaxies. This mode of growth has been proposed by several recent studies to explain the size difference between massive galaxies at high redshift and low redshift (e.g., {Bezanson} {et~al.} 2009; {Naab} {et~al.} 2009). Nearby ellipticals have gradients in their color and metallicity, such that they are bluer and more metal-poor at larger radii (e.g., {Franx}, {Illingworth}, \\& {Heckman} 1989). Interestingly, we can begin to address the origin of these gradients with the kind of data that we are now getting from HST. The relatively strong oxygen lines and weak H$\\beta$ of the infalling galaxy imply $\\log R_{23}\\sim 1$, and a metallicity that is $\\gtrsim 1/3$ times the Solar value ({Pilyugin} \\& {Thuan} 2005). The spectrum extracted from the disk of \\gala\\ has, by contrast, no detected oxygen lines and an unambiguous detection of H$\\beta$. It has $\\log R_{23}\\lesssim 0$, which implies a Solar or super-Solar metallicity. Qualitatively these results are consistent with the idea that the metallicity gradients of elliptical galaxies reflect a gradual increase with radius in the fraction of stars that came from infalling low-mass satellites. The apparent absence of star formation in the central regions of \\gala\\ might be related to its active nucleus. It has been suggested by many authors that AGN could prevent gas cooling and star formation (e.g., {Croton} {et~al.} 2006) and in this context the observed properties of \\gala, such as the lower limit on the ratio of its X-ray luminosity to [O\\,{\\sc iii}] and H$\\beta$, may provide constraints on the mechanism(s) of AGN feedback (see also {Fiore} {et~al.} 2008; {Kriek} {et~al.} 2007, 2009). In any case, the fact that both interacting galaxies host an AGN is remarkable, as it demonstrates that their black holes are undergoing a ``growth spurt'' prior to their merger. We note here that the only indication of the AGNs in the optical and near-IR is a faint emission line in the VIMOS spectra of \\gala, which we now identify\\footnote{The three erroneous redshifts for \\gala\\ were not due to a mis-identification of this line; the line was not recognized as a real feature in the GOODS-VIMOS analysis of the VIMOS spectrum.} as C\\,{\\sc iv}. There are several important caveats, uncertainties, and complications. First, \\gala\\ is not only growing through the accretion of \\galb, but also through star formation. There is evidence for star formation in the companion (although its emission lines could be influenced by its active nucleus) and also in the spiral arms of \\gala. In most models such ``residual'' star formation takes place in the center of the most massive galaxy (see, e.g., {Naab} {et~al.} 2009), but that is in fact the only place where we do {\\em not} see evidence for star formation. We note, however, that because of the large mass of \\gala\\ the specific star formation rate of the entire system is low at SFR\\,/\\,M$_{\\rm stellar}$\\,$\\lesssim 10^{-10}$\\,yr$^{-1}$. Second, although the spectrum of \\gala\\ resembles those of the compact galaxies studied in {Kriek} {et~al.} (2006) and {van Dokkum} {et~al.} (2008), the galaxy formed its stars at significantly lower redshift. As shown in \\S\\,3 its star formation rate probably was $\\sim 500$\\,\\msun\\,yr$^{-1}$ as recently as $150$\\,Myr prior to the epoch of observation, i.e., at $z\\approx 2$. It is therefore not a direct descendant of quiescent galaxies at $z\\sim 2.3$. Interestingly, star forming galaxies at $z>2$ are typically larger than \\gala\\ in the rest-frame optical (e.g., {Toft} {et~al.} 2007), which may imply that \\gala\\ is unusual or that a significant fraction of the star formation in massive galaxies at $z\\sim 2.5$ takes place in heavily obscured, compact regions. Third, the fact that the time since the truncation of star formation is similar to the dynamical time calls into question whether we are witnessing a ``two-stage'' galaxy formation process, with steady accretion of satellite galaxies following an initial highly dissipational star formation phase (e.g., {Naab} {et~al.} 2009; {Dekel} {et~al.} 2009). An alternative interpretation is that the companion galaxy is somehow related to the truncation, for example by triggering the AGN in \\gala\\ $\\sim 150$\\,Myr ago. Numerical simulations that aim to reproduce both the 2D spectrum and the morphological features might shed some light on these issues. As illustrated in this {\\em Letter} the WFC3 camera on HST has opened up a new regime of detailed spectroscopic and imaging studies of high redshift galaxies. The quality of the rest-frame optical continuum spectra shown in Fig.\\ \\ref{spec.plot} greatly exceeds what can be achieved from the ground (see, e.g., Kriek et al.\\ 2009), and the grism provides simultaneous spectroscopy of all 200-300 objects with $H\\lesssim 23$ in the WFC3 field. Future WFC3 spectroscopic and imaging surveys over large areas have the potential to robustly measure the evolution of galaxies over the redshift range $1$400\\,km/s), which could be an alternative origin for hypervelocity stars (HVSs), which are stars with a velocity so great that they are able to escape the gravitational pull of the Galaxy. Because SN Ia birthrates from the He star donor channel increase with metallicity, HVSs from the SN explosion scenario are more likely discovered in the high metallicity environment. Currently, some observations support the existence of WD + He star systems (e.g., KPD 1930+2752, V445 Pup, and HD 49798 with its WD companion), which are candidates of SN Ia progenitors. (1) Maxted et al. (2000) suggested that KPD 1930+2752 is likely to eventually result in a merger and produce an SN Ia (see also Geier et al. 2007). However, the DD model is not supported theoretically. Meanwhile, KPD 1930+2752 may also produce an SN Ia via the SD model, but the parameters of the binary system are not located in the contours of the He donor star channel for producing SNe Ia, i.e., KPD 1930+2752 will not produce an SN Ia via the SD model. (2) V445 Pup is an He nova (Ashok \\& Banerjee 2003; Kato \\& Hachisu 2003). Kato et al. (2008) presented a free-free emission dominated light curve model of V445 Pup, based on the optically thick wind theory (Kato \\& Hachisu 1994; Hachisu et al. 1996). The light curve fitting in their study shows that the mass of the WD is more than $1.35\\,M_{\\odot}$, and half of the accreted matter remains on the WD, leading to the mass increase of the WD. Thus, Kato et al. (2008) suggested that V445 Pup is a strong candidate of SN Ia progenitors (see also Woudt et al. 2009). However, we still do not know the orbital period of the binary system and the mass of the He donor star so far. This needs further observations of V445 Pup after the dense dust shell disappears. (3) HD 49798 is a H depleted subdwarf O6 star and also a single-component spectroscopic binary with an orbital period of 1.548\\,d (Thackeray 1970; Stickland \\& Lloyd 1994), which contains a X-ray pulsating companion (RX J0648.0-4418; Israel et al. 1995). The X-ray pulsating companion is suggested to be a massive WD (Bisscheroux et al. 1997). Based on the pulse time delays and the binary system's inclination, constrained by the duration of the X-ray eclipse, Mereghetti et al. (2009) recently derived the masses of the two components. The corresponding masses are 1.50$\\pm$0.05$\\,M_{\\odot}$ for HD 49798 and 1.28$\\pm$0.05$\\,M_{\\odot}$ for the WD. According to our binary evolution model, we found the massive WD can increase its mass to the Ch mass in future evolution. Thus, HD 49798 with its WD companion is a likely candidate of SN Ia progenitors. The He star donor channel with different metallicities can produce the young SNe Ia with delay times $\\sim$45$-$220\\,Myr. The young population of SNe Ia may have an effect on models of galactic chemical evolution, since they would return large amounts of iron to the interstellar medium earlier than previously thought. Especially, the high metallicity environments are much earlier to return iron to the interstellar medium, as SNe Ia from the He star donor channel occur systemically earlier for a high $Z$. In future investigations, we will explore the detailed influence of the young SNe Ia with different metallicity environments on the chemical evolution of stellar populations." }, "1003/1003.5489_arXiv.txt": { "abstract": "The impact that unrecognised differences in the chemical patterns of Galactic globular clusters have on their relative age determinations is studied. The two most widely used relative age-dating methods, horizontal and vertical, together with the more recent relative MS-fitting method, were carefully analyzed on a purely theoretical basis. The BaSTI library was adopted to perform the present analysis. We find that relative ages derived using the horizontal and vertical methods are largely dependent on the initial He content and heavy element distribution. Unrecognized cluster-to-cluster chemical abundance differences can lead to an error in the derived relative ages as large as $\\sim$0.5 (or $\\sim$6 Gyr if an age of 12.8 Gyr is adopted for normalization), and even larger for some extreme cases. It is shown that the relative MS-fitting method is by far the age-dating technique for which undetected cluster$-$to$-$cluster differences in the He abundance have less impact. Present results are used in order to pose constraints on the maximum possible spread in the He and CNONa elements abundances on the basis of the estimates - taken from the literature - of the Galactic globular clusters relative age dispersion obtained with the various relative age-dating techniques. Finally, it is shown that the age--metallicity relation found for young Galactic globular clusters by the GC Treasury program is a real age sequence and cannot be produced by variations in the He and/or heavy element distribution. ", "introduction": "Galactic globular clusters (GGCs) are the most ancient objects known for which reliable ages can be determined, and as the Universe cannot be younger than the oldest objects it contains, GGCs provide one of the most robust constraints that we have on cosmological models. However, absolute GGCs ages have to be estimated in order to apply this constraint. Although significant improvements in the absolute GGCs age estimates have been obtained in the last decade, they are still affected by both observational and theoretical uncertainties \\citep{VSB96, Cass98, Cass09} at the $\\approx20$\\% level. Nevertheless, it is possible to determine relative GGC ages with the accuracy required to address some outstanding problems, such as those related to the Milky Way's formation process. The pivotal importance of these problems and the need to improve age estimates as far as possible are the basis of the huge effort devoted in recent decades to the gathering of the relative ages of GGCs. As a consequence, there exists quite a rich literature dedicated to this fundamental topic \\citep[][and references therein]{Stet96, Sar97, R99, SW02, DeA05, MF09}. The relative age$-$dating techniques for GGCs - almost universally adopted - can be separated into two basic classes: those methods that are based on brightness difference measurements - the {\\sl vertical} method -, and those that are based on color difference measurements - the {\\sl horizontal} method - in the color--magnitude diagram (CMD). The most commonly adopted of the {\\sl vertical} methods is the magnitude difference between the main sequence turn-off (MSTO) and the zero age horizontal branch (ZAHB), usually estimated starting from the level of the RR Lyrae instability strip. The {\\sl horizontal} method is based on the measurement of the color difference between the MSTO and a point in the lower part of the red giant branch (RGB). It is clear that both approaches are independent of distance, reddening, and uncertainty in the adopted photometric zero points. However, apart from this common characteristic, the two methods present different advantages and drawbacks. In particular, from the point of view of their theoretical calibration, although the dependence of the ZAHB luminosity on the metallicity is still controversial, the vertical method seems to be more reliable than the horizontal one, which is strongly affected by the non-negligible uncertainties related to the treatment of superadiabatic convection and color-$T_{eff}$ relations. As a result, relative ages determined using the horizontal method are strongly model-dependent. On the other hand, from the observational point of view, the horizontal method seems to be largely unaffected - or in any case to quite a minor extent - by the difficulty of measuring the MSTO brightness due to its verticality in the CMD. An accurate determination of the ZAHB luminosity can also be a thorny problem, in particular for clusters at the extreme boundaries of the metallicity distribution, which generally have a red (or blue) horizontal branch and no RR Lyrae stars. In both methods, in order to minimize both observational and theoretical uncertainties in the relative age determination, the magnitude or the color difference estimated for a GGC is compared with that of a GGC of similar metallicity. A detailed discussion of the advantages and disadvantages of all the differential age-dating methods can be found in \\citet{Stet96} and we refer the interested reader to the quoted reference. \\citet{MF09} have quite recently performed a detailed analysis of the relative ages of a sizeable sample of GGCs. In order to increase the level of accuracy in their age estimates they have developed an independent method for estimating the MSTO brightness difference between two GGCs of similar metallicity. This method is based on a simultaneous fit of the fainter portion of the MS - whose location is quite insensitive to cluster age but largely dependent on cluster metallicity - and the lower portion of the RGB - whose location is strongly dependent on the cluster chemical composition (hereinafter, the {\\sl relative MS-fitting}, rMSF method). A careful analysis of the advantages and uncertainties of this method has been performed by \\citet{MF09} and will not be repeated here. It is important to note that all relative age-dating methods are based on the implicit assumptions that all GGCs, apart from the well known differences in their iron content and, eventually, in age, are relatively similar objects, i.e., with negligible - if any - differences in the helium content and/or heavy elements distribution. However, this common idea has been severely challenged in recent decades by the growing evidence, based on accurate spectroscopical measurements, that quite peculiar chemical patterns among various GGCs with the same iron content do exist, and quite often also among stars belonging to the same cluster \\citep[for a complete review on this issue we refer the reader to][]{Gratton04}. Even more surprising has been the recent discovery that many GGCs host multiple stellar populations \\citep{Bedin04, Piotto05, Piotto07, Milone08, Cass08}, whose photometric properties can be understood in terms of sometimes quite large differences in the initial He content and/or in the heavy elements distributions and/or age. On the basis of the quoted peculiar photometric and spectroscopic properties, a possibility now considered more plausible is that the multi-population evidence is not a peculiarity of a restricted number of objects, but on the contrary, that it could be a common feature in the GGC system. The main difference is that some clusters could have the capability of retaining the different stellar populations, whereas others could at present be formed completely \\citep{Derc08, DanCal08}, or at a significant level (between 50 and 70\\%), by second generation stars \\citep{Carret09}. It is evident that the presence of observationally deceptive differences in the chemistry of GGCs might have a huge impact on relative age measurements. In fact, observational differences between two GGCs of similar metallicity, as detected with any of the quoted techniques, could be erroneously interpreted as due to an age difference when - on the contrary - it might only be due to an unrecognized difference in the chemical abundance pattern. This issue has been recently investigated by \\citet{dantona09}, who have only considered the specific case of the two GGCs (NGC~1851 and NGC~6121) and showed that the brightness difference existing between the sub-giant branches of these two GGCs can be explained as due to a difference in the CNO element abundance, even under the hypothesis that the two clusters are perfectly coeval. In this article, we wish to address this issue in greater detail and more extensively, but only on purely theoretical grounds. It is worth noting that we do not investigate the robustness of the theoretical calibration of each independent age-dating method. We use a homogeneous set of stellar models in a completely differential approach for estimating how unrecognized differences in the chemical patterns of GGCs affect the relative age results obtained with the commonly adopted methods. The plan of this paper is as follows: in the next section, the adopted theoretical framework is presented; in Section~\\ref{methodsSec} we briefly review the relative age$-$dating techniques that we'll take into account in present work; in Section~\\ref{results} we present our analysis by discussing individually each age-dating method; and we close with our conclusions and final remarks. ", "conclusions": "In the present work, we have studied the impact that unrecognized differences in the chemical patterns of GGCs have on their relative age determinations. The two most widely used relative age-dating methods, horizontal and vertical, together with the more recent relative MS-fitting method described in \\citet{MF09}, were carefully analyzed on a purely theoretical basis. The BaSTI library of stellar models was adopted to perform the present analysis, supplemented by additional evolutionary computations for more extreme assumptions about the initial He content. Our main conclusions are summarized here: \\begin{itemize} \\item We find that relative ages derived using the horizontal-method are largely dependent on the initial Y value and CNONa mixture. Undetected differences in the He content and/or CNONa abundance translates in an unreal age determination. The difference between the measured relative age and the actual one is in the range from 0.08 to 0.3 (or 1 to 4 Gyr, if these values are transformed to absolute ages), this result worsens for high metallicity clusters. \\item For the vertical method, we find that the obtained relative ages are from 0.2 to 0.4 (or from 3 to 5 Gyr in absolute values) older than the isochrone (input) ones if a helium enhancement of Y=0.3 is considered, independently of the metallicity. This result worsens if larger - more extreme - values of Y are taken into account, reaching age differences of the order of more than 0.8 ($\\sim$ 10 Gyr) if extreme Y=0.35 or 0.4 values are considered. The vertical method is most sensitive to cluster-to-cluster undetected variations in Y. \\item We find that the vertical-method can be used to limit the possible cluster-to-cluster differences in He according to their measured relative ages. In particular, the Y dispersion in GGCs has been limited to $\\sim$0.01. \\item We find that the rMSF method is the relative age-dating technique for which undetected differences in the He content has less impact. It has been shown that neither the young group of GGCs' age--metallicity relation nor the 0.03 intrinsic relative age dispersion found by \\citet{MF09} is produced by a He dispersion in GGCs. \\item When considering the possibility of undetected differences in the CNONa mixture, our results allow us to constrain the maximum possible enhancement of CNONa elements, which should be of the order of 1.2; i.e., $\\sim0.18$~dex, with respect a \\lq{standard}\\rq\\ $\\alpha-$enhanced mixture. \\item The relative age dispersion found by \\citet{MF09} could be entirely produced by a cluster-to-cluster CNONa dispersion. \\item We reassert that the age--metallicity relation found in \\cite{MF09} is a real age sequence. \\item When taking also into account the advantages of the rMSF technique with respect to the other relative age dating methods, it appears that, so far, that the rMSF method is the approach that is much to be preferred for retrieving GGC chronology. \\end{itemize}" }, "1003/1003.1113_arXiv.txt": { "abstract": "X-ray and $\\gamma$-ray observations can help understand the origin of the electron and positron signals reported by ATIC, PAMELA, PPB-BETS, and Fermi. It remains unclear whether the observed high-energy electrons and positrons are produced by relic particles, or by some astrophysical sources. To distinguish between the two possibilities, one can compare the electron population in the local neighborhood with that in the dwarf spheroidal galaxies, which are not expected to host as many pulsars and other astrophysical sources. This can be accomplished using X-ray and $\\gamma$-ray observations of dwarf spheroidal galaxies. Assuming the signal detected by Fermi and ATIC comes from dark matter and using the inferred dark matter profile of the Draco dwarf spheroidal galaxy as an example, we calculate the photon spectrum produced by electrons via inverse Compton scattering. Since little is known about the magnetic fields in dwarf spheroidal galaxies, we consider the propagation of charged particles with and without diffusion. Extending the analysis of Fermi collaboration for Draco, we find that for a halo mass $\\sim10^{9}\\,\\mathrm{M}_{\\odot}$, even in the absence of diffusion, the $\\gamma$-ray signal would be above the upper limits. This conclusion is subject to uncertainties associated with the halo mass. If dwarf spheroidal galaxies host local magnetic fields, the diffusion of the electrons can result in a signal detectable by future X-ray telescopes. ", "introduction": "The nature of cosmological dark matter remains a tantalizing puzzle~\\cite{Bertone:2004pz}. If dark matter is made up of weakly interacting massive particles (WIMPs), their annihilation products may be observed and used for identification of the dark-matter particles. PAMELA~\\cite{PAMELA}, ATIC~\\cite{ATIC}, PPB-BETS~\\cite{Torii:2008xu}, and Fermi~\\cite{Abdo:2009zk} have observed unexpected features in the electron and positron spectra at high energies. The high-energy electrons and positrons could come from the annihilations or decays of dark matter particles~\\cite{Ibe:2008ye,Chen:2008md,Chen:2008qs,Fox:2008kb,ArkaniHamed:2008qn,Shirai:2009wi,Ibe:2009dx,Ibe:2009en,Shirai:2009wi,Shirai:2009kh}, but they could also be produced by astrophysical sources, such as pulsars, supernova remnants and Gamma-Ray Bursts (GRBs) \\cite{Yuksel:2008rf,Ioka:2008cv,Profumo:2008ms,Stawarz:2009ig,Calvez:2010fd,PhysRevD.80.123017,PhysRevLett.103.111302,PhysRevLett.103.051104}. Theoretical models of dark matter can accommodate a wide range of parameters (see for example Refs.~\\cite{ModelInd,FermiInt,Profumo:2008ms}). Astrophysical models of particle acceleration by pulsars are also uncertain, but they can also account for the observed signal. To distinguish between the two possibilities, it would be desirable to compare the electron and positron populations in the local neighborhood with that in some other parts of the galaxy, which are known to be devoid of pulsars and other potential astrophysical sources of high-energy particles. Dwarf Spheroidal Galaxies (dSphs) present such an opportunity, provided that one can infer the photon spectra generated by the interaction of electrons and positrons with the Cosmic Microwave Background radiation (CMB). Observations of dwarf spheroidal galaxies have recently been used in a dedicated search for {\\em decaying} dark matter in the form of sterile neutrinos~\\cite{Loewenstein:2008yi,Loewenstein:2009cm}. Sterile neutrinos are expected to undergo a two-body decay, producing a narrow line in the X-ray spectrum (for a recent review, see, e.g., Ref.~\\cite{Kusenko:2009up}). Detection of WIMP, which are much heavier and which annihilate rather than decay, presents a very different challenge~\\cite{Feng:2010gw}. For a number of WIMP models, the X-ray and $\\gamma$-ray signals would be too faint to observe in the foreseeable future, but the same models would predict the flux of high-energy electrons and positrons well below the levels observed by ATIC, PAMELA, PPB-BETS, and Fermi. A Breit-Wigner resonance~\\cite{Ibe:2008ye} or long-range interactions~\\cite{ArkaniHamed:2008qn,Lattanzi:2008qa} could increase the dark matter annihilation cross section, but it is difficult to reconcile Sommerfeld enhancement with the primordial relic abundance of WIMP or with the $\\mu$-type distortion of the CMB energy spectrum~\\cite{Feng:2009hw,Zavala:2009mi}. One does expect a boost factor from the small scale structure of dark matter, but the required values are well in excess of one's expectations based on numerical N-body simulations. In the absence of a compelling theoretical framework, we will not try to relate our predictions to any specific model of dark matter, but we will focus on a model-independent determination of whether the high-energy electrons originate from dark matter (which is abundant in both the local neighborhood and in a dwarf spheroidal galaxy), or from some astrophysical source candidates (whose population in a dwarf spheroidal galaxy is suppressed). ", "conclusions": "Current and future observations of dwarf spheroidal galaxies with $\\gamma$-ray detectors and with the next generation X-ray telescopes may be able to probe the origin of high-energy electrons and positrons observed by Fermi and PAMELA. Current results from Fermi, which we have extended to the possible case of non-diffusive propagation in a dwarf spheroidal galaxy, seem to be in conflict with the most conservative $\\gamma$-ray prediction from an almost model independent analysis of an assumed dark matter signal. The litmus test for dark matter annihilations as the possible origin of these signals comes from comparing the high-energy particle fluxes in a dwarf spheroidal galaxy and in the local neighborhood of the Milky Way disk. Dark matter dominated systems, such as dwarf spheroidal galaxies, should generate a predictable flux of the dark-matter annihilations products, subject to mass model uncertainties. At the same time, most of the astrophysical sources capable of producing high-energy particles in the sun's neighborhood are absent in dwarf spheroidal galaxies, which have very few stars and very little gas as compared with their dark matter content. The high-energy electrons and positrons produced in dwarf spheroidal galaxies can generate X-rays and $\\gamma$-rays by up-scattering CMB photons to higher energy bands. The predicted fluxes are above the upper limits set by Fermi for Draco; although this result could be interpreted as favoring an astrophysical source for the excess of high energy electrons and positrons in our galaxy, the large uncertainty in the mass of the halo limit the strength of this interpretation. Better mass measurements need to be obtained before conclusions can be made." }, "1003/1003.3116_arXiv.txt": { "abstract": "Distant clumpy galaxies are thought to be Jeans-unstable disks, and an important channel for the formation of local galaxies, as suggested by recent spatially-resolved kinematic observations of z$\\sim$2 galaxies. I study the kinematics of clumpy galaxies at z$\\sim$0.6, and compare their properties with those of counterparts at higher and lower redshifts. I selected a sample of 11 clumpy galaxies at z$\\sim$0.6 from the representative sample of emission line, intermediate-mass galaxies IMAGES. Selection was based on rest-frame UV morphology from HST/ACS images, mimicking the selection criteria commonly used at higher redshifts. Their spatially-resolved kinematics were derived in the frame of the IMAGES survey, using the VLT/FLAMES-GIRAFFE multi-integral field spectrograph. For those showing large-scale rotation, I derived the Toomre $Q$ parameter, which characterizes the stability of their gaseous and stellar phases. I find that the fraction of UV-selected clumpy galaxies at z$\\sim$0.6 is 20$\\pm$12\\%. Roughly half of them (45$\\pm$30\\%) have complex kinematics inconsistent with Jeans-unstable disks, while those in the remaining half (55$\\pm$30\\%) show large-scale rotations. The latter reveal a stable gaseous phase, but the contribution of their stellar phase makes them globally unstable to clump formation. Clumpy galaxies appear to be less unstable at z$\\sim$0.6 than at z$\\sim$2, which could explain why the UV clumps tend to vanish in rest-frame optical images of z$\\sim$0.6 clumpy galaxies, conversely to z$\\sim$2 clumpy galaxies, in which the stellar phase can substantially fragment. This suggests that the former correspond to patchy star-formation regions superimposed on a smoother mass distribution. A possible and widespread scenario for driving clump formation relies on instabilities by cold streams penetrating the dark matter halos where clumpy galaxies inhabit. While such a gas accretion process is predicted to be significant in massive, z$\\sim$2 haloes, it is also predicted to be strongly suppressed in similar, z$\\sim$0.6 haloes, which could explain why lowest-z clumpy galaxies appear to be driven by a different mechanism. Instead, I found that interactions are probably the dominant driver leading to the formation of clumpy galaxies at z$<$1. I argue that the nature of z$>$1 clumpy galaxies remains more uncertain. While cold flows could be an important driver at z$\\sim$2, I also argue that the \\emph{observed and cumulative} merger fraction between z=2 and z=3 is large enough so that every z$\\sim$2 galaxy might be the result of a merger that occurred within their past 1 Gyr. I conclude that it is premature to rule out mergers as a universal driver for galaxy evolution from z$\\sim$2 down to z=0. ", "introduction": "How galaxies formed, evolved, and built-up the local Hubble sequence is still an open and highly debated issue. For instance, there is now evidence for a significant evolution of the stellar mass density at z$<$1. Even between redshifts as low as z$\\sim$0.6 and z=0, the stellar mass density appears to be evolving by 0.2 dex (i.e., $\\sim$40\\%, see, e.g., \\citealt{PerezGonzalez08}). Such an evolution in stellar mass can also be seen in the evolution of the Tully-Fisher relation, which appears to be evolving by a factor two in stellar mass over the same redshift range \\citep{puech08,puech09b}. Such an increase in stellar mass needs to be fed by fresh gas. The absence of evolution in zero point of the \\emph{baryonic} Tully-Fisher relation between z$\\sim$0.6 and z=0 \\citep{puech09b} suggests that most of this gas was already gravitationally bound to galaxies at z$\\sim$0.6. What is the mechanism driving the conversion of this gas into stars? In spite of the impressive progress accomplished over the past years, mainly through spatially-resolved kinematics of distant galaxies, this issue still remains without a clear answer. At z$\\sim$0.6 (i.e., 6 Gyr ago), \\cite{yang08} gathered a representative sample of 63 galaxies with $M_{stellar} \\geq 1.5\\times 10^{10} M_\\odot$ and spatially-resolved kinematics derived from FLAMES/GIRAFFE observations at the VLT. They found that 40\\% of z$\\sim$0.6 intermediate-mass galaxies (i.e., the progenitors of local spirals), have chaotic velocity fields inconsistent with expectations from pure rotating disks. Subsequent analyzes suggest that most of them are likely associated with major mergers \\citep{hammer09b}. Mergers are also found to be a good driver for the large scatter seen in the distant Tully-Fisher relation \\citep{puech09b,covington10}. Actually, the remarkable co-evolution of the morphology, kinematics, star formation density, metal density, and stellar mass density, all could find a common and natural explanation in the frame of the spiral rebuilding scenario, according to which 50 to 75\\% of present-day spiral disks were rebuilt after a major merger since z=1, as proposed by \\cite{hammer05}. To this respect, the Milky Way appears to be exceptional, with a remarkable quite past history compared to other local spiral galaxies \\citep{hammer07}. There is now a growing body of evidence suggesting that disk rebuilding indeed took place at z$\\leq$1, which makes it as a viable driver for galaxy evolution at these epochs. On the theoretical side, numerical simulations showed how gas can be expelled in tidal tails during such mergers, and how it can be subsequently re-accreted \\citep{barnes02}. This re-accreted gas is expected to cool down and form new stars, re-building a new disk around a spheroidal remnant \\citep{springel05, robertson06}, which might correspond to morphologies as late as Sb galaxies \\citep{lotz08,hopkins09c}. Recent theoretical developments have shed light on the underlying process, which appears to be purely gravitational \\citep{hopkins09}. The requirement for a major merger to rebuild a new disk depends mainly on the gas fraction during the final coalescence, which needs to be at least 50\\% \\citep{robertson06}. Cosmological simulations are now also producing such re-processed disks at z$<$1, although the role of cosmological gas accretion is not totally understood, but taking this process into account could result in a lower gas fraction threshold for rebuilding a new disk \\citep{governato08}. On the observational side, first examples of rebuilt disks were recently detected at z$\\sim$0.6 \\citep{puech09,hammer09}, and the auto-consistency of the disk rebuilding process starts being investigated, both theoretically \\citep{hopkins09b,stewart09,hopkins09c} and observationally \\citep{hammer09b,kannappan09,bundy09,huertas10}. If the spiral rebuilding scenario appears to achieve encouraging successes in describing galaxy evolution at z$<$1, some points still need to be investigated. In particular, the impact of the expected numerous \\emph{minor} mergers on the survival of thin disks is still debated \\citep{toth92,hopkins08,purcell09,moster09b}. Furthermore, Luminous InfraRed Galaxies (LIRGs) account for about 80\\% of the star formation density reported at z$\\leq$1 \\citep{hammer05}, but their morphologies reveal that only half of z$>$0.5 LIRGs are compatible with mergers, while those in the other half appear to be spiral \\citep{melbourne05}. \\cite{marcillac06} showed that the large number density of LIRGs at these epochs suggests that they could experience between two and four star formation bursts until z=0, with typical timescales of $\\sim$0.1 Gyr, which is not consistent with a simply continuous star formation history. They concluded that minor mergers, tidal interactions, or gas accretion remain plausible triggering mechanisms in distant LIRGs harboring a spiral morphology. Interestingly, 75\\% of local LIRGs are barred, which could play a role in regulating star formation in such objects \\citep{wang06}. At higher redshifts (i.e., z$>$1), the co-moving density of galaxy appears to be dominated by clumpy irregular galaxies \\citep{elmegreen07}. Interest for such objects dates back to \\cite{cowie95}, who first noticed the unusual aspects of some high-z galaxies, dubbed as ``chain galaxies'' and described as ``linearly organized giant star-forming regions''. The large occurrence of blue star-forming knots in less edge-on objects was also later recognized as a general and intriguing feature of distant galaxies, probably linked to an early phase in the formation of local spiral galaxies \\citep{cowie95,vandenbergh96}, and were latter referred to as ``clump clusters'' by \\cite{elmegreen04}. The improved spatial resolution provided by the HST/ACS re-invigorated the interest for these objects, which were all suggested to be different incarnations of the same underlying population viewed along different inclination angles \\citep{oneil00,elmegreen04b}. Both kind of objects are therefore often referred to as ``clumpy galaxies''. They are found to be typically made of 5-10 kpc-sized clumps with stellar masses $\\sim 10^{7-9}$ M$_\\odot$ (i.e., $\\sim$100 times more massive than the largest star complexes in present-day spiral galaxies), and they typically account for one third of the total galaxy emission \\citep{elmegreen05}. Such clumps are thought to be linked to the formation of disks, as suggested by the increase of the inter-clump surface density, and the decrease of the mass surface density contrast between the clumps and the inter-clump regions, when going from clumpy galaxies with no evident inter-clump emission to clumpy galaxies with faint red disks, and spiral galaxies \\citep{elmegreen09b}. What is driving the formation of these clumps has been the subject of many attentions during the past decade. Numerical simulations suggested that clumps might originate from the local gravitational instability of very gas-rich disks of young galaxies \\citep{noguchi98,immeli04}. Due to their large masses, the clumps would experience strong dynamical friction and spiral towards the galaxy center within a few Gyr, which might lead to the formation of a bulge \\citep{noguchi99,elmegreen08,elmegreen09}, as well as a thick stellar disk through strong stellar scattering \\citep{bournaud09b}. Simulated clumps are found to show properties similar to observations \\citep{immeli04b,bournaud07}. The lifetime of these clumps is so short and the fraction of clumpy galaxies at z$>$1 so high, that making the clumpy phase a long-term phenomenon requires a continuous and rapid fresh supply of cold gas in order to feed the disk and regenerate new clumps \\citep{dekel09b}. Theoretical developments indeed suggested that early galaxy formation is fed by cold streams penetrating through dark matter halos \\citep{dekel09}. These cold streams are expected to maintain a dense disk that can undergo gravitational fragmentation into several giant clumps \\citep{dekel09b}. This possible link between high-z clumpy galaxies and the cosmological context was strengthened both by recent cosmological numerical simulations \\citep{agertz09,ceverino09}, and semi-analytic models \\citep{khochfar08}. Alternatively, it was proposed that clumps could also result from on-going mergers or interactions \\citep{taniguchi01,overzier08,dimatteo08}. Discriminating between the merger and fragmentation scenarii is not straightforward because it requires high-resolution integral field spectroscopy in high-z galaxies. Indeed, as stated by \\cite{noguchi99}, the most straightforward and powerful test for discriminating between the two hypothesis is to examine the kinematics of the clumps: in the merger scenario, a random orientation of clump spins is expected, while in the fragmentation scenario, clumps are expected to be coplanar. In particular, one interesting predictions of these simulations is that the large-scale rotation in the underlying disk should be preserved during the fragmentation phase \\citep{immeli04,bournaud07}. The achievement of integral field spectrograph working in the near infrared (e.g., SINFONI at the VLT, or OSIRIS at Keck), allowed several teams to gather spatially-resolved kinematic observations of z$>$1 distant galaxies (see \\citealt{forsterschreiber09,law09} and references therein, as well as \\citealt{wright07,wright08,bournaud08,vanstarkenburg08,epinat09}). Detection of an underlying rotation was claimed in several z$\\sim$2 clumpy galaxies, which has been used as a support to the fragmentation scenario \\citep{genzel08}. Surprisingly, analysis of integral field spectroscopy observations suggest that two different scenarii might take place at z$<$1, and at z$>$1. It therefore becomes necessary to start investigating whether both scenarii are consistent, and whether or not a transition between two different galaxy evolution drivers is occurring between these two epochs. To this aim, I take advantage of the IMAGES survey to study clumpy galaxies at z=0.6. The goal is to investigate which of the merger or fragmentation scenario is the most consistent with z$\\sim$0.6 clumpy galaxies, and investigate whether clumpy galaxies at z$<$1 and at z$>$1 are driven by a common physical process. This paper is organized as follows: In Sect. 2, I describe how the clumpy galaxies at z$\\sim$0.6 were selected; In Sect. 3, I present their kinematic and dynamical properties; Sect. 4 discusses the origin of the clumpiness in distant galaxies; Sect. 5 discusses the results, while conclusions are drawn in Sect. 6. Throughout the paper, I adopt $H_0=70$ km/s/Mpc, $\\Omega _M=0.3$, and $\\Omega _\\Lambda=0.7$, and the $AB$ magnitude system. ", "conclusions": "I selected and studied 11 clumpy galaxies at z$\\sim$0.6, drawn from the representative IMAGES sample of emission line, intermediate mass galaxies. Clumpy galaxies were selected mimicking the morphological UV selection criteria used at larger redshifts, and their fraction at z$\\sim$0.6 was found to be $\\sim$20\\%. Among the 11 clumpy galaxies, 5 were found to show complex kinematics compatible with major mergers, as suggested by \\cite{hammer09b}. The remaining clumpy galaxies, i.e., which show large-scale rotation, are found to be Toomre-stable in their gaseous phase, but unstable in their stellar phase, contrary to higher-z clumpy galaxies, which are found to be unstable in both phases, with a stronger level of effective instability. This could originate both in the likely higher fraction of old stars in z$\\sim$0.6 clumpy galaxies, which would not participate into the fragmentation process. This would naturally explain why z$\\sim$0.6 UV clumps tend to vanish when looking at reddest-band images, unlike z$\\sim$2 clumpy galaxies, which appear to be more persistent at longer wavelengths. While z$>$1 clumpy galaxies were largely associated with Jeans-unstable disks in the literature, I argue that current kinematic observations could actually support a fraction of only $\\sim$33\\% of such systems among z$\\sim$2 star-forming galaxies. While this is not incompatible with the cumulative fraction of major mergers in such objects over their past $\\sim$1 Gyr, theoretical and numerical works suggested a different channel for disk fragmentation, which could result from cold flows from the inter-galactic medium penetrating through the surrounding dark matter halos. These cold flows could maintain the gaseous phase high enough to fragment and regenerate clumps at a rate high enough to be consistent with observations. The progressive shut down of these cold flows with redshift in massive haloes could also explain why the formation of z$<$1 clumpy galaxies are found to be preferentially driven by interactions, since the duty cycle of merger would dramatically evolve between these two epochs. Within this theoretical frame, there is therefore no contradiction between the interpretation of kinematic observations of z$\\sim$2 and of z$\\sim$0.6 galaxies. However, it remains unclear whether interactions could not account for all or at least most of the evolution seen in intermediate and massive galaxies all the way from z$\\sim$2-3 down to z=0. Confirming the nature and the fraction of high-z clumpy galaxies will require better spatial resolutions on larger telescopes. Unfortunately, it is likely that such a debate will find a answer only with the advent of the Extremely Large Telescopes and their associated NIR integral field spectrographs such as EAGLE \\citep{puech08b,puech09c,puech09d}." }, "1003/1003.4500_arXiv.txt": { "abstract": "We continue our deep optical imaging survey of the Virgo cluster using the CWRU Burrell Schmidt telescope by presenting \\B-band surface photometry of the core of the Virgo cluster in order to study the cluster's intracluster light (ICL). We find ICL features down to \\mub$\\approx29$ \\magsec, confirming the results of Mihos \\etal (2005), who saw a vast web of low-surface brightness streams, arcs, plumes, and diffuse light in the Virgo cluster core using \\V-band imaging. By combining these two data sets, we are able to measure the optical colors of many of the cluster's low-surface brightness features. While much of our imaging area is contaminated by galactic cirrus, the cluster core near the cD galaxy, M87, is unobscured. We trace the color profile of M87 out to over 2000\\arcsec, and find a blueing trend with radius, continuing out to the largest radii. Moreover, we have measured the colors of several ICL features which extend beyond M87's outermost reaches and find that they have similar colors to the M87's halo itself, \\BV$\\approx0.8$. The common colors of these features suggest that the extended outer envelopes of cD galaxies, such as M87, may be formed from similar streams, created by tidal interactions within the cluster, that have since dissolved into a smooth background in the cluster potential. ", "introduction": "Massive galaxy clusters are known to contain a population of stars which reside outside of any of the cluster's galaxies, often referred to as intracluster light or ICL. ICL features typically have extremely faint surface brightnesses of $<1\\%$ of the brightness of the night sky, making their study extremely difficult. While the first indications of the existence of ICL came from observations by Zwicky (1951), only with the advent of modern CCD technologies have detailed studies of these stars been made possible (\\eg Uson \\etal 1991; V{\\'i}lchez-G{\\'o}mez et al. 1994; Bernstein \\etal 1995; Gregg \\& West 1998; Trentham \\& Mobasher 1998). The most straightforward method of detecting the ICL is through deep broadband imaging at optical wavelengths (\\eg Uson \\etal 1991; V{\\'i}lchez-G{\\'o}mez et al. 1994; Trentham \\& Mobasher 1998; Feldmeier \\etal 2002, 2004a; Mihos \\etal 2005, hereafter referred to as M05; Gonzalez \\etal 2005; Krick \\& Bernstein 2007). With such imaging, we can not only measure the luminosity of ICL the component, but we can study its spatial distribution and detect individual ICL features, such as streams, arcs, and plumes (\\eg Gregg \\& West 1998; Trentham \\& Mobasher 1998; Calc{\\'a}neo-Rold{\\'a}n \\etal 2000; White \\etal 2003; M05; Krick \\etal 2006; Yagi \\etal 2007) as well as any large-scale diffuse components (\\eg Gonzalez \\etal 2000; Feldmeier \\etal 2002, 2004a; Adami \\etal 2005; Zibetti \\etal 2005; Patel \\etal 2006; Krick \\& Bernstein 2007; Pierini \\etal 2008; Da Rocha \\etal 2008). ICL has also been detected using discrete stellar tracers (\\eg Ferguson \\etal 1998; Feldmeier \\etal 1998; Durrell \\etal 2002; Gal-Yam \\etal 2003; Arnaboldi \\etal 2004; Feldmeier \\etal 2004b; Gerhard \\etal 2005; Aguerri \\etal 2005; Neill \\etal 2005; Maoz \\etal 2005; Williams \\etal 2007; Castro-Rodrigu{\\'e}z \\etal 2009; McGee \\& Balogh 2010), although these methods often lack the ability to detect individual ICL features. Intracluster light is thought to form primarily by the tidal stripping of stars as galaxies interact and merge during the hierarchical accretion history of the cluster, causing the fraction of the cluster's luminosity found in the ICL to increase as it evolves (\\eg Napolitano \\etal 2003; Murante \\etal 2004; Willman \\etal 2004; Rudick \\etal 2006; Monaco \\etal 2006; Conroy \\etal 2007; Murante \\etal 2007; Purcell \\etal 2007; Yan \\etal 2009; Baria \\etal 2009). Numerous mechanisms for generating the ICL have been proposed, including the infall of groups into the cluster potential (Willman \\etal 2004; Rudick \\etal 2006), high speed encounters within the cluster (Moore \\etal 1996; Gnedin 2003), and galactic mergers during the buildup of the massive central galaxy (Murante \\etal 2007; Conroy \\etal 2007). All of these processes are likely to be occurring simultaneously and each will create distinct observable signatures (Rudick \\etal 2009). Thus, the formation of ICL is intimately linked to the dynamical history of the cluster, and the observable features of the ICL should contain a great deal of information about the evolutionary processes which have shaped both the cluster and its constituent galaxies. While ICL is generally thought of as stellar material found \\emph{outside} of any individual galaxy, in practice galaxies have no well-defined edge (Abadi \\etal 2006). Thus, the distinction between the ICL and the outer luminosity profiles of cluster galaxies, which display similar surface brightnesses in broadband imaging, is difficult and somewhat arbitrary (M05). In fact, simulations of the formation of clusters' most massive elliptical galaxies have shown that their extended stellar profiles form through similar merger and tidal stripping mechanisms as the more diffuse ICL (Dubinski \\etal 1998; Monaco \\etal 2006; Conroy \\etal 2007; Murante \\etal 2007; Ruszkowski \\& Springel 2009). We therefore prefer to use \\emph{intracluster light} as a qualitative description of cluster luminosity at low surface brightness, which may refer to individual tidal streams, extreme galactic outskirts, or any large-scale diffuse luminosity component, and which are all likely products of tidal stripping and disruption of galaxies during the dynamical evolution of the cluster. In addition to the quantity and morphology of the ICL, knowledge of the underlying stellar populations also provides a vital tool for understanding its formation. Giant elliptical galaxies are known to display radial color gradients, whereby the color index decreases, or becomes bluer, with increasing radius (\\eg Pettit 1954; de Vaucouleurs 1961; Carter \\& Dixon 1978; Strom \\& Strom 1978; Davis \\etal 1985; Vader \\etal 1988; Goudfrooij \\etal 1994; Bernardi \\etal 2003; Cantiello \\etal 2005; Liu \\etal 2005), primarily due to gradients in the stellar metallicities (\\eg Spinrad \\etal 1972; Strom \\etal 1976; Baum \\etal 1986; Carollo \\etal 1993; Tamura \\etal 2000; S{\\'a}nchez-Bl{\\'a}zquez \\etal 2007; Rawle \\etal 2010). Just as the strength of these gradients is an important clue in reconstructing the formation history of elliptical galaxies, the ages and metallicities of the intracluster stars will be highly dependent on the progenitor galaxies from which they were stripped. The results of Sommer-Larsen \\etal (2005) show that the metallicity of the ICL is expected to be on average similar to that of the outer envelope of the cD galaxy, while Purcell \\etal (2008) show that individual ICL streams formed by recent interactions should be more metal rich than the surrounding diffuse component. Murante \\etal (2004) have shown that the diffuse ICL stellar population is expected to be older, on average, than the galactic stars. Observationally, it is very difficult to break the age-metallicity degeneracy and make precise determinations of either quantity. Although Williams \\etal (2007) used HST ACS imaging of red giant branch stars in the Virgo cluster to do so, such observations are only possible in the most nearby systems and span very small fields of view. Their work found that the ICL is composed of stars with a wide variety of ages and metallicities, but that the dominant component is old ($\\simgt 10$ Gyr) and moderately metal poor ([M/H]$\\simlt -1.0$). A more common technique is to instead use broadband optical colors in order to compare the ICL stellar populations to those of galactic stars. Most studies have found results consistent with Williams \\etal (2007), where the average ICL color is similar to that of the outer halo of the cluster's brightest galaxy (\\eg Zibetti \\etal 2005; Pierini \\etal 2008; da Rocha \\etal 2008), which due to radial color gradients is bluer than the galaxy's interior regions. However, a number of clusters have been observed to have an ICL component with significantly redder colors, more similar to the brightest galaxy's interior (Krick \\& Bernstein 2007; da Rocha \\& Mendes de Oliveira 2005; Gonzalez \\etal 2000). As part of our ongoing deep imaging survey of the Virgo cluster using Case Western Reserve University's Burrell Schmidt telescope, this paper presents deep imaging of the Virgo cluster core in the \\B-band. By combining this data with the \\V-band results previously published in M05, we are able to measure the optical colors of Virgo's intracluster light. A detailed description of our data acquisition and reduction techniques is given in Section 2. Our \\B-band image is presented in Section 3, while Section 4 combines the two images in order to measure the colors of ICL features. Section 5 includes a summary of our results and a discussion of our interpretations. Finally, detailed error models for our photometric measurements can be found in the Appendix. ", "conclusions": "\\label{sec:discussion} In this paper, we have presented our \\B-band image of the Virgo cluster core, taken as part of our ongoing deep imaging survey of the Virgo cluster. All aspects of the survey, from the optical and mechanical design of the telescope system, to the acquisition, analysis, and reduction of the data have been optimized to detect diffuse light at low surface brightness. Our final photometric uncertainties are dominated by systematic errors resulting from sky subtraction and astrophysical backgrounds and not simple photon statistics. Our \\B-band imaging confirms the results of M05, which found a vast web of diffuse light features in the core of the Virgo cluster. These features include large-scale diffuse components, as well as a number of discrete features such as streamers, arcs, tails, and plumes which provide a wealth of information on the dynamical history of the cluster and its constituent galaxies. Just outside the cluster core, however, we have detected a large number of galactic cirrus features which prevent us from studying diffuse extragalactic sources which lie behind them. By combining these \\B-band results with the \\V-band imaging from M05, we have been able to measure the colors of a number of low surface brightness features in Virgo's core. We have measured the radial color profile of M87 out to very large radius, demonstrating that the color gradients seen in the inner regions of the galaxy extend to its outer stellar envelope. Furthermore, we have measured the colors of several tidal features which extend from M87's stellar envelope and find that the two populations have similar optical colors within the measurement uncertainties. These results are consistent with the majority of other measurements of ICL color, which suggest that the ICL colors should be similar to those of the outskirts of the cluster's brightest galaxies (\\eg Zibetti \\etal 2005; Sommer-Larsen \\etal 2005; Krick \\& Bernstein 2007). These results are consistent with the hypothesis that the outer envelopes of cD galaxies like M87 may be predominantly built-up by the tidal stripping and disruption of smaller satellite galaxies, the same mechanism which likely generates the tidal streams and other ICL features (Rudick \\etal 2009). In this scenario, the cD envelope is simply composed of tidal streams of ICL which have been mixed in the cluster potential, and dissolved to form a smooth distribution. While optical colors alone cannot definitively prove that any two stellar populations are identical due to the well-known age-metallicity degeneracy, the similar colors of these two populations is suggestive of a common origin. The optical colors of both the tidal features and M87's outer envelope are also consistent with the colors of the Virgo cluster's dwarf elliptical galaxy population (van Zee \\etal 2004). Thus, the disruption of these galaxies in the cluster potential may provide a ready source for the intracluster stars which make up the low-surface brightness outer envelope of M87 and the tidal streams. In fact, the thinness of two of the tidal streams studied in Section \\ref{sec:m87_streams} suggests that they originate from low-mass galaxies with small velocity dispersions (M05), such as dwarf ellipticals. Furthermore, the results of Williams \\etal (2007), who measured the age and metallicity of Virgo's intracluster stars at much larger radius from any of the cluster's large galaxies to be $\\simgt 10$ Gyr and [M/H]$\\simlt -1.0$, are consistent with the same \\bv$\\approx 0.8$ color that we measure for the streams and stellar envelope immediately beyond M87 (Bruzual \\& Charlot 2003). This color similarity further bolsters the case that M87's extended stellar envelope, the tidal streams from disrupted galaxies, and the truly intra-cluster stars all form a single population, formed through similar mechanisms of tidal stripping during the cluster's hierarchical assembly." }, "1003/1003.1055_arXiv.txt": { "abstract": "Discussion is given of non-linear soliton behavior including coupling between electrostatic and electromagnetic potentials for non-relativistic, weakly relativistic, and fully relativistic plasmas. For plasma distribution functions that are independent of the canonical momenta perpendicular to the soliton spatial structure direction there are, in fact, no soliton behaviors allowed because transverse currents are zero. Dependence on the transverse canonical momenta is necessary. When such is the case, it is shown that the presence or absence of a soliton is intimately connected to the functional form assumed for the particle distribution functions. Except for simple situations, the coupled non-linear equations for the electrostatic and electromagnetic potentials would seem to require numerical solution procedures. Examples are given to illustrate all of these points for non-relativistic, weakly relativistic, and fully relativistic plasmas. ", "introduction": "Since its discovery in 1959, the Weibel instability \\citep{wei59:wei,fri59:wei} has been, and continues to be, a subject of intense research. Part of this focus is due to the fact that the Weibel instability excites aperiodic fluctuations, i.\\,e., purely growing waves that do not propagate out of the system. Therefore, the Weibel instability has been considered to be responsible for the creation of seed magnetic fields in the early universe \\citep{sch03:cos,sak04:mag,sch05:ori}, which act as a progenitor to the large-scale magnetic field that we observe today in all (spiral) galaxies \\citep{bec96:mag}. Likewise, Weibel fluctuations can provide the necessary dissipation of bulk velocity that leads to the formation of shock waves \\citep{tau06:grb}. In highly relativistic jets that occur during events such as GRBs (Gamma-Ray Bursts) or in extreme objects such as AGNs (Active Galactic Nuclei), plasma instabilities \\citep{sch02:agn} and particularly the Weibel instability are, thus, ultimately responsible for the radiation observed at Earth. Weibel modes and their asssociated non-linear structures \\citep{byc03:wei} also play a role in the radiative cooling of relativistic particles in blazar and gamma-ray burst sources \\citep{sch07:coo}. Furthermore, Weibel modes can also be excited in quantum plasmas \\citep{haa08:qua,haa08:mac}, thus generalizing the classical Weibel instability equation; here, a connection has been made to (dark) soliton waves \\citep{shu06:las}. It is also worth noticing that there exist experimental verifications of Weibel \\citep{wei01:exp} and soliton \\citep{bor02:exp} modes in laser plasmas, thus emphasizing the broad applicability of the underlying mechanism, which converts the free energy from anisotropic distributions into magnetic field energy. In analytical studies, such soliton modes have been used to create magnetic turbulence \\citep{kin73:tur,wei77:pla}. The underlying analytical investigation of the non-linear aspects of the classic Weibel instability made use of the fact that the classic Weibel instability excites only transverse, electromagnetic fluctuations \\citep{usr06:nli}. For asymmetric distributions, it was shown that the range of unstable wavenumbers is reduced to one single unstable wave mode, which reminds one of solitary structures that are based on single wavenumbers, too. For the case of transverse electromagnetic modes, it was shown that, depending on the exact form of the distribution function, spatially limited structures are produced (solitons). From the radiation pattern of particles scattered in soliton modes \\citep{usr09:rad}, it is known that there are many similarities to synchrotron radiation. Motivated by the fact that the energy output of the particles is mostly referred to as synchrotron radiation and inverse Compton scattering, it is a necessity to take into account other radiation processes, also because synchrotron radiation requires the presence of a background magnetic field. Until now, the origin and themechanism for such a field have been discussed intensely \\citep{med99:grb,med07:wei,med07:jit}. The train of thoughts is the following: (i) in reality, particle velocity distributions should almost always be somewhat asymmetric; (ii) in this case, unstable plasma modes are monochromatic, as has been demonstrated analytically \\citep{tau06:is1}; (iii) such isolated structures can lead to soliton modes as has been shown for purely electromagnetic Weibel modes \\citep{usr06:nli}; (iv) in relativistic outflows such as for GRBs and AGN jets, all kinds of plasma instabilities are expected to arise. Hence, it is most important to discuss the radiation pattern for such scenarios that is generated by particles being scattered in such magnetic structures. In principle, the method of obtaining the differential frequency spectrum is similar to that for synchrotron radiation \\citep{usr09:rad}. The basic difference is that for synchrotron radiation the magnetic field is considered to be constant and the electron moves in circles perpendicular to the magnetic field, while in the case of the soliton the electrons move mostly linearly and are deflected via the Lorentz force. Thus, the radiation is not produced by acceleration through a constant background field but, instead, is caused by an interaction of the electrons with highly varying magnetic and electric fields. Inspired by the procedure shown in \\citet{usr06:nli}, a number of subsequent, more detailed, investigations revealed the exotic properties of the (linear) Weibel instability that are unfolded in the case of totally asymmetric distribution functions \\citep{tau06:is1,tau06:is2,tau07:evi,tau07:tun}. For such distributions, the electrostatic and electromagnetic wave modes are coupled, and it was shown that any unstable Weibel mode must be isolated, i.\\,e., restricted to a single discrete wavenumber. Specific examples for the distribution function illustrated that isolated Weibel modes are excited. Even if one allows for a small real part of the frequency, the isolated Weibel modes persist \\citep{tau07:wea}. Such weakly propagating unstable modes are important for oblique wave propagation, because for such cases the growth rate of unstable waves is maximal \\citep{bre04:bpi,bre05:bpi}. It is then appropriate to ask how non-linear soliton modes are influenced when one includes the coupling between electrostatic and electromagnetic potentials. The purpose of this paper is to explore that question. The paper is organized as follows. In Sec.~\\ref{tech}, the Vlasov equation is transformed and the non-linear wave equations are derived. In Sec.~\\ref{nonrel}, the non-relativistic soliton behavior is derived and examples are given for two basic types of distribution functions. In Sec.~\\ref{rel}, the weakly and the fully relativistic behaviors are derived, which requires approximations as regards the transformation of the volume element in momentum space. Furthermore, in Secs.~\\ref{nonrel} and \\ref{rel}, solutions are derived for two limiting cases regarding the relative values of the electrostatic and vector potentials. In Sec.~\\ref{summ}, the results are summarized and discussed. Detailed explanations of the integral transformations in the cases of non-relativistic and relativistic plasmas are given in \\ref{nr_transf} and \\ref{r_transf}, respectively. ", "conclusions": "\\label{summ} While the linear Weibel instability has been thoroughly investigated, regrettably the same cannot be said of the non-linear behavior including the coupling of electromagnetic and electrostatic effects. The exploration of the non-linear aspects given here has uncovered a variety of effects that are germane to future investigation for both non-relativistic and relativistic plasmas. While in both the non-relativistic and the relativistic case, the classical constants of motion are total energy and generalized momentum, the problem is to obtain a \\emph{closed form} expression for $\\gamma$ (or, in the notation used here, $p_x$) in terms of these constants. Because of the coupled effects of the magnetic fields and the electrostatic fields such appears not to be an easily tractible problem as shown in text. It has been shown previously that Weibel isolated modes (which, subsequently, can develop soliton modes) are retained in analytical calculations even if one allows for ``classic'' extended unstable wavenumber ranges. Because such structures develop only in asymmetric plasmas, they serve as an indicator for the asymmetry of the particle distribution function. Therefore, because precisely symmetric plasmas are difficult to achieve in nature, isolated Weibel modes will be ubiquitous, as shown recently \\citep{tau07:wea}. Furthermore, even if the unstable wave modes are allowed to have a ``weak'' propagating component, the isolated Weibel modes are still generated. Hence, soliton structures should \\emph{always} be taken into consideration when investigating: (i) instabilities in (relativistic) plasmas in general; (ii) non-linear behavior of the resulting unstable modes; (iii) particle radiation patterns due to scattering in such modes. Perhaps the most significant theme is that the occurrence of a non-linear Weibel-like soliton requires that the distribution functions be dependent on all three of the characteristic constants. Without such a dependence (and in particular with no explicit dependence on characteristic constants perpendicular to the spatial variation direction of the soliton) then there is no electromagnetic current and so no soliton. This point was demonstrated for non-relativistic, weakly relativistic, and fully relativistic plasma situations. Even then, the functional behavior of the distribution functions on the three characteristic variables was shown, by explicit examples, to play a fundamental role in determining the structure of the non-linear equations for the coupled electromagnetic and electrostatic fields. Cases were given where no soliton was possible, where solitons existed only for decoupled electromagnetic field with no electrostatic component, and where changes in the distribution functions altered the non-linear field equations so markedly that each situation had to be considered anew. The characteristic constants could be written down in closed form for the non-relativistic and weakly relativistic situations, and the constants could be approximated in the fully relativistic plasma situation for weak electrostatic (electromagnetic) effects on an electromagnetic (electrostatic) field. Nevertheless the complexity of the resulting non-linear field equations is daunting. Except for simple situations it has so far not proven possible to solve such non-linear equations in either particular cases or the general case for chosen distribution functions. One suspects that only numerical procedures will allow deeper insight into the classes of functional behavior for distribution functions that allow solitons, for the spatial structure of such solitons, and for the relative contributions of the electrostatic and electromagnetic fields to any such solitons. Future work should attempt to investigate the modifications of the radiation pattern due to particle scattering in such soliton structures. In doing so, the question can be explored if and to what extent the radiation spectrum in relativistic outflows deviates from pure synchrotron radiation. \\ack {\\it We thank the anonymous referee for scrutizining our manuscript ever so thoroughly. This work was partially supported by the Deutsche Forschungsgemeinschaft (DFG) through grant No.\\ \\mbox{Schl~201/17-1}, Sonderforschungsbereich 591, and by the German Academy of Natural Scientists Leopoldina Fellowship Programme through grant LDPS 2009-14 funded by the Federal Ministry of Education and Research (BMBF).} \\appendix" }, "1003/1003.1780_arXiv.txt": { "abstract": "We present the first results from a near-IR spectroscopic campaign of the Cl1604 supercluster at $z\\sim0.9$ and the cluster RX J1821.6+6827 at $z\\sim0.82$ to investigate the nature of [OII] $\\lambda$3727\\AA\\ emission in cluster galaxies at high redshift. Of the 401 members in Cl1604 and RX J1821+6827 confirmed using the Keck II/DEIMOS spectrograph, 131 galaxies have detectable [OII] emission with no other signs of current star formation activity, as well as strong absorption features indicative of a well-established older stellar population. The combination of these features suggests that the primary source of [OII] emission in these galaxies is \\emph{not} a result of star formation processes, but rather due to the presence of a Low-Ionization Nuclear Emission-Line Region (LINER) or Seyfert component. Using the NIRSPEC spectrograph on the Keck II 10-m telescope, 19 such galaxies were targeted, as well as six additional [OII]-emitting cluster members that exhibited signs of ongoing star formation activity. Nearly half ($\\sim$47\\%) of the 19 [OII]-emitting, absorption-line dominated galaxies exhibit [OII] to H$\\alpha$ equivalent width (EW) ratios higher than unity, the typical observed value for star-forming galaxies, with an EW distribution similar to that observed for LINERs at low redshift. A majority ($\\sim$68\\%) of these 19 galaxies are classified as LINER/Seyfert based primarily on the emission-line ratio of [NII] $\\lambda$6584\\AA\\ and H$\\alpha$. The fraction of LINER/Seyferts increases to $\\sim$85\\% for red [OII]-emitting, absorption-line dominated galaxies. The LINER/Seyfert galaxies in our Cl1604 sample exhibit average $L$([OII])/$L$(H$\\alpha$) ratios that are significantly higher than that observed in populations of star-forming galaxies, suggesting that [OII] is a poor indicator of star formation in a significant fraction of high-redshift cluster members. From the prevalence of [OII emitting, absorption-line dominated galaxies in both systems and the fraction of such galaxies that are classified as LINER/Seyfert, we estimate that at least $\\sim$20\\% of galaxies in high-redshift clusters with $M_{\\star}>10^{10}-10^{10.5}$ $M_{\\odot}$ contain a LINER/Seyfert component that can be revealed with line ratios. We also investigate the effect such a population has on the global star formation rate of cluster galaxies and the post-starburst fraction, concluding that LINER/Seyferts must be accounted for if these quantities are to be physically meaningful. ", "introduction": "At low redshift, the final result of galaxy processing and gas depletion in cluster galaxies is widely observed. Most galaxy populations in low-redshift clusters are dominated by bright early-type galaxies, primarily devoid of star formation (Dressler et al.\\ 1985, 2004; Balogh et al.\\ 1997; Hashimoto et al.\\ 1998; Lewis et al.\\ 2002; G{\\'o}mez et al. 2003; Pimbblet et al.\\ 2006). At higher redshifts ($z\\sim$ 0.4-1) where this processing has had less time to occur, the fraction of late-type and active or recently star-forming galaxies increases (Dressler \\& Gunn 1988; Couch et al.\\ 1994; Dressler et al.\\ 1997, 2004, 2009; van Dokkum et al.\\ 2000; Lubin et al.\\ 2002; Poggianti et al.\\ 2006; Oemler et al.\\ 2009). However, the physical processes that are responsible for the quenching of star formation and the transformation of disk galaxies to dormant spheroids over the last $\\sim$7 Gyr are still not well understood. To accurately quantify this evolution, it is essential to use diagnostics that are valid and accessible across a broad redshift range. To determine the rate at which a galaxy is forming stars, the H$\\alpha$ line at 6563\\AA\\ is typically used, as it is a relatively dust-independent measure of the star formation rate (SFR) in the last 10 Myr. As H$\\alpha$ moves out of the optical window other spectral lines must be used to determine galaxy SFRs. Many higher redshift surveys (0.6$\\leq z \\leq$1.4) use instead the [OII] doublet at 3727\\AA\\ as a proxy for H$\\alpha$ since it is traditionally associated with nebular regions of current star formation and is less sensitive to stellar absorption than higher order Balmer lines (e.g., Cooper et al.\\ 2006; Vergani et al.\\ 2008). However, a comprehensive study by Yan et al.\\ (2006; hereafter Y06) of 55,000 Sloan Digital Sky Survey (SDSS) galaxies at low redshift (0.07$\\leq z \\leq$ 0.1) suggests that [OII] emission is a poor indicator of the SFR in many galaxies. While a large fraction of blue star-forming galaxies in the sample have appreciable [OII] emission, approximately 40\\% of the red, early-type galaxies also show moderate to strong [OII] emission. In 91\\% of the latter the [OII] emission likely does not originate from normal star formation processes. Rather, the strengths of [NII] $\\lambda$6584\\AA\\ relative to H$\\alpha$ and [OIII] $\\lambda$5007\\AA\\ relative to H$\\beta$ $\\lambda$4861\\AA\\ (i.e., a BPT diagram; Baldwin, Phillips, \\& Terlevich 1981) for those [OII]-emitting red-sequence galaxies with all five features detected indicate that the line emission is related either to active galactic nuclei (AGNs) or other galactic processes not associated with star formation. Most commonly, the source of [OII] emission in these galaxies is related to processes associated with low-ionization nuclear emission line regions (LINERs). This result has significant consequences for galaxy evolution studies. For galaxy populations that are currently forming stars in addition to having LINER activity, using [OII] as an SFR indicator results in an overestimate of the global SFR. More importantly, since [OII] probes both the star formation and LINER activity, accurately identifying galaxies in a transitory phase following the truncation of a star formation event (i.e., ``K+A\" or ``E+A\" galaxies; Dressler \\& Gunn 1983; Dressler \\& Gunn 1992) becomes more difficult when [OII] is used as the sole SFR indicator. Although rare ($\\la$2\\%) among bright galaxies both in nearby clusters and in the local and distant field populations (Zabludoff et al.\\ 1996; Dessler et al.\\ 1999; Goto et al.\\ 2003; Tran et al.\\ 2003; Quintero et al.\\ 2004; Yan et al.\\ 2009), K+A galaxies typically comprise a significant fraction (15\\%-25\\%) of the galaxy populations in distant clusters (Dressler et al.\\ 1999, 2004; Tran et al.\\ 2003; Oemler et al.\\ 2009). The correct classification of such galaxies is a vital step in linking the large number of star-forming, disk galaxies seen in high-redshift clusters to the quiescent, early-type galaxies observed in their local counterparts (e.g.; Poggianti et al.\\ 1999; Oemler et al.\\ 2009; Wild et al.\\ 2009). The K+A classification is based on two physical properties of galaxies: significant recent star formation and the absence of current star formation. The second criterion as applied to high-redshift galaxies ($z\\ga0.3$) typically requires that the galaxy spectra be essentially devoid of [OII] emission (Dressler \\& Gunn 1992; Zabludoff et al.\\ 1996; Balogh et al.\\ 1999; Dressler et al.\\ 2004; Oemler et al.\\ 2009). If this phase of LINER emission is a typical stage of galaxy evolution, excluding all [OII] emitters from K+A samples severely underestimates the fraction of galaxies that are truly ``post-starburst\" or ``post-star-forming\". In the low redshift sample of Y06 a large fraction ($\\sim$80\\%) of the K+A population (selected to exclude current star formation on the basis of the H$\\alpha$ line) show appreciable levels of [OII] emission. If cluster galaxies at high redshift share similar properties, it is necessary to understand where galaxies that exhibit LINER emission lie along the evolutionary chain in clusters and what role, if any, LINER emission has in truncating star formation in cluster galaxies. To study the properties of this phenomenon at high redshift we use the extensive spectroscopic database from the Observations of Redshift Evolution in Large Scale Environments survey (ORELSE; Lubin et al.\\ 2009, hereafter L09). The ORELSE survey is an ongoing multi-wavelength campaign mapping out the environmental effects on galaxy evolution in the large scale structures surrounding 20 known clusters at moderate redshift ($0.6 \\leq z \\leq 1.3$). In particular this paper focuses on two structures, the Cl1604 supercluster at $z\\approx0.9$ and the X-ray-selected cluster RX J1821.6+6827 at $z\\approx0.82$. Combining the wealth of previous ORELSE observations in these fields with newly obtained Keck II Near-Infrared Echelle Spectrograph (NIRSPEC) spectroscopy of 25 galaxies, we investigate the pervasiveness of LINER emission in cluster galaxies at high redshift and the properties of galaxies whose optical emission lines are dominated by this phenomenon. The remainder of the paper is organized as follows: in \\S2 we discuss the two high redshift structures targeted by this survey. \\S3 describes the optical spectroscopy and discusses the target selection, observation, and reduction of the near-infrared spectroscopy. In \\S4 we describe our methods for equivalent width (EW), relative flux, and absolute flux measurements, including absolute spectrophotometric calibration and extinction corrections. In \\S5 we present our results and discuss their consequences for high redshift galaxy surveys. \\S6 presents our conclusions. We adopt a standard concordance $\\Lambda$CDM cosmology with $H_{0}$ = 70 km s$^{-1}$, $\\Omega_{\\Lambda}$ = 0.7, and $\\Omega_{M}$ = 0.3. All EW measurements are presented in the rest frame and all magnitudes are given in the AB system (Oke \\& Gunn 1983; Fukugita et al.\\ 1996). ", "conclusions": "In this study we have identified a population of [OII]-emitting, absorption-line dominated galaxies in high-redshift clusters that are primarily powered by LINER or Seyfert activity as evidenced by their optical and NIR spectroscopy. Of the 486 galaxies in the Cl1604 supercluster ($z\\sim0.9$) and the X-ray selected cluster RX J1821 ($z\\sim0.82$) for which we have obtained optical spectroscopy, 25 galaxies were selected for follow-up NIRSPEC \\emph{J}-band spectra. These galaxies were primarily selected to be a representative sample of a population of cluster galaxies that have optical spectra that exhibits moderately strong [OII] emission with no other spectral indicators of current star formation, as well as strong absorption-line features indicative of a well-established older stellar population. Galaxies with these spectral properties (which we have termed ``priority 1\") make up a third of the population in both structures. The main results of this investigation are given below: \\begin{itemize} \\item We find that nearly half ($\\sim$47\\%) of the [OII]-emitting, absorption-line dominated galaxies in this study have high levels of [OII] emission relative to the amount of H$\\alpha$ emission. \\item Nearly all of galaxies with high levels of [OII] emisssion relative to H$\\alpha$ and a majority ($\\sim$68\\%) of the targeted [OII]-emitting, absorption-line dominated galaxies have emission profiles dominated by a LINER or Seyfert component (referred to as LINER/Seyfert), primarily revealed by the flux ratio of H$\\alpha$ and [NII] $\\lambda$6584\\AA. \\item This LINER/Seyfert fraction has a strong dependence on color; $\\sim$85\\% of targeted [OII]-emitting, absorption-line dominated galaxies on the red sequence of the two structures have dominant LINER or Seyfert components as compared with only 33\\% of blue galaxies. \\item The bulk of our LINER/Seyfert population have observed EW([OII])/EW(H$\\alpha$) values significantly higher than unity, suggesting that a majority of these galaxies are powered by LINER and not Seyfert emission. The remainder are either powered by Seyfert emission or undergoing a transition phase in which both LINER/Seyfert and ongoing star-forming activity is occurring. \\item In addition to being primarily red in color, galaxies powered by LINER or Seyfert emission are almost exclusively compact early-type galaxies, contrasting sharply with the late-type morphologies of the star-forming galaxies observed in our sample. \\item The lower limit to the average extinction corrected \\emph{L}([OII])/\\emph{L}(H$\\alpha$) in galaxies classified as LINER/Seyfert in our Cl1604 sample is 1.51$\\pm$0.28, higher than that of star-forming galaxies at low redshift (i.e., 1.2; Kewley et al.\\ 2004) and that of the average star-forming galaxy observed in our sample (1.04$\\pm$0.27). We investigate various extinction schemes and metallicity differences in the samples and determine that the high levels of [OII] luminosity relative to H$\\alpha$ in LINER/Seyfert galaxies are not due to dust or metallicity effects and are rather the result of emission from the LINER/Seyfert itself. \\end{itemize} From the statistical properties of this sample we use the color distribution and prevalence of [OII]-emitting, absorption-line dominated galaxies in our entire Cl1604 and RX J1821 DEIMOS database to determine the fraction of cluster galaxies at high redshift that contain a LINER/Seyfert component. For galaxies with stellar masses equal to or greater than our stellar mass limit (roughly $M_{\\star}=10^{10}-10^{10.5} M_{\\odot}$) we estimate that $>20\\%$ of galaxies in these structures contain a LINER/Seyfert component. Additionally, the fraction of galaxies on the red-sequence that have appreciable [OII] emission (most of which is likely due to a LINER or Seyfert) is $\\sim$54\\% in both structures, similar to the fraction observed in red SDSS galaxies at low redshift. Since these two systems are in significantly different dynamical stages, these results imply that whatever mechanism is powering the emission in these galaxies is active for much longer than the dynamical timescale of the clusters and is not sensitive to the global environment in which a galaxy resides. We have established that a large fraction of high-redshift galaxies, especially those on the red sequence, have [OII] emission directly resulting from a process unrelated to star formation. This result has significant consequences for surveys of high-redshift galaxies that use [OII] as a star formation indicator. The global star formation rate as calculated from the extinction corrected [OII] line luminosity for the LINER/Seyfert galaxies is significantly higher than the same quantity derived from the H$\\alpha$ feature, itself an over-estimate of the actual SFR (due to H$\\alpha$ flux originating from the LINER/Seyfert component). We conclude that high-redshift galaxy surveys that rely on [OII] as an SFR indicator will be non-negligibly biased by LINER/Seyfert activity. While other recombination lines (e.g., H$\\alpha$, H$\\beta$) provide better estimates of the instantaneous SFR than [OII] when observable, the problem of residual LINER/Seyfert flux still remains. We also investigate the effect of the LINER/Seyfert population on the selection of transititory ``post-starburst\" galaxies, a population that is of considerable interest for many cluster and field evolutionary studies. We find that including H$\\delta$ strong LINER/Seyfert galaxies increases the percentage of post-starburst galaxies in the two structures to 18.8\\%, more than double the 9.3\\% obtained using traditional selection methods. While some LINER/Seyfert galaxies likely still have some residual star formation, the requirement that [OII] be absent for a galaxy to be classified as post-starburst is too conservative and will result in a post-starburst sample that is severely incomplete. Due to the prevalence of LINER/Seyfert activity across a large range of environments at both high and low redshift, we conclude that LINER/Seyferts must be carefully accounted for when interpreting post-starburst populations in the context of galaxy evolution.\\\\[14pt] ~~ ~~ We thank Jeff Newman and Michael Cooper for guidance with the \\emph{spec2d} reduction pipeline and for the many useful suggestions and modifications necessary to reduce our DEIMOS data. We thank Nick Konidaris for help with DEIMOS flux calibration and for useful discussions on equivalent width measurements in DEIMOS data. We also thank Chris Fassnacht for providing his notes on NIRSPEC observations and the Keck II support astronomers for their help with the observations. B.C.L. thanks Gary Creason for his patience and guidance. This material is based upon work supported by the National Aeronautics and Space Administration under Award NNG05GC34G for the Long Term Space Astrophysics Program. Additional support for this program was provided by NASA through a grant HST-GO-11003 from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc. BCL and LML would also like to acknowledge support from NSF-AST-0907858. The spectrographic data presented herein were obtained at the W.M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California, and the National Aeronautics and Space Administration. The Observatory was made possible by the generous financial support of the W.M. Keck Foundation. We thank the indigenous Hawaiian community for allowing us to be guests on their sacred mountain; we are most fortunate to be able to conduct observations from this site. \\newpage \\begin{sidewaystable*}[tp] \\centering \\begin{center} \\caption{Cl1604 NIRSPEC Observations} \\label{tab:obs1} \\begin{tabular}{ccccccccccc}\\hline \\hline $ID$ & Galaxy Number & $\\alpha_{2000}$ & $\\delta_{2000}$ & $z$ & \\emph{F}606\\emph{W} & \\emph{F}814\\emph{W} & \\emph{F}606\\emph{W}-\\emph{F}814\\emph{W} & Setup Number & $\\tau_{\\rm{exp}}$ & Class\\\\ \\hline J160344$+$432429 & 0 & 240.9322226 & 43.4079759 & 0.9023 & 23.60 & 21.97 & 1.63 & 1 & 5$\\times$900s\\tablenote{One exposure was not usable due to guider issues.} & Priority1\\\\ [4pt] J160344$+$432428 & 1 & 240.9325941 & 43.4077202 & 0.9024 & 22.11 & 23.16 & 1.05 & 1 & 5$\\times$900s$\\tablenotemark{a}$ & Priority1\\\\[4pt] J160345$+$432419 & 2 & 240.9375426 & 43.4051985 & 0.8803 & 24.01 & 22.46 & 1.55 & 1 & 5$\\times$900s$\\tablenotemark{a}$ & Priority1\\\\[4pt] J160342$+$432406 & 3 & 240.9247136 & 43.4016956 & 0.8986 & 25.10 & 23.45 & 1.64 & 2 & 4$\\times$900s & Priority1\\\\[4pt] J160342$+$432403 & 4 & 240.9250684 & 43.4006981 & 0.8959 & 24.12 & 22.38 & 1.74 & 2 & 4$\\times$900s & Priority1\\\\[4pt] J160330$+$432208 & 5 & 240.8732075 & 43.3687725 & 0.8983 & 22.84 & 21.40 & 1.44 & 3 & 4$\\times$900s & Priority1\\\\[4pt] J160329$+$432204 & 6 & 240.8697693 & 43.3676967 & 0.9045 & 23.54 & 22.12 & 1.42 & 3 & 4$\\times$900s & Priority1\\\\[4pt] J160416$+$431021 & 7 & 241.0657080 & 43.1725670 & 0.8999 & 24.91 & 23.75 & 1.16 & 4 & 4$\\times$900s & Priority3 (Filler)\\\\[4pt] J160416$+$431017 & 8 & 241.0648269 & 43.1713681 & 0.8999 & 22.99 & 21.49 & 1.49 & 4 & 4$\\times$900s & Priority1\\\\[4pt] J160404$+$432445 & 9 & 241.0150297 & 43.4124202 & 0.9017 & 23.60 & 22.72 & 0.88 & 5 & 4$\\times$900s & Priority3 (Filler) \\\\[4pt] J160403$+$432436 & 10 & 241.0108301 & 43.4099384 & 0.9015 & 23.34 & 21.82 & 1.52 & 5 & 4$\\times$900s & Priority1\\\\[4pt] J160429$+$431956 & 11 & 241.1195420 & 43.3321920 & 0.9185 & 24.48 & 23.76 & 0.72 & 6 & 2$\\times$900s & Priority3 (Filler)\\\\[4pt] J160428$+$431953 & 12 & 241.1171400 & 43.3312750 & 0.9198 & 23.37 & 21.75 & 1.61 & 6 & 2$\\times$900s & Priority1\\\\[4pt] J160406$+$431542 & 13 & 241.0276264 & 43.2615940 & 0.8674 & 23.44 & 21.57 & 1.88 & 8 & 4$\\times$900s & Priority1+Red Sequence\\\\[4pt] J160407$+$431539 & 14 & 241.0299022 & 43.2607188 & 0.8676 & 24.48 & 22.81 & 1.67 & 8 & 4$\\times$900s & Priority1+Red Sequence\\\\[4pt] J160426$+$431423 & 15 & 241.1100259 & 43.2397136 & 0.8676 & 23.43 & 21.48 & 1.95 & 9 & 3$\\times$900s & Priority1+Red Sequence\\\\[4pt] J160426$+$431419 & 16 & 241.1092254 & 43.2386527 & 0.8658 & 22.80 & 20.84 & 1.96 & 9 & 3$\\times$900s & Priority1+Red Sequence\\\\[4pt] J160427$+$431501 & 17 & 241.1104629 & 43.2503720 & 0.8601 & 23.77 & 21.80 & 1.97 & 10 & 4$\\times$900s & Priority1+Red Sequence\\\\[4pt] J160426$+$431439 & 18 & 241.1086670 & 43.2441610 & 0.8710 & 24.94 & 23.35 & 1.60 & 10 & 4$\\times$900s & Priority1+Red Sequence\\\\[4pt] J160406$+$431825 & 19 & 241.0243330 & 43.3068500 & 0.9189 & 25.59 & 23.78 & 1.81 & 11 & 5$\\times$900s$\\tablenotemark{a}$ & Priority2 (Filler)\\\\[4pt] J160406$+$431809 & 20 & 241.0266087 & 43.3024702 & 0.9195 & 24.29 & 22.36 & 1.93 & 11 & 5$\\times$900s$\\tablenotemark{a}$ & Priority1+Red Sequence\\\\[6pt]\\hline \\end{tabular} \\end{center} \\end{sidewaystable*} \\begin{sidewaystable*}[tp] \\caption{RX J1821 NIRSPEC Observations} \\label{tab:obs2} \\begin{center} \\centering \\begin{tabular}{ccccccccccc} \\hline \\hline $ID$ & Galaxy Number & $\\alpha_{2000}$ & $\\delta_{2000}$ & $z$ & $i\\arcmin$ & $z\\arcmin$ & $i\\arcmin-z\\arcmin$ & Setup Number & $\\tau_{\\rm{exp}}$ & Class\\tablenote{Note that no color cut was imposed on any RX J1821 targets.} \\\\ \\hline J182110+682350 & 21 & 275.29224260 & 68.39710400 & 0.7960 & 21.41 & 20.76 & 0.65 & 7 & 4$\\times$900s & Priority1\\\\[4pt] J182108+682329 & 22 & 275.28199650 & 68.39415620 & 0.8134 & 21.75 & 21.00 & 0.75 & 7 & 4$\\times$900s & Priority1\\\\[4pt] J182121+682715 & 23 & 275.33619440 & 68.45408210 & 0.8092 & 23.32 & 23.03 & 0.30 & 12 & 3$\\times$900s & Priority2 (filler)\\\\[4pt] J182123+682714 & 24 & 275.34600740 & 68.45392090 & 0.8093 & 23.72 & 23.36 & 0.36 & 12 & 3$\\times$900s & Priority2 (filler)\\\\[6pt]\\hline \\end{tabular} \\end{center} \\end{sidewaystable*} \\begin{sidewaystable*}[tp] \\caption{Spectral Properties of the NIRSPEC Sample} \\label{tab:spec} \\begin{center} \\centering \\begin{tabular}{ccccccccc} \\hline \\hline $ID$ & Galaxy Number & $z$ & EW([OII])\\tablenote{Measured in the rest-frame} & EW(H$\\alpha$)$\\tablenotemark{a}$ & Class EW([OII])/EW(H$\\alpha$)& F(H$\\alpha$) & F([NII]) & Emission Class \\\\ & & & (\\AA) & (\\AA) & & ($\\times10^{-18}$ erg s$^{-1}$ cm$^{-2}$) & ($\\times10^{-18}$ erg s$^{-1}$ cm$^{-2}$) & \\\\ \\hline J160344$+$432429 & 0 & 0.9023 & 18.8$\\pm$1.7 & 14.8$\\pm$2.1 & Low & 21.8$\\pm$2.9 & 17.5$\\pm$2.6& \\emph{LINER/Seyfert} \\\\[4pt] J160344$+$432428 & 1 & 0.9024 & 22.5$\\pm$1.3 & 13.3$\\pm$1.8 & Low & 18.9$\\pm$2.1 & 13.4$\\pm$2.1& \\emph{LINER/Seyfert} \\\\[4pt] J160345$+$432419 & 2 & 0.8803 & 10.7$\\pm$1.3 & 6.1$\\pm$1.7 & Low & 8.1$\\pm$2.0 & 6.5$\\pm$1.9& \\emph{LINER/Seyfert} \\\\[4pt] J160342$+$432406 & 3 & 0.8986 & 31.6$\\pm$1.7 & 6.7$\\pm$1.8 & High & 12.7$\\pm$1.6 & 17.4$\\pm$2.9& \\emph{LINER/Seyfert} \\\\[4pt] J160342$+$432403 & 4 & 0.8959 & 58.6$\\pm$1.4 & 17.0$\\pm$1.3 & Low & 44.5$\\pm$2.2 & 53.2$\\pm$3.3& \\emph{LINER/Seyfert} \\\\[4pt] J160330$+$432208 & 5 & 0.8983 & 16.1$\\pm$0.8 & 45.8$\\pm$1.1 & Low & 173.0$\\pm$3.6& 85.0$\\pm$2.9& \\emph{Star forming} \\\\[4pt] J160329$+$432204 & 6 & 0.9045 & 8.7 $\\pm$1.0 & 8.2$\\pm$1.1 & Low & 31.4$\\pm$4.1 & 13.3$\\pm$2.6 & \\emph{Star forming} \\\\[4pt] J160416$+$431021 & 7 & 0.8990 & 17.4$\\pm$2.3 & 38.6$\\pm$14.3 & Low & 6.4$\\pm$2.1 & $<$3.7\\tablenote{3$\\sigma$ upper limit} & \\emph{Star-Forming} \\\\[4pt] J160416$+$431017 & 8 & 0.8999 & 17.8$\\pm$1.4 & 33.3$\\pm$1.7 & Low & 119.0$\\pm$4.8 & 38.4$\\pm$3.1& \\emph{Star forming} \\\\[4pt] J160404$+$432445 & 9 & 0.9017 & 31.5$\\pm$0.9 & 41.8$\\pm$4.2 & Low & 35.6$\\pm$2.9 & 4.5$\\pm$1.5& \\emph{Star forming} \\\\[4pt] J160403$+$432436 & 10 & 0.9015 & 10.2$\\pm$1.9 & 3.2$\\pm$0.3 & High & 18.0$\\pm$3.4 & 20.7$\\pm$2.4& \\emph{LINER/Seyfert} \\\\[4pt] J160429$+$431956 & 11 & 0.9185 & 55.3$\\pm$1.9 & 52.3$\\pm$16.3 & Low & 24.8$\\pm$5.6 & 6.1$\\pm$3.8& \\emph{Star forming} \\\\[4pt] J160428$+$431953 & 12 & 0.9198 & 9.0$\\pm$1.1 & 25.4$\\pm$1.8& Low & 102.0$\\pm$6.9 & 55.5$\\pm$6.0& \\emph{Star forming} \\\\[4pt] J160406$+$431542 & 13 & 0.8674 & 9.9$\\pm$0.9 & 0.1$\\pm$0.7 & High & $<$14.3$\\tablenotemark{b}$ & 10.6$\\pm$2.4 & \\emph{LINER/Seyfert} \\\\[4pt] J160407$+$431539 & 14 & 0.8676 & 23.7$\\pm$3.3 & 7.2$\\pm$ 2.5 & Low & 19.4$\\pm$6.1 & 12.3$\\pm$4.1 & \\emph{LINER/Seyfert} \\\\[4pt] J160426$+$431423 & 15 & 0.8676 & 7.9$\\pm$0.6 & 0.4$\\pm$0.7 & High & $<$16.5$\\tablenotemark{b}$ & 14.1$\\pm$3.0 & \\emph{LINER/Seyfert} \\\\[4pt] J160426$+$431419 & 16 & 0.8658 & 6.3$\\pm$0.7 & 1.6$\\pm$1.1 & High & $<$22.0$\\tablenotemark{b}$ & $<$27.0$\\tablenotemark{b}$ & \\emph{Ambiguous} \\\\[4pt] J160427$+$431501 & 17 & 0.8601 & 6.3$\\pm$0.7 & 0.1$\\pm$ 1.6 & High & $<$8.1$\\tablenotemark{b}$ & $<$15.7$\\tablenotemark{b}$ & \\emph{Ambiguous} \\\\[4pt] J160426$+$431439 & 18 & 0.8710 & 56.7$\\pm$6.1 & 41.8$\\pm$6.0 & Low & 64.5$\\pm$5.2 & 23.4$\\pm$4.6 & \\emph{Star forming} \\\\[4pt] J160406$+$431825 & 19 & 0.9189 & 50.9$\\pm$4.5 & 60.2$\\pm$14.6 & Low & 28.4$\\pm$5.5 & 13.4$\\pm$5.1 & \\emph{Star forming} \\\\[4pt] J160406$+$431809 & 20 & 0.9195 & 73.7$\\pm$1.4 & 14.1$\\pm$3.8 & High & 29.7$\\pm$3.6 & 15.8$\\pm$3.9& \\emph{Star forming} \\\\[4pt] J182110$+$682350 & 21 & 0.7960 & 2.6$\\pm$0.3 & 0.3$\\pm$0.3 & High & $<$20.0$^{b}$ & 23.3$\\pm$4.1 & \\emph{LINER/Seyfert} \\\\[4pt] J182108$+$682329 & 22 & 0.8134 & 26.9$\\pm$0.5 & 1.9$\\pm$0.5 & High & 11.0$\\pm$2.6 & 44.4$\\pm$3.3 & \\emph{LINER/Seyfert} \\\\[4pt] J182121$+$682715 & 23 & 0.8092 & 17.3$\\pm$1.5 & 20.6$\\pm$3.5 & Low & 37.8$\\pm$6.0 & 18.1$\\pm$4.8 & \\emph{Star forming} \\\\[4pt] J182123$+$682714 & 24 & 0.8093 & 24.2$\\pm$2.0 & 39.9$\\pm$10.0 & Low & 32.9$\\pm$3.9 & 8.9$\\pm$2.9 & \\emph{Star forming} \\\\[6pt]\\hline \\end{tabular} \\end{center} \\end{sidewaystable*} \\begin{sidewaystable*}[tp] \\caption{Star Formation Properties of the Cl1604 NIRSPEC Sample} \\label{tab:starformation} \\begin{center} \\centering \\begin{tabular}{ccccccccccc} \\hline \\hline $ID$ & Galaxy & $z$ & F(H$\\alpha$) & F([OII]) & L(H$\\alpha$)$_{obs}$ & L([OII])$_{obs}$ & L($H\\alpha$)\\tablenote{Corrected using a constant extinction of E(B-V)=0.3}$_{corr}$ & SFR(H$\\alpha$)\\tablenotemark{a}$_{corr}$ & L(OII)\\tablenotemark{a}$_{corr}$ & SFR([OII])\\tablenotemark{a}$_{corr}$ \\\\ & & & & & $10^{40}$ erg s$^{-1}$ & $10^{40}$ erg s$^{-1}$ & $10^{40}$ erg s$^{-1}$ & $M_{\\odot}$ yr$^{-1}$ & $10^{40}$ erg s$^{-1}$ & $M_{\\odot}$ yr$^{-1}$ \\\\\\hline J160344$+$432429 & 0 & 0.9023 & 21.8$\\pm$2.9 & 11.4 $\\pm$5.1 & 8.7$\\pm$1.2 & 4.6$\\pm$2.1 & 22.2$\\pm$6.3 & 1.7$\\pm$0.5 & 23.2$\\pm$11.9 & 1.6$\\pm$1.0\\\\[4pt] J160344$+$432428 & 1 & 0.9024 & 18.9$\\pm$2.1 & 20.8$\\pm$9.3 & 7.7$\\pm$0.8 & 8.5$\\pm$3.8 & 19.2$\\pm$5.2 & 1.5$\\pm$0.4 & 42.7$\\pm$21.8 & 3.0$\\pm$1.8 \\\\[4pt] J160345$+$432419 & 2 & 0.8803 & 8.1$\\pm$2.0 & 8.9$\\pm$ 4.0 & 3.1$\\pm$0.7 & 3.4$\\pm$1.5 & 7.8$\\pm$2.7& 0.6$\\pm$0.2 & 17.1$\\pm$8.8 & 1.2$\\pm$0.7 \\\\[4pt] J160342$+$432406 & 3 & 0.8986 & 12.7$\\pm$1.6 & 11.6$\\pm$5.2 & 5.1$\\pm$0.6 & 4.7$\\pm$2.1 & 12.8$\\pm$3.6 & 1.0$\\pm$0.3 & 23.6$\\pm$12.0 & 1.6$\\pm$1.0\\\\[4pt] J160342$+$432403 & 4 & 0.8959 & 44.5$\\pm$2.2 & 43.6 $\\pm$19.3& 17.7$\\pm$0.9 & 17.4$\\pm$7.7 & 44.5$\\pm$11.2 & 3.5$\\pm$0.9 & 87.9$\\pm$44.6 & 6.1$\\pm$3.6 \\\\[4pt] J160330$+$432208 & 5 & 0.8983 & 173.0$\\pm$3.6 & 18.1$\\pm$8.0 & 69.5$\\pm$1.5 & 7.3$\\pm$3.2 & 174.2$\\pm$43.3 & 13.8$\\pm$3.4 & 36.7$\\pm$18.6 & 2.5$\\pm$1.5 \\\\[4pt] J160329$+$432204 & 6 & 0.9045 & 31.4$\\pm$4.1 & 17.1$\\pm$7.6 & 12.8$\\pm$1.7 & 7.0$\\pm$3.1 & 32.2$\\pm$9.0 & 2.5$\\pm$0.7 & 35.3$\\pm$18.0 & 2.5$\\pm$1.4 \\\\[4pt] J160416$+$431021 & 7 & 0.8999 & 6.4$\\pm$2.1 & 9.8 $\\pm$4.3 & 2.6$\\pm$0.8 & 3.9$\\pm$1.7 & 6.4$\\pm$2.6 & 0.5$\\pm$0.2 & 19.8$\\pm$10.1 & 1.4$\\pm$0.8 \\\\[4pt] J160416$+$431017 & 8 & 0.8999 & 119.0$\\pm$4.8 & 13.9 $\\pm$6.2 & 48.0$\\pm$1.9 & 5.6$\\pm$2.5 & 120.4 $\\pm$30.2 & 9.5$\\pm$2.4 & 28.4$\\pm$14.5 & 2.0$\\pm$1.2 \\\\[4pt] J160404$+$432445 & 9 & 0.9017 & 35.6$\\pm$2.9 & 30.4$\\pm$13.4 & 14.4$\\pm$1.2 & 12.3$\\pm$5.4 & 36.2$\\pm$9.4 & 2.9$\\pm$0.7 & 62.1$\\pm$31.5 & 4.3$\\pm$2.5 \\\\[4pt] J160403$+$432436 & 10 & 0.9015 & 18.0$\\pm$3.4 & 8.1$\\pm$3.7 & 7.3$\\pm$1.4 & 3.3$\\pm$1.5 & 18.3$\\pm$5.7 & 1.4$\\pm$0.5 & 16.6$\\pm$8.5 & 1.2$\\pm$0.7 \\\\[4pt] J160429$+$431956 & 11 & 0.9185 & 24.8$\\pm$5.6 & 23.1$\\pm$10.2 & 10.5$\\pm$2.4 & 9.8$\\pm$4.3 & 26.3$\\pm$8.8 & 2.1$\\pm$0.7 & 49.5$\\pm$25.1 & 3.4$\\pm$2.0 \\\\[4pt] J160428$+$431953 & 12 & 0.9198 & 102.0$\\pm$6.9 & 11.8$\\pm$5.3 & 43.4$\\pm$2.9 & 5.0$\\pm$2.3 & 108.9$\\pm$ 27.9 & 8.6$\\pm$2.2 & 25.4$\\pm$13.1 & 1.8$\\pm$1.1 \\\\[4pt] J160406$+$431542 & 13 & 0.8674 & $<$14.3\\tablenote{3$\\sigma$ upper limit} & 9.8$\\pm$4.4 & $<$5.3$\\tablenotemark{b}$ & 3.6$\\pm$1.6 & $<$13.2$\\tablenotemark{b}$ & $<$1.0$\\tablenotemark{b}$ & 18.3$\\pm$9.4 & 1.3$\\pm$0.8 \\\\[4pt] J160407$+$431539 & 14 & 0.8676 & 19.4$\\pm$6.1 & 4.9$\\pm$2.4 & 7.2$\\pm$2.3 & 1.8$\\pm$0.9 & 17.9$\\pm$7.2 & 1.4$\\pm$0.6 & 9.1$\\pm$5.1 & 0.6$\\pm$0.4 \\\\[4pt] J160426$+$431423 & 15 & 0.8676 & $<$16.5$\\tablenotemark{b}$ & 11.3$\\pm$5.0 & $<$6.1\\tablenotemark{b} & 4.2$\\pm$1.9 & $<$15.2\\tablenotemark{b} & $<$1.2$^{b}$ & 21.0$\\pm$10.7 & 1.5$\\pm$0.9 \\\\[4pt] J160426$+$431419 & 16 & 0.8658 & $<$22.0\\tablenotemark{b} & ...\\tablenote{No flux measurement was attempted for the LRIS object, due to the uncertainty in the flux calibration.} & $<$8.1\\tablenotemark{b} & ...\\tablenotemark{c} & $<$20.2\\tablenotemark{b} & $<$1.6\\tablenotemark{b} & ...\\tablenotemark{c} & ...\\tablenotemark{c} \\\\[4pt] J160427$+$431501 & 17 & 0.8601 & $<$8.1\\tablenotemark{b} & 9.7$\\pm$4.4 & $<$2.9\\tablenotemark{b} & 3.5$\\pm$1.6 & $<$7.3\\tablenotemark{b} & $<$0.6\\tablenotemark{b} & 17.7$\\pm$9.1 & 1.2$\\pm$0.8 \\\\[4pt] J160426$+$431439 & 18 & 0.8710 & 64.5$\\pm$5.2 & 12.2$\\pm$5.5 & 24.0$\\pm$1.9 & 4.6$\\pm$2.0 & 60.2$\\pm$15.7 & 4.8$\\pm$1.2 & 23.0$\\pm$11.8 & 1.6$\\pm$0.9 \\\\[4pt] J160406$+$431825 & 19 & 0.9189 & 28.4$\\pm$5.5 & 8.0$\\pm$3.6 & 12.0$\\pm$2.4 & 3.4$\\pm$1.5 & 30.3$\\pm$9.5 & 2.4$\\pm$0.8 & 17.2$\\pm$8.9 & 1.2$\\pm$0.7 \\\\[4pt] J160406$+$431809 & 20 & 0.9195 & 29.7$\\pm$3.6 & 7.2$\\pm$3.2 & 12.6$\\pm$1.5 & 30.6$\\pm$13.5 & 31.7$\\pm$8.7 & 2.5$\\pm$0.7& 154.4$\\pm$78.3 & 10.7$\\pm$6.3 \\\\[6pt]\\hline \\end{tabular} \\end{center} \\end{sidewaystable*} \\begin{sidewaystable*}[tp] \\caption{Imaging properties of the NIRSPEC sample} \\label{tab:imaging} \\begin{center} \\centering \\begin{tabular}{ccccccccccccccccc} \\hline \\hline $ID$ & Galaxy Number & $z$ & $\\alpha_{2000}$ & $\\delta_{2000}$ & $m_{F606W}$ & $m_{F814W}$ & r$\\arcmin$ & i$\\arcmin$ & z$\\arcmin$ & K$_s$ & Color\\tablenote{As defined in \\S5.3} & Morphology\\tablenote{Done by visual inspection, M = Merger, I = Interaction, C = Chaotic, S = Spiral, Asymm = Asymmetric Disk}\\\\ \\hline J160344$+$432429 & 0 & 0.9023 & 240.9322226 & 43.4079759 & 23.6011 & 21.9680 & 23.0289 & 21.6730 & 20.9049 & ...\\tablenote{Not detected in K$_{s}$} & Red & S M I\\\\[4pt] J160344$+$432428 & 1 & 0.9024 & 240.9325941 & 43.4077202 & 23.7290 & 22.1054 & 23.1568 & 21.8849 & 21.2576 & 20.5956 & Red & S C I\\\\[4pt] J160345$+$432419 & 2 & 0.8803 & 240.9375426 & 43.4051985 & 24.0052 & 22.4560 & 23.7045 & 22.6779 & 22.0024 & 20.7284 & Blue & S0 \\\\[4pt] J160342$+$432406 & 3 & 0.8986 & 240.9247136 & 43.4016956 & 25.0951 & 23.4488 & 23.9627 & 22.7284 & 22.2418 & 21.1584 & Blue & E \\\\[4pt] J160342$+$432403 & 4 & 0.8959 & 240.9250684 & 43.4006981 & 24.1215 & 22.3766 & 23.4728 & 22.3332 & 21.5753 & 20.4692 & Red & E elongated \\\\[4pt] J160330$+$432208 & 5 & 0.8983 & 240.8732075 & 43.3687725 & 22.8355 & 21.3960 & 22.4344 & 21.5230 & 20.8415 & 19.4092 & Blue & Sc I\\\\[4pt] J160329$+$432204 & 6 & 0.9045 & 240.8697693 & 43.3676967 & 23.5419 & 22.1219 & 22.9297 & 21.8290 & 20.9793 & 20.5688 & Blue & Sa Ring\\\\[4pt] J160416$+$431021 & 7 & 0.8990 & 241.0657080 & 43.1725670 & 24.9100 & 23.7452 & 24.9532 & 24.0943 & 22.4512 & ...$^{\\rm{c}}$ & Blue& Amorphous \\\\[4pt] J160416$+$431017 & 8 & 0.8999 & 241.0648269 & 43.1713681 & 22.9863 & 21.4923 & 22.4896 & 21.2475 & 20.4742 & 19.1758 & Blue & Sc I \\\\[4pt] J160404$+$432445 & 9 & 0.9017 & 241.0150297 & 43.4124202 & 23.5993 & 22.7161 & 23.8893 & 22.8681 & 22.7997 & ...$^{\\rm{c}}$ & Blue & S asymm M?\\\\[4pt] J160403$+$432436 & 10 & 0.9015 & 241.0108301 & 43.4099384 & 23.3368 & 21.8190 & 23.0297 & 21.9360 & 21.2114 & 19.7736 & Blue & Sa/S0 \\\\[4pt] J160429$+$431956 & 11 & 0.9185 & 241.1195420 & 43.3321920 & 24.4849 & 23.7607 & 25.1182 & 23.8742 & 23.8430 & ...$^{\\rm{c}}$ & Blue & S0 peculiar I? \\\\[4pt] J160428$+$431953 & 12 & 0.9198 & 241.1171400 & 43.3312750 & 23.3654 & 21.7545 & 23.1085 & 21.9486 & 21.2242 & 19.5014 & Blue & SBb \\\\[4pt] J160406$+$431542 & 13 & 0.8674 & 241.0276264 & 43.2615940 & 23.4426 & 21.5654 & 22.5928 & 21.5389 & 20.7046 & 19.2990 & Red & S0 asymm\\\\[4pt] J160407$+$431539 & 14 & 0.8676 & 241.0299022 & 43.2607188 & 24.4806 & 22.8149 & 23.9034 & 22.9086 & 22.1914 & 20.9843 & Red & elongated E \\\\[4pt] J160426$+$431423 & 15 & 0.8676 & 241.1100259 & 43.2397136 & 23.4286 & 21.4794 & 21.9561 & 20.9147 & 19.9575 & 18.7911 & Red & E I? \\\\[4pt] J160426$+$431419 & 16 & 0.8658 & 241.1092254 & 43.2386527 & 22.7983 & 20.8419 & 21.6547 & 20.3484 & 19.4208 & 18.3566 & Red & E \\\\[4pt] J160427$+$431501 & 17 & 0.8601 & 241.1104629 & 43.2503720 & 23.7716 & 21.8000 & 22.8843 & 21.5115 & 20.5948 & 19.2840 & Red & E I? \\\\[4pt] J160426$+$431439 & 18 & 0.8710 & 241.1086670 & 43.2441610 & 24.9447 & 23.3475 & 25.0073 & 23.4083 & 22.1090 & 21.3286 & Red & Amorphous?\\\\[4pt] J160406$+$431825 & 19 & 0.9189 & 241.0243330 & 43.3068500 & 25.5927 & 23.7835 & 25.4611 & 23.9588 & 22.9134 & 21.2724 & Red & S disturbed \\\\[4pt] J160406$+$431809 & 20 & 0.9195 & 241.0266087 & 43.3024702 & 24.2884 & 22.3585 & 22.7956 & 22.0227 & 21.1015 & 19.8396 & Red & E \\\\[4pt] J182110+682350 & 21 & 0.7960 & 275.2922426 & 68.3971040 & ...\\tablenote{ACS data not available for RX J1821} & ...$^{\\rm{d}}$ & 22.5616 & 21.4127 & 20.7620 & ...\\tablenote{$K_{s}$ magnitudes not available for RX J1821} & Red & ...$^{\\rm{d}}$ \\\\[4pt] J182108+682329 & 22 & 0.8134 & 275.2819965 & 68.3941562 & ...$^{\\rm{d}}$ & ...$^{\\rm{d}}$ & 22.8643 & 21.7501 & 20.9960 & ...$^{\\rm{e}}$ & Red & ...$^{\\rm{d}}$ \\\\[4pt] J182121+682715 & 23 & 0.8092 & 275.3361944 & 68.4540821 & ...$^{\\rm{d}}$ & ...$^{\\rm{d}}$ & 24.1828 & 23.3243 & 23.0264 & ...$^{\\rm{e}}$ & Blue & ...$^{\\rm{d}}$ \\\\[4pt] J182123+682714 & 24 & 0.8093 & 275.3460074 & 68.4539209 & ...$^{\\rm{d}}$ & ...$^{\\rm{d}}$ & 24.4237 & 23.7167 & 23.3609 & ...$^{\\rm{e}}$ & Red & ...$^{\\rm{d}}$ \\\\[6pt]\\hline \\end{tabular} \\end{center} \\end{sidewaystable*} \\begin{figure*}[p] \\plotone{NIRSPECspectralmosaic1.lowres.ps} \\caption{ACS \\emph{F}814\\emph{W} postage stamps of each galaxy targeted with NIRSPEC as well as the associated rest-frame DEIMOS (left) and NIRSPEC (right) spectra of each galaxy. The galaxy number is indicated next to each spectrum. The DEIMOS spectra are smoothed with a Gaussian of FWHM 15 pixels (roughly 2.7 \\AA\\ rest-frame at the redshift of the supercluster) and the NIRSPEC spectra are smoothed with a Gaussian of FWHM 1.7 pixels (roughly 2.7 \\AA\\ rest-frame). The error spectrum is plotted with a dashed line below each DEIMOS and NIRSPEC spectrum. All spectra are flux calibrated; however, no correction is made for internal extinction. For clarity the DEIMOS spectra are plotted with the zero flux level at the bottom of the plot. Due to the low level of continuum emission in some of the NIRSPEC targets, a dotted line shows the zero flux level for each NIRSPEC spectrum. The long dashed lines show the locations of important spectral features. NIRSPEC targets 0-5.} \\label{fig:mosaic1} \\end{figure*} \\begin{figure*}[p] \\plotone{NIRSPECspectralmosaic2.lowres.ps} \\caption{Same as in Figure \\ref{fig:mosaic1}. NIRSPEC targets 6-11.} \\label{fig:mosaic2} \\end{figure*} \\begin{figure*}[p] \\plotone{NIRSPECspectralmosaic3.lowres.ps} \\caption{Same as in Figure \\ref{fig:mosaic1}. NIRSPEC targets 12-17. The spectrum for galaxy 16 obtained with LRIS (center panel) is not flux calibrated. The spectrum is smoothed with a Gaussian with a 2 pixel FWHM (roughly 3.4 \\AA\\ rest-frame).} \\label{fig:mosaic3} \\end{figure*} \\begin{figure*}[p] \\plotone{NIRSPECspectralmosaic4.lowres.ps} \\caption{Same as in Figure \\ref{fig:mosaic1}. NIRSPEC targets 18-23. LFC $i\\arcmin$ postage stamps are used for galaxies 21-23 due to the lack of ACS data in the RX J1821 field.} \\label{fig:mosaic4} \\end{figure*} \\begin{figure*}[p] \\plotone{NIRSPECspectralmosaic5.lowres.ps} \\caption{Same as in Figure \\ref{fig:mosaic1}. NIRSPEC target 24. A LFC $i\\arcmin$ postage stamp is used for galaxy 24 due to the lack of ACS data in the RX J1821 field.} \\label{fig:mosaic5} \\end{figure*} \\clearpage" }, "1003/1003.1277_arXiv.txt": { "abstract": "We compute the expected luminosity function of GRBs in the context of the internal shock model. We assume that GRB central engines generate relativistic outflows characterized by the respective distributions of injected kinetic power $\\dot E$ and contrast in Lorentz factor $\\kappa=\\Gamma_{\\rm max}/\\Gamma_{\\rm min}$. We find that if the distribution of contrast extends down to values close to unity (i.e. if both highly variable and smooth outflows can exist) the luminosity function has two branches. At high luminosity it follows the distribution of $\\dot E$ while at low luminosity it is close to a power law of slope $-0.5$. We then examine if existing data can constrain the luminosity function. Using the $\\log{N}-\\log{P}$ curve, the $E_{\\rm p}$ distribution of bright BATSE bursts and the XRF/GRB ratio obtained by HETE2 we show that single and broken power-laws can provide equally good fits of these data. Present observations are therefore unable to favor one form of the other. However when a broken power-law is adopted they clearly indicate a low luminosity slope $\\simeq -0.6\\pm 0.2$, compatible with the prediction of the internal shock model. ", "introduction": "The isotropic luminosity of long gamma-ray bursts is known to cover a wide range from underluminous, nearby bursts such as GRB 980425 or GRB 060218 (with $L\\la 10^{47}$ erg.s$^{-1}$) to ultrabright objects like GRB 990123 ($L\\ga 10^{53}$ erg.s$^{-1}$). While it has been suggested that the weakest GRBs could simply be normal events seen off-axis \\citep{yamazaki:03}, this possibility has been recently discarded both from limits on afterglow brightness and for statistical reasons \\citep{soderberg:04,daigne:07}. The difference of six orders of magnitude between the brightest and weakest GRBs is therefore probably real. The parameters (stellar rotation, metallicity, etc.) which are responsible for this diversity in radiated power are not known. However, in the restricted range $10^{51}\\la L\\la 10^{53}$ erg.s$^{-1}$ the value of the isotropic luminosity is possibly fixed by the opening angle of the jet which may always carry the same characteristic energy \\citep{frail:01}. The purpose of this paper is to see how basic theoretical ideas and existing data can be used to constrain the GRB luminosity function (hereafter LF) $p(L)$. First, we should insist that $p(L)$ here represents the ``apparent\" LF which includes viewing angle effects and beaming statistics (i.e. bursts with narrow jets are more likely seen off-axis and therefore under-represented in the distribution). It is therefore different from the ``intrinsic\" LF $p_0(L)$ which would be obtained with all GRBs seen on-axis. In the lack of a complete, volume limited sample of GRBs with known redshift, only indirect observational indicators such as the $\\log{N}-\\log{P}$ plot can constrain the LF. These indicators however depend not only on $p(L)$ but also on the GRB rate and spectral energy distribution. The simplest possible form for $p(L)$ is a single power law $p(L)\\propto L^{-\\delta}$ between $L_{\\rm min}$ and $L_{\\rm max}$. Together with the parameters describing the GRB rate and spectral shape, $\\delta$, $L_{\\rm min}$ and $L_{\\rm max}$ can be adjusted to provide the best possible fit of the available indicators. Considering the mixing of the LF with other quantities in the fitting process it is remarkable that studies using different observational constraints have converged to a similar value of the slope $\\delta\\sim 1.5$ - $1.7$ \\citep[e.g.][]{firmani:04,daigne:07}. In a second step one can consider the more general case of a broken power law LF with now five parameters: $L_{\\rm min}$, $L_{\\rm max}$, the two slopes $\\delta_1$ and $\\delta_2$ and the break luminosity $L_{\\rm b}$. We will see in Sect.2 that there is some indication that the internal shock model of GRBs can produce a broken power law LF and we want to check if it is also favored by the existing observational data. As in our previous study we have used a Monte Carlo method to generate a large number of synthetic events where the parameters defining the burst properties are varied within fixed intervals. Preferred values of the parameters are those which yield the minimum $\\chi^2$ for a given set of observational constraints. We summarize these constraints and present the Monte Carlo simulations in Sect.3. We discuss our results in Sect.4 and Sect.5 is our conclusion. ", "conclusions": "We have demonstrated that in the framework of the internal shock model, a two branch LF is naturally expected, with a predicted low-luminosity branch which is a power-law of slope close to $ -0.5$. This result is robust as long as the central engine responsible for GRBs is capable to produce a broad diversity of outflows, from highly variable to very smooth.\\\\ Using a set of Monte Carlo simulations, we have then shown that current observations ($\\log{N}-\\log{P}$ diagram, peak energy distribution, fraction of XRRs and XRFs) are compatible with a broken power-law LF but still do not exclude a single power-law distribution. The low-luminosity slope of the broken power-law is strongly constrained to be $\\delta_{1}\\simeq 0.4-0.9$, compatible with the prediction of the internal shock model.\\\\ These results are encouraging but only preliminary. A better determination of the GRB LF would provide an interesting test of the internal shock model when the low luminosity branch becomes more easily accessible. This will however require the difficult task of detecting many bursts at the threshold of current instruments and measuring their redshift and spectral properties. \\\\" }, "1003/1003.3791_arXiv.txt": { "abstract": "We present {\\sl Spitzer} IRS mid-infrared (5--35 $\\mu$m) spectra of a complete flux-limited sample ($\\geq$ 3 mJy at 8 $\\mu$m) of young stellar object (YSO) candidates selected on the basis of their infrared colors in the Serpens Molecular Cloud. Spectra of 147 sources are presented and classified. Background stars (with slope consistent with a reddened stellar spectrum and silicate features in absorption), galaxies (with redshifted PAH features) and a planetary nebula (with high ionization lines) amount to 22\\% of contamination in this sample, leaving 115 true YSOs. Sources with rising spectra and ice absorption features, classified as embedded Stage I protostars, amount to 18\\% of the sample. The remaining 82\\% (94) of the disk sources are analyzed in terms of spectral energy distribution shapes, PAHs and silicate features. The presence, strength and shape of these silicate features are used to infer disk properties for these systems. About 8\\% of the disks have 30/13 $\\mu$m flux ratios consistent with cold disks with inner holes or gaps, and 3\\% of the disks show PAH emission. Comparison with models indicates that dust grains in the surface of these disks have sizes of at least a few $\\mu$m. The 20 $\\mu$m silicate feature is sometimes seen in absence of the 10 $\\mu$m feature, which may be indicative of very small holes in these disks. No significant difference is found in the distribution of silicate feature shapes and strengths between sources in clusters and in the field. Moreover, the results in Serpens are compared with other well-studied samples: the c2d IRS sample distributed over 5 clouds and a large sample of disks in the Taurus star-forming region. The remarkably similar distributions of silicate feature characteristics in samples with different environment and median ages -- if significant -- imply that the dust population in the disk surface results from an equilibrium between dust growth and destructive collision processes that are maintained over a few million years for any YSO population irrespective of environment. ", "introduction": "\\label{sintro} Newly formed stars are observed to have infrared (IR) excess due to their circumstellar disk composed of dust and gas (\\citealt{SS92,HI08}). Most older main-sequence (MS) stars, on the other hand, have photospheric emission with no excess in the IR. It is intuitive to conclude that the circumstellar disk evolves with time, gradually getting rid of the IR excess. One of the main questions in stellar astrophysics is how this happens. Observational studies, as well as theoretical simulations, have demonstrated the interaction between star and disk. The stellar radiation facilitates disk evolution in terms of photoevaporation (e.g., \\citealt{RI00,RA06,RA08,GH09}) or dust growth and settling (e.g., \\citealt{WE80,ST95,DT97,DD05,JO08}). In the other direction, mass is accreted from the disk to the star following magnetic field lines (e.g., \\citealt{MU03,WB03,NA04}). A diversity of stellar temperatures, luminosities and masses among young stars has been known and studied for decades. Facilitated by new IR and (sub-)millimeter observations, a great variety of disk shapes, structures and masses is now being actively studied. The next step is to try to connect stellar and disk characteristics in order to understand the evolution of these systems. The study of a single object, however, is unlikely to provide unambiguous information regarding the evolutionary stage of the associated disk. Most studies to date refer to samples of young stars scattered across the sky, or to sources distributed across large star-forming clouds like Taurus. In addition to evolutionary stage, the specific environment in which the stars are formed may influence the evolution of disks by dynamical and radiative interaction with other stars or through the initial conditions of the starting cloud, making it difficult to separate the evolutionary effects (e.g., \\citealt{RI98,RI00}). For this reason, clusters of stars are very often used as laboratories for calibrating the evolutionary sequence (e.g., \\citealt{LL95,HA01}). The power of this method, to gain statistical information on disk composition in coeval samples, was found to be very successful for loose associations of older, pre-main sequence stars such as the 8 Myr old $\\eta$ Cha \\citep{BO06} and the 10 Myr old TW Hydrae association \\citep{UC04}. Identifying clusters of even younger disk populations is a natural step towards the completion of the empirical calibration of the evolution of disks surrounding young low-mass stars. This paper analyzes the inner disk properties of a flux-limited, complete unbiased sample of young stars with IR excess in the Serpens Molecular Cloud (d = 259 $\\pm$ 37 pc, \\citealt{ST96}) which has a mean age $\\sim 5$ Myr, \\citep{OL09} with an YSO population in clusters and also in isolation. It has been recently argued (L. Loinard, private communication) that the distance to Serpens could be considerably higher than previously calculated. This would imply a rather younger median age for this cloud. The Spitzer Legacy Program ``From Molecular Cores to Planet-Forming Disks'' (c2d) has uncovered hundreds of objects with IR excess in five star-forming clouds (Cha II, Lupus, Ophiuchus, Perseus and Serpens), and allowed statistical studies within a given cloud \\citep{EV09}. The c2d study of Serpens with IRAC (3.6, 4.5, 5.8 and 8.0 $\\mu$m, \\citealt{FA04}) and MIPS (24 and 70 $\\mu$m, \\citealt{RI04}) data has revealed a rich population of mostly previously unknown young stellar objects (YSOs) associated with IR excess, yielding a diversity of disk SEDs (\\citealt{HA06,HB07,HA07}). Because of the compact area in Serpens mapped by Spitzer (0.89 deg$^2$), this impressive diversity of disks presents itself as an excellent laboratory for studies of early stellar evolution and planet formation. Indeed, the Serpens core (Cluster A, located in the northeastern part of the area studied by c2d) has been well studied in this sense (e.g., \\citealt{ZA88,EC92,TS98,KA04,EI05,WI07,WI09}), whereas only some of the objects in Group C (formerly known as Cluster C) were studied with ISOCAM data \\citep{DJ06}. Because of its wavelength coverage, sensitivity and mapping capabilities, the Spitzer Space Telescope has offered an opportunity to study many of these systems (star$+$disk) in unprecedented detail. Spitzer's photometry in the mid-IR, where the radiation reprocessed by the dust is dominant, gives information on the shape of the disks and, indirectly, its evolutionary stage (assuming an evolution from flared to flat disks). Follow-up mid-IR spectroscopic observations with the InfraRed Spectrograph (IRS, 5 -- 38, $\\mu$m \\citealt{HO04}) on-board Spitzer probe the physical and chemical processes affecting the hot dust in the surface layers of the inner regions of the disk. The shapes and strengths of the silicate features provide information on dust grain size distribution and structure (e.g., \\citealt{VB03,PY03,BO08,KE06,KE07,GE06,WA09,OF09}). These, in turn, reflect dynamical processes such as radial and vertical mixing, and physical processes such as annealing. A smooth strong Si--O stretching mode feature centered at 9.8 $\\mu$m is indicative of small amorphous silicates (like those found in the ISM) while a structured weaker and broader feature reveals bigger grains or the presence of crystalline silicates. Polycyclic aromatic hydrocarbon (PAH) features are a probe of the UV radiation incident on the disk, whereas their abundance plays a crucial role in models of disk heating and chemistry (e.g., \\citealt{GE06,DU07,VI07}). The shape and slope of the mid-infrared excess provides information on the flaring geometry of the disks \\citep{DD04}, while ice bands may form for highly inclined sources (edge-on) where the light from the central object passes through the dusty material in the outer parts of the disk \\citep{PO05}. Thus, the wavelength range probed by the IRS spectra enables analysis of the geometry of individual disks. It also probes the temperature and dust size distributions as well as crystallinity of dust in the disk surface at radii of 0.1 -- few AU. Statistical results from a number of sources help the understanding of the progression of disk clearing and possibly planet formation. Our group has been conducting multi-wavelength observing campaigns of Serpens. Optical and near-IR wavelength data, where the stellar radiation dominates, are being used to characterize the central sources of these systems \\citep{OL09}. Effective temperatures, luminosities, extinctions, mass accretion rates, as well as relative ages and masses are being determined. In this paper, we present a complete flux-limited set of Spitzer IRS spectra for this previously unexplored young stellar population in Serpens. We analyze these spectra in terms of common and individual characteristics and compare the results to those of Taurus, one of the best studied molecular clouds to date and dominated by isolated star formation. A subsequent paper will deal with the full SED fitting for the disk sources in Serpens and their detailed analyzes. Our ultimate goal is to statistically trace the evolution of young low-mass stars by means of the spectroscopic signatures of disk evolution, discussed above. Section \\ref{sdata} describes our Spitzer IRS data. In \\S~\\ref{sres} we divide our sample into categories based on their IRS spectra: the background contaminants are presented in \\S~\\ref{sbg} separated according to the nature of the objects; embedded sources are presented in \\S~\\ref{semb}, and disk sources in \\S~\\ref{sdisks} (with emphasis on PAH and silicate emission). In \\S~\\ref{sdis} we discuss disk properties in relation to environment and to another cloud, Taurus, for comparison. In \\S~\\ref{scon} we present our conclusions. ", "conclusions": "\\label{scon} We present Spitzer/IRS spectra from a complete and unbiased flux limited sample of IR excess sources found in the Serpens Molecular Cloud, following the c2d mapping of this region. \\begin{itemize} \\item Among our total of 147 IRS spectra, 22\\% are found to be background contamination (including background stars, redshifted galaxies and a planetary nebula candidate). This high number is not surprising given the position of Serpens, close to the galactic plane. \\item Excluding the background objects from the sample, the bona fide set of YSOs amounts to 115 objects. The embedded to disk source ratio is 18\\%, in agreement with the ratios derived from photometry \\citep{EV09} for the five c2d clouds. \\item Disks with PAH in emission amount to 2\\%\\ of the YSO population, but 3\\%\\ of the disk population. Only G and A star show PAH emission. In the disk population, 73\\%\\ show both silicate emission features, at 10 and 20 $\\mu$m, while 17\\%\\ show only the 20 $\\mu$m feature. 4\\%\\ of the disks show featureless mid-IR spectra. Only one source, \\#120, shows both PAH and silicate in emission. \\item Our YSO population in Serpens is very similar to that in Taurus, also studied with IRS spectra. In both regions about 70\\%\\ of the disk sources present both 10 and 20 $\\mu$m silicate features in emission and the 10 $\\mu$m feature is never seen without the 20 $\\mu$m feature. The only significant difference between the two populations is in the number of sources lacking the 10 $\\mu$m feature but showing the 20 $\\mu$m feature. This may be indicative of small holes ($\\lesssim$ 1 AU) in these sources. \\item The silicate features in the IRS spectra measure the grain sizes that dominate the mid-IR opacity. The relationships between shape and strength of these features present distributions very similar to those obtained from other large samples of young stars. Comparison with the models of \\citet{OF09} yield grains consistent with sizes larger than a few $\\mu$m. \\item No significant differences are found in the disk geometry or the grain size distribution in the upper layers of circumstellar disks (probed by the silicate features) in clustered or field stars in the cloud, indicating that the local environment where a star is born does not influence the evolution of its harboring disk. \\item Quantitatively, the shape and strength of the 10 $\\mu$m silicate feature were used to compare both Serpens and Taurus as well as a large sample of disks across five clouds, indicating remarkably similar populations. This implies that the dust population in the disk surface results from an equilibrium between dust growth and destruction processes that are maintained over a few million years. \\end{itemize}" }, "1003/1003.1870_arXiv.txt": { "abstract": "The dynamics of a tachyon field plus a barotropic fluid is investigated in spatially curved FRW universe. We perform a phase-plane analysis and obtain scaling solutions accompanying with a discussion on their stability. Furthermore, we construct the form of scalar potential which may give rise to stable solutions for spatially open and closed universe separately. ", "introduction": "The nature and the origin of dark energy is a fundamental puzzle in modern cosmology. Most dark energy models can be constructed by using a slowly rolling canonical scalar field, termed as quintessence. However, there has been increasing interests in alternative models with a non-canonical kinetic term. Among these models the most general formalism perhaps is k-essence \\cite{k-essence}. A more specific case is the tachyon field \\cite{tachyon}, which is motivated from string theory. It can be viewed as a special case of k-essence with Dirac-Born-Infeld (DBI) action \\cite{DBI}. Although the tachyon is an unstable field, its state parameter in the equation of state varies smoothly between $-1$ and $0$, thus many authors have already considered the tachyon field as a suitable candidate for a viable model of dark energy phenomenologically \\cite{tachyondeTP,tachyondeFP,tachyondeEJC,tachyondeBGTNJW,tachyondeST,tachyondeYG,tachyondeJM,tachyondeAAS,tachyondeBGJW,tachyondeGC}. For a review, we can refer to Ref.\\cite{DERE}. However, the dynamical dark energy models driven by a scalar field suffer from the so-called fine-tuning problem and coincidence problem. In order to address these problems, one may employ scalar field models exhibiting scaling solutions \\cite{SS1,SS2,SS31,SS32,SS33,SS34,SS35,SS36,SS37,SS38,SS39,SS310,SS311,SS312,SS313,SS314,SS315,tachyondeYG,SS4,tachyondeST,SS6,SS7}. The scaling solutions as dynamical attractors can considerably alleviate these two problems. Furthermore, by investigating the nature of scaling solutions, one can determine whether such behavior is stable or just a transient feature and explore the asymptotic behavior of the scalar field potential. Many authors have investigated a lot of scalar field models containing scaling solutions. For instance, a canonical scalar field with an exponential potential has scaling attractor solutions \\cite{SS2}. For quintessence dark energy model, there are two scaling solutions. One is fluid-scalar field scaling solution, which remains subdominant for most of the cosmic evolution. It is necessary that the scalar field mimics the background energy density (radiation/matter) in order to respect the nucleosynthesis constraint and can also alleviate the fine-tuning problem of initial conditions. The other is scalar field dominated scaling solution, which is a late time attractor and gives rise to the accelerated expansion. Since the fluid-scalar field scaling solution is non-accelerating, we need an additional mechanism exit from the scaling regime so as to enter the scalar field dominated scaling solution at late times. For the discussion on the exiting mechanism, we can refer to Refs. \\cite{DERE,exitingTB,exitingVS,exitingAA,exitingAAC,exitingSAK}. For tachyon field dark energy, the scaling solutions have also been investigated by many authors, for example Refs.\\cite{tachyondeTP,tachyondeFP,tachyondeEJC,tachyondeBGTNJW,tachyondeST,tachyondeYG,tachyondeJM,tachyondeAAS,tachyondeBGJW,tachyondeGC}. To be considered as a realistic model of dark energy, it is found that the fluid-scalar field scaling solutions are absent and only the scalar field dominated scaling solutions exist \\cite{tachyondeEJC}. This is very different from the quintessence case. Therefore, just as pointed out in \\cite{tachyondeGC}, tachyon models require more fine-tuning to be consistent with observations. Nevertheless, it is worthwhile to point out that the fluid-scalar field scaling solutions may be obtained when the Gauss-Bonnet coupling between tachyon field and fluid is considered in the tachyon dark energy model\\cite{SS7}. Although the latest results of WMAP5 have placed a constraint, $-0.063<\\Omega_{k}<0.017$ \\cite{WMAP5}, on the flatness of our observable universe, indicating that our observable universe is very close to flatness, it is still possible that our observable universe is spatially curved. Therefore, it is also interesting to investigate the dynamical behavior of dark energy models in a spatially curved FRW universe. Recently Copeland {\\it et.al.} have extended such investigations to the quintessence model in spatially curved FRW universe\\cite{CSQ}, following the strategy they have developed in Ref.\\cite{SS312}. Sen and Devi \\cite{tachyondeAAS} have also explored the scaling solutions with tachyon in modified gravity model employing the same method. In this Letter, we will closely follow this route to investigate the dynamics of tachyon dark energy model in spatially curved universe. Our Letter is organized as follows. In Section II we present the associated equations of motion for the tachyon field including the background fluid and obtain the scaling solutions. Then, we analyze the stability of these solutions. In Section III we turn to construct the scalar potential leading to such scaling solutions. In particular, its asymptotical forms are obtained in various circumstances for spatially open and closed universe respectively. ", "conclusions": "In this Letter we have investigated the dynamics of a tachyon field in FRW universe with spatial curvature. Following the scheme presented in \\cite{SS312,CSQ}, we denoted the modification of the cosmological equation due to the spatial curvature by a general function $L(\\rho)$, and then derived the conditions under which the system can enter a scaling solution. In particular, we obtain an attractor solution, where the tachyon field dominates over the fluid and the kinetic energy of the field scales with its potential energy. Furthermore, given the modification function for an open and closed universe respectively, we discussed the form of the scalar potential which gives rise to the late time attractor solution. For spatially open universe, we conclude that in regions where the curvature is negligible, the asymptotic form of the potential is $V\\propto\\phi^{-2}$, while in regions where it is dominant, the approximated potential will be $V \\propto\\phi^{-\\frac{2}{\\mu+1}}$. As far as the spatially closed universe is concerned, we find that the curvature is not allowed to be dominant among the total energy density of the universe, but only subdominant to the scalar field density, where the potential is a constant as $V\\approx Y_{c}^{2}(\\frac{A^{2}}{3})^{1/\\mu}$. When the curvature becomes negligible, the potential can be approximated by $ V \\sim \\phi^{-2} $. Moreover, it was noticed that in a contracting closed universe, the scalar dominated solutions are not stable, and the kinetic dominated solution will be a late time attractor if it exists. Comparing with the quintessence scalar field dark energy model with spatial curvature \\cite{CSQ}, the fluid-scalar field scaling solutions is absent in the tachyon dark energy model. In addition, in the quintessence scalar field dark energy model, there are the two forms of the potential: the exponential potential and the power-law potential. But in the tachyon dark energy model, only the power-law potential is required. In this Letter we only consider the case that the scalar field and the fluid do not interact with each other. It is a very interesting question if such an interaction term can be introduced in order to obtain a fluid-tachyon field scaling solutions in curved FRW universe. Effort has been made in present framework, unfortunately, the answer seems to be negative. However, it should be very desirable if we consider adding a Gauss-Bonnet coupling as in \\cite{SS7}. We expect to make further progress along this direction." }, "1003/1003.4208_arXiv.txt": { "abstract": "This deep, extended solar minimum and the slow start to Cycle 24 strongly suggest that Cycle 24 will be a small cycle. A wide array of solar cycle prediction techniques have been applied to predicting the amplitude of Cycle 24 with widely different results. Current conditions and new observations indicate that some highly regarded techniques now appear to have doubtful utility. Geomagnetic precursors have been reliable in the past and can be tested with 12 cycles of data. Of the three primary geomagnetic precursors only one (the minimum level of geomagnetic activity) suggests a small cycle. The Sun's polar field strength has also been used to successfully predict the last three cycles. The current weak polar fields are indicative of a small cycle. For the first time, dynamo models have been used to predict the size of a solar cycle but with opposite predictions depending on the model and the data assimilation. However, new measurements of the surface meridional flow indicate that the flow was substantially faster on the approach to Cycle 24 minimum than at Cycle 23 minimum. In both dynamo predictions a faster meridional flow should have given a shorter cycle 23 with stronger polar fields. This suggests that these dynamo models are not yet ready for solar cycle prediction. ", "introduction": "As each sunspot cycle wanes solar astronomers with widely different interests take their turn at predicting the size and timing of the next cycle. The average length of the previous 22 sunspot cycles is 131.7 months - almost exactly 11 years. However, with one exception, the last 8 cycles have been short cycles with periods closer to 10 years. The minimum preceeding Cycle 23 was in August or September of 1996 so many were expecting the minimum preceeding Cycle 24 to come in 2007 or even 2006. Instead, minimum came in November of 2008 (Fig. 1). This delayed start of Cycle 24, and the depth of the minimum (smoothed sunspot number at its lowest in nearly 100 years) stirred up additional interest and even more predictions \\citep{Pesnell08} including talk of an impending grand minimum like the Maunder Minimum \\citep[e.g.]{Schatten03}. \\begin{figure}[!ht] \\plotone{fig1.eps} \\caption{ Sunspots associated with Cycle 24 (black) began to dominate over those associated with Cycle 23 (white)in September of 2008. The smoothed sunspot number went through its minimum in December 2008. The smoothed number of spotless days per month went through its maximum in December 2008 as well. The average of these three traditional indicators of sunspot cycle minimum gives November of 2008 as Cycle 24 minimum.} \\end{figure} ", "conclusions": "Ohl's Geomagnetic Precursor, the Polar Field Precursors, the Amplitude-Period relation, and the Maximum-Minimum relation all indicate that Cycle 24 will be small with an amplitude of about 75. The other two geomagnetic precursor methods appear to be unduly impacted by the activity associated with the Halloween events of 2003 and give larger cycles. We conclude with \\citet{Wang09} that the more appropriate geomagnetic precursor is that of \\citet{Ohl66} - the minimum level of geomagnetic activity. The predictions based on Flux Transport Dynamos gave very different predictions but they both predict behavior in conflict with the observed meridional flow variations. The faster meridional flow after Cycle 23 maximum sould give a short cycle with strong polar fields acording to these models. Instead we find a long cycle with weak polar fields. We must conclude with \\citet{Tobias06} that these dynamo modes are not yet ready for cycle predictions." }, "1003/1003.3933_arXiv.txt": { "abstract": "Two-dimensional electrostatic turbulence in magnetized weakly-collisional plasmas exhibits a cascade of entropy in phase space [\\prl \\textbf{103}, 015003 (2009)]. At scales smaller than the gyroradius, this cascade is characterized by the dimensionless ratio $D$ of the collision time to the eddy turnover time measured at the scale of the thermal Larmor radius. When $D \\gg 1$, a broad spectrum of fluctuations at sub-Larmor scales is found in both position and velocity space. The distribution function develops structure as a function of $\\vperp$, the velocity coordinate perpendicular to the local magnetic field. The cascade shows a local-scale nonlinear interaction in both position and velocity spaces, and Kolmogorov's scaling theory can be extended into phase space. ", "introduction": "Fluid turbulence can be described mostly by Navier-Stokes equations, which consist of partial differential equations in the position space --- thus in 2D or 3D space. Such a description is possible because we may regard the fluid to be sufficiently collisional and to be locally in thermodynamic equilibrium. From the outer driving scales to the small dissipative scales, the energy (or the enstrophy in the 2D case) cascades in the wave-number space \\cite{Kolmogorov}. The {\\em inertial range} dynamics are basically governed by the nonlinear term. The fluxes are finally dissipated by the viscosity described by a diffusion operator, in the {\\em dissipation range} of wave-number space. There is a dimensionless number called the Reynolds number, which characterizes the scale separation of the turbulent dynamics. The ratio of the dissipation scale to the outer scale is in general determined by a fractional power of the Reynolds number. Kinetic turbulence is less well understood. We have some knowledge especially from the extensive study of fusion plasmas \\cite{WatanabeSugama-PoP04,Idomura-PoP06,Candy-PPCF07} and recently of space plasmas \\cite{Bale-PRL05,Howes-PRL08,GarySaitoLi,Sahraoui-PRL09,Alexandrova-PRL09}, but not to the same extent as fluid turbulence is known, particularly in the sense described above. Since most plasmas are in a collisionless or weakly-collisional environment, they are not in local thermodynamical equilibrium. In this case we have to invoke kinetic description in the phase space, which enables us to use Vlasov or Boltzmann theory, or their magnetized low-frequency limit, the gyrokinetics. Dissipation in the kinetic system occurs because of collisions. Without collisions, the entropy is conserved and the system is in principle reversible in a thermodynamic sense \\cite{Boltzmann}. Wave resonances or Landau damping do not by themselves constitute dissipation since no entropy is generated unless we include the effect of collisions. Irreversibility becomes the key issue here. In this context, the following question arises. Collisions are usually described by a diffusion operator in the \\textit{velocity space} \\cite{Helander}. How do collisions in a weakly collisional system determine the wave-number as well as velocity-space cutoff of the inertial range? In this paper, following Ref.\\ \\cite{Tatsuno-PRL09}, we show detailed numerical evidence of the coupling between position- and velocity-space structures, and the consequent achievement of the collisional dissipation described by the velocity-space diffusion operator. We introduce a new dimensionless number $D$, which, analogously to the Reynolds number in Navier-Stokes equations, characterizes the scale separation between the outer scale and the dissipation scale in the kinetic theory. We emphasize that we are concerned only with scales smaller than the Larmor radius. In order to focus on the nonlinear interaction, we omit Landau damping by ignoring variation along the field line. The system retains two collisionless invariants as the 2D Navier-Stokes equations do. These two invariants cannot share the same local-interaction space in a Kolmogorov-like phenomenology, and thus lead to a dual cascade (direct and inverse cascades) \\cite{Idomura-PoP06,Fjortoft-Tellus53,Kraichnan-PoF67,HasegawaMima-PoF78,Horton-PoP00}. In this paper we focus on the direct cascade only and the discussion of the inverse cascade is left for future publication. This paper is organized as follows. As a reduced kinetic model for plasma turbulence, we first introduce the gyrokinetic (GK) equation briefly and describe its basic nature in \\sect{sec:equations}. Scaling relations of the turbulent cascade are also briefly summarized. Section \\ref{sec:direct} shows the evidence of kinetic turbulent cascade by means of the numerical simulation of the decaying turbulence. In addition to the reproduction of the theoretical prediction, we show how much resolution is needed in both position and velocity space to achieve the proper dissipation and show the trend that the amount of dissipation is asymptotically independent of the collision frequency. We also present the characteristics of the nonlinear triad interaction in wave-number space and show that the assumption of local-scale interaction is supported. We finally summarize our results in \\sect{sec:summary}. ", "conclusions": "\\label{sec:summary} We presented electrostatic, decaying turbulence simulations for weakly collisional, magnetized plasmas using the gyrokinetic model in 4D phase space (two position-space and two velocity-space dimensions; the extension to three spatial dimensions is left for future work). Landau damping was removed from the system by ignoring variation along the background magnetic field. Nonlinear interactions introduce an amplitude-dependent perpendicular phase mixing of the gyrophase-independent part of the perturbed distribution function and create structure in $v_\\perp$ which is finer for higher $\\kperp$ (see \\fig{fig:2d-spectrum}). We found that the wave-number (Fourier) spectra of the perturbed distribution function and the resulting electrostatic fluctuations at sub-Larmor scales agreed well with theoretical predictions based on the interpretation of the nonlinear phase mixing as a cascade of entropy in phase space (see Figs.\\ \\ref{fig:k-spectra} and \\ref{fig:asymptotic spectra}) \\cite{Tatsuno-PRL09,Alex,Plunk-JFM}. The velocity-space (Hankel) spectra show a rough consistency with the theoretical scaling, although the agreement is not as good as that of the wave-number spectra (see Figs.\\ \\ref{fig:v-spectra} and \\ref{fig:2d-spectrum}). We introduced a dimensionless number (analogous to Reynolds number) that characterized the scale separation between the thermal Larmor scale and the collisional cutoff in phase space [see \\eqref{dimensionless number}], and showed that this number correctly predicted the resolution requirements for nonlinear gyrokinetic simulations (see Table \\ref{run table}). We also showed the trend that the entropy generation rate (or irreversibility) is independent of the collision frequency in the asymptotic limit of the weak collisionality (see \\fig{fig:invariants}). Finally we presented diagnostics of nonlinear transfer functions and showed that local-scale cascade of entropy is supported very well in both Fourier and Hankel spaces. We note that there are, in general, entropy cascades for each plasma species. Equations for the gyrokinetic turbulence at and below the electron Larmor scale are mathematically similar to the model simulated here and identical arguments apply \\cite{Alex,Plunk-JFM}. Similar considerations are also possible for ion-scale electromagnetic turbulence \\cite{Alex} and for minority ion species (with some differences to be discussed elsewhere). The structure of the small scales in phase space that we have presented is likely to be a universal feature of magnetized plasma turbulence. Understanding it theoretically and diagnosing it numerically is akin to the inertial-range studies for Kolmogorov turbulence, extended to the kinetic phase space. One should expect rich and interesting physics to emerge and it is likely that, just like in the case of fluid turbulence, predicting large-scale dynamics will require effective models for the small-scale cascade. An immediate physical implication of the existence of the entropy cascade is a turbulent heating rate independent of collisionality in weakly collisional plasmas." }, "1003/1003.2935_arXiv.txt": { "abstract": "{} {Our aim is to observationally investigate the cosmic Dark Ages in order to constrain star and structure formation models, as well as the chemical evolution in the early Universe. } {{ Spectral lines from atoms and molecules in primordial perturbations at high redshifts can give information about the conditions in the early universe before and during the formation of the first stars in addition to the epoch of reionisation. The lines may arise from moving primordial perturbations before the formation of the first stars (resonant scattering lines), or could be thermal absorption or emission lines at lower redshifts.} The difficulties in these searches are that the source redshift and evolutionary state, as well as molecular species and transition are unknown, which implies that an observed line can fall within a wide range of frequencies. The lines are also expected to be very weak. Observations from space have the advantages of stability and the lack of atmospheric features which is important in such observations. We have therefore, as a first step in our searches, used the Odin\\thanks{Odin is a Swedish-led satellite project funded jointly by the Swedish National Space Board (SNSB), the Canadian Space Agency (CSA), the National Technology Agency of Finland (Tekes) and Centre National d'Etudes Spatiales (CNES). The Swedish Space Corporation was the prime contractor and also is responsible for the satellite operation.} satellite to perform two sets of spectral line surveys towards several positions. The first survey covered the band 547\\,--\\,578\\,GHz towards two positions, and the second one covered the bands 542.0\\,--\\,547.5\\,GHz and 486.5\\,--\\,492.0\\,GHz towards six positions selected to test different sizes of the primordial clouds. Two deep searches centred at 543.250 and 543.100\\,GHz with 1\\,GHz bandwidth were also performed towards one position. The two lowest rotational transitions of H$_2$ will be redshifted to these frequencies from $z\\!\\sim\\!20\\!-\\!30$, which is the predicted epoch of the first star formation. } {No lines are detected at an rms level of 14\\,--\\,90 and 5\\,--\\,35\\,mK for the two surveys, respectively, and 2\\,--\\,7\\,mK in the deep searches with a channel spacing of 1\\,--\\,16\\,MHz. The broad bandwidth covered allows a wide range of redshifts to be explored for a number of atomic and molecular species and transitions. From the theoretical side, our sensitivity analysis show that the largest possible amplitudes of the resonant lines are about 1\\,mK at frequencies $\\lesssim$\\,200\\,GHz, and a few $\\mu$K around 500\\,--\\,600\\,GHz, assuming optically thick lines and no beam-dilution. However, if existing, thermal absorption lines have the potential to be orders of magnitude stronger than the resonant lines. We make a simple estimation of the sizes and masses of the primordial perturbations at their turn-around epochs, which previously has been identified as the most favourable epoch for a detection. This work may be considered as an important pilot study for our forthcoming observations with the Herschel Space Observatory. } {} ", "introduction": " ", "conclusions": "\\label{section summary} In order to constrain cosmological models of star and structure formation as well as the chemical evolution in the early Universe, we have performed spectral line surveys towards several positions without any known sources of emission in a search for primordial spectral lines from the Dark Ages. The first survey covered a broad band of 31\\,GHz between 547\\,--\\,578\\,GHz towards two positions with fixed reference-positions. The second survey covered 11\\,GHz in the bands 542.0\\,--\\,547.5\\,GHz and 486.5\\,--\\,492.0\\,GHz towards four positions. We also performed two deep searches towards one position with 543.100 and 543.250\\,GHz as centre frequencies. No lines were detected, and thus the results are upper limits in terms of noise level. Typical 1$\\sigma$ values are \\mbox{5\\,--\\,35\\,mK} in the 11\\,GHz survey, \\mbox{14\\,--\\,90\\,mK} in the 31\\,GHz survey, and 2\\,--\\,7\\,mK in the deep searches over a 1\\,GHz band. The major improvement made by the Odin observations compared to \\citet{1993A&A...269....1D} is the broad bandwidth covered allowing a wide range of redshifts to be explored for a number of atomic and molecular species. In addition, in the second survey we have taken into account the unknown sizes of the clouds by testing an observational strategy where we have observed towards four positions in a sequence with different angular distances between the reference and signal position. An important benefit of our observations is that we do no suffer from spectral line contamination from the terrestrial atmosphere. At low densities and in matter-radiation equilibrium conditions, the \\emph{only} expected signal is from resonant line scattering between CMB photons and matter moving with respect to the expansion of the Universe. These lines suffer, however, from the dependence of the low CMB temperature and the peculiar velocities of the moving primordial perturbations which is about 10$^{-3}$ at present and even lower at higher redshifts. In order to obtain an estimate of the \\emph{highest possible intensities} we assume optically thick resonant lines and no beam-dilution. The intensities will then be at most of the order of a few mK at frequencies below 100\\,--\\,200\\,GHz, and orders of magnitude lower at higher frequencies (5\\,$\\mu$K at 600\\,GHz). Since the lines most likely have $\\tau\\!\\ll\\!1$ this will further lower their intensities with orders of magnitudes. If existing, thermal absorption lines on the other hand have the potential to be observable at both high and low frequencies. Such lines do not suffer from the dependence of peculiar velocities of the moving primordial perturbations. The background radiation could also be considerably higher than the CMB, thereby producing stronger lines. The background could for instance be collapsing primordial perturbations, the first stars which are predicted to form at z\\,$\\sim$\\,20\\,--\\,30, or the remnants of pop~III supernovae. These objects probably emitted large amounts of energetic radiation which foreground primordial perturbations could absorb. Also here we find the \\emph{highest possible intensities} by assuming optically thick lines and no beam-dilution. Assuming that the CMB is the background radiation, we find that the largest possible intensity is 1\\,--\\,2\\,K below 100\\,GHz, and of the order of a few mK around 500\\,GHz. If the background radiation is higher than the CMB, the intensities will increase. Also in this case the opacity most likely is $\\ll\\!\\!1$ which will lower the absorption line intensities, in addition to beam-dilution. The strength and line-width of the primordial lines depend on the evolution of the primordial perturbations. In the first (linear) phase they are very broad and weak, but become increasingly stronger and more narrow during the evolution of the cloud. The turn-around phase, which produces the strongest and most narrow lines, and the beginning of the collapse have previously been identified as the most favourable evolutionary phases for observations \\citep{1996ApJ...457....1M}. We have made a simple estimation of the redshifts and sizes of the primordial perturbations at their turn-around phase, and find that the Odin beam size of 2$\\farcm$1 corresponds to a turn-around mass of about 6$\\times$10$^{12}$\\,M$_\\odot$ at a corresponding redshift of about 3 at a 1$\\sigma$ level. Smaller perturbations at higher redshift will thus suffer from beam-dilution. A beam-size of 40\\arcsec, which is approximately the beam-size of Herschel Space Observatory\\footnote{\\url{http://herschel.esac.esa.int/}}, corresponds to a 1$\\sigma$ turn-around mass and redshift of about 10$^{12}$\\,M$_\\odot$ and $z\\! \\sim\\!10$. The lowest rotational transitions of H$_2$ will fall in the Odin band around 500\\,GHz from z$\\sim$20\\,--\\,30, while the lowest transitions of other species such as HD, HD$^+$, and HeH$^+$ will fall in the band around 100\\,GHz from the same epoch. These transitions may be searched for using ground based antennas -- even though the foreground and atmospheric radiation will pose a problem. An important aspect of our work has been to test different observational strategies to prepare for our forthcoming observations with the much more sensitive telescope and receivers aboard the Herschel Space Observatory launched on May 14, 2009. The much lower noise level and the even broader band coverage with Herschel will increase the possibility of a detection. Other interesting facilities to consider in the searches for primordial molecules are for example the Atacama Large Millimeter Array\\footnote{\\url{http://www.alma.info/}} (ALMA), the Combined Array for Research in Millimeter-Wave Astronomy\\footnote{\\url{http://www.mmarray.org/}}~(CARMA), the IRAM Plateau de Bure Interferometer and the IRAM 30-m telescope \\footnote{\\url{http://www.iram.fr/}}, and the Very Large Array\\footnote{\\url{http://www.vla.nrao.edu/}} (VLA) at 20\\,--\\,45\\,GHz. Our observing methods and resulting limits, paired with a sensitivity analysis taking into account the evolution of primordial perturbations, should be a valuable input to the planning of these observations. Spectral lines from primordial atoms and molecules may very well be the only way to probe the epoch of the cosmic Dark Ages and its end when the first stars formed. The search for these primordial signals is a very difficult task indeed, but a detection could be possible with the use of future facilities and would introduce an additional important way to discriminate between models of the early universe as well as star and structure formation." }, "1003/1003.5724_arXiv.txt": { "abstract": "Wide field surveys will soon be discovering Type Ia supernovae (SNe) at rates of several thousand per year. Spectroscopic follow-up can only scratch the surface for such enormous samples, so these extensive data sets will only be useful to the extent that they can be characterized by the survey photometry alone. In a companion paper (Rodney and Tonry, 2009) we introduced the SOFT method for analyzing SNe using direct comparison to template light curves, and demonstrated its application for photometric SN classification. In this work we extend the SOFT method to derive estimates of redshift and luminosity distance for Type Ia SNe, using light curves from the SDSS and SNLS surveys as a validation set. Redshifts determined by SOFT using light curves alone are consistent with spectroscopic redshifts, showing a root-mean-square scatter in the residuals of $RMS_z=0.051$. SOFT can also derive simultaneous redshift and distance estimates, yielding results that are consistent with the currently favored $\\Lambda$CDM cosmological model. When SOFT is given spectroscopic information for SN classification and redshift priors, the RMS scatter in Hubble diagram residuals is 0.18 mags for the SDSS data and 0.28 mags for the SNLS objects. Without access to any spectroscopic information, and even without any redshift priors from host galaxy photometry, SOFT can still measure reliable redshifts and distances, with an increase in the Hubble residuals to 0.37 mags for the combined SDSS and SNLS data set. Using Monte Carlo simulations we predict that SOFT will be able to improve constraints on time-variable dark energy models by a factor of 2--3 with each new generation of large-scale SN surveys. ", "introduction": "\\label{sec:LightCurveModels} The shape of an observed SN light curve can be described by two sets of parameters. The first set, which we denote \\bfPhi, consists of all relevant physical parameters that determine the intrinsic shape and color of the light curve (e.g. the \\Ni\\ mass, the degree of mixing in the stellar interior, etc.). A second set of parameters, \\bftheta, relate to the location of the SN in space and time: redshift $z$, luminosity distance \\mue,\\footnote{The SOFT luminosity distance parameter \\mue\\ is the distance modulus residual relative to an empty universe. See equation 3 of Paper I.} host galaxy extinction \\Av, and time of peak \\tpk. These four {\\em location parameters} modify the intrinsic light curve shape by stretching and dimming it into the form that we eventually observe. We do not have a complete physical model for SN explosions that could efficiently and reliably translate a vector of physical parameters \\bfPhi\\ into a precise light curve prediction. Therefore, most \\TNSN\\ light curve fitters -- such as MLCS \\citep{Riess:1996,Jha:2002} and SALT \\citep{Guy:2005,Guy:2007} -- use one or more paramaters (e.g. $\\Delta$, or X1 and C) to define the shape and color of a broad-band SN light curve. By contrast, the SOFT method has no shape or color parameters, but instead uses a collection of SN light curve templates to provide examples of possible shapes and colors. These template light curves are then warped to reproduce the effects of a change in location \\bftheta. By comparing this adapted light curve model against a candidate SN object, we can derive the likelihood that the candidate has an intrinsic light curve shape similar to the template, and is at the assumed location \\bftheta. In Paper I we discussed how our template-based approach is best served by utilizing the framework of fuzzy set theory \\citep{Zadeh:1965}, and we applied SOFT to the task of SN classification. In this paper we will now explore how the application of fuzzy logic allows us to extend the SOFT method to the problem of estimating cosmologically useful parameters such as a SN's redshift and distance. ", "conclusions": "\\label{sec:Summary} In Paper I we introduced the SOFT program, using a set of fixed-shape light curve templates and the framework of fuzzy set theory for combining results from multiple templates. We applied SOFT as a SN classification tool in Paper I, and in this companion paper we have shown how SOFT can also be used to estimate parameters of cosmological interest: redshift $z$ and luminosity distance \\mue. The SOFT method is distinct from light curve fitters such as MLCS and SALT in that it does not describe the light curve shape with a parameterized model. Given this fundamental distinction, SOFT may provide a valuable addition to the parametric fitters in that it could have smaller $-$ or at least {\\em different} $-$ systematic uncertainties. Using Type Ia SN light curves from the SDSS and SNLS data sets, we have performed a set of verification tests to demonstrate the accuracy of SOFT. We found that the SOFT redshift estimates are comparable to optical photo-z measurements for SDSS galaxies: both methods yield $\\delta z=z_{SOFT}-z_{spec}$ residuals with an RMS scatter of approximately 0.05. Applying SOFT to derive (z,\\mue) coordinate pairs for a joint SDSS$+$SNLS sample, we find an RMS scatter around the $\\Lambda$CDM model of 0.18 mags in \\mue\\ when including spectroscopic redshift priors. When SOFT analyzes the SDSS$+$SNLS sample with no spectroscopic information at all, we find the RMS scatter of Hubble residuals increases to 0.37 mags. To investigate the near future of SN cosmology, we considered the degree to which a variable DE model can be constrained by larger SN samples containing thousands of light curves but very little spectroscopic follow-up. Using a bootstrap Monte Carlo approach, we simulated the distance and redshift estimates that might be obtained by SOFT for 9 different survey structures. We find that SOFT should be able to improve the DETF FoM by a factor of 2--3 with each new generation of SN surveys. The SOFT program has not yet realized the full potential of its fuzzy logic approach, but we have outlined several pathways for improving the method and limiting systematic biases. We have proposed a straightforward method for using a training set to calibrate the template library, and have also presented a ``Peer to Peer Cosmology'' approach in which SOFT can identify groups of similar SNe and do a direct comparison to determine their redshifts and relative distances. This latter method may be especially applicable in upcoming SN surveys such as Pan-STARRS and LSST, which will have an abundance of well-sampled multi-color light curves, but comparatively little spectroscopic followup. {\\bf Acknowledgments:} We would like to thank the anonymous referee for a thorough and critical reading of this work, which led to substantial improvements." }, "1003/1003.5662_arXiv.txt": { "abstract": " ", "introduction": "\\label{sec:intro} It is by now well established that about a quarter of the mass of the universe is in the form of some kind of non-luminous, non-baryonic matter. The standard cosmic history leading to thermal relic dark matter (DM), where DM annihilations freeze out at temperatures near its mass, is certainly an attractive possibility, but it is not the only one. Indeed, the observed baryon and lepton number densities did not arise due to thermal freeze-out of annihilations; they are due to a primordial asymmetry between the number densities of particles and anti-particles. A scenario of this type, ``Asymmetric Dark Matter'' (ADM), is also perfectly viable to explain the origin of the relic density of dark matter particles. One compelling idea is to have the dark matter asymmetry arise due to a direct connection to the asymmetry in the baryons and leptons (see e.g. \\cite{Kaplan:1991ah, farrar, ADM,Barr:1990ca,Thomas:1995ze,Kuzmin:1992up,Tytgat:2006wy,An:2009vq,Cai:2009ia,Kitano:2008tk,JMR}). In particular, one can imagine that the dark matter particle itself carries baryon number and/or lepton number. In that case the dark matter will share in the asymmetry of regular matter in the early universe. This general scenario was proposed in \\cite{ADM}, and it is the scenario that we will be assuming throughout this paper. One of its most appealing qualities is that it leads generically to a prediction for the mass of the dark matter particle; since the asymmetry in the dark sector ends up being roughly comparable to that in the baryons, the ratio of dark matter to baryonic energy densities is given approximately by the ratio of the dark matter mass $m_{\\tmop{DM}}$ to the mass of the proton. This leads to a prediction for $m_{\\tmop{DM}}$ of about $\\CO(1-10)$ GeV depending on the details of the model. Here we are simply assuming that all interactions which transfer the asymmetry between the dark and visible sectors have frozen out of equilibrium by the time the temperature has dropped below $m_{\\tmop{DM}} .$\\footnote{If the asymmetry transfer freezes out below the DM mass, the densities of DM and baryons are not related to each other in such a straightforward way, and depend on the size of the DM annihilation cross-section, much as in the standard thermal WIMP case. } It is an amusing coincidence that the prediction here for $m_{\\tmop{DM}}$ is fairly close to the weak scale, and thus could arise in models of electroweak symmetry breaking, much as in the case of the standard weakly interacting massive particle (WIMP). The fundamental test of the WIMP hypothesis will eventually be an issue of measuring the couplings of a given WIMP candidate, in order to check whether the associated annihilation cross section would indeed have led to an appropriate relic density. Such a confirmation would yield convincing evidence that the dark matter relic density did indeed have a thermal origin. Similarly, obtaining convincing evidence that the ADM mechanism is at play may also be possible. There would be two crucial ingredients which would be important to confirm; first, one would like to check that the dark matter particles do indeed carry a particular baryon number and/or lepton number, and second, one would like to check that $m_{\\tmop{DM}}$ is of an appropriate size. A difficulty here is that, since we are requiring that the operators responsible for carrying the baryon/lepton asymmetry between the dark and visible sectors froze out at temperatures above $m_{\\tmop{DM}}$, these operators must be fairly suppressed. On the other hand, it is precisely these operators which would reveal the baryon/lepton number carrying properties of the dark sector; the requirement of early freeze-out implies that signatures of these operators will in general be difficult to detect. In this paper, we will discuss one reasonably generic signature which the asymmetry transfering operators might have. If the dark matter carries baryon/lepton number, then decays or annhilations of this particle into the standard model leptons and/or baryons may in general be made possible. Since electric charge is conserved and the DM particle must be electrically neutral, the excess standard model baryon/lepton number will necessarily show up either in the neutrino sector, or in equal numbers of electrons and protons (or their anti-particles). Due to the requirements on $m_{\\tmop{DM}}$ in this scenario, neutrinos from DM decays would typically have energies of order a few $\\tmop{GeV}$, and could lead to a distinctive bump in cosmic neutrino data in this energy range. Moreover, due to the origin of the signal, this bump would be associated with anti-neutrinos rather than neutrinos (or vice versa), and this is a property which is potentially discernable at the Super Kamiokande detector \\cite{Ashie}, as well as MINOS \\cite{MINOS} and possible future detectors. This feature would make an observation of this type particularly compelling evidence in favor of the ADM mechanism. Unfortunately, in the case of operators leading to dark matter annihilations, early freeze-out requirements constrain the event rates to be considerably below the reach of both current and upcoming neutrino data. We relegate the details to appendix \\ref{app:ann}. As a result, we shall instead focus on operators leading to dark matter decays. As implied above, our most important limits will come from the Super Kamiokande (Super-K) neutrino detector. The MINOS detector has about 20\\% the fiducial volume of Super-K, though this is partly compensated for by its ability to directly distinguish between neutrinos and anti-neutrions, thereby cutting down on the atmospheric neutrino background in the ADM scenario. Although the IceCube \\cite{ICECUBE} experiment now typically sets more stringent limits on neutrino signals than Super-K, its low energy threshold of $\\sim 100 \\tmop{GeV}$ is too high for our purposes. We find that operators of dimension 6 leading to DM decays must be suppressed by at least $\\sim 10^{12} - 10^{13}$ GeV, due to existing neutrino constraints. In many models of the early universe, the reheating temperature is required to be less than about the GUT scale, and thus, upcoming improvements in the constraints will be probing into the remaining window where the operator responsible for decay is also capable of transfering the asymmetry between the SM and the DS. The outline of this paper is as follows. In section 2, we review the asymmetric dark matter (ADM) paradigm, and how one obtains quantitative predictions for the DM mass, focussing on a specific choice for the operator responsible for DM decay to neutrinos. In section 3, we discuss details of the production and observation of neutrinos from DM decays, and the resulting constraints on the decay spectrum. We also discuss possibilities for distinguishing neutrinos from anti-neutrinos at water-Cherenkov detectors and other possible future detectors. In section 4, we consider more general operators which may lead to the decay of the DM particle. For each operator, our analysis gives a bound on the size of the scale suppressing the interaction. In particular, we consider all possible interactions of dimension $\\le 6$ coupling the DM sector to a SM $(B-L)$-carrying gauge-singlet operator. Finally, in section 5, we speculate on future directions and model-building issues. ", "conclusions": "Asymmetric dark matter is a compelling and simple framework, alternative to the standard thermal WIMP paradigm. In particular, ADM favors lighter DM masses around $\\CO(1-10)$ GeV, which have received a recent boost in interest due to hints of a signal coming from the CoGeNT direct detection experiment \\cite{COGENT}, as well as possible explanations for the DAMA anomaly \\cite{DAMA,DAMAChannel, Feldstein, Fitzpatrick, Hooper, Chang, Petriello, Gondolo, WIMPless}. The phenomenological implications can differ in significant ways from those of a standard thermal WIMP, and we have noted one such possibility here, in which the DM particle carries lepton number and decays dominantly to anti-neutrinos. Such a scenario is motivated as a reasonably generic consequence of a mechanism for transferring the SM baryon asymmetry to the DS by virtue of the lepton number of the DM particle \\cite{ADM}. There are various model-building issues worth exploring in ADM. One would like to specify the mechanism that generates the asymmetry in the first place, and how the symmetric component is to be removed. Furthermore, one needs to generate a mass for the dark matter particle with the appropriate size to give the correct relic density. The fact that the typical masses required are ${\\cal O}(1-10$ GeV$)$ suggests a common origin with the electroweak scale, suppressed by additional loop factors or smaller couplings. If the dark matter is a scalar, the mass also needs to be protected from radiative corrections. This is clearly related to the issue of forbidding or suppressing the marginal coupling $|X|^2 |H|^2$ to the Higgs. Also, since conservation of lepton number is protecting the lifetime of the dark matter particle, it may not always be straightforward to take advantage of the see-saw mechanism for Standard Model neutrino masses. We will comment briefly on how these issues may be addressed in one example using mechanisms already suggested in the literature. As noted in \\cite{ADM}, the operator ${\\cal O}_2$ is attractive since it may easily be UV-completed to a model with Standard Model Majorana neutrino masses arising from the see-saw mechanism. For example, we may add two $SU(2)$ doublet fermions $d,d^c$ with $L=1+l_X, -1-l_X$ and hypercharge $Y=\\half, -\\half$ respectively, and take \\be \\Delta {\\cal L} &=& \\lambda \\psi d H + \\lambda' L d^c X + m_d d d^c. \\label{eq:simpmodel} \\ee The right-handed neutrino masses arise from a scalar vev $\\phi$ with $L=+2$. Thus no renormalizable couplings of $\\phi$ to the fields $X,\\psi,d,d^c$ are allowed as long as $ -\\frac{2}{3} < l_X < 0 $, and lepton number can then easily be an accidental symmetry in the dark sector. Note, however, that in this particular example, the operator ${\\cal O}_2$ is parametrically suppressed only by the scale $m_d$, which would be constrained by observations to be above the Planck scale according to Table 2 if the decay of the dark matter particle $X \\rightarrow \\psi \\nu$ were kinematically accessible. On the other hand, simple extensions may suppress ${\\cal O}_2$ without requiring such high scales. For example, consider a model with an extra gauged $U(1)$ broken by a scalar vev $\\langle \\Phi \\rangle$, and two copies $d_i, d_i^c$ of the doublet fields, where the non-zero $U(1)$ charges of the fields are as follows: \\begin{center} \\begin{tabular}{c|c} & U(1) \\\\ \\hline $X$, $d_2$ & 1 \\\\ $\\Phi$, $d_2^c$ & -1 \\end{tabular} \\end{center} \\noindent and 0 for the remaining fields. Then symmetry forces the lagrangian to be of the form \\be \\Delta {\\cal L} &=& \\lambda \\psi d_1 H + \\lambda' L d^c_2 X + m_{d,i} d_i d_i^c + c \\Phi d_2 d_1^c, \\ee which upon integrating out the $d_i$ fields leads to ${\\cal O}_2$ having a coefficient whose parametric suppresion is $\\frac{c \\lambda \\lambda' \\langle \\Phi \\rangle }{m_d^2}$. So far, in our discussions, the DM mass has been put in by hand. However, if the baryon-to-dark matter density ratio is truly to be explained, then $m_{\\rm DM}$ must be dynamically generated at the correct scale. As noted in \\cite{Morrissey:2009ur}, a mass of ${\\cal O}(1-10$ GeV$)$ arises quite generically in models of spontaneously broken supersymmetry from gravity-mediated effects, provided that the MSSM masses are generated by gauge mediation with a messenger scale at $M \\sim 10^{13}$ GeV. It would be more satisfying, though, to have a model where the mechanism that generates the DM mass is more predictive, and does not depend on a mass parameter that is free to vary over several orders of magnitude. Supersymmetry has the advantage that in theories with large $\\tan \\beta\\sim 20$, allowing a superpotential term of the form $W \\supset X H_d S$, where $S$ is a doublet, generates a mass for $X$ at the appropriate size. Furthermore, additional contributions to the dangerous marginal term $|X|^2 |H|^2$ can come only from superpotential terms, which may be easily controlled. There are many ways the asymmetry can be generated; for an incomplete list see \\cite{KolbTurner}. It is perhaps worth noting that the asymmetry does not have to originate in the SM and get transferred to the DS but could instead originate due to new sources of CP violation in the DS and then be transferred in the other direction. Finally, there are other possible ways that neutrino signatures of ADM could appear. One potentially interesting direction currently under investigation % involves neutrinos coming from DM annihilation in the sun. ADM is very interesting in that its annihilation cross-section is not a priori related to the cross-section for capture in the sun. Indeed, because of conservation of lepton number, the scattering process will never contribute radiatively to the annihilation process, and thus there is no theoretical barrier to taking the two cross-sections to be quite different.\\footnote{ It is conceivable that an additional potentially significant boost in the solar signal could come from exponential growth of the DM occupation in the sun due to WIMP-WIMP scattering \\cite{zentner}.}" }, "1003/1003.4791_arXiv.txt": { "abstract": "How to create planetesimals from tiny dust particles in a proto-planetary disk before the dust particles spiral to the central star is one of the most challenging problems in the theory of planetary system formation. In our previous paper \\citep{kato08}, we have shown that a steady angular velocity profile that consists of both super and sub-Keplerian regions is created in the disk through non-uniform excitation of Magneto-Rotational Instability (MRI). Such non-uniform MRI excitation is reasonably expected in a part of disks with relatively low ionization degree. In this paper, we show through three-dimensional resistive MHD simulations with test particles that this radial structure of the angular velocity indeed leads to prevention of spiral-in of dust particles and furthermore to their accumulation at the boundary of super-Keplerian and sub-Keplerian regions. Treating dust particles as test particles, their motions under the influence of the non-uniform MRI through gas drag are simulated. In the most favorable cases (meter-size dust particles in the disk region with a relatively large fraction of MRI-stable region), we found that the dust concentration is peaked around the super/sub-Keplerian flow boundary and the peak dust density is 10,000 times as high as the initial value. The peak density is high enough for the subsequent gravitational instability to set in, suggesting a possible route to planetesimal formation via non-uniformly excited MRI in weakly ionized regions of a disk. ", "introduction": "One of the most serious problems in planet formation is how the dust components grow up to larger objects in protoplanetary disks. The known key-stage is when the dust particles are meter-size \\citep{adachi76, wei77}. Since gas pressure gradient in a protoplanetary disk is negative, gas rotates at sub-Keplerian velocity. On the other hand, the dust particles tend to rotate at Keplerian velocity and therefore they feel headwind of gas. Losing their angular momenta, they spiral into a central star quickly on timescales of ~100 years. One possibility to overcome the meter-size barrier is the planetesimal formation via gravitational collapse of a cluster composed with small dust particles \\citep{saf69, gold73}. The dust would settle to midplane in the absence of turbulence. Early studies of dust settling in turbulent accretion disks were carried out by \\citet{cuzzi93}, \\citet{miyake95} and \\citet{dub95}. They presented that large dust particles settle down toward their equilibrium distribution in a turbulent diffusive time scale while small ones remain mixed throughout the whole gas disk. If dust density in midplane becomes high enough, Kelvin-Helmholtz instabilities between a sediment layer of dust and the overlying gas can drive turbulence and mix the dust layer enough to prevent gravitational instability \\citep{wei77, wei80, ishi03, gomez05}. \\citet{joh06b}, however, carried out two-dimensional simulation and found active turbulent concentration. Except the effect of Kelvin-Helmholtz instabilities, the high dust density in midplane leads to a streaming instability \\citep{good00, you05} because the dust is susceptible to preferential radial migration relative to the gas \\citep{naka86}. The streaming instability was shown to lead to turbulence and concentrate dust particles locally \\citep{you07, joh07a}. As one of other mechanisms of turbulent concentration, the effect of the turbulence driven by the magnetorotational instability (MRI; \\citealp{bal91, haw95}) is recognized. Planetesimal formation by gravitational instability in the MRI turbulence is addressed by \\citet{joh06a, joh07b} and \\citet{balsara09}. Turbulence level is the most fundamental parameter governing whether planetesimals can form by gravitational instability in these models. The turbulence level of MRI depends on the ionization degree of disk gas and the strength of magnetic field. The linear analysis showed that the growth rate of MRI is reduced by ohmic diffusion when collisions are so frequent that not only the ions and neutrals are well coupled but also electrical currents are damped \\citep{jin96, sam99}. At the same time, the growth wavelength becomes longer and transport rate of angular momentum is diminished. The nonlinear evolution of the MRI in nonideal MHD starting from the weak vertical magnetic field has been studied by \\citet{sano98}. They found that sustained MHD turbulence requires the magnetic Reynolds number $R_{\\rm m}=v^2_{\\rm A}/\\eta \\Omega \\gtrsim 1.0$, where $v_{\\rm A}$, $\\eta$ and $\\Omega$ are Alfven velocity, magnetic resistivity and Keplerian rotation frequency. The values of $\\eta$ are regulated by an ionization structure in a protoplanetary disk. The inner regions of a disk are thermally ionized \\citep{pne65, ume83}. At greater distances from the disk center, stellar X-rays and diffuse cosmic rays ionize the surface gas layer down to a certain column density \\citep{glass97, igea99}. In moderately distant regions in which disk column density is high enough, midplane layers are not highly ionized and MRI could be inactive there. This region is called \"dead zone\" \\citep{gam96}. The characteristics and interesting physics are prospective around dead zones. The magnetic turbulence outside the dead zone, which transports angular momentum and mass, creates gas density bump at the outer boundary of the dead zone. This density pile-up triggers the Rossby wave \\citep{li01, varni06} to form long-life anticyclonic vortices. \\citet{barge95} suggested that the dust particles are trapped in a center of vortex. By numerical simulation of Rossby vortices including solid particles, \\citet{ina06} confirmed the enhancement of dust density there. \\citet{lyra08} showed that gravitationally bound embryos can form in the regions of enhanced dust density. \\citet[][; hereafter referred to as Paper I]{kato08} showed that non-uniform excitation of MRI creates angular velocity profile that consists of super and sub-Keplerian regions locally in the disk (for details, see section 2). It is expected that this radial structure of the angular velocity leads to prevention of spiral-in of dust particles and their accumulation at the boundary of super and sub-Keplerian regions, since the super-Keplerian flow pushes the dust particles outward while the sub-Keplerian flow drags them inward. This interesting mechanism is based on the co-existence of MRI active and inactive regions. The ionization degree that regulates MRI depends on the relative abundance and size distribution of grains \\citep{wardle99, sano00}. It can also be changed by phase chemistry \\citep{ilgner08}. Slight changes in these quantities near boundaries at global active/dead zones of MRI may create the local inhomogenity in MRI growth. We will address this issue in a subsequent paper. Note that the local model in this paper is not a toy model that mimics the global inhomogeneous structure of dead and active zones. In this paper, we investigate dust accumulation induced by inhomogeneous MRI evolution by three-dimensional non-ideal MHD simulation including dust particles. The dust particles are represented by Lagrangian particles suffering gas drag and are assumed not to affect on gas. We consider both cases with radially non-uniform magnetic resistivity under uniform magnetic field direction and non-uniform magnetic field with radially uniform magnetic resistivity. MRI develops in low resistivity regions in the former case and in strong vertical magnetic field regions in the latter case. The latter model is the same as the model adopted in Paper I. We briefly summarize the results of Paper I in section~\\ref{sec:paper1}. We explain our simulation setting in section~\\ref{sec:model}. In section~\\ref{sec:result}, we present that the MRI leads to quasi-steady state and particles are locally concentrated as expected. We also analyze the distributions of particle radial velocity to understand the effect of weak remnant magnetic turbulence. We examine how high the particle density enhancement is and the possibility of planetesimal formation via gravitational instability. Finally, we summarize our results and discuss their implications in section~\\ref{sec:discussion}. ", "conclusions": "\\label{sec:discussion} We have investigated the dust concentration in protoplanetary gas disks that are in the quasi-steady state created by inhomogeneous magnetorotational instability. We set the inhomogeneous instability with the initial radial non-uniformity of either ionization degree or vertical component of the magnetic field. The result of our local three-dimensional resistive MHD simulations with Lagrangian particles clearly shows that the meter or decimeter sized dust particles are concentrated strongly and steadily in either setting except for the cases with $R_{\\rm m,ave} > 0.1$ (where $R_{\\rm m,ave}$ is the averaged magnetic Reynolds number averaged over the simulation box), which can allow planetesimal formation via gravitational instability. The MRI growth in our model generates the locally confined super-Keplerian region and this flow is quasi-stable due to the support of gas pressure gradient which is also formed by the MRI. Because unperturbed gas flow is sub-Keplerian due to global pressure gradient, the particles suffering gas drag are concentrated in the Keplerian domain at the outer-edge of the super-Keplerian area. Since the flow is in a quasi-steady state, particles are supplied from outer regions by gas drag migration and the particle density increases significantly. The process of particle concentration and the increasing rate of dust density depend on the particle size, the initial radial widths of unstable/stable regions, and the initial non-uniformity setting as follows: \\begin{description} \\item[ \\textit{The initial non-uniform setting, magnetic field or resistivity}:] The non-uniform resistivity model produces more stable and clean state of local rigid rotation than the non-uniform magnetic field model, because the high resistivity dissipates the magnetic perturbations more rapidly in the stable regions. Indeed, the velocity dispersion of particles in clumps is smaller in the non-uniform resistivity model. In the non-uniform magnetic field cases, the magnetic reconnection and the remnant weak instability create less smooth gas velocity field, so that particles tend to assemble at multiple radial positions. On the other hand, in the non-uniform resistivity case, all particles are concentrated at a specific radial position, and the peak values of the particle density in clumps is higher, although the particle density also depends on the size of numerical box. \\item[\\textit{The dust size, meter or decimeter}:] The meter-size particles are less coupled with gas motion than decimeter-size ones. Since their radial migration is faster, the meter-sized particles are more likely to be condensed in the radially narrow area without being scattered by the remnant gas turbulence. The peak dust density becomes $\\ga 1000$ times larger than the initial value. Even in the decimeter-size particle cases, the velocity dispersion of particles in the clump excited by the remnant turbulence is sufficiently smaller than radial drift velocity due to the quasi-steady gas flow and dust concentration is still observed in the simulations, although the enhanced dust density is at most $\\sim 100$. \\item[\\textit{The degree of remnant turbulence}:] When the stable region is initially smaller in the non-uniform magnetic field model (CASE2), that is, $R_{\\rm m,ave}$ is larger, stronger turbulence tends to remain. For $R_{\\rm m,ave} \\la 1$, however, the super-Keplerian area continues to exist and clumps with density $\\sim 100$ times larger than initial values are robustly created. Furthermore, for $R_{\\rm m,ave} \\la 0.1$, the clumps are persistent enough for following gravitational instability, while they are tentative and repeatedly produced in the cases of $0.1 \\la R_{\\rm m,ave} \\la 1$. \\end{description} We should address the following points in subsequent papers: \\begin{description} \\item[\\textit{consistent magnetic resistivity}:] We need to calculate the value of magnetic resistivity consistently as a time dependent value, because it is sensitive to both gas and dust densities. Around strong dust concentration areas, gas and dust density become larger and magnetic resistivity becomes higher, which may calm down the remnant turbulence. \\item[\\textit{Rossby wave instability (RWI)}:] RWI can be caused from a pressure bump on the radial-azimuthal plane. If we simulate in an enlarged box, vortex may grow, although in the non-uniform magnetic field cases, the growth of the instability may be reduced by the azimuthal magnetic field. If RWI occurs, it may serve as a mechanism of dust concentration in the azimuthal direction \\citep{ina06,lyra08}. \\item[\\textit{the vertical structure of a protoplanetary disk}:] In the upper disk regions, ionization degree may be so high that the MRI to occur \\citep{fle03, fro06, oishi07}. The turbulence there may affect the small dust dynamics near the midplane. \\item[\\textit{feedback from dust particles onto gas}:] The feedback may be the most important. A great concentration in our results here will inevitably change gas velocity field and it could affect the condition of dust concentration (density enhancement, velocity dispersion and so on) and planetesimal formation. We have already started simulations with this effect. \\item[\\textit{collisional destruction}:] While the velocity dispersions of dust particles in clumps are extremely low in some of our results, the dust particles can be collisionally disrupted into small fragments that are coupled to gas motion and the fragments may be dispersed by even weak turbulence. For clumps to gravitationally collapse, its timescale needs to be smaller than mean collision time. On the other hand, if the dust grains are also produced by the collisions, they might affect the gas ionization state through recombination of electrons onto their surface. We need to address these feedback effects by collisional disruption as well. \\end{description}" }, "1003/1003.4758_arXiv.txt": { "abstract": "Supra-arcade downflows (SADs) have been observed with \\textit{Yohkoh}/SXT (soft X-rays (SXR)), \\textit{TRACE} (extreme ultra-violet (EUV)), \\textit{SoHO}/LASCO (white light), \\textit{SoHO}/SUMER (EUV spectra), and \\textit{Hinode}/XRT (SXR). Characteristics such as low emissivity and trajectories which slow as they reach the top of the arcade are consistent with post-reconnection magnetic flux tubes retracting from a reconnection site high in the corona until they reach a lower-energy magnetic configuration. Viewed from a perpendicular angle, SADs should appear as shrinking loops rather than downflowing voids. We present XRT observations of supra-arcade downflowing loops (SADLs) following a coronal mass ejection (CME) on 2008 April 9 and show that their speeds and decelerations are consistent with those determined for SADs. We also present evidence for a possible current sheet observed during this flare that extends between the flare arcade and the CME. Additionally, we show a correlation between reconnection outflows observed with XRT and outgoing flows observed with LASCO. ", "introduction": "While the details of flare dynamics are debatable, the general energy release mechanism is widely accepted to arise from magnetic reconnection. Direct evidence of reconnection, however, has been scarce and often questionable. In all models of reconnection, loops flowing both towards and away from the reconnection site are an inevitable theoretical consequence (\\citeauthor{carm} \\citeyear{carm}; \\citeauthor{sturrock} \\citeyear{sturrock}; \\citeauthor{hirayama} \\citeyear{hirayama}; \\citeauthor{kopp} \\citeyear{kopp}). Observationally, though, these loops require very special circumstances in order to image since they are likely to be relatively devoid of emitting plasma and form high in the corona where they are viewed against a dark background. Supra-arcade downflows (SADs) have been observed in several flares and interpreted as the cross-sections of these shrinking loops as they retract through a bright fan (\\citeauthor{mck99} \\citeyear{mck99}; \\citeauthor{mck00} \\citeyear{mck00}; \\citeauthor{sadsI} \\citeyear{sadsI}). Cusped flares, like those predicted by the standard models, have been shown to have signatures of retracting loops (\\citeauthor{forbesacton96} \\citeyear{forbesacton96}; \\citeauthor{reeves08} \\citeyear{reeves08}). Imaging individual loops retracting above the flaring site with high enough temporal and spatial resolution has proven to be a challenge due to observational limitations. In order to observe the downflows which occur above the post-eruption arcade, the flare must occur near the limb; and long image exposures, which inevitably saturate the flaring site and are therefore not desirable for most flare observations, must be taken to provide proper contrast in the low signal to noise region above the arcade. Coronal mass ejections (CMEs) are frequently observed to be associated with eruptive flares. Current sheets are expected to extend between the arcade region and the CME \\citep{forbesacton96}. While the current sheets themselves are probably too narrow to be fully resolved with current instrumentation, calculations have shown that conduction fronts lead to the formation of a sheath of hot plasma surrounding the current sheet that widens the observable structure (\\citeauthor{reevesinpress} \\citeyear{reevesinpress}; \\citeauthor{seaton09} \\citeyear{seaton09}; \\citeauthor{yokshib98} \\citeyear{yokshib98}). Recent modeling has shown that this hot plasma sheath can be observed by a sensitive X-ray imager such as Hinode's X-ray Telescope (XRT) (\\citeauthor{reevesinpress} \\citeyear{reevesinpress}). Current sheet observations have been claimed and analyzed for several flares using EUV and white light coronagraphs (\\citeauthor{ciaray08} \\citeyear{ciaray08}; \\citeauthor{linetal07} \\citeyear{linetal07}; \\citeauthor{webb03} \\citeyear{webb03}; \\citeauthor{ciaetal02} \\citeyear{ciaetal02}; \\citeauthor{ko02} \\citeyear{ko02}). Because white-light coronagraphs measure polarization brightness, which is directly related to density, these measurements indicate that the density in the structures surrounding the current sheet is elevated compared with the background corona. In the following sections, we describe the ``Cartwheel CME\" flare as seen by XRT and LASCO. We describe the XRT observations in detail, which include a candidate current sheet, shrinking loops, and flows, and then show correspondences between XRT and LASCO flows. We also discuss a possible scenario for interpreting the observations based on magnetic modeling. ", "conclusions": "Downflowing oblong voids have been observed with SXT and TRACE above arcades resulting from long-duration flaring events. Similar flows have been observed on much larger scales with LASCO although not always in association with flare arcades (\\citeauthor{mck99} \\citeyear{mck99}; \\citeauthor{mck00} \\citeyear{mck00}; \\citeauthor{mck99} \\citeyear{mck99}; \\citeauthor{sadsI} \\citeyear{sadsI}; \\citeauthor{asai04} \\citeyear{asai04}; \\citeauthor{khan07} \\citeyear{khan07}; \\citeauthor{sheeley04} \\citeyear{sheeley04}). These features, known as supra-arcade downflows (SADs), have been interpreted as the cross-sections of individual shrinking magnetic loops as they retract from the reconnection site high above the arcade. The dark loops seen early in the image sequence are consistent with the ``shrinking empty loop\" discussion pertaining to SADs in \\cite{mck00} and the spectral measurements of \\cite{innes03}. The brightening of the flows in the middle of the sequence may have resulted from increased density in the current sheet as the flare progresses. Figure~\\ref{diagram}, adapted from \\cite{sadsI}, gives quantitative estimates of pertinent parameters that describe this scenario based on a small sample size. The viewing angle for SADs with respect to Figure~\\ref{diagram} (a) is from the side (i.e. perpendicular to the arcade). If this scenario is correct, then for the 2008 April 9 event, the viewing angle is face-on to the loops (i.e. along the axis of the arcade) so as to observe supra-arcade downflowing loops (SADLs) as shown in Figure~\\ref{diagram} (b). The perpendicular viewing angle is supported by the observation of shrinking loops in XRT and SECCHI as well as the magnetic field configuration derived with the PFSS modeling (see Figure~\\ref{tracks} (b)). The available relevant information obtained for this event has been labeled in the figure. \\begin{figure}[!ht] \\begin{center} \\includegraphics[width=.7\\textwidth]{f23a.pdf} \\includegraphics[width=.7\\textwidth]{f23b.pdf} \\caption{(a) Schematic diagram of supra-arcade downflows (SADs) resulting from 3-D patchy reconnection. Discrete flux tubes are created, which then individually shrink, dipolarizing to form the post-eruption arcade. The measured quantities shown are averages from the following events: 1999 January 20 (M5.2); 2000 July 12 (M1.5); 2002 April 21 (X1.5) \\citep{sadsI}. (b) Schematic diagram of supra-arcade downflowing loops (SADLs) also resulting from 3-D patchy reconnection. The measured quantities are averages from the flare occurring behind the western limb on 2008 April 9. Note that the viewing angle is perpendicular to that of the SADs observations.} \\label{diagram} \\end{center} \\end{figure} This flare is fairly unusual in that it occurred behind the limb in XRT allowing for better image quality high above the arcade with support from STEREO A observations of the flare on the limb. Under these circumstances, a relatively faint, thin candidate current sheet (CCS) is observed in XRT and the height at which the flows are first detected is increased from that of the previously-analyzed SADs thereby increasing the overall length of the downflow path as well. The striking similarity between the speeds and path lengths offers support to the argument that these two observational features are different views of the same phenomenon. Additionally, we observe a disconnection event which we interpret as a reconnection outflow pair. We believe that this is the first clearly-observed reconnection outflow pair observed so near to the solar surface along a directly observable SXR current sheet. The relatively narrow initial height range ($\\sim$100 Mm) with respect to the CCS extent (see Figure~\\ref{cs_qp_plot}) combined with the reconnection outflow pair observations seems to imply patchy reconnection with a relatively localized acceleration region. The possible correspondence of upflows observed to propagate from within the XRT FOV into the LASCO FOV, notably the upflow associated with the LASCO CME ``pinch-off point\", also support this statement. Reports of impulsive-phase RHESSI double coronal sources have been made by other authors (e.g. \\citeauthor{liu08} \\citeyear{liu08}; \\citeauthor{sui04} \\citeyear{sui04}; \\citeauthor{sui03} \\citeyear{sui03}) with the lower source corresponding to the top of the rising arcade and the upper source possibly corresponding to an ascending reconnection outflow or an ejected plasmoid. We note that these upper coronal sources have paths, speeds, and placements relative to the arcade similar to the bright plasmoid structure tracked as the eruption front for the ``Cartwheel CME\" flare (Figure 2). This is consistent with the extrapolated position of one coronal source from the 2002 April 15 flare \\citep{sui04} which was calculated to roughly track with a coronal loop observed by LASCO (similar to Figure 16 (top)). \\citeauthor{sui04} conjecture that ``the outward-moving coronal source is part of an ejected plasmoid (or a large-scale, helically twisted loop) with two ends anchored on the Sun...\". This inference of ``an ejected ... large-scale, helically twisted loop\" matches our interpretation of the 2008 April 9 eruption invoking an erupted flux rope (Figure~\\ref{cs_cartoon}). \\cite{liu08} also report that their source closer to the solar surface has a larger emission measure than the higher one. These results are consistent with our ``Disconnection Event\" observations (see Section 3.2) where the upflow portion is much dimmer and more diffuse than its downflowing counterpart. Indeed, all of the upflows for this flare are dim compared to the bright downflows. It is also worth noting that a non-radial, southward evolution of the loop-top source is reported for the flares in \\cite{sui04} (Figure 10 therein). The source of this divergence may have a similar mechanism as that proposed for the apparent southward drift of the CCS for this flare (Figure 20). We interpret the basic standard picture of this eruptive flare as being initiated by the release of a flux rope by some means. As the flux rope escapes into the outer corona, a current sheet forms in its wake. Stretched magnetic field lines are swept together into the current sheet where reconnection occurs. This reconnection results in the formation of pairs of cusped, looped field lines, each moving in opposite directions along the current sheet. The loops retracting toward the solar surface form the post-eruption arcade while their companion loops are swept into the outer corona along with the erupted flux rope. The post-eruption arcade follows the direction of the active region's polarity inversion line (PIL). The current sheet also follows this direction as it spans the top of the developing arcade. The ``Cartwheel CME\" flare offers a unique glimpse into nearly every facet of this flaring process due to long image exposures made possible by the limb occultation of the bright footpoints. The event is observed by several instruments (TRACE, \\textit{STEREO A}/SECCHI, \\textit{Hinode}/XRT, and LASCO). SECCHI observes the onset of the flux rope eruption near the solar surface. The formation of a candidate current sheet is seen by XRT, and we provide a possible explanation for its apparent southward progression being due to its position along an inclined PIL. (See \\cite{landi10} for additional analysis of the CME.) Shrinking loops are very clearly seen in the XRT and SECCHI observations. Although typically very difficult to detect due to the low signal to noise above the flaring region, XRT is able to observe flows moving outward along the current sheet with one even clearly associated with a downflow. These upflows track into the outer corona where they appear to correspond with outflows seen by LASCO. Finally, a post-eruption arcade develops within both the XRT and SECCHI FOV. Any one of these observations provides an argument in favor of 3-D patchy reconnection flare models; however, all of these taken together makes the ``Cartwheel CME\" flare a very compelling candidate as direct proof of reconnection. \\clearpage" }, "1003/1003.3178_arXiv.txt": { "abstract": "Connecting cosmological simulations to real-world observational programs is often complicated by a mismatch in geometry: while surveys often cover highly irregular cosmological volumes, simulations are customarily performed in a periodic cube. We describe a technique to remap this cube into elongated box-like shapes that are more useful for many applications. The remappings are one-to-one, volume-preserving, keep local structures intact, and involve minimal computational overhead. ", "introduction": "Numerical simulations have become an indispensable tool in modern cosmological research, used for investigating the interplay of complex physical processes, studying regimes of a theory which cannot be attacked analytically, generating high precision predictions for cosmological models, and making mock catalogs for the interpretation and analysis of observations. Such simulations traditionally evolve the matter distribution in periodic cubical volumes, which neatly allow them to approach the homogeneous Friedmann solution on large scales. The use of a periodic volume also allows the long-range force to be easily computed by fast Fourier transform methods in many popular algorithms (e.g.~particle-mesh, particle-particle-particle-mesh or tree-particle-mesh algorithms). Surveys, on the other hand, often cover cosmological volumes that are far from cubical in shape, and making a mock catalog which includes the full geometrical constraints of the observations is difficult. One approach is to simulate a sufficiently large volume that the survey can be embedded directly within the cube, but this often means large parts of the computational domain are unused. An alternative is to trace through the cube across periodic boundaries so as to generate the desired depth, with various rules for avoiding replication or double-counting of the volume (if desired). A third approach is to run a simulation in a non-cubical geometry. If the side lengths are highly disproportionate this can lead to its own numerical issues, and in addition it makes it difficult to reuse a given simulation for many applications. In this paper we present a new solution to this problem which allows one to embed a hypothetical survey volume inside a cosmological simulation while limiting wasted volume and artificial correlations. The method is based on the simple observations that (1) a cube with periodic boundary conditions is equivalent to an infinite 3-dimensional space with discrete translational symmetry, and (2) the primitive cell for such a space need not be a cube. We show that one may take the primitive cell to be a cuboid% \\footnote{To be precise, by \\emph{cuboid\\/} we mean here a rectangular parallelepiped, i.e.~a parallelepiped whose six faces are all rectangles meeting at right angles. For practical purposes, a cuboid is simply a box that is not a cube.} of dimensions $L_1 \\times L_2 \\times L_3$ for a discrete but large choice of values $L_1\\ge L_2\\ge L_3$. The possible choices, subject to the constraint $L_1 L_2 L_3 = 1$, are illustrated in Figure \\ref{fig:poss}. \\begin{figure} \\begin{center}% \\resizebox{7cm}{!}{\\includegraphics{possremap}}% \\end{center}% \\vspace{-0.15in} \\caption{Possible dimensions $L_1 \\times L_2 \\times L_3$ for cuboid remappings, subject to the conditions $L_1\\ge L_2\\ge L_3$, $L_1 L_2 L_3=1$, and $L_1<7$. The choice used for the Stripe 82 remapping described in Section \\protect\\ref{sec:examples} is indicated by a red circle.} \\label{fig:poss} \\end{figure} Our approach leads to a one-to-one remapping of the periodic cube which keeps structures intact, does not map originally distant pieces of the survey close together, and uses no piece of the volume more than once. It complements existing techniques for generating mock observations (e.g.~% \\citealt{Blaizot+05,Kitzbichler+07}) or for ray-tracing through simulations (e.g.~\\citealt{WhiHu00,Vale+03,Hilbert+09,Fullana+10}). Though not ideal in all cases, our remapping procedure may be seen as a general-purpose alternative that neatly skirts many of the complications involved with previous methods. We begin in Section \\ref{sec:math} with a mathematical description of the remapping, and explain what choices of dimensions are possible. In Section \\ref{sec:algorithm} we describe how to implement this remapping numerically. We present a few useful examples in Section \\ref{sec:examples}, and conclude in Section \\ref{sec:discussion} with a discussion of some of the advantages and limitations of this method. ", "conclusions": "\\label{sec:discussion} As surveys become increasingly complex and powerful and the questions we ask of them become increasingly sophisticated, mock catalogs and simulations which can mimic as much as possible the observational non-idealities become increasingly important. \\citet{Angulo+10} have shown that simulations of one cosmology can be rescaled to approximate those of a different cosmology. We have introduced a remapping of periodic simulation cubes which allows one simulation to take on the characteristics of many different observational geometries. The use of such techniques enhances the usefulness of cosmological simulations, which often involve a large investment of community resources. The methods we introduced in this paper lead to one-to-one remappings of the periodic cube which keep structures intact, do not map originally distant pieces of the survey close together, and use no piece of the volume more than once. The remapping can be done extremely quickly, meaning it can be included in almost any analysis tool with negligible overhead. Remapping the computational geometry is, however, not without its limitations. First and foremost, although the target geometry may have sides much longer than the original simulation, the structures will not contain the correct large scale power since it was missing from the simulation to begin with. This problem becomes less acute as the original simulation volume becomes a fairer representation of the Universe. Secondly, if the target geometry is too thin it is possible for points which are far apart in the survey to come from points close together in the simulation volume, leading to spurious correlations. These are analogous to the artificial correlations between survey ``sides'' that occur when it is embedded in a periodic cube. Simply excluding a boundary layer from the remapped volume can tame such correlations. In addition to the description of the algorithm in this paper, we have made Python and C++ implementations of the remappings, along with further examples and animations, publicly available at {\\tt http://mwhite.berkeley.edu/BoxRemap}. \\vspace*{1cm}" }, "1003/1003.3949_arXiv.txt": { "abstract": "We present a new efficient technique for measuring evolution of the galaxy luminosity function. The method reconstructs the evolution over the luminosity-redshift plane using any combination of three input dataset types: 1) number counts, 2) galaxy redshifts, 3) integrated background flux measurements. The evolution is reconstructed in adaptively sized regions of the plane according to the input data as determined by a Bayesian formalism. We demonstrate the performance of the method using a range of different synthetic input datasets. We also make predictions of the accuracy with which forthcoming surveys conducted with SCUBA2 and the Herschel Space Satellite will be able to measure evolution of the sub-millimetre luminosity function using the method. ", "introduction": "A cornerstone of any working model of the formation of structure in the universe is knowledge of the galaxy luminosity function (GLF). The GLF is a measure of the comoving space density of galaxies per interval in luminosity. Determining how the GLF changes with cosmological epoch provides far-reaching insights into the processes that dictate the means by which galaxies form and evolve. This places strong statistical constraints on evolutionary theories, enabling determination of key characteristics such as the epoch of galaxy formation, merger rates, transformations between population types and the history of the universe's global rate of formation of stars. The value of measuring the evolution of GLFs has been appreciated for several decades and as such, many techniques have been employed to achieve this aim. Without any loss of generality, the GLF can be expressed as the local GLF scaled by an evolution function. The simplest and most direct method of estimating this evolution function is a model independent one where the GLF is measured at different epochs and compared with the local GLF. This provides a set of discrete estimates of the evolution at specific epochs. However, adopting a model-based procedure yields significant advantages. A well-proven approach is to assume the evolution function depends only on luminosity and redshift (see next section). For a given model of the evolution function chosen a priori, the best fit model is determined, subject to a set of observational constraints which contribute to different regions in the luminosity-redshift ($L-z$) plane. The virtue of model-based techniques is that a model consistent with all available observations can be used to make predictions by extrapolation into regions where data are sparse or lacking. Furthermore, a model can be used to identify those datasets whose improvement would most efficiently increase the constraints on the evolution function. Model-based techniques are typically regarded as belonging to one of two types. Parametric evolution models adhere to some preconceived notion regarding the evolution of galaxies, for example, pure luminosity evolution or luminosity + density evolution \\citep[see][and references therein for recent examples]{wall08}. Although an advantage of this type of modelling is that smooth functions are readily extrapolated, a major disadvantage is that the real evolution function may take on a very different form. So called 'free-form' methods are not biased in this way and allow for a greater degree of flexibility over the $L-z$ plane when attempting to determine the evolution function. Freeform techniques date back over three decades. \\citet{robertson78,robertson80} developed an iterative freeform method to evaluate the evolution of the luminosity function for radio galaxies. The method was limited by allowing the evolution to vary only as a function of redshift (and not luminosity) and not giving a satisfactory indication of the range of solutions permitted by the observations. \\citet{peacock81} introduced a new method that incorporated luminosity dependence and measured the uncertainty on the evolution by considering the variation between different freeform model predictions. This improved technique saw successful application \\citep[e.g.,][]{peacock85,dunlop90} although by the authors' own admission, the full extent of the uncertainty still could not be verified by the range of models tested. Furthermore, the evolution function was modelled as a series expansion and the set of best-fit expansion coefficients determined by minimising $\\chi^2$ using a non-linear search. This is both inefficient and does not guarantee that the global minimum has been found. In this paper, we propose a new method, inspired by Peacock \\& Gull but offering some significant improvements. The method computes the evolution as a discretised (i.e., pixellised) function over the $L-z$. plane. This has the advantage that the evolution can be solved linearly, ensuring the best fit solution is always found with a single, highly efficient matrix inversion. Crucially, the full range of solutions permitted by the observations can be determined in a few simple extra computational steps. Despite being pixellised, we demonstrate how linear regularisation allows the evidence to be extrapolated into regions of the $L-z$ plane that are lacking data or where data are sparse. Finally, we show how the pixellisation can be adapted according to the constraints provided by the data. The original motivation for this work was to investigate evolution of the luminosity function in the sub-millimetre (submm) waveband. The far-infrared (far-IR) and submm wavebands are particularly important for investigating the evolution of galaxies, not least because approximately half of the energy emitted by stars and active galactic nuclei since the big bang has been absorbed by dust and then re-radiated in these wavebands \\citep{fixsen98}. In view of this, our interests lie in investigating not only the monochromatic luminosity functions at these wavelengths, but also the `dust luminosity function', i.e., the space density of galaxies as a function of the total luminosity emitted by the dust in a galaxy. Presently, observational data in the submm are particularly poor. However, this situation will begin to rapidly improve with the arrival of new submm instruments such as the second generation Submm Common User Bolometer Array (SCUBA2) and the Herschel Space Observatory ({\\sl Herschel}). Although the method we have developed is applicable to the luminosity function in any waveband, we describe and explore its use in this paper with an eye on its future application to forthcoming submm data. In section \\ref{sec_method} we present the method. Section \\ref{sec_demo} demonstrates the method with a range of synthetic datasets providing differing levels of constraints on the $L-z$ plane. Finally, we summarise in section \\ref{sec_summary} and discuss practical aspects of applying the method to real data. ", "conclusions": "\\label{sec_summary} In this paper, we have presented a new method for measuring evolution of the galaxy luminosity function. The method computes the evolution as a discretised function of redshift and bolometric luminosity. The advantage brought about by this discretisation is that the evolution can be solved linearly, ensuring the best fit solution is always found with a single, efficient matrix inversion. Formulating this as a linear problem also brings the additional advantage that the uncertainty in the reconstructed evolution can be obtained with a few simple extra computational steps. Furthermore, by introducing linear regularisation, the evolution can be extrapolated into regions of the $L-z$ plane where data are lacking. We have also developed a procedure for adaptively pixellising the evolution function. This allows the evolution to be reconstructed with higher resolution in regions of the $L-z$ plane where data constraints are stronger and lower resolution where constraints are weaker. The procedure uses the data to automatically set the optimal pixellisation via maximisation of the Bayesian evidence. We have applied the method to a range of synthetic datasets constructed with two different input evolution functions; one rising monotonically with redshift and luminosity and the other incorporating a redshift cutoff. Comparing the reconstructed evolution with the input evolution provides a means of testing how well the method performs with varying dataset type, coverage and signal-to-noise. Our findings indicate that redshift measurements are essential to locate the presence of a cutoff in the evolution and that these must lie beyond the cutoff. Number count data alone allow for a surprisingly reasonable reconstruction if the evolution function is known to be monotonic, but falsely indicate monotonic behaviour when a redshift cutoff is present. We have made predictions of the degree to which forthcoming submm surveys carried out with new instruments (SCUBA2 and Herschel) will allow evolution in the submm GLF to be determined. These simulations show that combining a mixture of number count data from both facilities and including measurements of redshifts down to the source confusion limit will allow for a very well resolved and reliable measurement of evolution. In particular, the predictions have demonstrated that much improved constraints are provided by pure number count data measured at a combination of different wavelengths, compared to counts measured at just one wavelength. This is a result of the cross-linking of flux bins on the $L-z$ plane due to the conversion between monochromatic and bolometric luminosity which is a function of both redshift and wavelength. In terms of applying the method to real data, there are a few additional complexities that must be considered. One complexity is that real data are typically incomplete. This will result in the reconstructed evolution being a lower limit on the real evolution. Indeed, our simulations indicate that, due to regularisation, the reconstructed evolution tends to be underestimated even when the data are complete, although by an amount which is within the derived uncertainties. However, depending on the magnitude of the incompleteness, this may give rise to an underestimate that is not within the derived evolution error budget. In addition, depending on the level of regularisation, severe incompleteness in one dataset could significantly affect the reconstructed evolution in other areas of the $L-z$ plane covered by near-complete data. Another consideration that must be made with real data is regarding the SED used for converting between monochromatic and bolometric luminosity. Clearly, if this is a poor match to the SEDs of the galaxies in a given dataset, then the assumed location of that dataset in the $L-z$ plane will be inaccurate and give rise to a systematic error in the recovered evolution. In the submm, this will be larger for fluxes near the peak of the observer-frame SED, but on average, should be a relatively small effect. Of course, instead of using a single SED as we have done in our simulations, source-specific SEDs, or dataset-specific SEDs could be used in practice. In this paper, we have demonstrated the behaviour of the reconstruction method with a combination of different datasets and evolution functions. Although we have concentrated on key characteristics, there are many more that could have been tested. Since an exhaustive investigation of all configurations that might be encountered in practice is well beyond the scope of this work, a sensible approach would be to carry out simulations when applying the method to datasets that differ greatly from those tested here. In this way, any potential systematics that might arise from the data or any non-uniqueness in the reconstructed evolution could be quantified. There are also several ways in which the method could be developed. For example, the adaptive pixellisation scheme is relatively simple and could be enhanced to offer greater adaptability. Another possible enhancement might be improving the use of redshift information. In this work, we incorporated redshifts by crudely binning number counts by redshift. This lowers the constraining power of the redshift data. With this in mind, we have begun development of a more sophisticated scheme that includes redshift information on a more efficient source by source basis. This will be presented in forthcoming work. We are about to enter a new era in submm cosmology with the arrival of new instruments such as {\\sl Herschel} and SCUBA2. Surveys conducted with these instruments will give an increase in the number of detected submm sources of several orders of magnitude compared to existing surveys. Their significantly improved sensitivities will allow much wider ranges in luminosity and redshift to be explored. Combined with new methods to measure evolution, over the next few years, this will bring about a revolution in our overall understanding of how galaxies form and evolve. \\begin{flushleft} {\\bf Acknowledgements} \\end{flushleft} SD acknowledges support by the Science and Technologies Facilities Council. We would like to thank the reviewer of this paper, Professor Steve Phillipps, for his helpful comments and suggestions." }, "1003/1003.3980_arXiv.txt": { "abstract": "The Lagrange point $L_1$ for the Sun-Earth system is considered due to its special importance for the scientific community for the design of space missions. The location of the Lagrangian points with the trajectories and stability regions of $L_1$ are computed numerically for the initial conditions very close to the point. The influence of belt, effect of radiation pressure due to Sun and oblateness effect of second primary(finite body Earth) is presented for various values of parameters. The collinear point $L_1$ is asymptotically stable within a specific interval of time $t$ correspond to the values of parameters and initial conditions. ", "introduction": "\\label{Intro} The circular restricted three body problem is modification of the three body problem where the third body is assumed to have very small mass which is infinitesimal in comparison to other two finite masses are called primaries. The restricted three body problem is generalized to include radiation pressure, oblateness of the second primary and influence of the belt. Further the primary bodies are moving in circular orbits about their center of mass. The well-known five equilibrium points(Lagrangian points) that appear in the planar restricted three-body problem are very important for astronautical applications. The collinear points are unstable and the triangular points are stable \\citet{Szebehely1967}. In the Sun-Jupiter system several thousand asteroids, collectively referred to as Trojan asteroids, are in orbits of triangular equilibrium points. But collinear equilibrium points are also made linearly stable by continuous corrections of their orbits(\\lq\\lq halo orbits\\rq\\rq). In other words the collinear equilibrium points are metastable points in the sense that, like a ball sitting on top of a hill. However, in practice these Lagrange points have proven to be very useful indeed since a spacecraft can be made to execute a small orbit about one of these Lagrange points with a very small expenditure of energy \\citet{Farquhar1967JSpRo,Farquhar1969AsAer}. Because of the its unobstructed view of the Sun, the Sun-Earth $L_1$ is a good place to put instruments for doing solar science. NASA's Genesis Discovery Mission has been there, designed completely using invariant manifolds and other tools form dynamical systems theory. In 1972, the International Sun-Earth Explorer (ISEE) was established , joint project of NASA and the European Space Agency(ESA). The ISEE-3 was launched into a halo orbit around the Sun-Earth $L_1$ point in 1978, allowing it to collect data on solar wind conditions upstream from the Earth \\citet{Farquhar1985JAnSc}. In the mid-1980s the Solar and Heliospheric Observatory (SOHO) \\citet{Domingo1995SoPh} is places in a halo orbit around the Sun-Earth $L_1$ position, about a million miles the Sun ward from the Earth. They have provided useful places to \\lq\\lq park\\rq\\rq a spacecraft for observations. The Chermnykh's problem is a new kind of restricted three body problem which was first time studied by \\citet{Chermnykh1987}. This problem generalizes two classical problems of Celestial mechanics: the two fixed center problem and the restricted three body problem. This gives wide perspectives for applications of the problem in celestial mechanics and astronomy. The importance of the problem in astronomy has been addressed by \\citet{Jiang2004IJBC}. Some planetary systems are claimed to have discs of dust and they are regarded to be young analogues of the Kuiper Belt in our Solar System. If these discs are massive enough, they should play important roles in the origin of planets\\rq orbital elements. Since the belt of planetesimal often exists within a planetary system and provides the possible mechanism of orbital circularization, it is important to understand the solutions of dynamical systems with the planet-belt interaction. Chermnykh's problem has been studied by many scientists such as \\citet{Papadakis2005Ap&SS}, \\citet{Jiang2006Ap&SS,YehJiang2006Ap&SSII}, \\citet{Papadakis2007Ap&SS} and reference their in. The goal of present paper is to investigate the nature of collinear equilibrium point $L_1$ because of the interested point to the mission design. Although there are two new equilibrium points due to mass of the belt(larger than 0.15) \\citet{JiangYeh2006Ap&SSI,YehJiang2006Ap&SSII} but they are left to examine. All the results are computed numerically with the help of computer because pure analytical methods are not suitable. The actual trajectories and the stability regions of $L_1$ however is more complicated than the discussed here. But for specific the time intervals, and initial values, these results provide new information on the behavior of trajectories around the Lagrangian point $L_1$ for different possible set values of the parameters. ", "conclusions": "The numerical computation presented in the manuscript provides remarkable results to design trajectories of Lagrangian point $L_1$ which helps us to make comments on the stability(asymptotically) of the point. We obtained the intervals of the time where trajectory continuously moves around the $L_1$, does not deviate far from the point but tend to approach (for some cases) it, the energy of perturbed point is negative for these intervals, so we conclude that the point is asymptotical stable. More over we have seen that after the specific time intervals the trajectory of perturbed point depart from the neighborhood and goes away from it, in this case the energy also becomes positive, so the Lagrangian point $L_1$ is unstable. Further the trajectories and the stability regions are affected by the radiation pressure, the oblateness of the second primary and mass of the belt." }, "1003/1003.0451_arXiv.txt": { "abstract": "The detailed interior structure models of super-Earth planets show that there is degeneracy in the possible bulk compositions of a super-Earth at a given mass and radius, determined via radial velocity and transit measurements, respectively. In addition, the upper and lower envelopes in the mass--radius relationship, corresponding to pure ice planets and pure iron planets, respectively, are not astrophysically well motivated with regard to the physical processes involved in planet formation. Here we apply the results of numerical simulations of giant impacts to constrain the lower bound in the mass--radius diagram that could arise from collisional mantle stripping of differentiated rocky/iron planets. We provide a very conservative estimate for the minimum radius boundary for the entire mass range of large terrestrial planets. This envelope is a readily testable prediction for the population of planets to be discovered by the Kepler mission. ", "introduction": "Since the first confirmed detection of an extrasolar planet in 1995, the catalog of known exoplanets has grown to surpass 400. Precision Doppler shift discovery techniques have now identified more than a dozen exoplanets in the mass range from about 2 to 10--15 Earth masses ($M_{\\oplus}$). They are expected to be terrestrial in nature and unlike the gas giant planets in our solar system. This new class of planets has been collectively termed ``super-Earths'' \\citep{Melnick:2001}. Recently two transiting super-Earths were discovered, providing for the first time radii, in addition to masses. One of them---CoRoT-7b \\citep{Leger:2009, Queloz:2009}---has high density and is likely rocky. The other one---GJ1214b \\citep{Charbonneau:2009}---has low density and is likely water rich and surrounded by a small hydrogen-rich envelope. Many more super-Earths and measurements of their radii are expected from the Kepler mission \\citep{Kepler:2009}. Theorists anticipate a rich diversity in the bulk composition and internal structure of super-Earths, which is reflected in the broad band of possible radii on a planetary mass--radius diagram. The band corresponds to the anticipated range of mean densities. There are four distinct types of materials that could make up a planet in this regime: silicates, iron alloys, volatiles/ices, and hydrogen--helium gas. The range of possible mixing ratios between these materials leads to degeneracies in the determination of bulk composition from radius and mass alone \\citep{Sasselov:2008, Adams:2008}. As shown by \\cite{Valencia:2007b}, in order to restrict the range of possible bulk compositions, precise radii and masses (to 5\\% and 10\\%, respectively) have to be complemented with knowledge of stellar abundance ratios (e.g., Si/Fe) and physical constraints on the maximum fraction of H$_2$O or iron in a planet. The latter constraints place limits on the maximum and minimum possible radii for solid planets, respectively. In this Letter, we consider the minimum possible radius a super-Earth could have at a given mass. Since we are interested in the limiting case, our discussion can be confined to rocky planets composed of iron and silicates with no ices/water or hydrogen--helium gas layers. Under physically plausible conditions around normal stars, planet formation will lead to differentiated super-Earths with an iron core and a silicate mantle, with the proportions of each determined by the local Si/Fe ratio \\citep{Grasset:2009}. The only way to significantly increase the mean density requires removal of the silicate mantle while preserving the iron core. The most efficient method to strip the mantle is by giant impacts, which are common in the final stages of planet formation \\citep[e.g.,][]{Chambers:2004}. Given the large gravitational potential of super-Earth planets, we suspect that complete mantle stripping is not possible, and unlike the case of asteroids, pure iron super-Earths do not exist around normal stars. In this Letter, we analyze the results from numerical simulations of planet--planet collisions and determine a theoretical lower limit on the planetary radii of super-Earths. We anticipate that observations by the Kepler mission will test our predictions. The initial conditions for our investigation are dependent on an understanding of planet formation. These results and observations could, in turn, help constrain theories for planet formation. In recent years, \\cite{Ida_LinI} and \\cite{Mordasini:2009a} have applied detailed models of planet formation to the generation of synthetic populations. Such synthetic populations can then be compared to the observed distribution of known exoplanets to place statistically significant constraints on planet formation models. For example, the analysis of \\cite{Mordasini:2009a, Mordasini:2009b} shows that the core accretion model of planet formation can reproduce observed populations. In this model, small planetesimals collide to form larger planetary embryos ($\\sim0.1M_{\\oplus}$), the most massive of which come to dominate accretion in a process known as runaway growth. This stage is followed by an oligarchic stage in which protoplanets become relatively isolated after consuming the surrounding planetesimals. At this stage, collisions between comparably sized large bodies, giant impacts, become important and dominate the end stages of the formation of terrestrial planets. Giant impacts have been extensively modeled in the planetary embryo size regime \\citep{Agnor_Asphaug:2004, Asphaug:2009}; however they had not been studied extensively up to Earth size, with the exceptions of the Moon forming impact (e.g., \\citealt{Canup:2004}). Recently, the first study focused on collisions between super-Earths determined the criteria for catastrophic disruption and derived a scaling law for mantle stripping \\citep{Marcus:2009}. ", "conclusions": "Using the results of giant impact simulations presented in Marcus et al. (2009), we have described the impact conditions necessary to strip away mantle material from a nearly fully accreted planet. Combining this with detailed interior structure models for super-Earths \\citep{Valencia:2006,Valencia:2007b, Fortney:2007}, we have constructed a mass--radius diagram for super-Earths and shown that the previous lower envelope in this relationship, corresponding to 100\\% iron super-Earths, is not consistent with collisional mantle stripping. Further, if the absence of super-Earths in the 10-100$M_{\\oplus}$ range seen in the planet formation models of \\cite{Ida_LinI} is correct, even the existence of super-Mercuries, $\\sim$ 70\\% iron by mass, may be limited to masses $\\lesssim$ 5$M_{\\oplus}$ (as with the top blue curve in Figure \\ref{fig:ironfrac}(a)). This restriction arises because $\\sim10M_{\\oplus}$ target bodies are required to make super-Mercuries larger than about $5M_{\\oplus}$. The Kepler mission will discover a few hundred planets in the mass--radius range of super-Earths shown in Figure 4. We predict that the lower envelope of the distribution that Kepler is going to measure will be significantly higher than our computed minimum based on collisional stripping. Our calculation derives a very conservative limit with very low probability for these extreme scenario to be realized. Hence, Kepler's limited sample of planets is unlikely to be large enough to include such rare events. If super-Earths violating the collisional stripping limit are actually confirmed around normal stars, then we would need to revisit basic assumptions about the planet formation process. The simulations in this Letter were run on the Odyssey cluster supported by the Harvard FAS Research Computing Group. \\begin{figure} \\figurenum{1} \\begin{center} \\includegraphics[scale=0.6]{fig1a.ps} \\includegraphics[scale=0.6]{fig1b.ps} \\end{center} \\caption{Critical impact velocity required to obtain the specified post-impact iron mass fraction vs. mass of the largest impact remnant. (a) Largest remnant with iron mass fraction of 70\\% (similar to Mercury) for various projectile-to-target mass ratios. (b) Largest remnant with various iron mass fractions for equal mass collisions at the lowest velocity for such an enrichment. Note that the solid line in both panels is for a 1:1 mass ratio collision yielding a largest remnant that is 70\\% iron.} \\label{fig:vcrit} \\end{figure} \\clearpage \\begin{figure} \\figurenum{2} \\begin{center} \\includegraphics[scale=0.6]{fig2a.ps} \\includegraphics[scale=0.6]{fig2b.ps} \\end{center} \\caption{Iron fraction of the largest remnant vs. mass of the largest remnant. (a) Impact velocities given in terms of the mutual escape velocity. (b) Impact velocities in km~s$^{-1}$. In both panels, line color indicates the projectile-to-target mass ratio: green, 1/10; red, 1/4; blue, 1/1. } \\label{fig:ironfrac} \\end{figure} \\clearpage \\begin{figure} \\figurenum{3} \\begin{center} \\includegraphics[scale=0.6]{fig3a.ps} \\includegraphics[scale=0.6]{fig3b.ps} \\end{center} \\caption{Mass-radius diagram for super-Earths. (a) Impact velocities given in terms of the mutual escape velocity. (b) Impact velocities given in km~s$^{-1}$. In both panels, the two solid black lines represent terrestrial composition (iron mass fraction of 0.33, upper line) and pure iron (lower line) \\citep{Valencia:2007b}. The colored lines are for super-Earths stripped of mantle material in a giant impact. Line color indicates the projectile-to-target mass ratio: green, 1/10; red, 1/4; blue, 1/1. In (b), the $\\times$ symbols represent potentially observable transiting planets with 5\\% uncertainty in the radii. Note that the difference in minimum radii between a pure iron planet and plausible densities achieved via collisional stripping is observable.} \\label{fig:mass_radius} \\end{figure} \\clearpage \\begin{figure} \\figurenum{4} \\begin{center} \\includegraphics[scale=0.6]{fig4.ps} \\end{center} \\caption{Mass-radius diagram for super-Earths. The dotted lines are pure ice (upper) and pure iron (lower) \\citep{Fortney:2007}. The solid line is a constraint placed on the mass-radius diagram from collisional mantle stripping, corresponding to a 80~km~s$^{-1}$ impact with a 1:1 projectile-to-target mass ratio. The extrapolation (dashed) corresponds to cases that require target masses $>15M_{\\oplus}$. The $\\times$ symbols indicate the masses and radii of Earth, Uranus, and Neptune.} \\label{fig:mass_radius2} \\end{figure} \\clearpage" }, "1003/1003.0667_arXiv.txt": { "abstract": "A model for the evolution of low-luminosity radio galaxies is presented. In the model, the lobes inflated by low-power jets are assumed to expand in near pressure-balance against the external medium. Both cases of constant external pressure and decreasing external pressure are considered. Evolution of an individual source is described by the power-size track. The source appears as its lobe is inflated and radio luminosity increases to above the detection level; the source then moves along the track and eventually disappears as its luminosity drops below the detection limit. The power-size tracks are calculated including the combined energy losses due to synchrotron radiation, adiabatic expansion, and inverse Compton scattering. It is shown that in general, the constant-pressure model predicts an excess number of luminous, small-size sources while underpredicting large-size sources in the power-size diagram. The predicted spectra are steep for most sources, which is inconsistent with observations. By comparison, the pressure-limiting model fits observations better. In this model, low-luminosity sources undergo substantial expansion losses in the initial phase and as a result, it predicts fewer luminous, small-size sources. The resultant spectra are flat for most sources except for the oldest ones, which seems consistent with observations. The power-size tracks, in contrast to that of high-luminosity radio galaxies, are characterized by a slow increase in luminosity for most of the source's life, followed by a rapid decline when the synchrotron or inverse Compton scattering losses set in. ", "introduction": "The generic model for radio galaxies assumes that twin jets emanating from an active galactic nucleus propagate outward in two opposite directions. The jets, which initially propagate at a relativistic speed, interact with the surrounding medium leading to formation of a diffuse emission region. Radio galaxies appear to have two classes: low- and high-luminosity radio galaxies, commonly referred to as FR I and II sources, respectively \\citep{fr74}. The jets in high-luminosity radio galaxies have relatively homogeneous morphology; they are well collimated and propagate through the surrounding medium---initially in the cores, then halos of their parent galaxies and then the intergalactic medium (IGM)---creating pair of large lobes. The jets are dim until the end of the lobes where there are bright hot spots. Classical double radio sources are a typical example of this class. By contrast, low-luminosity radio galaxies are characterized by jets that are bright close to the nucleus of their parent galaxy. The jets have diverse morphologies, a feature that can be interpreted as deceleration of jets due to entrainment of the external medium. The jets are initially laminar near the nucleus and then subject to turbulent disruption when passing through the flare region that is thought to be the main acceleration site for relativistic particles. The jets beyond the flare region spread out, resembling smoke arising from a chimney mixing with the ambient medium. The key issues in the understanding of radio galaxies include the evolution of radio galaxies and the underlining physics that distinguishes these two classes. One suggestion is that these two classes of source are intrinsically different, primarily in their jet dynamics, evolving along different tracks \\citep{jw99}. However, there are suggestions that some of the high-luminosity radio sources with weak jets may evolve into low-luminosity sources \\citep{gw87,gw88,kb07}. Some radio sources exhibit mixed features of FR Is and IIs. For example, there are souces with one-side jet showing the FR I features and the other showing the FR II features. This leads to an opinion that such classification may not be clear cut as previously thought \\citep{kb07}. It is well accepted that the radio emission in radio galaxies is due to synchrotron radiation by relativistic electrons (or positrons) injected from the jets. The total synchrotron power $P_\\nu$ evolves with time as the injection of a mixture of kinetic energy and magnetic energy competes against the losses due to volume expansion and radiation. When the losses dominate, the total power is a decreasing function of time as the source ages. Since the typical evolutionary time scale is $\\sim 10^8\\,\\rm yr$, it is not practical to measure how the total power changes in time directly by observations. One may study the temporal evolution of radio galaxies from the total spectral power, $P_\\nu$, as a function of the source's linear size \\citep{s63}. The linear size here is defined as the dimension of the lobe along the jet axis. Since the linear size $D$ increases as the source expands, a radio source should evolve along a particular track in the $P_\\nu$--$D$ diagram. There are many discussions in the literature on the time evolution of high-luminosity radio galaxies (or FR II sources) \\citep{ketal97,betal99,mk02}. In the existing models, there are three relevant regions where the physical processes determine the evolution of the source. These include the hot spots where particles are assumed to be accelerated and radiate, the head region that contains the hot spots, and the lobe---an emission volume inflated by the input of the jet, where relativistic particles are injected and the volume of the emitting plasma expands. Since the hot spots appear at the end of the jet, their distance to the center can be identified as the linear size of the lobe. Thus, the modeling of the time evolution of the lobe size is reduced to the problem of modeling of changes in the location of the hot spots. One of the widely-discussed models is the self-similar expansion model in which the jet creates a bow shock by sweeping up the ambient material \\citep{fz86,f91,kf98}. So, the location of the shock is completely determined by the jet power and the density of the surrounding medium. This allows one to establish the source size as a function of time. In this paper we consider the evolution of low-luminosity radio galaxies (or FR Is). We derive the radio power as a function of the source's size, $P_\\nu(D)$, which can be directly compared with the power-size diagram inferred from observations. In contrast to FRIIs, there are few discussions on the evolution of the low-luminosity sources, mainly because of the lack of quantitative models that relate the linear sizes to the jet dynamics. One of the defining features of FR Is is that the flare region, the hot-spot equivalent as compared to high-luminosity sources, is located close to the nucleus, which indicates that the jet is decelerated to the subsonic flow regime relatively close by the nucleus. Since the jet continues to expand well beyond the flare region, forming a diffuse emission region further beyond, the size of a FR I source is not directly related to the location of the flare region and should be determined separately from the expansion of the diffuse region. Here we continue to refer to this diffuse emission region as the lobe as in FR IIs despite the significant difference between their morphologies. For low-power jets, the physical conditions of the surrounding medium play a critical role in defining the change in the lobe size. If the medium is warm, the expansion proceeds at near pressure balance against the ambient pressure and the volume increases more slowly with time than the self-similar expansion in the high-luminosity sources. As a result, the low-luminosity sources grow in size, by comparison, much more slowly than the high-luminosity sources. In Sec 2, we discuss a generic model for both high-power and low-power jets. The evolution of the emitting plasma in the lobes is considered in Sec 3, with application to high-power jets in Sec 4. The evolution of low-luminosity radio sources is discussed in Sec 5. ", "conclusions": "We consider the evolution of low-luminosity radio galaxies with $z\\ll1$. The radio power as a function of the source's size is derived based on a generic model that uses global parameters of the jet-lobe system. In the global model discussed here, we assume that the lobe inflated by a low-power jet undergoes pressure-limiting expansion, in which the system is approximately in pressure-balance against the external medium. Thus, the energy equation (Eq [\\ref{eq:El}]) for the lobe can be solved by replacing the lobe pressure, $p_l$, with the external pressure, $p_{\\rm ex}$. The source's size as a function of time can be derived in both cases of constant external pressure and decreasing external pressure. In the calculation of the $P_\\nu$--$D$ track we use the standard theory of nonstationary spectra of relativistic particles that are subject to both adiabatic and radiative losses. The predicted $P_\\nu(D)$ can be compared with the $P_\\nu$--$D$ diagram from observations. The main conclusions are summarised as follows. \\noindent 1) The pressure-limiting expansion model (with $\\beta\\sim 3/2$) predicts $P_\\nu$--$D$ tracks that are generally consistent with observations; in particular, it predicts fewer small-size sources in the high luminosity region in the radio $P_\\nu$-$D$ diagram. Since the emitting particles suffer expansion losses, the luminosity increases gradually untill the synchrotron or ICS losses become dominant. \\noindent 2) The constant-pressure expansion model (with $\\beta=0$) generally predicts an overabundance of small-size FR I radio sources with a relatively high radio power and at the same time underpredicts the large-size FR I sources. These features are not consistent with observations (cf. Figure~\\ref{fig:PD2}). \\noindent 3) In the pressure-limiting expansion model, the physical conditions of the external medium, e.g. $p_{\\rm ex}$, play an important role in the evolution. By assuming equipartition, one expects to see a direct link between the radio power and the external pressure, i.e. the higher external pressure the higher radio power. The $P_\\nu$--$D$ track would also turn over at a much smaller size. It is interesting to compare the model prediction for low--luminosity sources with that of high-luminosity sources. For low-power jets, the lobes expand slowly in near pressure equilibrium with the external pressure. Thus, FR Is evolve much more slowly compared with FR IIs. Our model predicts that the radio power of FR Is may increase throughout most of their life time. By contrast, the radio power of FR IIs declines in most of their life time (except for the brief initial rise) and can be well described by two power-laws. The initial rise for FR IIs is very brief, less than 0.1 Myr for the parameters adopted in Figure~\\ref{fig:PD1}, while for FR Is, this can last more than 100 Myr. These differences suggest that these two types of source evolve differently. In our global model, both the details of the particle injection and the effect of spatial diffusion are not treated. The effect of spatial diffusion can be important in particle transport--this is particularly the case when particle acceleration is confined to a localized site, say the flare region (`hot spot') \\citep{eetal97}. The accelerated particles need to diffuse across the region surrounding the acceleration site. Note that this requirement may be relaxed if particles are accelerated by multiple weak shocks distributed over an extended region in the lobe. Particle diffusion should strongly depend on plasma turbulence in the region concerned. To obtain the solution for particle diffusion, one needs to deal with the full diffusion-loss equation that includes particle diffusion in plasma turbulence \\citep{l94}. One should emphasise that the pressure in the energy equation (\\ref{eq:El}) is the total pressure that may consist of radiating and nonradiating particles as well as magnetic pressure. For convenience, we consider only one species of particles (radiating particles). X-ray observations suggest that in some of the known FR I sources, the equipartition pressure in the lobe is substantially lower than the external pressure \\citep{cetal08}. One possible explanation is that the missing pressure may be provided by nonradiating particles due to entrainment. Although in principle, the case can be treated in a similar way to that presented here for single species of particles by choosing a low $\\eta$ and $\\Gamma\\to 5/3$, An extension of this model to including multi-component plasmas in the lobe will be considered in future work. Finally we comment on the equipartition assumption adopted in the calculation of the radio power. The magnetic fields are assumed to be completely entangled (i.e. we ignores the effect of the mean field). The equipartition assumption is reasonable as recent {\\em Chandra} X-ray observations suggest that the magnetic fields in lobes are close to the equipartition field \\citep{betal08}. It is worth noting that the frozen flux argument would suggest that the magnetic energy density in an expanding lobe decreases with time much faster than $\\propto1/t^{\\beta/(3-\\beta)}$. To maintain equipartition, one requires either that the magnetic turbulence be generated in the lobe \\citep{d80} or that the magnetic energy be injected into the lobe \\citep{es89}." }, "1003/1003.3601_arXiv.txt": { "abstract": "{The B2Vn star HR\\,7355 is found to be a He-rich magnetic star. Spectropolarimetric data were obtained with FORS1 at UT2 on Paranal observatory to measure the disk-averaged longitudinal magnetic field at various phases of the presumed 0.52\\,d cycle. A variable magnetic field with strengths between $\\langle B_z \\rangle=-2200$ and $+3200$\\,G was found, with confidence limits of 100 to 130\\,G. The field topology is that of an oblique dipole, while the star itself is seen about equator-on. In the intensity spectra the {He}{\\sc i}-lines show the typical equivalent width variability of He-strong stars, usually attributed to surface abundance spots. The amplitudes of the equivalent width variability of the {He}{\\sc i} lines are extraordinarily strong compared to other cases. These results not only put HR\\,7355 unambiguously among the early-type magnetic stars, but confirm its outstanding nature: With $v\\sin i = 320$\\,km\\,s$^{-1}$ the parameter space in which He-strong stars are known to exist has doubled in terms of rotational velocity. } ", "introduction": "\\citet{2008A&A...482..255R} suggested the B2Vn star {HR\\,7355} (HD\\,182\\,180, HIP\\,95\\,408) to be a He-rich magnetic star with a magnetosphere containing trapped gas that produces hydrogen line emission (\\citealt{2005MNRAS.357..251T}, \\citealt*{2007MNRAS.382..139T}). They based their suggestion on evidence derived from two FEROS echelle spectra and a Hipparcos light-curve, with a period of either single-wave 0.26\\,d or double-wave 0.52\\,d. However, a period of 0.26\\, could not be due to rotational modulation, since this would require a rotational speed well above the critical threshold. If confirmed, HR\\,7355 would have a unique position in the He-strong class: With $v\\sin i = 320$\\,km\\,s$^{-1}$ \\citep*{2002ApJ...573..359A} it would be the the most rapidly rotating magnetic star in the upper HR-diagram; by a factor of about two ahead of the current record-holder, {$\\sigma$\\,Ori~E}, which has $v\\sin i = 165$\\,km\\,s$^{-1}$ (\\citeauthor{2002ApJ...573..359A}, op.~cit.). This would make the star a show-case for the Rigidly Rotating Magnetosphere model, that so far, though with great success, was applied to only one star, $\\sigma$\\,Ori\\,E (\\citeauthor{2005MNRAS.357..251T}, op.~cit.). In order to test this claim, a spectropolarimetric campaign was carried out in 2008. \\begin{figure} \\centering \\includegraphics[angle=270,width=0.95\\columnwidth,clip=]{StokesV_3100G_nobin.ps} \\includegraphics[angle=270,width=0.95\\columnwidth]{B54669_33.ps}% \\caption[]{Normalized circular polarization $V/I$ of HR\\,7355 at MJD=54669.33, for which a longitudinal magnetic field component of 3200\\,G was derived. Top: The spectra around the ten Balmer lines from H$\\beta$ to H$_{13}$. Bottom: The $V/I$ vs.\\ the bracketed part of the right term in Eqn.~\\ref{eqn_reg} and the linear regression to the data as a solid line. The typical error per measurement is shown as a bar; the error in the abscissa is less than the symbol size. The mass of points clustering around (0,0) are due to the unpolarized continuum.} \\label{fig_FORSV} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=0.95\\columnwidth,clip]{MC_chi2.ps}% \\includegraphics[width=0.95\\columnwidth,clip]{MC_probdis.ps}% \\caption[]{Minimum of the two-dimensional probability density disctribution computed from the artificial test data (see Sect.~\\ref{sec_mcbs}) for the observed data shown in Fig.~\\ref{fig_FORSV}, and an overplot of the $\\chi^2$-contrours enclosing 68.3\\%, 95.4\\% and 99.7\\% probabilities (upper panel) and one dimensional probability distribution of the derived field strength (lower panel). } \\label{fig_FORSB} \\end{figure} ", "conclusions": "\\subsection{Ephemeris} {As we know from $\\sigma$\\,Ori~E \\citep{2000A&A...363..585R}, the phase most easy to pin down, both spectroscopically and photometrically, is the center of one of the two circumstellar clouds crossing in front of the star. We chose this point also for HR\\,7355 as $\\phi=0$. A corresponding epoch $T_0$ can then be identified as the time of strongest absorption in H$\\alpha$, i.e.\\ the maximal equivalent width. While our FORS data do not cover H$\\alpha$, and the FEROS data were not taken in this phase, a spectroscopic campaign with the echelle instrument UVES at Paranal that has been completed in the second half of 2009, suggests MJD$=54\\,940.33$ as epoch. However, since a full analysis of the UVES data is far beyond the scope of this discovery report, and the epoch is not as critical as the period for our purpose, we leave a detailed analysis and discussion of the H$\\alpha$ equivalent width curve to a later work (Rivinius et al, in prep.), and here just adopt the epoch. Nevertheless it is reassuring that additional absorption in H$\\beta$ in the FORS1 data is observed at $\\phi=0.006$ and $0.013$, when the data are sorted with the period derived below. These two occurrences, 25 cycles apart from each other and more than 500 cycles apart from the selected epoch, are in full agreement with this epoch derived from UVES data. For the period, \\citet{2008A&A...482..255R} gave a value of 0.521428(6)\\,d for the Hipparcos data, under the condition that the variations were double-wave sinsoidal. The single wave period, which in any case would have been to short to be rotational, is firmly excluded by the new magnetic measurements. { However, also the 0.52\\,d period is not able to satisfactorily phase all available data, and so its value had to be improved. In order to do so, we demand a period to sort all three data types, i.e.\\ the photometric data (1990-1993, double wave), the magnetic data (2008, single wave), and the equivalent widths (1999-2008, possibly double wave). With respect to the originally published value, the closest period soring all three data-types is $0.521442(4)$\\,d. Although the EW data suffers strong seasonal aliasing, already its most nearby alias is excluded by the almost completely scrambled photometric curve for this period. The most nearby period value for which EW and photometry could be reconciled is incompatible with the magnetic phase-curve. We thus conclude that the true rotational period of HR\\,7355 must be within the above value's uncertainties.} In a final step, we can assume that the photometric minima do have a certain phase relation to the spectroscopic curve. {There are two possiblities: First, the minima could be due to cloud eclipses in front of the star, then we can further require that one of the photometric minima occurs at phase $\\phi=0$. Second, the photometric might be due to photospheric flux modulation in the He-enriched parts \\citep{2007A&A...470.1089K,mikulasek}. In this case, the photometric minima would coincide with minimal {He}{\\sc i} equivalent width. However, in this particular star, the maximal $H\\alpha$ equivalent width and the minimal {He}{\\sc i} equivalent width are almost simultaneous, so that a final decision about this can only be made with new photometry more simultanous with recent observations.} in any case, already with the period derived from the equivalent widths alone one of the two photometric minima is very close to $\\phi=0$, so that under the assumption that the photometric variability is due to the circumstellar material obscuring the line of sight the period becomes $P=0.521444(3)$\\,d. We will use the latter value for the discussion, but note that this relies on the assumption that the photometric minima are due to eclipses, while the best period without this assumption is $P=0.521442(4)$\\,d The ephemeris used thus is \\begin{equation}\\label{eqn_eph} T_{\\rm {H}{\\rm I}~max.~absorption}({\\rm MJD}) = 54940.33 + 0.521444(3) \\times E \\end{equation} \\subsection{Periodic variations}\\label{subsec_per} With the above choice of epoch we expect to see photometric minima, as well as $B_z=0$, at phase $\\phi=0$, i.e.\\ when the magnetic equator is facing towards us (Fig.~\\ref{fig_phases}). Then using the above ephemeris, and assuming a sinusoidal variation of the magnetic field, we estimate the field curve as \\begin{equation}\\label{eqn_field} \\langle B_z \\rangle(t) = 350\\,{\\rm G} + 2850\\,{\\rm G} \\times \\sin\\left(2\\pi \\left(\\frac{ t \\rm -T_0}{P} - 0.02\\right)\\right) \\end{equation} where $t$ is the date, and $T_0$ and $P$ are epoch and period from Eqn.~\\ref{eqn_eph}, and the shift of 0.02 in phase is required to fulfill the above condition $B_{z,\\phi=0}=0\\,$G, due to the constant term of 350\\,G. \\begin{figure} \\centering \\includegraphics[viewport=0 215 543 745,width=\\columnwidth,angle=0.0,clip=]{phaseplot.ps} \\caption[]{Measured quantities (see Table~\\ref{tab_result}, also for typical errors) as function of the rotational phase (Eqn~\\ref{eqn_eph}). The Hipparcos double wave light-curve is shown in the uppermost panel, the {He}{\\sc i}\\,4388 and 4713 equivalent widths below as triangles. The $W_\\lambda$ measured in the FEROS spectra \\citep{2008A&A...482..255R} are shown as red (1999 data) and blue (2004 data) squares. The lowermost panels shows the $\\langle B_z \\rangle$ field measurements and the curve defined by Eqn.~\\ref{eqn_field}. } \\label{fig_phases} \\end{figure} The fact that $B_{\\rm max}\\approx - B_{\\rm min}$ confirms the assumption of $i\\approx 90^\\circ$ by \\citet{2008A&A...482..255R}, while the term of 350\\,G is easily explained either by a slight offset from this value, or by an off-center magnetic dipole. Such an off-center dipole would cause a non-sinusoidal field curve. {In fact, due to two datapoints being taken at identical phases, there are effectively only five points to constrain the three free parameters of Eqn.~\\ref{eqn_field}. The field curve, although it seems to fit very well, is thus not well constrained, which is why we cannot really give confidence limits for the parameters unless further data has been obtained. In particular, we stress that the exactly sinusoidal shape of the curve is rather an assumption than an actually observed property.}} When the magnetic poles face the observer, $\\langle B_z \\rangle$ becomes maximal. Although there are only few points, not sampling the curve in all detail and in particualr not necessarily the respective maxima and minima, it is clear that the {He}{\\sc i} absorption is much stronger in these phases than when the magnetic equator is visible (Fig.~\\ref{fig_phases}). This is in full agreement with the behavior observed in other He-strong stars, like $\\sigma$\\,Ori\\,E \\citep{2000A&A...363..585R}. {The amplitude of the {He}{\\sc i} EW variations is considerably larger than in $\\sigma$\\,Ori\\,E, however. There the maximal EW is only about a factor of 1.3 to 1.5 stronger than the minimal one, depending on the spectral line. In HR\\,7355 the lines strengthen, wrt.\\ their minima, by a factor of 2 for lines like {He}{\\sc i}4713, and even a factor of 3 for strong lines with significant broadening wings like {He}{\\sc i}4388. This strong modulation is indicative for two large Helium enhanced patches on the surface close the equator at opposite longitudes, which point to a large angle $\\beta$ between the rotational and magnetic axes.}" }, "1003/1003.1118_arXiv.txt": { "abstract": "We review the out-of-equilibrium properties of a self-gravitating gas of particles in the presence of a strong friction and a random force (canonical gas). We assume a bare diffusion coefficient of the form $D(\\rho)=T\\rho^{1/n}$, where $\\rho$ is the local particle density, so that the equation of state is $P(\\rho)=D(\\rho)\\rho$. Depending on the spatial dimension $d$, the index $n$, the temperature $T$, and whether the system is confined to a finite box or not, the system can reach an equilibrium state, collapse or evaporate. This article focuses on the latter cases, presenting a complete dynamical phase diagram of the system. ", "introduction": "There is currently a renewed interest in the study of long-range interacting systems, particularly outside the realm of astrophysics. These systems may exhibit inequivalence of statistical ensembles (\\textit{e.g.} canonical \\textit{vs} microcanonical) which affects their equilibrium properties (possibility of negative specific heat in the microcanonical ensemble) and even more dramatically, their dynamical properties. In the context of self-gravitating systems, the Newtonian dynamics is too complicated to permit an exhaustive analytical treatment. Hence, in the past few years, we have developed a model of self-gravitating particles in the presence of a strong friction and a random force \\cite{col1}, for which inertial effects are negligible. The main interest of this model is to be analytically tractable in many situations, while presenting dynamical phases reminiscent of their Newtonian counterparts. In addition, this model is intimately related to the Keller-Segel model of bacterial chemotaxis \\cite{col2}. We thus consider particles obeying the equations of motion $\\frac{d{\\bf x}_i}{dt}=-\\nabla \\Phi+\\sqrt{2D}{\\bf \\eta}_i,$ where ${\\bf \\eta}_i$ is a delta correlated random Gaussian force, and $\\Phi$ is the gravitational potential. In a proper mean-field limit, which becomes exact for an infinite number of particles, the density obeys a Fokker-Planck (or Smoluchowski) equation coupled to the Poisson equation (the gravitational constant is set equal to $G=1$): \\begin{equation} {\\frac{\\partial\\rho}{\\partial t}}=\\nabla \\cdot (\\nabla P+\\rho\\nabla\\Phi),\\quad \\Delta\\Phi=\\rho. \\label{dim1} \\end{equation} The pressure is related to the diffusion coefficient by $P(\\rho)=D(\\rho)\\rho$, and the isothermal case corresponds to $D=T$, while we will consider here the more general polytropic case $D=T\\rho^{1/n}$ \\cite{col3}. In addition, the system of total mass $M$ can be placed in a spherical bounded domain of radius $R$ or in an unbounded space. In the context of chemotaxis \\cite{col2}, $\\rho$ is the density of bacteria, and $c=-\\Phi$ is the density of a chemical that they secrete. The bacteria are attracted by the regions of high density of this chemical which generates an effective long-range interaction between them. This interaction exactly takes the form of gravity when neglecting the diffusion and the degradation of the chemical. Although we shall see that the actual phase diagram depends crucially on the value of the index $n$, the general physics is controlled by the parameter $T$ which we assimilate to the temperature. When $T$ is small, the kinetic pressure is not strong enough to compensate gravity and the system may collapse. For large $T$, the system is at equilibrium in a bounded domain or evaporates in an unbounded domain. In the next section, we present the main results that we have obtained in \\cite{col1,col2,col3,col4,col5,col6,virial,col7} concerning these different dynamical phases. ", "conclusions": "In Table~\\ref{Table1}, we summarize the static and dynamic phase diagram of a self-gravitating gas of particles with a bare diffusion coefficient $D=T\\rho^{1/n}$, where $T$ is the temperature and $d$ the spatial dimension. This table illustrates the crucial role played by the critical index $n_*=\\frac{d}{d-2}$. \\begin{table} \\centering \\tbl{Static and dynamic phase diagram of a self-gravitating gas.\\hfill{ }} { \\tiny \\begin{tabular}{@{}|c|c|c|c|@{}} \\Hline & & & \\\\ \\textbf{{Index $n$}} & \\textbf{{Temperature}} & \\textbf{{Bounded domain}} & \\textbf{{Unbounded domain}} \\\\ & & & \\\\ \\Hline & & Metastable equilibrium state & $\\bullet$ Evaporation : \\\\ & $T>T_c$ & (local minimum of free energy): & asymptotically free normal \\\\ $n=\\infty$ ($d>2$) & & box-confined isothermal sphere & diffusion (gravity negligible)\\\\ \\cline{2-3} & & Self-similar collapse with $\\alpha=2$ & $\\bullet$ Collapse: \\\\ & $TT_c$ & Equilibrium state: & Equilibrium state: \\\\ $0T_c$ & (local minimum of free energy): & asymptotically free anomalous \\\\ $n_*T_c$ & Equilibrium state: & Self-similar evaporation \\\\ & & box-confined (incomplete) polytrope & modified by self-gravity \\\\ \\cline{2-4} & & Pseudo self-similar collapse & \\\\ $n=n_*$ & & leading to a Dirac peak of & Collapse \\\\ & $T9.2$ kpc \\citep{dickey}, 8 kpc \\citep{predehl1} and $9^{+4}_{-2}$ \\citep{predehl2}. Moreover, it is quite possible that the accretion disc is enshrouded partially or wholly in the strong stellar wind of the WR companion (\\citealt{fender1}; \\citealt{szostek2}; \\citealt{vilhu2}). The nature of the compact object is not certain, but it is thought to be a black hole due to its spectral resemblance to other black hole XRB systems, such as GRS 1915$+$105 and XTE J1550$-$564 (see e.g. \\citealt{szostek3}, hereafter S08; \\citealt{hjalmarsdotter2}). Also there is no evidence of a neutron star system producing such massive radio outbursts (up to 20 Jy, \\citealt{waltman2}) as observed from Cyg X-3. Cyg X-3 is a unique system in our Galaxy, but recent observations of two XRBs in IC 10 \\citep{prestwich} and NGC 300 \\citep{carpano}, both containing a WR companion, show strong evidence of a black hole primary. It is worth noting that these systems may represent a crucial link towards the evolution of a double black hole binary. The X-ray spectra of Cyg X-3 are notoriously complex. Basically, Cyg X-3 exhibits the canonical X-ray states seen in other XRBs, namely the high/soft (HS) and low/hard (LH) states, in addition to the intermediate, very high and ultrasoft states (e.g. \\citealt{szostek1}, hereafter S04; \\citealt{hjalmarsdotter2}). However, in Cyg X-3, the strong radio emission is also classified into states of its own (\\citealt{waltman3,mccollough1}, hereafter M99). Recently these X-ray and radio states were implemented into a more unified picture as presented in S08: the radio/X-ray states. The X-ray emission has been found to be linked to radio emission in Cyg X-3. M99 found that during periods of flaring activity in the radio the hard X-ray (HXR) flux switches from an anti-correlation to a correlation with the radio. In addition, the HXR flux has been shown to anti-correlate with the soft X-rays (SXR) in both canonical X-ray states \\citep{mccollough}. In February 1997 a large radio flare ($\\sim$ 10 Jy) was observed by the Green Bank Interferometer (GBI) and the Ryle Telescope after a period of quenched emission. At the same time, the \\batsefan\\/ onboard the \\cgrofan\\/ detected a flare in the HXR that showed a strong correlation with the radio. The flare in the HXR was preceded by several days of very low HXR flux (below the \\cgrobatse\\/ one-day detection limit). The flare triggered a Very Long Baseline Array (VLBA) observation to obtain high resolution radio images of Cyg X-3 during the major flare. The resulting VLBA observations show an expanding one-sided jet \\citep{mio} that was found to have a velocity of $\\sim$ 0.81$c$ and an inclination of $\\sim$ 14 degrees to our line of sight. An important tool in understanding the nature of the transient black hole systems is the hardness-intensity diagram (HID, e.g. \\citealt{fender2} and references therein). Thus it should be natural to look at Cyg X-3's HID and compare it to other black hole systems, all the while bearing in mind that Cyg X-3 does not behave like a transient black hole XRB in outburst. Nevertheless, we find the HID to be a useful tool even in this case. \\citet{smit} have taken a first look at Cyg X-3's behaviour in a HID using EXOSAT data, where they identified two distinct branches with a possible third branch connecting these two. In this paper we further classify the different radio/X-ray states by constructing a HID for Cyg X-3 from the X-ray data and adding the radio dimension in order to shed more light on the disc/jet connection in the system. This is made possible by the large amount of data available from the \\rxtefan\\/ (\\textit{RXTE}) archive (totaling to $\\sim$ Msec) and simultaneous radio observations from the GBI, RATAN-600 and Ryle telescopes (their monitoring programs of microquasars). We describe the X-ray and radio data in Section 2. In Section 3 we summarize the disc-jet connection of Cyg X-3 and the HID of black hole systems. In Section 4 we summarize and refine the radio/X-ray states, in addition to presenting the HID including radio observations. We discuss the ramifications of our results in Section 5 and summarize the paper in Section 6. ", "conclusions": "As mentioned above, Cyg X-3's HID is remarkably similar \\textit{in shape} to that of transient black hole XRBs even though it is not strictly speaking a transient source. Despite the similarity, however, Cyg X-3 does not cycle through the HID like other black hole XRBs. Rather, it traces a path from right to left (indicating spectral softening) or then from left to right (spectral hardening), with changes in intensity within each state. By looking at Fig. \\ref{asmhid} it appears that the movement across the HID is continuous with no discernible gaps in between. This result differs from that obtained by \\citet{smit}, where they found two branches and possibly a third one connecting these two in the HID. However, the discrepancy could arise from the lack of observations or the slightly different hardness used in their work (5.5--19.0 keV/3.5-5.5 keV). From the spectra it appears that there are three major components in play: a thermal disc, a Comptonized thermal component and a Comptonized non-thermal component. When the source moves from right to left the contribution of the thermal disc increases. The changes in this component are simply mapping the accretion on to the compact object. At the same time the HXR component decreases, but there are two effects contributing to this change: the overall HXR flux is dropping and the X-ray spectrum is changing from one that is dominated by thermal Comptonization to one that is dominated by non-thermal Comptonization. The fact that the radio flux density increases when the source enters the flaring region and, at the same time, the HXR switches from an anti-correlation to a correlation during major flares, indicates a relationship between the radio and the non-thermal Comptonization. To further investigate the thermal/non-thermal Comptonization shift in the source, we have fitted the pointed \\rxte\\/ spectra above 20 keV with a power law model. The resulting photon index ($\\Gamma$) and reduced chi-square values ($\\chi_{red}^2$) are plotted in Fig. \\ref{curvature}. A larger chi-square value indicates more curvature in the spectra and therefore more thermal Comptonization. It is apparent from the figure that the thermal HXR component is present in the quiescent states and the HXR component is more non-thermal in the flaring states. Also, the photon index changes from $\\sim$ 3 in the quiescent states to decreasing steadily from 3 to 1.5 in the flaring states. \\begin{figure} \\begin{center} \\includegraphics[width=0.5\\textwidth]{fig6.eps} \\caption{Photon indices (top panel) and reduced chi-square values (bottom panel) from fits to pointed \\rxte\\/ spectra above 20 keV for different SXR hardnesses. The colours are as in Fig. \\ref{plotall}. } \\label{curvature} \\end{center} \\end{figure} \\subsection{Case Study: A Major Radio Flaring Event April-May 2000} In Fig. \\ref{flare}, a major radio flaring period consisting of two major flares is plotted together with the \\rxte\\/ pointings (P50062). This same event is part of the state transition curve in Fig. \\ref{hidroutes}, which shows the evolution of Cyg X-3 in the HID from quiescent states to flaring states and back. This same period was also analyzed in S08 where they show the succession of the X-ray spectra and corresponding models (in their Fig. 8b). From Figs. \\ref{flare} and \\ref{hidroutes} we can follow the succession of the radio/X-ray states of Cyg X-3 including periods of major radio flaring. Also in this case Cyg X-3 differs from the transient black hole XRBs that trace their HIDs anti-clockwise. Instead, Cyg X-3 starts in the quiescent state and moves to the hypersoft branch and after a major radio flare enters the FHXR state via the FSXR/FIM states. This flaring episode can repeat as can be seen in Fig. \\ref{flare}. Eventually Cyg X-3 will return to the quiescent state. It is hard to pinpoint any particular route that Cyg X-3 follows, except that the different radio/X-ray states are followed in succession. Most likely the spread in the intensity comes from X-ray flaring and hysteresis does not occur in the system. From Fig. \\ref{hidroutes} we can zoom on the major radio flaring events as shown in Fig. \\ref{flare} and see that from the FIM/FSXR state after the first major flare Cyg X-3 enters the hypersoft state just before the second major flare. S08 discussed that during this time of the HXR tail or radio major flare the non-thermal electrons may be accelerated to high enough energies to result in detectable emission in $\\gamma$-rays in the GeV/TeV range. Recently, there have been observations of GeV emission before the onset of a major radio flare and when Cyg X-3 exhibited similar properties to the hypersoft state \\citep{tavani,corbel}, i.e. high SXR flux and very low or non-existent radio emission, indicating a major ejection of highly relativistic particles which will then most likely form the radio jet a few days later. In this particular case, the decaying of the first major flare leads to a higher radio flux density ($\\sim$ 0.1 Jy) that usually is observed together with the hypersoft state. However, it is likely that the system is quenching based on the X-ray state. From the hypersoft state Cyg X-3 enters the FSXR state, and after the peak of the second radio flare, it enters the FIM state. Throughout this progression, the power in the non-thermal HXR increases. During the decline of the second major flare Cyg X-3 enters the FHXR state and thermal Comptonization starts to dominate the HXR spectra. The minor flaring also present in the FHXR state spectra might arise in the aftermath of a major flaring event representing subsequent but smaller ejections. We note that during our case the two major flares present behave very differently with the first flare exhibiting little HXR emission and the second flare exhibiting significantly more (see the discrepancy of hardnesses in Fig. \\ref{hidroutes}). A possible mechanism for creating the difference is discussed in Section 5.2.1. Cyg X-3 also has the added distinction of exhibiting radio flares \\textit{only} when crossing the jet line from left to right (spectral hardening), rather than from right to left (spectral softening) like other black hole XRBs. Fig. \\ref{flare} shows that Cyg X-3 produces multiple major radio flares during major flaring periods and additional quenching periods between the flares. However, their detection without accompanying X-ray observations may be difficult because the decaying initial major radio flare dominates the radio emission. We also note that the power density spectra (PDS) of Cyg X-3 are not a good indicator of state and that on the short time-scales all the soft state PDS can be fit by a power law with index approximately --2 and that basically there is no power above $\\sim$ 1 Hz (e.g. \\citealt{axelsson}). This most likely comes from the scattering in the WR wind \\citep{zdziarski} and the --2 power law is characteristic of a random walk process likely produced by flares in the X-ray. \\begin{figure} \\begin{center} \\includegraphics[width=0.5\\textwidth]{fig7.eps} \\caption{The top panel shows the \\asm\\/ 3--5 keV lightcurve and the bottom panel GBI 8.3 GHz (black) and Ryle 15 GHz radio (grey) lightcurves with simultaneous \\rxte\\/ pointings (P50062) marked in the upper part of the panel, coloured and numbered as in Fig. \\ref{plotall}.} \\label{flare} \\end{center} \\end{figure} \\begin{figure} \\begin{center} \\includegraphics[width=0.5\\textwidth]{fig8.eps} \\caption{The top panel shows the \\asm\\/ 3--5 keV lightcurve and the bottom panel the evolution of Cyg X-3 in the HID (marked as a grey area) following the system from quiescent state through HYS and flaring states back to quiescence including the major radio flaring event discussed in the Section 5.1. Both panels are color-coded to match each other.} \\label{hidroutes} \\end{center} \\end{figure} \\subsubsection{Spectral Analysis} \\begin{table*} \\caption{Model parameters for the average spectra of hypersoft state (HYS) as well as the individual consecutive HYS (50062-02-03-02), FSXR (50062-02-07-00), FIM (50062-01-01-00) and FHXR (50062-01-03-00)} spectra. \\label{compps} \\begin{center} \\begin{tabular}{lllccccc} \\hline\\hline Component & Parameter & Unit & Avg. HYS & Ind. HYS & Ind. FSXR & Ind. FIM & Ind. FHXR \\\\ \\hline \\texttt{phabs} & $nH_{0}$ & $10^{22}$ & 2f & 2f & 2f & 2f & 2f \\\\ \\texttt{pcfabs} &$nH_{1}$ & $10^{22}$ & 6f & 6f & 6f & 6f & 6f \\\\ & cov & & 0.9f & 0.9f & 0.9f & 0.9f & 0.9f \\\\ \\texttt{edge} & $E_{Fe}^{H-like}$ & keV & $8.99^{+0.06}_{-0.06}$ & $8.82^{+0.06}_{-0.05}$ & $9.16^{+0.08}_{-0.07}$ & $9.18^{+0.09}_{-0.08}$ & $9.41^{+0.07}_{-0.05}$ \\\\ & $\\tau_{Fe}^{H-like}$ & & $0.35^{+0.02}_{-0.03}$ & $0.30^{+0.02}_{-0.02}$ & $0.25^{+0.02}_{-0.02}$ & $0.18^{+0.02}_{-0.02}$ & $0.26^{+0.02}_{-0.02}$ \\\\ & $E_{Fe}^{He-like}$ & keV & -- & -- & -- & -- & $7.94^{+0.09}_{-0.09}$ \\\\ & $\\tau_{Fe}^{He-like}$ & & -- & -- & -- & -- & $0.18^{+0.03}_{-0.02}$ \\\\ \\texttt{gauss[Fe]} & $E_{Fe}$ & keV & $6.33^{+0.06}_{-0.07}$ & $6.22^{+0.05}_{-0.06}$ & $6.51^{+0.03}_{-0.03}$ & $6.45^{+0.03}_{-0.03}$ & $6.57^{+0.03}_{-0.03}$ \\\\ & $\\sigma_{Fe}$ & keV & $0.50^{+0.09}_{-0.09}$ & $0.40^{+0.08}_{-0.08}$ & $0.48^{+0.05}_{-0.05}$ & $0.51^{+0.06}_{-0.05}$ & $0.26^{+0.08}_{-0.09}$ \\\\ & $EW_{Fe}$ & eV & $415^{+73}_{-56}$ & $268^{+44}_{-35}$ & $558^{+52}_{-45}$ & $468^{+45}_{-38}$ & $230^{+25}_{-23}$ \\\\ \\texttt{diskbb} & $kT_{bb}$ & keV & $1.52^{+0.03}_{-0.03}$ & 1.52f & 1.52f & -- & -- \\\\ & $R_{in}$ & km & $14.1^{+1.2}_{-1.3}$ & $25.9^{+1.1}_{-1.3}$ & $14^{+3}_{-5}$ & -- & -- \\\\ \\texttt{compPS} & $kT_{e}$ & keV & $12.8^{+0.5}_{-0.5}$ & 12.8f & $13.3^{+1.0}_{-1.3}$ & $25.2^{+0.7}_{-0.7}$ & $16.7^{+0.2}_{-0.2}$ \\\\ & $kT_{bb}$ & keV & $0.75^{+0.03}_{-0.03}$ & $0.88^{+0.03}_{-0.03}$ & $1.03^{+0.02}_{-0.02}$ & $1.09^{+0.01}_{-0.01}$ & $1.05^{+0.01}_{-0.01}$ \\\\ & $R_{in}$ & km & $74^{+7}_{-5}$ & $83^{+4}_{-2}$ & $64^{+5}_{-6}$ & $73^{+2}_{-1}$ & $74^{+1}_{-2}$ \\\\ & $\\tau$ & & $1.01^{+0.17}_{-0.15}$ & $0.75^{+0.06}_{-0.05}$ & $1.1^{+0.2}_{-0.1}$ & $0.93^{+0.06}_{-0.05}$ & $>2.89a$ \\\\ & $\\Gamma$ & & 2f & 2f & 2f & $4f$ & -- \\\\ \\hline \\texttt{Fluxes} & $F_{bol,abs}^{b}$ & $10^{-9}$ erg cm$^{-2}$ s$^{-1}$ & $9^{+2}_{-1}$ & $28^{+2}_{-2}$ & $22^{+5}_{-5}$ & $44^{+2}_{-2}$ & $42^{+1}_{-2}$ \\\\ & $F_{bol}^{c}$ & $10^{-9}$ erg cm$^{-2}$ s$^{-1}$ & $13^{+2}_{-2}$ & $40^{+3}_{-3}$ & $31^{+8}_{-7}$ & $55^{+3}_{-2}$ & $51^{+2}_{-2}$ \\\\ & $L_{bol}^{d}$ & $10^{38}$ erg s$^{-1}$ & $1.2^{+0.2}_{-0.2}$ & $3.9^{+0.3}_{-0.3}$ & $3.0^{+0.7}_{-0.7}$ & $5.3^{+0.2}_{-0.2}$ & $5.0^{+0.2}_{-0.2}$ \\\\ & $F_{comp}/F_{bol}$ & & 0.44 & 0.41 & 0.75 & 0.96 & 0.98 \\\\ & $F_{disk}/F_{bol}$ & & 0.51 & 0.55 & 0.19 & -- & -- \\\\ & $F_{line}/F_{bol}$ & & 0.05 & 0.04 & 0.06 & 0.04 & 0.02 \\\\ \\hline \\texttt{Confidence} & $\\chi^{2}_{red}$/d.o.f. & & 1.02/41 & 0.98/32 & 1.24/36 & 1.04/74 & 1.24/70 \\\\ \\hline \\end{tabular} \\end{center} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] Only the lower limit is shown because the optical depth is pegged to the highest value. \\item[$^{\\mathrm{b}}$] Absorbed, bolometric flux (3--200 keV) of the model normalized to the \\pca\\/ data. \\item[$^{\\mathrm{c}}$] Unabsorbed, bolometric flux (3--200 keV) of the model normalized to the \\pca\\/ data. \\item[$^{\\mathrm{d}}$] Unabsorbed, bolometric luminosity (3--200 keV) of the model assuming a distance 9 kpc to the source normalized to the \\pca\\/ data. \\item[$^{\\mathrm{f}}$] Frozen in the fits. \\end{list} \\end{table*} \\begin{figure} \\begin{center} \\includegraphics[width=0.5\\textwidth]{fig9.eps} \\caption{The individual and consecutive data of the hypersoft (red), FSXR (green), FIM (cyan) and FHXR (magenta) states with best fit models. See details in the text and parameters in Table \\ref{compps}.} \\label{hysmodel} \\end{center} \\end{figure} In the following we fit a model to the average hypersoft spectra and consequent individual HYS, FSXR, FIM and FHXR spectra during the second major flaring event. With the exception of the hypersoft state the other spectral states of S04 have been discussed in S04, \\citet{hjalmarsdotter1} and \\citet{hjalmarsdotter2}. In the following we will examine the hypersoft state and compare it to the flaring states. We fit the average hypersoft spectra as well as individual, consecutive hypersoft (50062-02-03-02), FSXR (50062-02-07-00), FIM (50062-01-01-00) and FHXR (50062-01-03-00) spectra during a major radio flare event, this being the only major flare event in the data set where there is \\rxte\\/ coverage and showing the succession of these states. The model includes neutral absorption and partially covered absorption, a gaussian line for the iron complex (the dominant line being the He-like iron line, but including contributions also from the H-like line, see e.g. \\citealt{paerels}) and ionized iron edges (contributions most likely from the H-like and He-like edges). We model the continuum with \\compps\\/, which is a hybrid Comptonization model \\citep{svensson} including multicolour disc black body emission as seed photons for the Comptonization process. Additionally, we include an unscattered disc component in the model, but the best fits appear to suppress this component from the FIM and FHXR spectra and therefore it is excluded from those. The absorption components are frozen to average values found by Koljonen et al. (in prep.) based on multiple \\swiftxrt\\/ spectra, as these provide better coverage below 3 keV as opposed to the \\rxtepca\\/. The absorption edges of highly ionized iron (H- and He-like) improve the fits substantially as is noted in \\citet{hjalmarsdotter1} and therefore have been included in the model. Cyg X-3 harbors a large number of emission lines that are especially prominent in the high-resolution energy spectra, e.g. \\citet{paerels}, but are also visible in the \\swiftxrt\\/ data. These lines arise in the stellar wind of the WR companion ionized by the intense X-ray emission from the compact object. However, due to the poor spectral resolution of the \\pca\\/, only the iron complex region is prominent in the spectra ($E_{Fe}\\sim6.5$ keV in the fits). For \\compps\\/, we use the following settings: the electron distribution function is hybrid with a low temperature Maxwellian plus a power-law tail (Lorenz factor extending from $\\gamma_{min}=1.3$ to $\\gamma_{max}=1000$) for the HYS, FSXR and FIM spectra. As evidenced in Fig. \\ref{curvature} the FHXR spectrum shows curvature and therefore we fit only the Maxwellian distribution to the FHXR spectrum. The electron power law index is not well constrained in the fits, so we have frozen that to $\\Gamma=2$ for the HYS and FSXR spectra, a typical value for soft states \\citep{hjalmarsdotter1}, and $\\Gamma=4$ for the FIM spectrum. The seed photons arise in a multicolour disc black body (\\diskbb\\/, \\citealt{mitsuda}) with inner disc temperature as a characterizing parameter. We also consider unscattered disc black body emission and therefore include also an extra disc component to the model. We do not include any reflection component in the model. The inclination of the accretion disc to the line-of-sight is uncertain and dependent on the nature of the compact object, and in the wind-fed binary system it does not even need to lie in the orbital plane \\citep{vilhu2}. However, we find that exchanging from sphere-approximated Comptonization to slab Comptonization does not improve the fits. To summarize, we have only the scattered and unscattered inner disc temperatures, the electron temperature and the optical depth of the electron cloud as free parameters in addition to the gaussian line and the iron edges. We also allow for an intercalibration factor between \\pca\\/ and \\hexte\\/ data correcting for calibration differences. The best-fitting parameters can be found in Table \\ref{compps} and we see that this model fit the spectra well in the HYS and FIM states ($\\chi^{2}_{red}$ $\\sim$ 1.0) and acceptably in the FSXR and FHXR states ($\\chi^{2}_{red}$ $\\sim$ 1.2). We find from the modelling that the HYS and flaring states can be fit with a model consisting of an absorbed disc black body and hybrid Comptonization including spectral features from iron emission lines and edges. From Fig. \\ref{hysmodel} and the fit results we see that for the HYS, FSXR and FIM states, non-thermal Comptonization is needed to account for the HXR tail, and that the FHXR can be fit with only thermal Comptonization. We note that this model is not adequate for fitting the quiescent state spectra and most likely reflection \\citet{szostek2} and/or Compton downscattering \\citet{zdziarski} takes place in these states. From the succession of the spectral states we notice that the soft seed photon temperature for the Comptonization remains at $kT_{bb}\\sim1$ keV and the optical depth of the plasma remains at $\\tau\\sim1$ except when moving into the FHXR state where it pegs to the upper limit. However, high plasma optical depths of $\\tau\\sim3-6$ are suggested to explain the hard spectra of Cyg X-3 in \\citet{hjalmarsdotter2} and \\citet{zdziarski}. We find the electron temperature to be rather soft, $kT_{e}\\sim13-16$ keV, except in the FIM state where the best-fitting value is found to be $\\sim25$ keV. The FIM state is also the state where the highest radio flux densities are found, so this might indicate that the radio jet has an effect on the Comptonizing electron temperatures. On the other hand, in our case study the FIM state occurs after a $\\sim$ 10 Jy radio flare and the FHXR state after a $\\sim$ 2 Jy flare (see Ryle data in Fig. \\ref{flare}), which might have an impact on the order of these states occurring after a major radio flare. One would expect the FHXR state to occur right after the major flare peak exhibiting higher optical depth and therefore coming from a more compact area, likely closer to the base of the jet. As the flare evolves and becomes more optically thin one would expect the system to enter the FIM state exhibiting a more power law-like tail. We also note that for the HYS and FSXR states an unscattered disc component is required to achieve acceptable fits. The average HYS spectrum yields an inner disc temperature $kT_{bb}=1.52$ keV and this value is frozen for the individual spectral fits in order to gauge the level of unscattered disc flux in each state. We find that this disc component accounts for 50\\% of the X-ray flux in the HYS state and drops to 20\\% in the FSXR state; for the FIM and FHXR states all the disc photons appear to be scattered. We note that locking the two black body temperatures to each other in the model does not lead to an acceptable fit. The discrepancy of the inner disc temperature and the soft seed photon temperature for the Comptonization component in the HYS and FSXR state may or may not be a real effect, though it is required for acceptable fits. It is suggested in \\citet{szostek2} that there is strong wind emission present below 1 keV in the system. This emission might be the source of the soft seed photons that are Comptonized and causing the scattered disk temperature of $\\lesssim$ 1 keV in the model throughout the states (the wind being present at all times in the system). The unscattered disk component in the HYS and FSXR states would represent the accretion disk either forming during these states and launched in the jet (as this component is not present in the FIM and FHXR states), or then remaining invisible to us due to the WR wind. During a major flare event the disc would become visible as the jet plows through the wind. Of course, only extensive SEDs will enable us to determine the actual contribution to the X-ray emission of the relativistic jets, especially during the major flaring periods. \\subsection{Different Emission Scenarios in Cyg X-3} In the following we discuss briefly the different methods of producing the HXR emission in Cyg X-3. However, more studies are needed for determining the exact process(es) giving rise to the HXR emission. \\subsubsection{Disrupted/Attenuated Jet} The HID in Fig. \\ref{hidradio} indicates that flares with more prominent HXR emission are brighter in the radio than those with less contribution from the HXR (although this might simply be a selection effect). This could be a manifestation of a disrupted or an attenuated jet (\\citealt{perucho}, for an actual observation of a disrupted jet in AGN see e.g. \\citealt{evans}), i.e. a jet with a less significant HXR component, being therefore less energetic, may not be capable of plowing through the stellar wind of the companion. This effect is clearly present in our case study (Section 5.1) showing a major flare with very low HXR emission followed by a major flare with significant HXR emission. For this to be plausible, then the HXR emission must originate in the jet itself via synchrotron or synchrotron self-Compton (SSC), see e.g. \\citet{markoff,bosch}. However, the \\asm\\/ HID appears to be continuous in the flaring region as well, so the jet in Cyg X-3 would then go from being disrupted to unimpeded. Also, if the jet composition were hadronic, there might be hadronic jet-wind interactions producing the high-energy $\\gamma$-ray emission and possibly also a significant amount of neutrinos \\citep{romero}. However, the variable and steep $\\gamma$-ray spectrum of Cyg X-3 suggests an electron origin, especially if the electrons were located away from the accretion disc. \\subsubsection{Comptonization processes} Typically, the HXR emission in XRBs and microquasars is modelled by the Compton upscattering of a soft photon source, i.e. an accretion disc or some other SXR plasma component, by high-energy electrons in a corona surrounding the compact object. The nature of this corona is much debated and it could be e.g. patchy high-energy regions above the accretion disc (e.g. \\citealt{zdziarski1}), the base of the relativistic jet (e.g. \\citealt{markoff}), or an ADAF-like hot accretion flow (e.g. \\citealt{esin}). In the case of microquasars the relativistic particles could be created in or near the accretion disc and then be ejected through the jet. The HXR emission could then arise from the Compton upscattering of the soft accretion disc photons by these relativistic particles and the radio emission from the synchrotron radiation by those ejected in the jet. However, because the system is embedded in the dense stellar wind of the companion WR star, it brings another level of complexity to the system. In addition to the accretion disc around the compact object the surrounding plasma also affects the HXR photons in several ways including absorption, reflection \\citep{hjalmarsdotter1} and Compton downscattering \\citep{zdziarski}. The latter would occur in a cold ($\\sim 3 \\, {\\rm keV}$) plasma cloud that would reduce the HXR cut-off to $\\sim20 \\, {\\rm keV}$ \\citep{hjalmarsdotter1} as observed from Cyg X-3 as opposed to the usual value of $\\sim100 \\, {\\rm keV}$ observed in other XRBs, especially describing the quiescent X-ray spectra of Cyg X-3. This effect would be most likely present also in the flaring states, however since the thermal HXR emission changes to non-thermal emission in the flaring states as discussed above the effect is not so prominent in the X-ray spectra compared to the quiescent states. \\subsubsection{Shock-in-jet} Particle acceleration scenarios in the shock along the jet are very viable and offer an alternative way of producing the relativistic electrons as opposed to particle creation in or near the accretion disc. The so-called shock-in-jet model has been invoked successfully to describe Cyg X-3's radio behaviour during major flaring periods \\citep{lindfors} and low-level activity \\citep{millerjones}. These shocks would provide the non-thermal power law emission and the link to the radio and HXR during flaring episodes. The relativistic particles would be created in the internal shocks in the jet, and the HXR emission could arise either from synchrotron emission from the electrons in the jet or Comptonization of the disc photons by the high-energy electrons in the shock, depending on the location of the shock and emission regions. In fact, Comptonization of the disc photons by this shock could be in the form of bulk-motion Comptonization (BMC, e.g. \\citealt{titarchuk}; however in that paper the BMC is related to the accretion near the compact object and not to the outflow), since the shock might move along the jet with a significant fraction of the speed of light. This non-thermal shock emission may very well also be the source of the recently observed very high energy $\\gamma$-ray emission. If the low-level radio activity can be modelled also with the shock-in-jet model, can we apply this to the quiescent X-ray spectra? In fact, the Cyg X-3 quiescent state X-ray spectra can be modelled with an almost purely non-thermal electron distribution which contributes to the non-thermal Comptonization of SXR photons \\citep{hjalmarsdotter1}. In this case the peak of the Comptonized spectrum is no longer determined by the maximum electron temperature, but rather the energy where the majority of electrons are injected explaining the unusually low cut-off in the Cyg X-3 quiescent X-ray spectra. In this way the shock-in-jet model would provide the necessary elements of modelling both the quiescent and flaring X-ray and radio spectra." }, "1003/1003.3571_arXiv.txt": { "abstract": "{}{We report on a detailed abundance analysis of the strongly r-process enhanced giant star newly discovered in the HERES project, HE~2327$-$5642 ([Fe/H] = --2.78, [r/Fe] = +0.99).} {Determination of stellar parameters and element abundances was based on analysis of high-quality VLT/UVES spectra. The surface gravity was calculated from the non-local thermodynamic equilibrium (NLTE) ionization balance between \\ion{Fe}{i} and \\ion{Fe}{ii}, and \\ion{Ca}{i} and \\ion{Ca}{ii}.} {Accurate abundances for a total of 40 elements and for 23 neutron-capture elements beyond Sr and up to Th were determined in {\\LyudmilasStar}. For every chemical species, the dispersion of the single line measurements around the mean does not exceed 0.11~dex. The heavy element abundance pattern of HE~2327$-$5642 is in excellent agreement with those previously derived for other strongly r-process enhanced stars such as CS\\,22892-052, CS\\,31082-001, and HE\\,1219-0312. Elements in the range from Ba to Hf match the scaled Solar $r$-process pattern very well. No firm conclusion can be drawn with respect to a relationship between the fisrt neutron-capture peak elements, Sr to Pd, in {\\LyudmilasStar} and the Solar $r$-process, due to the uncertainty of the latter. A clear distinction in Sr/Eu abundance ratios was found between the halo stars with different europium enhancement. The strongly r-process enhanced stars reveal a low Sr/Eu abundance ratio at [Sr/Eu] = $-0.92\\pm0.13$, while the stars with $0 <$ [Eu/Fe] $< 1$ and [Eu/Fe] $< 0$ have 0.36\\,dex and 0.93\\,dex larger Sr/Eu values, respectively. Radioactive dating for {\\LyudmilasStar} with the observed thorium and rare-earth element abundance pairs results in an average age of 13.3~Gyr, when based on the high-entropy wind calculations, and 5.9~Gyr, when using the Solar r-residuals. {\\LyudmilasStar} is suspected to be radial-velocity variable based on our high-resolution spectra, covering $\\sim 4.3$ years.} {} ", "introduction": "\\label{Sect:intro} The detailed chemical abundances of Galactic halo stars contain unique information on the history and nature of nucleosynthesis in our Galaxy. A number of observational and theoretical studies have established that in the early Galaxy the rapid ($r$) process of neutron captures was primarily responsible for formation of heavy elements beyond the iron group \\citep[we cite only the pioneering papers of][]{Spite,Truran81}. The onset of the slow ($s$) process of neutron captures occurred at later Galactic times (and higher metallicities) with the injection of nucleosynthetic material from long-lived low- and intermediate-mass stars into the interstellar medium \\citep[see][and references therein]{Travaglio99}. Since 1994, a few rare stars have been found that exhibit large enhancements of the $r$-process elements, as compared to Solar ratios, suggesting that their observed abundances are dominated by the influence of a single, or at most very few nucleosynthesis events. The $r$-process is associated with explosive conditions of massive-star core-collapse supernovae \\citep{Woosley1994}, although the astrophysical site(s) of the $r$-process has yet to be identified. Observations of the stars with strongly enhanced $r$-process elements have placed important constraints on the astrophysical site(s) of their synthesis. \\citet{sne1994} found that the extremely metal-poor ([Fe/H]\\footnote{In the classical notation, where [X/H] = $\\log(N_{\\rm X}/N_{\\rm H})_{star} - \\log(N_{\\rm X}/N_{\\rm H})_{Sun}$.} $\\sim -3.1$) giant CS\\,22892-052 is neutron-capture-rich, [Eu/Fe]$\\simeq +1.6$ (following the suggestion of \\citealt{Beers2005}, we hereafter refer to stars having $\\mathrm{[Eu/Fe]} > +1$ and $\\mathrm{[Ba/Eu]} < 0$ as r-II stars), and that the relative abundances of nine elements in the range from Ba to Dy are consistent with a scaled Solar system $r$-process abundance distribution. Later studies of CS\\,31082-001 \\citep{hill02}, BD+17$^\\circ$\\,3248 \\citep{cowan2002}, CS\\,22892-052 \\citep{sne03}, HD\\,221170 \\citep{ivans2006}, CS\\,22953-003 \\citep{francois2007}, HE\\,1219-0312 and CS\\,29491-069 \\citep{hayek09} provided strong evidence for a universal production ratio of the second $r$-process peak elements from Ba to Hf during the Galaxy history. CS\\,31082-001 \\citep{hill02} provided the first solid evidence that variations in progenitor mass, explosion energy, and/or other intrinsic and environmental factors may produce significantly different $r$-process yields in the actinide region ($Z \\ge 90$). The third $r$-process peak ($76 \\le Z \\le 83$) is so far not well constrained, because, in most r-II stars, it is only sampled by abundance measurements of two elements, Os and Ir. The abundances of platinum and gold were obtained for CS\\,22892-052 \\citep{sne03}. The only detection of lead in a r-II star so far is in CS31082-001 \\citep{plez04}. \\citet{sne03} reported an underabundance of elements in the range of $40 < Z < 56$ relative to the scaled Solar $r$-process, which prompted a discussion of multiple $r$-process sites \\citep[see, for example,][]{Travaglio04, QW08, far09}. The detection of the radioactive elements thorium and uranium provided new opportunities to derive the ages of the oldest stars and hence to determine a lower limit for the age of the Universe \\citep[see the pioneering papers of][]{sne1996,cayr01}. It appears that all the r-II stars with measured Th (and U) can be devided into two groups: (a) stars exhibiting an actinide boost (e.g., CS\\,31082-001, HE\\,1219-0312), and (b) stars with no obvious enhancement of thorium with respect to the scaled Solar r-process pattern (e.g., CS\\,22892-052, CS\\,29497-004; for a full list of stars, see \\citealt{Roederer2009}). For the actinide boost stars, ages cannot be derived when only a single radioactive element, either Th or U, is detected. To make clear an origin of the heavy elements beyond the iron group in the oldest stars of the Galaxy, more and better measurements of additional elements are required. Currently, there are about ten r-II stars reported in the literature \\citep{hill02,sne03,HERESpaperI,honda2004,HERESpaperII,francois2007,frebel2007,Lai2008,hayek09}. Abundance pattern for a broad range of nuclei, based on high-resolution spectroscopic studies, have been reported for only six of these stars. Continuing our series of papers on the Hamburg/ESO R-process-Enhanced Star survey (HERES), we aim at extending our knowledge about synthesis of heavy elements in the early Galaxy by means of a detailed abundance analysis of the strongly $r$-process enhanced star {\\LyudmilasStar}. In our study we also investigate the reliability of multiple Th/$X$ chronometers for {\\LyudmilasStar}, where $X$ is an element in the Ba--Hf range. {\\LyudmilasStar} was identified as a candidate metal-poor star in the Hamburg/ESO Survey (HES; see \\citet{Christliebetal:2008a} for details of the candidate selection procedures). Moderate-resolution ($\\Delta\\lambda = 2$\\,{\\AA}) spectroscopy obtained at the Siding Spring Observatory (SSO) 2.3\\,m-telescope with the Double Beam Spectrograph (DBS) confirmed its metal-poor nature. Therefore, it was included in the target list of the HERES project. A detailed description of the project and its aims can be found in \\citet[][hereafter Paper~I]{HERESpaperI}, and methods of automated abundance analysis of high-resolution ``snapshot'' spectra have been described in \\citet[][hereafter Paper~II]{HERESpaperII}. ``Snapshot'' spectra having $R\\sim 20,000$ and $S/N\\sim 50$ per pixel at 4100\\,{\\AA} revealed that {\\LyudmilasStar} exhibits strong overabundances of the $r$-process elements, with [Eu/Fe] = $+1.22$ and $\\mathrm{[Ba/Eu]}=-0.56$ (Paper\\,II). This paper is structured as follows. After describing the observations (Section~\\ref{Sect:Observations}), we report on the abundance analysis of {\\LyudmilasStar} in Sections~\\ref{Sect:AbundanceAnalysis} and \\ref{Sect:Results}, based on high-quality VLT/UVES spectra and MAFAGS model atmospheres \\citep{Fuhr1}. The heavy element abundance pattern of {\\LyudmilasStar} is discussed in Section~\\ref{Sect:ssr}. Section~\\ref{Sect:ages} reports on the radioactive decay age determination. Conclusions are given in Section~\\ref{Sect:DiscussionConclusions}. ", "conclusions": "\\label{Sect:DiscussionConclusions} The high-quality VLT/UVES spectra of {\\LyudmilasStar} enabled the determination of accurate abundances for 40 elements, including 23 elements in the nuclear charge range between $Z = 38$--$90$. We confirmed that {\\LyudmilasStar} is strongly $r$-process enhanced, having $[r/\\mathrm{Fe}] = 0.99\\pm0.03$ where $r$ denotes the average of the abundances of seven elements (i.e., Eu, Gd, Tb, Dy, Ho, Er, and Tm), with an $r$-process contribution to the Solar system matter of more than 85\\,\\% according to the $r$-residuals of \\citet{rs99}. {\\LyudmilasStar} and three benchmark r-II stars, CS\\,22892-052 \\citep{sne03}, CS\\,31082-001 \\citep{hill02}, and HE\\,1219-0312 \\citep{hayek09}, have very similar abundance patterns of the elements in the range from Sr to Os. Hence, {\\LyudmilasStar} is a member of the small sample of currently- known r-II stars. The elements in the range from Ba to Hf in {\\LyudmilasStar} match the scaled Solar $r$-process pattern very well. We showed that the Solar $r$-residuals for the first $r$-process peak elements are rather uncertain. They may vary as much as 0.5\\,dex or even more, depending on the adopted Solar total abundances and $s$-process fractions. Therefore, no firm conclusion can be drawn with respect to a relation of the light trans-iron elements in {\\LyudmilasStar} and other r-II stars to the Solar $r$-process. We found a clear distinction in Sr/Eu abundance ratios between the halo stars with different europium enhancement and suggest using the [Sr/Eu] ratio in addition to [Eu/Fe] to separate the strongly $r$-process enhanced (r-II) stars from the other halo stars with dominant contribution of the $r$-process to heavy element production. The r-II stars, whose stellar matter presumably has experienced a single nucleosynthesis event, have $\\mathrm{[Eu/Fe]} > 1\\pm 0.1$, $\\mathrm{[Ba/Eu]} = -0.60\\pm 0.10$, and a low Sr/Eu abundance ratio of $\\mathrm{[Sr/Eu]} = -0.92\\pm 0.13$. Stars with very similar Ba/Eu ratios have two times (0.36\\,dex) larger Sr/Eu ratios if their Eu/Fe ratio is in the range $1.0 < \\mathrm{[Eu/Fe]} < 0.2$ (i.e., r-I stars), and nearly an order of magnitude (0.93\\,dex) higher Sr/Eu ratios if $\\mathrm{[Eu/Fe]} < 0$ (Eu-poor stars). The origin of the first neutron-capture peak elements in the r-I stars and Eu-poor stars is still unclear. Further theoretical studies will be needed to elucidate this problem. Only two elements, Os and Ir, of the third $r$-process peak were detected in {\\LyudmilasStar}. Iridium appears to be overabundant compared to the Ir abundance determined in other $r$-process enhanced stars. However, due to the uncertainty of the Ir abundance we cannot yet draw a firm conclusion on this point. The detection of thorium permitted an estimate of the radioactive decay age of {\\LyudmilasStar}, although the age uncertainty of 9.3\\,Gyr introduced by the uncertainty of the thorium abundance is rather large. Employing multiple Th/$X$ chronometers and initial production ratios based on the Solar $r$-residuals, an age of $5.9\\pm 2.8$\\,Gyr was obtained from nine Th/$X$ pairs, involving elements in the Sm--Os range. Using the predictions of the HEW $r$-process model, as given by \\citet{hayek09}, we obtained $\\tau = 13.3\\pm 6.2$\\,Gyr from 12 Th/$X$ pairs. Based on our high-resolution spectra, covering $\\sim 4.3$ years, we suspect that {\\LyudmilasStar} is radial-velocity variable with a highly elliptical orbit of the system. Determination of the orbital period would provide the unique opportunity to determine a lower limit for the mass of the secondary in this system. Scenarios for the site of the $r-$process include a high-entropy wind from a type-II supernova \\citep[e.g.][]{Woosley1994,takahashi1994}, ejecta from neutron star mergers \\citep[e.g.][]{freiburg1999}, or the neutrino-driven wind of a neutron star newly-formed in an accretion induced collapse (AIC) event \\citep[e.g.][]{woosley1992,qian2003}. According to these scenarios, it is expected that the secondary is a neutron-star. With a lower limit for the mass of the secondary it might be possible to constrain a scenario, because in the AIC case the neutron star is expected to have a mass just slightly above the Chandrasekhar mass, while core-collapse supernovae or neutron star mergers would result in remants of significantly higher mass." }, "1003/1003.4167_arXiv.txt": { "abstract": "{The rapidly increasing volume of asteroseismic observations on solar-type stars has revealed a need for automated analysis tools. The reason for this is not only that individual analyses of single stars are rather time consuming, but more importantly that these large volumes of observations open the possibility to do population studies on large samples of stars and such population studies demand a consistent analysis. By consistent analysis we understand an analysis that can be performed without the need to make any subjective choices on e.g. mode identification and an analysis where the uncertainties are calculated in a consistent way.\\\\\\\\ Here we present a set of automated asterosesimic analysis tools. The main engine of these set of tools is an algorithm for modelling the autocovariance spectra of the stellar acoustic spectra allowing us to measure not only the frequency of maximum power and the large frequency separation, but also the small frequency separation and potentially the mean rotational rate and the inclination. \\\\\\\\ The measured large and small frequency separations and the frequency of maximum power are used as input to an algorithm that estimates fundamental stellar parameters such as mass, radius, luminosity, effective temperature, surface gravity and age based on grid modeling. \\\\\\\\ All the tools take into account the window function of the observations which means that they work equally well for space-based photometry observations from e.g. the NASA {\\it Kepler} satellite and ground-based velocity observations from e.g. the ESO {\\it HARPS} spectrograph. } ", "introduction": "Asteroseismology is proving itself as a powerful tool to increase our understanding of stars. This progress is down due to both ground-based facilities such as {\\it Elodie} (Baranne et al. 1996), {\\it CORALIE} (Queloz et al. 2000), {\\it HARPS} (Mayor et al. 2003), {\\it UCLES} (Tinney et al.2001) and {\\it UVES} (Dekker et al. 2000) and space-based facilities such as {\\it WIRE} (Buzasi et al. 2000), {\\it MOST} (Walker et al. 2003), {\\it COROT} (Appourchaux et al. 2008) and {\\it Kepler} (Chaplin et al. 2010). And we do not expect the fun to stop here. For the future we have high expectations regarding the ground-based telescope network {\\it SONG} (Grundahl et al. 2009) and the ESA satellite {\\it PLATO} (Catala 2008). Solar-type stars are, due to their rich and structural oscillations spectra some of the best targets for asteroseismology. From the characteristic large frequency separation between oscillation modes with different radial order we can get a direct measurement of the (acoustic) radius of the stars (Christensen-Dalsgaard et al. 2007) and by combining this measurement with the characteristic small frequency separation between modes with different angular degree we can get estimates of the mass and the age of the stars (Christensen-Dalsgaard et al. 2007). The acoustic background that we measure in the stars provides us with interesting information about activity and convection (Karoff 2008). For the best observations it is possible to measure the frequencies of the individual oscillation modes which provide us with an inexhaustible source of information on stellar structure and evolution -- including: information on energy generation and transport, rotation and stellar cycles (Karoff et al. 2009). ", "conclusions": "" }, "1003/1003.2577_arXiv.txt": { "abstract": "{} {We study the correlations between the VLBA radio emission at 15 GHz, extended emission at 151 MHz, and optical nuclear emission at 5100\\,\\AA\\ for a complete sample of 135 compact jets.} {We use the partial Kendall's tau correlation analysis to check the link between radio properties of parsec-scale jets and optical luminosities of host active galactic nuclei (AGN).} {We find a significant positive correlation for 99 quasars between optical nuclear luminosities and total radio (VLBA) luminosities of unresolved cores at 15\\,GHz originated at milliarcseconds scales. For 18 BL Lacs, the optical continuum emission correlates with the radio emission of the jet at 15\\,GHz. We suggest that the radio and optical emission are beamed and originate in the innermost part of the sub--parsec-scale jet in quasars. Analysis of the relation between the apparent speed of the jet and the optical nuclear luminosity at 5100\\,\\AA\\ supports the relativistic beaming model for the optical emission generated in the jet, and allows the peak values of the intrinsic optical luminosity of the jet and its Lorentz factor to be estimated for the populations of quasars and radio galaxies. The radio-loudness of quasars is found to increase at high redshifts, which can be a result of lower efficiency of the accretion in AGN having higher radio luminosities. A strong positive correlation is found between the intrinsic kinetic power of the jet and the apparent luminosities of the total and the unresolved core emission of the jet at 15\\,GHz. This correlation is interpreted in terms of intrinsically more luminous parsec-scale jet producing more luminous extended structure which is detectable at low radio frequencies, 151\\,MHz. A possibility that the low frequency radio emission is relativistically beamed in superluminal AGN and therefore correlates with radio luminosity of the jet at 15\\,GHz can not be ruled out (abridged).} {} ", "introduction": "The orientation-based unification schemes of radio-loud active galactic nuclei (AGN) \\citep{barthel89,urry95} suggest that the continuum and broad-line emission from their central engine are seen directly in powerful FR II \\citep{fanaroff74} quasars, while in FR II radio galaxies the central emission can be partially/completely hidden by obscuring material (``dusty torus''). In this scheme the presence of relativistic jets implies that radio-loud quasars are the relativistically beamed counterparts of radio galaxies. The beamed synchrotron emission from the base of the jet may extend to optical wavelengths in quasars \\citep{impey90} and even in the relatively unbeamed radio galaxies \\citep{chiaberge02,hardcastle00}. We should expect that in quasars the emission from both the central engine (thermal?) and the jet (non-thermal) contribute to the total power, whereas in radio galaxies the bulk of the optical continuum emission may be attributed to the relativistic jet rather than the central engine hidden by the torus \\citep[e.g.,][]{arshakian10,tavares10}. One approach to investigate the physical processes in active galactic nuclei (AGN) at scales not reachable by present-day telescopes is to study the correlations between radiative energy in different wavebands. There is evidence that the beamed synchrotron emission from the base of the radio jet extends to visible wavelengths in BL Lacs and quasars \\citep{wills92}. Several authors \\citep{hardcastle00,chiaberge02,kharb04} investigated the correlations between the radio and optical regimes for radio-loud FRI-FRII radio galaxies with the sample size $\\la 65$ objects. On the basis of the correlations between unresolved optical core emission (in the high-resolution images with the \\emph{Hubble Space Telescope}) and the VLA (Very Large Array) radio core emission of the jet on scales of milliarcseconds \\citep[see][]{hardcastle00}, and color information, these authors argued that optical nuclei are due to synchrotron radiation from the jet. The VLBI (Very Long Baseline Interferometry) imaging of jets at 15\\,GHz reaches an unprecedented milliarcsecond resolution \\citep{ken98,zensus02,ken04} and the VLBI cores are resolved on submilliarcsecond scales \\citep{kovalev05}. The VLBA provides a high reliability in the results due to the excellent calibration properties and the repeatability of observations. Here we investigate the radio-optical correlation between the VLBA core emission at 15\\,GHz and the optical nuclear emission at 5100\\,\\AA\\ to test a single production mechanism for radio and optical continuum emission on scales of submilliarcseconds. The relativistic outflows of plasma material form near the central nucleus and trace the pc-scale broad-line region and kpc-scale narrow-line region transporting the kinetic energy preserved in the jet to hundreds of kiloparsecs away from the central engine. To test the link between the properties of the pc-scale jet and extended radio structure on kiloparsec scales, as well as the correlation between VLBI radio and optical nuclear luminosities we use the complete sample of 135 core-dominated AGN possessing relativistic jets \\citep{lister05}. In Sect.~\\ref{sec:rsample} we introduce the radio sample of AGN, and in Sect.~\\ref{sec:jetPar} we define the radio parameters of their compact jets. The optical nuclear luminosities are derived in Sect.~\\ref{sec:optLum}. Correlations between radio properties of the jet, and between radio properties and optical nuclear luminosities are discussed in Sect.~\\ref{sec:rCorr} and \\ref{sec:roCorr} respectively. Throughout the paper a flat cosmology model is used with $\\Omega_{m}=0.3$ ($\\Omega_{\\Lambda}+\\Omega_{m}=1$) and $H_0=70$ km\\,s$^{-1}$\\,Mpc$^{-1}$. ", "conclusions": "\\label{sec:conc} We used the statistically complete sample of 135 compact radio sources provided by the MOJAVE program to investigate correlations between the properties of the pc-scale jets at 15\\,GHz, the extended radio emission at 151\\,MHz, and the optical nuclear emission at 5100\\,\\AA . The main results are summarized as follows: \\begin{enumerate} \\item We determine the optical nuclei fluxes at 5100\\,\\AA\\ for 135 MOJAVE-1 AGN with their photometric optical fluxes available in the MAPS and USNO-B catalogs. There is a significant positive correlation for 99 quasars and 18 BL Lacs between optical nuclear luminosities and jet luminosities at 15\\,GHz originated in the jet at milliarcseconds scales. For quasars, correlations hold also between optical nuclear luminosity and luminosities of the unresolved core (at sub-milliarcseconds scales) and total radio (VLBA) of the jet, suggesting that the optical emission is non-thermal and originates in the innermost part of the jet at sub-parsec scales, while in the BL Lacs it is generated in the parsec-scale jet. \\\\ The generation of the relativistically beamed optical emission in the MOJAVE blazars is evident from the apparent speed -- optical nuclear luminosity relation plane. In this diagram, the data of quasars, BL Lacs and radio galaxies are fitted by an aspect curve derived from the relativistic beaming theory, suggesting that optical emission from these sources is relativistically Doppler-boosted. We estimate the peak values of the intrinsic optical luminosity and the Lorentz factor of the jet for each population of blazars: $L_{\\rm 5100,\\,int}=2\\times10^{20}$~W Hz$^{-1}$, $\\gamma=52$, and $p=3$ for quasars, $L_{\\rm 5100,\\,int}=9\\times10^{21}$~W Hz$^{-1}$, $\\gamma=20$, and $p=2$ for BL Lacs, and $L_{\\rm 5100,\\,int}=1.5\\times10^{21}$~W Hz$^{-1}$, $\\gamma=9$, and $p=2$ for radio galaxies. About 95\\,\\% of quasars have the peak values of $L_{\\rm 5100,\\,int}=4\\times10^{20}$~W Hz$^{-1}$, $\\gamma=33$, and $p=3$. The boosting factor $p=3$ favors the source of optical continuum emission for quasars to be discrete and optically thin. \\\\ \\item The kinetic power of 135 compact jets, determined from the flux density measured at 151\\,MHz, is on average higher in quasars than in BL Lacs, and it is lower in radio galaxies. We find a strong positive correlation between the intrinsic kinetic power of the jet and the apparent luminosities of the total and unresolved core emission of the jet at 15\\,GHz. This correlation can be a result of the correlation between kinetic power and intrinsic luminosity of the pc-scale jet, which reflects that intrinsically luminous compact jets deliver more relativistic electrons to kiloparsec-scales. In this way they accumulate more low-energy electrons in the radio lobes, which results in powerful extended radio lobes radiating at low radio frequencies. Another possibility to interpret the correlation found in the $Q_{\\rm j}-L_{\\rm jet}$ relation plane is that the low-frequency emission at 151\\,MHz is relativistically beamed and its variability is correlated with beamed radio emission at 15\\,GHz. Correlated variability at low and high frequencies was found in few superluminal sources. It is most likely that both effects, the beaming at low frequencies and the intrinsic luminosity of the jet at 15\\,GHz in beamed AGN, contribute to the relation between kinetic power and apparent luminosity of the jet. Monitoring of the MOJAVE AGN in a wide range of frequencies and over a period of few years is needed to distinguish and estimate the contribution of each effect. The kinetic power of quasars positively correlates with the radio-loudness, i.e., the ratio of the radio flux density at 15 GHz to the optical flux at 5100\\AA\\,, and negatively with the compactness of the core (the ratio between the unresolved core and the total VLBA emission). \\\\ \\item No correlation is found between the intrinsic radio luminosity at 15 GHz and the Lorentz factor of the jet for 48 quasars. A larger number of the MOJAVE-1 AGN with measured Doppler factors are needed to confirm this result. \\\\ \\item We find that the radio-loudness of quasars increases with increasing the redshifts. This positive correlation is due to the lack of compact AGN with small radio-loudness at high redshifts. This effect is interpreted as a tendency for strong AGN detected at high redshifts to have a high radio-loudness. This can happen if the Doppler factor of the jet in the radio regime is higher than that in the optical. \\end{enumerate}" }, "1003/1003.2741_arXiv.txt": { "abstract": "Neutrino emission in processes of breaking and formation of nucleon Cooper pairs is calculated in the framework of the Larkin-Migdal and the Leggett approaches to the description of superfluid Fermi liquids at finite temperatures. We explain peculiarities of both approaches and explicitly demonstrate that they lead to the same expression for the emissivity in pair breaking and formation processes. ", "introduction": "One of important mechanisms for cooling of superfluid neutron star interiors is nucleon Cooper pair breaking and formation (PBF) with a radiation of neutrino-antineutrino pairs, $\"N\"\\to \"N\" +\\nu+\\bar{\\nu}$ (see \\cite{FRS76,VS87,SV87,MSTV90,SVSWW97,Minimal,KHY,YKL99,YLS99,V01,BGV04,GV05,PGW,Sedr07,KR,LP,SMS,KV08,L08,SR09,LeinsonWrong} and references therein). Neutrinos are produced in weak interactions in which the lepton current $l_\\mu=\\bar{\\nu}(1-\\gamma_5)\\nu$ is coupled to the vector ($V_\\mu=g_V\\,\\bar{\\Psi}\\, \\gamma^\\mu\\, \\Psi$) and the axial-vector ($A_\\mu= g_A\\,\\bar{\\Psi}\\, \\gamma^\\mu \\gamma_5\\, \\Psi$) nucleon currents; $\\mathcal{L}=-\\frac{G}{2\\,\\sqrt 2}\\,(V^\\mu-A^\\mu)\\,l_\\mu$, where $G\\approx 1.2\\times 10^{-5}$~GeV$^{-2}$ is the weak interaction coupling, $g_V=-1$ and $g_A=1.26$. For the nuclear vector current holds $\\partial_\\mu V^\\mu=0$ that corresponds to the conservation of the baryon charge. Early works~\\cite{FRS76,VS87,SV87,MSTV90,SVSWW97,Minimal,KHY,YKL99,YLS99,V01,BGV04,GV05,PGW,Sedr07} did not care about the vector-current conservation. The latter is fulfilled only if in-medium renormalization of the vector current is performed together with a corresponding renormalization of Green's functions. This problem was tackled in Refs.~\\cite{KR,LP,SMS,KV08,L08,SR09,PLPS}. Reference~\\cite{LP} indicated that the emissivity of the $1S_0$ PBF processes on the vector current should be dramatically suppressed ($\\propto v_{\\rm F}^4$, where $v_{\\rm F}$ is the Fermi velocity of non-relativistic nucleons) provided the vector current conservation is fulfilled. Reference~\\cite{LP} used expressions derived within the standard BCS formalism of the superconductivity theory \\cite{Nambu,Schriffer} for low excitation energies $\\om, |\\vec{q}|\\ll \\Delta$, where $\\Delta$ is the nucleon pairing gap and $\\om$ is the net $\\nu\\bar{\\nu}$ energy, whereas the PBF reaction kinematics permits only $\\om >2\\Delta$, $|\\vec{q}|<\\om$. The correlation effects in the particle-hole channel were neglected and processes induced by the axial-vector current were disregarded. The consistent calculation of the PBF emissivity induced by the vector and axial-vector currents was performed in Ref.~\\cite{KV08} within the Larkin-Migdal-Leggett Fermi liquid approach. The latter takes properly into account of the correlation effects in both particle-particle and particle-hole channels. It was demonstrated that the neutrino emissivity is actually controlled by the axial-vector current and is suppressed only by the factor $\\propto v_{\\rm F}^2$, rather than $\\propto v_{\\rm F}^4$ as was stated in Ref.~\\cite{LP}. This result was supported in the subsequent work~\\cite{L08}, which however continued to work out the vector current contribution neglecting the correlation effects in the particle-hole channel. The convenient Nambu-Gorkov formalism developed for the description of metallic superconductors, cf. Refs.~\\cite{Nambu,Schriffer}, does not distinguish interactions in particle-particle and particle-hole channels. These interactions can be, however, essentially different in strongly interacting matter, like in nuclear matter and in liquid He$^3$. The adequate methods for Fermi liquids with pairing were developed for zero temperature by Larkin and Migdal in Ref.~\\cite{LM63} (see also \\cite{M67}) and for a finite temperature by Leggett in Ref.~\\cite{Leg65a,L66}. The problem of calculation of a response function of a Fermi system to an external interaction becomes tractable at cost of introduction of a set of Landau-Migdal parameters for quasiparticle interactions. Parameters can be either evaluated microscopically or extracted from analysis of experimental data, see \\cite{M67}. The technical difference of the mentioned approaches is that Larkin and Migdal worked out equations for full in-medium vertices, whereas Leggett calculated directly a response function. The former approach was aimed at the study of transitions in nuclei, and the latter on the analyzes of collective modes in superfluid Fermi liquid. The principal equivalence of both approaches was emphasized by Leggett in Ref.~\\cite{Leg65a,L66}. In Ref.~\\cite{KV08} we used the Larkin-Migdal approach. More specifically, we solved the Larkin-Migdal equations for the vertices induced by the weak vector and axial-vector currents and calculated the current-current correlation function $\\chi_{a}$, $a =\\{V,A\\}$, and the PBF emissivity. The explicit expression for $\\chi_{a}$ was obtained at $T=0$. It is sufficient to calculate emissivity of the process for $T\\ll 2\\Delta$ since small exponential factor $e^{-2\\Delta/T}$ comes already from the phase space volume and the temperature correction to $\\chi_{a}(T=0)$ is small in this limit. Despite Larkin-Migdal equations were derived in their original paper for $T=0$, actually the results can be generalized for arbitrary $T$. In Ref.~\\cite{KV08} we sketched which expressions should be modified at finite temperatures. Recent paper~\\cite{LeinsonWrong} tried to adopt the Leggett formalism to calculate the PBF emissivity for the $1S_0$ neutron pairing and came to different results compared to those derived in Ref.~\\cite{KV08}, even in case $T\\ll \\Delta$. One of the points of Ref.~\\cite{LeinsonWrong} was to find the PBF emissivity for arbitrary $T\\neq 0$. In view of the explicit claim by Leggett in Ref.~\\cite{L66} on the equivalence of his approach to that of Larkin and Migdal, this difference looks worrisome and requires a clarification. The aim of this paper is to reveal the correspondence between the Larkin-Migdal and Leggett approaches and to generalize the results of Ref.~\\cite{KV08} to arbitrary temperatures. We argue also that the results of the work \\cite{LeinsonWrong} are based on misinterpretation and wrong solution of the Leggett equations. The paper is organized as follows. In Section~\\ref{sec:KV} we introduce the Fermi liquid approach to the problem of the neutrino emissivity from nucleon matter. We focus on the emissivity via the PBF reactions. The main expressions are presented in the diagrammatic form valid at arbitrary temperatures. Then in Section~\\ref{sec:Correspondence} we demonstrate how the Larkin-Migdal equations and the Leggett equations follow from the same set of diagrams. In Section~\\ref{sec:Solution} we solve the Larkin-Migdal equations for the vertices induced by the vector and axial-vector currents and, then, present expressions for the PBF emissivity recovering the results of Ref.~\\cite{KV08} formulated now for arbitrary temperatures. In Section~\\ref{sec:Leggett} we demonstrate how the same results can be obtained within the Leggett formalism. In Section~\\ref{sec:Critic} we discuss the flaws in Ref.~\\cite{LeinsonWrong}. We conclude with Section~\\ref{sec:Conclude}. ", "conclusions": "\\label{sec:Conclude} In Ref.~\\cite{LM63} Larkin and Migdal extended the Fermi liquid approach onto Fermi systems with pairing. The equations for the full normal and anomalous vertices (the Larkin-Migdal equations) have been derived. They considered a particular case of s-wave paring at zero temperature, aiming at applications to atomic nuclei~\\cite{Migdal59,ML64,Migdal64}. In Refs.~\\cite{Leg65a,L66} Leggett generalized this approach to the Fermi liquid at non-zero temperature and applied it to study the low-frequency, low momenta collective excitations. In difference with Larkin and Migdal, Leggett formulated equations in a matrix form for symmetric and antisymmetric vertices (the Leggett equations). Explicit equations were formulated for vertices with the symmetry ($1$, $\\sigma_3$, $p_\\mu$, $p_\\mu \\sigma_3$)\\,. The present analysis provides necessary, although self-evident, extensions of the Larkin- Migdal formalism~\\cite{LM63} to arbitrary temperature and of the Leggett formalism~\\cite{Leg65a,L66} to arbitrary frequencies and momenta and to the vector and axial-vector weak current symmetry. Efficiency of the Larkin-Migdal and the Leggett approaches generalized in such a way is demonstrated on example of the calculation of the neutrino emissivity in the reactions of the nucleon Cooper pair breaking and formation. To be specific we considered s-state neutron pairing and included only zero harmonics in the Fermi liquid interaction both in the particle-particle and in the particle-hole channels. Compared to our previous work \\cite{KV08}, where explicit expressions were found in the low temperature limit $T\\ll 2\\Delta$, here we performed generalizations to arbitrary temperatures. Since recently there were published works, where presence of the formal difference in the Larkin-Migdal and Leggett approaches have led to misleading conclusions, we carefully analyzed both approaches demonstrating explicitly that they indeed lead to the very same results. First, from the diagrammatic equation for the current-current correlator and vertices we reproduced the Larkin-Migdal and the Leggett equations. Within the Matsubara techniques for $T\\neq 0$ we showed the correspondence between the vertices and the loop functions introduced in both approaches. It turned out possible to cast all necessary loop functions in terms of the one master function $g_T$. In passing we proved the validity of a relation between the loop functions at arbitrary frequencies and momenta, which has been previously proven by Leggett in the low frequency-momentum region. Then we solved first the Larkin-Migdal and then the Leggett equations for the renormalized vertices of the neutral weak currents at arbitrary temperature and found the current-current correlator for the neutrino-antineutrino pair in the neutron superfluids, imaginary part of which describes the neutrino-antineutrino pair production in reactions with the nucleon Cooper pair breaking and formation. Also we proved the exact conservation of the vector current. The results of recent publication~\\cite{LeinsonWrong} based on the Leggett approach are found to be invalid because of falsely interpretation of the Leggett equations and their wrong solution." }, "1003/1003.4601_arXiv.txt": { "abstract": "A nonabelian generalization of the neutral Witten current-carrying string model is discussed in which the bosonic current-carrier belongs to a two dimensional representation of SU(2). We find that the current-carrying solutions can be of three different kinds: either the current spans a U(1) subgroup, and in which case one is left with an abelian current-carrying string, or the three currents are all lightlike, travelling in the same direction (only left or right movers). The third, genuinely nonabelian situation, cannot be handled within a cylindrically symmetric framework, but can be shown to depend on all possible string Lorentz invariant quantities that can be constructed out of the phase gradients. ", "introduction": "Topological cosmic strings or superstrings of cosmological size are one-dimensional extended objects which are believed to have been formed in the early phases of cosmological evolution. They are of considerable interest because they may offer a observable window on the high energy physics of the primordial universe, \\ie at grand unified scales. \\par Topological strings are produced in phase transitions associated with spontaneous symmetry breaking. This is the standard Kibble mechanism~\\cite{kibble1,kibble2}. Almost all supersymmetric grand unified theories in which hybrid inflation~\\cite{hybrid2,Copeland:1994vg,Dvali:1994ms} can be realized lead to the formation of topological strings~\\cite{rachel1,rachel2,JRS03,john}. Besides, most classes of superstring compactification lead to a spontaneous breaking of a pseudo-anomalous U$(1)$ gauge symmetry producing local cosmic strings \\cite{BDP98}. Such strings also form in the case where the Higgs field has a non-minimal kinetic term~\\cite{Babichev:2008qv}. \\par The simplest kind of topological string is the Nambu-Goto string which is described by the Nambu-Goto action~\\cite{GN1,GN2}. The Nambu-Goto action is the worldsheet formulation counterpart of a field theory description in which the string arises as a solitonic solution of the abelian Higgs model~\\cite{NO}. Such a string has no internal structure and is described entirely in terms of a worldsheet Lagrangian and the tension per unit length of the string. \\par Most observational signatures in the gravitational sector expected from topological strings have been derived and simulated numerically for Nambu-Goto strings. There are five main possible observational effects (see~\\cite{VilenkinBook,PPJPU} and references therein): beamed gravitational wave bursts from kinks and particle acceleration; deflection, gravitational lensing effects and multiple image effects; Doppler shifting effects; background gravitational radiation from string loops; and string effects in the cosmic microwave background. The existence of kinks along the strings has been shown to occur also for current-carrying strings~\\cite{3D} and the electromagnetic effects of such strings, which are absent in the simpler Nambu-Goto string, have been investigated. An especially interesting observational consequence of the presence of cosmic string networks in the early universe potentially because it is susceptible to be detected in the cosmic microwave background is the Gott-Kaiser-Stebbins effect~\\cite{Gott:1984ef,Kaiser:1984iv}. This effect consists in a temperature shift that is due to the gravitational lensing of photons passing near a moving source. \\par Cosmic superstrings are formed by tachyon condensation at the end of brane inflation~\\cite{Sarangi:2002yt,Jones:2003da}. The tachyons are complex scalars [with a local U(1) gauge symmetry] identifiable with the ground state open string modes of the Neveu-Schwarz sector that end on coincident non-BPS branes and antibranes~\\cite{Green:1994iw,Banks:1995ch,Green:1996um,Lifschytz:1996iq}. There exist associated gauge fields living on the brane and antibrane so that there exists a U(1)$\\times$U(1) symmetry on the brane-antibrane configuration. A first linear combination of the U(1)'s is higgsed~\\cite{Dvali:2002fi,Dvali:2003zj} leading to the appearance of a first kind of cosmic superstrings that are D $p$-branes with $p-1$ dimensions compactified~\\cite{Majumdar:2002hy}. In type IIB superstring theory, and given a spacetime manifold $\\mcl M$, such stable $p$-branes, can, for example, be obtained by considering a $p+2$ brane-antibrane pair stretching over a submanifold $\\real^{p+3}\\subset \\mcl M$. The $p+2$ brane-antibrane pair will annihilate unless a topological obstruction exists. This obstruction can be obtained from K-theory~\\cite{Sen:1998tt,Witten:1998cd,Olsen:1999xx}. A second linear combination of the U(1)'s leads to the formation of F-strings~\\cite{Dvali:2002fi,Dvali:2003zj}. \\par All these types of strings have until recently been considered as structureless, so their dynamics is given by the Nambu-Goto action. Numerical simulations of networks (see \\cite{FRSB08} and references therein) of such strings have been produced with the result of scaling, a property thanks to which the string network never comes to dominate the Universe evolution, but neither are the string completely washed out of the Universe, so their effect, however small, is still detectable. \\par The Nambu-Goto string can be generalized to the case of a string with internal structure. Such a string can be obtained by including a coupling of the string forming Higgs field to additional (bosonic or fermionic, with global or local, abelian, or nonabelian symmetry) fields in the theory. In part of the parameter space, these fields condense onto the string (the symmetry gets broken) leading to the appearance of currents on the worldsheet in the form of Goldstone bosons propagating along string~\\cite{Witten:1984eb}. In such a case, the current-carrying string can be described using a worldsheet Lagrangian and a nontrivial equation of state relating the tension per unit length to the energy density of the string~\\cite{formal3,carterPLB89,formal2,formal4}; the actual form of this equation of state was discussed numerically \\cite{neutral,enon0,NoSpring} and analytically \\cite{models}. The presence of currents on the worldsheet modifies only slightly the gravitational properties of the long strings \\cite{GravStringPuyPP,GP94}, but it also halts cosmic string loop decay caused by dissipative effects, thereby yielding new equilibrium configurations \\cite{2D,3D} named vortons~\\cite{vortons0,vortons1,vortons2,vortons3,vortons4,vortons5}. Those can potentially change drastically the cosmological network evolution, at the point of ruling such strings out. \\par Although the current-carrying property of cosmic strings is in fact fairly generic \\cite{lowmass,DP95,PGdA03}, a possibility that has, until now, been completely disregarded is that for which the string would be endowed not only with many currents \\cite{LPXcoupled}, but also with currents of a nonabelian kind, as is to be expected in most grand unified theories. This natural extension of the Witten idea leads to numerous new difficulties, as in particular the internal degrees of freedom manifold is intrinsically curved, so that a local, flat, description of the string worldsheet manifold, turns out to be inappropriate \\cite{BCIII,BCIV}. This paper is devoted to the specific task of obtaining the equivalent microscopic structure of a nonabelian current-carrying cosmic string. \\par To do so, we restrict attention to the global situation in which, in a way similar to the so-called neutral Witten model \\cite{neutral}, we wish to capture the essential internal dynamics of the string without the undue complication of adding extra gauge vector fields. In the case of an abelian current, it was indeed shown that these contributions, although of potential great cosmological relevance (see, e.g. Ref.~\\cite{boucleem} and references therein), can however be treated in a perturbative way, not modifying in any essential way the actual microscopic structure \\cite{enon0}. We therefore assume, as a toy model, a U$(1)$ Higgs model whose breaking leads to the existence of the strings themselves, coupled to an SU$(2)$ doublet through a scalar potential with parameters ensuring a condensate. We first describe the fields and notation, derive their dynamical equations in full generality, and then discuss the condensate configuration. After having recovered the abelian cases as particular solutions of the general nonabelian situation, we concentrate on the strictly nonabelian solutions. We obtain an exact configuration, called trichiral, and show how this model makes explicit the obstruction theorem first obtained by Carter \\cite{BCIII,BCIV}. We then derive the stress energy tensor and its eigenvalues, namely the energy per unit length and tension, and show that they depend on all the possible two-dimensional Lorentz invariants that can be constructed from the phase gradients (and second derivatives) of the angular variables in the internal space. We conclude by discussing the possible cosmological consequences of this new category of objects. ", "conclusions": "Cosmic string are an almost generic prediction of most high energy theories, and they can have many observational cosmological consequences. They can also be current-carrying, and this property changes their dynamics drastically, as it has been argued that a network of current-carrying cosmic string could overproduce equilibrium loop configurations which, if stable, would overclose the Universe; such strings are clearly ruled out. The last case that was not yet studied is that for which the current carrier transforms according to some representation of a nonabelian group, and this is what has been presented above, in the particular (simplest) example of (global) SU$(2)$. By means of such a toy model, we have been able to derive the microscopic structure of a nonabelian current-carrying string, and exhibit the characteristic features of its stress energy tensor, out of which one obtains, through integration over the transverse degrees of freedom, the energy per unit length and tension. In principle, these quantities allow for a complete calculation of the dynamics of the strings, hence of the motion of a network. \\par We have found many differences between the abelian and the nonabelian situations. Where the abelian case involves a single state parameter, the simplest nonabelian model here developed contains far more parameters, namely at least 8. Besides, when the abelian current case, even with more than one current, involves only the phase gradients of the fields, the nonabelian case at hand exhibits implicit dependencies in the second derivatives with respect to the worldsheet coordinates of these phases. Those phases also acquire a profile, i.e. they must vary between the string core and the exterior: in accordance with the general Carter argument \\cite{BCIII,BCIV}, the path followed by the phases on the SU$(2)$ 3-sphere could not be smoothly projected onto the worldsheet itself, the later being flat while the former being intrinsically curved. Finally, whereas in the many current case the eigenvalues of the stress energy tensor depend only explicitly on the cross gradients, the microscopic structure - the profiles - depending only on the squared gradients, in the nonabelian case the profiles, and hence the energy per unit length and tension, depend on all the possible two dimensional Lorentz invariants that can be built out of the phase derivatives up to the second order. \\par If cosmic strings were ever formed, it is quite likely that they would be current-carrying, and in this category, since the well-tested standard electroweak theory already contains a broken SU$(2)$ with a Higgs field doublet as in our case \\cite{lowmass}, the model we developed here may be relevant, depending on the values of the unknown coupling parameters. At the cosmological level, abelian current carrying strings do intercommute in much the same way as non conducting ones \\cite{MatznerLaguna}. This is made possible because the currents in both pieces of the colliding strings can merely add up at the junction, being confined in the worldsheet through a linear interaction. In the nonabelian case, it is likely that the essentially nonlinear interaction terms would forbid such a simple readjustment of the phases: it is to be expected that the intercommutation probability is much lower than for ordinary strings. This, as is well known from the superstring case \\cite{Polchinski05}, can imply fundamentally different cosmological consequences. Another reason why one would expect intercommutation to be far less effective in the nonabelian current-carrying case is also related to extra dimensions: in the simplest Kaluza-Klein framework with a circular fifth dimension, the extra angular variable plays the role of the current-carrier phase and the equation of state can be calculated to be of the self-dual fixed trace kind \\cite{models} by projecting in the 4 dimensional base space \\cite{Carter90}; it can be conjectured that introducing many extra dimension with a complicated structure can lead to currents sharing many of the properties of the nonabelian ones discussed here. The intercommutation of nonabelian current-carrying cosmic string is therefore an important open problem that deserves further investigation." }, "1003/1003.3021_arXiv.txt": { "abstract": "We use stripped-down versions of three semi-analytic galaxy formation models to study the influence of different assumptions about gas cooling and galaxy mergers. By running the three models on identical sets of merger trees extracted from high-resolution cosmological $N$-body simulations, we are able to perform both statistical analyses and halo-by-halo comparisons. Our study demonstrates that there is a good statistical agreement between the three models used here, when operating on the same merger trees, reflecting a general agreement in the underlying framework for semi-analytic models. We also show, however, that various assumptions that are commonly adopted to treat gas cooling and galaxy mergers can lead to significantly different results, at least in some regimes. In particular, we find that the different models adopted for gas cooling lead to similar results for mass scales comparable to that of our own Galaxy. Significant differences, however, arise at larger mass scales. These are largely (but not entirely) due to different treatments of the `rapid cooling' regime, and different assumptions about the hot gas distribution. At this mass regime, the predicted cooling rates can differ up to about one order of magnitude, with important implications on the relative weight that these models give to AGN feedback in order to counter-act excessive gas condensation in relatively massive haloes at low redshift. Different assumptions in the modelling of galaxy mergers can also result in significant differences in the timings of mergers, with important consequences for the formation and evolution of massive galaxies. ", "introduction": "\\label{sec:intro} Understanding how galaxies form and the physics that drive their evolution has been a long-standing problem in modern astrophysics. A number of observational tests have recently succeeded in determining the fundamental cosmological parameters with uncertainties of only a few per cent, thus effectively removing a large part of the parameter space in galaxy formation studies. We are left, however, with the problem of dealing with our `ignorance' of complex physical processes, that are inter-twined in an entangled network of actions, back-reactions, and self-regulations. Over the past decades, different methods have been developed to study galaxy formation in a cosmological context. Among these, semi-analytic models have turned into a flexible and widely used tool to provide detailed predictions of galaxy properties at different cosmic epochs. These techniques find their seeds in the pioneering work by \\citet{White_and_Rees_1978}, were laid out in a more detailed form in the early 1990s \\citep*{White_and_Frenk_1991,Cole_1991,Kauffmann_White_Guiderdoni_1993}, and have been substantially extended and refined in the last years by a number of different groups \\citep[for a review, see][]{Baugh_2006}. In these models, the evolution of the baryonic components is modelled adopting simple yet physical and/or observationally motivated {\\it recipes}, coupled in a set of differential equations that describe the variation in mass as a function of time of different galactic components (e.g. stars, gas, metals). While it is relatively easy to compare results from different models\\footnote{A number of galaxy catalogues have been made publicly available by various groups; results from different versions of two of the models used in this study are available through a relational database accessible at: \\\\ http://www.mpa-garching.mpg.de/millennium/}, it is more complicated to understand the origin of any difference or similarity between them. This difficulty stems primarily from the fact that different groups adopt different sets of prescriptions (that are equally reasonable, given our poor understanding of the physics at play) and that, as mentioned above, the final results are given by a combination of these. There are, however, also a number of more subtle differences that are more `technical' in nature (e.g. the particular mass definition adopted, the cooling functions used, the use of analytic or numerical merger trees, etc.). The precise influence on models' results of at least some of these details is unclear. For example, it has been shown that the extended Press-Schechter (EPS) formalism \\citep{Bond_etal_1991,Bower_1991} does not provide an adequate description of the merger trees extracted directly from numerical simulations \\citep{Benson_Kamionkowski_Hassani_2005,Cole_etal_2008}. Although some of these most recent studies have provided `corrections' to analytic merger trees, many applications are still carried out using the classical EPS formalism, and little work has been done to understand to which measure this affects predictions of galaxy formation models. In this paper, we compare results from three independently developed semi-analytic models. Our goal is not to predict or reproduce any specific observation. Rather our aim is to understand the level of agreement between different semi-analytic models, with a minimum of assumptions or free parameters. This requires (1) that the models are implemented on {\\it identical} merger trees, such that results can be compared on a case-by-case basis and any differences can be attributed to {\\it specific parametrizations of physical processes}; and (2) that a minimum of physical processes are included in the models in order to avoid possible degeneracies, and to hopefully illuminate the effects of specific parametrizations or parameter choices. In our study, the first requirement has been satisfied by creating a standard set of halo merger trees extracted from $N$-body simulations, and running each model on these trees. The second requirement is somewhat more demanding as modern semi-analytic models contain treatment of numerous, coupled physical processes. We have chosen to simplify the models as much as possible by removing {\\it all} physics other than gas cooling and galaxy mergers. This allows us to focus on the influence of different assumptions typically made to model these two physical processes, that represent two basic ingredients of any galaxy formation model. Using large samples of identical haloes, we are able to compare results both in a statistical fashion and on a halo-by-halo basis. Previous studies \\citep[e.g.][]{Benson_et_al_2001, Yoshida_et_al_2002, Helly_etal_2003, Cattaneo_et_al_2007} have compared numerical predictions from stripped-down versions of semi-analytic models with those from hydrodynamical simulations, to verify whether these methods provide consistent predictions in the idealized case in which only gas cooling is included. The general consensus is that the cooling model usually employed in semi-analytic models is in good agreement with hydrodynamical simulations that adopt the same physics. More recent studies focused on object-by-object comparisons, however, have highlighted a number of important differences that were `hidden' in the relatively good agreement obtained by previous studies focusing on statistical comparisons \\citep{Saro_etal_2010}. A recent study by \\citet{Viola_etal_2008}, in particular, has shown that the cooling model implemented in {\\small MORGANA} (one of the models used in this study) predicts cooling rates that are significantly larger than those predicted from their implementation of the `classical' cooling model, proposed by \\citet{White_and_Frenk_1991}. In addition, Viola et al. have shown that {\\small MORGANA} provides results that are in good agreement with those of controlled numerical experiments of isolated haloes, with hot gas in hydrostatic equilibrium. While both SPH and semi-analytic techniques have their own weaknesses, making it unclear which of the two (if either) is providing the `correct' answer, these results appear confusing. It is therefore interesting to study how the different possible assumptions that can be made to model the evolution of cooling gas, propagate on predictions from galaxy formation models. Modelling of galaxy mergers has not been considered a major concern, but different assumptions about merging time-scales can be made, and these may have important consequences for the inferred stellar assembly history of massive galaxies, including brightest cluster galaxies. In addition, recent work has shown that the classical dynamical friction formula usually adopted in semi-analytic models tends to under-estimate merging times computed from controlled numerical experiments and high-resolution cosmological simulations (\\citealt*{Boylan-Kolchin_Ma_Quataert_2008}, \\citealt{Jiang_etal_2008}). Results from these studies, however, have not yet converged on the appropriate correction(s) that should be applied to the classical formula. In this paper, we will not address the issue of what is {\\it the best way} to model galaxy mergers or gas cooling. Instead, we will concentrate on the differences due to alternative implementations of these physical processes, with the aim of understanding their effects in a full semi-analytic model. We will also explore what these effects might imply for the importance of other physical processes. The numerical simulations and merger trees used in our study are described in Section \\ref{sec:simstrees}. In Section \\ref{sec:models}, we describe in detail how gas cooling and galaxy mergers are treated in each of the three models used in this work. Section \\ref{sec:mw} and Section \\ref{sec:scuba} present our results for two halo samples. In Section \\ref{sec:ressub}, we discuss the influence of numerical resolution and of different schemes for the construction of merger trees. Section \\ref{sec:mergtimes} compares the different implementations of merger times adopted in the three models considered. Finally, we summarize and discuss our results, and give our conclusions in Section \\ref{sec:disc}. ", "conclusions": "\\label{sec:disc} In this work we have compared results from three independently developed semi-analytic models, and we have focused on alternative implementations of gas cooling and galaxy mergers. Our model comparison is carried out using two large samples of {\\it identical} merger trees, which allows us to compare results on a case-by-case basis and to focus on differences due to different model assumptions and parametrizations. In particular, we have constructed two FOF-based merger tree samples. One is a set of 100 haloes from the MS-II with redshift zero dark matter halo masses similar to that of the Milky Way (MW-like haloes), while the other consists of 100 haloes from the MS with a number density similar to that measured for {\\small SCUBA} galaxies at $z\\sim 2$ ({\\small SCUBA}-like haloes). By using stripped-down versions of the models, we are able to avoid possible degeneracies and complications due to a different treatment of various physical processes and to concentrate on the influence of specific assumptions and/or parametrizations. As explained above, we have chosen to include only the processes of gas cooling and galaxy mergers. In the following, we summarize briefly the results obtained and discuss them. \\subsection{Gas cooling} Our results show that the different assumptions adopted about gas cooling lead to very similar results on mass scales similar to those of our own Galaxy, and to a generally good statistical agreement. Important differences arise, however, on larger mass scales. Two of the models used in this study (the Munich and Durham models) assume variations of the cooling model originally proposed by \\citet{White_and_Frenk_1991}. The other model used here ({\\small MORGANA}) adopts a more sophisticated model and, on the basis of previous published results \\citep{Viola_etal_2008}, was expected to provide systematically higher cooling rates for our {\\small SCUBA}-like haloes. Contrary to this naive expectation, however, results from the Munich model and from {\\small MORGANA} are very similar for this mass scale, while the Durham model gives systematically lower cooling rates. These results appear to be in contradiction with previously published tests, but the contradiction is only apparent as we discuss below. As explained in detail in Section \\ref{sec:models}, although the Munich and Durham models are variations of the same cooling model, they differ in a number of details, in particular for the assumption adopted for the hot gas distribution. Red lines in Fig.~\\ref{fig:ex_fofsubscuba} show results from the standard Durham model (that assumes a $\\beta$ profile for the hot gas) and from a model that uses the same cooling implementation but assumes an isothermal distribution for the hot gas (darker red lines), as in the Munich model. The figure shows clearly that by changing this assumption, the predicted cooling rates are larger, bringing model results closer to (although they are still lower than) those obtained from the Munich model (green lines in Fig.~\\ref{fig:ex_fofsubscuba}). Some differences are still apparent, however: the Munich model predicts systematically higher cooling rates than the Durham model, particularly at high redshift, where cooling is much more efficient in this model than in both of the Durham implementations considered here. These residual differences are due to a number of other different assumptions, in particular for the calculation of the cooling radius (see Section \\ref{sec:moddiff}). Statistically, the differences between the models are relatively {\\it small}, which reflects a general agreement in the underlying framework of these two models. Perhaps more surprising is the similarity between the results obtained from the Munich model and those from {\\small MORGANA}, at both mass scales analysed in this paper. As discussed earlier, \\citet{Viola_etal_2008} have claimed that the cooling model implemented in {\\small MORGANA} predicts cooling rates that are significantly larger than those obtained using the classical model by \\citet{White_and_Frenk_1991}. It should be noted, however, that their implementation of the classical model did not assume any special treatment for the `rapid cooling regime' (as is instead done in the original work by White and Frenk and all subsequent variants of this model), and assumed that the hot gas distribution is described by a polytropic equation of state with index $\\gamma_p = 1.15$, which is similar to the $\\beta$ profile assumed in the Durham model. In addition, results discussed in Viola et al. were obtained for isolated static haloes and it is not trivial to generalize them to the case of cosmological mass accretion histories, like those we have used here. Their conclusions is, however, valid when one adopts a gas distribution similar to that adopted in the Durham model used in this study (which is indeed similar to that they assumed in their implementation of the classical model). Adopting a steeper gas profile, as in the Munich model used here, changes results significantly in some mass regimes, bringing them in very good agreement with those from the {\\small MORGANA} model. It is important to realize that the differences highlighted above will have important consequences on predictions from galaxy formation models. For example, the Munich and {\\small MORGANA} models will need to assume a stronger feedback than the Durham model, to counter-act excessive cooling at low redshift in relatively large haloes. Although all models achieve this using very similar schemes (the AGN feedback), the relative importance of these additional physical processes will be different in these three models. Predicting much lower cooling rates for massive haloes, the Durham model will have difficulties in providing large numbers of galaxies with elevated star formation at high redshift. Indeed, this model is able to reproduce the observed galaxy number counts at $850\\mu$m only by assuming a top-heavy initial mass function (IMF) from the stars formed in bursts. This works because a top-heavy IMF has a much larger recycled gas fraction, which provides fuel for star formation. As illustrated above, {\\small MORGANA} predicts much larger cooling rates than the Durham model and is indeed able to reproduce the number counts of submillimiter sources (but the brightest ones) with a standard IMF \\citep{Fontanot_etal_2007}. The Munich model used in this study predicts even larger cooling rates on average, but its predictions for the submillimiter number counts have not been explored yet. \\subsection{Galaxy mergers} The three models used in this study adopt a different modelling for galaxy mergers. The Munich and Durham model assume variations of the classical dynamical friction formula. Results from these two models are somewhat correlated, but there is a large scatter and a large number of galaxies get significantly longer merger times in one model than in the other. As explained in Section \\ref{sec:mergtimes}, this is mainly due to the adoption of a different numerical factor in front of the dynamical friction formula employed, and to a different assumption about the Coulomb logarithm. {\\small MORGANA} uses formulae derived from numerical simulations and analytic models that take into account dynamical friction, mass loss by tidal stripping, tidal disruption of subhaloes, and tidal shocks \\citep{Taffoni_etal_2003}. As shown above, over the range of mass-ratios that provide merger times lower than the Hubble time, the merger times obtained using these formulae are systematically lower (by a factor $\\sim 10 \\, \\tdyn$) than those computed from the Munich model. Of the three models used in this study, the Munich model uses merger times that are closer to the fitting formula recently proposed by \\citet{Boylan-Kolchin_Ma_Quataert_2008}, when neglecting any orbital dependence. The differences just discussed have important consequences for the stellar assembly history of massive galaxies, and for the formation and evolution of the brightest cluster galaxies and of the intra-group and intra-cluster light. {\\small MORGANA} (and to some extent also the Durham model) will tend to assemble massive central galaxies earlier than the Munich model. To keep these galaxies red, these models will need to assume a somewhat stronger supernovae feedback so as to make most of the mergers driving their late stellar assemble {\\it dry} (i.e. avoid triggering late bursts that would rejuvenate the stellar population of these galaxies). A different balance between AGN cooling and tidal stripping of stars will also be required in these models to keep model predictions in agreement with observational results. These considerations are of course valid in the case all models would use the same treatment of all other physical processes at play. We remind, however, that as mentioned in Section \\ref{sec:intro}, these processes are usually treated in a different way complicating the comparison between different models. \\subsection{Numerical resolution and merger tree scheme} Taking advantage of the mini-MSII, run using the same initial conditions and volume as for the MS-II but lower resolutions, we have analysed how the results discussed above vary as a function of numerical resolution. The level of agreement between the three models used in our study is not affected by numerical resolution. None of the models used in this paper, however, achieves a good numerical convergence, with all of them predicting moderately larger cooling rates in lower resolution runs. On the other hand, results seem to be quite stable to alternative schemes for the construction of dark matter merger trees. In particular, we have compared results obtained using FOF-based trees with those obtained using subhalo-based trees, which represent the standard input of the Munich model used in this study. The small differences found by comparing results from these two schemes can be ascribed to the possibility of tracking the accreted haloes until they are stripped below the resolution limit of the simulation by tidal stripping and truncation. Interestingly, our results show that these processes act on relatively short time-scales, and that they are more efficient at higher redshift. \\subsection{What Next?} \\label{sec:concl} One question that this study does not address is: what is the {\\it best} way to model gas cooling and galaxy mergers? This is a very difficult question, particularly for the cooling model. It is clear that one needs to understand how the gas is distributed in dark matter haloes, and how this distribution is affected by heating from supernovae and/or AGN feedback. Although hydrodynamical simulations of galaxy formation are becoming increasingly sophisticated, these physical processes still need to be included as `sub-grid' physics, i.e. using prescriptions that are `semi-analytical' in nature. As a consequence, published hydrodynamical simulations offer little indication of appropriate modelling of the hot gas distribution and evolution. Our results have pointed out that modelling of the rapid cooling regime differ significantly in the implementations discussed in this paper. Substantial numerical work has been focused recently on this mode of accretion, although it was was discussed as early as \\citet{Binney_1977}, \\citet{Rees_and_Ostriker_1977}, and \\citet{White_and_Frenk_1991}. Interestingly, recent studies show that gas accretion during the `quasi-static regime' in hydrodynamical simulations is sensitive to different implementations of SPH \\citep[][see also \\citealt{Yoshida_et_al_2002}]{Keres_etal_2009}, while accretion rates in the rapid cooling regime are quite robust. One possible additional concern in using numerical results to inform semi-analytic models is the poor performance of SPH codes in resolving and treating dynamical instabilities developing at sharp interfaces in a multi-phase fluid \\citep{Argetz_etal_2007}. The situation is somewhat better for the modelling of galaxy mergers. Since the merging process is predominantly driven by gravity, it can be studied using controlled numerical experiments as done, for example, by \\citet{Boylan-Kolchin_Ma_Quataert_2008}. As noted earlier, however, recent work has not yet converged on the dynamical friction formula appropriate for galaxy formation models. Further work in this area is therefore needed. Gas cooling and galaxy mergers are two basic ingredients of any galaxy formation model that are relatively well understood. Also at this level, however, different assumptions have to be made when implementing these processes. These give rise to non-negligible differences that can have important implications on the weight that needs to be given to additional physical processes (e.g. AGN feedback, tidal stripping of stars, etc.). This paper highlights specific areas where further work is needed in order to improve our galaxy formation models, with the ultimate goal of improving our understanding of the physical processes driving galaxy formation and evolution." }, "1003/1003.2996.txt": { "abstract": "Our Universe is ruled by quantum mechanics and its extension Quantum Field Theory (QFT). However, the explanations for a number of cosmological phenomena such as inflation, dark energy, symmetry breakings, and phase transitions need the presence of classical scalar fields. Although the process of condensation of scalar fields in the lab is fairly well understood, the extension of results to a cosmological context is not trivial. Here we investigate the formation of a condensate - a classical scalar field - after reheating of the Universe. We assume a light quantum scalar field produced by the decay of a heavy particle, which for simplicity is assumed to be another scalar. We show that during radiation domination epoch under certain conditions, the decay of the heavy particle alone is sufficient for the production of a condensate. This process is very similar to preheating - the exponential particle production at the end of inflation. During matter domination epoch when the expansion of the Universe is faster, the decay alone can not keep the growing trend of the field and the amplitude of the condensate decreases rapidly, unless there is a self interaction. This issue is particularly important for dark energy. We show that quantum corrections of the self-interaction play a crucial role in this process. Notably, they induce an effective action which includes inverse power-law terms, and therefore can lead to a tracking behaviour even when the classical self-interaction is a simple power-law of order 3 or 4. This removes the necessity of having nonrenormalisable terms in the Lagrangian. If dark energy is the condensate of a quantum scalar field, these results show that its presence is deeply related to the action of quantum physics at largest observable scales. ", "introduction": "\\label{sec:intro} Observations of phenomena such as superconductivity and super fluidity in condense matter indicates that quantum particles can collectively behave like a classical self-interacting scalar field. The potential energy of this interaction plays an important role in breaking global and/or local (gauge) symmetries which usually are followed by a phase transition. The same phenomena is assumed to happen at fundamental level in particle physics where usually a quantum scalar field, e.g. Higgs boson is responsible for dynamical mass generation. Other phenomena, mostly cosmological such as inflation, leptogenesis, and many of candidate models for dark energy are based on the existence of a classical scalar field which is usually related to a fundamental quantum scalar field because the physics of the Universe and its content in its most elementary level is quantic. A classical field is more than just classical behaviour of a large number of scalar particles. In a quantum system particles can be in superposition states i.e. quantum mechanically correlated to each others. Decoherence which is generated by interaction of each particle or field with its environment remove the quantum superposition and correlation between quantum states, but this does not mean that after decoherence of scalar particles, they behave collectively like a classical field. The following simple example can demonstrate this fact: \\\\ Consider a closed system consisting of a macroscopic amount of unstable massive scalar particles which decay to a pair of light scalar particles with a global $SU (2)$ symmetry and a very weak coupling with each other. If the unstable particle is a singlet of this symmetry, the remnant particles are entangled by their $SU (2)$ state. After a time much larger than the lifetime of the massive particle, the system consists of a relativistic gas of pair entangled particles. If a detector measures this $SU (2)$ charge without significant modification of their kinetic energy, the entanglement of pairs will break i.e. the system decoheres and becomes a relativistic gas. The equation of state of a relativistic ideal gas is $w_{rel} = P/\\rho \\approx 1 / 3$ where the pressure $P$ and density $\\rho$ are defined as the expectation value of some operators acting on the Fock space of the system. By contrast, a classical scalar field $\\varphi (x)$ is a $C$-number and its density ${\\rho}_{\\varphi}$, pressure $P_{\\varphi}$, and kinetic energy are defined as: \\bea &&{\\rho}_{\\varphi} \\equiv K_{\\varphi} + V (\\varphi) \\label{rhophi} \\\\ &&P_{\\varphi} \\equiv K_{\\varphi} - V (\\varphi) \\label{pphi} \\\\ &&K_{\\varphi} = \\frac {1}{2}g^{\\mu\\nu} \\partial_\\mu \\varphi \\partial_\\nu \\varphi \\label{kphi} \\eea where $V (\\varphi)$ is a potential presenting the self-interaction of the field $\\varphi (x)$. When it is much smaller than kinetic energy $ K_{\\varphi}$ , we obtain $P_{\\varphi} \\approx {\\rho}_{\\varphi}$, and if $V (\\varphi) \\gg K_{\\varphi}$, $P_{\\varphi} \\approx -{\\rho}_{\\varphi}$. Therefore in general, a relativistic gas and a scalar field do not share the same equation of state, and the proof of decoherence in a system is not enough when a classical scalar field is needed to explain a physical phenomenon. Historically, the concept of a classical scalar field was first appeared in the context of scalar-tensor - Brans-Dicke - gravity theories (see~\\cite{scalarhistory} for a historical review). In these models the scalar field presents dilaton, the generator of conformal symmetry. Therefore it had a purely geometric nature. It was only later when people tried to quantize Einstein and other gravity models that this field got a {\\it particle} interpretation. The discovery of Higgs mechanism and other phenomena in condense matter in which scalar fields are present, encouraged this interpretation. More recently scalar field are found to be a principle ingredient in supersymmetric and superstring theories. In the classical limit quantum scalar fields are usually identified with classical fields and their differences are overlooked. When a classical system is quantized, according to canonical quantization procedure, classical observables are replaced by operators acting on a Hilbert or Fock space, respectively for a single particle and for a multi-particle quantum system. The expectation values of these operators are the outcome of measurements. Therefore, it is natural to define the classical observable related to a quantum scalar field as its expectation value: \\be \\varphi (x) \\equiv \\langle \\Psi|\\Phi (x)|\\Psi\\rangle \\label{classphi} \\ee where $|\\Psi\\rangle$ is the state of the quantum system i.e. an element of the Fock space of the system. In analogy with particles in the ground state in quantum mechanics, the classical field $\\varphi (x)$ is also called a {\\it condensate}. In fact a coherent state consisting of superposition of particles in the ground state behaves like a classical field i.e. $\\langle \\Psi|\\Phi (x)|\\Psi\\rangle \\neq 0$~\\cite{condwave}. This is an ideal and exceptional case in which the number of particles in the system is infinite. Nonetheless, in the cosmological context where the number of particles is very large it can be a good approximation. Thus, later in this work we use this state to calculate the evolution of a condensate in an expanding universe. Using canonical representation, it is easy to see that for systems with a limited number of scalar particles $\\langle \\Psi|\\Phi|\\Psi \\rangle = 0$. But in presence of an interaction, even after renormalization, a finite term can survive~\\cite{infrenorm} to play the role of a classical field (condensate) according to the definition (\\ref{classphi}). In fact $\\Phi$ can be considered to be {\\it dressed} and its expectation value even on the vacuum can be non-zero. Equivalently, $\\Phi$ can be considered as a free field. In this case $|\\Psi\\rangle$ must include infinite number of interacting particles. In both interpretations the presence of an interaction is a necessary condition for the condensation of a finite system~\\cite{condint} (see also Ref.~\\cite{noneqqft,nonequi} for a review). Although classical scalar fields play crucial roles in the modeling of many phenomena particle physics and cosmology, their existence is usually considered as granted and the efforts are concentrated on the relevant potentials, solutions of their dynamic equations, and quantization of small fluctuations around the classical background fields. For instance in the context of inflation and reheating of the Universe, fluctuations of inflaton are quantized around the uniform and classical background which is responsible for the exponential expansion of the Universe (see e.g. Ref.~\\cite{cosmoweinberg} for a review). Both in inflation and in ultra-cold matter the presence of a condensate is apriori justifiable. If the entropy of the system e.g. Universe before inflation was very small and inflatons were the dominant content, most of them had to be in their ground state - the zero mode - and therefore according to Ref.~\\cite{condwave} behaved as a condensate (see also Appendix \\ref{app:b} for a more general description of a condensate state). In other contexts such as in the reheating era, and in cosmological and lab phase transitions the entropy is not always small. Therefore, as the above example showed, in these cases the formation of a condensate from a quantum scalar field is not a trivial process and the necessary conditions for the existence of such coherent behaviour must be investigated. In quintessence models a classical scalar field is the basic content of the model and its energy density is interpreted as dark energy. Although in the framework of popular particle physics models such as supersymmetry, supergravity, and string theory many efforts have been concentrated on finding candidate scalar fields to play the role of quintessence~\\cite{quincand}, little work has been devoted to understand what are the necessary conditions for a quantum scalar field to condense in a manner which satisfies the very special characteristics of a quintessence field. For instance, such a condensate must initially have a very small density, much smaller than other content of the Universe (smallness problem). Present observations show that dark energy behaves very similar to a cosmological constant i.e. with the expansion of the Universe its energy density does not change or varies very slowly. Such a behaviour is not trivial. In the classical quintessence models usually the potential of the model is {\\it designed} such that a tracking solution do exist. Potentials with such property are usually non-normalizable. Moreover, they don't directly correspond to potentials (or kinetic terms) expected from fundamental theories such as supersymmetry, supergravity or string theory. Therefore one has to relate them ad-hocly to some sort of low energy effective model of a fundamental theory. The purpose of the present work is to fill the gap between quantum processes producing various species of particles/fields in the early Universe - presumably during and after reheating - and their classical component as defined in (\\ref{classphi}). In another word, we want to see how the microscopic properties of matter is related to macro-physics and vis versa. We are particularly interested in condensation at very large scales, relevant to dark energy models. For other phenomena such as baryo- and lepto-genesis and Higgs mechanism, if the energy scale was much larger than Hubble constant of the epoch, the process can be studied locally. This is not the case for dark energy which seems to be uniform at largest observable scales, and therefore the expansion of the Universe could play important role in its evolution. As the quantum physics of the epoch just after reheating is not well known, we consider a simple toy model with a light scalar as quintessence field in interaction with two heavy scalar fields. Our aim is to study the evolution of the classical component - the condensate - of the quintessence field. Between many possible types of quantum scalar field and interaction models, we concentrate on a class of models in which the heaviest of three particle decays to other fields. The motivation for such a model is the results obtained from the study of the effect of a decaying dark matter on the equation of states of the Universe~\\cite{houriustate}. It has been shown that a FLRW cosmology with a decaying dark matter and a cosmological constant behaves similar to a cosmology with a stable dark matter and a dark energy with $w = P/\\rho \\lesssim -1$. This is effectively what is concluded at least from some of present supernovae observations~\\cite{snobs}. More recently, the same effect has been proved to exist for the general case of interaction between dark matter and dark energy~\\cite{quindmint}. It has been also shown~\\cite{houridmquin} that if a decaying dark matter has a small branching factor to a light scalar field, the observed density and equation of state of dark energy can be explained without extreme fine-tuning of the potential or coupling constants. In other word, such a model solves both smallness and coincidence problems of dark energy. The present work should complete this investigation by studying the formation and evolution of classical component from a quantum point of view. More generally, it is believed that all the particles are produced directly or indirectly from the decay of inflaton or curvaton (in curvaton inflation models) oscillation. Quintessence field is not an exception and irrespective of the details of its physics, it has to be produced from the decay of the inflaton or another field. Although the toy model considered here basically assumes a long life heavy particle, in each step of calculation we also mention the differences in the results if the life time of the decaying particle is short. The main difference between these cases is the time duration in which the production of quintessence scalar by the decay is significant. In Sec. \\ref{sec:decay} we describe the Lagrangian of a decaying dark matter model and evolution equation of the condensate. We consider three decay modes for the heavy particle and use the closed time path integral method to calculate the contribution of interactions in the condensate evolution. The same methodology has been used for studying inflation models~\\cite{infquant}, late-time warm inflation~\\cite{warminf}, the effects of renormalization and initial conditions on the physics of inflation~\\cite{infrenorm}, baryogenesis~\\cite{baryo}, and coarse-grained formulation of decoherence~\\cite{coarsegrain} (see also~\\cite{nonequi} and references therein). In Sec.\\ref{sec:solevol} we solve field equations and discuss their boundary conditions. In Sec. \\ref{sec:classevol} we obtain an analytical expression for the asymptotic behaviour of the condensate and discuss the importance of the quantum corrections. We summarize the results in Sec. \\ref{sec:conclu}. In Appendix \\ref{app:a} we obtain non-vacuum Green's functions in presence of a condensate. In Appendix \\ref{app:b} we generalize the description of a condensate to a system in which not all the particles are in the ground state. Appendix \\ref{app:c} presents the solution of evolution equations in matter dominated era. Finally, in Appendix \\ref{app:d} propagators in a fluctuating background metric are determined. ", "conclusions": "\\label{sec:conclu} Although we have not yet observed any elementary scalar field, they are believed to play important roles in the foundation of fundamental forces and phenomena that have shaped our Universe. On the other hand, we have observed the composite scalar fields and their condensation in condense matter where they are responsible for various phenomena. such as symmetry breaking, mass acquisition of photons, quantization of flux tubes, formation of topological defects, and many other exotic behaviours of matter. From these findings we have learned that the condensate has usually a simple potential which can be related to the interactions in the original Lagrangian - at least qualitatively. For instance, in the case of Cooper pairs in superconductors, the presence of a $\\varphi^4$ potential can be schematically interpreted as an elastic scattering between electrons inside two Cooper pairs due to electromagnetic interaction. At low energies it is seen as a point-like scattering of 2 incoming scalar particles to two outgoing scalars - similar to self-interaction of a scalar field. The potential of the condensate of a fundamental scalar field should also trace back to the Lagrangian of the quantum field. Many particle physics applications of scalar condensates only consider the classical order which correspond to the potential in the Lagrangian. Because Lagrangians are usually local, the effect of the classical term is also local. However, as we have demonstrated in the previous sections, in some circumstances the effect of quantum corrections can be very crucial. In fact the properties of the condensate discussed in the previous sections, and in Appendices \\ref{app:a} and \\ref{app:b} are related to the non-local properties of quantum mechanics, and by extension quantum field theory. The descriptions of the quantum state of a condensate suggested in Ref.~\\cite{condwave} and its generalization in Appendix \\ref{app:b}, include infinite number of entangled particles. Non-locality of quantum mechanics assures that these particles {\\it feel} the presence of each others even at largest cosmological distances, and therefore behave collectively at large scales. At early times when the Universe is dense and the probability of scattering between particles is large, the local effect of the classical potential as well as the production of $\\Phi$ particles by the decay of $X$ particles, which at lowest order is like a classical scattering, controls the amplitude of the condensate and the distribution of free $\\phi$ particles. But due to the expansion of the Universe, the cross-section of interaction and scattering becomes smaller at late times, and non-local effects become dominant. The very small coupling of $\\Phi$ with other fields assures the stability of the condensate. It could be destroyed if interactions were strong enough to wash out the coherence of the condensate state, i.e. made particles to behave individually and semi-classically. Apriori it should be possible to design experiments or observations capable of distinguishing between a condensate of quantum origin and a fundamentally classical field as dark energy. If the latter case is true, dark energy must be due to a modification of the general relativity which is believed to be a classical theory at least for scales $k \\lesssim M_{Planck}$. In this case the classical field would have a purely geometrical origin. On the other hand, if dark energy is produced by a condensate we expect to see some quantum effects. For instance to be able to observe its quantum excitation, similar to the excitation of a Bose-Einstein condensate~\\cite{bec}. Evidently, due to the extremely small coupling of dark energy, the production of such excitations in the lab or their observation in cosmological environment is extremely difficult, if not impossible. Nonetheless, with the progress of our understandings in condense matter physics and related technological advance, there is hope that one day such an exploration become possible. If numerical simulations confirms our conclusions about the order of self-interaction potential, this would be a very significant result. In quantum field theories the dimension of an interaction term determines its renormalizability. For scalar fields in 4-dimension spacetimes $\\Phi^n$ with $n \\leqslant 4$ are renormalizable. In fact except this physically motivated requirement there is no other rule to constrain the self-interaction in a quantum field theory. Interestingly enough, these values for $n$ correspond exactly to models for which a late time cosmological condensate seems to excite. In the $n=3$ models the coupling has a mass dimension of one which is considered to be the vacuum expectation value of another field, or to be proportional to the Planck mass, the only dimensional fundamental constant we know. In the latter case the field $\\Phi$ is probably related to quantum gravity models. In addition, it must be a singlet or a 1-form in the group manifold of the symmetry group of the model, otherwise it breaks the symmetry. On other hand, in a $n=4$ model the coupling constant is dimensionless and $\\Phi$ can be in a nontrivial self-conjugate representation of the symmetry group. These observations help to constrain the candidate particle physics models of dark energy. We can also put a rough lower limit on the self-coupling of $\\Phi$. $\\Gamma_\\Phi$ the interaction width of $\\Phi$ must be larger than $H_0$ if $\\Phi$ is not yet freezed-out. For $n=4$ self-interaction this means $\\Lambda_\\Phi \\gtrsim \\sqrt{H_0 m_\\Phi}$, and for $n=3$ potential $\\Lambda_\\Phi/M_{Planck} \\gtrsim \\sqrt{H_0 m_\\Phi}/M_{Planck}$~\\cite{houridmquin}. In summary, we used quantum field theory techniques to study the condensation of a scalar field produced during the decay of a much heavier particle in a cosmological context. Such a process had necessarily happened during the reheating of the Universe. It can also happen at later times if the remnants of the decay don't significantly perturb primordial nucleosynthesis. We showed that one of the necessary conditions for the formation of a condensate is its light mass. By considering three decay models and a power-law self-interaction potential of arbitrary positive order, we showed that the self-interaction has an important role in the cosmological evolution of the condensate and its contribution to dark energy. In particular, we showed that only a self-interaction of order 3 or 4 can produce a stable condensate in matter domination epoch. These results are obtained analytically and by considering a number of simplifying approximations. Therefore, they need confirmation by a more precise calculation, which in the face of complexity of this model, must be numerical. With little modification or adaptation, most of the formulation and results of this work are applicable to other cosmological phenomena which are based on a scalar condensate. {\\bf We finish this section by reminding that if dark energy is the condensate of a scalar field, the importance of the quantum corrections in its formation and its behaviour found here is the proof of the reign of Quantum Mechanics and its rules at largest observable scales.} \\vspace{1.cm} \\begin {appendix} {\\Large \\bf Appendixes}" }, "1003/1003.3450.txt": { "abstract": "We measure the half-light radii of globular clusters (GCs) in 43 galaxies from the ACS Fornax Cluster Survey (ACSFCS). We use these data to extend previous work in which the environmental dependencies of the half-light radii of GCs in early type galaxies in the ACS Virgo Cluster Survey (ACSVCS) were studied, and a corrected mean half-light radius (corrected for the observed environmental trends) was suggested as a reliable distance indicator. This work both increases the sample size for the study of the environmental dependencies, and adds leverage to the study of the corrected half-light radius as a possible distance indicator (since Fornax lies at a larger distance than the Virgo cluster). We study the environmental dependencies of the size of GCs using both a Principal Component Analysis (PCA) as well as 2D scaling relations. We largely confirm the environmental dependencies shown in \\citet{J05}, but find evidence that there is a residual correlation in the mean half-light radius of GC systems with galaxy magnitude, and subtle differences in the other correlations --- so there may not be a universal correction for the half-light radii of lower luminosity galaxy GC systems. The main factor determining the size of a GC in an early type galaxy is the GC color. Red GCs have $\\langle r_h \\rangle = 2.8\\pm0.3 $~pc, while blue GCs have $\\langle r_h \\rangle = 3.4\\pm0.3$~pc. We show that for bright early-type galaxies ($M_B < -19$ mag), the uncorrected mean half-light radius of the GC system is by itself an excellent distance indicator (with error $\\sim 11$\\%), having the potential to reach cosmologically interesting distances in the era of high angular resolution adaptive optics on large optical telescopes. ", "introduction": "The launch of the Hubble Space Telescope (HST) revolutionized areas of astronomy which push the limits of high resolution imaging. One area which benefited has been the study of the half-light radii, $r_h$, of globular clusters (GCs). A typical GC has a half-light radius of a few parsecs. Ground-based imaging (before the era of adaptive optics) could resolve such objects only in galaxies in the Local Group (LG). This limited the statistics and confined the study to the GC systems in dwarf and late-type galaxies (\\eg,~ the sample in \\citealt{F00}). HST opened the study to GC systems in galaxies out to $\\sim 30$ Mpc, encompassing large numbers of all types of galaxies and including both the Virgo and Fornax clusters (at $\\sim 16.5\\pm1.1$ Mpc from \\citealt{M07} and $20.0\\pm 1.4$ Mpc from \\citealt{B09}, respectively). Interest in the half-light radius of GCs results both from the constraints they provide on the formation and evolution of GCs and also their possible use as a distance indicator, which dates back to an initial suggestion by \\citet{SS27}. The half-light radii of GCs, rather remarkably, are almost independent of GC mass \\citep{M00,J05,Barmby07,McLaugh08,Harris2009,H09b}, at least to $\\sim 10^6 M_\\odot$ \\citep{H05,M09}. In simulations they have been shown to be fairly constant as the GCs evolve \\citep{ST72,LS78,M90,AH98}, and may in fact trace the characteristic sizes of the proto-GC cloud \\citep{ML92,H09b}. In the Milky Way, it has long been known that $r_h$ increases systematically with galactocentric distance, although HST studies \\citep{KW01,J05,Harris2009} have shown that in early types, $r_h$ is much closer to being constant with galactocentric radius. In this paper we extend the work of \\citet[hereafter J05]{J05} which used ACS data on GCs in Virgo cluster galaxies (from the ACS Virgo Cluster Survey; ACSVCS, \\citealt{2004ApJS..153..223C}) to study the environmental impacts on the $r_h$ of GCs and calibrated a corrected mean half-light radius (i.e., corrected for the observed environmental trends) as a distance indicator. We add to this study similar data from the ACS Fornax Cluster Survey (ACSFCS; \\citealt{J07a}). This adds 43 galaxies to the sample, and also extends the lever arm for calibration of the corrected mean half-light radius as a distance indicator since the Fornax cluster is at a larger distance than the Virgo cluster. Other recent works have studied the half-light radius of GCs, extending both the total number and range of morphologies and environments studied. GC sizes in M31 and NGC~5128 were studied with emphasis on the fundamental plane of GCs by \\citet{Barmby07} and \\citet{McLaugh08}, respectively. \\citet{Barmby07} show (using data for M31, NGC~5128, the MW, the Magellanic Clouds and the Fornax dwarf spheroidal) that old GCs appear to have near-universal structural properties. Measurements of GC sizes in late type galaxies beyond the LG include the Sombrero galaxy \\citep{S06, H09b}, NGC~891 \\citep{H09a} and NGC~5190 \\citep{Forbes10}. Extremes in host galaxy luminosity are encompassed by the work on dwarf galaxies by \\citet{Georgiev09} and the study of six giant ellipticals of \\citet{Harris2009}. \\citet{DG07} studied the SB0 galaxy NGC~1533 in the Dorado group. At the limit of what can be currently done from the ground is the work \\citet{Gomez07}, who used IMACS at the Magellan telescopes to study GC sizes in NGC~5128. As we will comment below, these works extend many of the results on GC sizes we have obtained using the ACSVCS and ACSFCS to different host galaxy morphologies and environments. All in all, the structural properties of GCs seem to share many near-universal properties accross galaxy morphology and luminosity, but some differences seem to exist as well. The paper is organized as follows. In Section 2 we describe the data. In Section 3 we present the distributions of $r_h$ in Fornax early type galaxies. In Section 4 we describe the possible environmental dependencies which we explore with a principle components analysis (PCA) in Appendix A and traditional 2D trends in Section 5. In Section 6 we show the final corrections which are discussed in Section 7, both as tracers of the GC formation and evolution and for implications on the use of a corrected mean half-light radius as a distance indicator. Section 8 presents a summary of our conclusions. ", "conclusions": "We have used data on the half-light radii of GCs belonging to galaxies observed in the ACSFCS to extend the work of J05 in which the environmental dependencies of the half-light radii of GCs in early-type galaxies in the ACSVCS were studied, and a corrected mean half-light radius was suggested as a reliable distance indicator. By adding data from the ACSFCS, we increase the sample size for the study of the environmental dependencies, and add leverage to the study of the corrected half-light radius as a possible distance indicator (since Fornax is at a larger distance than the Virgo cluster). We find only subtle differences in the environmental dependencies of the sizes of GCs in Fornax cluster galaxies from what was found in the Virgo cluster by J05. In addition, we perform a Principal Component Analysis (PCA) to check that no major correlations are being hidden (in Appendix A). Looking at 2D relations, we again confirm the well known results that there is no correlation between GC size, $r_h$, and mass (for $M < 2\\times 10^6 M_\\odot$, but that blue GCs ($(g-z) < 1.05$ with $\\langle r_h\\rangle = 3.36\\pm 0.03\\pm0.25$ pc) are about 20\\% larger than red ones ($(g-z) > 1.15$ with $\\langle r_h\\rangle = 2.83 \\pm 0.02\\pm0.25$ pc). We show that the half-light radii of GCs in early-type galaxies in the Fornax cluster increase only slightly with galactocentric radius (or decreasing surface brightness) as was also found by J05 for systems in Virgo cluster early-types. In fact, the trend we find in Fornax cluster systems is slightly shallower than that seen in the Virgo cluster. As was found in J05, the trend of $r_h$ with galactocentric distance is significantly shallower (only $2\\sigma$ different from zero in Fornax galaxy GC systems) than that observed in our Galaxy perhaps pointing to a different formation scenario, or evolution of GCs in early types vs. late types (or in high density regions vs. low density regions). We discuss this observation in light of other recent studies of the GC systems of late-type galaxies and argue there is now some evidence for a general property of the GC systems of late-type galaxies in small groups to have a stronger trend of half-light radius with galactocentric distance than is seen in early types in clusters. This will provide stronger constraints on the differences in formation and evolution of GCs in different types of galaxies in different environments. We confirm the trend of mean half-light radius with galaxy color that was first observed in J05, but show suggestions that there is a residual correlation with galaxy luminosity in the mean half-light radius of GC systems of Fornax early-type galaxies which is larger than that seen in GC systems of Virgo early types. We revisit the two simple pictures of the origin of $r_h$ in GCs discussed by J05, arguing that the additional trend we observe for the mean half-light radius of GC systems in both Fornax and Virgo cluster early type galaxies to decrease with galaxy luminosity (as $\\langle r'_h \\rangle \\propto L^{-0.05\\pm0.02}$ in blue GCs and $\\langle r'_h \\rangle \\propto L^{-0.11\\pm0.04}$ in red GCs) may be providing some support for the ``pressure-confined proto-GC cloud'' model over tidal limitation of $r_h$. However since the size of this trend appears to depend on the details of the GC selection (as it was not present in the similar J05 sample) we suggest that it should not be over interpreted. We show that for the most luminous galaxies ($M_B < -19$ mag), the uncorrected mean half-light radius is by itself an excellent distance indicator, varying by around 10-15\\% across galaxies. This is especially true if we remove the unusual GC system of FCC 21 (Fornax~A). Once corrected for dependencies on GC and galaxy color and local surface brightness we find a constant value of $\\langle\\hat{r}_h\\rangle = 2.71 \\pm 0.05\\pm0.25$ pc for a GC with color $(g-z) = 1.2$, in a galaxy with color $(g-z)_{\\rm gal} = 1.5$ and at an underlying surface brightness of $\\mu_z = 21$ mag arcsec$^{-2}$ across giant galaxies in both the Virgo and Fornax clusters (for Virgo alone we find $\\langle\\hat{r}_h\\rangle = 2.67 \\pm 0.07\\pm0.25$ pc, in Fornax it is $\\langle\\hat{r}_h\\rangle = 2.78 \\pm 0.12\\pm0.25$ pc excluding FCC21). The same simple geometric calibration used by J05 to estimate an independent distance to the Virgo cluster of $D_{\\rm Virgo} = 16\\pm2.3$ Mpc gives a distance to the Fornax cluster of $D_{\\rm Fornax} = 18.4\\pm3.7$ Mpc (excluding GCs in FCC 21). This extension of the work of J05 to include GC systems in early-type galaxies in the Fornax cluster adds support to the idea of a constant mean $r_h$ in luminous early type galaxies, but suggests that the environmental dependencies may be subtly different in different environments which is especially important in the lower mass galaxies. The mean half-light radii in GC systems of massive early-type galaxies has the potential to provide a geometric distance measurement to bright early-type galaxies, which could reach cosmologically interesting distances (as quantified by J05) in the era of giant optical telescopes with adaptive optics." }, "1003/1003.6097_arXiv.txt": { "abstract": "Since the first limit on the (local) primordial non-Gaussianity parameter, $f_{\\rm NL}$, was obtained from the {\\sl Cosmic Background Explorer} ({\\sl COBE}) data in 2002, observations of the cosmic microwave background (CMB) have been playing a central role in constraining the amplitudes of various forms of non-Gaussianity in primordial fluctuations. The current 68\\% limit from the 7-year data of the {\\sl Wilkinson Microwave Anisotropy Probe} ({\\sl WMAP}) is $f_{\\rm NL}=32\\pm 21$, and the {\\sl Planck} satellite is expected to reduce the uncertainty by a factor of four in a few years from now. If $f_{\\rm NL}\\gg 1$ is found by {\\sl Planck} with high statistical significance, all single-field models of inflation would be ruled out. Moreover, if the {\\sl Planck} satellite finds $f_{\\rm NL}\\sim 30$, then it would be able to test a broad class of multi-field models using the four-point function (trispectrum) test of $\\tau_{\\rm NL}\\ge (6f_{\\rm NL}/5)^2$. In this article, we review the methods (optimal estimator), results ({\\sl WMAP} 7-year), and challenges (secondary anisotropy, second-order effect, and foreground) of measuring primordial non-Gaussianity from the CMB data, present a science case for the trispectrum, and conclude with future prospects. ", "introduction": "The physics of the very early, primordial universe is best probed by measurements of statistical properties of primordial fluctuations. The primordial fluctuations are the seeds for the temperature and polarization anisotropies of the CMB and the large-scale structure of the universe that we observe today. Therefore, both the CMB and the large-scale structure are excellent probes of the primordial fluctuations. In this article, we shall focus on the CMB. See the article by V. Desjacques and U. Seljak in this volume for the corresponding review on the large-scale structure as a probe of the primordial fluctuations. This article reviews a recent progress on our using the CMB as a probe of a particular statistical aspect of primordial fluctuations called ``{\\it non-Gaussianity}.'' Reviews on this subject were written in 2001 \\cite{komatsu:prep} and 2004 \\cite{bartolo/etal:2004}. The former review would be most useful for those who are new to this subject. In this article, we focus on the new discoveries that have been made since 2004. Particularly notable ones include: \\begin{itemize} \\item[1.] It has been proven that {\\it all} inflation models (not just simple ones \\cite{maldacena:2003,acquaviva/etal:2003}) based upon a single scalar field would be ruled out regardless of the details of models \\cite{creminelli/zaldarriaga:2004}, if the primordial non-Gaussianity parameter called $f_{\\rm NL}$ (more precisely, the ``local type'' $f_{\\rm NL}$ as described later) is found to be much greater than unity. \\item[2.] The optimal method for extracting $f_{\\rm NL}$ from the CMB data has been developed \\cite{komatsu/spergel/wandelt:2005,creminelli/etal:2006,smith/zaldarriaga:prep,yadav/etal:2008} and implemented \\cite{smith/senatore/zaldarriaga:2009}. The latest limit on the local-type $f_{\\rm NL}$ from the {\\sl WMAP} 7-year temperature data is $f_{\\rm NL}=32\\pm 21$ (68\\%~CL) \\cite{komatsu/etal:prep}. \\item[3.] The most serious contamination of the local-type $f_{\\rm NL}$ due to the secondary CMB anisotropy, the coupling between the Integrated Sachs-Wolfe (ISW) effect and the weak gravitational lensing, has been identified \\cite{goldberg/spergel:1999,verde/spergel:2002,smith/zaldarriaga:prep,serra/cooray:2008,hanson/etal:2009,mangilli/verde:2009}. However, note that the astrophysical contamination such as the Galactic foreground emission and radio point sources may still be the most serious contaminant of $f_{\\rm NL}$. These effects would pose a serious analysis challenge to measuring $f_{\\rm NL}$ from the {\\sl Planck} data. \\item[4.] The importance of distinguishing different triangle configurations of the three-point function of the CMB was realized \\cite{creminelli:2003,babich/creminelli/zaldarriaga:2004} and has been fully appreciated. It has been shown by many researchers that different configurations probe distinctly different aspects of the physics of the primordial universe. The list of possibilities is long, and a terribly incomplete list of references on recent work (since $\\sim$ 2004) is: \\cite{chen/etal:2007} on a general analysis of various shapes; \\cite{lyth/ungarelli/wands:2003,zaldarriaga:2004,lyth/rodriguez:2005,lehners:2008,bond/etal:2009,chambers/nurmi/rajantie:2010} on the local shape ($k_3\\ll k_1\\approx k_2$); \\cite{creminelli:2003,alishahiha/silverstein/tong:2004,arkani-hamed/etal:2004,seery/lidsey:2005} on the equilateral shape ($k_1\\approx k_2\\approx k_3$); \\cite{holman/tolley:2008,meerburg/etal:2009} on the flattend (or folded) shape ($k_1\\approx 2k_2\\approx 2k_3$); \\cite{senatore/smith/zaldarriaga:2010,huang:prep} on the orthogonal shape (which is nearly orthogonal to both local and equilateral shapes); \\cite{moss/chun:2007,arroja/mizuno/koyama:2008,renaux-petel:2009,meerburg/etal:2010,chen/wang:prep} on combinations of different shapes; and \\cite{chen/easther/lim:2007,chen/easther/lim:2008} on oscillating bispectra. Also see references therein. \\item[5.] The connected four-point function of primordial fluctuations has been shown to be an equally powerful probe of the physics of the primordial universe. In particular, a combination of the three- and four-point functions may allow us to further distinguish different scenarios. Many papers have been written on this subject over the last few years: \\cite{arroja/koyama:2008,arroja/etal:2009,chen/etal:2009,seery/sloth/vernizzi:2009} on single-field models; \\cite{boubekeur/lyth:2006,huang/shiu:2006,byrnes/sasaki/wands:2006,seery/lidsey:2007,seery/lidsey/sloth:2007,suyama/yamaguchi:2008,suyama/takahashi:2008,ichikawa/etal:2008,ichikawa/etal:2008b,cogollo/rodriguez/valenzuela-toledo:2008,rodriguez/valenzuela-toledo:2008,buchbinder/khoury/ovrut:2008,mizuno/etal:2009,mizuno/arroja/koyama:2009,gao/li/lin:2009,gao/hu:2009,byrnes/choi/hall:2009,byrnes/tasinato:2009,enqvist/takahashi:2008,enqvist/etal:2009,enqvist/takahashi:2009,kawasaki/takahashi/yokoyama:2009,renaux-petel:2009,huang:2008,huang:2009,chingangbam/huang:2009,chen/wang:prepb} on multi-field models; and \\cite{kawakami/etal:2009,takahashi/yamaguchi/yokoyama:2009} on isocurvature perturbations. CMB data are expected to provide useful limits on the parameters of the ``local-form trispectrum,'' $\\tau_{\\rm NL}$ and $g_{\\rm NL}$ \\cite{okamoto/hu:2002,kogo/komatsu:2006}. Preliminary limits on these parameters have been obtained from the {\\sl WMAP} data by \\cite{vielva/sanz:prep,smidt/etal:prep}. \\end{itemize} The number of researchers working on primordial non-Gaussianity has increased dramatically: Science White Paper on non-Gaussianity submitted to Decadal Survey Astro2010 was co-signed by 61 scientists \\cite{komatsu/etal:astro2010}. ", "conclusions": "Since the last review articles on signatures of primordial non-Gaussianity in the CMB were written in 2001 \\cite{komatsu:prep} and 2004 \\cite{bartolo/etal:2004}, a lot of progress has been made in this field. The current standard lore may be summarized as follows: \\begin{itemize} \\item[1.] {\\bf Shape and physics.} Different aspects of the physics of the primordial universe appear in different shapes of three- and four-point functions. \\item[2.] {\\bf Importance of local shape.} Of these shapes, the local shapes have special significance: a significant detection of the local-form bispectrum (with $\\fnlKS\\gg 1$) would rule out {\\it all} single-field inflation models, and the local-form trispectrum can be used to rule out a broad class of multi-field models (if not {\\it all} multi-field models) by testing $\\tau_{\\rm NL}\\ge (6\\fnlKS/5)^2$. \\item[3.] {\\bf Optimal estimators.} The optimal estimators of the bispectrum and trispectrum can be derived systematically from the expansion of the PDF. The optimal bispectrum estimator has been implemented. \\item[4.] {\\bf Secondary.} The most serious contamination of $\\fnlKS$ is due to the lensing-ISW coupling, which can be removed by using the template given in Section~4. \\item[5.] {\\bf Foreground.} The Galactic foreground contamination is minimal for $\\fnleq$ and $\\fnlor$, but it can be as large as $\\fnlKS\\sim 10$ for the local-form bispectrum. This must be carefully studied and eliminated in the {\\sl Planck} data analysis. The random (Poisson) point-source contamination can be removed by using the template given in Section~3.2. \\end{itemize} Some outstanding issues for the ``CMB and primordial non-Gaussianity'' include: \\begin{itemize} \\item[1.] {\\bf Second order.} (In Newtonian gauge) the products of the first-order terms and the intrinsically second-order terms in the sub-horizon limit do not contaminate the local-form bispectrum very much ($\\Delta\\fnlKS<1$). However, would the post-Newtonian effect give $\\Delta\\fnlKS\\sim 5$, as found by \\cite{pitrou/uzan/bernardeau:prep}? If so, we need to construct a template for this effect. \\item[2.] {\\bf More foreground.} How can we model the non-Poisson (clustered) point source bispectrum? How about the foreground (and secondary) contamination of the primordial trispectrum? \\item[3.] {\\bf Trispectrum estimators.} How can we implement the optimal trispectrum estimators for both local and non-local shapes? \\end{itemize} These issues would become important when the {\\sl Planck} data are analyzed in search of primordial non-Gaussianity. The {\\sl Planck} is expected to reduce the uncertainty in $\\fnlKS$ by a factor of four compared to the current limit, $\\fnlKS=32\\pm 21$ (68\\%~CL). If the {\\sl Planck} detected $\\fnlKS\\sim 30$, then the trispectrum would provide an important test of multi-field models. In particular, if $\\fnlKS\\gg 1$ and $\\tau_{\\rm NL}<(6\\fnlKS/5)^2$ are found, then {\\it all} single-field models and many (if not {\\it all}) multi-field models would be ruled out, and thus the standard paradigm of inflation as the origin of fluctuations would face a serious challenge. However, do not despair even if the {\\sl Planck} did not detect the primordial bispectrum or trispectrum - while the CMB may end its leading role as a probe of primordial non-Gaussianity (unless the next-generation, comprehensive CMB satellite which can measure both the temperature and polarization to the cosmic-variance-limited precision is funded \\cite{baumann/etal:2009}), the large-scale structure of the universe would eventually take over and substantially reduce the uncertainties in the local-form parameters such as $\\fnlKS$, $\\tau_{\\rm NL}$, and $g_{\\rm NL}$ (see Desjacques and Seljak's article in this volume)." }, "1003/1003.2211_arXiv.txt": { "abstract": "We report the discovery of \\hatcurb{}, a fairly massive transiting extrasolar planet orbiting the moderately bright star \\hatcurCCgsc\\ ($V = \\hatcurCCtassmv$), with a period of $P=\\hatcurLCP$\\,d. The transit is close to grazing (impact parameter \\hatcurLCimp) and has a duration of \\hatcurLCdur\\,d, with a reference epoch of mid transit of $T_c = \\hatcurLCT$ (BJD). The orbit is slightly eccentric ($e = \\hatcurRVeccen$), and the orientation is such that occultations are unlikely to occur. The host star is a slightly evolved mid-\\hatcurISOspec\\ dwarf with a mass of \\hatcurISOmlong\\,\\msun, a radius of \\hatcurISOrlong\\,\\rsun, effective temperature \\hatcurSMEteff\\,K, and a slightly metal-rich composition corresponding to $\\feh = \\hatcurSMEzfeh$. The planet has a mass of \\hatcurPPmlong\\,\\mjup\\ and a radius of \\hatcurPPrlong\\,\\rjup, implying a mean density of \\hatcurPPrho\\,\\gcmc. Its radius is well reproduced by theoretical models for the \\hatcurISOageshort\\,Gyr age of the system if the planet has a heavy-element fraction of about 50\\,\\mearth\\ (7\\% of its total mass). The brightness, near-grazing orientation, and other properties of \\hatcur{} make it a favorable transiting system to look for changes in the orbital elements or transit timing variations induced by a possible second planet, and also to place meaningful constraints on the presence of sub-Earth mass or Earth mass exomoons, by monitoring it for transit duration variations. ", "introduction": "\\label{sec:introduction} More than five dozen transiting extrasolar planets (TEPs) have been discovered to date, by both ground-based surveys and, recently, space-based surveys such as the {\\em CoRoT} and {\\em Kepler} missions \\citep{baglin:2006, borucki:2010}. The ground-based surveys are typically able to detect TEPs orbiting only the brighter stars. However this disadvantage has the compensatory advantage that those TEPs detected are amenable to a wide range of follow-up studies, which can provide valuable insight into their atmospheric properties and other physical conditions. Among the ground-based surveys, the Hungarian-made Automated Telescope Network \\citep[HATNet;][]{bakos:2004} survey has been one of the main contributors to the discovery of TEPs. In operation since 2003, it has now covered approximately 11\\% of the sky, searching for TEPs around bright stars ($8\\lesssim I \\lesssim 12.5$). HATNet operates six wide-field instruments: four at the Fred Lawrence Whipple Observatory (FLWO) in Arizona, and two on the roof of the hangar servicing the Smithsonian Astrophysical Observatory's Submillimeter Array, in Hawaii. Since 2006, HATNet has announced and published 13 TEPs. In this work we report our fourteenth discovery, around the relatively bright ($V = \\hatcurCCtassmv$) star we refer to as \\hatcur, also known as \\hatcurCCgsc{}. The layout of the paper is as follows. In \\refsecl{obs} we report the detection of the photometric signal and the follow-up spectroscopic and photometric observations of \\hatcur{}. In \\refsecl{bisec} we examine the spectroscopic evidence to confirm the planetary nature of the object. \\refsecl{analysis} describes the analysis of the data, beginning with the determination of the stellar parameters, and continuing with a description of our global modeling of the photometry and radial velocities. Results are presented in \\refsecl{results}, and discussed in \\refsecl{discussion}. ", "conclusions": "\\label{sec:discussion} At $V = \\hatcurCCtassmv$, \\hatcur{} is among the brightest TEP host stars presently known (only 11 of the other 66 are brighter), which should facilitate a wide range of follow-up studies. Our planetary mass determination of $\\mpl\\ = \\hatcurPPmshort$\\,\\mjup\\ places \\hatcur{b} in the category of the massive transiting planets discovered to date, which seem to be less common than those with smaller masses. In fact, a look at the mass distribution for the known transiting systems shows that it falls off precisely around 2\\,\\mjup. Since there is no obvious selection effect that would make these massive objects more difficult to detect or confirm (such as smaller radii, longer orbital periods, or fainter or hotter parent stars), it may be that they are indeed less common as a population. It has also been noticed \\citep[e.g.,][]{southworth:2009, joshi:2009} that the most eccentric cases among the transiting planets are all massive: all systems with $e > 0.2$ are more massive than 3\\,\\mjup, although the numbers are still small. A similar trend is seen among the non-transiting planets, bearing in mind the $\\sin i$ ambiguity in their masses. Among the transiting systems whose spin-orbit alignment has been measured through observation of the Rossiter-McLaughlin effect, it would also appear that the more massive ones are disproportionately misaligned compared to the ones with smaller-mass planets, considering their relative populations, although again the numbers are as yet too small to be conclusive. Obviously it would be useful to determine the spin-orbit alignment of the \\hatcur{} system, which should be readily measurable given that $v \\sin i = 8.4$\\,\\kms. The mass of \\hatcur{b} is quite similar to the recently discovered Kepler-5b \\citep{koch:2010}, but its radius is about 20\\% smaller. A comparison with the evolution models of \\cite{fortney:2007}, at an equivalent solar semimajor axis of $a_{\\rm equiv} = 0.0317$\\,AU, indicates that the radius of \\hatcur{b} is well reproduced for the \\hatcurISOageshort\\,Gyr age of the system if it has about 50\\,\\mearth\\ worth of heavy elements in its interior (about 7\\% of its total mass). This amount of metals is consistent with the correlation between core mass and metallicity of the parent star proposed by \\cite{Guillot:2006} and \\cite{Burrows:2007}, which seems to support the core-accretion mode of planet formation. The incident flux we compute for the planet averaged over its eccentric orbit, $\\langle F \\rangle = \\hatcurPPfluxavg \\times 10^9$ \\ergscmsq, places it in the proposed pM category of \\cite{fortney:2008}. These objects are expected to present temperature inversions in their atmospheres, and large day/night temperature contrasts. The equilibrium temperature of \\hatcur{b}, assuming zero albedo and full redistribution of the incident radiation, is \\hatcurPPteff\\,K. The non-zero eccentricity of the orbit for the relatively short period of 4.6 days raises the question of whether tidal forces have had enough time to circularize the orbit at the \\hatcurISOage\\,Gyr age of the system, or alternatively, whether there might be a second planet perturbing the orbit of \\hatcur{b} and pumping up its eccentricity. Following \\cite{adams:2006}, we estimate the timescale for tidal circularization to be $\\tau_{\\rm circ} \\approx 0.33$\\,Gyr, although this is a strong function of the poorly known tidal quality factor $Q_p$, for which we have adopted here the commonly used value of $10^5$. For $Q_p \\simeq 10^6$ the timescale is 10 times longer, and a recent study of the 0.78-day transiting planet WASP-19b has suggested $Q_p$ could be much larger still \\citep{hebb:2010}. Consequently we cannot rule out that the eccentricity we measure for \\hatcur{b} is primordial, at least in part, nor can we exclude the presence of a second planet in the system. The grazing orientation ($b = \\hatcurLCimp$) significantly enhances the detection sensitivity to additional planets in the system. These could manifest themselves through transit {\\it duration} variations (TDVs), or equivalently, changes in the orbital inclination angle as has been claimed for TrES-2b by \\cite{mislis:2009a} and \\cite{mislis:2009b}, and recently disputed by \\cite{scuderi:2010}. Perturbing planets may also induce variations in the times of mid-transit (transit {\\it timing} variations, or TTVs) through gravitational interaction \\citep[e.g.,][]{Dobrovolskis:1996, Holman:2005}. In general, there exists no exact closed-form expressions for evaluating the amplitudes of such signals since no general solution to the 3-body problem exists. One typical approach is to employ an approximation to simplify the problem. \\cite{Agol:2005} provided useful analytic expressions for the case of a coplanar system where the host star has a much larger mass than the two companion planets, for several different scenarios. The largest TTVs occur for planets in mean motion resonance, and we may use expressions [33] and [34] from \\cite{Agol:2005} to estimate the amplitude and libration period of the TTV signal in a 2:1 resonance. For the \\hatcur{} system we find that such a planet would induce a peak-to-peak deviation in the timings of 13.6\\,s and 126.3\\,s for a mass of $0.1\\,\\mearth$ and $1\\,\\mearth$, respectively, and the libration period would be $\\sim$175 days. Thus, an Earth-mass planet in this configuration would be eminently detectable with further observations. Out-of-resonance perturbers can be considered using equation [32] of \\cite{Agol:2005}, but we find they lead to sub-second amplitudes for various configurations, far beyond the current measurement capabilities. The properties of \\hatcur{b} also make it attractive for a future search for a companion satellite (`exomoon'). The Hill radius ($R_H$) extends to 8.1 planetary radii (9.1\\,$R_J$), and using the expressions of \\citet{domingos:2006}, which account for the planet's orbital eccentricity, the maximum distance at which an exomoon could be stable for this planet would be 0.436\\,$R_H$ and 0.824\\,$R_H$ for a prograde and retrograde exomoon, respectively. \\citet{barnes:2002} presented analytic approximations for the maximum stable exomoon mass for close-in extrasolar giant planets based upon tidal dissipation arguments, which were shown to provide excellent agreement with numerical integrations. Assuming again a tidal dissipation value of $Q_p \\simeq 10^5$ and a Love number of $k_{2p} \\simeq 0.51$, the maximum stable moon mass over the \\hatcurISOageshort\\,Gyr lifetime of the star is 0.002\\,\\mearth\\ and 0.12\\,\\mearth\\ for a moon at the maximum planet-moon separations for prograde and retrograde orbits, respectively. For $Q_p \\simeq 10^6$ these limits would be 10 times larger, or 1.2\\,\\mearth\\ in the retrograde case. Such exomoons would produce small TTVs, but the highly grazing orientation provides a better means of detection. \\citet{kipping:2009a, kipping:2009b} showed that an exomoon induces two types of transit duration variations on the host planet: a sky-projected tangential velocity variation (the V-component), and a spatially orthogonal transit impact parameter variation (the TIP component). While the TTV and TDV-V rms amplitudes of a 1.2\\,\\mearth\\ exomoon would be 4.5\\,s and 1.0\\,s, respectively, the TDV-TIP component could be as high as 16\\,s or 45\\,s peak-to-peak for the optimum configuration of a favorably inclined exomoon. For comparison, another near-grazing transiting planet is TrES-2b \\citep{odonovan:2006}, and repeating the same calculation (with $Q_p \\simeq 10^6$, as above) we find that the maximal TDV-TIP of the maximum-mass stable retrograde exomoon would be only 0.07\\,s. The much smaller value in this case is mostly due to the low maximum stable moon mass of 0.005\\,\\mearth, as a result of the lower planetary mass and longer system age. We suggest, therefore, that future monitoring of the transit duration of \\hatcur{b} could allow for interesting constraints not only on the presence of additional planets, but also of sub-Earth mass or Earth mass exomoons." }, "1003/1003.5613_arXiv.txt": { "abstract": "{ The Large Area Telescope (LAT), the main instrument of the Fermi Gamma-Ray Space Telescope, detects high energy gamma rays with energies from 20 MeV to more than 300 GeV. The two main scientific objectives, the study of the Milky Way diffuse background and the detection of point sources, are complicated by the lack of photons. That is why we need a powerful Poisson noise removal method on the sphere which is efficient on low count Poisson data. This paper presents a new multiscale decomposition on the sphere for data with Poisson noise, called Multi-Scale Variance Stabilizing Transform on the Sphere (MS-VSTS). This method is based on a Variance Stabilizing Transform (VST), a transform which aims to stabilize a Poisson data set such that each stabilized sample has a quasi constant variance. In addition, for the VST used in the method, the transformed data are asymptotically Gaussian. MS-VSTS consists of decomposing the data into a sparse multi-scale dictionary like wavelets or curvelets, and then applying a VST on the coefficients in order to get almost Gaussian stabilized coefficients. In this work, we use the Isotropic Undecimated Wavelet Transform (IUWT) and the Curvelet Transform as spherical multi-scale transforms. Then, binary hypothesis testing is carried out to detect significant coefficients, and the denoised image is reconstructed with an iterative algorithm based on Hybrid Steepest Descent (HSD). To detect point sources, we have to extract the Galactic diffuse background: an extension of the method to background separation is then proposed. In contrary, to study the Milky Way diffuse background, we remove point sources with a binary mask. The gaps have to be interpolated: an extension to inpainting is then proposed. The method, applied on simulated Fermi LAT data, proves to be adaptive, fast and easy to implement. } ", "introduction": "The Fermi Gamma-ray Space Telescope, which was launched by NASA in june 2008, is a powerful space observatory which studies the high-energy gamma-ray sky \\citep{Fermi}. Fermi's main instrument, the Large Area Telescope (LAT), detects photons in an energy range between 20 MeV to greater than 300 GeV. The LAT is much more sensitive than its predecessor, the EGRET telescope on the Compton Gamma Ray Observatory, and is expected to find several thousand gamma ray sources, which is an order of magnitude more than its predecessor EGRET~\\citep{egret:hartman99}. Even with its $\\mathbf{ \\sim 1 m^2}$ effective area, the number of photons detected by the LAT outside the Galactic plane and away from intense sources is expected to be low. Consequently, the spherical photon count images obtained by Fermi are degraded by the fluctuations on the number of detected photons. The basic photon-imaging model assumes that the number of detected photons at each pixel location is Poisson distributed. More specifically, the image is considered as a realization of an inhomogeneous Poisson process. This quantum noise makes the source detection more difficult, consequently it is better to have an efficient denoising method for spherical Poisson data. Several techniques have been proposed in the literature to estimate Poisson intensity in 2D. A major class of methods adopt a multiscale bayesian framework specifically tailored for Poisson data~\\citep{KolaczykNowak2000}, independently initiated by \\citet{Timmerman} and \\citet{Kolaczyk}. \\citet{Lefkimmiatis} proposed an improved bayesian framework for analyzing Poisson processes, based on a multiscale representation of the Poisson process in which the ratios of the underlying Poisson intensities in adjacent scales are modeled as mixtures of conjugate parametric distributions. Another approach includes preprocessing the count data by a variance stabilizing transform (VST) such as the \\citet{Anscombe} and the \\citet{Fisz} transforms, applied respectively in the spatial~\\citep{Donoho} or in the wavelet domain~\\citep{Fryzlewicz}. The transform reforms the data so that the noise approximately becomes Gaussian with a constant variance. Standard techniques for independant identically distributed Gaussian noise are then used for denoising. \\citet{Zhang} proposed a powerful method called Multi-Scale Variance Stabilizing Tranform (MS-VST). It consists in combining a VST with a multiscale transform (wavelets, ridgelets or curvelets), yielding asymptotically normally distributed coefficients with known variances. The choice of the multi-scale method depends on the morphology of the data. Wavelets represent more efficiently regular structures and isotropic singularities, whereas ridgelets are designed to represent global lines in an image, and curvelets represent efficiently curvilinear contours. Significant coefficients are then detected with binary hypothesis testing, and the final estimate is reconstructed with an iterative scheme. In \\citet{Starck09:fermi3d}, it was shown that sources can be detected in 3D LAT data (2D+time or 2D+energy) using a specific 3D extension of the MS-VST. There is, to our knowledge, no method for Poisson intensity estimation on spherical data. It is possible to decompose the spherical data into several 2D projections, denoise each projection and reconstitute the denoised spherical data, but the projection induces some caveats like visual artifacts on the borders or deformation of the sources. In the scope of the Fermi mission, we have two main scientific objectives: \\begin{itemize} \\item Detection of point sources to build the catalog of gamma ray sources, \\item Study of the Milky Way diffuse background, which is due to interaction between cosmic rays and interstellar gas and radiation. \\end{itemize} The first objective implies the extraction of the galactic diffuse background. Consequently, we want a method to suppress Poisson noise while extracting a model of the diffuse background. The second objective implies the suppression of the point sources: we want to apply a binary mask on the data (equal to $0$ on point sources, and to $1$ everywhere else) and to denoise the data while interpolating the missing part. Both objectives are linked: a better knowledge of the Milky Way diffuse background enables us to improve our background model, which leads to a better source detection, while the detected sources are masked to study the diffuse background. The aim of this paper is to introduce a Poisson denoising method on the sphere called Multi-Scale Variance Stabilizing Transform on the Sphere (MS-VSTS) in order to denoise the Fermi photon count maps. This method is based on the MS-VST~\\citep{Zhang} and on recent on multi-scale transforms on the sphere \\citep{starck2006,Abrial,starck:abrial08}. Section~2 recalls the multiscale transforms on the sphere which are used in this paper, and Gaussian denoising methods based on sparse representations. Section~3 introduces the MS-VSTS. Section~4 applies the MS-VSTS to spherical data restoration. Section~5 applies the MS-VSTS to inpainting. Section~6 applies the MS-VSTS to background extraction. Conclusions are drawn in Section~7. In this paper, all experiments were performed on HEALPix maps with $nside=128$~\\citep{Gorski}, which corresponds to a good pixelisation choice for the GLAST/FERMI resolution. The performance of the method is not dependent on the nside parameter. For a given data set, if nside is small, it just means that we don't want to investigate the finest scales. If nside is large, the number of counts per pixel will be very small, and we may not have enough statistics to get any information at the finest resolution levels. But it will not have any bad effect on the solution. Indeed, the finest scales will be smoothed, since our algorithm will not detect any significant wavelet coefficients in the finest scales. Hence, starting with a fine pixelisation (i.e. large nside), our method will provide a kind of automatic binning, by thresholding wavelets coefficients at scales and at spatial positions where the number of counts is not sufficient. ", "conclusions": "This paper presented new methods for restoration of spherical data with noise following a Poisson distribution. A denoising method was proposed, which used a variance stabilization method and multiscale transforms on the sphere. Experiments have shown it is very efficient for Fermi data denoising. Two spherical multiscale transforms, the wavelet and the curvelets, were used. Then, we have proposed an extension of the denoising method in order to take into account missing data, and we have shown that this inpainting method could be a useful tool to estimate the diffuse emission. Finally, we have introduced a new denoising method the sphere which takes into account a background model. The simulated data have shown that it is relatively robust to errors in the model, and can therefore be used for Fermi diffuse background modeling and source detection." }, "1003/1003.2802_arXiv.txt": { "abstract": "Effective energy windows (Gamow windows) of astrophysical reaction rates for (p,$\\gamma$), (p,n), (p,$\\alpha$), ($\\alpha$,$\\gamma$), ($\\alpha$,n), ($\\alpha$,p), (n,$\\gamma$), (n,p), and (n,$\\alpha$) on targets with $10\\leq Z\\leq83$ from proton- to neutron-dripline are calculated using theoretical cross sections. It is shown that widely used approximation formulae for the relevant energy ranges are not valid for a large number of reactions relevant to hydrostatic and explosive nucleosynthesis. The influence of the energy dependence of the averaged widths on the location of the Gamow windows is discussed and the results presented in tabular form. ", "introduction": "} Astrophysical reaction rates describe the change in abundances of nuclei due to nuclear processes in an astrophysical environment, such as a hot plasma composed of free electrons and atomic nuclei. The reaction rate per particle pair (or reactivity) is found by folding interaction cross sections $\\sigma$ with the appropriate energy distribution of the interacting particles in the plasma. For nucleons and nuclei interacting with each other in stars, the latter is the Maxwell-Boltzmann (MB) distribution, leading to the definition of the reactivity $\\mathcal{R}=F\\mathcal{I}$ with \\cite{ili07,rolfs} \\begin{eqnarray} F & = & \\sqrt{\\frac{8}{\\pi\\mu}}\\left(\\frac{1}{kT}\\right)^{\\frac{3}{2}}\\quad,\\\\ \\mathcal{I} & = & \\intop_{0}^{\\infty}\\sigma(E)Ee^{-\\frac{E}{kT}}\\, dE\\quad,\\label{eq:integrand}\\end{eqnarray} where $k$ denotes the Boltzmann constant, $T$ the plasma temperature, and $\\mu$ the reduced mass $\\mu=M_{1}M_{2}/(M_{1}+M_{2})$. Although the integration limits run from zero to infinity, the largest contributions to the integral $\\mathcal{I}$ stem from a narrowly confined energy range, depending on the energy-dependence of the cross sections and the MB distribution. This relevant energy range has been termed the Gamow window for charged particle reactions and is important for both nuclear experimentalists and theoreticians as it defines the energy window within which the reaction cross sections have to be known. Due to their importance, simple approximation formulas (see Eqs.\\ \\ref{eq:e0approx}, \\ref{eq:deltaapprox}, \\ref{eq:effneutron}, \\ref{eq:deltaneutron}) have been derived to estimate the effective energy windows for reactions (see next sections) and are frequently used. However, the derivations of these formulae make implicit assumptions which are not always valid and therefore they cannot be applied to a number of important cases. For example, it has been pointed out \\cite{ili07,newt07} that resonances below the conventionally computed Gamow window may contribute significantly to the reaction rate for narrow-resonance capture of charged particles on light targets. It will be shown in the following that this can be understood by a more appropriate treatment of the Gamow window calculation. The applicability of the approximation will be discussed in more detail and the appropriate energy windows derived quantitatively for charged-particle induced reactions (Sec.\\ \\ref{sec:Charged-particle-reactions}) and for neutral projectiles (Sec.\\ \\ref{sec:Reactions-with-neutrons}). Section \\ref{sec:Conclusion} concludes with a discussion of the validity of the present approach and a brief summary. ", "conclusions": "} It has to be pointed out that the preceding discussion made implicit assumptions which have to be scrutinized in any application of the results. An obvious assumption is that the cross sections used are correct. The results shown here were obtained with the 'FRDM' set given in \\cite{rath01}. Although transmission coefficients (and thus widths) are affected by low-lying nuclear states which can be populated in a reaction, the derived energy \\textit{windows} should be robust because the gross energy dependence of the cross section is the relevant quantity and not its absolute value. At astrophysical temperatures the MB distribution always favors energies below the Coulomb barrier. Therefore, the energy dependence of charged-particle widths is dominated by the Coulomb barrier penetration. Since the energy dependences of widths in all channels may contribute, the energy windows are also sensitive to the reaction $Q$ values. Accordingly, updates to masses may affect the conclusions close to the driplines. This also holds for neutron capture reactions. Although the selection of contributing partial waves strongly depends on the spectroscopy of a nucleus, the shape of the MB distribution is not strongly modified when folded with the energy dependence of any relevant partial wave, as explained in Sec.\\ \\ref{sec:Reactions-with-neutrons}. Therefore the relevant energy window for neutron capture is defined by the peak of the MB distribution. Another source of concern may be the use of statistical model Hauser-Feshbach cross sections for nuclei with low level density close to the driplines, especially for rates at low plasma temperature. As explained above, the energy windows are mostly determined by the Coulomb barrier penetration in the different channels or by the MB distribution. This will also hold for direct reactions. As explained below, the derived effective energy windows remain valid even when narrow resonances contribute and therefore this is not a limitation of the method. A final assumption seems to be that the cross sections are smooth without isolated resonance features. The notion of a single Gamow peak loses its validity when narrow, isolated resonances are dominating the reaction rate. In this case, the Gamow window would be fragmented into several Gamow peaks of different importance, depending on the resonance strengths. It can be shown, nevertheless, that the notion of an effective energy window remains valid and that only resonances within the energy window contribute significantly to the reaction rate (for details see, e.g., \\cite{ili07}). This applies provided the energy windows are derived as shown above. It does not make a statement about the relative importance of resonances because this depends on the actual resonance strengths, not just on the energy dependence of the widths. Summarizing, a complete numerical study of the effective energy windows for nuclear reaction rates has been performed for reactions induced by nucleons and $\\alpha$ particles. It has been shown that the actual energy range of relevant cross sections differs considerably from the ranges obtained by application of the standard formulae. The origin of this difference was explained and extensive tables of the actual energy windows were given. This will be important for further theoretical improvements of reaction rates as well as for helping to design experiments to measure cross sections at energies of astrophysical importance." }, "1003/1003.0036_arXiv.txt": { "abstract": "A new empirical formulae is given for estimating the masses of black holes in AGNs from the H$\\beta$ velocity dispersion and the continuum luminosity at 5100~\\AA\\.~ It is calibrated to reverberation-mapping and stellar-dynamical estimates of black hole masses. The resulting mass estimates are as accurate as reverberation-mapping and stellar-dynamical estimates. The new mass estimates show that there is very little scatter in the $M_\\bullet$ -- $L_{bulge}$ relationship for high-luminosity galaxies, and that the scatter increases substantially in lower-mass galaxies. ", "introduction": " ", "conclusions": "" }, "1003/1003.5425_arXiv.txt": { "abstract": "We analyze star forming galaxies drawn from SDSS DR7 to show how the interstellar medium (ISM) Na I $\\lambda\\lambda$5890, 5896 (Na D) absorption lines depend on galaxy physical properties, and to look for evidence of galactic winds. We combine the spectra of galaxies with similar geometry/physical parameters to create composite spectra with signal-to-noise $\\sim300$. The stellar continuum is modeled using stellar population synthesis models, and the continuum-normalized spectrum is fit with two Na~I absorption components. We find that: (1) ISM Na D absorption lines with equivalent widths EW $>$ 0.8~\\AA\\ are only prevalent in disk galaxies with specific properties -- large extinction ($\\rm A_V$), high star formation rates (SFR), high star formation rate per unit area ($\\Sigma_{\\rm SFR}$), or high stellar mass ($M_*$). (2) the ISM Na D absorption lines can be separated into two components: a quiescent disk-like component at the galaxy systemic velocity and an outflow component; (3) the disk-like component is much stronger in the edge-on systems, and the outflow component covers a wide angle but is stronger within $60\\,^{\\circ}$ of the disk rotation axis; (4) the EW and covering factor of the disk component correlate strongly with dust attenuation, highlighting the importance that dust shielding may play in the survival of Na~I. (5) The EW of the outflow component depends primarily on $\\Sigma_{\\rm SFR}$ and secondarily on $A_V$; (6) the outflow velocity varies from $\\sim120$ to 160~km~s$^{-1}$ but shows little hint of a correlation with galaxy physical properties over the modest dynamic range that our sample probes (1.2 dex in log$~\\Sigma_{\\rm SFR}$ and 1 dex in $\\log M_*$). ", "introduction": "\\label{sec:intro} Galactic-scale gaseous outflows (`galactic winds') are known to be ubiquitous in very actively star forming galaxies at all cosmic epochs \\citep{heckman90, pettini00, shapley03, menard09, weiner09}. Galactic winds play an vital role in the evolution of galaxies and the intergalactic medium (IGM). The ``baryon deficit\" in the Galaxy \\citep[e.g.,][]{silk03} indicates gas removal by outflows during past active episodes. The mass-metallicity and effective yield relations observed in local galaxies suggest that galactic winds transport highly metal-enriched gas out of galaxies and into the IGM \\citep[e.g.,][]{garnett02,tremonti04,dalcanton07}. The $\\Lambda$CDM model over-predicts the galaxy luminosity function at both the low and high luminosity ends if a constant mass-to-light ratio is assumed. The most natural way to reconcile the observations with theory is to invoke feedback processes, including supernova, stellar winds \\citep{Cole00, Benson03, stringer08, Oppenheimer09} and AGN activity \\citep{Silk98, Hopkins06}. However, it remains unclear which kind of feedback processes dominate as a function of luminosity and comic epoch. In actively star forming galaxies, galactic winds are driven by the mechanical energy and momentum imparted by stellar winds and supernovae \\citep{chevalier85, heckman90}. Young star clusters create over-pressured bubbles of coronal phase gas which expand and sweep up shells of ambient ISM until they `blow-out' of the disk into the halo. The collective action of multiple superbubbles drives a weakly collimated bi-polar outflow consisting of hot gas and cool entrained clouds. Radiation pressure is also likely to play a role in accelerating cool dusty material \\citep[e.g.,][]{murray05}. In this paper we use the term `galactic wind' to describe such outflows without regard for their eventual fate --- the gas may escape the halo potential well, or be recycled back into the disk in a process sometimes referred to as `galactic fountain' activity \\citep{shapiro76, bregman80, kahn81}. In the last two decades, there have been many attempts to directly observe galactic winds in galaxies at $z=0- 3$ \\citep{veilleux05}. In the local universe, outflows can be detected via X-rays which trace hot gas \\citep[e.g.,][]{dahlem98,martin02,strickland04}, optical nebular emission lines produced by warm gas \\citep[e.g.,][]{lehnert96} and ISM absorption lines (e.g., Na I, K I) from cold gas \\citep[e.g.,][]{heckman02}. Blue-shifted Na D absorption from the entrained cool gas is frequently detected in IR-bright starburst galaxies \\citep{heckman00}, LIRGs and ULIRGs \\citep{rupke02, rupke05a, rupke05b, martin05, martin06, martin09}, and it is sometimes evident in dwarf starbursts \\citep{schwartz04}. The velocity of this gas correlates weakly with star formation rate (SFR) and galactic rotation speed, with a factor of $\\sim30$ change in velocity observed over a range of 4 orders of magnitude in SFR \\citep[$v_{\\rm wind}\\propto {\\rm SFR}^{0.3}$;][]{rupke05b,martin05}. At intermediate redshift, the Mg II $\\lambda\\lambda$2796, 2803 ISM absorption line shifts into the observed-frame optical. The relatively wide separation of this doublet makes it a good choice for outflow studies. \\citet{tremonti07} detected $500-2000$ ${\\rm km~s}^{-1}$ outflow velocities in a small sample of $z \\sim 0.5$ post-starburst galaxies, and suggested that these outflows cold be the relics of AGN-driven winds. \\citet{weiner09} employed the spectral stacking technique to probe $z \\sim 1.4$ star forming galaxies drawn from the DEEP2 survey. This study, which is based on $\\sim1400$ galaxy spectra, demonstrated that blue-shifted Mg II absorption is ubiquitous in actively star forming galaxies. At high redshift, the observed-frame optical samples the rest-frame far ultraviolet, which is rich in strong ISM resonance absorption transitions, but lacking in stellar spectral features. Because of the difficulty of measuring the relative velocity of the stars and gas, outflows have only been detected in very luminous Lyman break galaxies \\citep{pettini00, pettini02} or in composite spectra \\citep{shapley03}. While much of our knowledge about galactic winds comes from studies of local star forming galaxies, the work to date has been based on relatively small samples of extreme objects (dwarf starbursts, ULIRGs). The properties of outflows in local normal star forming galaxies are still largely unknown. In this paper, we investigate Na D absorption in a sample of $\\sim$150,000 star forming galaxies drawn from the Sloan Digital Sky Survey. By stacking the spectra of galaxies selected to have similar physical attributes, we obtain very high S/N composite spectra. After carefully modeling and dividing out the stellar continuum, we are able to probe very low Na D Equivalent Widths (EWs) and extend our analysis over a wide range in galaxy physical parameters. This paper is arranged as follows. In \\S2, we introduce the sample selection criteria used in our study. Our method of creating composite spectra and measuring the Na D lines is developed in \\S3. We apply the method to the SDSS sample in \\S4 and summarize our results in \\S5. We use the cosmological parameters $H_0=70~{\\rm km~s^{-1}~Mpc^{-1}}$, $\\Omega_{\\rm M}=0.3$ and $\\Omega_{\\Lambda}=0.7$ throughout this paper. ", "conclusions": "In this paper, we study interstellar Na~I ``D'' $\\lambda\\lambda5890,5896$ absorption in a large sample of star forming galaxies drawn from SDSS DR7. At the low SFRs probed by our sample, cool stars with prominent stellar Na D lines make a significant contribution to the integrated light. We account for this by modeling the full galaxy continuum using the CB08 stellar population synthesis models. We find a high incidence of strong ISM Na D absorption lines (EW $> 0.8$~\\AA) in galaxies that are massive, heavily dust attenuated, and that have high $\\Sigma_{\\rm SFR}$(Fig.~\\ref{Fig_frac_NaD}.) We use the spectral stacking technique to increase the signal-to-noise ratio of the spectra and the completeness of the sample and stack spectra in various bins of galaxy physical properties. In the continuum-normalized spectra we identify two interstellar Na D absorption components --- one at the systemic velocity, which is most pronounced in edge-on galaxies, and one outflow (blue-shifted) component which is dominant in face-on galaxies. We use two-component absorption line profile fits to measure the outflow velocity, the absorption line EW, the line width, covering factor, and the optical depth. As highlighted in the appendix, our results are robust with respect to the modeling of the stellar continuum. We find that: \\begin{itemize} \\item The ISM Na D absorption arises from cool gas in the disk, and from cool gas entrained in a galactic wind that is perpendicular to the disk and has an opening angle of $\\sim60\\,^{\\circ}$. \\item For the systemic (disk) component, the Na D EW depends most strongly on the dust attenuation. Dust attenuation is undoubtedly important in shielding Na~I from the ionizing radiation of OB stars. At fixed $\\rm A_V$, there is an additional dependence on $\\Sigma_{\\rm SFR}$. Galaxies with higher $\\Sigma_{\\rm SFR}$ probably have a higher filling factor of cold clouds and thus a higher Na~D covering fraction. \\item For the outflow component the Na D EW depends strongly on $\\Sigma_{\\rm SFR}$ and secondarily on $\\rm A_V$. We hypothesize that the star formation surface density determines the amount of material lofted above the disk, and the dust attenuation influences the survival of Na I in the clouds. \\item The covering factor of Na~I in the outflow increases with increasing dust attenuation while the optical depth of the absorption decreases. This suggests that Na D is able to survive in clouds with lower column density as the dust-to-gas ratio increases. \\item The outflow velocity (line centroid) does not depend strongly on any galaxy physical parameters over the limited dynamic range in physical properties probed by our study. There is some evidence for a very shallow trend, $v_{\\rm off} \\propto \\Sigma_{\\rm SFR}^{0.1}$, which is consistent with theoretical expectations for the velocity of a swept-up shell of gas at the point where it blows out of the disk. \\item The line width of the outflow component is sensitive to $\\Sigma_{\\rm SFR}$, $\\rm A_V$, and $M_{*}$. These parameters are likely to influence the acceleration of halo clouds (by ram or radiation pressure), the ability of clouds to remain neutral, and the length of time a cloud can remain pressure confined. \\end{itemize} The question of how galactic scale outflows are influenced by AGNs is also of great interest. In a companion paper, we will study the Na D absorption properties in SDSS galaxies hosting type~2 AGNs using methods similar to those described here. The Na~D ISM absorption line provides a convenient probe of galactic winds due to its location in the optical. However, the low EW width of the line in typical galaxy spectra and the strength of the Na~D stellar feature make interstellar Na~D absorption a challenge to measure accurately. The strong sensitivity of Na~I to ionization by OB stars is also a very serious concern, in particular for studies that aim to compute mass outflow rates from Na~I column densities and velocities. Higher redshift studies that make use of higher ionization transitions such as Mg~II~$\\lambda2796, 2803$ may ultimately be able to provide more robust constraints over a larger range in galaxy physical properties. Theoretical models of galaxy evolution have begun to incorporate galactic winds, using various ad-hoc prescriptions based on our knowledge of the cool gas. However, the majority the energy and newly-synthesized metals in the outflow resides in the hot X-ray emitting phase of the wind (D.~K. Strickland \\& D.~C. Dinge, 2010, in prep.). To truly understand the impact of galactic winds on the evolution of galaxies and the intergalactic medium, it is crucial to measure the chemical composition and velocity of the hot gas. This will require a high sensitivity X-ray imaging spectrometer such as the planned {\\it International X-ray Observatory.}" }, "1003/1003.5755_arXiv.txt": { "abstract": "For the use of Gamma-Ray Bursts (GRBs) to probe cosmology in a cosmology-independent way, a new method has been proposed to obtain luminosity distances of GRBs by interpolating directly from the Hubble diagram of SNe Ia, and then calibrating GRB relations at high redshift. In this paper, following the basic assumption in the interpolation method that objects at the same redshift should have the same luminosity distance, we propose another approach to calibrate GRB luminosity relations with cosmographic fitting directly from SN Ia data. In cosmography, there is a well-known fitting formula which can reflect the Hubble relation between luminosity distance and redshift with cosmographic parameters which can be fitted from observation data. Using the Cosmographic fitting results from the Union set of SNe Ia, we calibrate five GRB relations using GRB sample at $z\\leq1.4$ and deduce distance moduli of GRBs at $1.4< z \\leq 6.6$ by generalizing above calibrated relations at high redshift. Finally, we constrain the dark energy parameterization models of the Chevallier-Polarski-Linder (CPL) model, the Jassal-Bagla-Padmanabhan (JBP) model and the Alam model with GRB data at high redshift, as well as with the Cosmic Microwave Background radiation (CMB) and the baryonic acoustic oscillation (BAO) observations, and we find the $\\Lambda$CDM model is consistent with the current data in 1-$\\sigma$ confidence region. ", "introduction": "Since an intrinsic relation between the peak luminosity and the shape of the light curve of SNe Ia has been found,\\cite{Phillips1993} SNe Ia has now been taken as near-ideal standard candles for measuring the geometry and dynamics of the universe. However, the maximum redshift of the SNe Ia which we can currently use is only about 1.7. On the other hand, the redshift of the last scattering surface of the cosmic microwave background (CMB) is at $z=1091.3$.\\cite{Komatsu2010} Recently, Gamma-Ray Bursts (GRBs) were proposed to be a complementary probe to SNe Ia and CMB to explore the early universe. As the most intense explosions observed in the universe so far, GRBs are likely to occur in high-redshift range up to at least $z=8.2$.\\cite{Tanvir2009,Salvaterra2009} Moreover, there are several luminosity relations of GRBs between the spectral and temporal properties which have been extensive discussed, such as the isotropic energy ($E_{\\rm iso}$) - peak spectral energy ($E_{\\rm peak}$) relation,\\cite{Amati2002} the luminosity ($L$) - spectral lag ($\\tau_{\\rm lag}$) relation,\\cite{Norris2000} the $L$ - variability ($V$) relation,\\cite{Fenimore2000,Riechart2001} the $L$ - $E_{\\rm peak}$ relation,\\cite{Schaefer2003a,Yonetoku2004} the $L$ - minimum rise time ($\\tau_{\\rm RT}$) relation,\\cite{Schaefer2002} and the collimation-corrected energy ($E_{\\gamma}$) - $E_{\\rm peak}$ relation;\\cite{Ghirlanda2004a} as well as several multiple relations such as the $E_{\\rm iso}$ - $E_{\\rm peak}$ - $t_{\\rm b}$ relation,\\cite{Liang2005} where $t_{\\rm b}$ is the break time of the optical afterglow light curves; the $L$ - $E_{\\rm peak}$ - $T_{0.45}$ relation,\\cite{Firmani2006} where $T_{0.45}$ is the rest-frame ``high-signal'' timescale; and the $L$ - $E_{\\rm peak}$ - $\\tau_{\\rm lag}$ (or $\\tau_{\\rm RT}$) relation.\\cite{Yu2009} Many authors have made use of GRB luminosity indicators as standard candles at very high redshift beyond SNe Ia redshift range for cosmological research.\\cite{Schaefer2003b,Takahashi2003,Bloom2003,Dai2004,Ghirlanda2004a,Ghirlanda2004b,Friedman2005,Firmani2006,Firmani2007,Liang2005,Xu2005,Wang2006,Bertolami2006,Ghirlanda2006,Schaefer2007,Wright2007,Wang2007,Basilakos2008,Cuesta2008a,Cuesta2008b,Daly2008,Qi2008a,Qi2008b,Qi2009} Due to the lack of the sample at low redshift which are cosmology independent, to calibrate the empirical GRB relations, one usually needs to assume a particular cosmological model with certain model parameters as a $priori$. As a result, the so-called circularity problem could prevent the direct use of GRBs for cosmology.\\cite{Ghirlanda2006} Many of works treat the circularity problem with statistical approach which carried out a simultaneous fit of the parameters in the calibration curves and the cosmology.\\cite{Schaefer2003b,Firmani2005,Li2008,Amati2008,WangY2008} However, the circularity problem can not be circumvented completely by means of the statistical approaches for an input cosmological model is still required in doing the joint fitting. More recently, Liang \\textit{et al.} proposed a new method to calibrate GRB luminosity relations in a cosmological model-independent way.\\cite{Liang2008a} The motivation of this calibration method is that objects at the same redshift should have the same luminosity distance in any cosmology. Thus the luminosity distance of a GRB at a given redshift can be obtained by interpolating directly from the Hubble diagram of SNe Ia, therefore GRB relations can be calibrated without assuming a particular cosmological model and the Hubble diagram of GRBs has been constructed. Following this cosmology-independent GRB calibration directly from SNe Ia, the derived GRB Hubble diagram can be used to constrain cosmological models at high redshift avoiding circularity problem \\cite{Liang2008b,Capozziello2008,Izzo2009,Wei2009a,Wei2009b,Wang2009,Liang2010,WLiang2010}. Capozziello \\& Izzo firstly used two GRB relations calibrated with the so-called Liang method to derive the related cosmography parameters which related to the derivatives of the scale factor.\\cite{Capozziello2008} Liang \\textit{et al.} combined the updated distance moduli of GRBs obtained by the interpolating method with the joint data to find the contribution of GRBs to the joint cosmological constraints in the confidence regions of cosmological parameters, and reconstructed the acceleration history of the universe with the distance moduli of SNe Ia and GRBs.\\cite{Liang2010} On the other hand, besides the interpolation method, the luminosity distance of a GRB can also be obtained directly from SNe Ia data by other mathematical approach. Liang \\& Zhang has proposed another approach to calibrate GRB relations by using an iterative procedure which is a non-parametric method in a model independent manner to reconstruct the luminosity distance at any redshift from SNe Ia.\\cite{Liang2008b} Similar to the interpolation method, Cardone \\textit{et al.} constructed an updated GRBs Hubble diagram calibrated by local regression from SNe Ia.\\cite{Cardone2009} Kodama \\textit{et al.} has proposed that the $L-E_{\\rm peak}$ relation can be calibrated with one empirical formula fitted from the luminosity distance of SNe Ia.\\cite{Kodama2008} However, according to the formula fitting approach, various possible formula can be fitted from the SNe Ia data which could give different calibration results of GRB relations. As the cosmological constraints from GRBs are sensitive to GRBs calibration results, and the fitting procedure depends seriously on the choice of the formula, the reliability of this method should be tested carefully. In other words, we should find one certain formula which is totally independent of any cosmological models and could accurately evaluate the Hubble relation. In Cosmography,\\cite{Visser2004} there is a well-known formula reflecting the Hubble relation between luminosity distance and redshift which can be extracted directly from basic cosmological principles and observation data, with cosmography parameters (the deceleration, jerk and snap parameters: $q$, $j$, and $s$) which are only related to the derivatives of the scale factor without any priori assumption on the underlying cosmological model. Recently, several authors have already used the cosmographic parameters fitting from SNe Ia and/or GRBs dataset to constrain cosmological parameters.\\cite{Cattoen2007,Cattoen2008,Capozziello2008,Vitagliano2010} If viewing this point from another angle, the cosmographic formula can be considered as a perfect fitting function to calibrate the GRB relations using SNe Ia data, as long as we take the same assumption that objects at the same redshift should have the same luminosity distance in any cosmology. In this paper, instead of the interpolation method using in Ref.~\\refcite{Liang2008a}, we propose another new approach to calibrate GRB luminosity relations with cosmographic fitting from SNe Ia data. The structure of this paper is arranged as follows. In section 2 we give a brief review of the cosmographic Hubble relation between luminosity distance and redshift. In section 3, we calibrate five GRB luminosity relations with cosmographic fitting results from SNe Ia data. In section 4, we construct the Hubble diagram of GRBs obtained by using the cosmographic methods and constrain the dark energy parameterization models of the Chevallier-Polarski-Linder (CPL) model,\\cite{Chevallier2001,Linder2003} the Jassal-Bagla-Padmanabhan (JBP) model\\cite{Jassal2004} and the Alam model\\cite{Alam2003} with GRB data at high redshift, as well as with the Cosmic Microwave Background radiation (CMB) and the baryonic acoustic oscillation (BAO) observations. Conclusions and discussions are given in section 5. ", "conclusions": "Due to the lack of the GRB sample at low redshift, there has been a so-called circularity problem which can always be a obstacle for applying GRBs data to constrain cosmological parameters. Based on the basic assumption that objects at the same redshift should have the same luminosity distance, Liang \\textit{et al.} proposed a new method to calibrate GRB relations in a completely cosmology-independent way, namely obtaining the distance modulus of a GRB by interpolating from the Hubble diagram of SNe Ia and then calibrate the GRB relations with these calculated distance moduli\\cite{Liang2008a}. There is a well-known fitting formula in cosmography, which can reflect the Hubble relation between luminosity distance and redshift with cosmographic parameters which can be fitted from SNe Ia. In this work, we propose another approach to calibrate GRB luminosity relations with cosmography fitting from SNe Ia data. We adopted the fitting results from the Union set of SNe Ia for the so-called Cosmography I and II,\\cite{Vitagliano2010} and calibrate five GRB relations by this cosmographic fitting method. The calibration results obtained using two cosmographic fitting are fully consistent with each other. Assuming that GRB luminosity relations do not evolve with redshift, we obtained the distance modulus of the GRB data at higher redshift $1.40$ such a theory possesses classical oscillons that are characterized by a small dimensionless parameter $\\epsilon$ and have a mass given by $M_{\\mbox{\\tiny{osc}}}\\sim\\frac{m}{\\lambda}l(\\epsilon) = \\frac{m_\\phi}{\\lambda\\,\\hbar}l(\\epsilon)$, where $l(\\epsilon)=\\epsilon^{2-d}$ in the standard expansion that we will describe in Section \\ref{Classical}, but can scale differently in other models (e.g., see \\cite{Fodor:2009kg}). So if $\\lambda\\,\\hbar\\ll l(\\epsilon)$, then $M_{\\mbox{\\tiny{osc}}}\\gg m_\\phi$. In such a regime we expect that various properties of the oscillon are well described classically, such as its size and shape. In this work we examine whether the same is true for the oscillon lifetime. We find that although the oscillon lifetime is exponentially long lived classically, it has a power law lifetime in the quantum theory controlled by the ``effective $\\hbar$\" -- in the above case this is $\\lambda_3^2\\,\\hbar$ or $\\lambda_4\\,\\hbar$. In this paper we treat the oscillon as a classical space-time dependent background, as defined by the $\\epsilon$ expansion, and quantize field/s in this background. We go to leading order in $\\hbar$ in the quantum theory. We find that oscillon's decay through the emission of radiation with wavenumbers $k=\\mathcal{O}(m)$. We show that the decay is a power law in $\\epsilon$ and the couplings, and we explain why this is exponentially suppressed classically. Our analysis is done for both single field theories, where the emitted radiation typically grows at a linear rate corresponding to $3\\,\\phi\\to2\\,\\phi$ or $4\\,\\phi\\to2\\,\\phi$ annihilation processes, and for multi-field theories, where the emitted radiation often grows at an exponential rate corresponding to $\\phi\\to2\\,\\chi$ decay or $2\\,\\phi\\to2\\,\\chi$ annihilation processess. We calculate the quantum decay rates in several models, which is supported by numerical investigations, but our work is also qualitative and of general validity. We also comment on collapse instabilities for $k=\\mathcal{O}(\\epsilon\\,m)$ modes, whose existence is model dependent. The outline of this paper is as follows: In Section \\ref{Classical} we start with a review of classical oscillons and describe their exponentially suppressed decay in Section \\ref{ExpSmall}. In Section \\ref{Quantum} we outline the semi-classical quantization of oscillons and derive the decay rate of oscillons in Section \\ref{HigherQuantum}. Having started with single field models, we move on to examine the effects of coupling to other fields in Section \\ref{TwoField}. In Section \\ref{ExpLinear} we discuss when the decay products grow linearly in time and when it is exponential. Here we demonstrate that coupled fields can achieve (depending on parameters) explosive energy transfer, which may have some cosmological relevance. We comment on collapse instabilities in Section \\ref{LeadingQuantum} and conclude in Section \\ref{Conclusions}. ", "conclusions": "We have found that even though an oscillon can have a mass that is much greater than the mass of the individual quanta, the classical decay can be very different to the quantum decay (this point does not appear to have been appreciated in the literature, for instance see the concluding sections of Refs.~\\cite{Graham:2006vy,Graham:2007ds}). The radiation of both classical and quantum oscillons can be understood in terms of forced oscillator equations, see eqs.~(\\ref{classFO},\\,\\ref{vq1}). We derived the frequency and wavenumber of the outgoing radiation, which were both $\\mathcal{O}(1)$ in natural units. Since a classical oscillon has a spread which is $\\mathcal{O}(1/\\epsilon)$ in position space, it has a spread which is $\\mathcal{O}(\\epsilon)$ in $k$-space. Its Fourier modes are therefore exponentially small at the radiating wavenumber and hence such radiation is exponentially suppressed. In the quantum theory, there simply {\\em cannot} be modes whose amplitudes are exponentially suppressed. Instead, zero-point fluctuations ensure that all modes have at least $\\mathcal{O}(\\hbar)$ amplitude-squared due to the uncertainty principle. We derived a formula for the quantum lifetime of an oscillon $\\sim 1/(\\lambda\\,\\hbar\\,\\epsilon^p)$. The power $p$ is model dependent: $p=4$ in the $\\phi^3+\\ldots$ theory (or $-\\phi^4+\\phi^5+\\ldots$), and $p=6$ in the $-\\phi^4+\\phi^6+\\ldots$ theory. Through a Floquet analysis, we explained why the growth of perturbations of small amplitude oscillons is linear in time, as opposed to exponential. The dimensionless $\\lambda\\,\\hbar$ controls the magnitude of the decay rate, as it should for a leading order in $\\hbar$ analysis. For example, the Standard Model Higgs potential has $\\lambda\\,\\hbar\\sim (m_{\\mbox{\\tiny{H}}}/v_{\\mbox{\\tiny{EW}}})^2\\sim 0.1\\,(m_{\\mbox{\\tiny{H}}}/100\\,\\mbox{GeV})^2$ and so this is not very small. On the other hand, the effective $\\hbar$ of the QCD axion potential is $\\lambda\\,\\hbar\\sim(\\Lambda_{\\mbox{\\tiny{QCD}}}/f_a)^4\\sim 10^{-48}\\,(10^{10}\\,\\mbox{GeV}/f_a)^4$ and so oscillons formed from axions, called ``axitons\" in \\cite{Kolb}, are governed by classical decay (ignoring coupling to other fields). We further considered the fate of an oscillon that is coupled to a second scalar $\\chi$ and found it to either decay or annihilate with a growth in $\\chi$ that can be exponentially fast, depending on parameters. Since oscillons may form substantially in the early universe \\cite{Farhi:2007wj,Copeland:1995fq} this may give rise to interesting phenomenology. At the very least, it presents a plausible cosmological scenario in which a parametric pump field exists that is qualitatively different to the homogeneous oscillations of the inflaton during p/reheating. This is a form of parametric resonance: explosive transfer of energy from a localized clump into bosonic daughter fields. We expect decay into fermions to be quite different (for discussion in the context of Q-balls, see \\cite{Cohen:1986ct}). This may have some cosmological relevance. It appears that if a field has a perturbative decay channel, then the oscillon will eventually decay through it. This is important because we expect most fields in nature to be perturbatively unstable, including the inflaton, p/reheating fields, Higgs, and most fields beyond the standard model. A good exception is dark matter. This conclusion may seem surprising given that the oscillon is a bound state of particles with a finite binding energy \\cite{Dashen:1975hd}. However, oscillons are formed from fields whose particle number is not conserved. One could imagine a situation in which $m_\\phi$ is only slightly greater than $2\\,m_\\chi$, and in this case the oscillon's binding energy may prevent direct decays into $\\chi$'s, but this requires fine tuning and will not forbid $2\\,\\phi\\to2\\,\\chi$ or $3\\,\\phi\\to2\\,\\phi$ or $4\\,\\phi\\to2\\,\\phi$ annihilations.% We conclude that in many scenarios an individual oscillon's lifetime will be shorter than the age of the universe at the time of production (this may prevent individual oscillons from having cosmological significance in such cases). Exceptions include the grand unified theory era \\cite{Copeland:1995fq}, inflation \\cite{Gleiser:2006te}, and axitons produced at the QCD phase transition \\cite{Kolb}. An interesting question for further study is whether oscillons can form and then decay, and then form again repeatedly, like subcritical bubbles in hot water. It is not implausible that such a process could continue over long time scales for cosmic temperatures of order the field's mass; similar to the production and disappearance of unstable particles in a relativistic plasma. This may modify cosmological thermalization." }, "1003/1003.4753_arXiv.txt": { "abstract": "There exists several modified gravity theories designed to reproduce the empirical Milgrom's formula (MOND). Here we derive analytical results in the context of the static weak-field limit of two of them (BIMOND, leading for a given set of parameters to QUMOND, and TeVeS). In this limit, these theories are constructed to give the same force field for spherical symmetry, but their predictions generally differ out of it. However, for certain realizations of these theories (characterized by specific choices for their free functions), the binding potential-energy of a system is increased, compared to its Newtonian counterpart, by a constant amount independent of the shape and size of the system. In that case, the virial theorem is exactly the same in these two theories, for the whole gravity regime and even outside of spherical symmetry, although the exact force fields are different. We explicitly show this for the force field generated by the two theories inside an elliptical shell. For more general free functions, the virial theorems are however not identical in these two theories. We finally explore the consequences of these analytical results for the two-body force. ", "introduction": "The current dominant paradigm is that galaxies are embedded in halos of cold dark matter. However, one observes that for gravitational accelerations below $a_0 \\sim 10^{-10}$~ms$^{-2}$, the total gravitational attraction $g$ in galaxy disks approaches $(g_N a_0)^{1/2}$ where $g_N$ is the usual Newtonian gravitational field as calculated from the observed distribution of baryonic matter. The successes of this recipe in galaxies could be an emergent phenomenon, linked with the complex feedback between baryons and cold dark matter, but a more radical explanation of these successes is a modification of gravity on galaxy scales: this paradigm is known as modified Newtonian dynamics \\citep[MOND, ][]{Mil83}. More precisely, within this paradigm, the Newtonian acceleration $\\vec{g}_N$ produced by the visible matter is linked to the true gravitational acceleration $\\vec{g}$ by means of an interpolating function $\\mu(x)$: \\begin{equation} \\mu\\left(g/a_{0}\\right)\\vec{g} = \\vec{g}_{N}, \\label{eq:1} \\end{equation} where $\\mu(x) \\sim x$ for $x \\ll 1$ and $\\mu(x) \\sim 1$ for $x \\gg 1$ (and $g=|\\vec{g}|$), or equivalently by means of an interpolating function $\\nu(y)$: \\begin{equation} \\vec{g} = \\nu\\left(g_N/a_{0}\\right)\\vec{g}_{N}, \\label{eq:2} \\end{equation} where $\\nu(y) \\sim y^{-1/2}$ for $y \\ll 1$ and $\\nu(y) \\sim 1$ for $y \\gg 1$. However, these expressions cannot be exact outside of spherical symmetry, since they do not respect usual conservation laws. There exists various flavors of modified gravity theories reproducing this relation in spherical symmetry, but all making slightly different predictions outside of it. For instance, in the Newtonian static weak-field limit of the generalized Einstein-Aether theories\\cite{Zlosnik,Halle,Zhao2008}, the gravitational potential $\\Phi$ obeys the following modified Poisson equation\\cite{BM84}: \\begin{equation} \\nabla \\cdot \\left[ \\mu(|{\\mathbf\\grad} \\Phi|/a_0) {\\mathbf\\grad} \\Phi \\right]=4\\pi G\\rho. \\end{equation} On the other hand, in Bekenstein's Tensor-Vector-Scalar (TeVeS) multifield theory\\cite{Bek04}, the gravitational potential in the static weak-field limit can be expressed as: \\begin{equation} \\Phi=\\Phi_N+\\phi, \\end{equation} where $\\Phi_N$ is the Newtonian potential obeying the usual Poisson equation, and $\\phi$ is a scalar field obeying an equation similar to Eq.~(3), but with a different $\\mu$-function \\citep[see, e.g.,][]{F07}, which can simply be, at least far from the strongest gravity regime, $\\mu(x)=x$. Finally, another possibility to recover Eqs.~(1) and (2) in the spherically symmetric static weak-field limit is the new BIMOND theory \\cite{Mil09_bim,Clif} where we can have $\\Phi=\\Phi_N+\\phi$, as in Eq.~(4), with $\\phi$ obeying \\citep[see also QUMOND,][]{Mil09_qu}: \\begin{equation} \\nabla^2 \\phi=\\nabla \\cdot \\left[ \\nu(|{\\mathbf\\grad} \\Phi_N|/a_0) {\\mathbf\\grad} \\Phi_N \\right], \\end{equation} where, e.g., $\\nu(y) = y^{-1/2}$ far away from the strongest gravity regime. Here, we intend to compare these various implementations of the MOND paradigm at the purely theoretical level, in the same philosophy as \\cite{Dai,Matsuo}. A useful way to compare the theoretical implications of these various theoretical frameworks is to compare the virial theorem ensuing from them. Indeed, the virial theorem is a very useful tool, e.g. to compute the 2-body forces in these theories. The scalar form of the virial theorem has been computed for Eq.(3) in the deep-MOND limit by \\cite{Mil94, Mil97} and in the deep-QUMOND case \\cite{Mil09_qu}. Hereafter, we extend this study to the whole intermediate regime (not only deep-MOND) in the BIMOND and TeVeS frameworks, and show that for peculiar choices of the $\\mu$ and $\\nu$ functions, the scalar form of the virial theorem is identical in these two theories and independent of the size and shape of the system (Sect.~II), although the exact force fields are different. We explicitly show how the force fields differ in the case of a test particle sitting inside an elliptical shell of matter (Sect.~III). We also explore the consequences for the 2-body force (Sect.~IV). ", "conclusions": "We conclude that, in general, the TeVeS and BIMOND force fields and virial theorems are different, except in spherical symmetry (with $\\nu=\\mu^{-1}$). However, for a specific choice of the $\\mu$ and $\\nu=\\mu^{-1}$ functions (observationally only valid in the intermediate gravity regime), the virial theorems are identical even outside of spherical symmetry. What is more, the potential energy of the system can then be expressed analytically, for the whole gravity regime, as a sum of the Newtonian potential energy and of a constant term proportional to the 3/2 power of the total mass of the system. However, while the virial theorems are then identical, the exact force fields are still different (i.e. the gravities have different curl-fields), as we showed in the case of elliptical shells. In this specific case, we can also get an analytic expression for the 2-body force under the approximation that the two bodies are very far apart compared to their internal sizes \\citep[see also][]{Mil94}. Since $2K+W_N-(2/3)\\sqrt{GM^3a_0}=0$, since the kinetic energy can be separated into the orbital energy $K_{\\rm orb}=M_1M_2 v^2_{\\rm rel}/(2M)$ and the internal energy of the bodies $K_{\\rm int}=-\\Sigma (1/2)W_{N_i}+\\Sigma (1/3) \\sqrt{GM_i^3a_0}$, and since the Newtonian potential energy can be separated into the interaction term (the mutual potential energy) and the internal Newtonian potential energies of the bodies $\\Sigma W_{N_i}$, we get: \\beq \\frac{M_1M_2 v^2_{\\rm rel}}{M} = \\frac{GM_1M_2}{r_{12}}+\\frac{2}{3} \\left[ (GM^3a_0)^{1/2} - \\sum_i (GM_i^3a_0)^{1/2} \\right]. \\eeq We can then assume an approximately circular velocity such that the 2-body force can be written \\bey & & \\vec{F}_{12} = M_1\\vec{a_1}=\\frac{M_1M_2v^2_{\\rm rel}\\vec{r}_{12}}{Mr^2_{12}} = \\frac{GM_1M_2\\vec{r}_{12}}{r^3_{12}} \\\\ & & + \\frac{2}{3}\\left[ 1 - \\sum_{i=1}^{i=2} \\left(\\frac{M_i}{M}\\right)^{3/2} \\right]\\frac{(GM^3a_0)^{1/2}\\vec{r}_{12}}{r^2_{12}}. \\nonumber \\eey For this specific choice of free function, the TeVeS and BIMOND 2-body forces are thus the same, provided the two bodies are very small compared to their mutual distance, but if they are not, the force will be different since we showed that the force fields are different in general. As a final remark, let us stress that this 2-body force depends on the mass ratio of the binary, meaning that the orbital period of an equal-mass binary would differ from that of a binary of extreme ratio but with the same total mass and separations, which could cause stars to segregate according to their masses." }, "1003/1003.1340_arXiv.txt": { "abstract": "The most usual tracer of molecular gas is line emission from CO. However, the reliability of that tracer has long been questioned in environments different from the Milky Way. We study the relationship between H$_2$ and CO abundances using a fully dynamical model of magnetized turbulence coupled to a chemical network simplified to follow only the dominant pathways for H$_2$ and CO formation and destruction, and including photodissociation using a six-ray approximation. We find that the abundance of H$_2$ is primarily determined by the amount of time available for its formation, which is proportional to the product of the density and the metallicity, but insensitive to photodissociation. Photodissociation only becomes important at extinctions under a few tenths of a visual magnitude, in agreement with both observational and prior theoretical work. On the other hand, CO forms quickly, within a dynamical time, but its abundance depends primarily on photodissociation, with only a weak secondary dependence on H$_2$ abundance. As a result, there is a sharp cutoff in CO abundance at mean visual extinctions $A_{\\rm V} \\simless 3$. At lower values of $A_{\\rm V}$ we find that the ratio of H$_2$ column density to CO emissivity $X_{\\rm CO} \\propto A_{\\rm V}^{-3.5}$. This explains the discrepancy observed in low metallicity systems between cloud masses derived from CO observations and other techniques such as infrared emission. Our work predicts that CO-bright clouds in low metallicity systems should be systematically larger or denser than Milky Way clouds, or both. Our results further explain the narrow range of observed molecular cloud column densities as a threshold effect, without requiring the assumption of virial equilibrium. ", "introduction": "Observed star formation takes place within giant molecular clouds (GMCs), so understanding how these clouds form and evolve is a key step towards understanding star formation on galactic scales. The main chemical constituent of any GMC is molecular hydrogen (H$_{2}$). However, it is extremely difficult to directly observe this molecular hydrogen {\\em in situ}. Radiative transitions in H$_{2}$ are weak, owing to the molecule's lack of a permanent dipole moment. Moreover, the lowest lying rotational energy levels of H$_{2}$ are widely spaced, and so are very rarely excited in gas with the temperatures typical of GMCs, $T \\sim 10$--20~K. For this reason it is common to use emission from carbon monoxide (CO), the second most abundant molecule in GMCs, as a proxy for the H$_{2}$. In order to use CO as a proxy for H$_{2}$, however, it is necessary to understand the relationship between the distributions of these two molecules. Both are readily dissociated by the absorption of ultraviolet photons with energies below the Lyman limit of atomic hydrogen, so in low density, low extinction gas such as the warm neutral component of the interstellar medium (ISM), the abundances of both molecules will be small. However, the two molecules form rather differently: H$_{2}$ forms predominantly on the surface of dust grains \\citep{gs63}, while CO forms almost exclusively in the gas phase, via any one of a number of chains of ion-neutral or neutral-neutral reactions (see e.g.\\ \\citealt{sd95} for a useful summary of CO formation chemistry). Moreover, H$_{2}$ can protect itself from ultraviolet radiation via self-shielding, which becomes effective for relatively low H$_{2}$ column densities \\citep{db96}. The corresponding process for CO is less effective \\citep{lee96}, requiring a higher column density, and the low abundance of carbon relative to hydrogen in the ISM means that the required column density typically corresponds to a situation in which a large fraction of the available carbon is already locked up in the form of CO. For these reasons, we would expect the CO/H$_{2}$ ratio to vary through a cloud, and this expectation is confirmed by detailed models of slab-like or spherical clouds \\citep[e.g.][]{pdr07}. These find that the transition from atomic hydrogen to molecular hydrogen occurs at a point closer to the cloud surface than the transition from ionized carbon, via neutral atomic carbon, to CO. Although spatial variations in the CO/H$_{2}$ ratio within a GMC would appear to make it difficult to use CO as a proxy for H$_{2}$, observations of Galactic GMCs show that in fact there appears to be a good correlation between the integrated intensity of the $J = 1 \\rightarrow 0$ rotational transition line of $^{12}$CO and the H$_{2}$ column density \\citep[see e.g.][]{dick78,sand84,sol87,sm96,dame01}. A number of independent studies have shown that GMCs in the Galactic disk all have CO-to-H$_{2}$ conversion factors that are approximately \\begin{equation} X_{\\rm CO} = \\frac{N_{\\rm H_{2}}}{W_{\\rm CO}} \\simeq 2 \\times 10^{20} {\\rm cm^{-2} \\: K^{-1} \\: km^{-1} \\: s}, \\end{equation} where $W_{\\rm CO}$ is the velocity-integrated intensity of the CO $J = 1 \\rightarrow 0$ emission line, averaged over the projected area of the GMC, and $N_{\\rm H_{2}}$ is the mean H$_{2}$ column density of the GMC, averaged over the same area. Although the former can be directly observed, the latter cannot. However, for nearby clouds, it can be inferred from measurements of the diffuse $\\gamma$-ray flux produced by interactions between high energy cosmic rays and atomic hydrogen, atomic helium and H$_{2}$. The $\\gamma$-ray flux along a given line of sight depends on the total hydrogen column density along that line of sight. Since the atomic hydrogen column density can be measured via its 21~cm emission, the H$_{2}$ column density can be inferred. As the values for $X_{\\rm CO}$ obtained in this way are consistent with those obtained by assuming that the GMCs are in virial equilibrium and using the observed linewidth-size relation to compute the cloud mass \\citep[e.g.][]{sol87}, it is generally accepted that CO emission is indeed a good proxy for H$_{2}$ mass in nearby GMCs. However, the issue of the environmental dependence of $X_{\\rm CO}$ remains highly contentious. In gas with a lower metallicity, or a higher ambient UV radiation field, CO photodissociation will be more effective, and the amount of CO in the cloud will be smaller, with the CO-rich gas occupying a smaller volume than in the Galactic case (e.g.\\ \\citealt{mb88}, Molina et~al., in prep.). The mean H$_{2}$ abundance may also be smaller, but the greater role played by H$_{2}$ self-shielding means that we do not expect the H$_{2}$ abundance to be nearly as sensitive to changes in the metallicity. Because of this, there are good theoretical reasons to expect the relationship between CO emission and H$_{2}$ mass to change as we change the metallicity. However, observational efforts to test this yield inconsistent results. Measurements of $X_{\\rm CO}$ that assume that extragalactic GMCs are in virial equilibrium and use a virial analysis to determine the cloud mass generally find values for $X_{\\rm CO}$ that are similar to those obtained in the Milky Way, with at most a weak metallicity dependence \\citep{wilson95,ros03,bolatto08}. On the other hand, measurements that constrain GMC masses using other techniques that do not depend on the CO emission, such as by measuring the far-infrared dust emission, consistently find values for $X_{\\rm CO}$ that are much larger than the Galactic value and that are suggestive of a strong metallicity dependence \\citep{israel97, rubio04, leroy07, leroy09}. Numerical simulations provide us with one way to address this observational dichotomy. If we can understand the distribution of CO and H$_{2}$ in realistic models of GMCs, then we may begin to understand why the different types of observation give such different results. However, until very recently, the ability of simulations to address this issue has been quite limited. Sophisticated treatments of gas and grain chemistry, radiative heating and cooling, and radiative transfer have been developed to model GMCs \\citep[see, for instance, the paper by][which compares results from a number of popular codes]{pdr07}, but the very complexity of these models limits their applicability to highly simplified geometries: typically the adopted cloud models are one-dimensional, assuming either spherical symmetry \\citep[e.g.][]{kosma96} or a semi-infinite, uniform slab \\citep[e.g.][]{meudon06}. Moreover, the modelling often assumes that the chemistry of the clouds is in equilibrium, which is valid only if all of the chemical timescales are much shorter than any dynamical timescale associated with the GMCs. Since real GMCs are observed to be highly inhomogeneous, and to be dominated by supersonic turbulent motions \\citep[see e.g.][and references therein]{mk04}, the applicability of the results from these simple spherical or slab models to real clouds is open to question. On the other hand, previous attempts to accurately model the turbulent dynamics of GMCs, and the inhomogeneous and intermittent density structure that is created by this turbulence typically have avoided modelling the cloud chemistry, rendering them also of limited use in addressing this question. The few studies that have attempted to model both the turbulence and the chemistry self-consistently \\citep[e.g.][]{jou98,les07,god09} have typically done so by reducing the dimensionality of the problem. For instance, \\citet{god09} present results on the chemical evolution of gas passing through magnetized, turbulent vortices, using two-dimensional simulations of these vortices, and then construct models for lines of sight through diffuse clouds by summing up the contributions from a number of different vortices. However, it is unclear whether the results obtained in this fashion are the same as those that would be obtained using a fully three-dimensional model for the turbulence. In a previous paper \\citep[][hereafter Paper I]{glo10}, we presented the first results from a project that aims to combine detailed chemical modelling with three-dimensional magnetohydrodynamical turbulence simulations in order to self-consistently model both the chemistry and the turbulent dynamics of the gas within a GMC. We showed that it is now computationally feasible to attempt this. Provided one makes a few simplifying assumptions regarding the extent of the chemistry to be treated, and the treatment of the UV radiation field, it is possible to model both the H$_{2}$ and the CO chemistry of a GMC with acceptable accuracy within a moderate resolution dynamical simulation. Paper~I presented results from a few trial simulations performed using only one set of cloud properties (mean density, metallicity, etc). In this paper, we present results from a much larger set of simulations that examine the sensitivity of H$_{2}$ and CO formation to changes in the mean densities and metallicities of GMCs, and explore the consequences of this for the CO-to-H$_{2}$ conversion factor and the dynamical structure of the observed clouds. In section~\\ref{sims}, we discuss our numerical approach and the initial conditions used for our simulations. In section~\\ref{res}, we present the main results of our simulations, and use them to derive the approximate dependence of the CO-to-H$_{2}$ conversion factor on cloud properties. In section~\\ref{dis}, we discuss the major consequences of our findings. In section~\\ref{cave}, we discuss a few potential caveats regarding our approach, and show why they are unlikely to significantly affect our main results. Finally, in section~\\ref{summ}, we close with a brief summary. ", "conclusions": "\\label{dis} We can draw a number of immediate conclusions from the results that we have presented in the previous section. First, it seems plain that when we talk about the formation of a molecular cloud, we should draw a distinction between the formation of a cloud in which the hydrogen is primarily in molecular form and the formation of a cloud from which CO emission can readily be detected, since the conditions for the former are not the same as those for the latter. To form an H$_{2}$-dominated cloud, the cloud must have a mean visual extinction of a few tenths of a magnitude (cf.\\ \\citealt{kmt08}, who find a critical value of approximately 0.5 for a spherical molecular cloud). More importantly, however, the cloud must survive for long enough to convert most of its atomic hydrogen into H$_{2}$. The time required is a strong function of $n_{0} {\\rm Z}$, the product of the mean density of the cloud and its metallicity, although the dependence is not linear, and the time required also depends upon the strength and nature of the turbulence (\\citealt{gm07b}; Milosavljevic et al., in prep.) On the other hand, to make a CO-bright cloud, we need a significantly higher mean extinction, $\\meanAV \\simgreat 3$, but do not require so much time, since the CO abundance in most of the gas comes into equilibrium within 1--2~Myr. In practice, given plausible cloud sizes, clouds that have a high enough mean extinction to be CO-bright will also be dense enough to form H$_{2}$ relatively quickly, but the converse is not necessarily true. Second, the existence of a visual extinction threshold above which clouds become CO-bright provides a simple explanation for the observation that in low metallicity systems cloud masses derived from CO observations are significantly smaller than those derived from techniques that do not depend on CO, such as infrared emission \\citep[e.g.][]{israel97, leroy07, leroy09}. In these systems, the CO largely traces the regions of a cloud (or cloud complex) that have mean extinctions greater than 2, but does not trace the H$_{2}$ in the surrounding envelope, which may be considerably more extensive and may contain a large fraction of the total mass of the cloud. Third, the fact that the CO fraction typically reaches equilibrium on a timescale comparable to or shorter than the crossing time of the cloud, and significantly shorter than the time required to assemble the cloud from warm atomic gas implies that CO emission will ``switch-on'' rapidly during the assembly of the cloud, as hypothesized by \\citet{hbb01}. In the time it takes compressive flows or gravity to double or triple the mean extinction of a cloud with $\\meanAV \\sim 1$ its CO content and luminosity can increase more than a hundredfold. The cloud will therefore quickly move from being effectively unobservable in CO to being readily observable. On the other hand, the long H$_{2}$ formation timescale implies that molecular clouds need not have equilibrium H$_{2}$ abundances, even once they become CO-bright, contrary to what is often assumed \\citep[see e.g.][]{kmt08,mk10}. How far from equilibrium the H$_{2}$ fractions are will depend on the cloud properties (mean density, metallicity, rms turbulent velocity, etc.), the age of the cloud, and also on the range of densities considered within the cloud, since overdense regions will reach equilibrium much faster than underdense regions. Fourth, we would expect the molecular clouds (or regions thereof) in low metallicity systems that we observe to be CO-bright to also be systematically larger and/or denser than their counterparts in the Milky Way. This conclusion follows from the fact that if $\\meanAV \\simgreat 3$ is required for CO to form, and we decrease ${\\rm Z}$, then we must increase either $n_{0}$ or $L$ to compensate, since $\\meanAV \\propto n_{0} L {\\rm Z}$. This conclusion also lends itself to a relatively simple observational test. If we observe a giant molecular cloud of size $L$ in a low metallicity system (where $L$ is the total size of the cloud, not of the CO-bright region), and compare it with a cloud of a similar size in the Milky Way, then the former should be systematically denser than the latter. This systematic difference in density may be detectable by examination of the CO(2-1)/CO(1-0) or CO(3-2)/CO(1-0) line ratios, or the ratio of of HCN emission to CO emission, all of which will be larger in denser systems. Finally, our results also suggest a relatively simple explanation for the fact that GMCs observed in the Milky Way span only a small range in column densities \\citep{blitz07}. The requirement that $\\meanAV \\simgreat 3$ for a CO-bright cloud implies that any cloud identified as a GMC must have a minimum mean column density of around $60 \\mbox{ M}_{\\odot} \\mbox{ pc}^{-2}$, since clouds with lower column densities will typically not contain much CO. The distribution of the mean column densities of GMCs above this minimum value depends upon the distributions of the size and mean density of the clouds. However, the cloud size distribution is typically found to be a sharply decreasing function of cloud size $L$. For example, \\citet{hcs01} find $n(L){\\rm d}L \\propto L^{-3.2 \\pm 0.1} {\\rm d}L$ for GMCs in the outer Galaxy. If the mean density of the GMCs spans only a small range of values, then the distribution of column densities $n(\\Sigma){\\rm d}\\Sigma$ will be a sharply decreasing function of $\\Sigma$. If we combine this fact with the existence of a minimum column density threshold for a CO-bright GMC, then it implies that most observed GMCs will have column densities that are close to the threshold. This argument relies on the mean cloud density not being a strong function of the size of the cloud. It is possible to show that this is a good assumption for at least the \\citet{hcs01} cloud sample. Given a cloud size distribution $n(L) \\propto L^{-3.2 \\pm 0.1}$, and a constant cloud mean density, then the cloud mass distribution is simply \\begin{eqnarray} n(M){\\rm d}M & = & n(L) \\frac{{\\rm d}L}{{\\rm d}M} {\\rm d}M, \\nonumber \\\\ & \\propto & L^{-3.2 \\pm 0.1} L^{-2} \\nonumber \\\\ & \\propto & M^{-(5.2 \\pm 0.1) / 3} \\nonumber \\\\ & \\propto & M^{-1.73 \\pm 0.03}, \\end{eqnarray} where we have assumed only that $M \\propto L^{3}$, consistent with our adoption of a constant cloud mean density. For comparison, the cloud mass distribution inferred by \\citet{hcs01} for this sample of clouds was $n(M) \\propto M^{-1.8 \\pm 0.03}$, demonstrating that even if the mean density of the clouds is not constant, it can at most be a very weak function of cloud size $L$. If this argument is correct, it suggests that most Galactic GMCs should have column densities close to $60 \\: {\\rm M_{\\odot}} \\: {\\rm pc^{-2}}$. This value is a factor of three smaller than the value of $170 \\: {\\rm M_{\\odot}} \\: {\\rm pc^{-2}}$ determined by \\citet{sol87} that is often taken to be canonical, but is in good agreement with the value recently derived by \\citet{hkdj09} from their re-examination of the \\citeauthor{sol87} clouds using new data from the Boston University-FCRAO Galactic Ring Survey. We note in conclusion that this argument makes no assumptions about the dynamical equilibrium of the clouds. This casts doubt on the frequently-made assumption that the small scatter in observed GMC column densities implies that the GMCs are in virial equilibrium. \\label{summ} We finish our discussion here by summarizing the key results of this paper. We have shown that the H$_{2}$ abundance in turbulent molecular clouds is controlled primarily by the time taken to form the H$_{2}$, and that it is relatively insensitive to the effects of UV photodissociation. Photodissociation becomes significant for determining the H$_{2}$ abundance only below a visual extinction threshold of a few tenths of a magnitude, in good agreement with previous observational and theoretical work. On the other hand, CO forms rapidly, but is strongly affected by photodissociation. The mean CO abundance falls off rapidly with decreasing mean extinction, particularly in clouds with $\\meanAV \\simless 3$. We have demonstrated that the CO-to-H$_{2}$ conversion factor is also determined primarily by the mean extinction of the cloud, and that it is almost constant for $\\meanAV \\simgreat 3$, but falls off as $X_{\\rm CO} \\propto A_{\\rm V}^{-3.5}$ for $\\meanAV \\simless 3$. Furthermore, only the clouds with visual extinctions greater than a few have sufficiently large CO integrated intensities to be detected by current observations. This therefore suggests a simple explanation for the discrepancy observed in low metallicity systems between cloud masses determined by CO observations and those determined by non-CO tracers, such as IR emission. CO observations inevitably select CO-bright clouds that have mean extinctions large enough to place them in the regime where $X_{\\rm CO}$ is approximately constant and equal to the Galactic value. On the other hand, observations that do not rely on CO have no such bias and so select clouds that lie on the power-law portion of the relationship. Therefore, CO observations find an $X_{\\rm CO}$ that does not vary significantly with metallicity, while other observations find a strong dependence on metallicity. Finally, we have shown that if we combine the requirement that $\\meanAV \\simgreat 3$ for a CO-bright cloud with the observational fact that the GMC size distribution is a steeply decreasing function of size, then we are lead quite naturally to the prediction that all GMCs should have near-constant column densities. Therefore, contrary to what is often assumed, the small scatter in observed GMC column densities does not necessarily imply that the GMCs are in virial equilibrium." }, "1003/1003.4615_arXiv.txt": { "abstract": "{We search for persistent extragalactic sources of \\gr s with energies above 100 GeV with the \\textit{Fermi} telescope.} {We construct a systematic survey of the extragalactic \\gr\\ sky at energies above 100 GeV. Such a survey has not been done before by the ground-based Cherenkov \\gr\\ telescopes, which have, contrary to \\textit{Fermi}, a narrow field of view.} {We study a map of arrival directions of the highest energy photons detected by \\textit{Fermi} at Galactic latitudes $|b|>10^\\circ$ and search for significant point-source-like excesses above the diffuse Galactic and extragalactic \\gr\\ backgrounds. We identify eight significant point-source-like excesses in this map. } { Seven of the eight sources are known TeV blazars. The previously unknown source is identified with the galaxy IC 310, which is situated in Perseus cluster of galaxies. The source is detected with a significance $6\\sigma$ above 30 GeV. We identify two possible scenarii for \\gr\\ emission from this source. One possibility is that emission originates from the base of relativistic outflow from the active nucleus, as in the BL Lacs and FR I type radio galaxies. Otherwise \\gr\\ photons could be produced at the bow shock that is formed as a result of the fast motion of the galaxy through the intracluster medium. The two models could be distinguished through studying of the \\gr\\ signal variability. } {} ", "introduction": " ", "conclusions": "" }, "1003/1003.5802_arXiv.txt": { "abstract": "{Do extrasolar planets affect the activity of their host stars? Indications for chromospheric activity enhancement have been found for a handful of targets, but in the X-ray regime, conclusive observational evidence is still missing.}{We want to establish a sound observational basis to confirm or reject major effects of Star-Planet Interactions (SPI) in stellar X-ray emissions.}{We therefore conduct a statistical analysis of stellar X-ray activity of all known planet-bearing stars within 30~pc distance for dependencies on planetary parameters such as mass and semimajor axis.}{In our sample, there are no significant correlations of X-ray luminosity or the activity indicator $L_X/L_{bol}$ with planetary parameters which cannot be explained by selection effects.}{Coronal SPI seems to be a phenomenon which might only manifest itself as a strong effect for a few individual targets, but not to have a major effect on planet-bearing stars in general.} ", "introduction": "The detection of extrasolar planets is one of the outstanding achievements in astronomy during the last 20 years. The first detected exoplanet revealed properties which were very surprising at that time: {\\em 51~Peg} \\citep{mayorqueloz1995} hosts a Jupiter-like planet at a distance of only 0.05~AU, thus the planet orbits its host star in less than five days. Since then, more than 400~other exoplanets have been found at the time of writing (see for example the Exoplanet Database at www.exoplanet.eu), both in very close orbits and in such more familiar from our own Solar system. With the existence of extrasolar planets established, the question arises what the environmental properties of such planets may be and if they might even allow the existence of life. The physical properties of planets, especially in close orbits, are crucially determined by the irradiation from their host stars. The evaporation rate of a planetary atmosphere depends on its exospheric temperature $T_{\\infty}$ , i.e., the regions where particles can escape freely \\citep{Lammer2003}. Thus, the host star's EUV and X-ray radiation is the key property determining a planet\u2019s exospheric temperature. Evaporation of the planetary atmosphere has been observed for the transiting planet {\\em HD~209458b} \\citep{Vidal-Madjar2003}: the planet loses hydrogen which is observable in absorption spectra during the transit. At very close distances, one might expect also planets to influence their host stars, in analogy to binary stars which show a higher activity level compared to single stars. Two different processes for Star-Planet-Interaction (SPI) have been put forward \\citep{CuntzSaar2000}. Planets can induce tidal bulges on the star with an interaction strength depending on the planetary semimajor axis ($\\propto a_{pl}^{-3}$), which might lead to enhanced coronal activity via increased turbulence in the photosphere. Planetary magnetic fields can also interact with the stellar magnetic field ($\\propto a_{pl}^{-2}$) and might also induce enhanced activity via Jupiter-Io-like interaction, i.e. flux tubes which connect star and planet and heat up their footpoints on the stellar surface, or magnetic reconnection. Some observational campaigns have been conducted to investigate the existence of possible SPI: \\cite{Shkolnik2005} monitored the chromospheric activity of 13 stars via \\ion{Ca}{ii} H and K line fluxes and found indications for cyclic activity enhancements in phase with the planetary orbit for two of these stars. The activity enhancements appeared once per planetary orbit, suggesting magnetic instead of tidal interaction. However, measurements obtained three years later \\citep{Shkolnik2008} showed that the activity enhancements had switched to a cycle in phase with the stellar rotation period instead. The {\\em coronal} activity of planet-bearing stars has been investigated in a first systematic study by \\cite{kashyapdrakesaar2008}. The authors claim an over-activity of planet-bearing stars of a factor of four compared to stars without planets, but their study had to include upper limits for a large number of stars since less than one third of the stars in their original sample were detected in X-rays at that time. A dedicated campaign to search for magnetic SPI in the case of {\\em HD~179949}, one of the stars which \\cite{Shkolnik2005} found to have cyclic activity changes in the chromosphere, was conducted by \\cite{SaarSPI2008}. These authors found spectral and temporal variability phased with the planetary orbit, but some of that might also be induced by intrinsic stellar activity variations, since the stellar rotation period is poorly known ($P_*=7-10$~d). Up to now, the observational basis of stellar coronal activity enhancements due to close-in planets is not sound enough to establish or reject the possibility of coronal SPI. In order to adress this issue we conducted an X-ray study of all planet-hosting stars within a distance of 30~pc with XMM-Newton which have not been studied with ROSAT before. In this fashion a volume-limited complete stellar sample can be constructed. ", "conclusions": "We analyzed a sample of all known planet-bearing stars in the solar neighborhood for X-ray activity and possible manifestations of coronal Star-Planet-Interactions (SPI) in dependence of the planetary parameters mass and semimajor axis. Our main results are summarized as follows: \\begin{enumerate} \\item In our sample of $72$ stars, there are no significant correlations of the activity indicator $L_X/L_{bol}$ with planetary mass or semimajor axis. \\item However, we do find a correlation of the X-ray luminosity with the product of planetary mass and reciprocal semimajor axis. Massive close-in planets are often found around X-ray brighter stars. \\item This dependence can be ascribed to selection effects: The radial velocity method for planet detections favors small and far-out planets to be detected around low-activity, X-ray dim stars. Additionally, if SPI induced an excess in $L_X$ without changing the $L_X/L_{bol}$ ratio, SPI would need to cause extremely high energy input in $L_{bol}$, leading to unrealistically short decay times for the planetary orbit. \\item There is no {\\em additional} effect detectable in $L_X$ which could be attributed to coronal manifestations of SPI. \\item Coronal SPI might still be observable for some individual promising targets, preferably in coordinated observations of the targets' coronae and chromospheres. \\end{enumerate}" }, "1003/1003.0420_arXiv.txt": { "abstract": "Low-amplitude Doppler-shift oscillations have been observed in coronal emission lines in a number of active regions with the EUV Imaging Spectrometer (EIS) on the \\textit{Hinode} satellite. Both standing and propagating waves have been detected and many periods have been observed, but a clear picture of all the wave modes that might be associated with active regions has not yet emerged. In this study, we examine additional observations obtained with EIS in plage near an active region on 2007 August 22--23. We find Doppler-shift oscillations with amplitudes between 1 and 2 km~s$^{-1}$ in emission lines ranging from \\ion{Fe}{11} 188.23~\\AA, which is formed at $\\log T = 6.07$ to \\ion{Fe}{15} 284.16~\\AA, which is formed at $\\log T = 6.32$. Typical periods are near 10 minutes. We also observe intensity and density oscillations for some of the detected Doppler-shift oscillations. In the better-observed cases, the oscillations are consistent with upwardly propagating slow magnetoacoustic waves. Simultaneous observations of the \\ion{Ca}{2} H line with the \\textit{Hinode} Solar Optical Telescope Broadband Filter Imager show some evidence for 10-minute oscillations as well. ", "introduction": "\\label{intro} Because the outer regions of the solar atmosphere are threaded by a magnetic field, they can support a wide range of oscillatory phenomena. Theoretical aspects of these waves and oscillations have been the subject of extensive investigations \\citep[e.g.,][]{Roberts1983,Roberts1984}. These initial studies were mainly driven by radio observations of short period oscillations \\citep[e.g.,][]{Rosenberg1970}. More recent detections of coronal oscillatory phenomena have resulted in considerable additional theoretical work. Recent reviews include \\citet{Roberts2000} and \\citet{Roberts2003}. Observational detections of coronal oscillatory phenomena include detections of spatial oscillations of coronal structures, which have been interpreted as fast kink mode MHD disturbances \\citep[e.g.,][]{Aschwanden1999}; intensity oscillations, which have been interpreted as propagating slow magnetoacoustic waves \\citep[e.g.,][]{DeForest1998,Berghmans1999}; and Doppler shift oscillations, which have been interpreted as slow mode MHD waves \\citep[e.g.,][]{Wang2002}. The interaction between the theoretical work and the growing body of oscillation observations provides fertile ground for testing our understanding of the structure and dynamics of the corona. The EUV Imaging Spectrometer (EIS) on the \\textit{Hinode} satellite is an excellent tool for studying oscillatory phenomena in the corona. \\citet{Culhane2007} provides a detailed description of EIS, and the overall \\textit{Hinode} mission is described in \\citet{Kosugi2007}. Briefly, EIS produces stigmatic spectra in two 40~\\AA\\ wavelength bands centered at 195 and 270~\\AA. Two slits (1\\arcsec\\ and 2\\arcsec) provide line profiles, and two slots (40\\arcsec\\ and 266\\arcsec) provide monochromatic images. Moving a fine mirror mechanism allows EIS to build up spectroheliograms in selected emission lines by rastering a region of interest. With typical exposure times of 30 to 90~s, however, it can take considerable time to construct an image. Higher time cadences can be achieved by keeping the EIS slit or slot fixed on the Sun and making repeated exposures. This sit-and-stare mode is ideal for searching for oscillatory phenomena. EIS Doppler shift data have already been used for a number of investigations of oscillatory phenomena. \\citet{VanDoorsselaere2008} have detected kink mode MHD oscillations with a period near 5 minutes. \\citet{Mariska2008} observed damped slow magnetoacoustic standing waves with periods of about 35 minutes. \\citet{Wang2009} have detected slow mode magnetoacoustic waves with 5 minute periods propagating upward from the chromosphere to the corona in an active region. \\citet{Wang2009a} have also observed propagating slow magnetoacoustic waves with periods of 12 to 25 minutes in a large fan structure associated with an active region. Analysis of oscillatory data with EIS is still just beginning. The amplitudes observed have all been very small---typically 1 to 2 km~s$^{-1}$. Thus a clear picture of the nature of the low-amplitude coronal oscillations has yet to emerge. In this paper, we add to that picture by analyzing portions of an EIS sit-and-stare active region observation that shows evidence for Doppler-shift oscillations in a number of EUV emission lines. We also use data from the \\textit{Hinode} Solar Optical Telescope (SOT) to relate the phenomena observed in the corona with EIS to chromospheric features and the magnetic field. ", "conclusions": "As we pointed out in \\S\\ref{intro}, a number of investigations have detected Doppler shift oscillations with EIS. Based on the phase differences between the Doppler shift and the intensity, we believe that the signals we have detected are upwardly propagating magnetoacoustic waves. The periods we detect are between 7 and 14 minutes. For a set of observations that begins at a particular time, there is considerable scatter in the measured periods and amplitudes. This is probably due to the relatively weak signal we are analyzing. But it may also be an indication that a simple sine wave fit is not a good representation of the data. It is likely that each line-of-sight passes through a complex, time-dependent dynamical system. While a single flux tube may respond to an oscillatory signal by exhibiting a damped sine wave in the Doppler shift, a more complex line-of-sight may display a superposition of waves. Coalignment of the EIS data with both SOT and MDI magnetograms shows that the portion of the EIS slit analyzed in this study corresponds to a unipolar flux concentration. SOT \\ion{Ca}{2} images show that the intensity of this feature exhibits 5-minute oscillations typical of chromospheric plasma, but also exhibits some evidence for longer period oscillations in the time range detected by EIS. Moreover, the \\ion{Ca}{2} intensity data show that the oscillations observed in EIS are related to significant enhancements in the \\ion{Ca}{2} intensities, suggesting that a small chromospheric heating event triggered the observed EIS response. \\citet{Wang2009} also detected propagating slow magnetoacoustic waves in an active region observed with EIS, which they associated with the footpoint of a coronal loop. While the oscillation periods they measured---5 minutes---were smaller than those detected here, many of the overall characteristics we see are the same. In each case, the oscillation only persists for a few cycles and the phase relationship indicates an upwardly propagating wave. In contrast with their results, however, we do not see a consistent trend for the oscillation amplitude to decrease with increasing temperature of line formation. Examination of both the Doppler shift data in Figure~\\ref{fig:doppler_average} and the results in Table~\\ref{table:oscillation_d_fits} shows that in one case the amplitude has a tendency to decrease with increasing temperature of line formation (oscillation beginning at 113.9 minutes) and in another case the amplitude clearly increases with increasing temperature of line formation (oscillation beginning at 190 minutes). Thus it does not appear that the results reported by \\citet{Wang2009} are always the case. \\citet{O'Shea2002} noted that for oscillations observed above a sunspot the amplitude decreased with increasing temperature until the temperature of formation of \\ion{Mg}{10}, which is formed at roughly 1~MK. They then saw an increase in amplitude in emission from \\ion{Fe}{16}. All the EIS lines we have included in this study have temperatures of formation greater than 1~MK. Combined EIS and SUMER polar coronal hole observations have also shown evidence propagating slow magnetoacoustic waves \\citep{Banerjee2009}. These waves have periods in the 10--30 minute range. These waves appear to be more like those we observe in that they are have periods longer than those studied by \\citet{Wang2009}. \\citet{Wang2009} suggested that the waves they observed were the result of leakage of photospheric p-mode oscillations upward into the corona. The longer periods we and \\citet{Banerjee2009} observe are probably not related to p-modes. Instead, we speculate that the periods of the waves are related to the impulsive heating which may be producing them. If an instability is near the point where rather than generating a catastrophic release of energy it wanders back and forth between generating heating and turning back off, waves would be created. The heating source could be at a single location, or, for example, locations near each other where instability in one place causes a second nearby location to go unstable and begin heating plasma. In this view, the periods provide some insight into the timescale for the heating to rise and fall and thus may be able to place limits on possible heating mechanisms. The behavior of slow magnetoacoustic oscillations as a function of temperature has been the subject of considerable theoretical work. It is generally believed that the damping of the waves is due to thermal conduction \\citep[e.g.,][]{DeMoortel2003,Klimchuk2004}. Because thermal conduction scales as a high power of the temperature, conductive damping should be stronger for oscillations detected in higher temperature emission lines \\citep[e.g.,][]{Porter1994,Ofman2002}. Earlier EIS observations of the damping of standing slow magnetoacoustic waves, however, show that this is not always the case \\citep{Mariska2008}. Thus, the temperature behavior of both the amplitude of the oscillations and the damping, differ from some earlier results. We believe that additional observations will be required to understand fully the physical picture of what is occurring in the low corona when oscillations are observed. Given the complex set of structures that may be in the line of sight to any given solar location under the EIS slit, we are not entirely surprised that different data sets should yield different results, which in some cases differ from models. For example, none of the current models for oscillations in the outer layers of the solar atmosphere take into account the possibility that what appear to be single structures in the data might actually be bundles of threads with differing physical conditions. Our observations along with others \\citep[e.g.,][]{Wang2009,Wang2009a,Banerjee2009} show that low-amplitude upwardly propagating slow magnetoacoustic waves are not uncommon in the low corona. The periods observed to date range from 5 minutes to 30 minutes. In all cases, however, the wave amplitudes are too small to contribute significantly to coronal heating. But understanding how the waves are generated and behave as a function of line formation temperature and the structure of the magnetic field should lead to a more complete understanding of the structure of the low corona and its connection with the underlying portions of the atmosphere. Instruments like those on \\textit{Hinode} that can simultaneously observe both the chromosphere and the corona, should provide valuable additional insight into these waves as the new solar cycle rises and more active regions become available for study." }, "1003/1003.2755_arXiv.txt": { "abstract": "We present a study on the clustering of a stellar mass selected sample of 18,482 galaxies with stellar masses $M_{*}>10^{10}$\\solmg at redshifts $0.410^{13}$\\solmm we find that this ratio is $<0.02$, much lower than the universal baryonic mass fraction. We show that the remaining baryonic mass is included partially in stars within satellite galaxies in these haloes, and as diffuse hot and warm gas. We also find that, at a fixed stellar mass, the stellar-to-total-mass ratio increases at lower redshifts. This suggests that galaxies at a fixed stellar mass form later in lower mass dark matter haloes, and earlier in massive haloes. We interpret this as a ``halo downsizing'' effect, however some of this evolution could be attributed to halo assembly bias. ", "introduction": "\\label{sec:intro} Astronomers in the last decade have made major progress in understanding the properties and evolution of galaxies in the distant Universe. Galaxies up to redshifts $z\\sim6$, and perhaps even higher, have been discovered in large numbers allowing statistically significant population characteristics to be derived \\citep[e.g.][]{2007ApJ...670..928B,2009MNRAS.395.2196M}. We in fact now have a good understanding of basic galaxy quantities, such as the luminosity function \\citep[e.g.][]{2003ApJ...595..698D,2005A&A...439..863I,2007MNRAS.380..585C,ciras10}, as well as how scaling relations, such as the Tully-Fisher relation evolve up to, at least, $z\\sim 1.2$ \\citep[e.g.][]{2004A&A...420...97B,2005ApJ...628..160C,2006MNRAS.366..308B,2006ApJ...653.1049W,2007ApJ...660L..35K,2007ApJ...668..846B,2007MNRAS.377..806C,2008A&A...484..173P,2009A&A...496..389F}. We have also begun to trace the stellar mass evolution of galaxies, as well as the star formation rate, determining when stars and stellar mass was put into place in the modern galaxy population \\citep[e.g.][]{2006A&A...453..869B,2006ApJ...651..120B,2006A&A...459..745F,2009ApJ...701.1765M}. Stellar masses are now becoming a standard measure for galaxies, and are being used to trace the evolution of the galaxy population in terms of star formation rates and morphologies \\citep[e.g.][]{2005ApJ...625..621B,2005ApJ...620..564C,2008MNRAS.383.1366C,2007ApJ...660L..47N,2007ApJ...663..834B,2008ApJ...686...72C}. However, stellar masses only trace one aspect of the masses of galaxies, and ideally and ultimately, we want to be able to measure galaxy total masses, that include contributions from stars, gas, and dark matter. Galaxies are believed to be hosted by massive dark matter haloes that make up more than 85\\% of their total masses, and thus clearly tracing the co-evolution of galaxies and their haloes is a major and important goal. Measuring dynamical or total masses for galaxies is, however, very difficult, as it requires observations that are very challenging to obtain, or requires unusual and rare circumstances such as gravitational lensing. For example, dynamical masses can be measured through rotation curves, or internal velocity dispersions, but this becomes difficult at higher redshifts, with reliable internal velocities existing for only a handful of galaxies at $z\\sim1.0$ and above \\citep[e.g.][]{2005ApJ...628..160C,2006ApJ...645.1062F,2006MNRAS.368.1631S,2007ApJ...668..846B,2008A&A...488...99V}. Furthermore, it is difficult to know whether these kinematic measures are tracing the total potential, or just the inner parts of galaxies. Likewise, using gravitational lensing to measure galaxy total masses is difficult, as it requires lensed background galaxies, and such instances are very rare. It is also not yet clear if these galaxies are representative, as they may have special mass profiles that are conducive to lensing. The weak galaxy-galaxy lensing technique provides another way to estimate the total masses of galaxies. No unusual circumstances are required, but individual measurements are extremely difficult to achieve. Stacking techniques give reliable results, but involve combining galaxies in haloes of different masses, which complicates the interpretation \\cite[e.g.][]{2006MNRAS.368..715M}. Stacked satellite kinematics can also be used to probe the total masses of galaxies, but face the same problems as the weak galaxy-galaxy lensing method \\citep{2009MNRAS.392..917M}. Therefore, obtaining total mass estimates for galaxies is difficult, and very few measures, or even estimates, have been produced for galaxies outside the local Universe. One very powerful method for measuring the total masses of galaxies is to measure their clustering. Clustering measurements are independent of photometric properties, and as such they can be used to highlight fundamental properties of galaxy populations without assumptions concerning stellar populations or mass profiles. Previously, it has been shown that galaxies with higher stellar masses cluster more than systems with lower stellar mass, with a very strong clustering above the characteristic stellar mass $M^{*}$ of the Schechter mass function \\citep{2006MNRAS.368...21L}. In the halo model of galaxy formation, the large-scale distribution of galaxies is determined by the distribution of dark matter haloes \\citep{MW02}. Therefore, halo clustering is a strong function of halo mass, with more massive haloes more strongly clustered, providing a means to study the relationship between galaxy properties, and dark matter halo masses. In fact, basic calculations allow one to convert a clustering strength for a sample of galaxies into a corresponding halo mass in which galaxies reside \\citep[e.g.][]{2004ApJ...611..685O,2008MNRAS.383.1131M,2008ApJ...679..269Y}. Clustering measurements have been performed on galaxy samples in both the nearby and distant Universe. In the local Universe, studies have established that clustering depends on the type of galaxy under consideration; for example, early-type, red, galaxies are more clustered than late-type, blue galaxies \\citep[e.g.][]{1997ApJ...489...37G,2002MNRAS.332..827N,2007MNRAS.379.1562C,2009MNRAS.392..682C}, and luminous galaxies are more clustered than faint galaxies \\citep[e.g.][]{2001MNRAS.328...64N, 2005ApJ...630....1Z,2006MNRAS.369...68S}. Moreover, distant massive galaxies selected by their extremely red colours have been shown to strongly cluster, and therefore are thought to inhabit massive dark matter haloes \\citep[e.g.][]{UDS-EDR-DRG,2007ApJ...654..138Q,2008MNRAS.383.1131M,hartley08}. In this paper, we present the first general study of the clustering properties for a stellar mass selected sample of galaxies up to $z\\sim2$. We carry this out by measuring the correlation length and amplitude for galaxies selected with stellar masses $M_*>10^{10}$\\solmg within the Palomar Observatory Wide-field Infrared Survey (POWIR) \\citep{2006ApJ...651..120B,2008MNRAS.383.1366C}, which covers the fields where spectroscopy and other multi-wavelength data are available through the Deep Extragalactic Evolutionary Probe 2 (DEEP2) survey and All-wavelength Extended Groth strip International Survey (AEGIS) \\citep{2003SPIE.4834..161D,2007ApJ...660L...1D}. In total we examine the clustering strength for 18,482 galaxies selected by stellar mass within 0.7 deg$^{2}$. We derive an estimate of the total masses for these galaxies, and study in detail their stellar-mass-to-total-mass ratio. This paper is presented as follows: our data-sets, catalogues and our photometric redshifts and stellar mass estimations are described in Section~\\ref{sec:data}; in Section~\\ref{sec:cf} we describe the methods used in measuring the clustering properties of our samples; Section~\\ref{sec:dmm} is dedicated to the methods we use to derive the masses of dark matter haloes for our samples; in Section~\\ref{sec:ratio} we compare our results with the literature and models; and finally we summerise our conclusions in Section~\\ref{sec:summ}. Throughout the paper, we assume a $\\Lambda$CDM cosmology with $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$, $h=H_0/70$~km~s$^{-1}$~Mpc$^{-1}$. To ease comparisons with previous work, we use a concordance model with fiducial values of $n_s=1.0$ and $\\sigma_8=0.9$. To determine stellar masses throughout this paper, we use the Initial Mass Function (IMF) from \\cite{chabrier03} and assume a Hubble constant of $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$. ", "conclusions": "\\label{sec:summ} We present in this paper an analysis of the clustering properties of a stellar mass selected sample of galaxies over 0.7 deg$^{2}$ taken from the Palomar Observatory Wide-Field Infrared Survey (POWIR). By utilising information from three separate fields, we measure the two-point correlation function, and based on this we derive the spatial correlation length and bias for our sample of galaxies within various stellar mass bins at $M_{*}>10^{10}$\\solmg. Our major results are: \\vspace{0.2cm} \\noindent I. We find that the correlation length ($r_{0}$) varies with both redshift and stellar mass. We find that the largest correlation lengths (most clustered systems) are found for galaxies at the highest stellar masses with $10^{11}$\\solmg$10^{10.5}$\\solmg are a factor of 30 to 170 times the stellar masses of each galaxy. This is in very good agreement with other methods for measuring halo masses, including lensing and kinematics. \\vspace{0.2cm} \\noindent III. We find a remarkable relation between dark matter halo mass ($M_{\\rm DM}$), and the ratio of stellar-to-dark-matter-mass ($M_{*}/M_{\\rm DM}$) for the central galaxy in the halo, and find that this relation does not vary much with redshift. This correlation is such that the more massive a dark matter halo mass $M_{\\rm DM}$ is, the lower the ratio of stellar-to-halo-mass ($M_{*}/M_{\\rm DM}$). We further find that this correlation exists over at least two orders of magnitude from halo masses $M_{\\rm DM}=10^{12.0}$ to $10^{14.0}$\\solmm. This is true for all methods of determining halo masses, from lensing, kinematics, and clustering. \\vspace{0.2cm} \\noindent IV. Within our massive galaxy sample at $z\\sim1-2$ we find that the ratio of stellar-to-halo-mass increases from high to low redshifts. This correlation implies that at a given stellar mass selection, the average underlying mass of the host halo decreases at lower redshifts. This may imply that the most stellar massive systems at the highest redshifts are hosted by very massive haloes, but at lower redshift they stop forming new stars, and therefore do not increase their stellar mass, while other systems in the lower mass dark matter haloes grow to reach similar stellar mass. This is an example of ``Halo downsizing''. The `Halo Assembly bias' implying that clustering depends not only on the mass of the dark matter halo, but also on the assembly history and environment of the galaxy, can also contribute partially to this observed increase in the stellar mass fraction with redshift. \\vspace{0.2cm} \\noindent V. We compare our results to the semi-analytic models built using the Millennium simulation. We find that there is roughly a good agreement between models and the data for the central galaxy. The predicted value of the total-stellar-to-total-halo mass ratio decreases only slightly with increasing halo mass, also in rough agreement with our observational results. We use the SDSS group catalogue from \\cite{2006MNRAS.366....2W} to examine whether the missing stellar mass can be accounted for by observed satellite galaxies. We find that, even though the overall agreement is good, the model marginally overpredicts the number of faint galaxies in comparison to the SDSS catalogue. Moreover the Millennium simulation fails to reproduce the ``halo downsizing'' effect we observe. We argue that this missing mass, in the more massive haloes, is most probably in the form of baryons in a warm/hot phase." }, "1003/1003.0566_arXiv.txt": { "abstract": "{} {We investigate the constraints imposed on the luminosity function (LF) of long duration Gamma Ray Bursts (LGRBs) by the flux distribution of bursts detected by the GBM at $\\sim1$~MeV, and the implications of the non detection of the vast majority, $\\sim$95\\%, of the LGRBs at higher energy, $\\sim1$~GeV, by the LAT detector. } {We find a LF that is consistent with those determined by BATSE and {\\it Swift}. The non detections by LAT set upper limits on the ratio $R$ of the prompt fluence at $\\sim1$~GeV to that at $\\sim1$~MeV. The upper limits are more stringent for brighter bursts, with $R<\\{0.1,0.3,1\\}$ for $\\{5,30,60\\}\\%$ of the bursts. This implies that for most bursts the prompt $\\sim1$~GeV emission may be comparable to the $\\sim1$~MeV emission, but can not dominate it. The value of $R$ is not universal, with a spread of (at least) an order of magnitude around $R\\sim10^{-1}$. For several bright bursts with reliable determination of the photon spectral index at $\\sim1$~MeV, the LAT non detection implies an upper limit to the $\\sim100$~MeV flux which is $<0.1$ of the flux obtained by extrapolating the $\\sim1$~MeV flux to high energy.} { For the widely accepted models, in which the $\\sim1$~MeV power-law photon spectrum reflects the power-law energy distribution of fast cooling electrons, this suggests that either the electron energy distribution does not follow a power-law over a wide energy range, or that the high energy photons are absorbed. Requiring an order unity pair production optical depth at 100~MeV sets an upper limit for the Lorentz factor, $\\Gamma\\lsim10^{2.5}$.} {} ", "introduction": "Gamma ray bursts (GRBs) are the most powerful explosions in the universe, with apparent (isotropic equivalent) energy output sometimes exceeding $10^{54}$ ergs. While it is widely accepted that GRBs are produced by the dissipation of energy in highly relativistic winds driven by compact objects (see, e.g., M\\'esz\\'aros 2006, Piran 2004 and Waxman 2003 for reviews) the physics of wind generation and radiation production is not yet understood. It is not known, for example, whether the wind luminosity is carried, as commonly assumed, by kinetic energy or by Poynting flux (e.g. Drenkhahn \\& Spruit 2002, Lyutikov et al. 2003), whether the radiating particles are accelerated by the dissipation of magnetic flux or by internal shocks dissipating kinetic energy, and whether the emission is dominated by synchrotron or inverse-Compton radiation of accelerated electrons, as commonly assumed, or by hadronic energy loss of accelerated protons (see Dermer \\& Fryer 2008 and references therein). GRBs were detected by instruments sensitive mainly in the 100 to 1000 keV range like BATSE (Paciesas et al. 1999) and the GRBM on BeppoSAX (Guidorzi et al. 2004). Measurements at high energy were limited by the lower sensitivity and/or the smaller field of view of higher energy instruments (e.g. CGRO/EGRET- Dingus 1995, Hurley et al. 1994, Gonzalez et al. 1994; AGILE/GRID-Marisaldi et al. 2009, Giuliani et al. 2008, Giuliani et al. 2010). The improved high energy, $\\gsim1$~GeV, sensitivity and field of view of the instruments on board the Fermi satellite (Atwood et al. 2009, Band et al. 2009) are expected to improve the quantity and quality of high energy GRB data, and may therefore provide qualitatively new constraints on models. The main goal of this paper is to determine the implications of the non-detection of the vast majority of long duration ($T>2$~s) GRBs by Fermi's LAT detector. We first show in \\S~\\ref{sec:LF} that the LF of LGRBs detected by Fermi's GBM at $\\sim1$~MeV is consistent with those inferred from BATSE and {\\it Swift} observations, and that these instruments sample the LF in a similar manner. We then derive in \\S~\\ref{sec:non_detect} the constraints on the $>100$~MeV emission implied by LAT non-detections. In \\S~\\ref{sec:EGRET} we compare our results with those inferred from the analysis of EGRET data (Gonzalez Sanchez 2005). Our results are summarized and their implications are discussed in \\S~\\ref{sec:discussion}. ", "conclusions": "\\label{sec:discussion} We have shown that the LGRB LF inferred from the sample of bursts detected by Fermi's GBM is consistent with those determined by BATSE and {\\it Swift}, see table~1, and that the GBM samples this LF in a manner similar to that of BATSE \\& {\\it Swift}, see figure~1. We have derived the upper limits implied by the LAT non-detections on the ratio $R$ of the 100~MeV to 1~MeV fluences, $\\nu F_\\nu=\\nu\\int dt f_\\nu(t)$ for $h\\nu=1,100$~MeV, during the prompt emission. The upper limits on $R$ obtained assuming $dN_\\gamma/dE\\propto E^{-2}$ at $E>100$~MeV (see fig.~2), also imply upper limits on the 0.1--1~GeV fluence, which is approximately given by $\\ln(10)R\\nu F_\\nu|_{1\\rm MeV}$ (the upper limit on the 1--10~GeV fluence is few times higher and depends on the assumed spectrum), and on the ratio of the 0.1--1~GeV fluence and the 0.1--1~MeV fluence, which is $\\approx R$. The upper limits on $R$ are more stringent for brighter bursts (see fig.~2), with $R<\\{0.1,0.3,1\\}$ for $\\{5,30,60\\}\\%$ of the bursts (see fig.~4). This implies that for most bursts the prompt $\\sim1$~GeV emission may be comparable to the $\\sim1$~MeV emission, but can not dominate it. For several bright bursts with reliable determination of the photon spectral index at $\\sim1$~MeV, the LAT non detection implies an upper limit to the $\\sim100$~MeV flux which is $<0.1$ of the flux obtained by extrapolating the $\\sim1$~MeV flux to high energy (see fig.~5). Examining figs.~2 and~4, we conclude that the ratio $R$ is not universal among GRBs. The detections and non-detection upper limits imply a spread in $R$ over at least an order of magnitude. The upper limits we obtain are similar to those inferred for the fluence at lower energy, 30--200~MeV, from EGRET's non-detections of BATSE bursts (see \\S~\\ref{sec:EGRET}). The upper limits on $R$ provide constraints on models for the prompt GRB emission. Models where the prompt $\\sim1$~MeV emission is produced by inverse-Compton scattering of optical synchrotron photons (e.g. Stern \\& Poutanen 2004, Panaitescu \\& Kumar 2007), typically predict $R\\ge1$. This is not supported by the data. Such models are not necessarily ruled out by the current data, as they might be modified to include a suppression of the $\\sim1$~GeV flux by pair production. Such modification may be required for all (widely discussed) models, in which the $\\sim1$~MeV power-law photon spectrum reflects the power-law energy distribution of fast cooling electrons. The suppression of the $\\sim100$~MeV flux, compared to that expected from an extrapolation of the $\\sim1$~MeV power-law spectrum, suggests that either the electron energy distribution does not follow a power-law over a wide energy range, or that the high energy photons are absorbed, probably by pair production. Requiring an optical depth of $\\sim1$ at $100$~MeV sets an upper limit to the expansion Lorentz factor $\\Gamma\\lsim10^{2.5}[(L/10^{52}{\\rm erg/s})/(t_v/10\\,{\\rm ms})]^{1/6}$ (e.g. eq. 7 of Waxman 2003). Significant compactness of the emission region has been suggested by several authors (e.g. Guetta et al. 2001, Pe'er \\& Waxman 2004). The spectrum is modified in this case, compared to the optically thin case, with 100~MeV to 1~MeV flux ratios in the range $0.01\\lsim R\\lsim0.1$ obtained for typical parameters (e.g. figs. 8 \\& 11 of Pe'er \\& Waxman 2004)." }, "1003/1003.1756_arXiv.txt": { "abstract": "We use one-dimensional two-zone time-dependent accretion disk models to study the long-term evolution of protostellar disks subject to mass addition from the collapse of a rotating cloud core. Our model consists of a constant surface density magnetically coupled active layer, with transport and dissipation in inactive regions only via gravitational instability. We start our simulations after a central protostar has formed, containing $\\sim$ 10\\% of the mass of the protostellar cloud. Subsequent evolution depends on the angular momentum of the accreting envelope. We find that disk accretion matches the infall rate early in the disk evolution because much of the inner disk is hot enough to couple to the magnetic field. Later infall reaches the disk beyond $\\sim$10 AU, and the disk undergoes outbursts of accretion in FU Ori-like events as described in Zhu et al. 2009c. If the initial cloud core is moderately rotating most of the central star's mass is built up by these outburst events. Our results suggest that the protostellar ``luminosity problem'' is eased by accretion during these FU Ori-like outbursts. After infall stops the disk enters the T Tauri phase. An outer, viscously evolving disk has structure that is in reasonable agreement with recent submillimeter studies and its surface density evolves from $\\Sigma \\propto R^{-1}$ to $R^{-1.5}$. An inner, massive belt of material-- the ``dead zone'' -- would not have been observed yet but should be seen in future high angular resolution observations by EVLA and ALMA. This high surface density belt is a generic consequence of low angular momentum transport efficiency at radii where the disk is magnetically decoupled, and would strongly affect planet formation and migration. ", "introduction": "In the picture of low-mass star formation from large, cold protostellar clouds, any small initial rotation of the cloud will lead to the formation of an accretion disk to conserve angular momentum; thus much, if not most, of the mass of low-mass stars is probably accreted from disks. However, if disks transport the infalling mass steadily to the central star, accretion luminosities are considerably higher than typically observed protostellar luminosities (Kenyon \\etal 1990, 1994; Enoch \\etal 2008; Evans \\etal 2009). This ``luminosity problem'' can be alleviated if disks spend most of their evolution accreting slowly, undergoing episodic accretion outbursts where most of the mass is added to the central stars. One implication of non-steady accretion is that mass may pile up in specific regions of disks, which could have significant consequences for understanding planet formation. Over the past decade, disk angular momentum transport by the magnetorotational instability (MRI) (Balbus \\& Hawley 1998) and by the gravitational instability (GI) (Durisen \\etal 2007) has become much better understood. It seems likely that both types of instability need to be considered to understand protostellar accretion. At early evolutionary stages the disk is likely to be quite massive with respect to the central protostar, suggesting that GI may be important; in addition, these disks are so cold that thermal ionization cannot sustain the MRI, which requires coupling of the magnetic field to the mostly-neutral gas through collisions (e.g., Reyes-Ruiz $\\&$ Stepinski 1995). On the other hand, external ionizing sources (cosmic and/or X rays) can provide the necessary ionization for MRI activitation up to a limiting surface density, ionizing the outer surfaces of the disk and leading to accretion in an ``active layer'', leaving a ``dead zone'' in the midplane (Gammie 1996). In addition, at high accretion rates, the inner disk can become thermally-ionized, as in FU Ori outbursts (Zhu \\etal 2007, 2009b). Mismatches in the transport rate between the GI and MRI can lead to outbursts of accretion (Armitage et al. 2001, Zhu et al. 2009 bc). In Zhu et al. 2010 (Paper I), we constructed one-dimensional, time-dependent disk models to study the protostellar unsteady accretion. These models assume both thermally-activated and layered MRI-driven accretion along with a local treatment of the GI. We further adopted steady mass addition at a specified outer disk radius to drive the system. We found that accretion in protostellar disks is unsteady over a wide range of parameters because of a mismatch between GI transport in the outer disk and MRI transport in the inner disk. Our results for outburst behavior in these one-dimensional model were sufficiently comparable to our previous two-dimensional simulations of outbursts (Zhu \\etal 2009c), confirming the utility of the faster 1D simulations to explore wider ranges of parameter space. In this paper we extend our results from Paper I to study long-term disk evolution from the protostellar phase to the T Tauri phase, using a more self-consistent treatment for infall from a collapsing, rotating cloud core. This allows us to study the effect of the initial core rotation and different disk accretion configurations (layered accretion, GI-only accretion, and constant-$\\alpha$ accretion) on the disk formation and evolution. Rice, Mayo, \\& Armitage (2009) have also investigated somewhat similar one-dimensional models but do not assume layered accretion. Our models treat the same phase of evolution as Vorobyov \\& Basu (2005, 2006, 2007) investigated with two-dimensional models. They also argue that protostellar accretion is generally non-steady, but for a different reason, specifically the accretion of clumps created by gravitational instabilities. Kratter \\etal (2009) have also investigated the protostellar accretion phase with three-dimensional numerical simulations. Although our treatment of the core collapse and the GI is much more schematic than used in the above investigations, we are able to treat radiative cooling more realistically, and can follow disk evolution to much smaller radii, where the (thermally-activated) MRI can become important. Furthermore, our two-zone disk model allows us to study the effects of different disk accretion configurations efficiently. The variety of disk structures predicted by our simulations with different initial cloud core rotations and disk accretion mechanisms can be tested by future EVLA and ALMA observations, which will help us to better understand disk accretion and planet formation. We describe our one-dimensional, two-zone model with self-consistent infall in \\S 2. In \\S 3 we explore the behavior of the disk models with infall. In \\S 4 we discuss the implications of our results, and summarize our conclusions in \\S 5. We defer all derivations to the Appendix. ", "conclusions": "In this paper, we have constructed a one-dimensional two-zone accretion disk model to study disk formation and long-term evolution under the collapse of a BE rotating core. The model evolution can be divided into three stages. At the early stage, when the mass falls to the inner disk within AU scale, the MRI can be sustained in the inner disk and efficiently and steadily transfers the infalling mass to the central star. Later, when the mass falls beyond AU scale, the disk goes to the outburst stage due to the accretion rates' mismatch by the MRI and GI as described in Paper I. After the infall completes, the disk enters the T Tauri phase and evolves on its own. Cores with higher initial rotation end up with a more massive disk and more disk episodic accretion events (outbursts). As long as the initial cloud core does not rotate extremely slowly to form a tiny disk (R$_{cmax}$$\\sim$1 AU), more than half of the star mass is built up by outbursts, which eases the ``luminosity'' problem. Disks exhibit a variety of behavior during the T Tauri phase. For a disk with accretion sustained only by gravitational instability, the disk evolves towards a Q=1 disk and the disk truncates at a radius slightly larger than the maximum centrifugal radius of the infall. If the disk has an active layer at the surface, however, the active layer can extend to a much larger radius and a sharp density drop develops at a characteristic radius R$_{D}$ that separates the marginally gravitationally stable dead zone and the MRI active but gravitationally stable outer disk. The density jump at R$_{D}$ may be observable by the EVLA and ALMA. The formation of a dense belt of material is associated with the failure of magnetically driven transport due to low ionization at intermediate radius in the disk; the only ways to avoid this are (1) if there is a separate, equally efficient hydrodynamic transport mechanism, or (2) if for some reason the MRI fails in the outer disk as well, perhaps due to dynamo failure." }, "1003/1003.3109_arXiv.txt": { "abstract": "{This is the third paper in a series in which we attempt to put constraints on the flattening of dark halos in disk galaxies. We observed for this purpose the \\HI\\ in edge-on galaxies, where it is in principle possible to measure the force field in the halo vertically and radially from gas layer flaring and rotation curve decomposition respectively. For this purpose we need to analyse the observed XV diagrams in such a way as to accurately measure all three functions that describe the planar kinematics and distribution of a galaxy: the radial \\HI\\ surface density, the rotation curve and the \\HI\\ velocity dispersion. In this paper, we first present the results of the modelling of our \\HI\\ observations of 8 \\HI\\ rich, late-type, edge-on galaxies. We find that in all of these we find differential rotation. Most systems display \\HI\\ velocity dispersions of 6.5 to 7.5 km s$^{-1}$ and all except one show radial structure in this property. There is an increase in the mean \\HI\\ velocity dispersion with maximum rotation velocity, at least up to 120 km s$^{-1}$. Next we analyse the \\HI\\ observations to derive the radial variation of the thickness of the \\HI\\ layer. The combination of these gas flaring measurements with the \\HI\\ kinematics measurements allow us to measure the total vertical force of each galaxy assuming hydrostatic equilibrium. We find that with the exception of the asymmetric IC5052, all of the galaxies in our sample are good candidates for 3D mass modelling to measure the dark halo shape. The flaring profiles are symmetric with respect to the galactic centres and have a common shape, increasing linearly inside the stellar disks and exponential outside where the gravitational potential is dominated by the dark halo. In the best example, UGC7321, we find in the inner regions small deviations from the midplane and accompanying increases in thickness of the \\HI\\ layer that are possibly a result of perturbations of the gravitational field by a relatively strong bar. } {} ", "introduction": "In paper I in this series \\citep{ofk2008a} we presented \\HI\\ observations of a sample of 8 edge-on, \\HI\\ rich, late-type galaxies. The aim of the project has been described there in detail. Briefly, we attempt to put constraints on the flattening of dark halos around disk galaxies by measuring the force field of the halo vertically from the flaring of the \\HI\\ layer and radially from rotation curve decomposition. For the vertical force field we need to determine in these galaxies both the velocity dispersion of the \\HI\\ gas (preferably as a function of height from the central plane of the disk) and the thickness of the \\HI\\ layer, all of these as a function of galactocentric radius. In addition we also need to extract information on the rotation curve of the galaxy and the deprojected \\HI\\ surface density, also as a function of galactocentric radius. To adequately constrain the halo shape these parameters needed to be measured out to as low surface densities as possible.Here we present the determination of these properties from our \\HI\\ observations. \\begin{figure*}[t] \\centering \\includegraphics[width=16cm]{12567fg1.eps} \\caption[ESO074-G015 (IC5052): Rotation curve, \\HI\\ velocity dispersion and face-on \\HI\\ surface density]{ESO074-G015 (IC5052). This figure shows the results of the modelling of the kinematics. The best fit measurements for each kinematic function are shown in as (red) crosses for the the \\HI\\ velocity dispersion in the top left-hand panel and in the lower left-hand plot for the deprojected \\HI\\ surface density. The top right-hand panel shows the derived rotation curve (dark/red) overlaid on the observed XV diagram. Also shown on the XV diagram are the rotation fits obtained using the less accurate envelope tracing method (dark grey/orange) and the peak flux method (light grey/yellow). The latter always shows a lower rotation velocity than the other two fits. The lower right-hand plot displays the derived rotation curve with the receding and approaching sides separately (in the electronic version red and blue respectively). The FWHM synthesised beam has dimensions $292\\times292$ pc.} \\label{fig:ch5-e074-g015-kine} \\end{figure*} In order to resolve the vertical structure, we selected a sample of nearby, \\HI\\ rich, late-type edge-on galaxies which are likely to exhibit large \\HI\\ flaring expected from high dark matter fractions. From paper I, we recall that the gradient of the force field in the vertical direction for a Gaussian gas layer can be written as \\begin{equation} \\label{eq:ch1-hydro} \\frac{\\partial K_z}{\\partial z} = - \\frac{\\sigma_{v,g}^2}{({\\rm FWHM}_{z,g}/2.35)^2}, \\end{equation} where $\\sigma_{v,g}^2$ is the vertical velocity dispersion of the gas and FWHM$_{z,g}$ the gas layer thickness. Using the flaring of the \\HI\\ distribution and the derived density distribution of the gas and stellar distributions, it is possible to measure the halo shape over the \\HI\\ extent of the luminous disk. For a given vertical gas velocity dispersion, the flaring will decrease and the gas surface density will increase, for increasingly flattened dark halos. Thus, in gas-rich late-type galaxies it is possible to measure the force field of the halo vertically from the gas layer flaring and radially from rotation curve decomposition. In the second paper \\citep{ofk2008b}, paper II, we presented a new method to accurately determine the rotation curve, deprojected \\HI\\ surface density and velocity dispersion at all radii in a edge-on gas disk. The superposition of velocity profiles from many radii in each sightline through an edge-on \\HI\\ disk tends to cause an overestimate of the velocity dispersion with most measurement methods (see paper II). Measuring the radial flaring profile requires a model of the galaxy rotation and face-on surface density; this also necessitates high accuracy rotation curve measurement and \\HI\\ surface density deprojection. In this third paper we present the rotation curve, \\HI\\ velocity dispersion and deprojected \\HI\\ surface density from the XV maps measured using the iterative XV diagram modelling program (see paper II). Previously, measurement of these parameters in edge-on galaxies has been difficult. Indeed, the gas velocity dispersion has been measured as a function of radius in only one edge-on galaxy NGC4244 \\citep{olling1996a}. If the data have suffiently high S/N, then our iterative XV modelling program is capable of deriving the variation of rotation, velocity dispersion and density as a function of both radius and height $z$ above the equatorial plane of these edge-on galaxies. For the data described in paper I, the S/N is not high enough to derive the variations in both $z$ and $R$. We use the XV diagrams integrated over $z$, and measure the gas surface density distribution and kinematics as a function of radius only, assuming that the gas velocity dispersion is isotropic and vertically isothermal. ", "conclusions": "\\label{sec:ch6-summ} Using our new XV modelling method detailed in paper II we have successfully fit the rotation curves of all eight galaxies in our sample; and the \\HI\\ velocity dispersion and surface density in all the galaxies except for the inner disk of IC2531. The rotation curve was the most reliably fit achieving an accuracy of roughly 2-4 \\kms. The surface density was also very reliably measured. The \\HI\\ velocity dispersion was more sensitive to local variations and noise in the XV map, however despite these challenges the uncertainty in the derived \\HI\\ velocity dispersion is only about 1--1.5 \\kms, which is comparable to the accuracy obtained from NGC4244 by \\citet{olling1996a}. The prevalence of bars in the larger galaxies and the increase in the mean velocity dispersion with maximum rotation speed suggests that bar buckling may be a common cause of widespread gas heating in disk galaxies. The measured rotation curves show that all galaxies in our sample display differential rotation, even the smallest galaxies IC5052 and ESO115-G021 which have peak rotation speeds of only $60$ \\kms. The increased velocity dispersion by 1--2.5 \\kms\\ observed in galaxies with bars suggests that shocks caused by disk instabilities in bars can cause gas heating. Iterative fitting of the \\HI\\ layer thickness enabled us to reliably measure the \\HI\\ flaring in 7 of the 8 galaxies in our sample, despite FWHM$_{\\theta}$ beamwidths of up to $1.3$ kpc and a low peak channel map signal-to-noise of around $13$ for most of the galaxies in our sample. The only galaxy we were unable to reliably fit was ESO138G-014 which was missing large scale \\HI\\ structure due to an absence of short baseline observations, and also suffered low peak signal-to-noise ($15$) and a large telescope beam ($0.965$ kpc). For galaxy observations with only two of these three handicaps it was possible to iteratively recover a flaring model that fit the observed \\HI\\ channel map distribution. UGC7321 is the most accurately measured flaring profile in our sample with errors of only 15-50 pc over most of the radial extent. The flaring profiles show an impressive degree of the symmetry between flaring measurements on either side of the galaxy --- this is most clearly seen in UGC7321. Axial symmetry of the gas layer thickness is expected as in the absence of local variations of the stellar surface density or gas velocity dispersion, the vertical flaring should reflect the smooth symmetric gravitational potential of the dark matter halo. The exceptionally smooth flaring profile of UGC7321 shows that, on the $0.7$ kpc scale of telescope beam, the stellar surface density and dark matter halo gravitational potential are also smooth. The most remarkable feature of \\HI\\ flaring profiles is a common shape displaying a linear increase at low radii and steepening to an exponential increase in the outer disk. The steepness of the inner flaring should be dominated by the radial surface density, while at the stellar disk truncation radius the flaring becomes exponential and is dominated by the dark halo shape. This typical flaring shape is found in all of the galaxies except ESO138-G014 which was incompletely imaged, and ESO274-G001 which appears to have a more linear flaring profile. With the exception of ESO146-G014, all of the less massive galaxies ($V_{max} < 100$ \\kms) display steep inner flaring profiles. The four larger galaxies ($V_{max} > 100$ \\kms), excluding ESO138-G014, display much shallower inner flaring profiles. The high sensitivity of the UGC7321 observations allowed us to discern variations from azimuthal symmetry of both the \\HI\\ midplane position and the \\HI\\ thickness. We show that at radii from approximately $2.5$ to $6$ kpc along the Y=0 axis, the FWHM$_z$ gas thickness varies by up to $500$ pc from the azimuthally averaged flaring in those radial annuli. At this same locations along the Y=0 axis the gas midplane tilts, dropping $\\sim$200 pc below the average disk midplane on the left-hand side, and rising $\\sim$200 pc above the midplane on the right-hand side. These local deviations are considered to be best observational evidence for a bar in an edge-on disk galaxy. The next paper in this series will be devoted to an analysis of the data on UGC7321 as a first attempt at deriving constraints on the flattening of a dark halo." }, "1003/1003.3423_arXiv.txt": { "abstract": "{Chemical models of dense cloud cores often utilize the so-called pseudo-time-dependent approximation, in which the physical conditions are held fixed and uniform as the chemistry occurs. In this approximation, the initial abundances chosen, which are totally atomic in nature except for molecular hydrogen, are artificial. A more detailed approach to the chemistry of dense cold cores should include % the physical evolution during their early stages of formation.} {Our major goal is to investigate the initial synthesis of molecular ices and gas-phase molecules as cold molecular gas begins to form behind a shock in the diffuse interstellar medium. The abundances calculated as the conditions evolve can then be utilized as reasonable initial conditions for a theory of the chemistry of dense cores.} {Hydrodynamic shock-wave simulations of the early stages of cold core formation are used to determine the time-dependent physical conditions for a gas-grain chemical network. We follow the cold post-shock molecular evolution of ices and gas-phase molecules for a range of visual extinction up to $A_{\\rm V} \\approx 3$, which increases with time. At higher extinction, self-gravity becomes important. } { As the newly condensed gas enters its cool post-shock phase, a large amount of CO is produced in the gas. As the CO forms, water ice is produced on grains, while accretion of CO produces CO ice. The production of CO$_{2}$ ice from CO occurs via several surface mechanisms, while the production of CH$_4$ ice is slowed by gas-phase conversion of C into CO. } {} ", "introduction": "The formation of molecular clouds from the diffuse atomic interstellar medium has been the subject of much interest. Proposed mechanisms have been reviewed in the recent literature \\citep{ABI09, BT07, MO07}. Despite this interest, most chemical models of cold dense interstellar clouds utilize the simple pseudo-time-dependent approximation, in which the physical conditions are homogeneous and fixed, beginning from an already dense, cold, and darkened state with all H in H$_2$. Moreover, the initial abundances for heavy elements are assumed to be atomic. Yet, the high abundance of molecular hydrogen in diffuse clouds suggests that molecules can be synthesized as clouds form as well as during their existence. One approach to dense cloud formation, that of hydrodynamic shock waves, has been studied by \\citet{Bergin04} (hereafter B04), who showed that the gas-phase molecules H$_2$ and CO are produced early in the formation of the cloud. In this paper, we revisit the approach of B04, and include a more complex treatment of the gas-grain chemistry that evolves in tandem with the post-shock physical conditions. We focus on the initial production of major solid phase and gaseous molecules in the post-shock material as the extinction gradually increases. Prior investigations have considered some aspects of the problem we discuss here. Early dynamic models of dense and diffuse clouds were reviewed by \\citet{Wi88}. Cyclic models involving shocks were studied by \\citet{NW92}. Our approach is similar to the constant-pressure collapse model of \\citet{Pineau91}, although that paper did not include grain surface chemistry. Such chemistry was included by \\citet{RH01III} in a pseudo-time-dependent model of the formation of mantle ices using an earlier version of our gas-grain network. These authors were mainly interested in explaining the formation of CO$_{2}$ ice, which had been under-produced in their previous models. A pseudo-time-dependent approach with surface chemistry was also used to study the water ice distribution at various conditions throughout the Taurus dark cloud \\citep{Nguyen02}. ", "conclusions": "\\label{concsection} The chemistry of cold dense cores has typically been treated by pseudo-time-dependent methods, in which physical conditions are homogeneous and time-independent. Also, the initial chemical abundances are artificially assumed to take the form of atoms except for molecular hydrogen. Although such gas-phase and gas-grain models of the chemistry are often in reasonable agreement with observations and even predictive in nature for the gas-phase species, the assumption of time-independent physical conditions ignores the formation of the core, a process that cannot be totally distinct from the chemistry that occurs. For these reasons, we have started a program to determine how the chemistry is altered by consideration of the physical evolution of dense cores. In this paper, we have reported our study of the chemistry that occurs during shock-induced formation of dark cloud material from the more diffuse background, with an emphasis on the composition of icy grain mantles and gaseous CO. In this picture, gas-phase and grain-mantle molecular species form after the post-shock material cools and becomes dense, and as the visual extinction increases, in contrast with the assumption of constant physical conditions. This stage can be considered as preliminary to the stage in which cold pre-stellar cores form. Since our approach does not extend to the phase in which the self-gravity governs the evolution of these cooled objects, we cannot continue to increase their size and visual extinction beyond an extinction of $\\approx 2-3$. Although our calculated column densities cannot then be compared with those of ordinary cold dense, or prestellar, cores, we have compared them with the observed abundances of diverse low-$A_V$ sources within the Taurus dark cloud, and a variety of other diffuse-to-translucent sources. There are several key results of our calculations: \\begin{enumerate} \\item{Despite some divergence at lower extinction, all model results for CO(g) tend to converge as the edge-to-center extinction reaches $\\sim 2.0$ (an H$_{2}$ column of $4 \\times 10^{21}$ cm$^{-2}$) where CO(g) becomes the dominant form of gaseous carbon. We conclude that homogeneous gas-grain dense core models should not use atomic carbon as an initial abundance, but would be more realistic with initial abundances from the low-extinction results reported here. In particular, the calculated CO(g) columns and [C$^+$ + C]/CO ratios for the different models are not drastically different from the values obtained from gas-phase models at steady-state as a function of visual extinction or molecular hydrogen column density. Our suggestion concerning the proper initial abundances to use has also appeared in a recent paper by \\citet{2009A&A...508..783L}, who advocated that ``mechanisms for the chemical evolution from diffuse to dark gas should be included in model calculations'', a view that is the basis for our paper.} \\item{By $A_V = 2.5$, all models show that H$_2$O(s) and CO(s) have grown above their minimum observable abundances. Models that form denser gas grow these ices at lower $A_V$. These ices thus become abundant before a full-fledged dense core can be produced.} \\item{The growth of CO(s) slows down the rate of synthesis of CH$_{4}$(s) because there is less of its precursor, neutral C, available. By an extinction of 3.0, the methane column densities for individual models tend to lie below those of CO(s), especially for the B parameter set, in which the dust temperature is higher.} \\item{The formation of CO$_2$(s) occurs efficiently only with the B parameter models, especially with the higher density models, 2-B and 4-B. The grain temperature is the most critical parameter for efficient CO$_2$(s) formation, with or without the Eley-Rideal mechanism. } \\item{Models excluding photodesorption rates would form multiple layers of H$_2$O(s) and CO(s) at lower $A_V$ than either thresholds or the lowest reported firm detections \\citep{dougco207}. This result indicates that the inclusion of measured photodesorption rates is critical for realistic models.} \\end{enumerate} Finally, and most importantly, our results show that if the shock model of B04 is correct, the early stages of both gas-phase and grain-surface chemistry occur as dense material is being formed, so that the chemistry of the cold cores that form from the final collapse of dense material should betray some indications of the physics of core formation. Models that follow the chemistry as the dark cloud material formed in our model collapses to produce dense pre-stellar cores are desirable. In addition, other dynamical scenarios than our shock model do exist, and we hope to explore them." }, "1003/1003.3615_arXiv.txt": { "abstract": "A survey of young bipolar outflows in regions of low-to-intermediate-mass star formation has been carried out in two class I methanol maser transitions: $7_0-6_1A^+$ at 44~GHz and $4_{-1}-3_0E$ at 36~GHz. We detected narrow features towards NGC~1333I2A, NGC~1333I4A, HH25MMS, and L1157 at 44~GHz, and towards NGC~2023 at 36~GHz. Flux densities of the lines detected at 44~GHz are no higher than 11 Jy and the relevant source luminosities are about $10^{22}$~erg~s$^{-1}$, which is much lower than those of strong masers in high-mass star formation regions. No emission was found towards 39 outflows. All masers detected at 44~GHz are located in clouds with methanol column densities of the order of or larger than a few $\\times 10^{14}$~cm$^{-2}$. The upper limits for the non-detections are typically of the order of 3--5~Jy. Observations in 2004, 2006, and 2008 did not reveal any significant variability of the 44 GHz masers in NGC~1333I4A, HH25MMS, and L1157. ", "introduction": "Bright and narrow maser lines of methanol (CH$_3$OH) have been found towards many star-forming regions~\\citep{hmb,menten91a,kurtzetal}. According to the classification of~\\citet{menten91b}, methanol masers can be divided into two classes, I and II, with each class characterized by a certain set of transitions. The Class I maser transitions are the $7_0-6_1A^+$ transition at 44~GHz, $4_{-1}-3_0E$ transition at 36~GHz, $5_{-1}-4_0E$ transition at 84 GHz, $8_0-7_1A^+$ transition at 95~GHz etc., while the Class II transitions are the $5_1-6_0A^+$ transition at 6.7~GHz, $2_0-3_{-1}E$ transition at 12~GHz, the $J_0-J_{-1}E$ series of transitions at 157~GHz, etc. A list of the most powerful Class I and II transitions is presented in e.g.,~\\citet{cragg}. Both Class I and Class II masers are often overlaid upon broad thermal lines. Many methanol transitions, e.g., the series of $2_K-1_K$ transitions near 96~GHz, $3_K-2_K$ transitions near 145~GHz, $5_K-4_K$ transitions near 241~GHz etc. never exhibit maser features. These transitions are often called ``purely thermal\". The nature of methanol masers is still unknown. \\citet{plammen} suggested that Class I masers arise in postshock gas in the lobes of bipolar outflows, where the abundance of methanol is enhanced due to grain mantle evaporation. This hypothesis has further support in the fact that in a number of star-forming regions Class I masers appear to be associated with outflows~\\citep{kurtzetal,chen}. However, this hypothesis is not generally accepted because there are no high-velocity Class I masers and the apparent association between the masers and the outflows may be caused by the fact that both of them arise in the same regions of star formation rather than by a physical association between these objects. The difficulties in the exploration of methanol masers partly appear because until recently the masers have been observed in regions of massive star formation, which are relatively distant (2--3 kpc from the Sun or farther) and highly obscured at the optical and even NIR wavelengths. In addition, high mass stars usually form in clusters. These properties make it difficult to resolve maser spots and to associate masers with other objects in these regions. In contrast, regions of low-mass star formation are much more widespread and many of them are only 200--300 pc from the Sun; they are less heavily obscured than regions of high-mass star formation, and there are many isolated low-mass protostars. Therefore, the study of masers in these regions might be more straightforward compared to that of high-mass regions, and hence, the detection of Class I masers in low-mass regions might have a strong impact on maser exploration. Bearing this in mind, we performed in 2004 a~``snapshot\" search for Class I methanol masers towards bipolar outflows driven by low-mass YSOs~\\citep{paperI} (Paper I) at 44, 84, and 95~GHz. The source list consisted of five so-called chemically rich outflows, where the abundances of methanol and some other molecules are increased as a result of grain mantle evaporation. The search proved successful: three maser candidates, NGC~1333I4A, HH25, and L1157 were found at 44~GHz. VLA observations of L1157 at 44~GHz confirmed that this source is really a maser~\\citep{kalen10}. Therefore a further work in the field looks promising. In order to obtain a general idea about the main properties of the Class I masers in the regions of low-mass star formation, we performed a more extended search for these objects. The new survey was carried out at the frequency of the $7_0-6_1A^+$ transition, but most sources were additionally observed in another Class I maser line, the $4_{-1}-3_0E$ line at 36~GHz. The physical relation between low-mass protostellar outflows and Class I methanol masers is poorly understood. Hence, it is tempting to make a comprehensive survey of such outflows. Unfortunately, the enormous amount of observing time makes such a survey unrealistic. The naive expectation, supported by the successful search of ~\\citet{paperI} is to find methanol masers towards bright thermal sources of methanol; therefore the basis of our source list consists of outflows where~\\citet{kalen07} detected thermal emission in the $5_{-1}-4_0E$, $8_0-7_1A^+$, and/or $2_K-1_K$ methanol lines. We included also three outflows, IRAS03282, Serpens S68FIRS1 and Serpens SMM4, where~\\citet{bach95} and~\\citet{garay02} found a significant enhancement of methanol abundance relative to that in quescent gas. Because methanol enhancement has been detected in young, well-collimated outflows from Class 0 and I sources, we included several such objects in our list regardless of whether methanol enhancement had been found there. A subsample of our list consisted of YSOs with known outflows and/or H$_2$O masers located in Bok globules~\\citep{yc92,cbh2o}. Like other objects from our list, these YSOs are typically isolated objects of low or intermediate mass, located in nearby ($<$500~pc) small and relatively simple molecular clouds. In total, our source list consisted of 37 regions which harbor 46 known outflows driven by Class 0 and I low-mass protostars, taken from the literature. The observed sources, positions, and the relevant literature are given in Tables~\\ref{gauss} and~\\ref{negative}. In addition to the survey we performed second- and third-epoch 44~GHz observations of the maser candidates detected in 2004. \\begin{table} \\centering \\begin{minipage}{80mm} \\caption{The parameters of the observed lines and those of the OSO 20-m at the line frequencies. The parameters of the lines, observed in the preliminary survey~\\citep{paperI} are included for completeness. } \\label{freqtabl} \\begin{tabular}{|l|c|c|c|c|} \\hline Transition &Frequency &$S\\mu^2\\;^a$& HPBW & G \\\\ & (GHz) & (Debye) &$('')$ & (Jy/K) \\\\ \\hline $7_0-6_1A^+$ & 44.069476 & 6.1380& 82 &20.5\\\\ $5_{-1}-4_0E$ & 84.521206 & 3.0830& 44 &22\\\\ $8_0-7_1A^+$ & 95.169516 & 7.2211& 39 &25\\\\ $2_{-1}-1_{-1}E$ & 96.739393 & 1.2133& 39 &25\\\\ $2_0-1_0A^+$ & 96.741377 & 1.6171& 39 &25\\\\ $2_0-1_0E$ & 96.744549 & 1.6167& 39 &25\\\\ $2_1-1_1E$ & 96.755507 & 1.2443& 39 &25\\\\ $4_{-1}-3_0E$ & 36.169290 & 2.5184& 105 &18\\\\ \\hline \\end{tabular} \\medskip $^a$--The product of the permanent dipole moment and the line strength from~\\cite{muller}\\\\ \\end{minipage} \\end{table} \\begin{figure*} \\includegraphics[width=150mm]{kalenskii_1.eps} \\caption{Upper panels: spectra of the regions of low-mass star formation in which maser candidates in the $7_0-6_1A^+$ line were detected. Shown from top to bottom are the $7_0-6_1A^+$, $8_0-7_1A^+$, $5_{-1}-4_0E$, $4_{-1}-3_0E$, and $2_K-1_K$ lines at 44, 95, 84, 36, and 96~GHz, respectively. The horizontal axis plots the radial velocity in km~s$^{-1}$ and the vertical axis the spectral flux density in Jansky. The 84~GHz, 95~GHz, and 96~GHz spectra are taken from Paper I for all sources except L1157, for which the new spectra at all these frequencies except 96 GHz have been taken towards the stronger maser position. Lower panels: spectra of other sources detected at 44 and/or 36~GHz.} \\label{mas} \\end{figure*} \\begin{table*} \\centering \\begin{minipage}{150mm} \\caption{The gaussian parameters of detected lines. Here 44 denotes the $7_0-6_1A^+$ at 44~GHz, 36, the $4_{-1}-3_0E$ line at 36~GHz. For NGC~2023, the numbers in parentheses show the R.A. and DEC. offsets in arcsec from the position given in the second and third columns.} \\label{gauss} \\begin{tabular}{|l|c|r|c|c|c|c|r|c|c|l|} \\hline Source &R.A. & DEC &Line&$\\int S_{\\nu}dV$ &$V_{\\rm LSR}$& FWHM &$S_{\\nu}$&Obs.&Notes$^a$&Refs\\\\ &(J2000)&(J2000)&$ $&(K$\\cdot$km~s$^{-1}$)&(km~s$^{-1}$)&(km~s$^{-1}$)&(Jy)&(year)&&\\\\ \\hline NGC~1333I2A &03 29 01.0& 31 14 20 &44& 4.92(1.64)& 9.24(0.12)&1.66(0.42)& 2.81 & 2007 & r &1,2\\\\ & & & & 5.74(1.64)& 11.57(0.31)&2.71(0.79)& 1.97 & & & \\\\ % & & &36& 13.7(1.44)&--7.10(0.37)&7.88(1.04)& 1.64 & 2007 & & \\\\ % & & & & 13.7(2.70)& 8.07(0.15)&2.00(0.26)& 6.43 & & & \\\\ % & & & & 10.4(4.14)& 12.48(1.25)&6.72(1.36)& 1.46 & & & \\\\ & & & & 33.7(4.86)& 10.54(0.08)&2.49(0.25)&12.71 & & & \\\\ NGC~1333I4A &03 29 10.3& 31 03 13 &36& 10.6(1.98)& 2.73(0.38)& 4.61(1.22)& 2.14 & 2007 & c &1,3\\\\ % & & & & 7.56(1.44)& 7.11(0.16)& 1.97(0.34)& 3.60 & & & \\\\ & & &44& 3.28(0.41)& 6.49(0.26)& 5.10(0.59)& 1.95 & 2007 & & \\\\ & & & & 1.85(0.41)& 7.51(0.02)& 0.33(0.05)& 5.13 & & & \\\\ NGC~2023(0,0)&05 41 28.5&--02 19 19&44& & & &$<3.60$& 2007 & b & 4\\\\ % (0,0) & & &36&15.3(0.7) & 6.46(0.02)& 0.92(0.05)&15.73 & 2008 & b & \\\\ (20,20) & & & & 18.5(0.82)& 6.39(0.02)& 0.99(0.05)&17.39 & & b & \\\\ ($-60,-60$) & & & & & & &$<3.6$& & q & \\\\ ($60,-60$) & & & & & & &$<3.0$& & c & \\\\ (60,60) & & & & 0.63(0.04)& 6.55(0.04)& 1.36(0.11)& 0.44 & & q & \\\\ ($-60,60$) & & & & & & &$<3.6$& & q & \\\\ & & & & & & & & & & \\\\ HH25MMS &05 46 06.5&--00 13 54&44& 5.33(0.82)&10.51(0.04)& 0.48(0.09)& 10.41& 2007 & r;mo&1,5\\\\ & & &36& 11.9(1.08)&10.56(0.12)& 2.93(0.28)& 3.82 & 2007 & & \\\\ S68N &18 29 47.5& 01 16 51 &44& 12.3(1.64)& 8.86(0.32)&5.18(0.83)& 2.23 & 2007 & c;mo& 6 \\\\ % Serpens CB2 &18 29 58.4& 01 13 35 &44& 8.61(1.63)& 8.13(0.64)&6.10(0.87)& 1.32 & 2007 & b;mo& 7 \\\\ % L1157 &20 39 08.1& 68 01 14 &44& 12.0(0.60)& 0.69(0.08) &3.82(0.24)& 3.4 & 2004 & &1,8,\\\\ & & & & 2.40(0.20)& 0.75(0.01) &0.37(0.03)& 6.2 & & &9,10 \\\\ &20 39 10.0& 68 01 42 &36& 36.6(3.60)&--0.70(0.20)&3.97(0.31)& 8.6 & 2008 & m & \\\\ & & & & 10.8(3.20)& 1.40(0.12)&1.66(0.27)& 6.2 & & & \\\\ & & &44& 15.1(0.57)& 0.61(0.08)&3.90(0.20)& 3.6 & & & \\\\ & & & & 2.0(0.32)& 0.91(0.03)&0.53(0.08)& 3.5 & & & \\\\ \\hline \\end{tabular} \\end{minipage} \\flushleft $^a$--r, red wing; b, blue wing; c, central position; q, quescent gas; mo, multiple outflows; m, maser position determined with the VLA. 1--Kalenskii et al. (2007); 2--Bachiller et al. (1998); 3--Blake et al. (1995); 4--Sandell et al. (1999); 5--Gibb \\& Davis (1998); 6--Garay et al. (2002); 7--Davis et al. (1999); 8--Bachiller et al. (1995); 9--Bachiller et al. (2001); 10--Benedettini et al. (2007); \\end{table*} ", "conclusions": "A survey of young bipolar outflows in regions of low-to-intermediate-mass star formation has been carried out in two Class I methanol maser transitions, $7_0-6_1A^+$ at 44~GHz and $4_{-1}-3_0E$ at 36~GHz. As a result of the survey we detected narrow features at 44~GHz towards NGC~1333I2, NGC~1333I4A, HH25MMS, and L1157. One more maser candidate was detected at 36~GHz towards the blue lobe of a bipolar outflow driven by a low-mass YSO in the NGC~2023 region. Flux densities of the lines detected at 44~GHz are no higher than 11 Jy and their luminosities are about $10^{22}$~erg~s$^{-1}$, which is much lower than those of strong maser lines in regions of high-mass star formation. No emission was found towards 39 outflows. The upper limits for the non-detections are typically of the order of 3--5~Jy. Thus, new masers in regions of low-mass star formation should be searched for with a sensitivity of 1~Jy or better. Observations at 44 GHz in 2004, 2006, and 2008 did not reveal a significant variability of the masers in NGC~1333I4A, HH25MMS, and L1157. All masers at 44~GHz in these low-mass star formation regions were found in clouds with methanol column densities of several times $10^{14}$~cm$^{-2}$ at linear scales of about 0.06 pc. Even higher methanol column densities have been reported towards stronger masers in regions of massive star formation. Therefore, the following trend seems to exist: molecular clouds with methanol column densities less than $10^{14}$~cm$^{-2}$ cannot produce 44~GHz masers; clouds with methanol column densities $10^{14}-10^{15}$~cm$^{-2}$ can produce weak masers with luminosities about $10^{22}$~erg~s$^{-1}$; clouds with methanol column densities higher than $10^{16}$~cm$^{-2}$ can produce strong masers with luminosities $10^{24}-10^{25}$~erg~s$^{-1}$." }, "1003/1003.2012_arXiv.txt": { "abstract": "We present a flexible interactive 3D morpho-kinematical modeling application for astrophysics. Compared to other systems, our application reduces the restrictions on the physical assumptions, data type and amount that is required for a reconstruction of an object's morphology. It is one of the first publicly available tools to apply interactive graphics to astrophysical modeling. The tool allows astrophysicists to provide a-priori knowledge about the object by interactively defining 3D structural elements. By direct comparison of model prediction with observational data, model parameters can then be automatically optimized to fit the observation. The tool has already been successfully used in a number of astrophysical research projects. ", "introduction": "The interpretation of astrophysical data often depends strongly on the knowledge of depth information along the line of sight. In most cases, however, this is the least well known information. That is true for the distance and especially for the position of substructure {\\em within} an object. The development of effective methods for the reconstruction of the 3D structure of astrophysical objects is therefore an issue of growing importance in astronomy. Photographic images only provide a two-dimensional integration of the emission and absorption along the line of sight. The depth information is therefore flattened. Sometimes, symmetry properties combined with a favorable orientation of an object provide sufficient information to visually deduce what the structure must be. This can be the case for planetary nebulae and has been used to automatically reconstruct the 3D structure (Leahy~\\cite{Leahy91:DEA}, Magnor et al.~\\cite{Magnor04:CIVRPN,Magnor05:RVPN}, Lin\\c{t}u et~al.~\\cite{Lintu07:MDGDPN,Lintu07:3DEAPN}). If no such symmetries are present, then the depth information must come from other types of information, which usually depend on a fundamental physical model for the object class that is considered. This information could be the velocity field, e.g. in a radially expanding nebula a mapping between velocity and position exists. However, for some objects -- such as turbulent interstellar clouds -- such a mapping is not possible. Much of observational astrophysics research involves physical modeling with limited constraints to deduce physical properties of the observed objects. Astrophysicists measure a limited number of physical properties of an object, via electromagnetic waves, to which a physical model of the phenomenon is then fitted. Such models are usually not unique. Most of the astrophysical modeling effort tends to gravitate towards massive parallel supercomputing for dynamical simulations. Analysis and visualization of such simulations are done separately and often are complex and computationally intense processes themselves~\\cite{Henney09}. While such simulations produce insight into generic astrophysical processes, they are rarely suitable for elucidating the properties and structure of particular objects. Knowing the properties of individual objects is essential when a single object class, e.g. planetary nebulae, shows a large variety of presentations in images. A serious difficulty for the modeling of particular objects is our fixed vantage point on Earth which restricts all observations to be along a single direction (up to the parallax provided by the Earth's orbit, which is negligible % for typical distances to astronomical nebulae). This is in strong contrast to, e.g., medical imaging where 3D information is recovered from observations from multiple directions around the subject. For any astronomical object beyond the solar system, we are able to observe only one 2D projection of its actual 3D volumetric shape. For the correct physical interpretation of observational data, information about the object's actual 3D shape has to be available~\\cite{Lucy74:ITR,Bremer95:TDI,Palmer94:DAG,Leahy94:IDA,Volk93:DPN}. Obtaining new structural information and insight on particular objects is the main purpose of the application that we present in this paper. Our approach to modeling individual objects is very different from previous methods. The application that we present (called {\\it Shape}) becomes essential when automatic reconstruction methods fail because theoretical or observational constraints are insufficient. A lack of constraints for an automatic reconstruction is at least partially compensated by scientific user judgement. Often the available constraints are sufficient to test one or more hypotheses about the structure of an object. In such cases, rather than a reconstructive, a constructive morpho-kinematical modeling approach is more suitable to reconstruct the structure and velocity field. The term morpho-kinematical is applied to modeling that involves only structural (morphological) and velocity (kinematic) information. This is in contrast to dynamical simulations, which include the effects of forces and temporal evolution from a set of simpler initial and boundary conditions. In general, the outcome of dynamical simulations is not predictable in detail and very hard to tune to a specific object. Conventional morpho-kinematical modeling uses hard-coded mathematical descriptions of the objects, processes and boundary conditions. Therefore, the user needs at least basic programming skills in the particular language of the code. Modern 3D modeling software of the graphics industry shows that such modeling can be done effectively without user programming intervention. Although such software can visualize gas-dynamical processes, it is inefficient and its usefulness for astrophysical processes is limited. The general workflow of such systems seemed, however, very suitable for modeling particular astrophysical objects~\\cite{Steffen06}. Following the technique of modern interactive 3D graphics systems, we have developed {\\it Shape} with specialized functionality for interactive astrophysical modeling on single desktop or laptop computers. The primary purpose of {\\it Shape} is to interactively generate 3D models. However, in contrast to conventional astrophysical modeling tools, it integrates the visualization and analysis of the model into the same system. Direct access to the model data at any stage of the modeling process allows for effective comparison with the observed data in a variety of ways within the feedback loop of the iterative workflow (Figure~\\ref{workflow}). In artistic work on astronomical topics, commercial tools like Maya or 3D Studio Max are frequently being employed~\\cite{Yu05:DFD,Davis05:STA,Matthews05:DDF}. Professional animation tools are designed to assist in creating realistic 3D scenes of familiar environments. Unfortunately, when used for scientific work they display serious shortcomings. Especially volume rendering with mesh structures and particle systems are very different from the physical correct radiation transfer needed for reliable interpretation of astrophysical phenomena. A qualitative and quantitative comparison of such models with real objects is not possible. The key problem preventing their use in astrophysics research is the inability to produce the type of renderings that are comparable to the observations obtained with telescopes and other scientific instrumentation like spectrographs (e.g., for Doppler-shift measurements). With \\emph{Shape} we remedy most of the shortcomings of previous astrophysical reconstruction systems by applying the powerful structure modeling techniques of commercial animation suites, while adding the information output and processing systems that are necessary for astrophysical research applications. We go beyond the current commercial rendering techniques by using physically more accurate modeling of the radiation transfer from the sources to the observer. In this paper we first comment on previous related work in Section \\ref{related}, on the type of observational data that are used for this work in Section \\ref{data} and then introduce the \\emph{Shape} system. In Section \\ref{results} we show three examples of previously published research applications of \\emph{Shape}, before giving an outlook on future developments and our conclusions in Sections \\ref{future} and \\ref{conclusion}, respectively. \\begin{figure}[t] \\centering \\includegraphics[width=0.45\\textwidth]{Figure1} \\caption{The \\emph{Shape} modeling tool that we present here provides a unified framework for the whole modeling and visualization workflow. The user employs his physical knowledge to construct an initial model which can be visualized and compared to observational data in several ways, allowing for easy interactive and iterative refinement of the model. When all necessary physical information is reflected in the model, its parameters can be automatically optimized, minimizing the difference between the model and the observational data. The final model can then be used to generate various types of graphical output.} \\label{workflow} \\end{figure} ", "conclusions": "\\label{conclusion} We have presented a novel 3D application for the modeling and reconstruction of astrophysical objects that incorporates interactive modeling tools. It considerably extends the capabilities of conventional reconstruction systems. This is achieved through the support of a system of construction and ``modifier'' tools that allow extremely complex structures and velocity fields to be assembled without the need for user programming. A number of visualization styles that are common for astronomical observations can be used to compare the model with the observed data. The workflow is enhanced by the ability to continuously compare the model to observational data during the modeling process as well as by an automatized optimization algorithm. The tool has been shown to cover the entire modeling and visualization pipeline of a common morpho-kinematical modeling task, supporting the scientifically accurate reconstruction of a wide class of astrophysical objects while keeping a convenient and user-friendly interface. The software we have presented has been thoroughly tested and applied in a number of astronomical research projects, some of which we have quoted as examples. In addition to scientific research, it may prospectively be applied for physically plausible artistic works or for the generation of astronomical animations for educational purposes, e.g. in digital planetariums. Non-astrophysical uses can also be imagined wherever velocity information is observed, e.g. in the field of Doppler radar observations of tornados and other weather phenomena. \\emph{Shape} is freely available as a Java WebStart application from its website at \\url{http://www.astrosen.unam.mx/shape/} \\begin{figure}[t] \\centering \\includegraphics[width=\\columnwidth]{Figure10} \\caption{The top panels show the observed image (center) and two different model renderings of nova RS Ophiuchi~\\cite{Ribeiro2009}. The model on the right includes a special rendering filter that corresponds to the transmission filter on the \\emph{Hubble Space Telescope}. It excludes some of the emission with very high velocities along the line of sight. The bottom panels show the spectral line profiles from the observations (noisy red line) and the final model (left, continuous line). The right panel compares the observed profile with models at different viewing angles that are just consistent with the observations (dashed and dotted lines). [Figure reproduced with permission of the authors (Ribeiro et~al.~\\cite{Ribeiro2009}).]} \\label{ribeiro_7} \\end{figure} \\begin{figure}[t] \\centering \\includegraphics[width=\\columnwidth]{Figure11} \\caption{The outburst of nova RS Ophiuchi has been modeled by Ribeiro et~al.~\\cite{Ribeiro2009} using an expanding bipolar nebula (mesh structure) with a nearly inert waist (shaded structure). Using the time modifier feature of \\emph{Shape} they were able to model the expansion between the first and second epoch observations. [Figure reproduced with permission of the authors (Ribeiro et~al.~\\cite{Ribeiro2009}).]} \\label{ribeiro_6} \\end{figure} \\newpage" }, "1003/1003.5318_arXiv.txt": { "abstract": "If gamma-ray bursts are sources of ultra-high energy cosmic rays, then radiative signatures of hadronic acceleration are expected in GRB data. Observations with the {\\it Fermi Gamma-ray Space Telescope (Fermi)} offer the best means to search for evidence of UHECRs in GRBs through electromagnetic channels. Various issues related to UHECR acceleration in GRBs are reviewed, with a focus on the question of energetics. ", "introduction": "{\\it Fermi } observations of GRBs provide a new probe of particle acceleration in the relativistic outflows of GRBs. Some generic features of the high-energy behavior of 10 Large Area Telescope (LAT) GRBs consisting of 8 long-soft and 2 short-hard GRBs at the time of this conference have been identified, as summarized by Omodei \\cite{omo09}. These are: \\begin{enumerate} \\item A delayed onset of the $\\gtrsim 100$ MeV radiation observed with the LAT compared with the start of the keV/MeV GBM radiation; \\item Long-lived LAT emission extending well after the Gamma ray Burst Monitor (GBM) radiation has fallen below background, as known previously for long duration GRBs from EGRET \\cite{hur94}; \\item Existence of a distinct hard spectral component in addition to a component described by the Band function in both long and short GRBs, confirming the discovery of such a component from joint BATSE/EGRET TASC analyses \\cite{gon03}. \\end{enumerate} In this contribution, the possibility that these features can be explained by UHECRs in GRBs is considered. Because protons and ions are weakly radiative compared to electrons, even with escaping energies $E \\approx 10^{20}$ eV needed to explain UHECRs, large amounts of nonthermal hadronic energy are required. We consider energetics arguments to constrain models for UHECRs from GRBs. ", "conclusions": "In this contribution, we have sketched the energy requirements for GRBs to be sources of UHECRs, and for electromagnetic signatures of ultrarelativistic hadrons to be found in the {\\it Fermi} data. The large amounts of energy needed has been noted many times in the past, whether from proton synchrotron \\cite{tot98,zm01} or photohadronic processes. Although the large energy requirements make uncomfortable demands on $\\gamma$-ray emission models, an internal consistency is found insofar as the baryon load in long GRBs must be large given their relative rarity within the GZK radius, and that enormous energies are available from the rotational and accretion energy in the newly forming black holes. Future {\\it Fermi} observations and the possibility of detecting PeV neutrinos from GRBs with IceCube could establish whether GRBs are the sources of UHECRs. \\vskip0.2in \\noindent {\\bf Acknowledgments:} I would like to thank Guido Chincarini and Bing Zhang for their kind invitation to the Venice Shocking Universe conference, Soeb Razzaque for discussions, and Justin Finke for a careful reading of the manuscript. This work is supported by the Office of Naval Research." }, "1003/1003.2965_arXiv.txt": { "abstract": "{We studied the stellar population in the central 6.6\\arcmin $\\times$ 6.6\\arcmin\\,region of the ultra-deep (1Msec) {\\em Chandra} Galactic field--- the ``Chandra bulge field'' (CBF) approximately 1.5 degrees away from the Galactic Center ---using the {\\em Hubble Space Telescope} ACS/WFC blue (F435W) and red (F625W) images. We mainly focus on the behavior of red clump giants -- a distinct stellar population, which is known to have an essentially constant intrinsic luminosity and color. By studying the variation in the position of the red clump giants on a spatially resolved color-magnitude diagram, we confirm the anomalous total-to-selective extinction ratio, as reported in previous work for other Galactic bulge fields. We show that the interstellar extinction in this area is $\\langle A_{\\rm F625W} \\rangle = 4$ on average, but varies significantly between $\\sim3-5$ on angular scales as small as 1 arcminute. Using the distribution of red clump giants in an extinction-corrected color-magnitude diagram, we constrain the shape of a stellar-mass distribution model in the direction of this ultra-deep {\\em Chandra} field, which will be used in a future analysis of the population of X-ray sources. We also show that the adopted model for the stellar density distribution predicts an infrared surface brightness in the direction of the ``Chandra bulge field'' in good agreement (i.e.~within $\\sim$15\\%) with the actual measurements derived from the {\\em Spitzer}/IRAC observations. ", "introduction": "Our Galaxy hosts discrete X-ray sources of various types. The X-ray luminosities of Galactic sources range from $L_x\\sim10^{39}-10^{40}$ erg s$^{-1}$ for very luminous accreting neutron star and black hole binaries down to $10^{26}-10^{30}$ erg s$^{-1}$ for coronally active stars. Statistical properties of these different populations of sources are poorly known with large allowed ranges or uncertainties, which in some cases might lead to an underestimation of their contribution to the global X-ray emission of our Galaxy and other galaxies. Constraining the luminosity functions of the X-ray emitting populations is particularly relevant for understanding the faint, unresolved X-ray emission of the Galaxy that is distributed quite smoothly along the Galactic plane (the so called Galactic ridge X-ray emission, GRXE), which was discovered in early X-ray experiments \\cite[e.g.][]{worrall82}. Until recently, the nature of the GRXE remained unexplained, to a large extent due to our scarce knowledge of the statistical properties of populations of Galactic X-ray sources. Only recently it was shown that at least the majority of this unresolved X-ray emission arises from the cumulative emission of a large number of faint discrete sources, predominantly accreting white dwarfs and coronally active stars \\citep{revnivtsev06,revnivtsev09}. Understanding the origin of the GRXE would have been impossible without the progress in understanding the statistical properties of faint Galactic X-ray sources, which was achieved through X-ray all-sky surveys \\citep{sazonov06}. However, due to the limiting sensitivity of currently available all-sky surveys, these studies can probe only sources near the Sun, and the number of sources that are used to construct luminosity functions is small. This leads to significant uncertainties in the inferred properties and leaves the question about possible variations of populations throughout the Galaxy open. The resulting uncertainties have direct consequences for our studies of distant galaxies. As has been shown, the GRXE-type emission is a dominant component (after subtracting the contribution from low mass and high mass X-ray binaries with luminosities $L >10^{35}$ erg sec$^{-1}$) for a significant fraction of non-starburst galaxies \\cite[see e.g.][]{revnivtsev07,revnivtsev08}. Therefore it is important to understand the level of universality of the X-ray luminosity per unit solar mass inferred from studies of nearby objects. Another non-local sample of faint Galactic X-ray sources is required to address these questions. The brightest ($L_x\\gtrsim10^{32}$ erg s$^{-1}$) end of the luminosity functions of various classes of faint Galactic sources can be probed over distances of several kilo-parsecs by surveys of the plane with a limiting sensitivity $F_x\\approx10^{-15}-10^{-14}$ erg s$^{-1}$ cm$^{-2}$ and this is now extensively explored \\cite[e.g.][]{hands04,grindlay05}. But for fainter sources one needs to perform much deeper observations (with limiting sensitivity $\\sim10^{-16}-10^{-17}$ erg sec$^{-1}$ cm$^{-2}$) and choose a region with minimal obscuring dust and gas to avoid the effects of interstellar absorption. These observations with a total exposure time of 1~Msec were recently performed with the {\\em Chandra X-ray Observatory} in a direction approximately 1.5 degree away from the Galactic center. This field, located at the galactic coordinates ($l$,$b$)=(0.10$^{\\circ}$,$-$1.43$^{\\circ}$) and chosen for its low extinction \\cite[$A_V\\approx4$; see e.g.][]{drimmel03} and proximity to the Galactic Center, was first observed with {\\em Chandra} in 2005 for 100 ks (``LimitingWindow'' target) to study the nature and radial distribution of hard (2--10 keV) X-ray point sources in the inner Galactic bulge \\citep{hong09}. Simultaneously, images with the Advanced Camera for Surveys Wide Field Camera (ACS/WFC) on the {\\em Hubble Space Telescope} ({\\em HST}) were obtained of the central 6.6\\arcmin $\\times$ 6.6\\arcmin\\ region (centered at RA=267.86968 deg, Dec=-29.581033 deg , J2000) of the field observed by {\\em Chandra} to search for optical counterparts \\citep{vandenberg09}. Analogous to other inner-bulge fields included in this survey (e.g.\\, Baade's Window and Stanek's Window, at $\\sim$3.8-degree and $\\sim$2.2-degree offsets from the Galactic center, respectively) this field was called the Limiting Window, which refers to the rapidly declining opportunities for optical identification when moving closer to the heavily-obscured Galactic center region. In 2008, Revnivtsev et al., who referred to this region as the ``1.5degree field'', indicating its approximate angular distance from the Galactic Center---revisited this field with {\\em Chandra} for 900 ks to derive constraints on the resolved fraction of the GRXE \\citep{revnivtsev09}. Below we refer to this field as the \\1p5 -- ``Chandra Bulge field''. \\begin{figure}[htb] \\includegraphics[width=0.9\\columnwidth]{./13527_fig1.pdf} \\caption{Three-color optical images of the field containing the ultra-deep {\\em Chandra} observations of the \\1p5. The upper panel shows the mosaic image obtained by the 1.5m Russian-Turkish telescope (RTT150 telescope) (red shows the {\\em i'} spectral band, green - {\\em r'}, blue - {\\rm g'}), while the lower panel shows a smaller area observed with the {\\em HST} ACS/WFC (red - F625W, green - F658N, blue - F435W). From the top image it is evident that the interstellar extinction is strongly variable across the field (the regions close to boundaries of the figure are darker than the central regions). It is seen that the majority of stars are relatively faint and yellow -- these are the old bulge population of the Galaxy.} \\label{image} \\end{figure} In comparison with the Solar neighborhood, the analysis of the origin of the X-ray emission from the \\1p5 has its own complications: \\begin{itemize} \\item First of all, we need to know the stellar density enclosed by the volume studied by {\\em Chandra}. \\item We need to understand the importance of the absorption that results from cold interstellar matter along the line of sight. \\end{itemize} To study these issues, we used observations of the \\1p5 in the optical band. In addition to the deep narrow-field images obtained with {\\em HST}, large mosaics ($\\sim 25^\\prime\\times 25^\\prime$) of the field were obtained with the ground-based 1.5m Russian-Turkish telescope RTT150 (see Fig.~\\ref{image}) and the 6.5-m Magellan/Baade telescope (see Fig.\\ref{comparison_magellan}). In this paper we concentrate on the properties of the red clump giants (Sect.~\\ref{sec_rcg}) in the \\1p5 as derived from the {\\em HST} images (Sect.~\\ref{sec_rcghst}) to infer properties of the interstellar extinction in this direction (Sect.~\\ref{sec_ext}) and constrain the stellar density distribution along the line of sight (Sects.~\\ref{sec_model} and \\ref{sec_distribution}). We also used results of the measurements with {\\em Spitzer}/IRAC to compare the observed infrared emission towards the \\1p5 with that predicted by the adopted stellar density model (Sect.~\\ref{sec_spitzer}). Our conclusions are summarized in Sect.~\\ref{sec_conclusion}. Other properties of the stellar population like age and metalicity will be explored in separate works. The ultimate goal, also to be addressed in a follow-up paper, is to derive the specific X-ray emission per unit mass and use this to study the unresolved X-ray emission from other galaxies. ", "conclusions": "\\label{sec_conclusion} The study of populations of Galactic X-ray sources requires knowledge of the properties of the underlying population of normal stars, because only then one can estimate the statistical properties of X-ray sources and their systematic dependencies on the age, metallicity, density etc. of the stellar population. One of the largest projects performed recently by the {\\em Chandra X-ray Observatory} is the ultra-deep observation of a region close to the Galactic Center - the \\1p5. This observation provides us with a wealth of data on the X-ray population in this area. To connect these properties to the properties of normal stars we must study this area in the optical and near-infrared bands. We analyzed {\\em HST}/ACS images of the central 6.6\\arcmin $\\times$ 6.6\\arcmin\\, field of the \\1p5 in the F435W and F625W bands and found the following: \\begin{itemize} \\item Interstellar extinction in the direction of the \\1p5 is $\\langle A_{\\rm F625W}\\rangle \\sim 4$ and variable on an angular scale of $\\sim$1\\arcmin. Variations of the extinction in this area on larger scales were previously noticed by \\cite{dutra03} ($4^\\prime$ scales), \\cite{revnivtsev09b} ($1.8^\\prime$ scales), \\cite{vandenberg09} ($1.6^\\prime\\times3.2^\\prime$ scales). \\item The color-magnitude diagram corrected for extinction, as determined in angular bins of $0.96^\\prime\\times0.96^\\prime$, suggests that there is still residual extinction on smaller scales at the level of $A_{\\rm F625W}\\sim 0.2$ mag or less. This would be difficult to measure with RCGs due to their decreasing numbers. \\item We determined the extinction law (in particular, the ratio of the total-to-selective extinction $\\Delta \\langle m_{\\rm F625W} \\rangle /\\Delta \\langle m_{\\rm F435W}-m_{\\rm F625W}\\rangle$) in the direction of the \\1p5 -- $1.3\\pm0.1$ -- and found that it is significantly different from the canonical one $\\sim1.9$. The discrepancy shows a similar trend as that determined by various authors using OGLE data of the bulge \\cite[e.g.][]{popowski00,udalski03,sumi04}. This suggests that we might expect deviations from the standard correspondence between the photoabsorption column $n_HL$, measurable in the X-ray band, and extinction in the optical/NIR bands. This likely differing X-ray (vs. IR) absorption in the bulge region will affect derived intrinsic X-ray spectra and thus (in some cases) the inferred source types (e.g., with less X-ray absorption for a given color excess, soft spectral components will be decreased, etc.). \\item We compared a model of the stellar distribution along the line of sight towards the \\1p5 with the observed distribution of red clump giants in the color-magnitude diagram. We show that the adopted model adequately describes the data. In the model, the majority of the stellar mass in the direction of the \\1p5 is provided by Galactic bulge stars with only a small contribution from the stellar disk and nuclear stellar disk components ($\\sim5$\\% and $\\sim$9\\%, respectively). Small discrepancies between the model and the data could be the result of variations in the extinction on small angular scales ($\\lesssim$1\\arcmin), differences between the local and bulge luminosity functions of red clump giants, uncertainties in the distribution of the extinction along the line of sight, or an oversimplified model of the stellar mass profile of the bulge. \\item We showed that the adopted stellar model predicts an integrated 3.5 $\\mu$m surface brightness of this part of the Galaxy that is only 13\\% higher than the surface brightness actually measured by {\\em Spitzer}/IRAC in the 3.6 $\\mu$m band. \\end{itemize}" }, "1003/1003.0198_arXiv.txt": { "abstract": "\\baselineskip 11pt Beam asymmetries result in statistically anisotropic cosmic microwave background (CMB) maps. Typically, they are studied for their effects on the CMB power spectrum, however they more closely mimic anisotropic effects such as gravitational lensing and primordial power asymmetry. We discuss tools for studying the effects of beam asymmetry on general quadratic estimators of anisotropy, analytically for full-sky observations as well as in the analysis of realistic data. We demonstrate this methodology in application to a recently detected $9\\sigma$ quadrupolar modulation effect in the WMAP data, showing that beams provide a complete and sufficient explanation for the anomaly. ", "introduction": "\\label{sec:model} In a realistic CMB observation, the effective sky signal at each point in the time-ordered data (TOD) is a convolution of the true sky signal with the experimental beam, oriented according to the scan strategy. Schematically, we have \\be T_i = \\int_{S^2} d \\Omega\\, r_{i}(\\Omega) \\Theta(\\Omega) + n_i \\label{eqn:obs_real} \\ee where $T_i$ is the temperature for time-step $i$ in the TOD, $\\Theta(\\Omega)$ is the underlying CMB signal, $r_{i}(\\Omega)$ is the beam response and $n_i$ is the instrumental noise. For the purposes of compact notation we will abbreviate $\\int_{S^2} d \\Omega$ as $\\int$ for the remainder of this paper. The integral in Eq.~\\ref{eqn:obs_real} can be performed by brute force in real-space using interpolation on pixelized maps of the beam and sky \\cite{Wehus:2009zh}. For this approach to be computationally feasible, the beam must be assumed zero outside some small patch where its response is peaked, and so it is difficult to study sidelobe effects with this approach, although it can be quite fast. In this work, we will find it more useful to work in harmonic space, where the effects of beams are easier to study analytically. We begin by writing the beam response as a harmonic sum. If we center the beam at the north pole, with some fiducial beam axis aligned along the $+x$-axis (the $\\phi=0$ meridian), and expand it in spherical harmonics $b_{\\ell m}$, then the beam at location $\\Omega$ for the $i^{\\rm th}$ TOD observation is given by (e.g. \\cite{Souradeep:2001ds}) \\be r_{i}(\\Omega) = \\sum_{s=-s_{\\rm max}}^{s_{\\rm max}} \\sum_{\\ell=|s|}^{\\ell_{\\rm max}} \\sum_{m=-\\ell}^{\\ell} D^{\\ell}_{m s}(\\phi_i, \\theta_i, \\alpha_i) b_{\\ell s} \\, \\yslm{0}{\\ell}{m}(\\Omega). \\label{eqn:beam} \\ee For the purposes of compact notation we will drop the summation limits in what follows. The limits themselves will be discussed later. The action of the Wigner-$D$ matrix can be visualized as follows: imagine fixing the coordinate system in space and performing right-handed rotations of the beam image about the $z$-axis by an angle $\\alpha_i$, then about the $y$-axis by an angle $\\theta_i$, and finally about the $z$-axis by an angle $\\phi_i$. The first rotation gives the beam its orientation: $\\alpha_i$ is the angle of the fiducial beam axis, measured from the southern side of the meridian which passes through the pixel location $(\\theta_i, \\phi_i)$ assigned to the observation. We use $\\yslm{s}{\\ell}{m}$ to denote a spin-weighted spherical harmonic, of which the standard spherical harmonics are a special case with $s=0$. Unless otherwise noted, the harmonics should be taken as functions of $\\Omega$. We can then rewrite Eq.~\\eqref{eqn:obs_real} as \\ba T_i &=& \\sum_{\\ell m s} D^{\\ell}_{-m s}(\\phi_i, \\theta_i, \\alpha_i) b_{\\ell s} (-1)^{m} \\Theta_{\\ell m} + n_i \\nn \\\\ &=& \\sum_{\\ell m s} e^{-i s \\alpha_i} B_{\\ell s} \\Theta_{\\ell m} \\, \\yslm{s}{\\ell}{m}(\\theta_i, \\phi_i) + n_i. \\label{eqn:obs} \\ea In the second step we have used the close relationship between Wigner-$D$ functions and the spin-weighted spherical harmonics \\cite{VarQuant}, and introduced the beam transfer function $B_{l s}$ given by \\be B_{\\ell s} = \\sqrt{ \\frac{4\\pi}{2l+1} } b_{ls}. \\ee For a beam which is normalized to have unit response to a monopole, $B_{00} = 1$. We shall refer to the $s\\!=\\!0$ coefficients of the beam as the symmetric part, as they represent the component of the beam which depends only on the radial distance from its center. The $s\\!\\ne\\!0$ coefficients encapsulate beam asymmetry. On scales much smaller than the beam size, $B_{ls}$ becomes very small, which effectively band-limits $B_{ls} \\Theta_{lm}$. This scale determines $\\ell_{\\rm max}$. The evaluation of Eq.~\\eqref{eqn:obs} may then be performed in $\\clo(s_{\\rm max} l_{\\rm max}^2)$. To perform this convolution for each timestep in the TOD is in general prohibitively expensive, and some approximations must be made. There are several possible approaches: $\\mathbf{(1)}$ The convolution can be computed over a grid covering the full rotation group with fast-Fourier-transforms for the Euler angles $\\phi$ and $\\alpha$ and, optionally, for $\\theta$~\\cite{Wandelt:2001gp}. The TOD can be obtained by interpolation off this grid. $\\mathbf{(2)}$ % If the beam may be represented using a small number of symmetric basis functions, each of these may be rapidly convolved in harmonic space and then sampled based on the location and orientation of these basis functions for each sample in the TOD \\cite{Tristram:2003ee}. The difficulty here is the ability to represent the beam as a sum of symmetric functions. Note that in the limit that the beam is represented as a sum of delta functions, this approach is conceptually the same as real-space integration. We note the above approaches for completeness. In this work, we will use a popular \\cite{Smith:2007rg,Hirata:2004rp,Hirata:2008cb,Mitra:2004nx} map-based approach based on the assumption that the TOD noise is uncorrelated on the timescales which separate pixel visits. In this case, it is a good approximation to the mapmaking process (in the absence of beam deconvolution \\cite{Armitage:2004pk}) to take \\be \\tilde{\\Theta}(\\Omega_p) + n(\\Omega_p) = \\sum_{ i \\in p} T_i / H_p, \\label{eqn:beam_asym_map} \\ee where $\\tilde{\\Theta}$ is an effectively observed sky, $n(\\Omega_p)$ is a noise map, and $H_p$ is the number of elements in the sum, which is taken over all hits assigned to pixel $p$, with center at $\\Omega_p$. This approach can also be used for differencing experiments, which suppress correlations between TOD samples by mapmaking from the difference between two nearly-identical detectors, to remove common mode fluctuations. In this case, one can use an effective beam which is a hit-weighted sum of the two beams which are differenced \\cite{Hinshaw:2006ia}. In conjunction with Eq.~\\eqref{eqn:obs}, the sum of Eq. \\eqref{eqn:beam_asym_map} can be seen to effect a Fourier transform of the distribution of orientation angles, with an effective observed sky given by \\cite{Smith:2007rg} \\be \\tilde{\\Theta}(\\Omega_p) = \\sum_{s} w(\\Omega_p,-s) \\left[ \\sum_{\\ell m} B_{\\ell s} \\Theta_{\\ell m} \\, \\yslm{s}{\\ell}{m}(\\Omega_p) \\right], \\label{eq:obs_sky} \\ee where the details of the scan strategy are contained in the spin $-s$ field \\be w(\\Omega_p,-s) = \\sum_{i \\in p} e^{-i s \\alpha_i} / H_p. \\label{eq:spin_weight} \\ee Since ${}_s Y_{lm}$ involves $s$ (spin-weighted) derivatives of the $Y_{lm}$, each term in the $s$ sum is the real-space product of the scan strategy and beam-filtered derivatives of the CMB. For a beam which is approximately azimuthally symmetric, or a scan strategy which broadly distributes the orientation angles, $w(\\Omega_p,-s) B_{\\ell s}$ falls off sharply with $s$, and it follows that calculation of only the lowest $s$ terms suffice to give a good approximation to the beam-convolved map. This determines an effective $s_{\\rm max}$ which can be much less than that naively required to describe accurately the beam in Eq.~\\eqref{eqn:beam}. Given a scan strategy, Eq.~(\\ref{eq:obs_sky}) provides an $\\clo(s_{\\rm max} l_{\\rm max}^3)$ method to compute effectively the sky observed by an experiment with an asymmetric beam $B_{ls}$ and given scan strategy $w(\\Omega_p,s)$. This approximation is useful not only for its speed, but also to gain an intuitive analytical understanding of beam effects, which we proceed to discuss in the following sections. ", "conclusions": "\\label{sec:conclusions} Beam asymmetry effects can be more important for anisotropy estimators than for power spectrum analysis, because their effects generally enter at lower order in the asymmetry. Beam asymmetries fit nicely into the larger formalism of quadratic anisotropy estimators. They result in a mean-field bias which directly traces the scan strategy $\\slm{w}{S}{L}{M}$, and can be calculated analytically on the full sky or determined from Monte-Carlo simulations on the cut sky. Beam effects appear to provide a sufficient explanation for the $9\\sigma$ detection of an apparent quadrupolar modulation of the primordial power spectrum in the WMAP data. We note that the WMAP team already incorporates the effects of beam asymmetry into their power spectrum analysis (where it is a much smaller effect in any case), and so the resolution of this anomaly should not have any effect on their cosmological parameter constraints. All of this work will apply directly to the Planck experiment, which has a less symmetrizing scan strategy than WMAP. Planck's increased sensitivity also opens up the field for the precision analysis of interesting astrophysical secondaries such as the anisotropic signal from gravitational lensing. The tools and techniques which we have discussed here may also be extended straightforwardly for use with polarization data." }, "1003/1003.5774_arXiv.txt": { "abstract": "{The low mass protostar IRAS16293-2422 is a prototype Class 0 source with respect to the studies of the chemical structure during the initial phases of life of Solar type stars.} {In order to derive an accurate chemical structure, % a precise determination of the source physical structure is required. The scope of the present work is the derivation of the structure of IRAS16293-2422.} {We have re-analyzed all available continuum data (single dish and interferometric, from millimeter to MIR) to derive accurate density and dust temperature profiles. Using ISO observations of water, we have also reconstructed the gas temperature profile.} {Our analysis shows that the envelope surrounding IRAS16293-2422 is well described by the Shu ``inside-out'' collapsing envelope model or a single power-law density profile with index equal to 1.8. In contrast to some previous studies, our analysis does not show evidence of a large ($\\geq 800$ AU in diameter) cavity.} {Although IRAS16293-2422 is a multiple system composed by two or three objects, our reconstruction will be useful to derive the chemical structure of the large cold envelope surrounding these objects and the warm component, treated here as a single source, from single-dish observations of molecular emission. } ", "introduction": "\\label{sec:introduction} Understanding how our Sun and Solar System formed is arguably one of the major goals of modern astrophysics. Many different approaches contribute to our understanding of the past history of the Solar System. Analyzing the relics of the ancient eons, comets and meteorites, is one. Studying present day objects similar to what the Sun progenitor is another. Here we pursuit the latter approach and analyze in detail the case of one of the best studied solar type protostars, IRAS16293-2422 (hereinafter IRAS16293). IRAS16293 is a Class 0 protostar in the $\\rho$ Ophiuchus complex at 120 pc from the Sun \\citep{Loi08} and has played the role of a prototypical solar-type protostar for astrochemical studies, just as Orion~KL has done for high-mass protostars. This is beause of its proximity and the resulting line strength of molecular emission \\citep[e.g.][to mention just a few representative works from the previous decades]{Wal86,Mun92,Bla94,Van95}. It is in this source that the phenomenon of the ``super-deuteration''\\footnote{The super-deuteration refers to the exceptionally high abundance ratio of D-bearing molecules with respect to their H-bearing isotopologues found in low mass protostars, with observed D-molecule/H-molecule ratios reaching the unity \\citep[see e.g. the review in][]{Cec07}.} has been first discovered, with the detection of surprising abundant multiply deuterated molecules: formaldehyde, hydrogen sulfide, methanol and water \\citep{Cec98,Vas03,Par03,But07}. It is also in this source that the first hot corino \\citep{Bot04,Cec07} has been discovered, with the detection of several abundant complex organic molecules in the region where the dust grain mantles sublimate \\citep{Cec00a,Caz03,Bot04}. Not surprisingly, therefore, IRAS16293 has been the target of several studies to reconstruct its physical structure, namely the dust and gas density and temperature profiles \\citep{Cec00b,Sch02,Sch04,Jor05}, the mandatory first step to correctly evaluate the abundance of molecular species across the envelope. \\citet{Cec00b} used water and oxygen lines observations obtained with the Infrared Space Observatory (ISO) to derive the gas and dust density and temperature profile. Conversely, \\citet{Sch02} used the dust continuum observations to derive the structure of the envelope. Furthermore, while Ceccarelli et al. assumed the semi-analytical solution by Shu \\& co-workers \\citep{Shu77,Ada86} to fit the observational data, \\citet{Sch02} assumed a single power law for the density distribution, and a posteriori verified that the Shu's solution also reproduced the observational data. The two methods lead to similar general", "conclusions": "\\label{sec:conclusions} The new analysis of the single dish and interferometric continuum observations of the envelope of IRAS16293 confirms that an envelope of about 2 M$_\\odot$ surrounds the proto-binary system of IRAS16293. The envelope can be described with a Shu-like density distribution, corresponding to the gas collapsing towards a 2 M$_\\odot$ central star. The luminosity of IRAS16293 has been re-evaluated to be 22 L$_\\odot$ for a distance of 120 pc. Both the single dish and interferometric data can be reproduced by an envelope with an inner radius between 20 and 30 AU, equivalent to about 0.4$''$, and smaller than the radius at which the dust temperature reaches 100 K (the ice mantle sublimation temperature, namely 75 to 85 AU). We found that our new analysis can reproduce the full SED, including the Spitzer MIR data, without the necessity of a central cavity of 800 AU radius \\citep{Jor05}. The difference between our models and the previous ones (based on the \\citealt{Sch02} initial model) is the larger bolometric luminosity (kept as a free parameter in our models) and the lower optical thickness of the envelope ( $\\tau_{100\\mu m}$=2, namely twice smaller than in the Sch{\\\"o}ier's models). These differences contribute to make that the predicted MIR radiation flux agrees with the observed one. Finally, the interferometric data, being dominated by the two components of the binary system, do not provide significant constraints on the envelope structure, with one exception. They exclude the case of an envelope with a single power-law density profile centered on the mid-way point between the two sources. Based on this assumption, \\citet{Sch04} suggested the presence of a cavity 800 AU in diameter. Since no data constrain where the center of the envelope is located, we favor the solution with the envelope centered on one of the two sources. In addition, the Shu-like model fits slightly better the PDBI data and is thus our preferred solution. As already noted by \\citet{Cec00b}, the ISO data do not allow to constrain the inner H$_2$O abundance because the detected lines are optically thick and cover a relatively low range of upper-level energies ($\\leq 300$ cm$^{-1}$). Also the outer-envelope water abundance is relatively poorly constrained. In addition, the relatively low spectral resolution of ISO does not allow to determine whether some lines are contaminated by the emission from the outflow. The future observations with the HIFI spectrometer aboard the Herschel Space Observatory, launched in May 2009, will certainly constrain better the water abundance profile across the IRAS16293 envelope, helping to understand the distribution of water in protostars similar to the Sun progenitor. On the one hand, the water abundance in the outer envelope (0.7--20$\\times$10$^{-7}$) derived from the ISO observations here is consistent with some previous estimates of water abundance in cold gas \\citep[e.g.][]{Cer97} but only marginally consistent with other low estimates \\citep[e.g.][]{Sne00}, so the new Herschel/HIFI observations, with their largely improved spatial and spectral resolution, will be crucial in settling the question. On the other hand, the inner envelope abundance ($<$2$\\times$10$^{-6}$) is lower than expected if all ice in the mantles sublimates \\citep[see][]{Cec00b}. Also in this case, the Herschel/HIFI observations will help to understand this point. Finally, as stated in the Introduction, the major scope of the present work is to provide as accurate as possible estimates of the dust and gas temperature profiles of the cold envelope of IRAS16293 and its warm inner component, also known as the hot corino, to interpret the data observed in two large projects, TIMASSS and CHESS (see Introduction). We are aware that the proposed description has the intrinsic and clear limit of not taking into account the multiple nature of the IRAS16293 system. But it has the merit in allowing the interpretation of the single-dish observations of the upcoming projects, within this limitation. A more detailed analysis will only be possible once the relevant molecular emission is observed with interferometers, resolving the two components of the system. In absence of that, the analysis based on a single warm component and cold envelope is the only vialable and allows a first understanding of the chemical composition of a system which eventually will form a star and planetary system like our own." }, "1003/1003.4544_arXiv.txt": { "abstract": "{Strong lensing is one of the most direct probes of the mass distribution in the inner regions of galaxy clusters. It can be used to constrain the density profiles and to measure the mass of the lenses. Moreover, the abundance of strong lensing events can be used to constrain the structure formation and the cosmological parameters through the so-called \"arc-statistics\" approach. However, several issues related to the usage of strong lensing clusters in cosmological applications are still controversial, leading to the suspect that several biases may affect this very peculiar class of objects.}{With this study we aim at better understanding the properties of galaxy clusters which can potentially act as strong lenses.}{We do so by investigating the properties of a large sample of galaxy clusters extracted from the N-body/hydrodynamical simulation {\\sc MareNostrum Universe}. We perform ray-tracing simulations with each of them identifying those objects which are capable to produce strong lensing effects. We explore the correlation between the cross section for lensing and many properties of clusters, like the mass, the three-dimensional and projected shapes, their concentrations, the X-ray luminosity and the dynamical activity.}{We quantify the minimal cluster mass required for producing both multiple images and large distortions. While we do not measure a significant excess of triaxiality in strong lensing clusters, we find that the probability of strong alignments between the major axes of the lenses and the line of sight is a growing function of the lensing cross section. In projection, the strong lenses appear rounder within $R_{200}$, but we find that their cores tend to be more elliptical as the lensing cross section increases. As a result of the orientation bias, we also find that the cluster concentrations estimated from the projected density profiles tend to be biased high. The X-ray luminosity of strong lensing clusters tend to be higher than that of normal lenses of similar mass and redshift. This is particular significant for the least massive lenses. Finally, we find that the strongest lenses generally exhibit an excess of kinetic energy within the virial radius, thus indicating that they are more dynamically active than usual clusters.} {We conclude that strong lensing clusters are a very peculiar class of objects, affected by many selection biases which need to be properly modeled when using them to study the inner structure of galaxy clusters or to constrain the cosmological parameters.} ", "introduction": "Gravitational lensing is one of the most powerful tools for studying the formation of cosmic structures in the universe. The light from distant sources, traveling in space and time, is deflected by the matter along its path before being collected by the observers. Thus, we measure an integrated effect which contains a wealth of information about the cosmic structures at different epochs. Depending on the impact parameter on the intervening matter and on the mass of the deflectors encountered by the light along its path, gravitational lensing manifests itself in the weak and in the strong regimes. In the weak lensing regime the shapes of distant galaxies, which happen to be at large angular distances from the largest mass concentrations on the sky, are slightly changed, such that this effect can be revealed only though statistical measurements. Nevertheless, these tiny distortions can be used for tracing the large scale structure of the universe both in two and in three dimensions \\citep[see e.g.][for some recent results]{2008A&A...479....9F,2007MNRAS.381..702B,2007ApJS..172..239M}, from which important cosmological constraints can be derived \\citep{BA01.1}. This is a field of research which will have extraordinary improvements in the next decades, thanks to some upcoming missions \\citep{2006AAS...209.8610W,2007AAS...210.5102K,2006AAS...209.9809J,2008arXiv0807.4036R}. Weak lensing allows to reconstruct the mass distribution up to the outskirts of galaxy clusters \\citep[see e.g.][]{2006ApJ...653..954D,2006A&A...451..395C,2007MNRAS.379..317H}. Strong lensing is an highly non-linear and relatively rare effect which is observable in the central regions of galaxies and clusters. In this regime, the background sources can be multiply imaged and/or highly distorted to form very elongated images, the so called \"gravitational arcs\". They are powerful cosmological probes for many reasons. First, such events can be used to investigate the inner regions of the lenses. Thus, they can be used to test the predictions of the Cold-Dark-Matter paradigm on the inner structure of dark matter halos \\citep{ME01.1,2004ApJ...604L...5M,2004ApJ...604...88S,BA04.1,2007MNRAS.381..171M,2008A&A...489...23L}. Second, they can be used to recover the mass distribution in the centre of the lenses, providing complementary informations to those obtained from weak lensing \\citep{2005A&A...437...49B,2005MNRAS.362.1247D,2006A&A...458..349C,2007ApJ...668..643L,2009A&A...500..681M}. Lensing masses can then be used for measuring the cluster mass function. Third, the position and the distortions of the strongly lensed images as a function of the source redshift reflect the geometry of the universe \\citep{2004A&A...417L..33S,2005MNRAS.362.1301M}. Finally, the probability of observing strong lensing events is deeply connected to the abundance, the mass and the formation epoch (through the concentration) of the lenses. This makes statistical lensing a potentially powerful tool to study the structure formation \\citep{BA98.2,LI05.1,ME05.1}. In this paper, we focus on the properties of the most massive and therefore most efficient strong lenses in the universe: the galaxy clusters. In the framework of the hierarchical scenario of structure formation, these are the youngest bound systems in the sky. About $85\\%$ of their mass is believed to be in the form of cold-dark-matter \\citep{2007ApJ...664..117G}. The remaining $15\\%$ is made of a diffuse gas component, the Intra-Cluster-Medium, and of other baryons in the form of stars, the vast majority of which is inside the cluster galaxies. Being relatively young structures, the interaction between the baryons and the dark matter is less strong than in older systems like galaxies. For this reason, clusters are important laboratories for studying the properties of the dark matter \\citep{2004ApJ...606..819M}. However, there are several issues that we need to take into account when studying these systems. In particular, clusters where gravitational arcs are observed are a limited fraction of the total number \\citep{2005ApJ...627...32S,LU99.1,GL03.1,ZA03.1,2008AJ....135..664H} and therefore a particular class of objects. Broadly speaking they are the most massive clusters, but there are several other properties that boost the cluster ability to produce strong lensing events. For example, we know that substructures, asymmetries and projected ellipticity of the lenses are all contributing to the strong lensing cross section of clusters \\citep{ME03.1,ME07.1}. Both observations and simulations agree that strong lenses have high concentrations \\citep{HE07.1,2007A&A...473..715F,2003ApJ...598..804K,2003A&A...403...11G,2008ApJ...685L...9B}. Triaxiality is also relevant, because clusters seen along their major axis are more efficient lenses \\citep{OG03.1}. Although cluster galaxies statistically do not change the distributions of the arc properties \\citep{ME00.1,FL00.1,2008MNRAS.386.1845H}, cD galaxies sitting at the bottom of the cluster potential well increase the ability for producing long and thin arcs by $\\sim 30-50\\%$ \\citep{ME03.2}. The gas physics, in particular cooling, could also affect the strong lensing properties of clusters \\citep{PU05.1,2008ApJ...676..753W,2010arXiv1001.2281M}. Finally, \\cite{TO04.1} showed that the cluster ability to produce gravitational arcs can also be enhanced by the dynamical activity in the lens. By studying with high time resolution how the lensing cross section changes during an edge-on collision between the main cluster clump and a substructure, these authors found that the strong lensing efficiency is boosted by a factor of 10 during the merging phase. Later, \\cite{FE06.1}, using semi-analytic methods, showed that the arc optical depth produced by clusters with moderated and high redshifts is more than doubled by mergers. Although an extensive work has been done in the past decade, a better characterization of the strong lens cluster population is mandatory. Given the complexity of clusters and the great importance that many of their properties have for strong lensing, the only reliable way to do that is through the ray-tracing analysis of a large number of simulated clusters. A first important work in this framework was done by \\cite{HE07.1}, who analyzed a sample of 878 clusters from an N-Body pure dark-matter cosmological simulation. Important properties like concentrations, axis ratios and substructures were discussed. In this work, we aim at extending the analysis of \\cite{HE07.1} in three ways. First, we include a much larger number of clusters (now 49366 systems), taken from a larger cosmological volume ($500\\,h^{-1}$Mpc vs $320\\,h^{-1}$Mpc). Second, the clusters used here are obtained from an N-body-hydrodynamical simulation where the evolution of the gas component is also considered. Thus, we can correlate the lensing properties of clusters with some important X-ray observables. Third, we study in detail the possible correlations between cluster dynamics and strong lensing efficiency, which was made only through analytical models so far. The plan of the paper is as follows. In Sect. 2 we summarize the main characteristics of the cosmological simulation {\\sc MareNostrum Universe}. In Sect. 3 we discuss the simulation methods and we define several lensing quantities useful for the following analysis. In Sect. 4 we discuss the correlation between lens masses and strong lensing ability. Sect. 5 is dedicated to the statistical analysis of the shapes and orientations of strong lensing clusters. We discuss the biases in the concentrations in Sect. 6. In Sect. 7 we focus on the X-ray properties of strong lensing clusters. Finally, in Sect. 8 we correlate the strong lensing efficiency with the dynamical state of the lenses. We summarize the main results and the conclusions of this study in Sect. 9. ", "conclusions": "In this paper, we have investigated the properties of $\\sim 50000$ strong lensing clusters at different redshifts in the {\\sc MareNostrum Universe} cosmological simulation. Projecting each of these clusters along three orthogonal lines of sight, we have considered almost $150000$ lens realizations in total. With such a big number of objects, we can statistically characterize the population of strong lensing clusters much better than it was done in the past. Moreover, the {\\sc MareNostrum Universe} includes gas, thus it allows to investigate the correlation between strong lensing and X-ray observables. We have classified the strong lenses into two categories, namely clusters which are critical for sources at $z_{\\rm s}=2$ and clusters which can induce large distortions in the images of these sources, i.e. form giant arcs. We have explored several structural properties of the strong lensing clusters, like the masse, the shapes, the concentrations, the X-ray luminosity, and the dynamical activity. Our main results can be summarized as follows: \\begin{itemize} \\item strong lensing clusters are typically massive objects. Their masses vary over two orders of magnitude. The minimal mass for strong lensing depends on both the redshift of the lenses and of the sources. For sources at $z_{\\rm s}=2$, we find that clusters can develop critical lines down to masses of $\\sim 10^{13}\\;h^{-1}\\;M_\\odot$. Requiring that clusters are also able to produce giant arcs, increases the mass limit by almost one order of magnitude. The lensing cross section further depends on the mass, so that we can estimate that the minimal mass required for a cluster, in order that the expected number of giant arcs in a deep HST observation is $\\sim 1$, is $\\sim 2 \\times 10^{14}\\;h^{-1}\\;M_\\odot$; \\item the three-dimensional shape of strong lensing clusters does not seem to be significantly different from that of the general cluster population. However, strong lensing clusters tend to have their major axes oriented along the line of sight. This ``orientation bias\" is larger for clusters which produce giant arcs than for clusters which only possess critical lines, and it becomes larger by increasing the lensing cross section; \\item due to the orientation bias and to the fact that their halos are generally well described by prolate triaxial models, strong lensing clusters tend to appear rounder when projected on the sky. However, zooming over their central regions, we noticed that their projected mass maps are described by rather elongated distributions, evidencing the presence of substructures projected near cluster cores; \\item the concentrations measured by fitting the density profiles of strong lensing clusters stacked in mass and redshift bins do not differ significantly from the concentrations of the general cluster population. Nevertheless, due to the orientation bias, the concentrations of the same objects inferred from the analysis of the projected density profiles are generally larger than in 3D. For clusters with large lensing cross sections for giant arcs, the 2D-concentrations can be larger by more than a factor of two. These results may provide a viable explanation of the large concentrations reported for some of the strongest lenses observed so far; \\item the X-ray luminosity-mass relation of strong lensing clusters is likely to differ from that of the general cluster population, especially at the lowest masses. We found that at a fixed mass, strong lenses with increasingly larger cross sections for lensing have higher X-ray luminosities, indicating that some process occurring in these objects enhances both the X-ray luminosity and the strong lensing cross section; \\item the strong lensing efficiency is certainly correlated with the dynamical activity in clusters. We found that clusters with large lensing cross sections are characterized by a systematic departure from virial equilibrium. A similar departure from virial equilibrium is found for clusters which are in the process of accreting substructures, which accidentally transit across their cores perpendicularly to the line of sight. \\end{itemize} In conclusion, our results show that strong lensing clusters are likely to be a very peculiar class of objects, characterized by several selection biases, which need to be properly taken into account in many applications. For example, due to the orientation bias, it is very likely that the 3D-masses inferred from strong lensing models of observed clusters are biased high. In arc statistics studies, the statistical modeling of the strong lensing cluster population need to include mergers, triaxiality, and asymmetries in the projected mass distributions, as also suggested by previous studies. Due to the intrinsic difficulties at modeling analytically all of these effects, numerical simulations again seem to be the only viable way to describe strong lensing clusters. In this sense, large statistical samples of numerically simulated lenses, as that extracted from the {\\sc MareNostrum Universe}, are fundamental tools for interpreting the current strong lensing observations. A major effort is now necessary to clarify the existing inconsistencies between the properties of simulated and observed galaxy clusters, especially in the central regions." }, "1003/1003.1227_arXiv.txt": { "abstract": "Several binary systems consisting of a massive star and a compact object have been detected above 100~GeV in the Galaxy. In most of these sources, gamma-rays show a modulation associated to the orbital motion, which means that the emitter should not be too far from the bright primary star. This implies that gamma-ray absorption will be non negligible, and large amounts of secondary electron-positron pairs will be created in the stellar surroundings. In this work, we show that the radio emission from these pairs should be accounted for when interpreting the radio spectrum, variability, and morphology found in gamma-ray binaries. Relevant features of the secondary radio emission are the relatively hard spectrum, the orbital motion of the radio peak center, and the extended radio structure following a spiral-like trajectory. The impact of the stellar wind free-free absorption should not be neglected. ", "introduction": "Five galactic sources of very high-energy (VHE) gamma-rays have been associated so far to binary systems consisting of a compact object and a massive bright star: PSR~B1259$-$63\\cite{aha05a}; LS~5039\\cite{aha05b}; LS~I~+61~303\\cite{albert06}; Cygnus~X-1\\cite{albert07}; HESS~J0632$+$057\\cite{hinton09}. Three of them show a modulation in gamma-rays associated to the orbital motion (PSR~B1259$-$63\\cite{aha05a}; LS~5039\\cite{aha06}; LS~I~+61~303\\cite{albert09}), which implies that the gamma-ray emitter cannot be too far from the star. Since the spectrum of the star radiation peaks in the ultraviolet, efficient absorption of VHE photons will take place through interaction with stellar photons, leading to the creation in the stellar surroundings of secondary electron-positron pairs with energies $\\ge 10$~GeV, which is about a half of pair creation gamma-ray energy threshold. Once created, secondary pairs cool down trapped in the stellar wind due to slow diffusion and radiate mainly via synchrotron and inverse Compton (IC) emission. Depending on the medium magnetic field, either synchrotron or IC will dominate (for the latter, see Ref.~\\refcite{pelliza09}). The gamma-ray binaries found to date emit also non-thermal radio and X-ray radiation (see Ref.~\\refcite{bosch09} and references therein). In principle, despite the primary gamma-ray emitter could be also producing radiation at lower energies, the component generated by the secondary pairs in the stellar wind could even dominate the whole non-thermal output of the source\\cite{bosch08}. ", "conclusions": "" }, "1003/1003.4128_arXiv.txt": { "abstract": "{Stellar flares affect all atmospheric layers from the photosphere over chromosphere and transition region up into the corona. Simultaneous observations in different spectral bands allow to obtain a comprehensive picture of the environmental conditions and the physical processes going on during different phases of the flare.} {We investigate the properties of the coronal plasma during a giant flare on the active M dwarf CN~Leo observed simultaneously with the UVES spectrograph at the VLT and \\emph{XMM-Newton}.} {From the X-ray data, we analyze the temporal evolution of the coronal temperature and emission measure, and investigate variations in electron density and coronal abundances during the flare. Optical \\ion{Fe}{xiii} line emission traces the cooler quiescent corona.} {Although of rather short duration (exponential decay time $\\tau_{LC} < 5$ minutes), the X-ray flux at flare peak exceeds the quiescent level by a factor of $\\approx$100. The electron density averaged over the whole flare is greater than $5 \\cdot 10^{11}$~cm$^{-3}$. The flare plasma shows an enhancement of iron by a factor of $\\approx$2 during the rise and peak phase of the flare. We derive a size of $<9000$~km for the flaring structure from the evolution of the the emitting plasma during flare rise, peak, and decay. } {The characteristics of the flare plasma suggest that the flare originates from a compact arcade instead of a single loop. The combined results from X-ray and optical data further confine the plasma properties and the geometry of the flaring structure in different atmospheric layers.} ", "introduction": "Flares are among the most prominent signs of stellar activity. They can be observed over the entire electromagnetic spectrum from X-ray to radio wavelengths, demonstrating that all layers of the stellar atmosphere are affected, with a wide range of plasma temperatures and densities involved. A flare is a highly dynamic event, revealing the complex behavior of the stellar atmosphere with rapidly changing physical conditions, i.\\,e. its interplay with varying magnetic fields in time and its response to the sudden release of large amounts of energy. In the commonly accepted picture of a stellar flare \\citep{Haisch_Flares, Priest_Forbes}, a magnetic reconnection event in the corona drives the acceleration of particles that radiate in nonthermal hard X-rays and radio gyrosynchrotron emission as they spiral down along the magnetic field lines. Lower atmospheric layers are strongly heated by the impact of these particles and immediately start cooling by radiation in the optical and UV continuum as well as in chromospheric and transition region emission lines. Chromospheric evaporation then brings ``fresh material'', emitting soft thermal X-ray emission, into the corona. The strong X-ray and UV radiation field finally induces further chromospheric emission \\citep{Hawley_Fisher}. Flares on the Sun and on stars show a wide variety of amplitudes, from the smallest micro/nanoflares to giant flares with luminosity increases by orders of magnitude, and timescales ranging from a few seconds up to several days. Depending on energy budget, decay time and shape of the lightcurve, impulsive and gradual flares can be discerned. On the Sun, the former have been associated with compact emission regions, i.\\,e., single loops, while the latter typically involve a series of eruptions in a whole arcade of loops \\citep[so-called two-ribbon flares, see][]{Pallavicini_flaretypes}. For stellar flares, the loop geometry can usually not be observed directly, but scaling laws based on the hydrostatic case, relating the loop temperature and pressure with the size of the loop, as derived and tested for the quiet Sun \\citep{RTV} can be adapted to stellar flares to give an estimate of the dimensions of the involved coronal structures \\citep{Aschwanden_laws}. Additionally, hydrodynamic loop modeling approaches allow to assess the loop length from the analysis of the decay of the flare in X-rays and to estimate the amount of additional heating \\citep{Serio_flaremodel, Reale_loops}. This approach has recently been extended to an analysis of the rise phase of the flare \\citep{Reale_diagnostics}. Individual flare events can thus be characterized with respect to to the physical properties of the flaring structure from observational quantities also available for stellar events. The greater the spectral coverage and the better the spectral and temporal resolution, the more information can be gained and the better are the constraints on the physical properties of the flare plasma and the geometry of the flaring active region. Simultaneous observations in multiwavelength bands allow to study the flare throughout all atmospheric layers from the photosphere into the corona. Strong flares dominate the overall flux level in all wavelength bands; the X-ray luminosity can increase by more than two orders of magnitude during the strongest events \\citep[e.\\,g.][]{Algol_nature, Algol_Flare, Favata_EV_Lac}, which allows the strongest events even on relatively faint targets to be investigated in great detail. For outstanding flares like the long-lasting giant flare observed on Proxima~Cen with \\emph{XMM-Newton}, the temporal evolution of plasma temperatures, densities and abundances during different flare phases could be observed directly \\citep{Guedel_Proxima_Cen_1, Guedel_Proxima_Cen_2}; detailed and specific hydrodynamic loop modeling revealed a complex flare geometry with two loop populations and differing heating mechanisms \\citep{Reale_Proxima_Cen}. Even for much smaller flares, low- and intermediate-resolution X-ray spectroscopy makes it possible to follow the evolution of plasma temperature and emission measure during flare rise and decay. The temperature of the flare plasma is typically observed to peak before the maximum emission measure is reached, verifying the overall framework of initial heating, evaporation, and conductive and radiative cooling as predicted by hydrodynamic loop models \\citep[e.\\,g.][]{Reale_diagnostics}. High-resolution spectroscopy is needed to directly measure coronal densities during flares and thus to obtain the weights of changes in plasma density or volume. In general, coronal densities deduced from the ratio of the forbidden and intercombination lines in He-like triplets, especially from \\ion{O}{vii} and \\ion{Ne}{ix} at formation temperatures of 2~MK and 3.5~MK, respectively, during smaller flares are consistent with the quiescent value within the errors but tend to be higher \\citep[see e.\\,g.][]{Mitra_Kraev_densities}; unambiguously increased coronal densities have so far only been measured during the giant flare on Proxima Centauri \\citep{Guedel_Proxima_Cen_1}. Another aspect of the evaporation scenario is the observation of abundance changes during flares, leading directly to the question to what extent flares influence the overall coronal abundance pattern. \\citet{Nordon_flares, Nordon_flares2} systematically analyzed abundance trends in a sample of flares observed with \\emph{Chandra} and \\emph{XMM-Newton} and found that the most active stars with a pronounced inverse FIP effect during quiescence compared to solar photospheric values showing a shift to a FIP effect compared to the corresponding quiescent values, and vice versa for the less active stars showing the solar-like FIP effect during quiescence. This confirms the existence of the inverse FIP effect and of strong discontinuities in the abundance patterns of corona and chromosphere/photosphere also for stars where photospheric abundances are difficult to measure, however, the question what causes this bias is left unanswered . We present the observation of a giant flare on the active M~dwarf CN~Leo with \\emph{XMM-Newton} and the UVES spectrograph at the VLT. An extremely short impulsive outburst at the beginning of the flare has already been discussed by \\citet{CN_Leo_onset}. This paper is the concluding third paper of a series analyzing the whole flare event in greater detail. Paper~I \\citep{flare_chromos} reports the observational results on the chromospheric emission as diagnosed by the UVES spectra, while Paper~II \\citep{flare_model} describe model chromospheres and synthetic spectra obtained with the atmospheric code \\emph{PHOENIX}. Here we focus on the temporal evolution of the coronal plasma temperature, emission measure and abundances determined by the X-ray data and their implications on the geometry of the flaring structure, on the physical conditions dominating the flare environment, and on the abundance patterns predominating in different atmospheric layers. ", "conclusions": "We have analyzed simultaneous X-ray and optical data covering a giant flare on the active M dwarf CN~Leo with regard to the plasma properties in the corona and the geometry of the flaring structure. As it could be expected from other events of similar strength observed in X-rays, temperature and density are enhanced by more than an order of magnitude compared to the quiescent corona at the flare peak. The flare plasma, which can be considered to consist mostly of material evaporated from chromosphere or photosphere, shows a different composition. Despite its high amplitude, the flare was only of short duration, indicating a rather compact flaring structure, which is also comfirmed by simple loop modeling. Sustained heating during the decay as well as substructure in the lightcurve and hardness ratio however suggest an arcade-like structure much more complex than a single flaring loop. Broad multiwavelength coverage allowed us to characterize this exceptional event in great detail." }, "1003/1003.3014_arXiv.txt": { "abstract": "We show non-perturbatively that the power spectrum of a self-interacting scalar field in de Sitter space-time is strongly suppressed on large scales. The cut-off scale depends on the strength of the self-coupling, the number of $\\ee$-folds of quasi-de Sitter evolution, and its expansion rate. As a consequence, the two-point correlation function of field fluctuations is free from infra-red divergencies. ", "introduction": " ", "conclusions": "" }, "1003/1003.2632_arXiv.txt": { "abstract": "Identifying the population of young stellar objects (YSOs) in high extinction regions is a prerequisite for studies of star formation. This task is not trivial, as reddened background objects can be indistinguishable from YSOs in near-infrared colour-colour diagrams. Here we combine deep JHK photometry with J- and K-band lightcurves, obtained with UKIRT/WFCAM, to explore the YSO population in the dark cloud IC1396W. We demonstrate that a colour-variability criterion can provide useful constraints on the star forming activity in embedded regions. For IC1396W we find that a near-infrared colour analysis alone vastly overestimates the number of YSOs. In total, the globule probably harbours not more than ten YSOs, among them a system of two young stars embedded in a small ($\\sim 10000$\\,AU) reflection nebula. This translates into a star forming efficiency SFE of $\\sim 1$\\%, which is low compared with nearby more massive star forming regions, but similar to less massive globules. We confirm that IC1396W is likely associated with the IC1396 HII region. One possible explanation for the low SFE is the relatively large distance to the ionizing O-star in the central part of IC1396. Serendipitously, our variability campaign yields two new eclipsing binaries, and eight periodic variables, most of them with the characteristics of contact binaries. ", "introduction": "Our knowledge about the origin of low-mass stars is primarily based on observations of a few ($\\sim 10$) nearby star forming regions. Although many of these regions are quite similar (for example, most of them do not harbour massive stars), there are indications for environmental differences in the outcome of star formation \\citep[e.g.][]{2000prpl.conf..121M,2008ApJ...688..377S,2009ApJ...703..399L}. Moreover, most of the extensively studied clouds are similar in star formation efficiency \\citep{2009ApJS..181..321E}, indicating a comparable evolutionary state \\citep[see][for an exceptional case]{2009ApJ...704..292F} Hence, there is an incentive to extend our observational studies to more distant and diverse regions. The problem is to find efficient means to identify clusters of embedded T Tauri stars. Since extinction is often too high for optical observations, the surveys are best carried out in the near/mid-infrared. The established method is to look for excess emission from circumstellar disks using colour-colour diagrams. At least in the near-infrared, however, reddenened background giants and extragalactic objects can occupy the same regions in colour-colour space as young stellar objects (YSOs), as extinction and disks can redden the colour in a similar way. Also, this method will only probe for classical T Tauri stars with disks (CTTS) and miss the population of diskless weak-line T Tauri stars (WTTS). Thus, in order to select a more complete and less contaminated sample for follow-up spectroscopy, additional criteria have to be used to separate background/foreground objects from YSOs. One good option here is variability: T Tauri stars exhibit characteristic and well-studied photometric variations due to spots, activity, accretion, and disks \\citep[e.g.][]{1994AJ....108.1906H,2001AJ....121.3160C,2002AJ....124.1001C,2008A&A...485..155A}, which can easily be probed with relatively small telescopes. We present here a case study for a combined colour-variability criterion to identify YSOs in a high extinction region. \\begin{figure} \\includegraphics[width=8.5cm,angle=0]{f100.ps} \\caption{K-band image of the cloud IC1396W. This image has been created by co-adding the individual K-band frames from our time series observations. The positions and directions of the three outflows are indicated with dashed lines \\citep{2003A&A...407..207F}. The position of IRAS 21246+5743 is marked with an ellipse. The YSO candidates identified in this paper (see Table \\ref{t2}) are marked with circles; one of them is further north and not seen in the image. The square indicates the position of the small nebula with two embedded YSOs, see Sect. \\ref{comb} and Fig. \\ref{f10}. \\label{f100}} \\end{figure} Our target region is the dark cloud IC1396W\\footnote{$\\alpha = 21^h26^m$ $\\delta = +58^o00^m$, galactic coordinates $b = 98.3^o$, $l=+05.2^o$} in the vicinity, but outside, of the large HII region IC1396 which is ionized by the O6.5V star HD 206267. IC1396 contains several prominent examples of star forming activity, for example IC1396N \\citep{2001A&A...376..271C,2001A&A...376..553N} and IC1396A \\citep{2004ApJS..154..385R}, see \\citet{2005A&A...432..575F} for a summary. If IC1396W is associated with this region, the distance would be $\\sim 750$\\,pc \\citep{1979A&A....75..345M}, but so far no independent distance measurement is available. Given its relatively simple structure and radiation environment, the sample of globules in IC1396 is an ideal testbed to probe the effect of an ionizing star on the surrounding gas and embedded star forming activity. The available data is consistent with a scenario in which the radiation from the O star impacts both the globule masses and the star forming activity in the globules: With progressively larger distances from the ionizing star, the globule masses increase, while the star forming activity drops \\citep{2005A&A...432..575F}. IC1396W contains the red source IRAS 21246+5743 and at least three molecular hydrogen outflows \\citep[Fig. \\ref{f100},][]{2003A&A...407..207F}. Associated with these outflows are the features MHO 2267-2772 in the recently released Catalogue of Molecular Hydrogen Emission-Line Objects \\citep{2010A&A...511A..24D}. Several of the H$_2$ emission features are detected in the optical as Herbig-Haro objects HH 864A-C by \\citet{2005A&A...432..575F}. The same authors estimate the cloud mass to be about 500$\\,M_{\\odot}$ and the size to be $\\sim 2$\\,pc, based on the distance of 750\\,pc, which makes it one of the largest clouds in the IC1396 region. The IRAS source seems to be a young Class-0 object of about 16$\\,L_{\\odot}$ and 33\\,K \\citep{2003MNRAS.346..163F} and also the driving source of the main outflow HH864. The ISO maps at 160 and 200$\\,\\mu m$ give evidence for the presence of two more embedded sources close to the central object ($2.5'$ SW and NE, respectively). Thus, this is a site of active, ongoing star formation.\\footnote{CO(1-0) observations by \\citet{2006NewA...12..111Z} do not detect the outflows nor dense cores associated with the possible driving sources in the cloud. Their beam size of $106\"\\times 70\"$, however, is much larger than typical Class-0 sources, i.e. insufficient spatial resolution can explain this result.} In a shallow near-infrared survey, we found that IC1396W harbours a population of reddened objects, most of them clustered in a region coinciding with a clump of gas \\citep{2003A&A...407..207F}. We present here new near-infrared observations to constrain the number and characteristics of young stars in this globule. The idea is to combine colours and variability in the near-infrared as indicator for youth. After discussing our observations and data reduction in Sect. \\ref{data}, we search for YSOs in this region based on colour, variability, and complementary H$\\alpha$ narrow band photometry in Sect. \\ref{ident}. As it turns out, IC1396W harbours only a small number ($\\la 10$) of YSOs, indicating a low star formation efficiency (see Sect. \\ref{yso}). Our variability database yields a number of serendipitous discoveries, which are discussed in Sect. \\ref{seren}. ", "conclusions": "" }, "1003/1003.0895_arXiv.txt": { "abstract": "This is the first of a series of papers aimed at characterizing the populations detected in the high-latitude sky of the {\\it Fermi}-LAT survey. In this work we focus on the intrinsic spectral and flux properties of the source sample. We show that when selection effects are properly taken into account, {\\it Fermi} sources are on average steeper than previously found (e.g. in the bright source list) with an average photon index of 2.40$\\pm0.02$ over the entire 0.1--100\\,GeV energy band. We confirm that FSRQs have steeper spectra than BL Lac objects with an average index of 2.48$\\pm0.02$ versus 2.18$\\pm0.02$. Using several methods we build the deepest source count distribution at GeV energies deriving that the intrinsic source (i.e. blazar) surface density at F$_{100}\\geq10^{-9}$\\,ph cm$^{-2}$ s$^{-1}$ is 0.12$^{+0.03}_{-0.02}$\\,deg$^{-2}$. The integration of the source count distribution yields that point sources contribute 16$(\\pm1.8)$\\,\\% ($\\pm$7\\,\\% systematic uncertainty) of the GeV isotropic diffuse background. At the fluxes currently reached by LAT we can rule out the hypothesis that point-like sources (i.e. blazars) produce a larger fraction of the diffuse emission.\\\\ ", "introduction": "The origin of the extragalactic gamma-ray background (EGB) at GeV $\\gamma$-rays is one of the fundamental unsolved problems in astrophysics. The EGB was first detected by the SAS-2 mission \\citep{fichtel95} and its spectrum was measured with good accuracy by the Energetic Gamma Ray Experiment Telescope \\citep[EGRET][]{sreekumar98,strong04} on board the Compton Observatory. These observations by themselves do not provide much insight into the sources of the EGB. Blazars, active galactic nuclei (AGN) with a relativistic jet pointing close to our line of sight, represent the most numerous population detected by EGRET \\cite{hartman99} and their flux constitutes 15\\,\\% of the total EGB intensity (resolved sources plus diffuse emission). Therefore, undetected blazars (e.g. all the blazars under the sensitivity level of EGRET) are the most likely candidates for the origin of the bulk of the EGB emission. Studies of the luminosity function of blazars showed that the contribution of blazars to the EGRET EGB could be in the range from 20\\,\\% to 100\\,\\% \\citep[e.g.][]{stecker96,chiang98,muecke00}, although the newest derivations suggest that blazars are responsible for only $\\sim20$--$40$\\,\\% of the EGB \\citep[e.g.][]{narumoto06,dermer07,inoue09}. It is thus possible that the EGB emission encrypts in itself the signature of some of the most powerful and interesting phenomena in astrophysics. Intergalactic shocks produced by the assembly of Large Scale Structures \\citep[e.g.][]{loeb00,miniati02,keshet03,gabici03}, $\\gamma$-ray emission from galaxy clusters \\citep[e.g.][]{berrington03,pfrommer08}, emission from starburst as well as normal galaxies \\citep[e.g.][]{pavlidou02,thompson07}, are among the most likely candidates for the generation of diffuse the GeV emission. Dark matter (DM) which constitutes more than 80\\,\\% of the matter in the Universe can also provide a diffuse, cosmological, background of $\\gamma$-rays. Indeed, supersymmetric theories with R-parity predict that the lightest DM particles (i.e., the neutralinos) are stable and can annihilate into GeV $\\gamma$-rays \\citep[e.g.][]{jungman96,bergstrom00,ullio02,ahn07}. With the advent of the {\\it Fermi} Large Area Telescope (LAT) a better understanding of the origin of the GeV diffuse emission becomes possible. {\\it Fermi} has recently performed a new measurement of the EGB spectrum \\citep[also called isotropic diffuse background,][]{lat_edb}. This has been found to be consistent with a featureless power law with a photon index of $\\sim$2.4 in the 0.2--100\\,GeV energy range. The integrated flux (E$\\geq$100\\,MeV) of 1.03$(\\pm0.17)\\times10^{-5}$\\,ph cm$^{-2}$ s$^{-1}$ sr$^{-1}$ has been found to be significantly lower than the one of 1.45($\\pm0.05$)$\\times10^{-5}$\\,ph cm$^{-2}$ s$^{-1}$ sr$^{-1}$ determined from EGRET data \\citep[see][]{sreekumar98}. In this study we address the contribution of {\\it unresolved} point sources to the GeV diffuse emission and we discuss the implications. Early findings on the integrated emission of {\\it unresolved} blazars were already reported in \\cite{lat_lbas} using a sample of bright AGN detected in the first three months of {\\it Fermi} observations. The present work represents a large advance, with $\\sim$4 times more blazars and a detailed investigation of selection effects in source detection. This work is organized as follows. In $\\S$~\\ref{sec:spec} the intrinsic spectral properties of the {\\it Fermi} sources are determined. In $\\S$~\\ref{sec:sim} the Monte Carlo simulations used for this analyses are outlined with the inherent systematic uncertainties (see $\\S$~\\ref{sec:syst}). Finally the source counts distributions are derived in $\\S$~\\ref{sec:logn} and $\\S$~\\ref{sec:bands} while the contribution of point sources to the GeV diffuse background is determined in $\\S$~\\ref{sec:edb}. $\\S$~\\ref{sec:discussion} discusses and summarizes our findings. Since the final goal of this work is deriving the contribution of sources to the EGB, we will only use physical quantities (i.e. source flux and photon index) averaged over the time (11 months) included in the analysis for the First {\\it Fermi}-LAT catalog \\citep[1FGL,][]{cat1}. ", "conclusions": "" }, "1003/1003.0771_arXiv.txt": { "abstract": "The study of ionized gas morphology and kinematics in nine eXtremely Metal-Deficient (XMD) galaxies with the scanning Fabry-Perot interferometer on the SAO 6-m telescope is presented. Some of these very rare objects (with currently known range of O/H of 7.12 $<$ 12+$\\log$(O/H) $<$7.65, or Z\\sunn/35 $<$ Z $<$Z\\sunn/10) are believed to be the best proxies of `young' low-mass galaxies in the high-redshift Universe. One of the main goals of this study is to look for possible evidence of star formation (SF) activity induced by external perturbations. Recent results from HI mapping of a small subsample of XMD star-forming galaxies provided confident evidence for the important role of interaction-induced SF. Our observations provide complementary or new information that the great majority of the studied XMD dwarfs have strongly disturbed gas morphology and kinematics or the presence of detached components. We approximate the observed velocity fields by simple models of a rotating tilted thin disc, which allow us the robust detection of non-circular gas motions. These data, in turn, indicate the important role of current/recent interactions and mergers in the observed enhanced star formation. As a by-product of our observations, we obtained data for two LSB dwarf galaxies: Anon~J012544+075957 that is a companion of the merger system UGC~993, and SAO 0822+3545 which shows off-centre, asymmetric, low SFR star-forming regions, likely induced by the interaction with the companion XMD dwarf HS 0822+3542. ", "introduction": " ", "conclusions": " \\begin{enumerate} \\item The ionized gas kinematics in very low-metallicity star-forming galaxies, studied with FPI H$\\alpha$ observations are significantly disturbed and rarely can be well fitted by only single disc with regular rotation. \\item Several of our galaxies show more or less clear evidence from morphology and H$\\alpha$ velocity fields for various stages of mergers (UGC~993, SDSS~J1044+0353, SBS~1116+517, SBS~1159+545, HS~2236+1344). In two cases we have good evidence for two independently rotating discs (UGC~993 and HS~2236+1344). \\item In the other galaxies, the rotation component of the overall velocity field is important, but large disturbances appear. The residuals of the `best-fitting' rotation model imply either a recent merger, or sufficiently strong disturbance by nearby galaxies (HS~0822+3542 and UGC~772). \\item Probable starburst-induced shells were identified in several galaxies (UGC 772, UGC~993, SBS~0335--052E, SBS~1116+517) through their ionized gas velocity fields and velocity dispersion maps. We presented the results of simple simulations of the expected velocity patterns which one can observe in dwarf galaxies with expanding shells. To first order, this allowed a feel for the effect of shells on the results of the tilted-ring model. \\item The interacting/merging nature of the binary system of the well separated XMD galaxies SBS~0335--052E and W is best evident from their HI morphology and kinematics data. Despite a relatively large mutual distance, the tidal action of each component to the other clearly affects the gas dynamics of these very gas-rich objects and triggers the current SF burst and very likely the previous major SF episode \\citep[as emphasized by][] {SBS0335_GMRT}. To understand this unique interacting system (a nice representative of high-redshift young galaxies) in more detail, one needs a wide grid of models of interacting gas-rich galaxies, like, e.g., \"Identikit\" \\citep[][]{Barnes2009}, but including SF processes. \\item The H$\\alpha$ images and velocity fields for two LSB dwarfs, Anon~J012544+075957 and SAO~0822+3545, companions of two program XMD galaxies, are obtained and analysed. They can be used in statistical studies of SF in LSB dwarfs. The star formation in SAO~0822+3545 is highly asymmetric and takes place mainly in two knots at the northern edge, that is likely induced by the recent interaction with the nearby XMD BCG HS~0822+3542. \\end{enumerate} The statistics of XMD starbursting galaxies, for which kinematics of gas were studied in detail, is still insufficient. Nevertheless, the results of our FPI study of the ionized gas kinematics in the subsample of nine XMD star-forming galaxies, along with the complementary results of the GMRT HI study of a part of these and other XMD galaxies \\citep{Ekta08,SBS0335_GMRT,EC10}, indicate that strong interactions and mergers of very metal-poor dwarf galaxies are one of the major or significant factors triggering their current and recent starbursts. In particular, this is valid for the six most metal-poor (12+$\\log$(O/H)$<$7.30) dwarf starbursting galaxies. Merger-induced starbursts are consistent with the idea that the progenitors of such rare objects either are old and have been evolving very slowly on the cosmological timescale before the current starburst have occurred, or they are extremely metal-poor because they are comparatively young, and thus began their star formation and chemical enrichment with large delay." }, "1003/1003.4687_arXiv.txt": { "abstract": "We have developed a method to compute the possible distribution of radio emission regions in a typical pulsar magnetosphere, taking into account the viewing geometry and rotational effects of the neutron star. Our method can estimate the emission altitude and the radius of curvature of particle trajectory as a function of rotation phase for a given inclination angle, impact angle, spin-period, Lorentz factor, field line constant and the observation frequency. Further, using curvature radiation as the basic emission mechanism, we simulate the radio intensity profiles that would be observed from a given distribution of emission regions, for different values of radio frequency and Lorentz factor. We show clearly that rotation effects can introduce significant asymmetries into the observed radio profiles. We investigate the dependency of profile features on various pulsar parameters. We find that the radiation from a given ring of field lines can be seen over a large range of pulse longitudes, originating at different altitudes, with varying spectral intensity. Preferred heights of emission along discrete sets of field lines are required to reproduce realistic pulsar profiles, and we illustrate this for a known pulsar. Finally, we show how our model provides feasible explanations for the origin of core emission, and also for one-sided cones which have been observed in some pulsars. ", "introduction": "Of the various aspects relevant for solving the unresolved problem of radio emission from pulsars, two are probably the most significant : the actual mechanism of the emission itself, which is still not fully understood (e.g. Zhang 2006, Melrose 2006); and the effects of viewing geometry and pulsar rotation, which can significantly alter the properties of the pulsar profiles that the observer finally samples. The latter aspect has received significant attention in the recent years, but much still remains to be investigated. Blaskeiwicz et~al. (1991, hereafter BCW91) were the first to work out the basic effects of rotation, and showed that the observed asymmetry between the leading and trailing parts of pulsar radio profiles can be due to rotation effects. Further improvements were carried out by Hibschman \\& Arons (2001), who analyzed the first order effects of rotation on the polarization angle sweep. Later, Peyman \\& Gangadhara (2002), adapting the method of BCW91, refined the formulation and showed that the asymmetries due to rotation can be ascribed to the differences in the radius of curvature of the particle trajectories on the leading and trailing sides of the magnetic axis. By analysing the pulse profiles of some selected pulsars, which clearly show the core-cone structure in the emission beam, Gangadhara \\& Gupta (2001, hereafter GG01) and Gupta \\& Gangadhara (2003, hereafter GG03) showed that the asymmetry in the locations of the conal components around the central core component can be interpreted in terms of aberration and retardation (A/R) effects (combined effects of rotation and geometry), leading to useful estimates of emission heights of the conal components. Further refinements of these concepts have been carried out by Dyks et~al. (2004), and Gangadhara (2004 \\& 2005, hereafter G04 \\& G05), Dyks (2008) and Dyks et~al. (2009). All of the above said works have established that rotation effects are of significant importance in understanding the observed emission profiles of radio pulsars. What has been found wanting is a detailed, quantitative treatment that couples the rotation effects in the pulsar magnetosphere to the possible emission physics and to the emission and viewing geometries, to produce observable radio profiles. Some impediments to this have been recently overcome by Thomas and Gangadhara (2007, hereafter TG07), who have considered in detail the dynamics of relativistic charged particles in the radio emission region, and obtained analytical expressions for the particle trajectory and it's radius of curvature. In this paper, we describe a scheme for simulation of pulsar profiles that encompasses a detailed treatment of all the effects mentioned above. We start with describing the background and motivation for the work (\\S\\ref{sec:significance}), then go on to the profile simulation method (\\S\\ref{sec:prof}). We describe the main results from our study and their dependence on pulsar parameters in \\S\\ref{sec:results-discussion}, and discuss how realistic pulsar profiles may be obtained from our model. We also address the issues of core emission, one-sided or partial cones, and extension of our method to other models of emission physics. Our final conclusions are summarized in \\S\\ref{sec:summary}. ", "conclusions": "\\label{sec:results-discussion} \\subsection {General results} We first discuss the general results and trends that are deduced from our simulation studies, as illustrated in the results displayed in the plots given in \\S\\ref{app:figs}. \\subsubsection{Emission heights} The heights for the allowed emission spots have a minimum value near the magnetic meridian ($\\phi'=0$), with smoothly increasing values on the leading and trailing sides. However, the variation of height with rotation phase is asymmetric such that the increment with rotation phase is always faster for the trailing side than for the leading side. Whereas this increase with rotation phase is purely geometric, the asymmetry in this is due to the modification of particle trajectories produced by rotation. On the leading side, the emission beam bends in the direction favourable to rotation and hence advances in azimuthal phase by $\\delta\\phi'_{\\rm aber}$ from that of the corresponding field line tangent. Hence at a fixed $\\phi',$ an emission spot located at a lower emission height than in the non-rotating case will satisfy Eq.\\ref{eq_los_vel}. In contrast, on the trailing side, the bending of the emission beam causes a lag in azimuthal phase by $\\delta\\phi'_{\\rm aber}$ from the corresponding field line tangent. This lag can be compensated by an emission spot located at a different altitude than the non-rotating case. Hence an emission spot located at a different emission height will satisfy Eq.\\ref{eq_los_vel} and contribute radiation to the observer. Both the value of the minimum height (at $\\phi'=0$), as well as the asymmetry, are larger for the inner field lines as compared to the outer field lines. This supports the intuitive expectation that outer field lines would be visible out to much larger pulse longitude ranges than inner field lines. Further, it is seen that the minimum height and asymmetry of the variation increase with geometry, being more for larger values of $\\alpha$ and $\\beta$. Also, pulsars with larger values of $\\alpha$ will have a larger range of variation of allowed emission heights, for the same field line. It is reasonable to surmise that a similar value of emission height can be seen recursively for several combinations of larger and smaller values of $\\alpha$ and $S_{\\rm L},$ for a certain $\\phi'$. Thus we find that the values of $\\alpha$ and $S_{\\rm L}$ dominantly decide the range of emission heights. Even otherwise, the $\\alpha$ dependence of the emission height can be directly understood from Eq.(\\ref{eq_solution_0_order}), since the expression $r(t)$ is explicitly dependent on the value of $\\alpha.$ \\subsubsection{Radius of curvature} The radius of curvature inferred for the possible emission spots also varies signficantly with rotation phase, and it is significantly asymmetric between leading and trailing sides. If rotation effects were not considered, this radius of curvature would be same as that of the corresponding field line and be symmetric around zero pulse phase. Given that the observed radius of curvature of the particle trajectory is a combination of the curvature of the field line and curvature introduced due to rotation, the observed asymmetry is a rotation effect and can be understood as a combination of two effects: first, as the allowed heights are different on the leading and trailing sides (for the same phase $\\pm \\phi'$ on either side of the zero phase), the radius of curvature of the field line itself would be different; second, on the leading side, the curvature of the field line gets combined with that induced by rotation (TG07), resulting in a lower value of the $\\rho$ for the particle trajectory. Whereas on the trailing side, the two curvatures are opposed and hence the net curvature is reduced, leading to larger values of $\\rho.$ In fact, for some field lines, rotation induced curvature can cancel the curvature due to the field lines, at some points on the trailing side, resulting in sharply peaked curves for $\\rho$ for certain values of $S_{\\rm L},$ as seen in the plots in Figs. \\ref{fig_A30-B3}, \\ref{fig_A60-B1} \\& \\ref{fig_A90-B1}. The variation of $\\rho$ on the leading side is more or less steady, while on the trailing side it often varies very rapidly due to the aforesaid reasons. The trend is that the asymmetry seen in $\\rho$ will be higher for larger values of $\\alpha$ and smaller values of $S_{\\rm L}$ and hence it is a combined effect of both. As mentioned earlier, the rotation effects are more at higher $\\alpha$ and hence the larger asymmetry. Since the range of emission altitudes covered by the emission spots for inner field lines are higher than that of the outer field lines, the corresponding values of $\\rho$ also will have a higher range and a higher asymmetry than the ones for the outer field lines. However, there are variations in the amount of asymmetry with in the field lines when the $\\alpha$ varies, and hence a steady variation of asymmetry with $\\alpha$ will not be observed as for the emission heights. Comparing the $\\rho/\\rho_{\\rm L}$ plots for figs. \\ref{fig_A60-B1} \\& \\ref{fig_A90-B1}, we can find on the leading side that the radius of curvature gets reduced when $\\alpha$ increases from $60^{\\circ}$ to $90^{\\circ}$ for all field lines. But on the trailing side the behaviour is slightly different. The inner field lines ($S_{\\rm L}=0.1 ~\\&~ 0.3 $) have the radius of curvature slightly reduced on the trailing side, while an increment in radius of curvature is seen for outer field lines ($S_{\\rm L}=0.1 ~\\&~ 0.3 $), when $\\alpha$ increases from $60^{\\circ}$ to $90^{\\circ}.$ Another comparison of the $\\rho/\\rho_{\\rm L}$ plots for the figs. ~\\ref{fig_A30-B1} ~\\&~ \\ref{fig_A60-B1} obviously shows the same trend for the leading side. However, the variation of $\\rho/\\rho_{\\rm L}$ on trailing side shows a slightly different behaviour, varying among field lines. The variation of $\\rho $ on the trailing side is not in a steady pattern as on the leading side. The competing curvatures due to rotation induced curvature and the intrinsic field line curvature gives a highly varying pattern for $\\rho $ on the trailing side. \\subsubsection{Spectral Intensity, $I_{\\omega}$} The derived spectral intensity curves reflect the asymmetry inherited from the variation of $\\rho$ with pulse phase, combined with the effects of the value of $\\gamma$. In particular, it is readily seen that the intensity dramatically evolves with $\\gamma.$ For lower values of $\\gamma$, the $I_{\\omega}$ has a stronger leading part while for higher values of $\\gamma$ the $I_{\\omega}$ has a stronger trailing part. This effect can be better understood by considering the parameter $\\rho_{\\rm p}$ (defined in Eq.~(\\ref{eq_rho_p})), which gives the value of $\\rho$ at which the spectral intensity peaks, for given values of frequency and $\\gamma$. For values of $\\rho$ greater than or less than $\\rho_{\\rm p}$, the spectral intensity falls monotonically. Further, the peak value of the spectral intensity also depends on the specific value of gamma, as per Eq.~(\\ref{eq_spec_Int}). Hence, the variation of spectral intensity with pulse longitude can be inferred from that of $\\rho$ with longitude, for different values of gamma. This is illustrated in fig. ~\\ref{fig_A30-B1-SPEC} where $\\rho/\\rho_{\\rm p}$ and the corresponding $I_{\\omega}(\\phi')$ are plotted side by side as a function of the rotation phase, for specific combinations of parameters. Three different cases are useful to consider. For situations where $ \\rho/\\rho_{\\rm p}$ is greater than 1.0 for the entire pulse window, the spectral intensity curve shows a maximum at $\\phi'=0$ and falls asymmetrically on either side, with the reduction in intensity being larger for the trailing side, due to the faster increase of $\\rho$. This effect, which is seen for relatively small values of $\\gamma$ (less than 400-600), naturally leads to asymmetric pulse profiles, with possibilities for sharply one-sided profiles. It is interesting to note that in some cases, the intensity on the trailing side can drop to negligible values, compared to its value at the corresponding longitude on the leading side. This could be a natural explanation for the one-sided cones reported in literature, and is discussed in more detail in \\S\\ref{sec:partial-cones}. For situations where $ \\rho/\\rho_{\\rm p}$ is less than 1.0 for the entire pulse window, the spectral intensity curve shows a minimum at $\\phi'=0$ and rises asymmetrically on either side, with the increase being larger on the trailing side. However, the contrast in the intensity levels between leading and trailing sides is typically not as high as for the first kind. This behaviour is seen for relatively large values of $\\gamma$. For intermediate cases, where values of $\\rho/\\rho_{\\rm p}$ less and greater than 1.0 can occur at different points in the pulse longitude window, we see more complicated shapes for the spectral intensity curves, including multiple maxima at different pulse longitudes. For a given geometry and $s_{\\rm L}$ value, the transition through these 3 different cases can take place as $\\gamma$ is varied over a range of values. Thus, lower values of $\\gamma$ tend to produce profiles with strong leading and weak trailing intensities which get converted to weak leading and strong trailing kind as the $\\gamma$ increases to very high values (e.g. 2nd and 3rd panels from the top in fig. \\ref{fig_A30-B1-SPEC}). Though the contrast in intensity is less for the latter, the absolute value of the spectral intensity is higher, due to the $\\gamma$ dependence in Eq.~(\\ref{eq_spec_Int}). Most of the asymmetric intensity profile effects become more dramatic for inner field lines and for more orthogonal rotators (larger $\\alpha$) and smaller values of $\\beta.$ \\subsubsection{$\\theta$ and $\\phi$} The values of $\\theta$ and $\\phi$ are asymmetric on the leading and the trailing sides of the profiles, while they are symmetric in the non-rotating case (G04). This asymmetry is also an effect of rotation. Since $\\theta$ and $\\phi$ are functions of $\\phi'$ their values are affected by the aberration phase shift which is different on leading and the trailing sides. Their values are dominantly decided by $\\alpha$ and $\\beta$. As expected, the shape of the $\\phi$ curve closely resembles the S-shape of the polarization angle curve (see fig.~\\ref{fig_A30-B1}). \\subsubsection{Mis-alignment angle, $\\eta_{\\rm mis}$} The Mis-alignment angle $\\eta_{\\rm mis},$ defined as $\\eta_{\\rm mis}=\\cos^{-1}({\\hat \\textbf n}\\cdot{\\hat \\textbf v}),$ gives an estimate of the offset between $\\hat \\textbf n$ and the estimated $\\hat \\textbf v.$ In principle, for a perfect estimation of the emission point, the line of sight should exactly coincide with the velocity vector, and hence $\\eta_{\\rm mis}$ should be zero. However, in actual computations, $\\eta_{\\rm mis}$ always has a small, finite value. A quick check of the accuracy of the computation is provided by the value of $\\eta_{\\rm mis}$ : a lower value implies a higher precision of estimation of the emission spot, and vice versa for a higher value. A rough classification that can be taken for the precision of the estimation is: a value of $\\eta_{\\rm mis} \\ll 1$ indicates a highly precise estimation of the emission spot, and vice versa for $\\eta_{\\rm mis} \\gg 1.$ By this scheme, we find that there is satisfactory precision for all estimations for $r$ within 20 \\% of $r_{\\rm L}$ (see fig.~\\ref{fig_MISANG-INNER}). In some cases the $\\eta_{\\rm mis}$ exceeds 1, but only when $r/r_{\\rm L}> 0.2.$ However, according to established observational results radio emission heights are restricted within 10 \\% of $r_{\\rm L}$ for normal pulsars (e.g. Kijak 2001), and hence our method is quite satisfactory in this regime of interest. \\subsection{Effects of Parameters} The above described behaviour of the height of emission spots, radius of curvature and spectral intensity are strongly dependent on the parameters like geometry, field line location, radio frequency and Lorentz factor of the particles. In some cases, there is a complex interplay between the dependencies on these different parameters. Here, we explore some of these effects in detail. The generic effects of $\\gamma,$ $\\alpha$ and $\\beta,$ $\\omega$ and $S_{\\rm L}$ are listed briefly below in an order that may characterize the hierarchy of their effects on total intensity profiles. \\subsubsection{Inclination angle, $\\alpha$} \\label{alpha-effect} The parameter $\\alpha $ is a major driver of the effects of rotation, and has the strongest influence on our results and conclusions. The rotation effects (leading-trailing asymmetry of $r,$ $\\rho$ and $I_{\\rm \\omega}$) are more prominent for large values of $\\alpha$ and less for small values of $\\alpha.$ The range of emission altitudes is found to be relatively high for lower values of $\\alpha$, and relatively low for higher values of $\\alpha,$ being the lowest for $\\alpha=90^{\\circ}.$ Like wise, $\\rho$ appears to reach higher values for higher $\\alpha$ and vice versa for lower $\\alpha.$ \\subsubsection{Normalized foot value of the field lines, $ S_{\\rm L} $} The effect of moving from inner to outer regions of the magnetosphere (increasing $S_{\\rm L}$ values) also has a very dramatic effect on the results. Rotation effects are strongest for the innermost field lines, and decrease significantly for larger $S_{\\rm L}$ values. For relatively small values of $S_{\\rm L}$ (usually $\\le$ 0.3), the leading part has emission heights that vary relatively gently with $\\phi'$, whereas the trailing part shows steeply rising emission height. The emission heights become less asymmetrical for increasing values of $S_{\\rm L}.$ The values of $\\rho$ steadily increase with decreasing $S_{\\rm L}$ (i.e. inner field lines) on the trailing side. Dramatic effects such as very large, peaked values of $\\rho$ on the trailing side, owing to the mutual cancellation of intrinsic and rotation induced curvatures, are seen only on inner field lines. This peak shifts closer towards zero pulse phase as the value of $ S_{\\rm L}$ becomes smaller. For outer field lines, the profiles are much more symmetric, and since $\\rho/\\rho_{\\rm p} < 1$ for a significant range of pulse phase on either side of $\\phi'=0$, the profiles more often exhibit minima at $\\phi'=0$, even for moderate values of $\\gamma$. \\subsubsection{The lorentz factor, $\\gamma$} \\label{gamma-effect} The effect of $\\gamma$ has a very clear signature on the asymmetry of the spectral intensity profiles. For lower values of $\\gamma$, there can be strong asymmetries with leading side stronger than the trailing side, and maxima at $\\phi'=0$. For larger values of $\\gamma$, the sense of this asymmetry can reverse, with trailing side becoming stronger than the leading, and a minima at $\\phi'=0$ ; however, the degree of the asymmetry, as measured by the ratio of the intensities at corresponding longitudes, is generally less than that for the case for the low $\\gamma$ values. \\subsubsection{Emission frequency, $\\omega$} The frequency of emission, $\\omega$, acts as a counter to the effect of $\\gamma$, though in a relatively weak manner, as can be understood from Eqs. ~(\\ref{eq_spec_Int}) and ~(\\ref{eq_rho_p}). Thus, an increase in $\\omega$ produces changes which can be compensated by a corresponding change in $\\gamma$ by a factor proportional to $\\omega^{1/3}$. In certain cases, this could result in profiles where the sense of asymmetry between leading and trailing sides could reverse over a large enough range of radio frequencies. Such effects are seen sometimes in some real profiles. \\subsection{Realistic profiles} One of the significant results from our simulation studies is that the possible regions of emission associated with a given annular ring of field lines (characterised by a constant value of $S_{\\rm L}$) are visible over a wide range of pulse phase, albeit with different intensity levels. This aspect, combined with the results for field lines with different values of $S_{\\rm L}$, leads to the conclusion that a very large fraction of the pulsar magnetosphere is potentially visible to us. This results in simulated pulsar profiles that are very different from the observed profiles of real pulsars which appear to have well defined emission components, restricted in pulse phase extent to occupy only some fraction of the on-pulse window. This disparity with the observed profiles persists even after we incorporate into our simulations the models of discrete, annular conal rings of emission. Hence, in order to reproduce realistic profiles matching with observations, we need some additional constraints for the emission regions. In the most general case, such non-uniformities in the distribution of emission regions can exist in any of the three coordinate directions, viz. $r$, $\\theta$ and $\\phi$. Non-uniformity of emission in the $\\theta$ direction is achieved in some sense at the basic level, by considering only discrete sets of $S_{\\rm L}$ values for active emission regions. As discussed, this is not enough to constrain the intensity variations to reproduce realistic pulsar profiles. The possibility of non-uniform emission along the $\\phi$ coordinate could help produce discrete emission components in the observed profiles. As seen in fig. ~\\ref{fig_A30-B1}, for a given value of $S_{\\rm L},$ the emission at different pulse longitudes originates at widely different $\\phi$ locations. If the sources of charged particles were located only at fixed $\\phi$ points along the ring of constant $S_{\\rm L},$ then these could be arranged to modulate the simulated intensity pattern with a suitable ``window'' function, to obtain discrete emission components in the observed profile. However, the well known phenomenon of sub-pulse drift argues against this being a viable option. Sub-pulse drift, which is now believed to be fairly common in known pulsars (e.g. Velterwede et~al. 2006), wipes out any azimuthal discretization of sources of emission in the pulsar magnetosphere -- it would lead to ``filling up'' of the intensity average profile over a given range of pulse longitude, as is observed in drifting sub-pulses that occur under any discrete emission component in known pulsars. The third option is to have non-uniform emission in the radial direction, i.e. preferred heights of emission for a given set of field lines. Since the contributions at different pulse phases are from different heights, this would naturally lead to non-uniform distribution of intensity in pulse phase, resulting in realistic looking pulse profiles. The idea of preferred heights of emission in the pulsar magnetosphere is not entirely new -- the ``radius to frequency mapping'' model for pulsar emission postulates different heights for different frequencies, with the height of emission increasing for decreasing frequency values (Kijak \\& Gil 1997,1998; GG01, GG03). \\begin{figure} \\begin{center} \\epsfxsize= 6.5 cm \\epsfysize=0cm \\rotatebox{0}{\\epsfbox{f2.eps}} \\caption[short_title]{\\small PSR B2111+46 at 333 MHz : The emission height for each $S_{\\rm L}$ value associated with a particular component is plotted in panel (a). The simulated and un-modulated spectral intensity curve corresponding to each $S_{\\rm L}$ value is plotted in panel (b). The simulated sub-components, after applying Eq.(\\ref{Eq-Mod-spec}) for the best fit values given in Table.\\ref{tabsim} are shown in panel (c). The sum-total of the simulated sub-components giving the final profile (solid curve) is shown superposed with the observed profile (dotted curve) in panel (d). \\label{2111+46-333}} \\end{center} \\end{figure} Preferred emission heights of emission with a spread in the $r$ direction, can be modeled as a multiplication of the spectral intensity ${I}_{\\omega}(\\phi')$ with a modulating function, $F^{\\rm i}_{\\omega} (\\phi')$. For modelling a profile as a sum total of emissions from a core and several discrete conal regions, the modulated spectral intensity can be expressed as \\begin{eqnarray}\\label{eq_mod_funct} {I}_{\\omega}(\\phi') &=& \\sum_{i}^{N} {I}^{\\rm i}_{\\omega} (\\phi')\\,\\, F^{\\rm i}_{\\omega} (\\phi')~, \\label{Eq-Tot-spec} \\\\ F^{\\rm i}_{\\omega} (\\phi')&=& A^{\\rm i}_{\\omega} \\exp\\left[-\\left( \\frac{r^{\\rm i}_{0}(\\phi')- H^{\\rm i}_{\\omega }} {2\\, \\Delta H^{\\rm i}_{\\omega}}\\right)^2 \\right]~, \\label{Eq-Mod-spec} \\end{eqnarray} where the index $i$ represents, a corresponding pair of leading-trailing components presumed to be arising from a particular ring of field lines; while $i=1$ exclusively represents the central core component. Here $N$ is the total number of such discrete emission components (for example, $N=3$ would correspond to a 5 component profile forming a central core and two pairs of conal components), ${I}_{\\omega}(\\phi')$ is the total spectral intensity from all field lines combined, while ${I}^{\\rm i}_{\\omega} (\\phi')$ is the spectral intensity from the $i$th ring of field lines. For a given emission region, $H^{\\rm i}_{\\omega}$ represents the mean height, and $\\Delta H^{\\rm i}_{\\omega}$ represents the spread of the region. The variable $r^{\\rm i}_{0}(\\phi')$ represents the values of emission altitude for each value of $\\phi'$ estimated by the simulation method as described earlier, for the $i$th ring of field lines. To map this intensity as a function of rotation phase as seen in the observer's frame, the effects of retardation and aberration need to be included explicitly. The effect of aberration is estimated by default and the $\\phi'$ is inclusive of the aberration phase shift. The retardation phase shift is to be estimated from the value of $r$ corresponding to the emission spot. The rotation phase corresponding to ${I}_{\\omega}$ is updated after adding the retardation phase shift $d\\phi'_{\\rm ret}$ with $ \\phi',$ and this is represented by the mapping of the ordered pair $(\\phi',{I}_{\\omega})\\rightarrow (\\phi'+d\\phi'_{\\rm ret},{I}_{\\omega}).$ The $d\\phi'_{\\rm ret}$ can be estimated as (G05) : $$d\\phi'_{\\rm ret}=\\frac{1}{r_{\\rm L}}({\\vec r}\\cdot{\\hat n}),$$ where ${\\vec r}=r{\\hat e}_{\\rm r}$ and the expression for ${\\hat e}_{\\rm r}$ is given in \\S\\ref{app:approx}. The height of emission ($ H^{\\rm i}_{\\omega }$) and the normalized foot value ($S^{\\rm i}_{\\rm L}$) corresponds to the peak of the $i$th component of the profile. The $\\Delta H^{\\rm i}_{ \\omega}$ and $A^{\\rm i}_{\\omega}$ are model parameters. \\begin{table} \\centering \\begin{minipage}{100mm} \\caption{The parameters for simulating the profiles of PSR B2111+46 }\\label{tabsim} \\begin{tabular}{ccccccccccccccccccccccccccccccc} \\hline \\\\ \\multicolumn{1}{c}{Frequency} & ${\\rm i}^{a}$ & \\multicolumn{1}{c}{$H^{i}_{ \\omega}$} & \\multicolumn{1}{c}{$\\Delta H^{i}_{ \\omega}$} & \\multicolumn{1}{c}{$A^{i}_{\\omega}$ } & \\multicolumn{1}{c}{$S_{\\rm L}$ } & \\multicolumn{1}{c}{Lorentz factor} \\\\ \\multicolumn{1}{c}{MHz} & &\\multicolumn{1}{c}{Km} &\\multicolumn{1}{c}{Km} & \\multicolumn{1}{c}{ } & \\multicolumn{1}{c}{ } & \\multicolumn{1}{c}{$\\gamma$} \\\\ \\hline \\hline 333 & 1 & 1500 & 600 & 1.8 & 0.08 & 750 \\\\ & 2 & 1834 & 500 & 0.14 & 0.18 & 750 \\\\ & 3 & 3800 & 500 & 0.7 & 0.3 & 500 \\\\ 408 & 1 & 300 & 850 & 2.2 & 0.09 & 750 \\\\ & 2 & 1200 & 750 & 0.1 & 0.22 & 750 \\\\ & 3 & 3000 & 400 & 0.9 & 0.35 & 500 \\\\ 610 & 1 & 200 & 700 & 1.7 & 0.13 & 750 \\\\ & 2 & 500 & 450 & 0.3 & 0.31 & 700 \\\\ & 3 & 2600 & 500 & 1.5 & 0.35 & 550 \\\\ \\hline \\end{tabular}\\\\ \\\\ \\small{$^a$ $i=1$ represents the core component, \\\\ $i=2$ represents the inner conal component, \\\\ $i=3$ represents the outer conal component.} \\end{minipage} \\end{table} \\begin{figure*} \\begin{center} \\epsfxsize= 14 cm \\epsfysize=0cm \\rotatebox{0}{\\epsfbox{f3.eps}} \\caption[short_title]{\\small PSR B2111+46 at 408 \\& 610 MHz: In panels (a) and (c) , the dotted curves shows the simulated sub-components and the continuous curve shows their sum-total. In panels (b) and (d), this final simulated profile (solid curve) is superposed with the observed profile (dotted curve). \\label{2111+46-ALL}} \\end{center} \\end{figure*} \\subsection{Profiles for PSR B2111+46 : a test case} Using the afore said methods for simulation of pulse profiles, we have attempted to reproduce the intensity profiles of PSR B2111+46 obtained from EPN data base and GMRT data, at multiple radio frequencies. This pulsar has a multi-component profile, with a well identified core component and 2 cones of emission (e.g. Zhang et~al. 2007). It has a rotation period of 1.014 sec and $\\alpha=14$ and $\\beta=-1.4$ (Mitra \\& Li 2004). The other parameters used in the simulation are listed in Table \\ref{tabsim}. The values of emission heights $H^{\\rm i}_{\\omega}$ and field line locations $S^{\\rm i}_{\\rm L}$ for the discrete emission components are the values from estimates employing the method given in Thomas and Gangadhara (2009). The zero phase of the profile is fixed on the basis of the analysis of core emission of this pulsar, using the method developed in the the same work. The simulation method is illustrated in fig.~\\ref{2111+46-333}. Values of $S_{\\rm L}$ corresponding to the core and conal components are employed in generating the emission height plots in panel (a) of fig.~\\ref{2111+46-333}. The corresponding spectral intensity plot for each of these $S_{\\rm L}$ values, for the final best fit choice of $\\gamma$ (in Table \\ref{tabsim}) is shown in panel (b) of this figure. The individual components generated after applying the best fit height function are shown in panel (c) and the sum total intensity profile is shown in panel (d), along with the observed profile. Best fits of these profiles to the observed data were obtained by varying $ A^{\\rm i}_{\\omega}$ and $\\Delta H^{\\rm i}_{\\omega}$ in the function $F^{\\rm i}_{\\omega}(\\phi')$, and by varying the value of $\\gamma$ in the range 100 to 1000. The same procedure is repeated for 408 MHz and 610 MHz profiles and the results are shown in fig.~\\ref{2111+46-ALL}. All the final parameters and best-fit results are summarised in Table.\\ref{tabsim}. An encouraging first order match between the simulated and observed profiles has been achieved (see panel (d) of fig.~\\ref{2111+46-333} \\& panels (b) and (d) of fig.~\\ref{2111+46-ALL}). The core component is quite well fit for most of the cases, and so are the leading conal components. There is some mismatch in the widths of the conal components, especially for the trailing side, where the real data shows a smoother blending of the components, compared to the simulated profile where the components appear more narrow and relatively well separated. It is remarkable that with a single value of $A^{\\rm i}_{\\omega}$ for a leading-trailing pair of cones of emission, the ratios of the peak values of the intensity of the leading and trailing components of the cones match so well with the real data. It is also interesting to note that the best fit values for $\\gamma$ are very similar for a given emission component, at different frequencies, supporting a model of a common bunch of accelerated charged particles being responsible for the emission at different frequencies. Further, that the spread of $\\gamma$ values across the different emission components is also quite small, indicates very similar operating conditions over most of the magnetosphere. The best fit values for $\\Delta H^{\\rm i}_{\\omega}$, though reasonable, are somewhat large in amplitude, indicating somewhat extended emission regions in the magnetosphere. We note that these relatively large values of $\\Delta H^{\\rm i}_{\\omega}$ and some of the limitations of the fits may be due to the lack of some generalizations in our model. These include factors like coherency of emission, a realistic spread of $\\gamma$ values around the mean values obtained here, as well as a realistic spread in the values of $S_{\\rm L}$ due to finite thickness of the rings of emission on the polar cap. Whereas a detailed treatment of all of these is beyond the scope of this work and will be taken up later, some basic inferences can still be drawn. For example, if a small range of $S_{\\rm L}$ values around the mean is considered, it is easy to argue that much of the width of a profile component can be filled up by radiation from such a bunch of field lines. This can be understood from panel (a) of fig. ~\\ref {2111+46-333}, where a line of constant height intersects the curve for a given field line at two points, one each on the leading and trailing side. The phase of this point of intersection will move systematically as we go to neighbouring field lines. Hence, wider profile components can be achieved with smaller values of $\\Delta H^{\\rm i}_{\\omega}$. Furthermore, due to the asymmetry in the emission height curves, the shift in phase with change of $S_{\\rm_L}$ is more on the trailing side, which would naturally lead to broader component widths and better `blending' of the components in the profile, something that is not as easily achievable by having a large range of emission heights (as the shift of phase for a given separation of heights on a given field line is lesser on the trailing side). One indicator of the significance of the spread of $S_{\\rm_L}$ values is the amount by which $S_{\\rm_L}$ needs to be changed to move the peak of one conal component to the point half-way to the peak of the next conal component. Not very surprisingly, our rough estimates show that the required change in $S_{\\rm_L}$ is close to the half-way point to the $S_{\\rm_L}$ value of the next cone, which would indicate a closely packed structure of concentric rings. The component profiles could be further influenced by considering a distribution of $\\gamma$ values associated with the emitting particles. We have found that significant shifts in the peaks of the leading and trailing pair of components for simulated profiles are obtained for lower $\\gamma$ values ($\\gamma < 500$), while the peak positions appear almost frozen for increasing $\\gamma$ values. Thus it is realistic to assume that a spread of $\\gamma$ values can broaden the emission components. This factor also may reduce the $\\Delta H^{\\rm i}_{\\omega}$ required to effect a good fit. Nevertheless, we would like to point out that there is only one unique combination of the parameters that can produce a profile which is similar to the observed one. We have not found any degenerate combination of values for the parameters that are shown in Table.\\ref{tabsim}. Thus, the similarity of the simulated profiles with the observed ones gives an assurance that, we should be able to simulate the observed profiles with greater similitude with a model overcoming the above-said limitations. \\subsection{Core emission} The generation of the profile components for PSR B2111+46 described above naturally leads to a discussion on the core emission. In fact, the study of the phenomena of core emission has spawned enormous amount of literature. Perhaps the most notable ones are the landmark work by Rankin (1983) that systematized the pulsar emission profile into `core' and `cone', and the succeeding works by Rankin (1993a \\& 1993b) that further developed the core-cone classification scheme. The hollow cone model was invoked to explain the geometry (e.g. Taylor \\& Stinebring 1986) and the origin of core emission. Radhakrisnan and Rankin (1990) have conjectured that the emission mechanism for cores might be different from that of cones, owing to the behaviour of polarization position angle curve near the core being different from the rotating vector model. However, there are no satisfactory theoretical grounds for postulating diverse mechanisms for cores and cones. A major difficulty that curvature radiation encounters in explaining the core emission is the insufficient curvature of the almost straight field lines in a region relatively close to the pulsar polar cap. Since the intensity of emission is proportional to $1/\\rho^2$, the values of $\\rho$ provided by the intrinsic curvature of the field lines is too large and hence insufficient to generate enough intensity of emission typically observed for core component. This factor even prompted invoking other emission mechanisms for explaining core emission (e.g. Wang et~al. 1989). In our simulation studies, the presence of the core component comes about quite naturally. It can be seen from all the plots of spectral intensity in \\S\\ref{app:figs} that the emission from regions near the profile centre is comparable (for higher $\\gamma $ values) or even somewhat higher (for lower $\\gamma $ values) than that from regions in the wings of the profile. This happens because we get low values of $\\rho$ for inner field lines near $\\phi' \\approx 0$, which are comparable to that of outer field lines, and this occurs consistently for all the combinations of $\\alpha$ and $\\beta$ (see the panels for $\\rho/r_{\\rm L}$ in \\S\\ref{app:figs}). The reason is that rotation induces significant curvature into the trajectory of particles, even though they are confined to move along the nearly straight inner field lines. The forces of constraint act in such a way that the particle is hardly allowed to deviate away from the field line on which it is moving, and the resulting scenario is discussed in detail in TG07. Due to the co-rotation of the field lines and the action of the aforesaid forces of constraint the charged particles are added with a velocity component in the direction of rotation, which is nearly perpendicular to the velocity component parallel to the field line, in the observer's frame. This induces an additional curvature in the trajectory of the particle and makes it significantly different from that of the field line curvature in the observer's frame of reference (TG07). Hence the trajectory of charge particles moving on almost straight field lines near the magntic axis can have a highly curved trajectory and hence a relatively low value of radius of curvature that is significantly different from that of the field lines. This scenario allows for significant emission near the central region of the profiles. By applying Eq.~(\\ref{Eq-Tot-spec}) and Eq.~(\\ref{Eq-Mod-spec}) appropriately, as described earlier, profile shapes resembling strong core components can be easily generated. Hence by applying our method, we provide a natural explanation for core emission, that circumvents the issue of too high $\\rho$ that precludes a strong core component with curvature emission. In the simulation of the profiles for PSR B2111+46, we find that the core originates from have relatively inner field lines and lower emission heights than the cones. Assuming the same mechanism of emission, viz. curvature radiation, for the core and conal component, we are able to produce a simulated core component that matches quite well with the observed profile. We notice that the best-fit values for the amplification factor $A^{i}_{\\omega}$ found in the simulation for the core component (Table ~\\ref{tabsim}) are comparable to those of the cones. These values are not unduly high, considering the situation that the density of plasma should be relatively higher for the lower altitude and hence an additional factor for relatively stronger emission at lower altitudes. Our profile-matching of PSR B2111+46 thus shows that strong core emission can originate from inner field lines due to curvature emission. \\subsection{Partial cones}\\label{sec:partial-cones} According to LM88, partial cone profiles are the ones in which one side of a double component conal profile is either missing or significantly suppressed. These are recognised by the characteristic that the steepest gradient of the polarization position angle is observed towards one edge of the total intensity profile, instead of being located more centrally in the profile as the rotating vector model postulates. LM88 speculated that this happens when the polar cap is only partially (and asymmetrically) active. It is significant that, out of the 32 pulsars listed in LM88 that display the partial cone phenomena, as many as 22 have the steepest gradient point occurring in the trailing part of the profile. In other words, most of the partial cone profiles show a strong leading component and an almost absent trailing component. \\begin{table*} \\centering \\begin{minipage}{100mm} \\caption{The parameter values employed in fig.\\ref{PARTIAL} }\\label{tabpar} \\begin{tabular}{cccccccccccccccccccccccccc} \\hline \\\\ \\multicolumn{1}{c}{Panel} & $\\alpha$ & \\multicolumn{1}{c}{ $\\beta$} & \\multicolumn{1}{c}{$S_{\\rm L}$ } &\\multicolumn{1}{c}{$H^{i}_{ \\omega}$} &\\multicolumn{1}{c}{$\\Delta H^{i}_{ \\omega}$} \\\\ \\multicolumn{1}{c}{No.} &\\multicolumn{1}{c}{$[^o]$} &\\multicolumn{1}{c}{$[^o]$} &\\multicolumn{1}{c}{ } & \\multicolumn{1}{c}{$Km$ } & \\multicolumn{1}{c}{$Km$} \\\\ \\hline \\hline 1 & 90 & 1 & 0.3 & 1000 & 200 \\\\ \\\\ 2 & 90 & 1 & 0.5 & 1200 & 150 \\\\ \\\\ 3 & 60 & 1 & 0.3 & 1000 & 200 \\\\ \\\\ 4 & 60 & 1 & 0.5 & 1400 & 150 \\\\ \\\\ 5 & 30 & 1 & 0.1 & 2500 & 500 \\\\ \\\\ 6 & 30 & 1 & 0.3 & 1000 & 150 \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*} There are two possible scenarios that have been postulated to explain partial cones : (1) only a part of the polar cap is active (this works for both kinds of partial cones) or (2) the A/R effects are so large as to shift the entire active region of the intensity profile towards the leading side (this works for the strong leading type partial cones, which are the majority). However, Mitra et~al (2007) studied several pulsars with partial cones with very high sensitivity observations and found that the almost-absent parts of the cones do flare up occasionally and show emission for about a few percentage of the total time. This tends to rule out both the scenarios above, and requires an explanation where the intensity is naturally suppressed in one side of the cone. In our simulation studies, one sided cones appear as a natural by-product. We notice that for smaller values of $\\alpha,$ inner field lines and lower $\\gamma$ values, the intensity profile is almost always significantly suppressed on the trailing side, as compared to the leading side. The reason for this is quite obvious. As explained earlier, the $\\rho$ for inner field lines is highly asymmetrical between the leading and trailing sides -- it remains more or less steady on the leading side, while on the trailing side it shoots up to a high value and then falls. Whenever the $\\rho$ shoots up such that $\\rho/\\rho_{\\rm p}\\gg 1$, the spectral intensity is significantly reduced. On the other hand, on the leading side we mostly have $\\rho/\\rho_{\\rm p}\\approx 1$ and hence the spectral intensity is significant there. For relatively lower values of $\\gamma,$ $\\rho_{\\rm p}$ reduces and hence there is a greater chance of having $\\rho/\\rho_{\\rm p}\\gg 1$, while for higher values of $\\gamma$, $\\rho/\\rho_{\\rm p}$ drops down and eventually becomes closer to 1. Hence the intensity plots shown in \\S\\ref{app:figs} are stronger on the leading side at low $\\gamma$, and stronger on the trailing side at high $\\gamma$. However, it is to be noted that (i) the values of $\\gamma$ required to achieve stronger trailing side profiles are very high -- usually significantly more than 1000; whereas, for more typical values of $\\gamma$, we get the stronger leading side profiles and (ii) the intensity contrast obtained for the stronger leading side profiles is much larger and striking, compared to that for the stronger trailing side profiles. Both these facts argue naturally for a strong preponderance of one sided cones with stronger leading side profiles, as is statistically seen in the results of LM88. To further illustrate the idea, we have generated profiles as shown in fig. ~\\ref{PARTIAL} that resemble partial cone profiles by using our simulation technique for specific combinations of parameters, which are listed in Table ~\\ref{tabpar}. The thin line curve represents the un-modulated profile, which is simulated by assuming that the emission is uniform all along the field line, while the thick line shows the final modulated profile. The active region of the final profile is clearly shifted to the leading side, due to the afore said behaviour of $\\rho.$ The suppression of the intensity on the trailing part in comparison to the leading side is seen in all the plots and is most dramatic for the inner field lines, for low values of $\\gamma$, and for large values of $\\alpha$. \\subsection{Studying the mechanisms of emission : future prospects} In this section, discuss some of the future possibilities from the present work. Though we have developed a model under certain specific conditions and demonstrated some useful results from the same, it has significant potential for applicability under diverse circumstances and conditions. The $r$ and $\\rho$ of the emission spot are two of the fundamental ingredients for computing the intensity of emission within any model of radio emission for pulsars. These values, along with other parameters that we have calculated in our method after explicitly taking into account of effects of rotation and geometry, are applicable for any model of radiation that precepts the condition embodied in Eq.(\\ref{eq_los_vel}), i.e. having the radiation beam aligned with the velocity vector and line-of-sight. Thus the present method of computing $r$ and $\\rho$ is well suited for studying curvature radiation models in vacuum approximation. The profiles simulated by these models can be compared with the observed ones to check their veracity. Though we have employed single particle curvature radiation formulation, it is well known that this cannot explain the extremely high luminosities seen in typical pulsar radio emission. Coherent emission from bunches of charged particles have been argued to be necessary for explaining the high luminosities (eg. Ginzburg et~al. 1969, Melrose 92, Melrose 2006). The model of coherent emission constructed by Buschauer and Benford (1976) considered relativistic charge and current perturbations propagating through the bunches with N number of charges, which boosted the emission much above typical $ N^2$ factor. Further, they have shown that the characteristic frequency will be significantly shifted to higher values than the typical $\\approx 1.5 \\gamma^3 c/\\rho.$ However, the extremely short lifetime of these moving sheets of plasma (bunches) made it implausible to radiate, and due to this reason these emission models were almost forgotten. In later years, the possibility of formation of Langmuir micro-structures (solitons) due to the collective behavior of the plasma brought back the possibility of bunched radiation (Asseo 1993). It was shown that the radiation from such a bunch could be expressed by just using the classical formula for curvature radiation (Asseo 1993). Melikidze et~al. (2000) considered the three component structure of charge distribution for solitons in the pulsar magnetosphere and obtained a different spectral intensity distribution from that of the classical formula for curvature radiation. However, Gil et al. (2004) used single charged bunches of charge Q as equivalent to a single particle with the same charge Q, to explore the effects of the surrounding plasma on the curvature emission, and showed that sufficient luminosity could be produced from curvature emission that matches with the observed luminosity of pulsars. As mentioned in the earlier \\S\\ref{sec:comp-intensity} our estimates and results corresponding to altitude, radius of curvature, magnetic azimuth and magnetic colatitude are equally valid for the case of coherent and incoherent emission, as long as the vacuum approximation is invoked. This is because the peak of the emitted beam will be aligned with the direction of velocity for emission from a source moving at ultra-relativistic speeds, de-facto in vacuum approximation. Hence the premise contained in Eq.(\\ref{eq_los_vel}) for the computation of these quantities will remain valid for both of the cases. For the case of a simple model of coherence for a bunch of net charge Q, the spectral intensity profile estimated will be similar to that of the emission from a single particle with charge Q, and likewise the relative intensity will also be the same. Hence the results that we have drawn upon spectral intensity are valid for the simple case of coherent emission too. However, invoking models of coherent radiation with additional features apart from a simple coherent model, may push the intensity estimates to significantly different values and the resulting shape of the intensity profile will be considerably altered. Two such examples are mentioned in the following. Buschauer and Benford (1976) has shown that both intensity profile and characteristic frequency will be altered if the allowance is made for the propagation of a charge and current density wave through the coherent bunch.Considering this model we find that it can alter the shape of the computed spectral intensity curve corresponding to a given field line, from that of the present results. This is mainly because of the reason that characteristic frequency $\\omega_{\\rm c}$ will be shifted to a higher value than in the case of single particle curvature radiation. Another case is the spectral intensity formula for emission from solitons having a three component charge structure (Eq.(12) in Melikhidze et~al. 2000) which also will yield significantly different estiamtes for spectral intensity, from that of the spectral intensity estimated for the single particle emission. Both of these models are treated in the vacuum approximation and hence they satisfy the condition embodied in Eq.(\\ref{eq_los_vel}), i.e. having the radiation beam aligned with the velocity vector and line-of-sight. This ensures that the method of estimation and hence the results corresponding to altitude, radius of curvature, magnetic azimuth, magnetic colatitude etc. will be applicable for these two cases also. The only quantity that is altered by the inclusion of these models, from a single particle case, is the spectral intensity estimate. Nevertheless, these models can be quite easily incorporated into our simulation studies, simply by modifying the form of the spectral intensity expression that is used. In the emission models where the effects of the surrounding plasma are considered, the peak of the radiation beam may be offset from the velocity vector by a finite angle (Gil et~al. 2004). This requires a modification to the condition in Eq.(\\ref{eq_los_vel}) such that ${\\hat \\textbf {n}}\\cdot {\\hat \\textbf {v}} = \\eta_{\\rm max } $ where $ \\eta_{\\rm max } $ is the value of the angle of offset by which the peak of the emission beam is offset from the velocity vector. Coupling this with some modifications to our method can deliver the values of $r$ and $\\rho$ appropriate for this case too. The analysis and results that ensue from all of the above said considerations will be discussed in our forthcoming works.\\\\ We have developed a method to compute the probable locations of emission regions in a pulsar magnetosphere that will be visible at different pulse longitudes of the observed profile. The effects of geometry and rotation of the pulsar are accounted in a detailed manner in this method, which is a very useful new development. Our method includes `exact ' and `approximate' techniques for carrying out the estimation of the relevant emission parameters. The `approximate' method is useful for certain extreme regimes of parameter space, and for faster computation of the results. The misalignment angle, which provides a good check of the accuracy of the computations, shows that our method achieves satisfactory precision. Besides the exact location of possible emission regions, we are able to compute several other useful parameters like the height of emission, and the radius of the particle trajectory at the emission spot, the azimuthal location of the associated field line etc., for different combinations of pulsar parameters like $\\alpha$ and $\\beta.$ Further, using the classical curvature radiation as the basic emission mechanism (which is apt for a debut level analysis), we are able to compute the spectral intensity from any emission spot. By assuming a uniform emission all along the field lines, we have estimated the spectral intensity for a range of pulse phase that the line of sight sweeps through. We have discussed how realistic looking pulsar profiles can be generated from these generalized intensity curves, by assuming specific range of emission heights along specific rings of field lines. We have illustrated the capabilities of these methods by generating simulated profiles for the test case of the pulsar PSR B2111+46, and have shown that fairly good match with observed profiles can be achieved. We have also shown how further detailed (and practical) considerations can help improve this match. We have shown how our results offer a direct and natural explanation for the puzzling phenomena of partial cones that are seen in some pulsar profiles. Our simulations also provide a direct insight into the generation of the core component of pulsar beams. Finally, we have indicated how our method can be extended to incorporate more sophisticated models for the emission mechanism and produce intensity profiles for the same. These, as well as extension to polarized intensity profiles, will be taken up as future extensions of the work reported here." }, "1003/1003.4825_arXiv.txt": { "abstract": "{} { We study the mean profiles of the multi--component pulsars PSRs~B1839+09, B1916+14 and B2111+46. We estimate the emission height of the core components, and hence find the absolute emission altitudes corresponding to the conal components. } { By fitting Gaussians to the emission components, we determine the phase location of the component peaks. Our findings indicate that the emission beams of these pulsars have the nested core--cone structures. Based on the phase location of the component peaks, we estimate the aberration--retardation (A/R) phase shifts in the profiles. Due to the A/R phase shift, the peak of the core component in the intensity profile and the inflection point of the polarization angle swing are found to be symmetrically shifted in the opposite directions with respect to the meridional plane in such a way that the core shifts towards the leading side and the polarization angle inflection point towards the trailing side. } { We have been able to locate the phase location of the meridional plane and to estimate the absolute emission altitude of both the core and the conal components relative to the neutron star center, using the exact expression for the A/R phase shift given by Gangadhara (2005).} {} ", "introduction": "Pulsar radio emission is understood to be emitted by the relativistic plasma accelerated along the dipolar magnetic field lines (e.g., Ruderman \\& Sutherland 1975). Among the various models proposed for pulsar emission, the coherent curvature radiation has turned out to be an effective mechanism for explaining some of the important pulsar radiation properties. The common occurrence of an odd number of components in the mean pulsar profiles has lead to the nomenclature of a nested conal structure for the pulsar emission beam (e.g., Rankin 1983a; Rankin 1993). However, Lyne \\& Manchester (1988) suggested that the emission within the beam is patchy, i.e., the distribution of component locations within the beam is random rather than organized in one or more hollow cones. Also studies by Mitra \\& Deshpande (1999) indicate that the structure of the pulsar emission beam is more likely to be nested hollow cones. Gangadhara \\& Gupta (2001, hereafter GG01), and Gupta \\& Gangadhara (2003, hereafter GG03) showed that the prevalent picture of emission cones axially located around the central core component is a suitable model for explaining the core-cone structure of the pulsar emission beam. A long--standing question in pulsar astronomy has been the location of the radio emission region in the magnetosphere. In the literature, there are mainly two types of methods proposed for estimating the radio emission altitudes: (1) {\\it a purely geometric method,} which assumes the pulse edge is emitted from the last open field lines (e.g., Cordes 1978; Gil \\& Kijak 1993; Kijak \\& Gil 2003), (2)~{\\it a relativistic phase shift method,} which assumes that the asymmetry in the conal components phase location relative to the core is due to the aberration-retardation phase shift (e.g., GG01, Gangadhara 2005, hereafter G05). Both methods have merits and demerits: the first method has an ambiguity in identifying the last open field lines, while the latter is restricted to the profiles in which the core-cone structure can be clearly identified. The emission heights of PSR B0329+54 given in GG01, six other pulsars in GG03 and the revised ones by Dyks, Rudak \\& Harding (2004, hereafter DRH04) are all relative to the emission height of the core, which is assumed to be zero. However, the core emission is believed to originate from lower altitudes than that of the conal components (e.g., Blaskiewicz et~al. 1991; Rankin 1993). Hoensbroech \\& Xilouris (1997) estimated the emission heights at high frequency radio profiles for a set of pulsars. They suggested that the emission heights at high frequency can set an upper limit for the core emission height. By assuming a fixed emission altitude across the pulse, Blaskiewicz et~al.~(1991, hereafter BCW91) presented a relativistic rotating vector model. The results of this purely geometric method are found to be in rough agreement with those of BCW91. However, the relativistic phase shift method clearly indicates that the emission altitude across the pulse window is not constant (GG01; GG03; DRH04; Johnston \\& Weisberg 2006; Krzeszowski et al. 2009). By considering the relativistically beamed radio emission in the direction of the magnetic field line tangents, Gangadhara (2004, hereafter G04) solved the viewing geometry in an inclined and slowly rotating dipole magnetic field. A more exact expression for the relativistic phase shift is given in (G05), which also includes the phase shift due to polar cap currents. In the present work, we analyze the mean profiles of PSRs B1839+09 and B1916+14 at 1418 MHz, and PSR~B2111+46 at 610 MHz and 1408 MHz, to estimate the absolute emission height of the pulse components. In Sect.~2, we give a method for estimating the absolute emission height of pulse components. ", "conclusions": "Based on the A/R method, we estimated the absolute emission height of the core as well as cones in three pulsars: PSRs B1839+09, B1916+14 and B2111+46. Though this method is based on the existing standard models in literature, the combination of the A/R phase shift and the delay-radius relation of BCW91 for estimating the core height is novel. The geometrical method, involving a comparison of the measured pulse widths with geometrical predictions from dipolar models, is believed to yield absolute emission heights. However, the estimation of emission height, using the geometrical method, is based on the assumption that the pulse edges originate from the last open field lines of the polar cap. In general, the edge of the on-pulse region may not originate from the last open field line, and hence the assigning of the edges of the intensity profile to the last open field lines can be misleading. For example, the range of magnetic foot-colatitude for field-lines that are associated with components in PSR B2111+46 are in the range from $S/S_{\\rm lof}\\sim 0.13$ to $0.5,$ whereas the last open field line is at $ 1.$ This means that the boundary of the active region of emission can lie anywhere from $\\approx 0.5$ to $1.$ According to DRH04, the A/R phase shift advances the centroid of the intensity profile to an earlier phase by $\\delta\\phi'_c=2 r_{\\rm lof}/r_{\\rm LC},$ while the PPAIP is delayed to a later phase by $\\delta\\phi'_{\\rm PPAIP}\\sim 2\\,r_{\\rm core}/r_{\\rm LC},$ where $r_{\\rm lof}$ is the emission height from the last open field-line and $r_{\\rm core}$ is the emission height of the core. Then $\\Delta \\phi'= 2 (r_{\\rm lof}+ r_{\\rm core})/ r_{\\rm LC},$ and the emission height $r= r_{\\rm LC}\\,\\,\\Delta \\phi'/4 = (r_{\\rm lof}+ r_{\\rm core})/2 ,$ gives only an average of the emission height for the core and the pulse edge, which is far from the true value. This emission height cannot represent any specific pulse sub-component of the profile, and can be misleading in cases where $r_{\\rm lof}$ and $r_{\\rm core}$ are significantly different. Further more, this will introduce large systematic errors in the emission heights estimated from geometrical methods, due to the aforesaid assumption of identifying the last open field lines with the pulse edges. Rankin (1983a) has argued that the pulsar emission cones are quasi-axial, i.e., the conal components are not exactly axially located with respect to the magnetic axis. Mitra \\& Deshpande (1999) have suggested that the pulsar emission beams are nearly circular in the aligned configuration ($\\alpha\\sim 0^\\circ$) and change to elliptical in the orthogonal configuration $(\\alpha\\sim90^\\circ)$. The majority of the pulsar observations indicate that the beam geometry is likely to be nested cones, distributed in a nearly non-coaxial fashion about the magnetic axis. A likely case is that the cones, which are coaxial in the co-rotating frame, will appear non-coaxial in the laboratory frame because of the A/R phase shifts (GG01; GG03). In the works GG01 and GG03, the emission height of the core was neglected by assuming that it is considerably smaller than that of the components. However, we find that the emission height of the core is quite significant and cannot be neglected in comparison to the emission height of the components. We identify the meridional plane M as being located at the mid point between the centroid of the intensity profile and the PPAIP, owing to the A/R effects. By recognizing this, we were able to estimate the absolute emission heights of both the core and the conal components. As mentioned before, we restricted the region of the fit of the BCW91 (relativistic RVM) curve to the section of the PPA data falling within the FWHM of the core component for estimating the core emission height, and the justification for doing so is given now. The expression for the BCW91 was derived by assuming that the emission altitude across the active region of the pulse profile is a constant. Thus in a BCW91 fitting, a single $r$ value was taken to characterize the emission height of the full region of the PPA data. But later observational results (e.g., GG01, GG03) established that the emission altitude corresponding to the subpulse components in multi-component profiles spans over a large range of emission heights. This elicits the fact that in multi-component profiles the $r,$ found by fitting the BCW91 curve to the full PPA region of the active profile, might give an emission altitude that can be significantly different from those obtained from the A/R method for the subpulse components. The best-fit value of $r$ in the BCW91 model, which is the weighted average of $r_{\\rm i}$ that characterize the emission height at each point of pulse phase, is given by Eq.~(\\ref{eq_rbyrL2}) in \\ref{subsec:BCW}. Hence a single value of $r$ found from the BCW91 fit cannot be closer to the true emission height corresponding to the core or cone peak if $r_{\\rm i}$ varies significantly within the region of the fit. For example, one can compare the emission altitudes in our Tables \\ref{tabcore} and \\ref{tabcomp} with those given in Table~3 of ML04 for PSR B2111+46. It can be surmised that a single value of $r$ cannot characterize the emission altitude across the entire active region of a multi-component profile. We can think of two viable alternatives in this scenario: either (1) adapt or modify the BCW91 formulation for a variable emission altitude $r$ (Dyks 2008) or (2) fit the BCW91 curve for regions of the PPA profile having a relatively constant value of $r.$ We prefer the latter alternative to evade the modification of the theory behind the BCW91 model. Here we note that \\emph{$\\alpha$ and $\\beta$ are not invoked as fit parameters}; instead we used their published values in Eq.~(\\ref{eq_{BCW}}). This is expected to further reduce the ambiguity of the fit results and to aid in counteracting an obvious disadvantage in this `restricted' fit method, i.e., having a reduced number of fitted PPA data points than a fit for the `full range' of PPA data. It is remarkable that some of the fit statistics (e.g., reduced $\\chi^2$) given in Table~\\ref{tabcore} reveal that the present method of fitting is comparable (in a few cases even better) to the existing ones in the literature (e.g., compare the $\\chi^2$ value given in Table~\\ref{tabcore} with Table~2 of ML04 for PSR B2111+46). We estimated the standardized residuals (SR) and found the percentage of SR that falls within -2 and +2 as given in Col.~(8) of Table~\\ref{tabcore}. As it is known, a good fit is expected to have a threshold 95 \\% of the standardized residuals to fall within -2 and +2. Several previous works found that $\\alpha$ and $\\beta$ are highly covariant in PPA fits (e.g., Everette \\& Weisberg 2001). But this covariance of $\\alpha$ and $\\beta$ with the $r$ parameter was not mentioned by any of them. The fit statistics do not reveal any significant covariance of $\\alpha$ and $\\beta$ with $r.$ This gives us a further clue for finding the $r$ parameter without invoking a concurrent fit for $\\alpha$ and $\\beta$ (see \\ref{app:Psi} for the fitting procedure). Owing to the extreme difficulties encountered in determining $\\alpha $ and $\\beta$ through RVM fitting, a larger range of PPA data have always been preferred for a better fit (e.g., Everette \\& Weisberg 2001). The justification for doing so is that $\\alpha $ and $\\beta$ must remain constant throughout the entire PPA profile. But in the present scenario, as described earlier, the selection of a large range of PPA data for fitting does not always translate into a better estimation of $r$ because of the variation of the emission height with pulse phase. So, owing to all of the above said reasons we restrict the fit of the BCW91 curve to the PPA profile, falling around the core component, which is expected to yield an emission altitude characterizing the core height. A section of inner cones often lapses over the core as is seen in the Gaussian fits (panel (a) of Figs.~\\ref{freq_1839}--\\ref{freq_1408}) of the total intensity profiles. The inner cones may contribute to the core polarization near the edges of pulse phase of the FWHM region that we bracketed. Hence the PPA corresponding to the bracketed region will be `contaminated ' by the adjacent conals, and this has to be accounted for. We estimated and accounted for the error induced because of this effect in the estimation of the core emission heights (see \\ref{app:Psi1} and \\ref{subsec:BCW}). The possibility that the A/R phase shift may be reduced by the rotational distortion of the magnetic field line due to a sweep-back of the vacuum dipole magnetic field lines has to be considered. The sweep-back of dipole magnetic field lines was first treated in detail by Shitov (1983). Further, Dyks \\& Harding (2004) investigated the rotational distortion of pulsar magnetic field by making the approximation of a vacuum magnetosphere. For $\\phi'=30^\\circ,$ $\\beta=-1.6^\\circ$ and $\\alpha = 14^\\circ$ we computed the phase shift $\\delta\\phi'_{\\rm mfsb}$ due to the magnetic field sweep-back (Dyks \\& Harding 2004; also see Eq. (49) in G05). It is found to be $<0.0001$~rad for $r/r_{\\rm LC}\\leq 0.06, $ which is much smaller than the aberration, retardation and polar cap current phase shifts in PSR~B2111+46. Hence we neglect the magnetic field sweep-back effect. The field-aligned polar-cap current does not introduce any significant phase shift into the phase of the PPAIP. But it introduces a positive offset into the PPA, though it roughly cancels due to the negative offset by aberration (Hibschman \\& Arons 2001). The phase shift of pulse components due to the polar cap current was estimated recently by G05, and found to be quite small compared to the A/R phase shift." }, "1003/1003.1280_arXiv.txt": { "abstract": "{The location of B supergiants in the Hertzsprung-Russell diagram (HRD) represents a long-standing problem in massive star evolution. Here we propose their nature may be revealed utilising their rotational properties, and we highlight a steep drop in massive star rotation rates at an effective temperature of 22\\,000 K. We discuss two potential explanations for it. On the one hand, the feature might be due to the end of the main sequence, which could potentially constrain the core overshooting parameter. On the other hand, the feature might be the result of enhanced mass loss at the predicted location of the bi-stability jump. We term this effect ``bi-stability braking'' and discuss its potential consequences for the evolution of massive stars. ", "introduction": "\\label{s_intro} The large number of B supergiants as well as their location in the Hertzsprung-Russell diagram represents a long-standing problem in massive star evolution (e.g. Fitzpatrick \\& Garmany 1990). Even the most basic question of whether B supergiants are core hydrogen (H) burning main sequence (MS) or helium burning objects has yet to be answered. Here we propose their nature may be revealed utilising their rotational properties. On the MS, O-type stars are the most rapid rotators known (with $\\varv$$\\sin$$i$ up to 400 \\kms), but B supergiants rotate much more slowly (with $\\varv$$\\sin$$i$ $\\la$ 50\\kms), which has been attributed to the expansion of the star after leaving the MS. Hunter et al. (2008) noted a steep drop in rotation rates at low gravities (log $g$ $<$ 3.2) and suggested the slowly rotating B supergiants to be post-MS. The steep drop was also used to constrain the core overshooting parameter $\\alpha_{\\rm ov}$ in massive star models (Brott et al. 2010). The slowly rotating B supergiants are also cooler (with $\\teff$ below $\\sim$22\\,000 K) and $\\varv$$\\sin$$i$ is observed to drop steeply below this $\\teff$. Here we introduce an alternative explanation for the slow rotation of B supergiants: wind-induced braking due to bi-stability, or bi-stability braking (BSB). Mass loss plays a crucial role in the evolution of massive stars. Whilst a large amount of attention has been directed towards the role of stellar winds in terms of the {\\it loss of mass}, as winds ``peel off'' the star's outer layers (Conti 1976), much less effort has been dedicated to understanding the associated {\\it loss of angular momentum} (but see Langer 1998, Meynet \\& Maeder 2003). Yet the angular momentum aspect of these winds may be equally relevant for understanding massive stars as the loss of mass itself, possibly in a mass range as low as $\\sim$10-15\\msun. We first recapture the physics of bi-stable winds and BSB (Sect.~\\ref{s_bsb}), before presenting the current knowledge of rotational velocities of massive stars. We note a steep drop at $\\sim$22\\,000 K (Sect.~\\ref{s_hook}) and propose two possible explanations for it. In the first one, the drop is due to the separation of MS objects from a second population of slow rotators (Sect.~\\ref{s_eoms}), whilst in the second one the slow rotation is the result of BSB (Sect.~\\ref{s_cbsb}). ", "conclusions": "\\label{s_disc} In principle it is possible that both effects of ``two populations'' and BSB occur simultaneously at 22\\,000 K, with BSB occurring above a certain critical mass, and the ``two population'' scenario taking over in the lower mass (10-20 \\msun) range, but this situation might appear somewhat contrived. The strongest argument for the ``two population scenario'' are the large N abundances of the B supergiants, whilst the strongest argument for BSB is that the drop is observed at the correct location (whilst no such coincidence would be expected for the alternative interpretation). Using our standard models, BSB can only operate above a certain critical mass and would not be able to explain the steep drop in rotational velocities of stars below the critical mass. The reason BSB does not operate at lower masses in our standard models (of Fig.\\,\\ref{VROT_GG}) is that the drop feature has been used to constrain the core overshooting parameter of $\\alpha_{\\rm ov}$ = 0.335. The applicability of BSB could be pushed to lower masses if the MS lifetime were extended. This could be achieved by increasing $\\alpha_{\\rm ov}$. When we enlarge $\\alpha_{\\rm ov}$ to 0.5, BSB also occurs at 20\\msun\\ for our Galactic and LMC models. What is clear is that the critical mass is model-dependent. For instance, the solar-metallicity models of Meynet \\& Maeder (2003) show BSB in the lower ($\\sim$15-20\\msun) range. We point out that if BSB were the correct explanation for the drop feature all the way down to $\\sim$10$\\msun$, we would require a very large core overshooting parameter, and the consequences would be far-reaching. For instance, it would imply that B (and even A) supergiants are MS objects burning H in their cores. This would potentially solve the long-standing problem of the presence of such a large number of B supergiants. Moreover, if BSB could work for the entire mass range, it would also have profound implications for the Blue to Red (B/R) supergiant ratio that has been used to constrain massive star models as a function of metallicity for decades. Furthermore, if the absence of rapidly rotating B supergiants is due to BSB, one might wonder what this would imply for the evolutionary state of the presumably rapidly rotating B[e] supergiants. The rapid rotation of these extreme objects could possibly be related to close binary evolution or merging (Pasquali et al. 2000), but this requires future investigation. If BSB would indeed occur in the lower mass range (down to $\\sim$10$\\msun$), one should be aware that the derived overshooting parameter of 0.335 becomes a lower limit and that the real value becomes larger. Although this would be consistent with the suggested increase in $\\alpha_{\\rm ov}$ with stellar mass (Ribas et al. 2000), such a large value of $\\alpha_{\\rm ov}$ might be considered uncomfortable, as the highest mass data-point in Ribas et al. is based on one binary star, V380 Cyg, for which the results have been challenged (Claret 2003). To summarise, we have presented two potential explanations for the steep drop in rotation rates at 22\\,000 K. Currently, we have insufficient information to decide which one is correct. In any case, our study demonstrates the important role of mass loss for massive star evolution, and especially the importance of {\\it specifics} in its dependence on the stellar parameters. Furthermore, we have highlighted the significant influence of mass loss on the angular momentum transport in massive stars. Last but not least, BSB may offer a novel method of diagnosing the effects of mass loss via its influence on the angular momentum. Current analyses yield controversial results with respect to the existence of a BS jump. On the one hand, the predicted drop in terminal wind velocity across the BS range has been confirmed (Crowther et al. 2006). On the other hand, for temperatures below the BS-Jump, the mass-loss rates obtained from spectral modelling are generally much lower than predicted (Vink et al. 2000, Crowther et al. 2006). A simultaneous investigation of the abundances, mass loss, and rotational properties of a large sample of massive stars, e.g. with the {\\sc flames ii} Tarantula survey (Evans et al. 2009), would be most helpful to settle these issues." }, "1003/1003.3285_arXiv.txt": { "abstract": "We present the results of timing analysis of the low-frequency Quasi-Periodic Oscillation (QPO) in the Rossi X-Ray Timing Explorer data of the black hole binary XTE J1550--564 during its 1998 outburst. The QPO frequency is observed to vary on timescales between $\\sim$100 s and days, correlated with the count rate contribution from the optically thick accretion disk: we studied this correlation and discuss its influence on the QPO width. In all observations, the quality factors ($\\nu_0$/FWHM) of the fundamental and second harmonic peaks were observed to be consistent, suggesting that the quasi-periodic nature of the oscillation is due to frequency modulation. In addition to the QPO and its harmonic peaks, a new 1.5$\\nu$ component was detected in the power spectra. This component is broad, with a quality factor of $\\sim$0.6. From this, we argue what the peak observed at half the QPO frequency, usually referred to as ``sub-harmonic'' could be the fundamental frequency, leading to the sequence 1:2:3:4. We also studied the energy dependence of the timing features and conclude that the two continuum components observed in the power spectrum, although both more intense at high energies, show a different dependence on energy. At low energies, the lowest-frequency component dominates, while at high energies the higher-frequency one has a higher fractional rms. An interplay between these two components was also observed as a function of their characteristic frequency. In this source, the transition between low/hard state and hard-intermediate state appears to be a smooth process. ", "introduction": "Since the launch of \\emph{Rossi X-ray Timing Explorer} (RXTE), an extraordinary progress has been achieved in the knowledge of the variability properties of black-hole candidates (BHCs) in X-ray binaries. Different types of Quasi-Periodic Oscillations (QPO) have been observed in these systems. While only a few binaries show high-frequency QPOs (HFQPOs, 50--450 Hz, see \\citealt{rem06,bel06}), low-frequency QPOs (LFQPOs, mHz to $\\sim$10 Hz, see \\citealt{rem06,cas05}) are detected in virtually all observed BHCs. As both low and high-frequency QPOs are thought to arise in the accretion flow close to the black hole, the study of their properties and behavior can provide important clues on the physics of accretion onto BHCs. In the case of LFQPOs, several distinct types showing different properties have been observed. Three main types, dubbed types A, B, and C respectively, stand out in the present scenario. \\citet{wij99} and \\citet{hom01} first reported type A and type B LFQPO in XTE J1550--564, while \\citet{rem02} dubbed the ubiquitous LFQPO that appears together with band-limited noise type-C. The detailed properties of the different types of LFQPOs were investigated by \\citet{cas05}. However, the physical difference among types remains unknown. The X-ray transient XTE J1550--564 was discovered on 1998 September 7th \\citep{smi98} with the \\emph{RXTE} All Sky Monitor \\citep{woo99}. The discovery prompted a follow-up series of almost daily pointed RXTE/PCA (Proportional Counter Array; \\citealt{jah96}) observations, which revealed a hard power-law dominated spectrum. Two weeks later, it reached a peak intensity of 6.8 Crab at 2--10 keV. The marked softening of the spectrum during this period indicates a transition from low/hard state (LS) to hard-intermediate state (HIMS), reaching the soft-intermediate state (SIMS) in occasion of the 6.8 Crab peak (see \\citealt{hom05,bel09}). After the bright peak, XTE J1550--564 remained in the HIMS for more than three weeks. Strong LFQPOs were observed, with frequency changing quite dramatically during the first $\\sim40$ d of the outburst but then stabilizing from day 40--52. The type of QPO also changes at day 40 (from C/C' to B: see \\citealt{rem02}). Additional outbursts of the systems followed in 2000, 2002 and 2003. The optical \\citep{ors98} and radio \\citep{cam98} counterparts were identified shortly after the discovery of the source. Subsequent optical observations showed that the dynamical mass of the compact object is $10.5\\ \\pm\\ 1.0\\ M_\\sun$, indicating a black-hole nature. Its binary companion was found to be a low-mass star, and the distance to the source was estimated to be about 5.3 kpc \\citep{oro02}. \\citet{cor02} discovered a large-scale, relativistically moving and decelerating jet emitting in radio and X-rays. In this paper, we concentrate on the power density spectra found at the beginning of the 1998 outburst of XTE J1550--564 and analyze in detail the type-C QPO and noise components, focussing on their relative properties. Particular attention is given to the ``harmonic'' peaks, which reveal information about the nature of the observed signal. ", "conclusions": "We have analyzed 47 \\emph{RXTE} observations with type C LFQPO acquired during the first half of XTE J1550--564 1998 outburst. Satisfactory fits to the power spectra of the 6th--47th observations were obtained with a model consisting of flat-top noise $L_{ft}$, peaked noise $L_{pn}$, fundamental QPO $L_F$, sub-harmonic QPO $L_s$, second harmonic QPO $L_h$ and sometimes a third harmonic QPO components. We identify the peaked noise $L_{pn}$ as a new harmonic component at 1.5$\\nu$ the fundamental. This suggests that what we called previously sub-harmonic quasi-periodic oscillation (QPO) may actually be the fundamental. A similar Q-factor between the fundamental $L_F$ and the harmonic $L_h$ was observed, suggesting a frequency modulation as cause of their width, whereas the sub-harmonic $L_s$ is broader, and the $1.5\\nu_{0}$ feature $L_{pn}$ is even broader. We also found a significant interplay among both QPOs and broad-band components with both frequency and energy. As the disk flux increases, $L_F$ frequency increases, $L_{ft}$ and $L_s$ ratio increase, while $L_{pn}$ and $L_h$ ratio decrease. A similar interplay also takes place as the energy increase. The nature of the interplay remains unknown, since we didn't even know the nature of the components. However, from the interplay we may explain the PDS difference in the first 5 observations, and suggest the transition from LS to HIMS is a smooth process." }, "1003/1003.2001.txt": { "abstract": "We present the result of a study of the X-ray emission from the Galactic Centre (GC) Molecular Clouds (MC) within 15 arcmin from Sgr A*. We use \\xmm\\ data (about 1.2 Ms of observation time) spanning about 8 years. The MC spectra show all the features characteristic of reflection: i) intense Fe K$\\alpha$, with EW of about 0.7-1 keV, and the associated K$\\beta$ line; ii) flat power law continuum and iii) a significant Fe K edge ($\\tau\\sim0.1-0.3$). The diffuse low ionisation Fe K emission follows the MC distribution, nevertheless not all MC are Fe K emitters. The long baseline monitoring allows the characterisation of the temporal evolution of the MC emission. A complex pattern of variations is shown by the different MC, with some having constant Fe K emission, some increasing and some decreasing. In particular, we observe an apparent super-luminal motion of a light front illuminating a Molecular nebula. This might be due to a source outside the MC (such as Sgr A* or a bright and long outburst of a X-ray binary), while it cannot be due to low energy cosmic rays or a source located inside the cloud. We also observe a decrease of the X-ray emission from G0.11-0.11, behaviour similar to the one of Sgr B2. The line intensities, clouds dimensions, columns densities and positions with respect to Sgr A*, are consistent with being produced by the same Sgr A* flare. The required high luminosity (about 1.5$\\times$10$^{39}$ erg s$^{-1}$) can hardly be produced by a binary system, while it is in agreement with a flare of Sgr A* fading about 100 years ago. The low intensity of the Fe K emission coming from the 50 and the 20 km \\s\\ MC places an upper limit of 10$^{36}$ erg \\s\\ to the mean luminosity of Sgr A* in the last 60-90 years. The Fe K emission and variations from these MC might have been produced by a single flare of Sgr A*. ", "introduction": "Sgr A*, the supermassive black hole (BH) at the center of the Milky Way, now radiates at a rate about 8 orders of magnitude lower than the Eddington luminosity for its estimated mass of M$_{BH}\\sim$4~$\\times10^6$ M$_{\\odot}$ (Sch\\\"odel et al. 2002; Eisenhauer et al. 2003; Ghez et al. 2003; 2005; Gillesen et al. 2009). Such a low luminosity is difficult to reconcile with the dense environment that is present in the center of the Galaxy and has motivated the development of several radiatively inefficient accretion/ejection models (Melia \\& Falcke 2001). Although Sgr A* is known to display flares in X-rays (Baganoff et al. 2001; Goldwurm et al. 2003) and near-infrared (Genzel et al. 2003; Ghez et al. 2004), during which the X-ray intensity increases by factors up to 160 (Porquet et al. 2003) from the quiescent value, the bolometric luminosity still remains extremely low during these events compared to the Eddington one or even to the accretion power expected from the capture of stellar wind material from the nearby stars. On the other hand, one may wonder whether Sgr A* has always been so underluminous or if it experienced, in the past, long periods of high energy activity, that would make the massive black hole of our Galaxy more similar, than appears today, to typical low-luminosity Active Galactic Nuclei. Indication of Sgr A* past activity can be sought in the interstellar medium surrounding the black hole. Sunyaev et al. (1993) were the first to interpret the X-ray emission, seen with GRANAT to roughly follow the distribution of the Molecular Clouds (MC) of the region, as scattering by the molecular material of emission from a past outburst of Sgr A* and predicted, at that time, a correlation of the X-ray fluorescent line of neutral iron with the MC. Koyama et al. (1996) with ASCA and Murakami et al. (2001b) with Chandra did in fact find such a correlation, particularly evident with the most massive MC complex of the region, Sgr B2, and proposed, using parameters derived from this cloud, that Sgr A* underwent, about 300 years ago, a major outburst of X-ray emission, with a luminosity of the order of few 10$^{39}$ ergs s$^{-1}$. The fluorescence line at 6.4 keV (K$\\alpha$) is produced by the extraction of an electron from the inner shell (K) of neutral or low-ionized iron atoms and the following electron transition from the second shell (L). Such line (actually a close doublet) is generally associated with another line (K$\\beta$) due to the transition from the upper (M) level. Collisionally-ionized iron atoms in a hot plasma preferentially produce lines in the 6.5-6.9 keV range, associated with a plasma continuum spectrum. Thus, the origin of the 6.4 keV line is most probably associated with either a large irradiation by photons having energies higher than 7.1 keV or by energetic particles, most probably electrons. Diffuse X-ray (2-10 keV) emission in the galactic center region is complex and still under intense investigation (Park et al. 2004; Goldwurm 2008; Koyama et al. 2009) but it certainly consists of at least the following components: a uniformly distributed soft emission well described by a low temperature ($\\approx$1 keV) plasma, a little less uniform but centrally peaked 6.7 keV line associated with continuum emission described by a hot (kT $\\approx 7$ keV) plasma model, and a clumpy 6.4 keV iron line component well correlated with molecular material. The soft component can be fully explained by the supernova (SN) activity of the region, while the origin of the other components is more uncertain. The interpretation of hot plasma emission for the 6.7 keV line and associated hard component is problematic because such plasma cannot be confined in the region and its regeneration would require a too large amount of energy (but see Belmont \\& Tagger (2006) for a heating mechanism of a helium dominated hot plasma at the galactic center). An alternative is that the hot component may contain an important contribution from faint sources. This interpretation is supported by the similarity in the X-ray and near-infrared surface brightness distribution (Revnivtsev et al. 2006b) and by the large fraction of weak point sources in the population that the Chandra deep survey of the galactic center unveiled (Muno et al. 2004, Revnivtsev et al. 2007). For the 6.4 keV component, two models of emission are competing: the reflection model quoted above (Sunyaev \\& Churazov 1998), and one that attributes the 6.4 keV emission to low energy particles, most probably electrons (Valinia et al. 2000; Yusef-Zadeh et al 2002; 2007). In the frame of the latter model, the impact by low energy protons (Dogiel et al. 2009) or by fast moving SN ejecta (Bykov 2003) with neutral material have also been considered. This emission line has been detected in other, but not all, molecular clouds of the central region (Murakami et al. 2001a; Yusef-Zadeh et al. 2002; Nakajima et al. 2009), for which a reflection from single Sgr A* event interpretation is more problematic and for some of which has also been found a correlation with non-thermal radio filaments, indicating local particle acceleration (Yusef-Zadeh et al. 2007). On the other hand, the detection of hard X-ray emission up to 100 keV from Sgr B2 obtained with the Integral observatory by Revnivtsev et al. (2004) supports the reflection nebula interpretation. These authors also demonstrated that the emission line intensity was constant until about 2000 and therefore that the original outburst must have lasted at least 10 years. However, the most convincing evidence that supports the photon-ionisation model from an external source, comes now from the recent detections of variability of the line and continuum emission, a signature predicted and modelled in detail by Sunyaev \\& Churazov (1998). Up to now two different claims of variability detection have been published. Muno et al. (2007) observed variation in the continuum (not in the line emission) flux and morphology from two 6.4 keV nebulae at about 6 arcmin from Sgr A*. Then using data from several satellites (ASCA, \\chandra, \\xmm\\ and Suzaku), Koyama et al. (2008) and Inui et al. (2009) showed that the Sgr B2 6.4 keV line emission is changing in a way that it would be produced by a wave front passing through the different components of the Sgr B2 complex. Maybe the more compelling evidence, up to now, is the discovery of time evolution of the hard X-ray emission from Sgr B2 observed over 7 years by the same instrument on the Integral satellite (Terrier et al. 2010). Indeed the evolution of the hard X-ray emission of Sgr B2 is best explained by an X-ray reflection nebula scenario in which the fading of the reflection component (for the first time measured in the Compton hump region) is due to the propagation of the decay part of the outburst and rules out competing models based on irradiation by low energy cosmic ray electrons. However, Sgr B2 is quite a special object and one may wonder whether the reflection nebula model works only for this cloud or if it holds for the other 6.4 keV features of the region. If the origin of the scattered emission is external to the cloud, signatures should be seen elsewhere. Sgr B2 is indeed not the only MC close to Sgr A*. The supermassive BH sits on the middle of the Central Molecular Zone (CMZ; Morris \\& Serabyn 1996), a condensation of MCs right in the center of the Galaxy. The detailed study of this region can therefore validate the reflection model and even narrate the past history of Sgr A* emission (Sunyaev et al. 1993; 1998; Cramphorn et al. 2002). Moreover, the light front due to a major flare may be used as a tool to scan the distribution of the MC material in the CMZ. The main focus of this work is the study of the X-ray emission from the molecular clouds located around Sgr A* using the 8-year Sgr A* monitoring program carried out with the \\xmm\\ satellite. In section 2 the different observations and data reduction are presented. Section 3 shows the GC images in the Fe K$\\alpha$ band and the analysis of the CS maps, with the aim to correlate the Fe K emission with the MC disposition, so that we can infer the location/distribution and column density of the different MC within the \\xmm\\ field of view. Section 4 presents the mean spectra from the different regions selected through the CS maps. Section 5 shows the time evolution of the Fe K emission from the MC. Section 6 shows the discovery of a super-luminal echo in a X-ray reflecting nebula. In Section 7 the results are discussed, in particular in subsection 7.2 the possibility that all MC are illuminated by a single flare from Sgr A* is discussed. The conclusions are summarised in Section 8. ", "conclusions": " \\begin{itemize} \\item{} The Fe K$\\alpha$ emission is asymmetrically distributed and is concentrated in features that are correlated to some high density molecular clouds of the CMZ. These X-ray Fe K bright MC show all the features of a reflection spectrum: neutral Fe K lines with EW of about 0.7-1 keV, Fe K edges ($\\tau\\sim0.1-0.3$) and flat continua ($\\Gamma\\sim$0.6-1.7). \\item{} We discovered an apparent super-luminal motion corresponding to a MC called \"the bridge\". The illumination front spans at least 15 light years in a time-scale of about 2-4 years. This phenomenon cannot be due to a source internal to the molecular cloud or to propagation of low-energy cosmic rays. A series of events is also unlikely because a similar pattern of variability is observed in causally disconnected regions. The most probable cause of this variation is the illumination of the molecular material from an external source. One possible geometry is that the illuminating fronts are nearly parallel to the bridge. To produce an apparent super-luminal motion and the observed pattern of variations, the X-ray radiation must originate far away from the reflector (more than about 18 pc). The source must thus have had a luminosity higher than about $1.3\\times10^{38}$ erg \\s\\ for several years. A prediction of this picture is that the variations observed in the bridge region 1 propagate to region 2 and then 3 (maybe bridge region 4). Assuming that Sgr A* is the primary source and that it had a luminosity of about 10$^{39}$ erg \\s, it would place the bridge about 60 pc behind Sgr A*. This would imply either a long (about 400 years) or, to take intermittency to the extreme, it could imply that there have been multiple (i.e., two) brief events of activity of Sgr A* in the past. \\item{} We observe that the Fe K emission from G0.11-0.11 is decreasing, as if the MC is responding to the same flare that is illuminating Sgr B2, whose emission is also fading (Inui et al. 2009; Terrier et al. 2010). Considering the sizes and column density of the MC we estimate for G0.11-0.11 a position of about 17 pc behind Sgr A*, from the intensity of the Fe K line only. Surprisingly, in that position, G0.11-0.11 satisfies another completely independent constraint. In fact, it is consistent with being illuminated by the same light front hitting Sgr B2. These two MC are, thus, experiencing the decay of a period of activity of Sgr A* (activity characterised by a luminosity of about $1.4\\times10^{39}$ erg \\s) that lasted for at least 10 years (Revnivtsev et al. 2004) and that ended about 100 years ago (between 70 and 150 years ago). Other possibilities for the Fe K emission are not excluded, but in those scenarios the observed properties would result from pure chance. \\item{} The upper limit on the neutral Fe K emission from the 20 and 50 km \\s\\ molecular clouds allow us to further constrain the mean level of activity of Sgr A* in the recent past, being lower than 8$\\times10^{35}$ erg \\s\\ in the last 60-90 years. The Fe K emission from Sgr B2, G0.11-0.11 and the 50 km \\s\\ cloud suggest that the emission (presumably induced by Sgr A*) from Sgr B2 and G0.11-0.11 is destined to switch off in the next decades. At that point the emission produced by low energy cosmic rays, if present, might become dominant. \\item{} We also observe two bright molecular clouds with intense but stable Fe K emission. Arbitrarily assuming that they reflect a flare of 10$^{39}$ erg \\s, they would be placed at about 35 and 40 pc behind Sgr A*. \\item{} The emission and variations of the MC in the central part of the Galaxy fit the scenario of a past period of activity of Sgr A* (either continuous or with bursts of emission). This period might have started a few hundreds of years ago and lasted until about 70-150 years ago. Since then, Sgr A* experienced a long period of low activity until now. This idea is in agreement with the observation of many Fe-K-bright MC as well as many X-ray weak MC and with the (poorly) known 3-d distribution of MC within the CMZ. \\item{} The analysis of the long \\xmm\\ monitoring of the Galactic Centre shows that, in the past, Sgr A* had activity levels of the order of $\\eta=10^{-5}$ of its Eddington luminosity (as compared to its current $\\eta=10^{-8}$). In those conditions it appeared more similar to the classical low luminosity AGN and the Galactic Black Holes in the quiescent state. A deeper study of that period is important for the comprehension of the accretion mechanism, in view of the unification schemes. \\end{itemize} Continuous deep X-ray monitoring of the region, as well as an improvement of the molecular maps of the CMZ are mandatory to fully understand the origin of the X-ray emission from the MC, in particular, to disentangle the contribution to the observed variations to be ascribed to the supermassive black hole of the galactic center and to X-ray binaries of the region. The connections between the X-ray Fe K maps, their variations and the measurements of the column densities and positions of the MC within the CMZ will allow to precisely reconstruct the history of the emission of Sgr A* and of the bright X-ray sources inside the CMZ. The flares from these sources will provide a tool to scan the CMZ and reconstruct the small scale distribution of MC." }, "1003/1003.0623_arXiv.txt": { "abstract": "{} {Radio galaxies with a projected linear size $\\gtrsim$\\,1 Mpc are classified as giant radio sources. According to the current interpretation these are old sources which have evolved in a low-density ambient medium. Since radiative losses are negligible at low frequency, extending spectral ageing studies in this frequency range will allow to determine the zero-age electron spectrum injected and then to improve the estimate of the synchrotron age of the source. } {We present Very Large Array images at 74 MHz and 327 MHz of two giant radio sources: 3C35 and 3C223. We performed a spectral study using 74, 327, 608 and 1400 GHz images. The spectral shape is estimated in different positions along the source. } {The radio spectrum follows a power-law in the hot-spots, while in the inner region of the lobe the shape of the spectrum shows a curvature at high frequencies. This steepening is in agreement with synchrotron aging of the emitting relativistic electrons. In order to estimate the synchrotron age of the sources, the spectra have been fitted with a synchrotron model of emission. Using the models, we find that 3C35 is an old source of 143$\\pm$20\\,Myr, while 3C223 is a younger source of 72$\\pm$4\\,Myr.} {} ", "introduction": "\\begin{table*} \\caption{Summary of radio observations and images. Col. 1: source name; Col. 2, 3: radio pointing position; Col. 4: observing frequency; Col. 5: bandwidth; Col. 6: date of observations; Col. 7: time of integration; Col. 8: VLA configuration; Col. 9: resolution; Col. 10: Position Angle; Col. 11:rms.} \\label{obs} \\begin{tabular}{ccccccccccc} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Name &$\\alpha$(J2000) & $\\delta$(J2000) &${\\nu}$& $\\Delta\\nu$ & Date & Duration &Array& HPBW&PA&rms\\\\ & ($h$ $m$ $s$) & (\\degr\\ \\arcmin\\ \\arcsec)& MHz& MHz& &hours&&\\arcsec$\\times$\\arcsec&\\degr&mJy/beam\\\\ \\hline 3C35 & 01 12 02.20 & +49 28 35.00 & 73.8 & 1.562 &23-11-2003&6&B&93$\\times$64&-85&95 \\\\ & & &327.5&3.125&23-11-2003&6&B&23$\\times$17&-85&1.3 \\\\ & & &327.5-321.5&3.125&21-03-2004&3.5&C&55$\\times$50&15&2.3\\\\ & & &327.4&3.125&&9.5&B+C&27$\\times$21&-88&1.0\\\\ \\hline 3C223 &09 39 52.74&+35 53 58.20& 73.8 & 1.562 &16-11-2004&5&A&25$\\times$24&-58&43\\\\ &&& 73.8 & 1.562 &03-03-2005&5&B&83$\\times$73&-62&98\\\\ &&&73.8& 1.562 &&10&A+B&26$\\times$25&-61&40\\\\ &&&327.3&6.25&16-12-2004&5&A&6$\\times$5&-77&0.7\\\\ &&&328.9&6.25&03-03-2005&5&B&19$\\times$16&-82&1.3\\\\ &&&327.3&6.25&&10&A+B&7$\\times$6&-78&0.6\\\\ \\hline \\noalign{\\smallskip} \\end{tabular} \\end{table*} \\begin{table*} \\caption{Sources properties.} \\begin{center} \\begin{tabular}{cccccccc} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Name & $\\alpha$(J2000) & $\\delta$(J2000) & z & kpc$\\slash\\arcsec$&LAS& LLS & L$_{178\\,MHz}$ \\\\ & ($h$ $m$ $s$) & (\\degr\\ \\arcmin\\ \\arcsec) & &&\\arcmin &kpc& W\\,Hz$^{-1}$ \\\\ \\hline 3C35 &01 12 02.23 &$+$49 28 35.2 & 0.0673 &1.273&12.5&950& 10$^{26.09}$\\\\ 3C223 &09 39 52.74 &$+$35 53 58.2 &0.1368 &2.393&5.4 &780&10$^{26.89}$\\\\ \\hline \\noalign{\\smallskip} \\label{source} \\end{tabular} \\end{center} \\begin{list}{}{} \\item[] Col. 1: source name; Col. 2: and 3: source coordinates from NASA/IPAC extragalactic database (NED) ; Col. 4: redshift by \\cite{Burbidge1972} and \\cite{SDSS}; Col. 5: arcsec to kpc conversion; Col. 6: largest angular size; Col. 7: largest linear size; Col. 8: radio luminosity at 178 MHz \\cite[]{Laing1980}. \\end{list} \\end{table*} \\begin{figure*}[th] \\begin{center} \\includegraphics[width=14cm]{fig1.pdf} \\end{center} \\caption[]{Radio images of 3C35, all contours start at (3$\\sigma$) and are scaled by $\\sqrt{2}$. {\\it Left}: VLA image at 74 MHz; the resolution is 93\\arcsec$\\times$64\\arcsec\\, with a PA=$-$85\\degr, and the first two levels are at $-$285 and 285\\,mJy/beam. {\\it Right}: 327 MHz VLA image. The image is obtained by combining the B and C configuration data, and the resolution is 27\\arcsec$\\times$21\\arcsec\\, with a PA=$-$88\\degr. The first two levels of contours are $-$3 and 3\\,mJy/beam.} \\label{3c35cntrnat} \\end{figure*} \\begin{figure*}[th] \\begin{center} \\includegraphics[width=14cm]{fig2.pdf} \\end{center} \\caption[]{Radio images of 3C223, all contours start at (3$\\sigma$) and are scaled by $\\sqrt{2}$. {\\it Left}: VLA image at 74 MHz. The image is obtained by combining the data of A and B configurations; the resolution is 26\\arcsec$\\times$25\\arcsec\\, with a PA=$-$61\\degr. The first two levels of contours are $-$120 and 120\\,mJy/beam. {\\it Right}: 327 MHz VLA image, obtained by combining the data of A and B configurations; the resolution is 7\\arcsec$\\times$6\\arcsec\\, with a PA=$-$78\\degr. The first two levels of contours are $-$1.8 and 1.8\\,mJy/beam.} \\label{3c223cntrnat} \\end{figure*} According to the standard model of active galactic nuclei (AGNs), at the center of active galaxies resides a super-massive black hole. The AGN is powered by an accretion disk surrounded by a torus of gas and dust \\cite[]{Blandford&Rees1974}. The powerful radio emission observed in classical double radio sources is produced by a bipolar pair of jets; relativistic outflows of matter which originate in the AGN. They first propagate in the interstellar medium (ISM) and then in the intergalactic medium (IGM) for a typical time of 10$^8$~yr \\cite[]{Scheuer1974} . The hot-spots are the regions where the energy carried by the jets is diffused into the radio lobes. The observed diffuse radio emission is produced in the 'cocoon' or lobe, which is formed by the built up jet material and/or energy in the region between the core and the hot-spots. The energy evolution of the cocoon can be traced by observations and spectral studies of the radio lobes. Radio lobes expand, and, assuming the source is in the equipartition regime, the pressure of the relativistic plasma in the lobe equals the pressure of the external environment \\cite[]{Begelman1984}. The radio spectrum of radio galaxies is initially described by a power-law. The final shape of the spectrum moves away from the power-law showing a steepening at higher frequencies. This is due to the competition between processes of energy injection and losses due to adiabatic expansion, synchrotron emission and inverse Compton scattering with the CMB photons, \\citep{K,Kellermann1964,P}. The initial models developed to interpret these spectra assumed a uniform and constant magnetic field and an isotropic injection of electrons \\citep[e.g.,][hereafter KP and JP]{K,P,JP}. If the former assumptions are satisfied it is possible, in theory, to use the synchrotron spectrum to estimate the age of the radiating particles. Many authors \\citep[e.g.,][]{Tribble1993,Eilek1997,Blundell2000} argue that the observed filamentary structures in the radio lobes \\citep[e.g. for Cygnus A,][]{Carilli1991} can be interpreted as the effect of inhomogeneous magnetic fields on the synchrotron emission. However, \\cite{Kaiser2000} demonstrated that the spatial distribution of the synchrotron radio emission can be used to estimate the age for FRII sources \\cite[]{Fanaroff1974}. Furthermore, based on the dynamical and radiative self-similar models in \\cite{Kaiser&Alexander1997} and \\cite{KaiserThorpe1997}, \\cite{Kaiser2000} developed a 3-dimensional model of the synchrotron emissivity of the cocoons of powerful FRII radio sources. The projection along the line of sight (LOS) of the 3D model can be easily compared with radio observations. X-ray emission related to the lobes has been detected in a number of radio sources. This is attributed to the inverse Compton (IC) scattering of microwave background photons. The direct estimates of the magnetic field ($B_{IC}$) obtained from the combination of the X-ray IC flux and the radio synchrotron spectrum, give values near to those found with the equipartition ($B_{eq}$) assumption \\citep[e.g.,][]{Croston2004,Croston2005}. Moreover, the above-mentioned comparisons suggest that lobes are not overpressured at the late stages of evolution of radio galaxies \\citep{Croston2004,Croston2005,Konar2009}. On the other hand, a strong variation of the X-ray/radio flux ratio across the lobes has been found \\citep{Isobe2002,Hardcastle2005,Goodger2008}. This cannot be explained with models in which either the electron energy spectrum or the magnetic field vary independently as function of position in the lobes, but it is consistent with models in which both vary together as function of position. \\noindent Among radio galaxies (RG), those with a projected linear size $\\gtsim$~1\\,Mpc\\footnote{Throughout we adopt H$_0$\\,=\\,71\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_m$\\,=\\,0.27, $\\Omega_{\\Lambda}$\\,=\\,0.73 \\cite[]{Spergel2003}. Many radio galaxies have been classified as giant in the past using a different set of cosmological parameters. For this reason some GRG could have a linear size slightly less than 1Mpc.} are defined as giant radio galaxies (GRG). In the complete sample of 3CR radio sources \\cite[]{Laing1983} around 6$\\%$ of the sources are giants; there are about 100 known. GRG typically have radio powers below $10^{26.5}$ W~Hz$^{-1}$~sr$^{-1}$, have linear sizes less than 3 Mpc, and are observed at redshifts $z<$~0.25, even though $z<$~0.5 could be assumed as an upper limit \\citep{Ishwara-Chandra1999,Schoenmakers2000,Lara2004,Saripalli2005,Machalski2007}. The P-D diagram \\citep{Lara2004,Ishwara-Chandra1999} shows a dearth of high luminosity GRG, as predicted by evolutionary models \\citep{Blundell1999,KaiserThorpe1997} and a maximum GRG linear size cut off of 3 Mpc. An estimate of the predominant process of radiative losses, obtained by separating the contributions of the inverse Compton and synchrotron losses, shows that the ratio of the estimated B$_{CMB}$/B$_{eq}$ increases with linear size, and IC losses dominate the radiative losses in GRG \\cite[]{Ishwara-Chandra1999}. As argued by many authors, the observed physical characteristics mentioned above could be the result of selection effects introduced by the selection criteria or by biases due to the low sensitivity of typical radio images. The faintest regions of GRG are well detected, even with a modest angular resolution, only with low frequency interferometric observations. The low-frequency spectral index information is crucial to derive the energy distribution of the radiating electrons, and to study the energy transport from the nucleus to the lobes in these exceptionally large radio sources. Multifrequency spectral aging studies of GRG have been recently presented by Jamrozy and collaborators \\citep{Jamrozy2004,Jamrozy2005,Jamrozy2008}. The median value for the estimated spectral ages is 23-24 Myr. The injection spectral index ranges from 0.55 to 0.88; it appears to increase with luminosity and redshift but shows an inverse correlation with linear size. In this paper we present a multifrequency spectral analysis of two classical double giant radio galaxies 3C35 and 3C223. In Sect. 2 radio observations and data analysis at 74 and 327\\,MHz are described. In Sect. 3 we present radio images of 3C35 and 3C223 at 74 and 327\\,MHz. In Sect. 4 we show the spectral index maps and the spectral analysis obtained by combining images at 74, 327, 608 and 1400 \\,MHz. Results are discussed and summarized in Sect. 5. \\begin{figure*}[ht] \\begin{center} \\includegraphics[width=14cm]{fig3.pdf} \\end{center} \\caption[]{ 3C35: Spectral index maps are shown in color; pixels whose brightness was below 3$\\sigma$ have been blanked. Contour levels are the radio brightness at 327\\,MHz, start at (3$\\sigma$) and are scaled by $\\sqrt{2}$. {\\it Left} Spectral index map between 74\\,MHz and 327\\,MHz, with a resolution of 95\\arcsec$\\times$95\\arcsec. {\\it Right} Spectral index map between 327\\,MHz and 1.4\\,GHz, with a resolution of 45\\arcsec$\\times$45\\arcsec (the image at 1.4\\,GHz was taken from the NVSS \\cite[]{NVSS}).} \\label{3c35spixmap} \\end{figure*} \\begin{figure*}[ht] \\begin{center} \\includegraphics[width=14cm]{fig4.pdf} \\end{center} \\caption[]{3C223: Spectral index maps are shown in color; pixels whose brightness was below 3$\\sigma$ have been blanked. Contour levels are the radio image at 327\\,MHz , start at (3$\\sigma$) and are scaled by $\\sqrt{2}$. {\\it Left} Spectral index map between 74\\,MHz and 327\\,MHz, with a resolution of 26\\arcsec$\\times$26\\arcsec. {\\it Right} Spectral index map between 327\\,MHz and 1.48\\,GHz, with a resolution of 7.5\\arcsec$\\times$7.5\\arcsec.} \\label{3c223spixmap} \\end{figure*} ", "conclusions": "In this paper we present new VLA images of the sources 3C35 and 3C223 at the observing frequencies of 327 and 74\\,MHz. By combining our images with those at 1.4\\,GHz available in the literature, we produced spectral index distribution maps between 74--327\\,MHz and 327\\,MHz--1.4\\,GHz for both sources. The spectral index across the sources are more constant in the low frequency range, while in the high frequency range the spectral indices increase from the hot-spots to the inner region of the lobes near to the core. In particular, for the source 3C35, $\\alpha$ ranges between 0.6 and 0.8 in the interval of frequencies 74--327\\,MHz, while between the frequencies 327\\,MHz--1.4\\,GHz the values of $\\alpha$ change from 0.6 in the hot-spot's region to 1.7 in the inner region of the lobes. On the other hand, for the source 3C223 the value of $\\alpha$ is on average 0.6 in the range 74--327\\,MHz , but it could reach extreme values which range between 0.4 and 1.6. In the range between 327\\,MHz--1.4\\,GHz $\\alpha$ varies from 0.7 in the hot-spots to 1.5 in the inner region of the lobes. By considering the two radio sources in a minimum energy condition, i.e. in the equipartition regime, we estimated the magnetic field of the two sources. The estimate was made using two different approaches often adopted in the literature, a fixed frequency range and a fixed energy range. Moreover, two different plasma populations has been considered (see Tab. \\ref{35flux} and \\ref{223flux}): one in which the energy is equally divided between relativistic protons and electrons and another one in which all the energy is provided by a plasma of relativistic electrons--positrons. For both sources the resulting equipartition magnetic field ranges between values of 0.5--1.6 $\\mu$G, in concordance with typical values of the measured IC magnetic fields. In particular, in the case of 3C223, the value of the equipartition magnetic field is within a factor of two in agreement with the measured IC magnetic field \\cite[]{Croston2004}. By using our images with those at higher frequencies available in the literature, we obtained the spectral shape of the radio spectrum in many different positions along the lobes. The hot spots of the source 3C35 are well described by power--laws, while the hot spots of 3C223 show quite curved spectra. The inner regions of the lobes for the two sources present a break in the range of frequency around 1.0\\,GHz. Since for both the sources the magnetic field is low, the inverse Compton losses are as important as the synchrotron losses, and we can assume an isotropic electron population. Therefore we fitted the spectra with a JP \\cite[]{JP} model to estimate the frequency break $\\nu_{break}$; in the case of 3C35 we also estimated $\\alpha_{inj}$ while for 3C223 we used a fixed value. In the case of 3C35 we found that $\\nu_{break} \\simeq$ 800\\,MHz and $\\alpha_{inj}$ is on average 0.5. For 3C223 the $\\nu_{break}$ is about 1.4\\,GHz with a fixed $\\alpha_{inj}$=0.5. The break frequency $\\nu_{break}$ is a time-dependent function, by assuming that there is no expansion and the magnetic field is constant, from the frequency break we calculated the radiative age of the source $t_{syn}$. \\cite{Blundell2000} claimed an anomalous diffusion of relativistic particles which implies that no information about the age of the source can be inferred from the shape of the emission radio spectrum. They discussed about the discrepancy between the estimates of the spectral and dynamical ages for sources older then $10^{7}$ yr. But \\cite{Kaiser2000} demonstrated that diffusion will not alter the distribution of relativistic particles, therefore the spatial distribution of the synchrotron radio emission can be use to estimate the age for FRII sources \\cite[]{Fanaroff1974}. Moreover, as we discussed above, the magnetic field of these particular sources is low with respect to the inverse Compton equivalent magnetic field, therefore a spatially variable magnetic field has a minor impact on the energy losses of the relativistic electrons. For 3C35 the estimated age is about 143$\\pm$20\\,Myr while for 3C223 is about 72$\\pm$4\\,Myr. The radiative age confirms that the two sources are rather old. However, these estimates must be considered upper limits if adiabatic losses can not be neglected. A high resolution analysis of the spectral index behaviour has been made by fitting the two frequency spectral index $\\alpha_{1.4}^{0.3}$ with the synchrotron model as function of the distance from the core. We fitted with a law $\\nu_{break} \\propto 1/d^{\\beta}$. For the source 3C35 the frequency break is about 800\\,MHz and parameter $\\beta$ is about 2.5$\\pm$0.7, in agreement with an expansion with a constant speed and/or unimportant adiabatic losses. For the source 3C223, $\\nu_{break} \\simeq$ 1.4\\,GHz, while the parameter $\\beta$ is about 3.6$\\pm$1.1; as discussed before this could be explained if adiabatic losses play an important role in the energy balance of the source and/or if the expansion velocity of the source is not constant." }, "1003/1003.3831_arXiv.txt": { "abstract": "We present a novel method to significantly speed up cosmological parameter sampling. The method relies on constructing an interpolation of the CMB-log-likelihood based on sparse grids, which is used as a shortcut for the likelihood-evaluation. We obtain excellent results over a large region in parameter space, comprising about 25 log-likelihoods around the peak, and we reproduce the one-dimensional projections of the likelihood almost perfectly. In speed and accuracy, our technique is competitive to existing approaches to accelerate parameter estimation based on polynomial interpolation or neural networks, while having some advantages over them. In our method, there is no danger of creating unphysical wiggles as it can be the case for polynomial fits of a high degree. Furthermore, we do not require a long training time as for neural networks, but the construction of the interpolation is determined by the time it takes to evaluate the likelihood at the sampling points, which can be parallelised to an arbitrary degree. Our approach is completely general, and it can adaptively exploit the properties of the underlying function. We can thus apply it to any problem where an accurate interpolation of a function is needed. ", "introduction": " ", "conclusions": "" }, "1003/1003.1719_arXiv.txt": { "abstract": "I present results from the modeling of stellar bars in nearly 300 barred galaxies in the local universe through parametric multi-component multi-band image fitting. The surface brightness radial profile of bars is described using a S\\'ersic function, and parameters such as bar effective radius, ellipticity, boxiness, length and mass, and bar-to-total luminosity and mass ratios, are determined, which is unprecedented for a sample of this size. The properties of bars in galaxies with classical bulges and pseudo-bulges are compared. For a fixed bar-to-total mass ratio, pseudo-bulges are on average significantly less massive than classical bulges, indicating that, if pseudo-bulges are formed through bars, further processes are necessary to build a classical bulge. I find a correlation between bar ellipticity and boxiness, and define a new parameter as the product of these two quantities. I also find correlations between this product and normalised bar size, between the sizes of bars and bulges, and between normalised bar size and bulge-to-total ratio. Bars with different ellipticities follow parallel lines in the latter two correlations. These correlations can arise if, starting off with different normalised sizes and ellipticities, bars grow longer and stronger with dynamical age, as a result of angular momentum exchange from the inner to the outer parts of galaxies, consistent with previous theoretical predictions. A plausible consequence is that bar pattern speeds should become lower with bar dynamical age, and towards galaxies with more prominent bulges. ", "introduction": "\\label{sec:intro} Many recent studies, from observational and theoretical viewpoints, have established that stellar bars in disc galaxies can play an important role in galaxy evolution \\citep[see][for reviews]{SelWil93,KorKen04,Gad09a}. Theoretical work indicates that the redistribution of angular momentum, induced by the bar, in the galaxy interstellar medium, as well as in the stellar and dark matter components, has a number of important consequences \\citep[e.g.][]{AthMis02,Ath03,MarShlHel06,BerShlJog06}. Gas lying beyond the bar ends is driven outwards, whereas gas lying within the bar ends is driven to the central regions \\citep[e.g.][]{Sch81,ComGer85,Ath92,FriBen93,FriBenKen94,PinStoTeu95}. This secular evolution scenario has been partially confirmed, at least qualitatively, with observational evidence that barred galaxies show flatter chemical abundance (O/H) radial gradients \\citep[][further, \\citealt{MarRoy94} find that the stronger the bar the flatter the gradient]{ZarKenHuc94} and higher central concentrations of molecular gas (CO -- \\citealt{SakOkuIsh99}). This movement of gas to the centre might in principle help build a young and kinematically cold stellar bulge component, i.e. a disc-like bulge \\citep[see][]{Ath05b}. Indeed, observations suggest that disc-like bulges exist and have formation processes linked to dynamical disc instabilities, such as bars, as opposed to the old and kinematically hot classical bulges \\citep[e.g.][and references therein]{CarStideZ97,GadDos01,ErwBelGra03,Fis06,DroFis07,FisDro08,Gad09b}. Theory also suggests how bars evolve with time. Broadly speaking, bars slow down their pattern rotation speed, and get longer and thinner (i.e. more eccentric and stronger) during the course of their evolution, capturing stars from the disc \\citep[see][]{Ath03}. Observations suggest that the strong bar in NGC 4608 has increased in mass by a factor of $\\approx1.7$, through the capture of $\\approx13\\%$ of the disc stars \\citep{Gad08}. In addition, more evolved bars also show more rectangular-like face-on isophotal shapes, i.e. they are more boxy \\citep[as in the $N$-body simulations of][]{AthMis02}. In detail, however, simulated bars can become abruptly shorter and thicker a few Giga-years after their formation, due to the onset of dynamical vertical instabilities that originate box/peanut bulges. (These seem to be simply the inner parts of bars that buckle off the disc plane and can be seen in inclined systems.) Then, about 1 Gyr later, they recover the original evolutionary trends. Furthermore, a substantial gas component in the disc can also complicate the picture of the evolution of bar properties, halting how these properties change, and in some cases even reversing the trends \\citep[e.g.][]{BouCom02,BouComSem05,DebMayCar06,BerShlMar07}. To date, there is no study aimed directly at providing observational evidence on how bar properties change with time. This is partially due to the difficulty of estimating bar dynamical ages. Although some first steps have been done in this direction \\citep[see][]{GadDes05,GadDes06,PerSanZur09}, the results are as yet inconclusive, and the methods developed require large amounts of telescope time. Moreover, bars are found very often in disc galaxies, and the fraction of disc galaxies hosting bars seems to increase with time, i.e. the fraction is lower at redshift $z\\sim1$, as compared to $z\\sim0$ \\citep[see][and references therein, but see also \\citealt{BarJogMar08} and \\citealt{JogBarRix04}]{SheElmElm08}. In addition, weaker, but also global non-axisymmetric structures, such as oval distortions in the disc, can as well efficiently produce such redistribution of angular momentum. Therefore, the consequences of the presence of such structures should be conspicuous, and studies on the properties of bars and their host galaxies at $z\\approx0$, as well as higher redshifts, can give direct clues on galaxy evolution. The first and easier step in observational studies of barred galaxies is to identify bars, and there are a number of studies on the fraction of disc galaxies with bars \\citep[e.g.][]{EskFroPog00,BarJogMar08,MarJogHei09}. Other studies have used ellipse fits to the images of barred galaxies to obtain bar properties such as ellipticity and length \\citep[e.g.][]{MarJog07,BarJogMar08,BarJabDes09}, although it has been shown that ellipse fits can lead to an underestimation of the bar ellipticity \\citep[or an overestimation of the bar axial ratio,][]{Gad08}. In \\citet{Mar95}, bar axial ratios and lengths were visually assessed, and a relation was found between bar length and the normalised diameter of the bulge. Ten years later, \\citet{Erw05} measured bar lengths using ellipse fits and found that bar size scales with disc size, and confirmed previous results that bars in early-type disc galaxies are clearly larger than those in late-type galaxies \\citep[see also][]{AguMenCor09}. He went further and argued that this observational evidence can be qualitatively consistent with the simulations that show bars getting larger with time, if indeed secular evolution produces early-type disc galaxies from late-type ones \\citep[see also][]{FriBen95,MarFri97,GadDes05,GadDes06}. Applying ellipse fits to 2MASS images of 151 spiral galaxies, \\citet{MenSheSch07} obtained measurements of bar length and axial ratio, and found a weak trend of higher ellipticities for larger bars, which is also consistent with the results from simulations. More recently, \\citet[][see also \\citealt{deJ96c,LauSalBut05,LauSalBut07,Gad08}]{DurSulBut08} used 2D bulge/bar/disc decompositions of 97 Sb, Sbc and Sc galaxies, employing $i$-band images from the Sloan Digital Sky Survey (SDSS) to obtain parameters such as bar length and bar-to-total luminosity ratio. \\citet{WeiJogKho09} did a similar work using $H$-band images of 143 spirals and explored also the stellar mass content in bars. About 60\\% of the galaxies in both samples are barred. In this paper, I explore the results of 2D bulge/bar/disc decompositions of 291 barred galaxies, using SDSS images in the $g$ and $i$ bands, from \\citet[][hereafter Paper I]{Gad09b}. This allows me to study bar properties in a level of detail which is unprecedented for a sample of this size. A thorough characterization of bars in massive galaxies in the local universe is thus put forth in Sect. \\ref{sec:comp}, after a description of the data at hand in the next section. In Sect. \\ref{sec:clapse}, the properties of bars in galaxies with classical and pseudo-bulges are compared. In Sect. \\ref{sec:corr}, I explore correlations between bar properties in order to test the predictions from simulations on the secular growth of bars. These results are discussed in Sect. \\ref{sec:disc}, while Sect. \\ref{sec:conc} summarises the paper. ", "conclusions": "\\label{sec:conc} I have explored the results from detailed 2D image decomposition of nearly 300 barred galaxies with stellar masses above $10^{10}~{\\rm M}_\\odot$, at $z\\sim0$, concerning mainly the structural properties of bars. This results in a thorough description of bonafide stellar bars in the local universe, including distributions of bar S\\'ersic index, ellipticity, boxiness, length and bar-to-total luminosity and mass ratios. The interplay between bars and the bulges and discs in their host galaxies was also examined. Such detailed characterization of local bars can be compared with similar results from studies with samples at higher redshifts, in order to directly investigate how the structure of barred galaxies evolve in time. It can also be compared to results from theoretical work, in order to assess how well theory can describe the structural properties of such bars. Furthermore, a complete description of the properties of local bars is useful to insert ad hoc models of bars in a theoretical framework. Bars in galaxies hosting classical and pseudo-bulges share similar properties, except that bars in the former are on average larger than those in the latter, considering both absolute and normalised sizes. This is consistent with previous results comparing bar sizes in early- and late-type disc galaxies, as usually the former host classical bulges, whereas the latter host pseudo-bulges. For a fixed bar-to-total mass ratio, pseudo-bulges are on average significantly less massive than classical bulges. This indicates that, if pseudo-bulges are formed through disc instabilities such as bars, then more than that is necessary to build a classical bulge. Normalised bar size is correlated with the product of the bar ellipticity and boxiness ($\\epsilon\\times c$, which is related to bar strength) and with bulge-to-total ratio, $B/T$. If higher $B/T$ indicates more evolved bars, then these results can be interpreted as qualitatively consistent with general expectations from theoretical work, which suggests that evolved bars grow longer and stronger in time. This would come along with a decrease in bar pattern speed, and assumes that effects caused by gas dynamics are sufficiently small. Bars with different ellipticities follow parallel tracks in the trend found between normalised bar size and $B/T$, suggesting that bars could form with different normalised sizes and ellipticities but still follow similar evolutionary paths." }, "1003/1003.5283_arXiv.txt": { "abstract": "In our earlier work on the development of a model--independent data analysis method for reconstructing the (moments of the) time--averaged one--dimensional velocity distribution function of Weakly Interacting Massive Particles (WIMPs) by using measured recoil energies from direct Dark Matter detection experiments directly, it was assumed that the analyzed data sets are background--free, i.e., all events are WIMP signals. In this article, as a more realistic study, we take into account a fraction of possible residue background events, which pass all discrimination criteria and then mix with other real WIMP--induced events in our data sets. Our simulations show that, for the reconstruction of the one--dimensional WIMP velocity distribution, the maximal acceptable fraction of residue background events in the analyzed data set(s) of ${\\cal O}(500)$ total events is $\\sim$ 10\\% -- 20\\%. For a WIMP mass of 50 GeV with a negligible uncertainty and 20\\% residue background events, the deviation of the reconstructed velocity distribution would in principle be $\\sim$ 7.5\\% with a statistical uncertainty of $\\sim$ 18\\% ($\\sim$ 19\\% for a background--free data set). ", "introduction": "Currently, direct Dark Matter detection experiments searching for Weakly Interacting Massive Particles (WIMPs) are one of the promising methods for understanding the nature of Dark Matter and identifying them among new particles produced at colliders as well as reconstructing the (sub)structure of our Galactic halo \\cite{Smith90, Lewin96, SUSYDM96, Bertone05}. In our earlier work \\cite{DMDDf1v}, we developed methods for reconstructing the (moments of the) time--averaged one--dimensional velocity distribution of halo WIMPs by using (a functional form of) the recoil spectrum as well as the measured recoil energies directly. This analysis requires {\\em no} prior knowledge about the WIMP density near the Earth {\\em nor} about their scattering cross section on nucleus, the unique required information is the mass of incident WIMPs. We therefore turned to also develop the method for determining the WIMP mass model--independently by combining two experimental data sets with two different target nuclei \\cite{DMDDmchi-SUSY07, DMDDmchi}. In the work on the development of these model--independent data analysis procedures by using measured recoil energies from direct detection experiments directly, it was assumed that the analyzed data sets are background--free, i.e., all events are WIMP signals. Active background discrimination techniques should make this condition possible. For example, the ratio of the ionization to recoil energy, the so--called ``ionization yield'', used in the CDMS-II experiment provides an event--by--event rejection of electron recoil events to be better than $10^{-4}$ misidentification \\cite{Ahmed09b}. By combining the ``phonon pulse timing parameter'', the rejection ability of the misidentified electron recoils (most of them are ``surface events'' with sufficiently reduced ionization energies) can be improved to be $< 10^{-6}$ for electron recoils \\cite{Ahmed09b}. Moreover, as demonstrated by the CRESST collaboration \\cite{CRESST-bg}, % by means of inserting a scintillating foil, which causes some additional scintillation light for events induced by $\\alpha$-decay of $\\rmXA{Po}{210}$ and thus shifts the pulse shapes of these events faster than pulses induced by WIMP interactions in the crystal, the pulse shape discrimination (PSD) technique can then easily distinguish WIMP--induced nuclear recoils from those induced by backgrounds% \\footnote{ For more details about background discrimination techniques and status in currently running and projected direct detection experiments, see e.g., Refs.~\\cite{Aprile09a, EDELWEISS-bg, % Lang09b} % }. However, as the most important issue in all underground experiments, the signal identification ability and possible residue background events which pass all discrimination criteria and then mix with other real WIMP--induced events in our data sets should also be considered. Therefore, in this article, as a more realistic study, we follow our first work on the effects of residue background events on the determination of the WIMP mass \\cite{DMDDbg-mchi} and want to study how well we could reconstruct the WIMP velocity distribution model--independently by using ``impure'' data sets and how ``dirty'' these data sets could be to be still useful. The remainder of this article is organized as follows. In Sec.~2 I review the model--independent method for reconstructing the time--averaged one--dimensional velocity distribution function of halo WIMPs by using data from direct detection experiments directly. In Sec.~3 the effects of residue background events in the analyzed data sets on the measured energy spectrum as well as on the reconstructed WIMP mass will briefly be discussed. In Sec.~4 I show numerical results of the reconstructed WIMP velocity distribution by using mixed data sets with different fractions of residue background events based on Monte Carlo simulations. I conclude in Sec.~5. Some technical details will be given in an appendix. ", "conclusions": "In this paper we reexamine the model--independent data analysis method introduced in Ref.~\\cite{DMDDf1v} for the reconstruction of the one--dimensional velocity distribution function of Weakly Interacting Massive Particles from data (measured recoil energies) of direct Dark Matter detection experiments directly by taking into account a fraction of residue background events, which pass all discrimination criteria and then mix with other real WIMP--induced events in the analyzed data sets. This method requires neither prior knowledge about the WIMP scattering spectrum nor about different possible background spectra; the unique needed information is the recoil energies recorded in direct detection experiments and (eventually) the mass of incident WIMPs. For the mass of incident WIMPs required as an unique input information in this method, we first assumed that it could be known precisely with a negligible uncertainty from other (e.g., collider) experiments. Our simulations show that, assuming an exponential form for the residue background spectrum, with a data set of $\\sim$ 500 total events, and a background ratio of $\\sim$ 10\\% -- 20\\%, the WIMP velocity distribution function could in principle be reconstructed with an \\mbox{$\\sim$ 6.5\\%} (for a 25 GeV WIMP mass, 20\\% background events) -- $\\sim$ 38\\% (for a 250 GeV WIMP mass, 10\\% background events) deviation. If the WIMP mass is $\\cal O$(50 GeV), the maximal acceptable background ratio could be risen to $\\sim$ 40\\% with a deviation of only $\\sim$ 14\\%. Moreover, for lighter/heavier WIMP masses, since the relatively flatter/sharper background spectrum could contribute relatively more events to high/low energy ranges, the reconstructed velocity distribution could therefore be shifted to higher/lower velocities. However, since for (very) light WIMPs ($\\mchi~\\lsim~40$ GeV), the kinematic maximal cut--off of the recoil energy due to the Galactic escape velocity is (much) lower than the experimental cut--off, a (large) fraction of background events in high energy ranges could thus in practice be neglected, and the shift could not be very significant for WIMPs lighter than $\\sim$ 50 GeV. Since a model--independent method for determining the WIMP mass by using two experimental data sets with two different target nuclei has also been developed \\cite{DMDDmchi-SUSY07, DMDDmchi}, we considered in this paper also the case that the velocity distribution function is reconstructed with a reconstructed WIMP mass from other direct detection experiments. Our simulations show that, since lighter/heavier WIMP masses could be over-/underestimated by using this method with background--mixed data sets \\cite{DMDDbg-mchi}, the reconstructed points of the velocity distribution would thus be shifted to lower/higher velocities, the opposite direction of the shift due purely to the background contribution to high/low energy ranges. Although this second effect shifts the reconstructed velocity distribution (much) more strongly, with $\\sim$ 5\\% -- 10\\% background events mixed in the analyzed data sets, the WIMP velocity distribution function could in principle still be reconstructed with an $\\sim$ 7\\% (for 25 GeV WIMPs, 10\\% backgrounds) -- $\\sim$ 16\\% (for 250 GeV WIMPs, 5\\% backgrounds) deviation. If the WIMP mass is $\\lsim~{\\cal O}$(100 GeV), the maximal acceptable background ratio could even be as large as $\\sim$ 20\\% with a deviation of only $\\sim$ 9\\%. Furthermore, in order to check the need of a prior knowledge about an (exact) form of the residue background spectrum, a constant spectrum for residue backgrounds has also been considered. Since the WIMP mass would always be overestimated \\cite{DMDDbg-mchi}, the reconstructed WIMP velocity distribution could thus be (strongly) shifted to lower velocity ranges. However, data sets with background fractions of $~\\lsim~5\\%$ could in principle be used to at least give a rough outline of the WIMP velocity distribution (for $\\mchi~\\gsim~100$ GeV), or even reconstruct the distribution pretty well (for $\\mchi~\\lsim~100$ GeV). Finally, for rather next--to--next generation detectors, we considered also the case of 5,000 total events and extended the maximal experimental cut--off energies for WIMP signals and backgrounds to 150 keV. Assuming a maximal background ratio of $\\lsim~5\\%$, our results show that the WIMP velocity distribution function could in principle be reconstructed in the velocity range $v~\\lsim~{\\cal O}$(500 km/s) with a deviation of $\\lsim~6\\%$ (for a WIMP mass of 100 GeV). Once WIMPs are light ($\\mchi~\\lsim~{\\cal O}$(50 GeV)), this deviation could even be reduced to $\\lsim~2.5\\%$. In summary, as the second part of the study of the effects of residue background events in direct Dark Matter detection experiments, we considered the reconstruction of the velocity distribution function of halo WIMPs. Our results show that, with projected experiments using next--generation detectors with $10^{-9}$ to $10^{-11}$ pb sensitivities \\cite{Baudis07a, Drees08, Aprile09a, Gascon09} and $< 10^{-6}$ background rejection ability \\cite{CRESST-bg, % EDELWEISS-bg, % Lang09b, % Ahmed09b}, % once one or more experiments with different target nuclei could accumulate a few hundreds events (in one experiment), we could in principle at least give a rough outline of the WIMP velocity distribution function, e.g., an approximate estimate of the location of its peak, even though there could be some background events mixed in our data sets for the analysis. After that, by means of increased number of observed (WIMP--induced) events and improved background discrimination techniques \\cite{CRESST-bg, % EDELWEISS-bg}, % the shape and properties of the velocity distribution of halo Dark Matter should be understood more clearly. \\subsubsection*" }, "1003/1003.0974_arXiv.txt": { "abstract": "We present a new method for constructing maps of the secondary temperature fluctuations imprinted on the cosmic microwave background (CMB) radiation by photons propagating through the evolving cosmic gravitational potential. Large cosmological N-body simulations are used to calculate the complete non-linear evolution of the peculiar gravitational potential. Tracing light rays back through the past lightcone of a chosen observer accurately captures the temperature perturbations generated by linear (the integrated Sachs-Wolfe or ISW effect) and non-linear (the Rees-Sciama or RS effect) evolution. These effects give rise to three kinds of non-linear features in the temperature maps. (a) In overdense regions, converging flows of matter induce cold spots of order 100~Mpc in extent which can dominate over the ISW effect at high redshift, and are surrounded by hot rings. (b) In underdense regions, the RS effect enhances ISW cold spots which can be surrounded by weak hot rings. (c) Transverse motions of large lumps of matter produce characteristic dipole features, consisting of adjacent hot and cold spots separated by a few tens of Megaparsecs. These non-linear features are not easily detectable; they modulate the ISW sky maps at about the 10~per~cent level. The RS effect causes the angular power spectrum to deviate from linear theory at $l \\sim 50$ and generates non-Gaussianity, skewing the one-point distribution function to negative values. Cold spots of similar angular size, but much smaller amplitude than the CMB cold spot reported by Cruz et al. are produced. Joint analysis of our maps and the corresponding galaxy distribution may enable techniques to be developed to detect these non-linear, non-Gaussian features. Our maps are available at http://astro.dur.ac.uk/$\\sim$cai/ISW. ", "introduction": "The integrated Sachs-Wolfe (ISW) effect \\citep{Sachs67} arises from the decay of the large-scale potential fluctuations as they are traversed by cosmic microwave background (CMB) photons and induces secondary temperature perturbations in the CMB radiation. It occurs in both open cosmological models ($\\Omega_{\\rm{m}}<1$) and models containing a cosmological constant ($\\Omega_\\Lambda>0$) or dark energy. It provides an independent way of measuring the dynamical effect of dark energy, but it is challenging to detect: at large scales, the ISW signal suffers from cosmic variance, as there are too few independent modes, while at somewhat smaller scales, it is entangled with the Rees-Sciama (RS) effect \\citep{Rees68}, which arises from the non-linear evolution of gravitational potential perturbations. Since it is impossible to evade the cosmic variance at large scales, the information at small scales becomes valuable. To extract cosmological information from such scales requires disentangling the ISW and RS effects. Studies using N-body simulations have found that the combined ISW plus RS power spectra of CMB temperature fluctuations deviate from linear theory expectation at very large scales, $k \\sim 0.1$ $h$~Mpc$^{-1}$ at $z=0$, with the scale of this transition becoming even larger at higher redshift. This discovery that the non-linear effect was more important at higher redshift was made by \\citet{Cai09} and later confirmed by \\citet{Smith09}. Recent studies using both perturbation theory and fitting formulae indicate that the combined ISW and RS effect makes a non-trivial contribution to the overall CMB non-Gaussianity as measured by the bispectrum \\citep[e.g.][]{Verde02,Giovi03, Boubekeur09, Mangilli09}. The primordial-lensing-RS correlation contribution may yield an effective non-Gaussianity parameter, $f_{\\rm NL}$, of 10 \\citep{Mangilli09}. All these results suggest that the RS effect is a very important supplement to the ISW effect even at very large scales. Observational evidence for a possible contribution from the RS effect at large scales has come from a combined study of the SDSS LRG samples and the CMB \\citep{Granett09}. These authors report a $4$-$\\sigma$ detection of the ISW and RS signal at the scale of 4~degrees, somewhat higher than expected from the standard flat $\\Lambda$CDM universe. There have also been attempts to attribute an extreme cold spot in the CMB to the ISW and RS effects \\citep[e.g.][] {Martinez-Gonzalez90b,Martinez-Gonzalez90a, Rudnick07,Inoue06,Inoue07,Tomita08, Masina09b, Masina09a}. However, on theoretical grounds such an explanation seems unlikely because the estimated sizes of the non-linear structures responsible for the cold spot are typically $>100$~Mpc, which seems too large to occur in a $\\Lambda$CDM universe which assumes Gaussian initial conditions \\citep[see also][for discussion of non-Gaussianity.] {Cruz05,McEwen05, Cruz06,McEwen06, Cruz07,McEwen08}. However, it is still unclear whether or not the combined ISW and RS effects can generate cold spots of a few degrees with the right amplitudes and whether such large-scale non-Gaussianity can arise from large-scale structure. A full understanding of the ISW and Rees-Sciama effect could be crucial in explaining the oddities of these observations. Meanwhile, the increase in sensitivity of forthcoming CMB experiments opens possibilities of exploiting CMB temperature fluctuations down to arcmin scales (Planck\\footnote{www.sciops.esa.int/PLANCK/}, ACT\\footnote{http://www.physics.princeton.edu/act/}, SPT\\footnote{http://pole.uchicago.edu/} and APEX-SZ\\footnote{http://bolo.berkeley.edu/apexsz}), at which the ISW and RS effect may also entangle with other large-scale astrophysics of interest, i.e. lensing \\citep[e.g.][]{Verde02, Nishizawa08,Mangilli09} and the Sunyaev-Zel'dovich (SZ) effect \\citep{Sunyaev72} \\citep[e.g.][]{Cooray02a, Fosalba03,Bielby09}. To disentangle all these effects is the key to thoroughly exploiting the information encoded in these upcoming CMB measurements. N-body simulations are the ideal tool for investigating the phenomena discussed above since they treat the non-linear regime accurately and permit the construction of a full sky map of the ISW and RS effects and the full underlying 3-dimensional light-cone. Maps of the ISW effect have been constructed from both simulations and observations \\citep{Barreiro08,Granett09}. Most of them simply adopt the linear approximation, using only the density field to estimate the time derivative of the potential. We will show in this paper that these linear maps are far from accurate. There are also maps constructed from ray-tracing through simulations \\citep{Tuluie95,Puchades06}, but these are limited by small simulation box sizes and are not adequate to explore very large-scale structures. A full sky map of the RS effect has been constructed using a constrained high-resolution hydrodynamical simulation to model the RS effect in the very local universe \\citep{Maturi07a}. This map is useful for understanding the RS effect from within the radial distance of $110$ Mpc, which is a very small volume. Maps from the ray-tracing of large cosmological volumes are still missing. In this paper, we develop a new method of constructing a full sky light-cone of the time derivative of the potential, $\\dot{\\Phi}$, using a large N-body simulation. Our method of computing $\\dot{\\Phi}$ is fully non-linear and so should model the complete RS effect as well as the ISW component. Our Gpc box size simulation provides a sufficient number of independent large scale modes to investigate the ISW effect fully. We ray-trace through the light-cone to produce maps of temperature fluctuations induced by the ISW and RS effects. Our maps cover a large range of scales and cosmic time with high accuracy and allow us to investigate the ISW and RS effects. The maps will also be a valuable source for understanding CMB secondary non-Gaussianity arising from large-scale structure. They may also prove useful for disentangling lensing and SZ effects from the ISW and RS effects. The paper is organized as follows. In \\S2, we present the basic physics and the mathematical description of the ISW and RS effects. In \\S3, we describe our method of computing CMB temperature perturbations from our N-body simulation and ray-tracing to produce full sky maps from light-cone data. In \\S4, we identify and discuss three characteristic non-linear features of the temperature perturbations. Full sky maps are presented in \\S5. Finally, in \\S6, we discuss our results and draw conclusions. ", "conclusions": "Our aim in this paper is to investigate fully the gravitational effect that CMB photons suffer when passing through the evolving non-linear gravitational potential, $\\Phi$, of intervening large-scale structure. We have developed a method of using large cosmological N-body simulations to compute the time derivative of the potential, $\\dot{\\Phi}$, along the past light-cone of an observer. By integrating along light rays we have created full sky maps, both using the full non-linear calculation of $\\dot{\\Phi}$ and an alternative in which the dark matter velocity field is assumed to be related to its density field by the normal linear theory relation. By comparing the results of the two calculations we were able to assess fully the linear Integrated Sachs-Wolfe (ISW) and non-linear Rees-Sciama (RS) contributions to the induced CMB temperature fluctuations. In general, in a $\\Lambda$CDM universe, the linear ISW effect is dominant at low redshift where the accelerating expansion, driven by dark energy, causes the decay of perturbations in the gravitational potential. The propagation of CMB photons through these decaying potential wells and hills produces hot and cold regions respectively with $\\Delta T$ of the order of $10\\mu$K on scales of hundreds of Megaparsecs. At low redshift the non-linear (RS) effect produces only small scale perturbations to this large scale pattern. However, with increasing redshift the RS contribution decreases in amplitude much more slowly than does the contribution from the ISW effect, which vanishes as $\\Omega_{\\rm m}(z)$ approaches unity. Hence, the importance and scale of the RS effect becomes larger at higher redshift, confirming the conclusions of \\citet{Cai09} and \\citet{Smith09}. We have shown that the origin of the RS effect is primarily the non-linear relation between the velocity and density field rather than the non-linearity of the density field itself.\\footnote{The fact that to model $\\dot \\Phi$ and the growth of large-scale structure requires more accurate modelling of the large-scale velocity field than is provided by linear theory may suggest that using redshift space distortions to measure the growth rate of density perturbations, $\\beta$, may also not be sufficiently accurate.} Our investigation of the RS effect has revealed three distinct non-linear phenomena that give rise to corresponding characteristic features in the temperature perturbation maps. \\begin{itemize} \\item Dipoles are produced by the transverse motion of large lumps of dark matter, with typical sizes of tens of Megaparsecs, much larger than the scale of individual galaxy clusters. \\item Convergent flows, on the scale of up to 100$h^{-1}$Mpc, around non-linear overdense regions give rise to cold spots of order $\\mu$K surrounded by hot filamentary shells. At high redshift these can be strong enough to dominate over the linear ISW effect and change the sign of the temperature perturbation centred on overdense regions. \\item Divergent flows around void regions are characterised by RS contributions consisting of hot rings around cold central regions, where the density contrast of the void is nearly saturated ($\\delta \\approx -1$). This is a small effect at low redshift, but acts to strengthen the cold spots produced by the linear ISW effect. \\end{itemize} Unfortunately, none of these effects can be easily detected individually. At low redshift these phenomena make only 10\\% changes to the temperature perturbations predicted by the linear ISW effect. At very high redshift they are completely dominant, but their amplitudes are very low. We find that they contribute to cold spots of comparable physical scale to those reported in the literature, but their amplitudes are many times smaller. It may be possible to detect these large scale features induced by the non-linear velocity field by employing stacking techniques. The detectability of RS kinematic features produced by merging clusters has been discussed by \\citet{Rubino-Martin04} and \\citet{Maturi07b}. They conclude that around a thousand clusters would be needed to detect the RS signal above the contaminating effects of the primordial CMB temperature fluctuations and instrument noise. However, we have found imprints of the RS perturbations on scales of a few tens of Megaparsecs, much larger than the merging cluster scale, and with slightly larger amplitudes. Thus, these large-scale features might be more easily detected, requiring the co-addition of fewer objects. From our all sky maps we find that the RS contributions to the overall power spectrum of temperature perturbations become important for $l> 80$ (a few degrees) and completely dominate for $l > 200$. At still smaller scales the kinetic SZ effect is expected to dominate and at such scales a full treatment would have to incorporate additional modelling of this contribution. The RS contribution to the temperature maps is strongly non-Gaussian with a skewed one-point distribution. In future work it will be interesting to investigate the RS contribution to higher order statistics as its non-Gaussian characteristics might limit the ability to detect primordial non-Gaussianity in the underlying primary CMB fluctuations. Combining our full-sky maps with mock galaxy catalogues built from the same N-body simulations will be a powerful tool for developing cross-correlation techniques aimed at extracting the ISWRS signal from redshift surveys." }, "1003/1003.0145_arXiv.txt": { "abstract": "Deep Very Large Array imaging of the binary X-ray source SS\\,433, sometimes classified as a microquasar, has been used to study the intrinsic brightness distribution and evolution of its radio jets. The intrinsic brightness of the jets as a function of age at emission of the jet material $\\tau$ is recovered by removal of the Doppler boosting and projection effects. We find that intrinsically the two jets are remarkably similar when compared for equal $\\tau$, and that they are best described by Doppler boosting of the form $D^{2+\\alpha}$, as expected for continuous jets. The intrinsic brightnesses of the jets as functions of age behave in complex ways. In the age range $60 < \\tau < 150$~days, the jet decays are best represented by exponential functions of $\\tau$, but linear or power law functions are not statistically excluded. This is followed by a region out to $\\tau \\simeq 250$~days during which the intrinsic brightness is essentially constant. At later times the jet decay can be fit roughly as exponential or power law functions of $\\tau$. ", "introduction": "The galactic X-ray binary source SS\\,433, consisting of a stellar-mass black hole in close orbit about an early-type star \\citep{BBS2008,HG2008,B2010}, is a miniature analogue of an AGN \\citep{MR99}, and is often classified as a microquasar. Two mildly relativistic jets emerge from opposite sides of the compact object at speeds of $0.26~c$. Modeling of the optical spectrum shows that the jet system precesses with a period of 162 days about a cone of half-angle $20^{\\circ}$ \\citep{AM79, Fabian, Milgrom, M84}. Imaging by the Very Large Array (VLA) confirmed this picture, as the radio images showed helical jets on both sides of the source \\citep{Spencer79,HJ81a,HJ81b}. Higher resolution images made by VLBI reveal the structure down to a scale of a few AU \\citep{Ver87,Ver93}. Analysis of the 15~GHz VLA-scale structure of the jets in SS\\,433, with angular resolution of about $0.1\\arcsec$, was presented in \\citet{PaperI}. Multi-epoch dual-frequency analysis of SS\\,433 during the summer of 2003 will be presented in \\citet{PaperIII}, hereafter Paper~III, and for the summer of 2007 in \\citet{PaperIV}, hereafter Paper~IV. In this paper we use high dynamic range VLA images of SS\\,433 to study the radiative intensity of the two jets as a function of the material's {\\em birth epoch} $t$ and {\\em age at emission} $\\tau$ (hereafter simply ``age''); see Appendix~1 (\\S\\ref{s:app1}) for our definitions of these quantities. Our goals are (i) to determine if the two jets are intrinsically the same, and (ii) to learn if the jets behave as individual non-interacting components or as a continuous stream. SS\\,433 offers a unique opportunity to answer these questions because it presents two jets with ever-changing mildly-relativistic velocities known as functions of time and position on the sky from their optical properties. In \\S\\ref{s:obs} we describe the observations and data reduction. In \\S\\ref{s:jets} we determine the properties of the jets and we discuss the physical implications of these results in \\S\\ref{s:Discussion}. Our conclusions are summarized in \\S\\ref{s:conclude}. ", "conclusions": "\\label{s:conclude} The principal results in this paper are: \\begin{enumerate} \\item We have used a deep VLA A-array image of SS\\,433 at 4.86~GHz to study the intrinsic brightness profiles of the twin jets. \\item Radiation from both jets is detected out to at least $6\\arcsec$ from the core, corresponding to jet ages of about 800~days. \\item The observed brightnesses of the jets are strongly affected by projection effects and Doppler boosting. \\item Intrinsically the two jets are remarkably similar, and they are best described by Doppler boosting of the form $D^{2+\\alpha}$, as expected for a continuous jet. \\item The intrinsic brightness of the jets behaves in a complex way that is not well described by single linear, exponential, or power law decay. \\item During their first $\\sim$150~days, the jet decays are well represented by linear or exponential functions of age, with linear half-lives or exponential half-lives of about 40~days, the same for the two jets. Power law fits to the data in this age range give exponents of about $-1.8$. \\item There is a transition region, corresponding to jet ages between about 150 and 250~days, during which the jets maintain roughly constant intrinsic brightnesses. This represents nearly one complete precession period. This also corresponds to about $ 150 < t < 250 \\mbox{ days}$ in either jet. \\item At later times the jet decay can be roughly fit as exponential functions of age, with exponential half-lives of about 80~days, or as power laws with indices of $a \\leq 4$. \\end{enumerate}" }, "1003/1003.2320_arXiv.txt": { "abstract": "Explosive astrophysical systems, such as supernovae or compact star binary mergers, provide conditions where strange quark matter can appear. The high degree of isospin asymmetry and temperatures of several MeV in such systems may cause a transition to the quark phase already around saturation density. Observable signals from the appearance of quark matter can be predicted and studied in astrophysical simulations. As input in such simulations, an equation of state with an integrated quark matter phase transition for a large temperature, density and proton fraction range is required. Additionally, restrictions from heavy ion data and pulsar observation must be considered. In this work we present such an approach. We implement a quark matter phase transition in a hadronic equation of state widely used for astrophysical simulations and discuss its compatibility with heavy ion collisions and pulsar data. Furthermore, we review the recently studied implications of the QCD phase transition during the early post-bounce evolution of core-collapse supernovae and introduce the effects from strong interactions to increase the maximum mass of hybrid stars. In the MIT bag model, together with the strange quark mass and the bag constant, the strong coupling constant $\\alpha_s$ provides a parameter to set the beginning and extension of the quark phase and with this the mass and radius of hybrid stars. ", "introduction": "The future FAIR facility at GSI, Darmstadt, will explore the equation of state (EoS) of strongly interacting matter for intermediate temperatures $T$ and high baryon densities $n_b$ around isospin symmetry, that is for proton fractions $Y_p \\sim 0.5$. Supernovae (SNe) and binary mergers hold environments with similar conditions for $T$ and $n_b$ but with $Y_p \\leq 0.3$. As will be discussed in the scope of this article, core-collapse SNe with matter at a low value of $Y_p$ and dynamical timescales in the range of ms, provide conditions suitable for a phase transition to strange quark matter. Such a scenario was recently studied in \\cite{Sagert09} applying the MIT bag approach for the EoS of quark matter and using low critical densities for its onset. Simulations with different progenitor models and two different bag constants led to SN explosions accompanied by a significant neutrino burst which can be observed by present and future neutrino detectors \\cite{Dasgupta09}. In the following, we will introduce in more detail the hybrid EoS used in the above work and analyze its influence on the dynamics of the PNS evolution. We will discuss the compatibility with heavy ion (HI) data and pulsar observations. Furthermore we will include first order corrections from the strong interaction constant $\\alpha_s$ and study its influence on the maximum mass of the cold hybrid star configurations. ", "conclusions": "Because of the different proton fractions of matter in terrestrial and astrophysical laboratories, such as the future FAIR facility at GSI and supernovae or compact star mergers, on the one hand and their similarities in $T$ and $n_b$ on the other, the study of heavy ion collisions and explosive astrophysical scenarios can complement each other in probing the phase diagram of strongly interacting matter, also in regard to the phase transition from hadronic to quark matter. However, the study of possible observable signals and impacts of quark matter in astrophysical systems requires hydrodynamical simulations with an input of an appropriate quark-hadron equation of state. In this article we present such an approach where a quark matter phase transition has been implemented in a hadronic equation of state for a large range of temperatures, proton fractions and densities. Applying the latter to simulations of core-collapse supernovae, we find that a quark matter phase transition can cause the formation of a second shock wave which leads to the explosion of the star, accompanied by a second neutrino burst. The latter is dominated by anti-neutrinos, which can be observed by future and present neutrino detectors. If found, the second neutrino peak can give correlated information about the progenitor mass and the critical density for the onset of quark matter. \\paragraph{Acknowledgement} The project was funded by the Swiss National Science Foundation grant. no. PP00P2- 124879/1 and 200020-122287 and the Helmholtz Research School for Quark Matter Studies, the Italian National Institute for Nuclear Physics, the Graduate Program for Hadron and Ion Research (PG-HIR), the Alliance Program of the Helmholtz Association (HA216/EMMI) and the DFG through the Heidelberg Graduate School of Fundamental Physics. The work of G.~P. is supported by the Deutsche Forschungsgemeinschaft (DFG) under Grant No. PA 1780/2-1. The authors are additionally supported by CompStar, a research networking program of the European Science Foundation, and the Scopes project funded by the Swiss National Science Foundation grant. no. IB7320-110996/1. ~" }, "1003/1003.5220.txt": { "abstract": "We explore the cosmic evolution of massive black hole (MBH) seeds forming within `quasistars' (QSs), accreting black holes embedded within massive hydrostatic gaseous envelopes. These structures could form if the infall of gas into the center of a halo exceeds about 1 $\\msun$ yr$^{-1}$. \u00e6The collapsing gas traps its own radiation and forms a radiation pressure-supported supermassive star. When the core of the supermassive star collapses, the resulting system becomes a quasistar. We use a merger-tree approach to estimate the rate at which supermassive stars might form as a function of redshift, and the statistical properties of the resulting QS and seed black hole populations. \u00e6We relate the triggering of runaway infall to major mergers of gas-rich galaxies, and to a threshold for global gravitational instability, which we link to the angular momentum \u00e6of the host. This is the main parameter of our models. \u00e6Once infall is triggered, its rate is determined by the halo potential; the properties of the resulting supermassive star, QS and seed black hole depend on this rate. \u00e6After the epoch of QSs, we model the growth of MBHs within their hosts in a merger-driven accretion scenario. We compare MBH seeds grown inside quasistars to a seed model that derives from the remnants of the first metal-free stars, and also study the case in which both channels of MBH formation operate simultaneously. \u00e6We find that a limited range of supermassive star/QS/MBH formation efficiencies exists that allows one to reproduce observational constraints. Our models match the density of $z=6$ quasars, the cumulative mass density accreted onto MBHs (according to So\\l tan's argument), and the current mass density of MBHs. \u00e6\u00e6The mass function of QSs peaks at $M_{\\rm QS}\\simeq 10^6 \\msun$, and we calculate the number counts for the {\\it JWST} $2-10\\ \\mu$m band. We find that {\\it JWST} could detect up to several QSs per field at $z\\simeq 5-10$. \u00e6 ", "introduction": "While there is ample evidence that supermassive black holes \u00e6populate the nuclei of most large galaxies and that some black holes with masses exceeding $10^9 \\msun$ formed as early as $z \\ga 6$ \\citep[e.g.,][]{Fanetal2001a,Barthetal2003,Djorgovski2008,Willott2009,Jiang2009}, there is little consensus as to the progenitors of these holes. \u00e6Two schools of thought have persisted since \\cite{Rees1978} first devised a flow chart outlining possible routes of massive black hole (MBH\\footnote{We refer here generically to MBHs when the hole mass is above the limit for black hole formation in today's stars, $\\simeq 50 \\msun$. This definition comprises both seed black holes and supermassive black holes in galaxies and quasars.}) formation. According to one line of argument, supermassive black holes grew from the remnants of an early population of massive stars, the so-called Population III (Pop III), which is believed to have formed in pregalactic minihalos at $z\\ga 20$. \u00e6According to the other, the precursors of supermassive black holes could have formed by the `direct collapse' of large amounts of gas in much larger halos at later times. \u00e6 Each scenario has both positive and negative attributes. \u00e6Although Pop III remnants were unlikely to have been more massive than a few hundred $\\msun$ each, they would have formed relatively early and thus their growth process would have had a considerable head start. \u00e6They could have congregated and merged in the cores of merging minihalos, while simultaneously growing by accretion. \u00e6Models for the growth of supermassive black holes from stellar-mass seeds \\citep{MadauRees2001,VHM,Rhook2006,Monaco2007,Somerville2008,VLN2008} are moderately successful in reproducing the current-day population of supermassive black holes, but have rather more difficulty in producing enough $10^9 \\msun$ holes at $z \\ga 6$ to explain the earliest known quasars \\citep{VR2005,Shapiro2005,VR2006}. \u00e6Additional worries about the Pop III scenario include the possibility that too many of the remnants would have been ejected from the cores of merging halos and that their accretion rates would be depressed by the shallow potential wells of the host mini-halos and heating of the ambient gas by stellar radiation and winds \\citep[and references therein]{Milos2009}. Seed black holes that formed by direct collapse (e.g., at $z \\la 15$) would have had less time to grow, but this would have been partly compensated by their larger initial masses. \u00e6Various direct collapse models for seed formation have been proposed \\citep{LoebRasio1994, Eisenstein1995,HNR1998, BrommLoeb2003,Koushiappas2004,Begelman2006,LN2006,Begelman2010}, but so far only limited attempts have been made to place these models into the context of structure formation theories \\citep[e.g.,][]{VLN2008, Lagos2008,svanwas2010}. \u00e6Because direct collapse models draw only indirectly on star formation lore, there is much less consensus about the initial conditions for direct collapse and the details of how it might have occurred. In a series of recent papers, we have described a sequence of events that we believe represents a plausible route to MBH formation via direct collapse \\citep{Begelman2006,Begelman2008,Begelman2010}. \u00e6In this picture, the main triggering event is the infall of gas into the center of a halo at a rate exceeding about 1 $\\msun$ yr$^{-1}$. \u00e6Such large rates of infall are possible in halos with virial temperatures in excess of about $10^4$ K, which only become common at $z \\la 10-15$. \u00e6Questions remain about the ability of gas to accumulate at such a high rate without too much of it fragmenting into stars, but as we argue below, recent simulations as well as analytic calculations suggest that the importance of fragmentation may have been severely overestimated in past work. \u00e6The collapsing gas traps its own radiation and forms a quasistatic, radiation pressure-supported supermassive star, which burns hydrogen for about a million years while growing to a mass $\\ga 10^6 \\msun$ (Begelman 2010; note that the earlier claim in Begelman et al. 2006 that H-burning is unable to postpone collapse is erroneous and is corrected in the later paper). \u00e6 Because of rotation, the black hole that forms initially probably comprises only a small fraction of the core, with a mass $\\la 10^3 \\msun$, but it grows rapidly at a rate set by the Eddington limit for the massive gaseous envelope. \u00e6A novel feature of our model is the prediction that the envelope swells by a factor of $\\ga 100$ in radius as it absorbs the energy liberated by black hole growth. \u00e6The resulting object, which we have dubbed a `quasistar', (QS) resembles a red giant with a luminosity comparable to a Seyfert nucleus. \u00e6As the black hole grows inside it, its photosphere expands and cools until it hits a minimum temperature associated with the Hayashi track, at which point it disperses, leaving behind the naked seed black hole. In this paper, we use a merger-tree approach to estimate the rate at which supermassive stars might form as a function of redshift, and the statistical properties of the resulting QSs and seed black holes. \u00e6In Section 2 we summarize the properties of supermassive stars, QSs and seed black holes, as a function of time and gas infall rate. \u00e6The existence of runaway infall depends on a threshold for global gravitational instability, which we model as a function of the gas mass and angular momentum following each halo merger --- we discuss these criteria in Section 3. \u00e6Once infall is triggered, its rate is determined by the halo potential; the properties of the resulting supermassive star, QS and seed black hole depend on this rate. \u00e6 To relate the properties of the seed black holes to the observable distributions of active and quiescent black holes at different redshifts, we apply a simple set of rules for their subsequent growth by accretion and mergers. These are described in Section 4, and in Section 5 we present our results. \u00e6We compare our direct collapse seed model to a pure Pop III seed model, and also study the case in which both channels of MBH formation operate simultaneously. \u00e6We summarize our conclusions and discuss the prospects of testing this model observationally in Section 6. ", "conclusions": "We have used a merger-tree approach to estimate the rate at which supermassive ($\\ga 10^6 \\msun$) stars might have formed as a function of redshift, and the statistical properties of the resulting quasistars and seed black holes. \u00e6Key to the formation of supermassive stars is a large gas infall rate ($\\dot M \\simgt 1 \\msun$ yr$^{-1}$), driven by global gravitational instability in the potential of a recently merged halo. \u00e6We use the analyses presented by Begelman et al. (2008) and Begelman (2010) to estimate the properties of the resulting supermassive star, QS and seed black hole as a function of this rate. \u00e6 We relate the properties of the seed black holes to the observable distributions of active and quiescent black holes at different redshifts, by applying a simple set of rules for their subsequent growth by accretion and mergers. We compare our direct collapse seed model to a pure Pop III seed model, and also study the case in which both channels for MBH formation operate simultaneously. \u00e6 Our results can be summarized as follows. \\begin{itemize} %\\item The efficiency of supermassive star/QS formation depends mainly on one free parameter, the angular momentum threshold below which inflows are triggered, $\\lambda_{\\rm thr}$. \\item A limited range of supermassive star/QS/MBH formation efficiencies exists that allows one to reproduce observational constraints (the density of $z=6$ quasars, the cumulative mass density accreted onto MBHs, \u00e6the current mass density of MBHs, reionization). \u00e6These constraints translate into $0.01 \\leqslant \\lambda_{\\rm thr} \\leqslant 0.02$. \\item The mass function of QSs peaks at $M_{\\rm QS}\\simeq 10^6 \\msun$, and decreases almost as a power-law with slope $\\simeq -(1.5-2)$. The more efficient QS and MBH formation is, the steeper the mass function. This is because more seeds form at the highest redshifts when the hosts have relatively low mass and circular velocity. \\item Modeling QS emission as a blackbody at $T_{\\rm QS}=4000$ K, we calculate the number counts for the {\\it JWST} $2-10\\ \\mu$m band. Assuming a sensitivity of 10 nJy at 2$\\mu$m, we find that {\\it JWST} could detect up to several QSs per field at $z\\simeq 5-10$. \\item The redshift of formation of QSs increases with the cosmic bias of their hosts. We therefore expect that the latest-forming QSs should be found in the field, rather than in a high-density environment. \\item The number density of MBHs can be dominated by either \u00e6the descendants of Pop III remnants or \u00e6QS MBHs. In the low-efficiency case ($\\lambda_{\\rm thr}=0.01$) the number density of MBHs is dominated by the Pop III channel, if available; the opposite is true for the high-efficiency case ($\\lambda_{\\rm thr}=0.02$). In contrast to the number density of MBHs, the mass density is always dominated by QS MBHs. \\item \u00e6If the only channel of MBH formation is via QS seeds, then the mass function of MBHs cuts off below $\\log(M_{\\rm BH}/\\msun)=4.5$, where the seed mass function drops as well. If Pop III remnants offer an alternative route to MBH formation, then the mass function is double-peaked, with each peak tracing \u00e6a different seed formation mechanism. \u00e6The peak at $\\log(M_{\\rm BH}/\\msun)\\simeq4.5$ becomes more pronounced, the more efficient QS formation is. \\end{itemize} This model differs from previous proposed mechanisms (e.g. Lodato \\& Natarajan 2006) in various respects. First, gas accumulation in the central regions of protogalaxies is described by global, rather than local, instabilities (e.g., via the Toomre stability criterion). This implies that the spin parameter threshold for collapse cannot be derived by demanding that the Toomre parameter, $Q=\\frac{c_{\\rm s}\\kappa}{\\pi G \\Sigma}$ (where $\\Sigma$ is the surface mass density, $c_{\\rm s}$ is the sound speed, $\\kappa=\\sqrt{2}V_{\\rm h}/R$ is the epicyclic frequency, and $V$ is the circular velocity of the disc), approaches a critical value, $Q_c$, of order unity. We instead assume a fixed spin parameter threshold (cf. BVR2006). Second, we have here relaxed the assumption that large gas infall rates can occur only at zero metallicity to avoid fragmentation. The predicted mass functions have distinctly different shapes (compare Figure 3 with Figure 2 in Lodato \\& Natarajan 2007), and MBH seeds tend to form later in the model presented here ($z\\simeq 5-10$ instead of $z\\simeq 16-18$, when zero metallicity is required). Differences between models can possibly be looked for at the highest redshifts, as at later times the growth of MBHs by accretion and mergers likely washes out dissimilarities. The most direct prediction of this work is shown in Figure~\\ref{JWST}, where we estimate the detectability of QSs in a {\\it JWST} field of view. \u00e6At low metallicities, QSs at redshifts of a few may resemble featureless blackbodies, with colors reminiscent of brown dwarfs. \u00e6Direct redshift measurements may not be feasible, in which case quasistars would have to be identified via their massive hosts. \u00e6Their low numbers --- a consequence of their short lifetimes --- will make it even more challenging to find them with a telescope having a relatively small field of view. \u00e6It is possible that some QSs could have formed in metal-enriched regions, and even at relatively low redshifts ($z \\ga 2$). \u00e6These might be relatively easy to detect, but extremely rare. Other, secondary characteristics of QSs might conceivably aid in their detection, e.g., if accretion onto the black hole deep in the core leads to the formation of a jet that pierces the stellar surface. \u00e6 The shape of the MBH mass function below $\\log(M_{\\rm BH}/\\msun)=4.5$ is an additional diagnostic. \u00e6Direct dynamical measurement of MBH masses at the low-mass end is extremely difficult, due to the necessity of resolving the dynamics of stars within the sphere of influence of the putative MBH, i.e., the region where the Keplerian potential of the MBH dominates over the overall galaxy potential. This region typically subtends much less than an arcsecond in nearby galaxies, \u00e6at or below the resolution limit of existing 8--10 m class telescopes. Future 20--30 m telescopes are likely to increase the sample of low-mass MBHs. \u00e6 Gravitational waves produced during the inspirals of compact objects into MBHs --- extreme-mass-ratio inspirals (EMRIs), in particular --- are expected to provide accurate constraints on the mass function of black holes at low redshift. \u00e6The proposed space-based gravitational wave detector, the {\\it Laser Interferometer Space Antenna} ({\\it LISA}), can probe the mass function in the $10^4M_{\\odot}$--$10^7M_{\\odot}$ range. \\cite{Gair2010} find that with as few as $10$ events, LISA should constrain the slope of the mass function below $\\sim 10^6 \\msun$ to a precision $\\sim0.3$, which is the current level of observational uncertainty in the low-mass slope of the black hole mass function \\citep{GreeneHo}. \u00e6The combination of electromagnetic and gravitational wave observations in the coming years will improve the currently limited constraints on what route, or routes, lead to MBH seed formation. Numerical simulations \\citep[e.g.,][]{Mayer2009} can test how strongly the `bars within bars' inflow mechanism is tied to the angular momentum of the merging galaxies, and either validate our hypothesis that a single parameter $\\lambda_{\\rm thr}$ can describe the efficiency of the cascade process, or indicate that additional parameters influence the formation of quasistars and their nested black holes. \u00e6Even the existence of a second parameter may not have a strong effect on the results: for instance, we found that our results are not very sensitive to varying the threshold gas fraction, as long as it is above $f_{\\rm gas}\\simeq 0.025$. \u00e6It will also be necessary to study the global dynamical behavior of self-gravitating inflows in the inner regions, where the total gravitational potential is expected to approach a Keplerian shape and the dominant means of angular momentum transport may change from one involving large-scale bars ($m=2$) to a flow resembling an eccentric disk ($m=1$: e.g., Hopkins and Quataert 2009). It is even more critical to check how much of the infalling gas actually reaches the small radial scales necessary to build the supermassive star. \u00e6We have suggested that fragmentation and star formation may be less important than previously thought in quenching such inflows. \u00e6Testing these hypotheses will require numerical experiments with high spatial resolution and a large dynamic range. \u00e6" }, "1003/1003.0113_arXiv.txt": { "abstract": "The noise of a device under test (DUT) is measured simultaneously with two instruments, each of which contributes its own background. The average cross power spectral density converges to the DUT power spectral density. This method enables the extraction of the DUT noise spectrum, even if it is significantly lower than the background. After a snapshot on practical experiments, we go through the statistical theory and the choice of the estimator. A few experimental techniques are described, with reference to phase noise and amplitude noise in RF/microwave systems and in photonic systems. The set of applications of this method is wide. The final section gives a short panorama on radio-astronomy, radiometry, quantum optics, thermometry (fundamental and applied), semiconductor technology, metallurgy, etc. This report is intended as a tutorial, as opposed to a report on advanced research, yet addressed to a broad readership: technicians, practitioners, Ph.D. students, academics, and full-time scientists. ", "introduction": "\\label{sec:xsp-introduction} Measuring a device under test (DUT), the observed spectrum contains the DUT noise, which we can call \\emph{signal} because it is the object of the measurement, and the background noise of the instrument. The core of the cross-spectrum measurement method is that we can measure the DUT simultaneously with two equal instruments. Provided that experimental skill and a pinch of good luck guarantee that DUT and instruments are statistically independent, statistics enables to extract the DUT spectrum from the background. \\begin{figure}[b] \\centering\\namedgraphics{scale=0.6}{xsp-correl-basics}{\\textwidth} \\caption{Basics of the cross-spectrum method.} \\label{fig:xsp-correl-basics} \\end{figure} \\begin{figure} \\centering\\namedgraphics{scale=0.8}{mce-sqrt-law}{\\textwidth} \\vspace*{-2em} \\caption{Average and deviation of the cross spectrum $|\\left_m|$, as a function of the number $m$ of averaged realizations of white Gaussian noise. Since the statistical properties of $S_{yx}(f)$ are the same at any frequency, only one point (i.e., one frequency) is shown and the variable $f$ is dropped. The DUT noise is 10 dB lower than the background.} \\label{fig:mce-sqrt-law} \\end{figure} The two-channel measurement can be modeled as the block diagram of Fig.~\\ref{fig:xsp-correl-basics}, where $a(t)$ and $b(t)$ are the background of the two instruments, and $c(t)$ the DUT noise, under the hypothesis that $a(t)$, $b(t)$ and $c(t)$ are statistically independent. Thus, the observed signals are \\begin{align*} x(t)&=c(t)+a(t)\\\\ y(t)&=c(t)+b(t)~. \\end{align*} We are interested in the power spectral density\\footnote{The PSD as a statistical concept will be defined afterwards. Newcomers can provisionally use $S_{yx}(f)=\\frac1TY(f)X^\\ast(f)$, which is the is the readout of the FFT analyzer. $T$ is the measurement time.} (PSD), which is a normalized form of spectrum that expresses the power per unit of bandwidth, denoted with $S(f)$. It will be shown that the average cross-PSD $\\left$ converges to the DUT PSD $S_{cc}(f)$, which is what we want to measure. The idea of the cross-spectrum method is explained in Fig.~\\ref{fig:mce-sqrt-law}. This figure builds from the output of the free-running analyzer, after selecting one frequency ($f_0$). This is a sequence of $|S_{yx}(f_0)|$ called realizations, which we average on contiguous groups of $m$ values $|\\left_m|$. The averages form a (slower) sequence whose statistical properties depend on $m$. So, Fig.~\\ref{fig:mce-sqrt-law} plots the average and the variance of the sequence of averages, as a function of $m$. At small values of $m$, the background is dominant and decreases as $m$ increases. Beyond $m\\approx100$, we observe that $|\\left_m|$ stops decreasing and approaches the value of 0.1 ($-10$ dB), which is the DUT noise in this example. The standard deviation further decreases. The background is dominant below $m\\approx100$. Beyond, the DUT noise shows up and the estimation accuracy increases, as seen from the deviation-to-average ratio. Notice that the choice of $|\\left_m|$ as an estimator of $S_{yx}(f)$ is still arbitrary and will be further discussed. All this report is about how and why the cross-spectrum converges to the DUT noise $S_{cc}(f)$, and about how this fact can be used in the laboratory practice. The scheme of Fig.~\\ref{fig:xsp-correl-basics} is analyzed from the following standpoints \\begin{description} \\item[Normal use.] All the noise processes [$a(t)$, $b(t)$ and $c(t)$] have non-negligible power. We use the statistics to extract $S_{cc}(f)$. \\item[Statistical limit.] In the absence of correlated phenomenon, thus with $c=0$, the average cross spectrum takes a finite nonzero value, limited by the number of averaged realizations. \\item[Hardware limit.] After removing the DUT, a (small) correlated part remain. This phenomenon, due to crosstalk or to other effects, limits the instrument sensitivity. \\end{description} Though the author is inclined to use phase and amplitude noise as the favorite examples (Section \\ref{ssec:xsp-pm-noise} and \\ref{ssec:xsp-am-noise}), the cross-spectrum method is of far more general interest. Examples from a variety of research fields will be discussed in Section~\\ref{ssec:xsp-other-applications}. As a complement to this report, the reader is encouraged to refer to classical textbooks of probability and statistics, among which \\cite{Feller:probability,Papoulis:probability,Cramer:statistics,Davenport-Root:noise} are preferred. ", "conclusions": "" }, "1003/1003.0263_arXiv.txt": { "abstract": "As neutron stars spin-down and contract, the deconfinement phase transition can continue to occur, resulting in energy release(so-called deconfinement heating) in case of the first-order phase transition. The thermal evolution of neutron stars is investigated to combine phase transition and the related energy release self-consistently. We find that the appearance of deconfinement heating during spin-down result in not only the cooling delay but also the increase of surface temperature of stars. For stars characterized by intermediate and weak magnetic field strength, a period of increasing surface temperature could exist. Especially, a sharp jump in surface temperature can be produced as soon as quark matter appears in the core of stars with a weak magnetic field. We think that this may serve as evidence for the existence of deconfinement quark matter. The results show that deconfinement heating facilitates the emergence of such characteristic signature during the thermal evolution process of neutron stars. ", "introduction": " ", "conclusions": "" }, "1003/1003.4992_arXiv.txt": { "abstract": "We present results from the first cosmological simulations which study the onset of primordial, metal-free (population III), cosmic star formation and the transition to the present-day, metal-rich star formation (population II-~I), including molecular (H$_2$, HD, etc.) evolution, tracing the injection of metals by supernov{\\ae} into the surrounding intergalactic medium and following the change in the initial stellar mass function (IMF) according to the metallicity of the corresponding stellar population. Our investigation addresses the role of a wide variety of parameters (critical metallicity for the transition, IMF slope and range, SN/pair-instability SN metal yields, star formation threshold, resolution, etc.) on the metal-enrichment history and the associated transition in the star formation mode. All simulations present common trends. Metal enrichment is very patchy, with rare, unpolluted regions surviving at all redshifts, inducing the simultaneous presence of metal-free and metal-rich star formation regimes. As a result of the rapid pollution within high-density regions due to the first SN/pair-instability SN, local metallicity is quickly boosted above the critical metallicity for the transition. The population III regime lasts for a very short period during the first stages of star formation ($\\sim 10^7\\,\\rm yr$), and its average contribution to the total star formation rate density drops rapidly below $\\sim 10^{-3}-10^{-2}$. ", "introduction": "\\label{sect:intro} The standard paradigm of cosmic structure formation relies on the classical approach of Jeans' theory \\cite[]{Jeans1902} applied to primordial matter fluctuations in the frame of an expanding Universe. Cosmological models for structure formation have been developed since several decades \\cite[e.g.][]{GunnGott1972,Peebles1974,WhiteRees1978} and the overall picture agrees with a ``flat'' Universe where ``cold-dark matter'' is the dominant fraction of matter and the ``cosmological constant'', $\\Lambda$, is the dominant fraction of the cosmological energy density. Baryonic structures arise from in-fall and condensation of gas into dark-matter potential wells. In particular, it seems \\cite[e.g. ][]{Maiolino_et_al_2007} that molecular gas could account for a significant fraction of the dynamical mass of early objects. Recent determinations of the cosmological parameters \\cite[]{wmap7_2010} suggest a present-day expansion rate $H_0\\simeq 70\\rm~km/s/Mpc$ (in units of $\\rm 100~km/s/Mpc$ this parameter becomes $h\\simeq 0.70$), a total-matter density parameter $\\Omega_{0m}\\simeq 0.272$, with ``baryonic'' component $\\Omega_{0b}\\simeq 0.0456$, and a cosmological-constant density parameter $\\Omega_{0\\Lambda}\\simeq 0.728$. The primordial power spectrum of perturbation is well fitted by a power law with index $n\\simeq 0.96$ and normalization via mass variance within $8$~Mpc/$h$ radius $\\sigma_8\\simeq 0.8$. As a reference, it is common to define the standard $\\Lambda$CDM model the one with the following parameters: $H_0=\\rm 70~km/s/Mpc$, $\\Omega_{0m}=0.3$, $\\Omega_{0b}=0.04$, $\\Omega_{0\\Lambda}=0.7$, $\\Omega_{0tot}=1.0$, $n=1$, $\\sigma_8=0.9$. In the frame of cosmic evolution, it is believed that structure formation takes place from the growth of primordial fluctuations in matter density. These would contract and collapse allowing gas cooling and the subsequent build-up of stars. The determination of the properties of early stars and their effects on the following baryonic-structure formation episodes is still a problem under debate \\cite[for a complete review see e.g.][]{Ciardi_Ferrara_2005}. Nevertheless, it is reasonably well established that the very first generation of stars should be characterized by massive objects with typical masses much larger than the presently observed ones \\cite[][]{SS1953, WW1995, Larson1998, Chiosi2000, HegerWoosley2002,HegerWoosley2008}. These primordial stars (population III stars) are formed out of a pristine environment, where the cooling agents are limited to primordial, H-based molecules only, i.e. H$_2$ and HD, which are able to cool the gas down to temperatures of $\\sim 10^2~\\rm K$. Therefore, the mass of primordial stars should be relatively large and their spectrum, biased towards such large objects, is commonly referred to as ``top-heavy'' initial mass function (IMF) \\cite[][]{Larson1998}. These features imply very short lifetimes (up to $\\sim 10^6$ years only) and final death mostly into black holes \\cite[][]{HegerWoosley2002}. The only mass range where primordial stars can explode as pair-instability supernov{\\ae} (PISN) and pollute the surrounding medium is [140-260]~M$_\\odot$ \\cite[][]{HegerWoosley2002}. Nevertheless, also at masses of $\\sim 100-140~\\rm M_\\odot$ there can be some mass loss before collapse (pulsational pair SN). The key uncertainty here is primary nitrogen production and the dredge up of carbon and oxygen. In particular, if the stellar atmosphere is highly CNO enhanced there may be substantial mass loss, but zero metallicity should still be a good first approximation for such stellar flows (S.~Woosley, private communication). \\\\ Despite the many uncertainties on their characteristics, population III (popIII) stars have an important impact on the evolution of the intergalactic medium (IGM), since they initiate the metal pollution of the IGM, with consequent change of its chemical composition and cooling properties (chemical feedback). Therefore, star formation events in enriched regions will happen in completely different conditions, because metals allow further cooling and fragmentation to smaller scales. This results in an initial stellar-mass function peaked at lower masses, similar to the nowadays observed Salpeter-like IMF \\cite[]{Salpeter1955} for population II-I (popII-I) stars.\\\\ A very debated issue is the understanding of the transition from the primordial popIII star formation regime to the standard popII-I regime. In this respect, there are evidences for the existence of a critical metallicity, $Z_{crit}$, at which the modalities of star formation allow such transition \\cite[][]{Bromm_et_al_2001,Schneider_et_al_2002}: $Z_{crit}$ is the metallicity at which the metal cooling function dominates over the molecular one. In this case, the fragmentation process becomes highly enhanced, but the exact value is not well-established, yet. Different studies suggest discrepant values with $Z_{crit}$ varying between $\\sim 10^{-6}~Z_\\odot$ \\cite[e.g.][]{Schneider_et_al_2006} and $\\sim 10^{-3}~Z_\\odot$ \\cite[e.g.][]{Bromm_Loeb_2003}\\footnote{ We adopt $Z_\\odot\\simeq 0.0201$ \\cite[][]{Anders_Grevesse_1989, Grevesse_Sauval_1998}. See, however, \\cite{Asplund_et_al_2009} for a recently updated value of $Z_\\odot\\simeq 0.0134$}. \\\\ The main uncertainty is the presence of dust at high redshift. To our knowledge, dust is produced and injected in the ISM mainly by the low--mass stars in the AGB phase. This would imply no dust production at very-high redshift, due to ``long'' stellar lifetimes: a 8~M$_\\odot$ star has a life of $\\sim 0.1$ Gyr, comparable with the age of Universe at $z\\sim 15-20$. None the less, the presence of large amounts of dust and heavy elements has been detected at moderately-high redshift, i.e. at $z>6$, when the Universe is younger than $\\sim 1~\\rm Gyr$ \\cite[e.g. ][]{Bertoldi2003,Maiolino_et_al_2004}. This suggests that dust production must have occurred primarily in the ejecta of supernova explosions, which are the final fate of massive, short-lived stars. If, indeed, supernov{\\ae} could be able to induce dust production then star formation in the early Universe and the level of $Z_{crit}$ would be strongly influenced, as well. Enrichment by dust would not impact the thermal structure of the IGM, but its presence would alter the whole star formation process. Indeed, when the metallicity of star forming regions is still below $\\sim 10^{-6}\\,Z_\\odot$ the only relevant coolants are molecules, mainly H$_2$ and HD, while above $\\sim 10^{-4}\\,Z_\\odot - 10^{-3}\\,Z_\\odot$ gas cooling is fully dominated by metal fine-structure transitions \\cite[e.g. ][]{Maio2007} and cloud fragmentation can happen down to sub-solar scales \\cite[][]{Bromm_et_al_2001,Schneider_et_al_2002,SFS2004}. In between, cooling capabilities depends mainly on the amount of metals depleted onto dust grains. The mass content locked into dust grains, almost independently from the exact mass of SNII/PISN progenitor, reaches $2\\%-5\\%$ of the parent mass for type-II SN \\cite[][]{Kozasa1991,TodiniFerrara2001,Nozawa_et_al_2003,BianchiSchneider2007}, and $15\\%-30\\%$ for PISN \\cite[][]{Nozawa_et_al_2003,SFS2004}, in a few hundred days after the beginning of the explosion. Despite the huge difference between the dust mass fractions deriving from SNII and PISN, the dust-to-metal mass ratios (or ``depletion factors'') turn out to be $\\sim 0.3 - 0.7$ in both cases and, for metal poor SNII progenitors with masses of $\\rm 25\\,M_\\odot - 30\\, M_\\odot$, it is up to $\\sim 1$. The natural conclusion is that first massive stars could spread in the early Universe a lot of dust during their final phases (as either SNII or PISN). \\\\ In the present work, we aim at investigating the birth of the first stars, the following cosmic metal enrichment from their explosive death and the transition to the standard, presently observed, star formation regime. We focus on how such transition and the associated features are affected by different choices of $Z_{crit}$, IMF, post-supernova metal yields, star formation density thresholds, box dimension and resolution. \\\\ Throughout the paper we will refer to ``population III regime'' and ``population II-I regime'' when the metallicity is below or above the critical level, respectively. In addition, we will adopt the standard $\\Lambda$CDM cosmological model.\\\\ The paper is organized as follows: we describe the code and the simulations in Sect. \\ref{sect:sims}, results are given is Sect. \\ref{sect:results}, and parameter dependence is addressed in Sect. \\ref{sect:params} (\\ref{sect:imf}, \\ref{sect:threshold}, and \\ref{sect:largerboxes}). We discuss and conclude in Sect. \\ref{sect:disc}. ", "conclusions": "\\label{sect:disc} In this work, we perform numerical simulations of early structure formation, including both primordial, molecular evolution and metal enrichment from stellar death, to study the transition from an early, massive star formation mode to a more standard one, regulated by the gas metallicity. We follow \\cite[see details in][]{Maio2007,TBDM2007} the abundances of e$^-$, H, H$^+$, H$^-$, He, He$^+$, He$^{++}$, H$_2$, H$_2^+$, D, D$^+$, HD, HeH$^+$, C, O, Mg, S, Si, Fe, and use different initial stellar-mass functions, metal yields, and critical metallicities for the transition from a popIII to a popII star formation mode, $Z_{crit}$. At the present, our work is the only one dealing with detailed chemical evolution, from early molecule creation to the later stages of star formation, and, at the same time, allowing to trace simultaneous, different stellar populations, according to the underlying metallicity. \\\\ Indeed, early structure formation can be accurately modeled only with a proper treatment of both the chemistry of primordial molecules and the metal enrichment. In fact, the main influence of chemical evolution on the following generations of structures is via metal pollution (chemical feedback). This event can completely alter the cooling properties of the gas and thus the modalities of star formation, inducing a transition from a top-heavy to a standard IMF. The transition is believed to happen when the gas is enriched above $Z_{crit}$, which allows fragmentation below the typical scales determined by primordial molecular cooling. Because of our ignorance of the features of early dust formation and the lack of precision of many atomic and molecular data, the exact determination of $Z_{crit}$ is still elusive; reasonable values should range between $10^{-6}\\,Z_\\odot$ and $10^{-3}\\,Z_\\odot$ \\cite[][]{Bromm_et_al_2001,Schneider_et_al_2002}. \\\\ In our simulations we assume four different values of $Z_{crit}$ in the above range: $10^{-6}, 10^{-5}, 10^{-4}$, and $10^{-3}\\,Z_\\odot$. From our investigation, some common features emerge. In general, metal pollution, independently from the parameters adopted in the simulations, is very patchy, with excursions of orders of magnitudes at all redshifts. This is consistent with any simulation including metal enrichment \\cite[][]{Raiteri1996,Gnedin1998,Mosconi2001,Lia2002a,Lia2002b,KawataGibson2003,Kobayashi2004,RicottiOstriker2004,Cecilia_et_al_2005,TBDM2007,Oppenheimer_et_al_2010}. In addition, the gas is easily enriched above $Z_{crit}$. For this reason, the average contribution from pristine, metal-free (or $Z < Z_{crit}$) stars to the total cosmic star formation density is dominant only in the very early phases of structure formation, while it drops below $\\sim 10^{-3}$ quite rapidly, after the explosion of the first pair-instability supernov{\\ae} and their metal ejection. In fact, PISN explosions which follow the death of the first, metal-free or very-metal-poor stars, are the main responsibles for enriching the surrounding medium up to a minimum level of $\\sim 10^{-4}\\,Z_\\odot$. This means that nearby star forming regions have a very high probability of being polluted above $Z_{crit}$ (as seen also in Fig. \\ref{fig:Z_z}), while popIII star formation can still occur farther away, in rare, isolated regions with pristine or low-metallicity gas (e.g. Fig. \\ref{fig:maps_metallicity_comparison_2col}). None the less, this would not be the dominant star formation regime. \\\\ Our findings hold regardless of the value of $Z_{crit}$, the popIII IMF adopted, and the numerical parameters involved in modeling star formation (as shown in Sect. \\ref{sect:imf}, \\ref{sect:threshold}, and \\ref{sect:largerboxes}). Differences are found if different IMF mass ranges for primordial stars are used. Because of the short life and the high metal yields of early, massive SN, the popIII regime contributes, in any case, only slightly to the global SFR \\cite[as expected by e.g. ][]{RicottiOstriker2004}, since the early pollution events quickly raise $Z$ above $Z_{crit}$, independently from the detailed prescriptions. \\\\ The simulations were performed using a standard $\\Lambda$CDM cosmology, but slightly different parameters would not change the general picture. We have checked this, by running simulations with $\\Omega_{0m}=0.26$, $\\Omega_{0\\Lambda}=0.74$, $\\Omega_{0b}=0.0441$, $h=0.72$, $\\sigma_8=0.796$, and $n=0.96$ \\cite[WMAP5 data, ][]{wmap5_2008} and found that the same results hold, albeit shifted by a $\\Delta z\\sim$ of a few to lower redshift (because of the smaller $\\Omega_{0m}$). \\\\ One last comment about the role of dust production from SNII and/or PISN: one of the main uncertainties in determining $Z_{crit}$. According to our findings, given the strength and rapidity of metal enrichment (see e.g. Figs. \\ref{fig:Z_z}, \\ref{fig:Z_z_largerbox5} and \\ref{fig:Z_z_largerbox10}), details about dust and its impact on $Z_{crit}$ are not relevant for the general process and for our understanding of the transition from primordial, popIII regime to present-day-like, popII-I regime (even if they could be of some interest on very local scales). The entire process, in fact, is dominated by metal pollution and the strong yields of early, massive stars. \\\\ We have to stress some {\\it caveats}, though. Ejection of particles into the IGM is an unknown process. We have assumed winds originated from stars (kinetic feedback), but different mechanisms (like gas stripping, shocks, thermal heating from stellar radiation, etc.) could play a role, as well, mostly at high redshifts, when objects are small and can easily loose part of their baryonic content. In addition, diffusion and conduction will probably alter the smoothness of metal and molecule distribution, but these phenomena have not been extensively studied, yet, and probably will depend on many parameters: e.g., the way metal or gas particles are ejected and mixed, how they are transferred away from the production sites, how strong is the efficiency of such processes, just to mention a few. Some attempts to address such issues have been done \\cite[][]{Spitzer1962,CowieMcKee1977,Brookshaw1985,Sarazin1988,Monaghan1992,ClearyMonaghan1999,KlessenLin2003,Jubelgas_et_al_2004,Monaghan_et_al_2005,Wadsley_et_al_2008,Greif2009}, but much more realistic and detailed analyses are still strongly needed. \\\\ To conclude, we have followed the structure formation process from very early times to first star formation and subsequent metal pollution. In the 1~Mpc side box simulations we have seen that, after $\\sim 2\\times 10^8\\,\\rm yr$, molecular evolution leads the very first bursts of star formation (popIII), but metal enrichment is extremely fast (Figs. \\ref{fig:maps_metallicity_zcrit}, \\ref{fig:maps_metallicity_comparison}, and \\ref{fig:maps_metallicity_comparison_2col}) in inducing the nowadays observed star formation mode (popII-I). In fact, we observe a steep increase of $Z$, with local values rapidly reaching and overtaking $Z_{crit}$ (Fig. \\ref{fig:Z_z}). Metal pollution proceeds from the densest cores of star formation outwards, because of supernova ejections from high-density to lower-density environments (Figs. \\ref{fig:maps_metallicity_comparison}, and \\ref{fig:maps_metallicity_comparison_2col}). Rare, unpolluted regions can still survive, determining the simultaneous presence of the two star formation regimes, and $Z_{crit}$ can affect the level of residual popIII star formation. As a result of this rapid pollution, we find that the average contribution of the popIII component to the total star formation rate density is of a few times $\\sim 10^{-4}-10^{-3}$ (with a maximum of $\\sim 10^{-2}-10^{-1}$) at $z \\sim 11$ (Figs. \\ref{fig:sfr_ratio},\\ref{fig:imfrange}, and \\ref{fig:imfyields}). This general picture is preserved regardless of the precise value of the metallicity threshold $Z_{crit}$ the slope of the popIII IMF, and their yields, but is quite sensitive to the popIII mass range. The change of the critical density threshold for star formation influences little these conclusions, but can affect the onset and the decrement of the popIII contribution to the total SFR (e.g. Fig. \\ref{fig:threshold}). We have found similar results in larger-box simulations, either in the metal enrichment features (Figs. \\ref{fig:Z_z_largerbox5} and \\ref{fig:Z_z_largerbox10}) and in the star formation behaviour (Fig. \\ref{fig:sfr_resolution_ratio})." }, "1003/1003.4440_arXiv.txt": { "abstract": "{It is now more than 40 years since the discovery of gamma-ray bursts (GRBs) and in the last two decades there has been major progress in the observations of bursts, the afterglows and their host galaxies. This recent progress has been fueled by the ability of gamma-ray telescopes to quickly localise GRBs and the rapid follow-up observations with multi-wavelength instruments in space and on the ground. A total of 674 GRBs have been localised to date using the coded aperture masks of the four gamma-ray missions, \\textit{BeppoSAX}, \\textit{HETE II}, \\textit{INTEGRAL} and \\textit{Swift}. As a result there are now high quality observations of more than 100 GRBs, including afterglows and host galaxies, revealing the richness and progress in this field. The observations of GRBs cover more than 20 orders of magnitude in energy, from $10^{-5}$\\,eV to $10^{15}$ eV and also in two non-electromagnetic channels, neutrinos and gravitational waves. However the continuation of progress relies on space based instruments to detect and rapidly localise GRBs and distribute the coordinates.} \\FullConference{The Extreme sky: Sampling the Universe above 10 keV - extremesky2009,\\\\ October 13-17, 2009\\\\ Otranto (Lecce) Italy} \\begin{document} ", "introduction": "Gamma-ray bursts (GRBs) are powerful flashes of gamma-rays appearing in the sky with durations ranging from a fraction of a second to over 1000 seconds. The isotropic equivalent energy radiated is between $10^{48}$ and $10^{55}$ ergs. The gamma-rays are thought to originate from a highly relativistic outflow with Lorentz factor $\\Gamma$\\,$>$100 that is ejected from a central source. The content of the jet, the level of magnetisation and the mechanism generating the gamma-rays are highly uncertain. In this short overview we cover a small number of topics, such as durations, spectral lags, high-energy emission and afterglows, and refer the reader to much longer recent review articles \\cite{gehrels:2009,meszaros:2006,nakar:2007,piran:2004,zhang:2004} and the excellent book on GRBs by Vedrenne and Atteia (2009) \\cite{vedrenne:2009}. It is interesting to compare the results obtained with recent missions between July 1996 and December 2009 (Table~\\ref{table:ags}). There is a total of 674 GRBs detected by these 4 missions and most of them are better localised with X-ray measurements of the afterglow. The X-ray, optical and radio afterglow detections are listed in Table~\\ref{table:ags} and the data are taken from the webpage maintained by Jochen Greiner\\footnote{http://www.mpe.mpg.de/$\\sim$jcg/grbgen.html.}. It should be noted the \\textit{Swift} has detected the largest number of GRBs and has the highest percentage of detections of X-ray and optical afterglows because it has an onboard X-ray detector and an optical/UV telescope that immediately slews to the position of the burst and provides accurate coordinates for ground-based observations. \\textit{Swift} has set the standard in this field. New GRB space missions should have either an X-ray telescope or infrared telescope on board to localise the GRB and enable rapid follow-up observations with other space and ground-based telescopes. \\begin{table}[b] \\begin{center} \\caption{Afterglow detections for GRBs localised by recent $\\gamma$-ray missions between July 1996 and December 2009.} \\label{table:ags} \\begin{tabular}{@{}l c c c c @{}} \\hline\\hline & \\textbf{\\textit{BeppoSAX}} & \\textbf{HETE II} & \\textbf{\\textit{INTEGRAL}} & \\textbf{\\textit{Swift}} \\\\ \\hline \\textbf{GRBs} & 55 & 79 & \\textbf{66} & 474 \\\\ \\textbf{X-ray} & 31 & 19 & \\textbf{34} & 397 \\\\ \\textbf{Optical} & 17 & 30 & \\textbf{21} & 247 \\\\ \\textbf{Radio} & 11 & 8 & \\textbf{8} & 34 \\\\ \\hline \\end{tabular} \\end{center} \\end{table} ", "conclusions": "" }, "1003/1003.1265_arXiv.txt": { "abstract": "The curvature of a relativistic blast wave implies that its emission arrives to observers with a spread in time. This effect is believed to wash out fast variability in the light curves of GRB afterglows. We note that the spreading effect is reduced if emission is anisotropic in the rest-frame of the blast wave (i.e. if emission is limb-brightened or limb-darkened). In particular, synchrotron emission is almost certainly anisotropic, and may be strongly anisotropic, depending on details of electron acceleration in the blast wave. Anisotropic afterglows can display fast and strong variability at high frequencies (above the `fast-cooling' frequency). This may explain the existence of bizarre features in the X-ray afterglows of GRBs, such as sudden drops and flares. We also note that a moderate anisotropy can significantly delay the `jet break' in the light curve, which makes it harder to detect. ", "introduction": "GRB afterglows are likely produced by relativistic blast waves propagating from the center of the explosion. This model is, however, challenged by recent observations. In particular, the {\\it Swift} satellite revealed several puzzling features in the X-ray afterglow. It observed an early plateau stage and flares with fast rise and decay times (Nousek et al. 2006; Burrows et al. 2005). Less frequent but even more bizarre are sudden drops in the X-ray light curve (as steep as $\\tobs^{-10}$ in GRB~070110, Troja et al. 2007). These behaviors are inconsistent with the standard model of afterglow production. Can the emission from the forward or reverse shock of the blast wave show strong variations on timescales $\\dtobs\\ll\\tobs$? It is usually argued that this is impossible: the spherical curvature of the emitting surface (of radius $R$ and Lorentz factor $\\Gamma$) implies a spread in arrival times of its emission, which washes out variability on timescales shorter than \\begin{equation} \\label{eq:Dtobs} \\Dtobs = \\frac{R}{2c\\Gamma^2}. \\end{equation} For a relativistic blast wave, this duration is comparable to the observed time passed since the beginning of the explosion, $\\Dtobs\\sim\\tobs$. This appears to prohibit any rapid and strong variations in the light curve (see Ioka et al. 2005 for discussion). Therefore, the observed fast variability in afterglows is usually associated with additional emission from radii much smaller than the blast-wave radius. This model invokes a late activity of the central engine (Zhang et al. 2006). The material ejected at large $\\tobs$ and emitting at radii $R\\ll\\Gamma^2\\tobs c$ will have $\\Dtobs\\ll\\tobs$ and can produce flares with $\\dtobs\\ll\\tobs$. Note however that (i) it is unclear in this model why the observed flares have the approximately universal $\\dtobs/\\tobs\\sim 0.1$ (Chincarini et al. 2007; Lazzati \\& Perna 2007), (ii) the very steep drops at the end of some plateaus can hardly be explained by this model unless it assumes that the entire plateau is produced at small radii inside the ejecta and the emission from the blast wave is negligible (Kumar, Narayan \\& Johnson 2008). Another difficulty for GRB theory is that many afterglows lack the predicted `jet breaks' (Burrows \\& Racusin 2006; Sato et al. 2007): only a small fraction of afterglow light curves show a clear achromatic break that is expected from jets (Willingale et al. 2007).\\footnote{Many afterglows show chromatic breaks, which occur either in the X-ray or in the optical, but not in both bands.} Some bursts show X-ray light curves extending for tens to hundreds of days with a constant temporal slope (Grupe et al. 2007). The interpretation of these observations is difficult and often leads one to assume large jet opening angles, implying in some cases extremely high energy for the explosion (Shady et al. 2007). An implicit assumption in the general discussion of these puzzling features is that the emission is isotropic in the rest frame of the relativistically moving source (however, see Lyutikov 2006). In this paper, we discuss the effects of a possible anisotropy and suggest that they can help explain observations. In \\refsec{sec:response} we write down a general formula for the observed flux from a flashing sphere when the emission is anisotropic in the source rest frame. In \\refsec{sec:effects} we list the consequences of anisotropy for the curvature effect, the jet break in the afterglow light curve and the size of the radio image of the blast wave. In \\refsec{sec:meca} we consider the standard radiative mechanism of afterglows -- synchrotron emission -- and discuss its anisotropy. The results are summarized in \\refsec{sec:discussion}. ", "conclusions": "\\label{sec:discussion} The usual assumption of isotropic emission in the rest frame of the blast wave is likely to be invalid. Even the standard synchrotron model with isotropic electron distribution produces anisotropic, limb-darkened radiation (\\refsec{sec:meca}). This fact is a consequence of the preferential orientation of the magnetic field in the blast wave. Strong limb-brightening is also possible if the radiating electrons are preferentially accelerated along the magnetic field. Anisotropy may resolve a few puzzles encountered in afterglow modeling: \\\\ (i) The usual argument that the curvature effect filters out fast variability, prohibiting strong variations in the light curve on timescales $\\Delta t <\\tau=R/2\\Gamma^2 c$, is not valid for anisotropic emission. An anisotropic variable spherical source can produce fast changes in the light curve, similar to observed bizarre features in GRB afterglows. This result holds for both limb-darkened and limb-brightened types of anisotropy. It suggests that the X-ray flares observed by {\\it Swift} with $\\Delta t/\\tobs\\la 0.1$ do not necessarily imply an additional component of internal origin. Instead they may be produced, e.g. by the reverse shock in the blast wave, whose emission may suddenly brighten and weaken as the reverse shock propagates into the inhomogeneous ejecta of the explosion. This model may also explain sudden steep drops in the afterglow light curve as observed in GRB~070110 (Troja et al. 2008). This explanation assumes that the X-ray radiating particles are cooling fast compared with the jet expansion timescale, as slow cooling would suppress short time-scale variations of the source luminosity. Examples of such short times-cale behaviors are given by the toy model in \\reffig{fig:lc}. It shows the synchrotron emission produced by a thin shell with Lorentz factor $\\Gamma(R)=\\Gamma_0=300$ at $RR_\\mathrm{dec}$. This approximately describes a blast wave decelerating in a uniform medium. The shell is assumed to radiate with bolometric power proportional to the dissipation rate in the blast wave, which gives $\\dot{E}(R)=\\dot{E}_0\\left(R/R_\\mathrm{dec}\\right)^2$ at $RR_\\mathrm{dec}$. A realistic blast wave has two shocks -- forward and reverse -- and both can produce a long-lived afterglow. Our toy model may describe the emission from either shock, although it is very much simplified. To illustrate the curvature effect on variability, we add two features: a sudden brief increase in $\\dot{E}(R)$ at $R=3\\,R_\\mathrm{dec}$ (which simulates a flare) and the abrupt cutoff of $\\dot{E}(R)$ at $15\\,R_\\mathrm{dec}$. For comparison, we show the light curves produced for three cases: isotropic emission in the rest frame of the shell, limb-brightened emission and limb-darkened emission described in Section~4. \\begin{figure} \\centerline{\\includegraphics[width=0.5\\textwidth]{newfig2.eps}} \\caption{ Bolometric light curve for a toy afterglow model (see text). The result is plotted for three cases: isotropic emission in the rest frame of the blast wave $A(\\theta)=1$ (solid curve), limb-brightened synchrotron emission $A_1(\\theta)$ with $a=0.03$ (dotted curve; see Section~4 for the description of the synchrotron model) and limb-darkened synchrotron emission $A_2(\\theta)$ with $a=0.03$ (dashed curve). } \\label{fig:lc} \\end{figure} (ii) If a relativistic source is limb-darkened, most of its emission in the fixed lab frame is confined within an angle {\\it smaller} than $\\Gamma^{-1}$. This effect suggests a possible explanation for the lack of jet-break detections in GRBs, as the increased beaming significantly delays the jet break in the observed light curve (Section~3.2). We also discussed in Section~3.3 the consequences of such anisotropy for the apparent size of the radio afterglow source. Although the strong limb-darkening appears to be impossible for standard synchrotron afterglows, it may be possible for a different radiative mechanism. For example, limb-darkening may be expected for the jitter mechanism (Medvedev \\& Loeb 1999), as the electrons are preferentially accelerated perpendicular to the shock plane and radiate preferentially in the radial direction. While this paper was focused on afterglow, the source of prompt GRB emission may also be intrinsically anisotropic. This may impact models that propose the curvature effect to control the steep X-ray decay at the end of the prompt emission (see e.g. Genet \\& Granot 2009; Zhang et al. 2009). The effect can be seen in Figure~4. Anisotropy of the prompt emission may also change the optical depth of the source to high-energy photons, $\\tau_{\\gamma\\gamma}$, as the cross section for $\\gamma\\gamma$ reaction strongly depends on the angle between photons. This may affect the constraints on the Lorentz factor of the jet that are inferred from $\\tau_{\\gamma\\gamma}<1$. The effect is especially strong for emission without front-back symmetry in the source frame; such asymmetric emission would be a more radical assumption compared with the ordinary limb-brightening or limb-darkening considered in this paper." }, "1003/1003.4395_arXiv.txt": { "abstract": "We show how the motion of cosmic superstrings in extra dimensions can modify the gravitational wave signal from cusps. Additional dimensions both round off cusps, as well as reducing the probability of their formation, and thus give a significant dimension dependent damping of the gravitational waves. We look at the implication of this effect for LIGO and LISA, as well as commenting on more general frequency bands. ", "introduction": " ", "conclusions": "" }, "1003/1003.4676_arXiv.txt": { "abstract": "Active galactic nuclei present continuum and line emission. The emission lines are originated by gas located close to the central super-massive black hole. Some of these lines are broad, and would be produced in a small region called broad-line region. This region could be formed by clouds surrounding the central black hole. In this work, we study the interaction of such clouds with the base of the jets in active galactic nuclei, and we compute the produced high-energy emission. We focus on sources with low luminosities in the inner jet regions, to avoid strong gamma-ray absorption. We find that the resulting high-energy radiation may be significant in Centaurus A. Also, this phenomenon might be behind the variable $\\gamma$-ray emission detected in M87, if very large dark clouds are present. The detection of jet-cloud interactions in active galactic nuclei would give information on the properties of the jet base and the very central regions. ", "introduction": "Active galactic nuclei (AGN) are extragalactic sources composed by a super-massive black hole (SMBH), an accretion disk and bipolar relativistic jets. Some AGNs present continuum emission in the whole electromagnetic spectrum, from radio to $\\gamma$-rays. Besides the continuum radiation, AGNs also have optical and ultra-violet line emission. Some of these lines are broad, emitted by gas moving at velocities $v_{\\rm g} > 1000$~km~s$^{-1}$ and located in a region close ($d \\sim 10^{17}$~cm) to the SMBH. The structure of the matter in the broad line region (BLR) is not well known but some models assume that the gas could be clumpy. Dense clouds, confined by the hot external medium or by magnetic fields, would be ionized by photons from the accretion disk producing the emission lines broadened by the movement of the clouds around the SMBH. In the particular case of Faranoff Riley (FR) I galaxies, where the accretion disks have low luminosities, the photoionization of the clouds will be inefficient to produce lines and the clouds might be dark. Centaurus A (Cen A) and M87 are the closest AGNs, located at distances of $\\sim 4$ and $\\sim 16$~Mpc, respectively. These AGNs are classified as FR I radio sources and in the case of Cen A the nuclear region is obscured by a dense torus of gas and dust. Although the BLRs of Cen A and M87 have not been detected~\\cite{Alex}, clouds with similar characteristic to those detected in FR II AGNs may surround the SMBH~\\cite{wang}. We are interested in the high-energy emission produced by the interaction of these possibly dark clouds with the jets of the AGN. We focus here on Cen A and M87, since their moderate accretion rate would imply reduced photon-photon opacities in the interaction region, allowing $\\gamma$-rays to escape. Assuming standard parameters for the clouds and jets, we study the main physical processes that take place as a consequence of the interaction, and calculate the expected high-energy emission. ", "conclusions": "The total luminosity produced by jet-cloud interactions depends on the number of clouds inside the jet, each one producing a spectral energy distribution (SED) with similar characteristics and luminosity levels to those shown in Fig. 2. The number of clouds inside the jet is $N_{\\rm cj} = f V_{\\rm j}/V_{\\rm c}$, where $f$ is the filling factor of dark clouds, and $V_{\\rm j}$ and $V_{\\rm c}$ are the jet and the cloud volume, respectively. Considering $V_{\\rm j}$ up to $z \\sim 10^{16}$~cm and $f \\sim 5\\times10^{-6}$ (in FR II galaxies $f \\sim 10^{-6}$), we obtain $N_{\\rm cj} \\sim 10$. In the whole sphere of size $\\sim 10^{16}$~cm the number of clouds is $\\sim 5\\times10^3$ for the considered value of $f$. The simultaneous interaction of $\\sim$ 10 clouds with the jet will produce more luminosity than the one produced by only one interaction. If all the clouds have the same properties (i.e. $n_{\\rm c}$, $R_{\\rm c}$ and $v_{\\rm c}$) and are located at $z \\sim z_{\\rm int}$ in the jet, then the contribution of 10 clouds will produce a SED with a similar appearence than that shown in Fig.~\\ref{f2}, but with a luminosity $\\sim$ 10 times larger, being now detectable by HESS and \\emph{Fermi} telescopes in the case of Cen A, as is shown in Fig.~\\ref{f3}. The emission detected by these instruments from Cen~A is larger than the luminosity predicted by our model. However, if clouds are larger than $10^{13}$~cm, the luminosity level detected by HESS\\cite{hess_CenA} and \\emph{Fermi}\\cite{fermi_CenA} could be achieved. A more detailed calculation of the emission produced by many clouds interacting with the jet at different $z$ will be presented in a future work\\cite{Araudo}. In the case of M87, the jet base is expected to be at $z_0 \\sim 50 R_{\\rm Sch} \\sim 4\\times10^{16}$~cm. At such a height on the jet, the jet-cloud interaction will be inefficient producing high-energy emission due to the small cloud to jet section ratio. In order to obtain a detectable luminosity, clouds with a radius $\\sim 10^{14}$~cm would be necessary. In the case of a very big cloud entering the jet close to $z_0$ in M87, the interaction might produce the variable luminosity detected by HESS\\cite{hess_m87}. \\begin{figure}[] \\centerline{\\psfig{file=fig3.ps,width=7cm, angle=270}} \\vspace*{8pt} \\caption{The same as in Fig.~\\ref{f2} but for $\\sim$~10 clouds.\\label{f3}} \\end{figure}" }, "1003/1003.3110_arXiv.txt": { "abstract": "{This is the first paper of a series in which we will attempt to put constraints on the flattening of dark halos in disk galaxies. We observe for this purpose the \\HI\\ in edge-on galaxies, where it is in principle possible to measure the force field in the halo vertically and radially from gas layer flaring and rotation curve decomposition respectively. In this paper, we define a sample of 8 \\HI\\ rich late-type galaxies suitable for this purpose and present the \\HI\\ observations. ", "introduction": "Since the 1970's, it has been known that the curvature of the universe is remarkably flat. This implies that the ratio of the total density to critical density is $\\Omega_{tot} \\approx 1$. At this time, it was also known from measurements of stars and gas that the mass density of luminous matter is $\\Omega_{lum} \\lesim 0.005$ --- less than $0.5\\%$ of that required for a flat $\\Omega_{tot}\\sim$1 universe. Indeed, the missing matter controversy had begun in the 1930's when \\citet{oort1932}, following \\citet{kapteyn1922}, and \\citet{zwicky1937} independently found evidence for vast amounts of unseen matter on different scales. Oort's analysis of the spatial and velocity distribution of stars in the Solar neighbourhood concluded that luminous stars comprised approximately half the total mass indicated by their motion, assuming gravititational equilibrium. Zwicky's analysis of the velocity dispersions of rich clusters found that approximately $90-99\\%$ of the mass was unseen, if the systems were gravitationally bound. Later on, the inventory of the mass in the Solar neighbourhood showed that the dark matter in the galactic disk at the solar galactocentric radius was mostly fainter stars and interstellar gas \\citep[e.g.][]{hf00}. Zwicky's value of the velocity dispersion in the Coma cluster is close to the current one \\citep[e.g.][]{cd96}, and, despite the presence of hot X-ray emitting gas, there must be dark matter to cause $\\Omega_{matter} \\sim 0.2 - 0.3$. \\begin{table*}[t] % \\centering \\caption[HI observations]{HI observations} \\label{tab:obs-HI} \\begin{small} \\begin{tabular}{lllccllc} \\hline Galaxy & Other & Date & Telescope & Array & Project & Observer & Integration \\\\ & name &&&& \\ \\ ID && time (hr) \\\\ \\hline\\hline ESO074-G015 & IC5052 & 11-12 FEB 2001 & ATCA & 375 & C934 & Ryan & 1.67 \\\\ ESO074-G015 & IC5052 & 13-14 APR 2001 & ATCA & 750D & C934 & Ryan & 2.21 \\\\ ESO074-G015 & IC5052 & 24-25 FEB 2002 & ATCA & 1.5A & C894 & O'Brien & 10.23 \\\\ ESO074-G015 & IC5052 & 01 DEC 2002 & ATCA & 6.0A & C894 & O'Brien & 10.63 \\\\ ESO074-G015 & IC5052 & 17 OCT 1992 & ATCA & 6.0C & C212 & Carignan & 10.05 \\\\ \\hline ESO109-G021 & IC5249 & 20-22 MAR 2003 & ATCA & EW352 & CX043 & Dahlem & 8.54 \\\\ ESO109-G021 & IC5249 & 03-04 FEB 2003 & ATCA & 750D & C894 & O'Brien & 8.01 \\\\ ESO109-G021 & IC5249 & 28 NOV 2002 & ATCA & 6.0A & C894 & O'Brien & 13.43 \\\\ ESO109-G021 & IC5249 & 18 OCT 1992 & ATCA & 6.0C & C212 & Carignan & 11.11 \\\\ \\hline ESO115-G021 & & 09 FEB 2005 & ATCA & EW352 & C1341 & Koribalski & 10.09 \\\\ ESO115-G021 & & 06 JAN 2005 & ATCA & 750B & C1012 & Hoegaarden & 11.17 \\\\ ESO115-G021 & & 23 JUN 1995 & ATCA & 750B & C073 & Walsh & 7.57 \\\\ ESO115-G021 & & 08 SEP 2002 & ATCA & 6.0C & C894 & O'Brien & 6.51 \\\\ ESO115-G021 & & 03 DEC 2002 & ATCA & 6.0A & C894 & O'Brien & 10.05 \\\\ ESO115-G021 & & 13 DEC 2002 & ATCA & 6.0A & C894 & O'Brien & 5.72 \\\\ \\hline ESO138-G014 & & 08-09 NOV 2002 & ATCA & 1.5A & C894 & O'Brien & 10.43 \\\\ ESO138-G014 & & 29-30 NOV 2002 & ATCA & 6.0A & C894 & O'Brien & 9.79 \\\\ \\hline ESO146-G014 & & 17 JAN 2002 & ATCA & 750A & C894 & O'Brien & 1.37 \\\\ ESO146-G014 & & 27-29 DEC 2000 & ATCA & 750C & C894 & O'Brien & 10.70 \\\\ ESO146-G014 & & 11 JAN 2001 & ATCA & 750C & C894 & O'Brien & 1.67 \\\\ ESO146-G014 & & 31 JUL 2001 & ATCA & 1.5A & C894 & O'Brien & 7.78 \\\\ ESO146-G014 & & 14-15 APR 2002 & ATCA & 6.0A & C894 & O'Brien & 10.19 \\\\ ESO146-G014 & & 27-28 JAN 2002 & ATCA & 6.0B & C894 & O'Brien & 10.51 \\\\ \\hline ESO274-G001 & & 29 AUG 1993 & ATCA & 1.5B & C073 & Walsh & 6.07 \\\\ ESO274-G001 & & 07 OCT 1993 & ATCA & 1.5D & C073 & Walsh & 6.80 \\\\ ESO274-G001 & & 28-29 NOV 2002 & ATCA & 6.0A & C894 & O'Brien & 9.63 \\\\ \\hline ESO435-G025 & IC2531 & 12 JAN 2002 & ATCA & 750A & C894 & O'Brien & 9.63 \\\\ ESO435-G025 & IC2531 & 17 JAN 2002 & ATCA & 750A & C894 & O'Brien & 2.02 \\\\ ESO435-G025 & IC2531 & 07 MAR 1997 & ATCA & 1.5D & C529 & Bureau & 9.07 \\\\ ESO435-G025 & IC2531 & 06 APR 1996 & ATCA & 6.0A & C529 & Bureau & 9.94 \\\\ ESO435-G025 & IC2531 & 13-14 SEP 1996 & ATCA & 6.0B & C529 & Bureau & 9.96 \\\\ ESO435-G025 & IC2531 & 17-18 OCT 1992 & ATCA & 6.0C & C212 & Carignan & 7.45 \\\\ \\hline UGC07321 & & 26,30 MAY 2000 & VLA & C & AM649 & Matthews & 16.00 \\\\ UGC07321 & & 01 NOV 2000 & ATCA & 1.5D & C894 & O'Brien & 2.49 \\\\ \\hline \\end{tabular} \\end{small} \\end{table*} Modern dark matter research began in 1970 with several papers which found that galaxies contain more gravitating matter than can be accounted for by the stars. \\citet{freeman1970} noted that the atomic hydrogen (\\HI) rotation curves of the late-type disk galaxies NGC300 and M33 peaked at a larger radius than expected from the stellar light distribution. This implied that the dark matter was more extended than the stellar distribution and that its mass was at least as great as the luminous mass. However, these data were of poor angular resolution. At the same time \\citet{rubin1970} published a rotation curve of M31 based on optical data, which did not seem to decline in the outer parts. In the 1970s radio observations of increasingly better angular resolution and better sensitivity\\footnote{KCF and PCvdK recall influential colloquia and other presentations by M.S. Roberts in the early 1970s on a flat \\HI\\ rotation curve of the Andromeda galaxy which helped to steer the evolution of this subject.} showed that flat rotation curves are typical in spiral galaxies \\citep{shostak1971,rs1972,rr1973,bosma1978} and that the HI extent of a spiral galaxy can be far greater than the extent of the optical image. This, combined with surface photometry, leads to very high mass-to-light ratios in the outer parts, as shown clearly for several galaxies in \\citet{bosma1978} and \\citet{bk1979}. In 1974, dark matter halos were found to extend even further, when \\citet{oyp1974} and \\citet{eks1974} tabulated galaxy masses as a function of radius and found that galaxy mass increased linearly out to at least $100$ kpc and $10^{12}$ M$_{\\odot }$ for normal spirals and ellipticals. Despite this large dark matter fraction inferred in galaxies, it was still not sufficient to reach the critical $\\Omega_{tot}\\sim$1 value of a flat universe. See e.g. \\citet{sr01} and \\citet{rob2008} for a brief history of dark matter in galaxies. Around the same time, application of nuclear physics to Big Bang theory showed that big bang nucleosynthesis (BBNS) in the early universe produced specific abundances of the light elements, and predicted the total baryon density was $\\Omega_b \\approx 0.044$. In the last decade, the combined observations of high redshift SNIa, \\citep{reissetal1998,perlmutteretal1999}, the 2-Degree Field Galaxy Redshift Survey (2dFGRS) \\citep{percivaletal2001} and WMAP microwave background measurements \\citep{spergel_wmap2003} confirmed the baryonic mass density $\\Omega_b$ found by BBNS and established that dark energy comprises about $75\\%$ of the critical density. Consequently, the scale of the missing dark matter is now reduced. The mass density $\\Omega_m$ is now only $0.25$, but the problem remains: the measured baryonic mass density is still only $\\sim$18\\%\\ of the total mass density $\\Omega_m$. Thus, dark matter accounts for $82\\%$ of mass in the universe, and baryonic matter is only 4.5\\%\\ of the total content of the universe.\\footnote{For this illustration we used the parameters adopted in the Millenium Simulation \\citep{swjetal2005}, which are $\\Omega_m = \\Omega_{dm} + \\Omega_b = 0.25$, $\\Omega_b = 0.045$, $\\Omega_{\\Lambda} = 0.75$.} More than 1000 galactic rotation curves have now been measured, and very few display a Keplerian decline with radius. \\HI\\ and \\Ha\\ rotation curves of spiral galaxies show that the total-mass-to-light ratios are typically M/L = 10-20 M$_{\\odot}$/L$_{\\odot}$, and the luminous matter therefore accounts for only $5-10\\%$ of the total mass inferred from the rotation curves. For low surface brightness (LSB) disk galaxies \\citep{deblok1997} and dwarf irregular (dI) galaxies \\citep{swaters1999}, the M/L values increase to 10-100 M$_{\\odot}$/L$_{\\odot}$, with an extreme of $220$ for ESO215-G009, a gas-rich LSB galaxy with a very high gas mass to light ratio of $M_{\\rm HI}/L_B = 21$ and low recent star formation \\citep{wjk2004}. Dwarf spheroidal (dSph) galaxies, with typical total masses of only $\\sim$10$^{7}$ M$_{\\odot }$ within radii of a few hundred parsecs, are even more extreme: several have very large dark matter fractions with mass-to-light ratios in the range $200-1000$ M$_{\\odot}$/L$_{\\odot}$. In these faint, small galaxies, the dynamical mass is estimated from the line-of-sight velocities of individual stars. The Ursa Major dSph \\citep{willmanetal2005}, is one of the most dark matter dominated galaxies known to date with a central mass-to-light ratio $M/L\\sim$500 M$_{\\odot}$/L$_{\\odot}$, which is believed to increase further at larger radii \\citep{kweg2005}. These systems appear to have only very small baryonic mass fractions. The rotation curves of disk galaxies are important probes of the equatorial halo potential gradient. By decomposing the observed rotation curve into contributions from the visible mass components, the radial potential gradient of the halo can be measured, assuming the system is in centrifugal equilibrium. The dark halo mass component is typically fitted by a pseudo-isothermal halo model with density \\begin{equation} \\label{eq:pISO} \\rho(R) = \\rho_0 \\left[ 1 + \\left(\\frac{R}{R_c}\\right)^2 \\right]^{-1}, \\end{equation} where the halo is characterised by its central density $\\rho_0$ and core radius $R_c$. Pseudo-isothermal halos have an asymptotic density $\\rho \\propto R^{-2}$ at large radii which is consistent with commonly flat rotation curves. \\citet{kf2004} compiled the published dark halo density distributions for a large sample of Sc-Im and dSph galaxies, and found well-defined scaling relationships for the dark halo parameters for galaxies spanning a range of 6 decades of luminosity.\\footnote{E-Sbc galaxies were not included to avoid the larger uncertainties associated with stellar bulge-disk decomposition and the relatively larger contribution of the stellar mass of varying stellar ages.} They found that halos of less luminous and massive dwarf spheroidals have higher central densities, up to $\\sim$1 M$_{\\odot}$ pc$^{-3}$ and core radii of $\\sim$0.1 kpc, compared to $\\sim10^{-3}$ M$_{\\odot}$ pc$^{-3}$ and $\\sim$30-100 kpc respectively, for large bright Sc galaxies. The observed correlations suggest a continuous physical sequence of dark halo population in which the properties of the underlying dark halo scale with the baryon luminosity \\citep{kf2004}. We now consider the flattening of the dark halo density distribution, defined by $q=c/a$, where $c$ is the halo polar axis and $a$ is the major axis in the galactic plane. The vertical distribution of the halo is much more difficult to measure than the radial distribution in the equatorial plane, as most luminous tracers of the galaxy potential gradient lie in the plane of the galaxy and offer little indication of the vertical gradient of the potential. Past measurements that were obtained with a variety of different methods gave a large range of $q$ from $0.1$ to $1.37$, with no concentration on any particular value. For our Galaxy, the halo shape has been measured more than ten times by four different methods, that yielded $q-$values ranging from $0.45$ to $1.37$. In this series of papers, we will use the flaring of the HI gas layer to measure the vertical flattening of the dark halo, because we believe this method to be the most promising for late-type disk galaxies. Like all steady state mass components of a galaxy, the gaseous disk of an isolated disk galaxy can be assumed to be in hydrostatic equilibrium in the gravitational potential of the galaxy, unless there are visible signs from the gas distribution and kinematics that the HI layer is disturbed (e.g. by mergers or local starbursts). In the vertical direction, the gradient of gas pressure is balanced by the gravitation force (ignoring any contribution from a magnetic pressure gradient). From the observed distribution of the gas velocity dispersion and the gas density distribution, we can in principle measure the total vertical gravitational force. \\begin{table}[t] % \\centering \\caption[Spatial \\& spectral resolution of \\HI\\ observations]{Resolution of \\HI\\ observations} \\label{tab:resol-HI} \\begin{tabular}{lccc} \\hline & \\multicolumn{2}{c}{Spatial} & Spectral \\\\ Galaxy & \\multicolumn{2}{c}{resolution} & resolution \\\\ \\cline{2-3} & (arcsec) & (kpc) & (km s$^{-1}$)\\\\ \\hline\\hline ESO074-G015 & 9.0 $\\times$ 9.0 & 0.292 $\\times$ 0.292 & 3.298 \\\\ ESO109-G021 & 8.0 $\\times$ 8.0 & 1.179 $\\times$ 1.179 & 3.298 \\\\ ESO115-G021 & 8.9 $\\times$ 8.9 & 0.194 $\\times$ 0.194 & 3.298 \\\\ ESO138-G014 & 10.7 $\\times$ 10.7 & 0.965 $\\times$ 0.965 & 3.298 \\\\ ESO146-G014 & 7.6 $\\times$ 7.6 & 0.793 $\\times$ 0.793 & 3.298 \\\\ ESO274-G001 & 9.8 $\\times$ 9.8 & 0.162 $\\times$ 0.162 & 3.298 \\\\ ESO435-G025 & 9.0 $\\times$ 9.0 & 1.305 $\\times$ 1.305 & 3.298 \\\\ UGC07321 & 15.0 $\\times$ 15.0 & 0.727 $\\times$ 0.727 & 5.152 \\\\ \\hline \\end{tabular} \\end{table} Euler's equation for a steady-state fluid of density $\\rho$, velocity $\\mathbf{v}$ and pressure $p$ in a gravitational potential $\\Phi$ is $-\\mathbf{(v.\\nabla)v} = \\rho^{-1} \\mathbf{\\nabla}p + \\mathbf{\\nabla}\\Phi$. In the case of a vertically Gaussian gas density distribution with a vertically isothermal gas velocity dispersion, the gradient of the total vertical force $K_z$ in the $z$-direction can be calculated directly from the gas velocity dispersion $\\sigma_{v,g}(R)$ and the gas layer thickness FWHM$_{z,g}(R)$, each of which are functions of radius: \\begin{equation} \\label{eq:ch1-hydro} \\frac{\\partial K_z}{\\partial z} = - \\frac{\\sigma_{v,g}^2}{({\\rm FWHM}_{z,g}/2.35)^2}. \\end{equation} It is possible to measure the halo shape over the entire \\HI\\ extent of the luminous disk using the flaring of the \\HI\\ distribution, which typically extends in radius to $2-3 R_{25}$, and by measuring the density distribution of the gas and stellar distributions. For a given vertical gas velocity dispersion, a more flattened dark halo requires decreased flaring and increased gas surface density. The relatively high gas content of late-type galaxies allows measurement of both the halo vertical force field from the gas layer flaring and the halo radial force field from rotation curve decomposition. In this series of papers we will attempt to measure the halo flattening using the flaring of \\HI\\ disk in eight small late-type disk galaxies. \\begin{table}[t] % \\caption[Noise of \\HI\\ channel maps]{Noise of \\HI\\ channel maps} \\label{tab:noise-HI} \\begin{tabular}{lcccc} \\hline Galaxy & \\multicolumn{3}{c}{Noise} & Max Signal-\\\\ \\cline{2-4} & (Jy/beam) & (atoms cm$^{-2}$) & (K) & to-noise \\\\ \\hline\\hline ESO074-G015 & 0.00135 & 1.8348 $10^{19}$ & 10.0 & 14.5 \\\\ ESO109-G021 & 0.00125 & 2.1464 $10^{19}$ & 11.7 & 10.1 \\\\ ESO115-G021 & 0.00113 & 1.5774 $10^{19}$ & 8.6 & 18.6 \\\\ ESO138-G014 & 0.00197 & 1.9008 $10^{19}$ & 10.4 & 15.2 \\\\ ESO146-G014 & 0.00143 & 2.7190 $10^{19}$ & 14.9 & 11.3 \\\\ ESO274-G001 & 0.00151 & 1.7209 $10^{19}$ & 9.4 & 14.6 \\\\ ESO435-G025 & 0.00123 & 1.6789 $10^{19}$ & 9.2 & 10.9 \\\\ UGC07321 & 0.00038 & 1.8579 $10^{18}$ & 1.0 & 93.0 \\\\ \\hline \\end{tabular} \\end{table} The main advantage of this method is that it can be used for all gas-rich spiral galaxies inclined sufficiently to measure accurate kinematics, unlike some methods which are applicable only to specific kinds of galaxies like polar ring galaxies. This minimum inclination was determined to be $i \\gesim 60^{\\circ}$ by \\citet{olling1995}. This method was first tried by \\citet{crb1979} on the Galaxy, and early development was undertaken by \\citet{vdkruit1981}, who applied it to low resolution observations of NGC891 and concluded that the halo was not as flattened as the stellar disk. It was then used for several galaxies in the 1990's, most notably the careful study of the very nearby Sc galaxy NGC4244 which was found to have a highly flattened halo with $q=0.2^{+0.3}_{-0.1}$ out to radii of $\\approx 2 R_{25}$ \\citep{olling1996b}. All applications of the \\HI\\ flaring method to date have derived highly flattened halo distributions ($q \\leq 0.5$). With the exclusion of NGC4244, we suspect that the assumption of a radially-constant gas velocity dispersion led to errors in the measured total vertical force, and thereby to the derived flattening of the halo. The other difficulty for all these galaxies except NGC4244 is that they are large galaxies with peak rotation speeds $>170$ \\kms. Given that {\\it (i)} spiral galaxies typically have \\HI\\ velocity dispersions in the relatively small range $6-10$ \\kms \\citep{tamburro09} and {\\it (ii)} the halo shape $q$ is roughly proportional to the gradient of the vertical force, we see from Eqn.~\\ref{eq:ch1-hydro} that the \\HI\\ thickness is to first order inversely proportional to the peak rotation speed. Consequently, the \\HI\\ flaring method should work best for small disk galaxies with relatively low rotation speeds, as these galaxies should exhibit more \\HI\\ flaring. \\citet{bosma1994} already showed by calculations that flaring is relatively more important in galaxies with low circular velocity. The flaring method has also been applied to the Galaxy by \\citet{om2000}. However, uncertainty in the position and rotation speed of the Sun resulted in a large uncertainty of the measured halo shape $0.5 \\leq q \\leq 1.25$. \\begin{table*}[t] % \\centering \\caption[HI measurements]{HI measurements} \\label{tab:meas-HI} \\begin{tabular}{lcccrrr} \\hline Galaxy & RA & Dec & PA & Dist\\ \\ & $v_{sys}$\\ \\ \\ \\ & $v_{max}$\\ \\ \\ \\ \\\\ & & & ($^{\\circ}$) & (Mpc) & (km s$^{-1}$) & (km s$^{-1}$) \\\\ \\hline\\hline ESO074-G015 & 20 52 05.57 & --69 12 05.9 & 141.3\\ \\ & 6.70 & 590.7 & 93.4 \\\\ ESO109-G021 & 22 47 06.07 & --64 50 00.6 & 14.9 & 30.40 & 2360.1 & 112.4 \\\\ ESO115-G021 & 02 37 47.28 & --61 20 12.1 & 43.4 & 4.93 & 512.2 & 64.4 \\\\ ESO138-G014 & 17 06 59.22 & --62 04 58.3 & 134.4\\ \\ & 18.57 & 1508.4 & 120.4 \\\\ ESO146-G014 & 22 13 00.08 & --62 04 05.5 & 47.3 & 21.45 & 1691.1 & 70.2 \\\\ ESO274-G001 & 15 14 13.84 & --46 48 28.6 & 36.9 & 3.40 & 522.8 & 89.4 \\\\ ESO435-G025 & 09 59 55.77 & --29 37 01.1 & 74.5 & 29.89 & 2477.0 & 236.3 \\\\ UGC07321 & 12 17 34.56 & +22 32 26.4 & 81.8 & 10.00 & 401.9 & 112.1 \\\\ \\hline\\hline Galaxy & $W_{20}$ & $W_{50}$ & FI & $M_{\\rm HI}$\\ \\ & \\multicolumn{2}{c}{Diam.} \\\\ & (km s$^{-1}$) & (km s$^{-1}$) & (Jy km s$^{-1}$) & (M$_{\\odot}$)\\ \\ & (arcmin) & (kpc) \\\\ \\hline\\hline ESO074-G015 & 160.8 & 186.8 & 37.9 & 4.0 10$^{8}$ & 9.350 & 18.2 \\\\ ESO109-G021 & 211.4 & 224.8 & 10.9 & 2.4 10$^{9}$ & 5.020 & 44.4 \\\\ ESO115-G021 & 112.2 & 128.9 & 72.0 & 3.4 10$^{8}$ & 13.600 & 19.5 \\\\ ESO138-G014 & 226.3 & 240.8 & 51.0 & 4.2 10$^{9}$ & 9.210 & 49.8 \\\\ ESO146-G014 & 127.7 & 140.4 & \\ \\ 8.4 & 9.1 10$^{8}$ & 3.840 & 24.0 \\\\ ESO274-G001 & 167.7 & 178.8 & 152.6\\ \\ & 4.2 10$^{8}$ & 16.550 & 16.4 \\\\ ESO435-G025 & 460.9 & 472.7 & 13.1 & 2.8 10$^{9}$ & 7.300 & 63.5 \\\\ UGC07321 & 209.1 & 224.2 & 38.3 & 9.0 10$^{8}$ & 9.580 & 27.9 \\\\ \\hline \\end{tabular} \\end{table*} In this first paper we discuss the selection of our sample and present \\HI\\ observations. In paper II we will review methods to derive the information we need for our analysis from these data, namely the rotation curves and the \\HI\\ distributions, \\HI\\ velocity dispersion and the flaring of the \\HI\\ layer, all as a function of radius in the deprojected galaxy plane. Paper III will be devoted to applying this to our data and presenting the results for each individual galaxy. In paper IV will analyse the data of one galaxy in our sample, namely UCG7321 and we will set limits on the flattening of its dark halo. ", "conclusions": "\\label{sec:disc-HI} The resolution of our \\HI\\ data along the major axis is high, with the number of independent beamwidths on each side of the galaxy centre ranging from $15$ to over $50$ in our galaxies. In the vertical direction all galaxies are spanned by at least $5$ beamwidths. Unfortunately three of the galaxies in the sample suffer from incomplete imaging, resulting in missing information about the extended spatial structure due to the lack of observations along short baselines. For two of these galaxies ESO074-G015 and ESO274-G001, the images still contain substantial information due to the small linear size of the beam, $290$ pc and $160$ pc, respectively. For these galaxies the iterative \\HI\\ modelling methods used to measure the deprojected \\HI\\ density distribution and kinematics (Paper II) should provide good measurements of the flaring. But the \\HI\\ images of ESO138-G014 indicate missing extended \\HI\\ emission and low spatial resolution which will probably prevent reliable measurement of the flaring in this galaxy. All the other galaxies in our sample are promising candidates for accurate measurement of the \\HI\\ kinematics and vertical flaring. The \\HI\\ column density distribution of the galaxies in our sample varies quite substantially. The four galaxies with maximum rotation speeds $\\gesim 100$ \\kms\\ all have \\HI\\ disks that extend to greater than $5$ kpc away from the plane. Three of them (ESO435-G25, UGC7321 and ESO138-G14) also have the ``figure-8'' signature in the XV diagram suggestive of an edge-on bar in the gas distribution. ESO435-G25 \\citep{bf1999} and UGC7321 \\citep{pbld2003} both have boxy-peanut shaped stellar bulges consistent with a bar seen edge-on. Combined with the figure-8 signature in the \\HI\\ gas structure this is strong evidence for a bar \\citep{km1995,bf1999,ab1999,athanassoula2000}. E138-G014, IC5249 and IC2531 all have large \\HI\\ disks with radii larger than $20$ kpc, while UGC7321 is much smaller with a radius of only $14.0$ kpc. The other four galaxies are not much smaller than UGC7321 in radial extent, but they are considerably smaller in total mass. They also lack the vertical extensions in the central disk that suggest additional sources of heat in the disk. This suggests that one of the causes of extended high latitude \\HI\\ filaments is heating related to star formation and other processes associated with the bar. The peak brightness profiles of each galaxy show that the inferred \\HI\\ opacity (assuming a constant \\HI\\ spin temperature) varies substantially across the major axis of each galaxy. The high spatial resolution of our images has made it possible to measure high \\HI\\ brightness temperatures. The maximum brightness temperature in our galaxies ranges from $94.4$ K for UGC7321 to $168.6$ K for ESO146-G14. In addition to ESO146-G14, two other galaxies have high \\HI\\ brightness temperatures $> 150$ K. Assuming a mean \\HI\\ spin temperature of $300$ K, the maximum inferred opacity for these \\HI\\ bright galaxies is $\\sim$0.7. This is comparable to the maximum \\HI\\ opacity of $0.85$ measured in NGC891 \\citep{kvdkdb2004}. Inspection of the peak brightness as a function of major axis position for these three galaxies shows that, in each case, the region of increased opacity is localised spanning a projected radius of $\\sim$1-2 kpc. In ESO115-G21, this increased opacity region occurs at the galactic centre. However for ESO138-G14 and ESO146-G14, the regions of increased opacity occur in the outer disk. Three of our eight galaxies appear to be warped, and one is lopsided. This is consistent with results from larger samples, as e.g. mentioned in the review by \\citet{sancisi2008}." }, "1003/1003.1115_arXiv.txt": { "abstract": "This paper performs a semi-analytic study of relativistic blast waves in the context of gamma-ray bursts (GRBs). Although commonly used in a wide range of analytical and numerical studies, the equation of state (EOS) with a constant adiabatic index is a poor approximation for relativistic hydrodynamics. Adopting a more realistic EOS with a variable adiabatic index, we present a simple form of jump conditions for relativistic hydrodynamical shocks. Then we describe in detail our technique of modeling a very general class of GRB blast waves with a long-lived reverse shock. Our technique admits an arbitrary radial stratification of the ejecta and ambient medium. We use two different methods to find dynamics of the blast wave: (1) customary pressure balance across the blast wave and (2) the ``mechanical model''. Using a simple example model, we demonstrate that the two methods yield significantly different dynamical evolutions of the blast wave. We show that the pressure balance does not satisfy the energy conservation for an adiabatic blast wave while the mechanical model does. We also compare two sets of afterglow light curves obtained with the two different methods. ", "introduction": "\\label{section:introduction} The afterglow emission of a gamma-ray burst (GRB) is believed to be produced by a relativistic blast wave (M\\'esz\\'aros \\& Rees 1997). The relativistic blast wave is driven by an ``ejecta'', which is ejected by the central engine of the GRB explosion. As the ejecta interacts with a surrounding ambient medium, two (forward and reverse) shock waves develop (e.g., Piran 2004). The forward shock (FS) wave sweeps up the ambient medium, and the reverse shock (RS) wave propagates through the ejecta. As the blast wave has high Lorentz factors $10^2 - 10^3$ (e.g., M\\'esz\\'aros 2006), the FS wave is highly relativistic and an equation of state (EOS) with a constant adiabatic index 4/3 may well describe the gas in the FS-shocked region. However, the strength of the RS wave varies as the blast wave propagates. In the case of a constant-density ambient medium, the RS wave is initially non-relativistic and then transitions to a mildly relativistic or relativistic regime (Kobayashi 2000; Sari \\& Piran 1995). Thus, an EOS with a constant adiabatic index is not adequate for the gas in the RS-shocked region; a variable adiabatic index needs to be considered to account for change in the gas temperature. Although the EOS with a constant adiabatic index has been widely used in analytical and numerical studies of relativistic hydrodynamics, it is valid only for the gas of either non-relativistic (with the index 5/3) or ultra-relativistic temperature (with the index 4/3). The correct EOS for a relativistic ideal gas is formulated in terms of modified Bessel functions (e.g., Synge 1957), and its equivalent adiabatic index varies from 5/3 to 4/3 as the temperature increases. As it is not convenient to deal with modified Bessel functions, there has been effort to find simpler EOSs that closely reproduce the correct EOS of a relativistic ideal gas. Taub (1948) showed that the choice of EOS is not arbitrary and must satisfy a certain inequality (Taub's inequality). By taking the equal sign in Taub's inequality, Mignone et al. (2005) derived a simple form of EOS that has correct limiting values 5/3 and 4/3. The same EOS as in Mignone et al. (2005) was previously introduced by Mathews (1971), considering a relativistic ``monoenergetic'' gas where all particles have the same energy. The validity of this EOS was addressed by Blumenthal \\& Mathews (1976) for the cases of both infinite mean free collision times and very short mean free collision times. This EOS was also adopted by Meliani et al. (2004) and Mignone \\& McKinney (2007). In particular, Mignone \\& McKinney (2007) demonstrated in their relativistic numerical simulations that use of an EOS with a constant adiabatic index can significantly endanger the solution when transitions from cold to hot gas (or vice versa) are present. We use the same above EOS in this paper. Following Mathews (1971), we consider a relativistic monoenergetic gas and show that it closely reproduces the correct EOS of a relativistic ideal gas. Then we use this EOS to find a simple form of jump conditions for relativistic hydrodynamical shocks. This simple set of jump conditions applies to shocks of arbitrary strength. A short-lived RS was proposed to explain a brief optical flash (M\\'esz\\'aros \\& Rees 1999; Sari \\& Piran 1999a, 1999b). A dynamical evolution of such a short-lived RS was studied analytically by assuming an equality of pressure across the blast wave (Kobayashi 2000; Sari \\& Piran 1995). The RS wave here is short-lived since the ejecta is assumed to have a constant Lorentz factor. However, in general, the ejecta is expected to emerge with a range of the Lorentz factors. The shells with lower Lorentz factors will gradually ``catch up'' with the blast wave as it decelerates. Thus, the RS wave is long-lived. Such a long-lived RS was studied for a power-law ejecta by assuming a constant ratio of the two pressures at the FS and RS (Rees \\& M\\'esz\\'aros 1998). In this paper, we present a detailed description of our blast-wave modeling technique for even more general class of explosions where the ejecta and the ambient medium have an arbitrary radial structure or stratification. More specifically, we study analytically the spherical expansion of such a stratified ejecta and find the trajectory of the RS wave through the ejecta self-consistently. In order to find a dynamical evolution of the blast wave, we use two different methods: (1) customary pressure balance and (2) the ``mechanical model'' (Beloborodov \\& Uhm 2006). Using a simple example model, we demonstrate that, although the customary assumption of pressure balance for the blast wave yields an estimated evolution, it is not rigorously accurate. In particular, the energy conservation is not satisfied for an adiabatic blast wave; the total energy is decreased by a factor of 5 in the case of the example model. The mechanical model was developed for relativistic blast waves, by relaxing the pressure balance (or proportionality) and applying the conservation laws of energy-momentum tensor and mass flux on the blast between the FS and RS. Using the same example model, we show that the energy conservation is satisfied for the mechanical model. We also show that dynamical evolutions found by the two methods differ significantly. Finally, we present the afterglow light curves in X-ray and optical bands. We compare two sets of light curves corresponding to the two different dynamical evolutions mentioned above. In Section~\\ref{section:shock}, we derive a simple set of jump conditions for relativistic hydrodynamical shocks. In Section~\\ref{section:blast_waves}, we describe in detail our blast-wave modeling technique. We also provide a simple method of evaluating the blast energy, employing a Lagrangian description for the blast wave. In Section~\\ref{section:mechanical_model}, we review the mechanical model including more detailed equations. ", "conclusions": "As the blast wave propagates, the strength of the RS wave exhibits a transition from non-relativistic to mildly relativistic or relativistic regime (or vice versa). Thus, an EOS with a constant adiabatic index is not adequate for the RS-shocked region. We address that a more realistic EOS with a variable adiabatic index needs to be used for the gas in the RS-shocked region. Following Mathews (1971), we consider a relativistic monoenergetic gas and find its EOS. We show that there is only 4.8 \\% of maximal difference in the quantity $\\kappa$ (pressure divided by internal energy density) when compared to a relativistic ideal gas. Then we show that jump conditions of relativistic hydrodynamical shocks simplify significantly for the monoenergetic gas (see Section~\\ref{section:shock}). The simple form of jump conditions presented here is exact for a monoenergetic gas and applies to shocks of arbitrary strength (relativistic, mildly relativistic, or non-relativistic). We emphasize that its usage is not to be restricted to GRB blast waves; it can be applied to other areas of relativistic hydrodynamical shocks. Then we present a semi-analytic formulation for relativistic blast waves with a long-lived RS. We describe in detail a complete set of tools for finding a dynamical evolution of the blast wave for a very general class of explosions. The ambient medium can have an arbitrary radial profile, and the explosion ejecta can also be arbitrary as long as $\\Gej^\\prime(\\tau) \\leq 0$. We provide two different methods of finding dynamics of the blast wave: (1) customary pressure balance and (2) the mechanical model (Beloborodov \\& Uhm 2006). Using a simple example model, we show that the pressure balance across the blast wave does not satisfy the energy conservation for an adiabatic blast wave; the total energy is decreased by a factor of 5 in the case of the example model. The mechanical model does not assume a pressure balance or proportionality across the blast wave (neither $\\pf=\\pr$ nor $\\pf/\\pr={\\rm const.}$ is assumed). Instead, it finds the dynamics of the blast wave from a set of coupled differential equations that express the conservations of energy-momentum tensor and mass flux applied on the blast between the FS and RS. Using the same example model, we show that the energy conservation is satisfied for the mechanical model as expected. We also show that the two methods yield very different dynamical evolutions of the blast wave and, as a result, very different afterglow light curves. We conclude that the customary prescription of pressure balance poorly describes the dynamics of the blast wave with a long-lived RS and resulting afterglow light curves are inaccurate in a significant manner." }, "1003/1003.6029_arXiv.txt": { "abstract": "{HD 49798 is a hydrogen depleted subdwarf O6 star and has an X-ray pulsating companion (RX J0648.0$-$4418). The X-ray pulsating companion is a massive white dwarf. Employing Eggleton's stellar evolution code with the optically thick wind assumption, we find that the hot subdwarf HD 49798 and its X-ray pulsating companion could produce a type Ia supernova (SN Ia) in future evolution. This implies that the binary system is a likely candidate of SN Ia progenitors. We also discussed the possibilities of some other WD + He star systems (e.g. V445 Pup and KPD 1930+2752) for producing SNe Ia. ", "introduction": "% \\label{sect:intro} Type Ia supernova (SN Ia) explosions are among the most energetic events observed in the Universe. They appear to be good cosmological distance indicators owing to their high luminosities and remarkable uniformity, and have been applied successfully in determining cosmological parameters (e.g. \\textbf{$\\Omega_{M}$} and \\textbf{$\\Omega_{\\Lambda}$}; Riess et al. 1998; Perlmutter et al. 1999). However, the exact explosion mechanism and the nature of progenitors are still poorly understood (Hillebrandt \\& Niemeyer 2000; Podsiadlowski 2010; Wang et al. 2008a, 2010), and no SN Ia progenitor system before the explosion has been conclusively identified (Wang \\& Han 2009, 2010a; Meng \\& Yang 2010a). It is widely accepted that SNe Ia arise from thermonuclear explosions of carbon--oxygen white dwarfs (CO WDs) in binaries (Nomoto et al. 1997; Livio 2000). Over the past few decades, two groups of SN Ia progenitor models have been proposed, i.e. the double-degenerate (DD) and single-degenerate (SD) models. The DD model involves the merger of two CO WDs (Tutukov \\& Yungelson 1981; Iben \\& Tutukov 1984; Webbink 1984; Han 1998). Although the DD model might be able to account for the explosion of a few overluminous SNe Ia (Howell et al. 2006; Gilfanov \\& Bogd$\\acute{\\rm a}$n 2010), it is still suffering from the theoretical difficulty that the mergers of two WDs may lead to an accretion-induced collapse rather than thermonuclear explosion (Nomoto \\& Iben 1985; Saio \\& Nomoto 1985; Timmes et al. 1994). For the SD model, the companion could be a main-sequence (MS) star or a slightly evolved star (WD + MS channel), or a red-giant star (WD + RG channel) (Hachisu et al. 1996; Li $\\&$ van den Heuvel 1997; Langer et al. 2000; Fedorova et al. 2004; Han $\\&$ Podsiadlowski 2004, 2006; Chen $\\&$ Li 2007; Ruiter et al. 2009; L\\\"{u} et al. 2009; Meng \\& Yang 2010b; Wang, Li \\& Han 2010; Wang \\& Han 2010b). Observationally, there is increasing evidence indicating that at least some SNe Ia may come from the SD model (Hansen 2003; Ruiz-Lapuente et al. 2004; Voss \\& Nelemans 2008; Wang et al. 2008b; Justham et al. 2009). Moreover, the detections of variable Na I D lines (Patat et al. 2007; Blondin et al. 2009; Simon et al. 2009) and derivation of smaller absorption ratio $R_{\\rm V}$ that is characteristic of circumstellar material (CSM) dust (Wang et al. 2009c), perhaps suggests the presence of CSM around a subclass of SNe Ia. Recently, Wang et al. (2009a) studied the WD + He star channel of the SD model to produce SNe Ia, in which a CO WD accretes material from an He MS star or a slightly evolved He star to increase its mass to the Chandrasekhar (Ch) mass. The study derived the parameter spaces for the progenitors of SNe Ia. By using a detailed binary population synthesis approach, Wang et al. (2009b) found that the Galactic SN Ia birthrate from this channel is $\\sim$$0.3\\times 10^{-3}\\ {\\rm yr}^{-1}$, and that this channel may account for SNe Ia with short delay times ($\\sim$45$-$140\\,Myr) from the star formation to SN explosion (see also Wang \\& Han 2010c). Hot subdwarf stars are near the blue end of the horizontal branch in the Hertzsprung-Russell diagram. A particularly interesting member of this class is a subdwarf O star, HD 49798, which is one of the brightest subdwarfs and also a single-lined spectroscopic binary with an orbital period of 1.548\\,d (Thackeray 1970; Stickland \\& Lloyd 1994). This hydrogen-deficient subdwarf of spectral type O6 has been studied extensively (Kudritzki \\& Simon 1978; Hamann et al. 1981). Bisscheroux et al. (1997) showed that the hot subdwarf star HD 49798 must have a degenerate C-O core, and is in the He-shell burning phase, which can explain its high luminosity. Israel et al. (1997) reported the detection of a pulsating soft X-ray source (RX J0648.0$-$4418) with a pulsation period of $\\sim$13\\,s from this binary. The X-ray spectrum of the source is very soft, but has a high-energy excess. The discovery of the regular X-ray pulsations must arise from a compact companion of either a neutron star or a WD (e.g. Israel et al. 1997; Bisscheroux et al. 1997). Bisscheroux et al. (1997) excluded a neutron star as the companion of HD 49798, and showed that all observations are consistent with a weakly magnetized massive WD, which is accreting material from the wind of its subdwarf companion. Kudritzki \\& Simon (1978) estimated a mass of 0.7$-$2.7$\\,M_{\\odot}$ for the hot subdwarf HD 49798. This system is consistent with a double spectroscopic binary, favored by the detection of a fast X-ray pulsar source, for which all the orbital parameters (including the masses of the two components) may be derived. With this purpose, Mereghetti et al. (2009) recently observed HD 49798/RX J0648.0$-$4418 in may 2008 with XMM-Newton satellite. They confirmed that the 13\\,s pulsation in the X-ray binary HD49798/RXJ0648.0$-$4418 is due to a rapidly rotating WD. From the pulse time delays and the system's inclination, constrained by the duration of the X-ray eclipse discovered in this observation, they derived the masses of the two components. The corresponding masses are 1.50$\\pm$0.05$\\,M_{\\odot}$ for HD 49798 and 1.28$\\pm$0.05$\\,M_{\\odot}$ for the WD. The existence of WD + He star systems is supported by some observations (Wang et al. 2009a). The hot subdwarf HD 49798 with its WD companion is such a system, and may be a candidate of SN Ia progenitors. The goal of this paper is to investigate the evolution and final fate of the hot subdwarf HD 49798 and its WD companion, and to explore whether this binary system could produce an SN Ia. In Section 2, we describe the numerical code of the binary evolution and the input physics. In Section 3, we give the binary evolutionary results. Finally, discussion and conclusion are given in Section 4. ", "conclusions": "\\label{5:DISCUSSION AND CONCLUSIONS} The observations indicate that the X-ray source (RX J0648.0-4418) is a weekly magnetized massive WD which is accreting matter from the wind of its subdwarf companion. The mass loss of HD 49798 from the wind in our calculations is about $3\\times10^{-9}\\,M_{\\odot}\\,\\rm yr^{-1}$. By using the Bondi-Hoyle formalism as described by Davidson \\& Ostriker (1973), we can make an estimate of mass $\\dot M_{\\rm acc}$ captured from the wind by the gravitational field of the WD, and the luminosity $L_{\\rm acc}$ converted by the potential energy in the process of accretion. We find that a wind velocity between 800$-$1350\\,km\\,s$^{-1}$ with $\\dot M_{\\rm wind}=3\\times10^{-9}\\,M_{\\odot}\\,\\rm yr^{-1}$ will result in an accretion luminosity between $10^{30}-10^{31}$\\,erg\\,s$^{-1}$, consistent with that of the observed X-ray luminosity $\\sim$$10^{31}$\\,erg\\,s$^{-1}$ in the 0.2$-$10\\,keV energy range (e.g. Mereghetti et al. 2009). Mereghetti et al. (2009) confirmed that RX J0648.0$-$4418 is a rapidly rotating WD. The maximum stable mass of a rotating WD may be above the standard Chandrasekhar (Ch) mass (e.g. Uenishi et al. 2003; Yoon \\& Langer 2005; Chen \\& Li 2009), and the maximum possible mass a CO WD can reach by mass accretion is about 2.0$\\,M_{\\odot}$ (see Yoon \\& Langer 2005). According to our calculations, the maximum mass that the WD RX J0648.0$-$4418 can reach is about $1.62\\,M_\\odot$ (see Fig. 2), which is larger than the standard Ch mass ($1.4\\,M_\\odot$) we set in this paper. Thus, we point out that the WD companion of HD 49798 might evolve towards a thermonuclear explosion of the super-Ch mass WD, producing an overluminous SN Ia (e.g. Howell et al. 2006). We also note that, if rotation is taken into account, He burning is much less violent than that without rotating (see Yoon et al. 2004). This may significantly increase the He-accretion efficiency (i.e. $\\eta _{\\rm He}$ in our parametrization). Therefore, more He-rich matter can be converted into C and O, increasing the chance for a WD to survive above the Ch mass limit. The process that leads to the formation of this peculiar system is still poorly understood. It is suggested that HD 49798/RX J0648.0-4418 corresponds to a previously unobserved evolutionary stage of a massive binary system, after the common-envelope phase and spiral-in (e.g. Israel et al. 1997; Bisscheroux et al. 1997). A primordial binary system with a primary mass $M_{\\rm 1,i}\\sim5.0-8.0\\,M_\\odot$ and a secondary mass $M_{\\rm 2,i}\\sim2.0-6.5\\,M_\\odot$ may produce a system like HD 49798 and its WD companion (Wang et al. 2009b). We note that the primordial binary system has a short delay time ($\\sim$100\\,Myr) from the star formation to SN explosion. Thus, HD 49798 and its WD companion can contribute to the young population of SNe Ia revealed by recent observations (Mannucci et al. 2006; Aubourg et al. 2008). The young population of SNe Ia may have an effect on models of galactic chemical evolution, since they would return large amounts of iron to the interstellar medium much earlier than previously thought. The WDs usually have masses in a narrow range centered at about 0.6$\\,M_\\odot$ (Kepler et al. 2007). However, a few examples of WDs with very high mass ($>$1.2$\\,M_\\odot$) have recently been reported (Dahn et al. 2004; Vennes \\& Kawka 2008). And the X-ray source RX J0648.9-4418 may be such a massive WD. These massive WDs in binary systems are good candidates for the formation of SNe Ia if the mass transfer can occur, since a small amount of accreted mass could drive them above the Ch mass limit. Besides HD 49798 and its WD companion, there are also some other WD + He star systems, e.g. V445 Pup and KPD 1930+2752, which are also good candidates of SN Ia progenitors. V445 Pup is an He nova (Ashok \\& Banerjee 2003; Kato \\& Hachisu 2003). Kato et al. (2008) recently presented a free-free emission dominated light curve model of V445 Pup, based on the optically thick wind theory (Hachisu et al. 1996). The light curve fitting showed that the mass of the WD is above $1.35\\,M_{\\odot}$, and half of the accreted matter remains on the WD, resulting in the mass increase of the WD. Thus, V445 Pup is suggested to be one of the best candidate of SN Ia progenitors (Kato et al. 2008; see also Woudt et al. 2009). However, the orbital period of the binary system and the mass of the He star are still uncertain so far. To clarify the above parameters, further observations of V445 Pup are needed when the dense dust shell disappears. KPD 1930+2752 is regarded as another candidate of SN Ia progenitors, giving rise to SN Ia explosion in the form of merging WDs (Maxted et al. 2000; Geier et al. 2007). Note that, the DD model is not supported theoretically, as it may lead to an accretion-induced collapse rather than to an SN Ia (Nomoto \\& Iben 1985; Saio \\& Nomoto 1985; Timmes et al. 1994). On the other hand, KPD 1930+2752 may also produce an SN Ia through the SD model. However, the mass of the He donor star in KPD 1930+2752 is limited to the range between 0.45$\\,M_{\\odot}$ and 0.52$\\,M_{\\odot}$ (Geier et al. 2007), which is below the minimum mass (0.95$\\,M_{\\odot}$) for producing SNe Ia (Wang et al. 2009a). Thus, KPD 1930+2752 may not be a good candidate to produce an SN Ia via the SD model. In this paper, by using the optically thick wind model (Hachisu et al. 1996) and adopting the prescription of KH04 for the mass accumulation efficiency of the He-shell flashes onto the WD, we showed that the hot subdwarf HD 49798 and its X-ray pulsating companion could produce an SN Ia in future evolution. We also discussed the possibilities of some other WD + He star systems for producing SNe Ia. To further study the WD + He star channel of SNe Ia, large samples of WD + He star systems are expected in future observations. \\normalem" }, "1003/1003.4730_arXiv.txt": { "abstract": "X-ray images of galaxy clusters often display underdense bubbles which are apparently inflated by AGN outflow. I consider the evolution of the magnetic field inside such a bubble, using a mixture of analytic and numerical methods. It is found that the field relaxes into an equilibrium filling the entire volume of the bubble. The timescale on which this happens depends critically on the magnetisation and helicity of the outflow as well as on properties of the surrounding ICM. If the outflow is strongly magnetised, the magnetic field undergoes reconnection on a short timescale, magnetic energy being converted into heat whilst the characteristic length scale of the field rises; this process stops when a global equilibrium is reached. The strength of the equilibrium field is determined by the magnetic helicity injected into the bubble by the AGN: if the outflow has a consistent net flux and consequently a large helicity then a global equilibrium will be reached on a short timescale, whereas a low-helicity outflow results in no global equilibrium being reached and at the time of observation reconnection will be ongoing. However, localised flux-tube equilibria will form. If, on the other hand, the outflow is very weakly magnetised, no reconnection occurs and the magnetic field inside the bubble remains small-scale and passive. These results have implications for the internal composition of the bubbles, their interaction with ICM -- in particular to explain how bubbles could move a large distance through the ICM without breaking up -- as well as for the cooling flow problem in general. In addition, reconnection sites in a bubble could be a convenient source of energetic particles, circumventing the problem of synchrotron emitters having a shorter lifetime than the age of the bubble they inhabit. ", "introduction": "\\label{sec:intro} The gravitational potential wells of galaxy clusters are filled with hot ($10^{7-8}$ K), hydrostatically-settled gas which emits X-rays via thermal bremsstrahlung (e.g. \\citealt{Molendi:2004}). Many galaxy clusters, viewed in X rays, display dark cavities of size $\\sim10$ kpc at various distances from the cluster centre {\\mk \\citep{Boehringer_etal:1993,Carilli_etal:1994,Dunn_Fabian:2004,McNamara_Nulsen:2007,Birzan_etal:2008}.} They are dark because they have a lower density than the surrounding intra-cluster medium (ICM), but precisely how much less dense is uncertain, except that they are at least a factor of three or so less dense. {\\mk Observationally constraining the density is difficult because the line of sight contains also surrounding material; this problem becomes more severe where the bubble is at a larger distance from the cluster centre (see \\citealt{Ensslin_Heinz:2002} for details).} The bubbles are apparently inflated by an Active Galactic Nucleus (AGN) at the cluster centre and then rise buoyantly through the ICM. In addition, we infer the presence of an internal magnetic field and cosmic rays from observed radio synchrotron emission. There is a growing consensus that negative feedback from AGN could solve the cooling flow problem: the accretion of gas onto a supermassive black hole in the central galaxy releases energy to heat the ICM, preventing it from cooling and collapsing towards the centre of the cluster. Observationally, there is a strong correlation between those clusters which require heating (i.e. have a short cooling timescale) and the presence of optical-line emission and radio emission from AGN as well as star formation \\citep{Burns:1990,Rafferty_etal:2008,Cavagnolo_etal:2008}. The means by which this energy might be transferred to the ICM is not yet understood, but an interaction between the AGN outflow and the surroundings does seem very likely {\\mk (e.g. \\citealt{Brueggen_Kaiser:2002,Reynolds_etal:2002,Churazov_etal:2005,Brueggen_etal:2005,Brighenti_Mathews:2006}).} For this reason, it is important to gain some understanding of how AGN-inflated bubbles interact with their surroundings. Rising bubbles in a fluid tend to lose their spherical shape after rising a distance comparable to their radius. First, bubbles tend to flatten while the surrounding medium is flowing around them because the material flowing past their sides is moving with greater velocity than either in front of or behind them; we know from Bernoulli's principle that the pressure at the sides must therefore be lower and so the bubble expands laterally. {\\mk Then, the bubble is shredded into many smaller bubbles and eventually becomes completely mixed into the surrounding medium. In general there is more than one instability responsible for this shredding: the Rayleigh-Taylor (R-T) instability appears at the leading edge of the bubble where a dense fluid (the ICM) lies above a less dense fluid (the bubble) and the Kelvin-Helmholtz (K-H) instability appears at the sides of the bubble where there is a discontinuity in velocity and density. In the absence of magnetic fields, the growth time of the longest wavelength mode (i.e. the bubble radius) of the R-T instability is comparable to the time the bubble takes to rise a distance equal to its own size. The growth time of the K-H instability is likely to be somewhat longer if there is a large density contrast between the bubble and its surroundings. However, in many clusters we see large bubbles which have risen distances many times greater than their own size -- some mechanism must be inhibiting the instabilities \\citep{JonandDeY:2005,Ruszkowski_etal:2007}.} An obvious candidate is a magnetic field, coherent on the length scale of the bubble, either in the ambient medium (`magnetic draping', see \\citealt{Lyutikov:2006,DurandPfr:2008}) or inside the bubble, or both. Alternatively, by analogy with smoke rings it seems plausible that there is some purely hydrodynamical process responsible -- for instance \\citet{ScaandBru:2009} and \\citet{Brueggen_etal:2009} find that adding a subgrid-turbulence model to hydro simulations could encourage the bubble to stay in one piece. Here, I concentrate on the magnetic field inside the bubble, and show how an arbitrary `turbulent' magnetic field in a new-born bubble could reconnect into a large-scale equilibrium and thus provide the necessary rigidity. This process is similar to that taking place in stars which make a transition from convective to non-convective, for instance in proto-neutron stars \\citep{BraandSpr:2004,Braithwaite:2008}. In section \\ref{sec:analytic} I look at the process of relaxation to equilibrium, finding a relation between the initial magnetic helicity and the equilibrium field strength as well as comparing the relevant timescales. In section \\ref{sec:sims} I present numerical simulations of the reconnection process, before looking in some detail in section \\ref{sec:structure} at the structure of the equilibria found. In sections \\ref{sec:disc} and \\ref{sec:conc} I discuss the results and then summarise and conclude. ", "conclusions": "\\label{sec:conc} I have considered the evolution of the magnetic field inside AGN-inflated bubbles which are observed as dark cavities in X-ray images of galaxy clusters. It is found that the magnetic field undergoes relaxation to a global-scale equilibrium filling the entire bubble, consisting of twisted flux tube(s) arranged in some pattern. The relaxation process inevitably involves magnetic reconnection -- the reconnection regions could provide energetic synchrotron-emitting particles via X-point and Fermi acceleration {\\mk(see e.g. \\citealt{Parker:1957,Miller_etal:1997}).} The timescale on which this relaxation takes place, or in other words the stage during this relaxation we are likely to observe, depends crucially on various parameters: the magnetic field strength, mass density, Lorentz factor and size of the outflow as well as the properties of the ambient intra-cluster medium into which the bubble expands. Given the uncertainly in these parameters, it is impossible at this stage to distinguish between the following eventualities (see section \\ref{sec:timescales}). In the following, the radius of the bubble and the dominant length scale of its magnetic field structure are $r$ and $l$ respectively; the Alfv\\'en speed is $v_{\\rm A}=B/\\sqrt{4\\pi\\rho}$ and there is a reconnection timescale $\\tau_{\\rm rec}=l/(\\alpha v_{\\rm A})$ where $\\alpha\\approx0.1$ is the reconnection speed parameter. \\begin{enumerate} \\item The AGN outflow is weakly magnetised and little reconnection occurs; the observed field is small-scale and evolves passively in response to the bubble's interaction with the ICM. Measurement of the relevant parameters would show that $l\\ll r$ and $\\tau_{\\rm rec}>\\tau_{\\rm age}$ where $\\tau_{\\rm age}$ is the age of the bubble. \\item The AGN outflow is strongly magnetised and the magnetic field relaxes towards a global equilibrium. However, because the helicity of the field is low, the bulk of the magnetic energy is dissipated and no global equilibrium is reached. At the time of observation the reconnection is still ongoing and $\\tau_{\\rm rec} \\approx \\tau_{\\rm age}$. The field may consist of local-equilibrium flux tubes of size $l \\tau_{\\rm age}$. The field may consist of large-scale twisted flux tube(s) arranged in figure-of-eight patterns or as a single torus configuration, similar to the spheromak shape found in laboratory experiments. In this case, the magnetic field will give the bubble some rigidity, helping keep it intact as it moves through the ICM. \\end{enumerate} To illustrate this with plausible parameters, if we measure a density $10^{-5}m_p$g cm$^{-3}$ and field strength $20\\mu$G in a bubble of radius $10$kpc then $\\tau_{\\rm rec}\\approx (l/r)\\,7$Myr; if the bubble is older than $7$Myr then we have the global-equilibrium case (iii). To reach this situation the AGN outflow must have had high helicity; this is likely if the accretion disc is fed material with a consistent net flux. Fluctuating or vanishing net flux through the accretion disc will result in case (ii) even if the outflow is strongly magnetised. During reconnection to equilibrium, the shape of the bubble may change in response to plasma flow inside the bubble on the order of the Alfv\\'en speed. However, if the density of the bubble is much less than the density of the surrounding ICM, the effect on the shape of the bubble will be rather modest. Finally, it is shown that the difference in gas pressure between a bubble and its surroundings is equal to one third of the magnetic energy density, i.e. the magnetic field produces an `isotropic magnetic pressure' $P_{\\rm mag}=B^2/24\\pi=P_{\\rm o}-P_{\\rm i}$ where the subscripts o and i denote pressure outside and inside the magnetised volume. In this and other contexts this is more useful than the $B^2/8\\pi$ which is more common in the literature; this is a general feature of three-dimensional problems. {\\it Acknowledgements.} The author would like to thank Marcus Br\\\"uggen, Eugene Churazov, Peter Goldreich, \\AA ke Nordlund, Christoph Pfrommer and Henk Spruit for assistance and useful discussions." }, "1003/1003.5927_arXiv.txt": { "abstract": "An overview of some recent progress on magnetohydrodynamic stability and current sheet formation in a line-tied system is given. Key results on the linear stability of the ideal internal kink mode and resistive tearing mode are summarized. For nonlinear problems, a counterexample to the recent demonstration of current sheet formation by Low \\emph{et al}. {[}B. C. Low and \\AA. M. Janse, Astrophys. J. \\textbf{696}, 821 (2009){]} is presented, and the governing equations for quasi-static evolution of a boundary driven, line-tied magnetic field are derived. Some open questions and possible strategies to resolve them are discussed. ", "introduction": "In plasma physics, line-tying refers to the presence of conducting walls enclosing the magnetized plasma of interest with a non-vanishing component of the magnetic field normal to the wall. Line-tying is usually employed as an idealization of the boundary condition in some astrophysical plasmas, where the plasma density varies several orders of magnitude over a short distance. For example, the magnetic field lines in the coronae of stars and accretion disks are rooted in the dense, highly conducting gas below. In the limit of infinite density contrast, the dense gas may be treated as a conducting wall at which the field lines are tied. Line-tying imposes a strong constraint on the plasma motion, therefore in general is stabilizing. This is particularly true in the case of ideal magnetohydrodynamics (MHD). In ideal MHD, magnetic field lines are frozen-in to the plasma flow. If the footpoints of magnetic field lines are not moving, i.e., the wall is rigid, then the footpoint mapping following the field lines from one end to another is conserved. In resistive MHD, magnetic field lines are allowed to slip though the plasma on a resistive time scale. In this sense, the line-tying becomes imperfect, and the stabilizing effect becomes less stringent. It has long been speculated that solar coronal loops can remain stable for longer than typical MHD instability time scale due to line-tying stabilization. \\cite{Raadu1972,Hood1992} Line-tying may also be used in a broader context with a flow imposed on the boundary, to model the convection on the solar surface or the rotation of accretion disks. The imposed flow shuffles the footpoints of magnetic field lines, twists up and entangles them, thereby converting the kinetic energy of the flow into the magnetic energy in the field. Release of the stored magnetic energy powers some of the most splendid plasma phenomena such as solar flares and coronal mass ejections (CMEs). Dissipation of the magnetic energy on smaller scales is thought to be the energy source of the coronal heating. In a seminal paper, Parker argued that when the magnetic field line topology becomes complicated due to the footpoint shuffling, tangential discontinuities (current sheets) inevitably will form.\\cite{Parker1972} In the presence of dissipative mechanisms such as resistivity, tangential discontinuities will be smoothed out and the magnetic energy will be turned into heat. Parker coined the term {}``topological dissipation'' to describe this process. Following Parker's original suggestion, many authors \\cite{MikicSV1989,GalsgaardN1996,DmitrukGD1998,RappazzoDEV2006,RappazzoVED2007,RappazzoVED2008} have shown by direct numerical simulation that random shuffling of the coronal magnetic field lines by photospheric motions progressively increases the current density, resulting in sporadic energy release. The key ingredient of Parker's scenario of coronal heating is the ubiquitous presence of current sheets in a magnetized plasma. To illustrate his point, Parker considered a uniform magnetic field $B_{0}\\mathbf{\\hat{z}}$ line-tied to conducting plates at $z=\\pm L/2$, which represent the photosphere. Motions in the photosphere displace the field-line footpoints and entangle the field lines. Suppose we freeze the footpoints and let the system relax to the minimum energy state under the condition of preserving all the ideal topological constraints. Parker argues that in general the minimum energy state must contain tangential discontinuities. Since its first appearance in Ref. \\cite{Parker1972}, this simple setting has become the {}``standard model'' of the Parker problem. The property that magnetostatic equilibria tend to form tangential discontinuities is stated as a \\textquotedblleft{}Magnetostatic Theorem\\textquotedblright{} in Parker's book.\\cite{Parker1994} Although Parker gives compelling physical arguments in support of the Theorem, no rigorous proof is given. Parker's claim has stimulated considerable debate that continues to this day\\cite{vanBallegooijen1985,ZweibelL1987,Antiochos1987,LongcopeS1994a,NgB1998,LongbottomRCS1998,CraigS2005}, without apparent consensus. The point of contention is whether or not a magnetic field with neither null points, nor separatrices, nor closed field lines can develop tangential discontinuities. It is well established that null points, separatrices, or closed field lines are the preferential sites for current sheet formation when a system undergoes some instabilities or is deformed via boundary conditions.\\cite{Syrovatskii1971,RosenbluthDR1973,PriestR1975,Syrovatskii1981,Low1991,CowleyLS1997,ScheperH1998} In the Parker problem, however, none of these structures are present. Van Ballegooijen was the first to question Parker's claim. He demonstrated by an analytical argument that the equilibrium will always be continuous, unless the footpoint mapping is discontinuous.\\cite{vanBallegooijen1985,LongcopeC1996} Subsequently, several authors have found smooth equilibria using either linear theories for small footpoint displacements or Lagrangian numerical relaxation schemes for nonlinear problems. \\cite{ZweibelL1987,LongbottomRCS1998,CraigS2005,Wilmot-SmithHP2009} However, as pointed out by Ng and Bhattacharjee,\\cite{NgB1998} van Ballegooijen's argument is valid only for simple current sheets; it fails if a current sheet has multiple branches joined together (see Figure 2 of Ref. \\cite{NgB1998}). Ng and Bhattacharjee also proved a theorem which asserts that there is at most one smooth equilibrium for any given smooth footpoint mapping. It follows that if a smooth equilibrium becomes unstable, there is no other smooth equilibrium the system can relax to; therefore the system must relax to an equilibrium with tangential discontinuities. The theorem of Ng and Bhattacharjee is within the framework of reduced magnetohydrodynamics (RMHD), which assumes the existence of a strong guide field.\\cite{Kadomtsev1975,Strauss1976} It also assumes double-periodicity in the plane perpendicular to the guide field. It is not known at the present time whether the result can be generalized to full MHD or more general boundary conditions, although Aly \\cite{Aly2005} appears to have taken an important step towards that goal. Lagrangian relaxation studies indeed suggest current sheet may form under certain circumstances,\\cite{LongbottomRCS1998,CraigS2005} although it is rather difficult to draw decisive conclusions from numerical studies due to limits of spatial resolution. This paper gives an overview of our recent progress on MHD stability and current sheet formation in a line-tied system. Previously published results are summarized in Sec. \\ref{sec:Linear} and Sec. \\ref{sec:nonlinear}, with new perspectives added. Sec. \\ref{sec:Linear} summarizes the results on linear stability regarding line-tying effects on the ideal internal kink mode and resistive tearing mode. Parker's problem on current sheet formation is discussed in Sec. \\ref{sec:nonlinear}. The discussion is divided into two parts. The first part discusses why the nonlinear evolution of ideal kink mode is a relevant setting to test Parker's hypothesis. The second part is a discussion about the recent demonstration of current sheet formation by Low \\emph{et al}. \\cite{Low2006a,Low2007,JanseL2009,LowJ2009} and our objection to it.\\cite{HuangBZ2009} Some outstanding open questions will be discussed along the way. Sec. \\ref{sec:QSevol} contains the new results, where the governing equations for quasi-static evolution of a boundary driven, line-tied magnetic field are derived, and a possible strategy of applying these equations to fully settle an open question regarding the Parker problem as well as the demonstration of Low \\emph{et al}. is outlined. Finally, some future perspectives are discussed in Sec. \\ref{sec:openissue}. ", "conclusions": "" }, "1003/1003.3447.txt": { "abstract": "TWA 30 is a remarkable young (7~$\\pm$~3 Myr), low-mass (0.12~$\\pm$~0.04 M$_{\\sun}$), late-type star (M5~$\\pm$~1) residing 42~$\\pm$~2 pc away from the sun in the TW Hydrae Association. It shows strong outflow spectral signatures such as [S II], [O I], [O II], [O III], and Mg I], while exhibiting weak H$\\alpha$ emission ($-$6.8~$\\pm$~1.2 \\AA). Emission lines of [S II] and [O I] are common to T Tauri stars still residing in their natal molecular clouds, while [O III] and Mg I] emission lines are incredibly rare in this same population; in the case of TWA 30, these latter lines may arise from new outflow material colliding into older outflow fronts. The weak H$\\alpha$ emission and small radial velocity shifts of line emission relative to the stellar frame of rest (generally $\\lesssim$10 \\kms) suggest that the disk is viewed close to edge-on and that the stellar axis may be inclined to the disk, similar to the AA Tau system, based on its temporal changes in emission/absorption line strengths/profiles and variable reddening (A$_V$~=~1.5--9.0). The strong Li absorption (0.61~$\\pm$~0.13 \\AA) and common kinematics with members of the TWA confirm its age and membership to the association. Given the properties of this system such as its proximity, low mass, remarkable outflow signatures, variability, and edge-on configuration, this system is a unique case study at a critical time in disk evolution and planet-building processes. ", "introduction": "Over the past 30 years, a universal but complex paradigm of star formation has been developed through many observations and theoretical breakthroughs \\citep[e.g.,\\ ][]{1987ARA&A..25...23S, 1998apsf.book.....H}. Briefly, the stellar birth process begins with the gravitational collapse of material within giant molecular cloud complexes (Class 0 objects; \\citealt{1987ApJ...312..788A}) and progresses as the embedded protostar accretes material from an enshrouding infall envelope (Class I), with this transition lasting $\\sim$200 kyr \\citep{2008ApJ...684.1240E}. As the gas clears, the revealed protostar continues to accrete from a circumstellar disk of material (Class II, Classical T Tauri stars, cTTS; \\citealt{1989A&ARv...1..291A}). By $\\sim$10 Myr, accretion ceases \\citep{2009arXiv0911.3320F,2003MNRAS.342..876W}, and the object is now a pre--main sequence star, possibly with a debris disk (Class III and later) capable of assembling terrestrial planets over timescales of 10--100 Myr \\citep{2004E&PSL.223..241C}. This process from cloud core to pre--main sequence star is marked by a wide array of dynamic phenomena, including outflows and accretion. Recent evidence suggests that the accretion processes at work in Class II systems operate in the regime of very low-mass stars and brown dwarfs as well (e.g., \\citealt{2003ApJ...582.1109W,2005ApJ...626..498M}). Accretion rates scale roughly as $\\dot M \\sim M^2$ \\citep{2003ApJ...592..266M,2004A&A...424..603N,2004AJ....128.1294C,2005ApJ...626..498M,2006ApJ...639L..83A}, meaning lower mass systems should display weaker signatures of accretion. Forbidden emission lines (FELs) in optical spectra, arising from low-density optically thin gas, have been used to identify a handful of outflows originating from very low-mass stars and brown dwarfs (e.g., \\citealt{2001A&A...380..264F,2003ApJ...592..266M, 2004ApJ...616.1033L,2005ApJ...626..498M}). Only five young brown dwarfs have had their outflows spatially resolved -- three optical jets (via spectroastrometry; \\citealt{2005Natur.435..652W,2007ApJ...659L..45W,2009ApJ...691L.106W}) and two molecular outflows (via direct imaging; \\citealt{2008ApJ...689L.141P,2005ApJ...633L.129B}). Identification of such systems is hampered not only by their intrinsic faintness but also by the distances (d~$>$~120 pc) to the nearest star-forming regions (e.g., Sco-Cen, Taurus, Orion), making it difficult to resolve structure on the scales of disks ($\\sim$10--100 AU) and jets (a few hundred AU). Hence, progress in understanding the size, morphology, and energetics of jets powered by very low-mass stars and brown dwarfs is limited. Fortunately, the recent identification of new low-mass members of the nearby TW Hydrae Association (TWA) may provide more ideal systems for studies of disk and jet structures in this mass regime. Indeed, one of the three brown dwarfs with a resolved optical jet, 2MASSW J1207334$-$393254 (2M1207AB, also known as TWA 27AB; \\citealt{2002ApJ...575..484G,2004A&A...425L..29C}), is a member of the TWA. The namesake of the TWA, TW Hydrae, was the first cTTS found in isolation \\citep{1976ApJS...30..491H} -- located only 54~$\\pm$~6 pc from the Sun \\citep{2007A&A...474..653V}, 23 degrees above the Galactic Plane and far from any molecular cloud \\citep{1978ppeu.book..171H,1983A&A...121..217R,2009PASJ...61..585T}. More than a decade after the discovery of TW Hydrae, \\cite{1989ApJ...343L..61D} and \\cite{1992AJ....103..549G} found four more T Tauri stars in the same vicinity by using the IRAS Point Source Catalog and targeting stars with infrared excess. Based on their common X-ray activity, \\cite{1997Sci...277...67K} postulated that these five objects formed a physical association of young stars, which they termed the TW Hydrae Association. Soon after, targeted surveys using ROSAT All-Sky Survey (RASS) data \\citep{1999A&A...346L..41S,1999ApJ...512L..63W,2001ApJ...549L.233Z}, kinematic surveys \\citep{2003ApJ...599..342S,2005A&A...430L..49S}, and photometric near-infrared (NIR) surveys \\citep{2002ApJ...575..484G,2007ApJ...669L..97L} discovered several more members, bringing the total to 23 confirmed systems (\\citealt{2005ApJ...634.1385M}), of which five contain brown dwarfs. These systems share similar kinematics and indicators of youth such as strong chromospheric activity and, for late-K to M type members, Li I $\\lambda$6708 absorption. Isochronal ages derived from the Hertzsprung-Russell diagram yield age estimates of $\\sim$8 Myr for the TWA \\citep{1998ApJ...498..385S,1999ApJ...512L..63W,2000ApJ...530..867W,2006A&A...459..511B}. At an average distance of 53 pc (Mamajek, in preparation), the TWA is the nearest association containing actively accreting young stars\\footnote{Although, a recently identified young brown dwarf, 2MASS J0041353$-$562112, shows signs of accretion and may be a possible $\\beta$ Pic or Tuc-Hor member, placing it at $\\sim$35--50 pc \\citep{2009ApJ...702L.119R}.}, making it an attractive target for studying planet formation and the evolution of circumstellar disks. Moreover, members of the TWA span a range of disk evolutionary stages, with some members showing signs of actively accreting disks, passive/non-accreting disks, or debris disks, while other members show no signs of circumstellar material \\citep[e.g., ][]{1999ApJ...521L.129J,2005ApJ...631.1170L}. We have undertaken a survey to identify additional low-mass members to the TWA, which could provide important new case studies of accretion and outflow in substellar-mass objects residing at close distances ($\\sim$50 pc). In this paper, we report the discovery of a new, low-mass member of the TWA, 2MASS J11321831$-$3019518, which we term TWA 30. Its membership is confirmed by its kinematics, Li I $\\lambda$6708 absorption strength, and signatures of a low-gravity photosphere. TWA 30 has strong forbidden lines of [O I], [O II] and [S II] within 30 \\kms\\ of the stellar rest velocity, possibly indicating an outflow from a nearly edge-on system. The presence of [O III] and Mg I] in the spectrum of TWA 30 are extremely rare for a cTTS and have not been observed before in a young system residing outside of its natal molecular cloud. The fluctuating continua levels from both the optical and the NIR spectra suggest that they may be affected by highly variable reddening on timescales of a day to several weeks. In $\\S$ 2, we describe the discovery and observations of TWA 30; in $\\S$ 3, we analyze the kinematics, spectral morphology, emission/absorption lines, X-ray activity, and estimate an age and mass of TWA 30; in $\\S$ 4, we discuss evidence for this system having an inclined stellar axis to an edge-on disk and update the disk fraction of the TWA; finally in $\\S$5, we give our conclusions. ", "conclusions": "We have identified a new and unusual member of the TW Hydrae Association, TWA 30, which has emission lines of [O I], [O II], [O III], [S II], and Mg I] near the stellar rest velocity, indicating it is powering an outflow. The temporal changes in reddening, absorption/emission EWs, and line profiles, particularly for Na I D, suggests a stellar magnetospheric axis inclined with respect to a disk that is viewed nearly edge-on. The presence of forbidden line emission [OIII] and Mg I] in the spectra of TWA 30 marks the first time these lines have been seen in a young star not residing in a molecular cloud and are rarely present even amongst cTTS in star formation regions. Both the optical and NIR spectral types appear to vary in accordance with the reddening from M4--M5.25 and from M4.5--M6, respectively, so we have adopted a spectral type of M5~$\\pm$~1. We suggest that the earlier spectral types are seen at periods of high accretion and include continuum excess emission. The spectral type of TWA 30 therefore might be as late as M5.25--M6. From evolutionary models, we estimate an age of 7~$\\pm$~3 Myr and a mass of 0.12~$\\pm$~0.04 M$_{\\sun}$. The close proximity of TWA 30, 42~$\\pm$~2 pc, makes it an excellent target for follow-up studies to spatially resolve the outflow and circumstellar disk. With the inclusion of TWA 30, we have updated the disk census of the TWA, finding that 35$^{+11}_{-8}$\\% of observed TWA members still retain their circumstellar disks, a higher ratio than previous estimates." }, "1003/1003.2500_arXiv.txt": { "abstract": "Numerous authors have suggested that the ultra-high energy cosmic rays (UHECR) detected by the Pierre Auger Observatory and other cosmic-ray telescopes may be accelerated in the nuclei, jets or lobes of radio galaxies. Here I focus on stochastic acceleration in the lobes. I show that the requirement that they accelerate protons to the highest observed energies places constraints on the observable properties of radio lobes that are satisfied by a relatively small number of objects within the Greisen-Zat'sepin-Kuzmin (GZK) cutoff; if UHECR are protons and are accelerated within radio lobes, their sources are probably already known and catalogued radio galaxies. I show that lobe acceleration also implies a (charge-dependent) upper energy limit on the UHECR that can be produced in this way; if lobes are the dominant accelerators in the local universe and if UHECR are predominantly protons, we are unlikely to see cosmic rays much higher in energy than those we have already observed. I comment on the viability of the stochastic acceleration mechanism and the likely composition of cosmic rays accelerated in this way, based on our current understanding of the contents of the large-scale lobes of radio galaxies, and finally discuss the implications of stochastic lobe acceleration for the future of cosmic ray astronomy. ", "introduction": "\\label{intro} It has been known for many years (e.g. Hillas 1984) that the large-scale structures of radio-loud active galaxies are possible sites for the acceleration of the highest-energy cosmic rays yet to be detected, the ultra-high-energy cosmic rays (UHECR) with energies above a few $\\times 10^{19}$ eV. Radio galaxy jets, hotspots and lobes are particularly interesting to modellers, both because the synchrotron emission by which we see them in the radio already implies the presence of a high-energy particle population (albeit leptonic and of much lower energies) and therefore of a particle acceleration process, and because the physical conditions, in particular the magnetic field strength $B$, can either be estimated from equipartition or minimum energy arguments (Burbidge 1956) or, more recently, determined directly from observations of inverse-Compton emission (e.g. Hardcastle \\etal\\ 2002). It is thus reasonably easy to say whether any given component of a radio galaxy is capable of confining an energetic particle of a given energy and charge, a necessary precondition in almost all models of particle acceleration. The idea that radio-loud AGN might be the origin of the UHECR receives some tentative support from, or is at least consistent with, recent results from the Pierre Auger Observatory (PAO) suggesting that the spatial arrival directions of UHECR above $6 \\times 10^{19}$ eV are correlated with local AGN (Abraham \\etal\\ 2007). The imposition of this high low-energy cutoff on the cosmic rays ought to imply that they have a relatively local (within $\\sim 100$ Mpc) origin, since UHECR at these energies coming from larger distances would suffer strong attenuation due to interactions with the photons of the cosmic microwave background radiation (the so-called Greisen-Zat'sepin-Kuzmin or GZK cutoff; Greisen 1966) and also means that these UHECR undergo the smallest possible deflection in the Galactic and intergalactic magnetic fields. A particularly striking effect in the PAO data released in 2007 was the spatial coincidence between several of the UHECR and the position of the closest radio galaxy to us, Centaurus A (e.g. Moskalenko \\etal\\ 2009). While it is not yet clear whether the correlation with local AGN remains significant in the PAO data collected since 2007, updated versions of the Abraham \\etal\\ (2007) map appear to show a continued overdensity of UHECR around the position of Cen A (e.g. Fargion 2009). Meanwhile, several authors have suggested that the correlation between the arrival directions of UHECR in the original PAO dataset and the positions of local radio-loud AGN is at least as good as that with AGN in general (Nagar \\& Matulich 2008; Hillas 2009). How can specifically radio-loud AGN accelerate UHECR? It is of course possible that they are accelerated on sub-parsec scales, comparable to the scale of jet generation or initial collimation. The high photon and magnetic field energy densities expected close to the active nucleus provide important loss processes, but the acceleration efficiency might also be higher. Many authors have discussed mechanisms by which UHECR can be accelerated in the nuclear regions of Cen A and of radio galaxies in general (e.g.\\ Kachelrie\\ss, Ostapchenko \\& Tom\\`as 2009) but these necessarily rely on assumptions about the physical conditions close to the nucleus that are hard to test observationally. In what follows I therefore focus on the larger-scale components of radio-loud AGN. {\\it Direct} information about the leptonic particle acceleration processes in radio galaxies, derived from observations in the optical and X-ray where the synchrotron loss timescales are shorter than the transport timescales from the nuclei so that {\\it in situ} particle acceleration is required, implies that particle acceleration must be taking place in the hotspots of powerful double (Fanaroff \\& Riley 1974 class II, hereafter FRII) radio galaxies, and in the kpc-scale jets of the lower-power FRI class. FRII hotspots have traditionally been modelled as the terminal shocks of the relativistic, internally supersonic jet that extends up to Mpc scales in these objects (e.g. Blandford \\& Rees 1974; Heavens \\& Meisenheimer 1987; Meisenheimer \\etal\\ 1989), and, while optical and X-ray synchrotron evidence complicates this picture (e.g. Prieto \\etal\\ 2002; Wilson, Young \\& Shopbell 2001; Hardcastle \\etal\\ 2007a) it seems clear that they are particle acceleration sites. Moreover, their sizes and their magnetic field strengths, which can be measured very well via the inverse-Compton process in the most luminous systems where X-ray synchrotron emission is not a contaminant (e.g. Harris \\etal\\ 1994; Hardcastle \\etal\\ 2004) are certainly sufficient to allow UHECR to be confined (Hillas 1984). However, the space density of FRIIs is very low: we expect only a few within the GZK cutoff (for example, the nearest FRII in the northern sky, 3C\\,98, is at a distance of 134 Mpc) and so their effect on the PAO sky above $6 \\times 10^{19}$ eV is negligible. The numerically dominant population of radio galaxies, by several orders of magnitude, within 100 Mpc is composed of low-power FRI objects. Here the resolved particle acceleration region is typically the 100-pc to kpc-scale inner jet. Several nearby FRI radio galaxies, including Cen A (e.g., Hardcastle \\etal\\ 2003, 2007c; Goodger \\etal\\ 2010) and M87 (e.g., Perlman \\& Wilson 2005; Harris \\etal\\ 2006) have jets that are comparatively strong sources of X-ray synchrotron emission, allowing their particle acceleration properties to be studied in detail, while the evidence is consistent with the idea that all powerful FRI jets can accelerate leptons to the $>$ TeV energies required for X-ray synchrotron emission (e.g. Worrall \\etal\\ 2001). The picture that emerges from the X-ray observations is of a combination of strongly localized particle acceleration, which may be due to small-scale shocks, and a more diffuse process, which produces a different X-ray spectrum (and therefore a different electron energy spectrum) and which may therefore have different underlying acceleration physics. It has been argued, most recently by Honda (2009), that the Cen A jet is capable of accelerating protons to energies comparable to those of the PAO UHECR, which of course implies acceleration of heavy nuclei to even higher energies. This work relies on rather generous assumptions about the sizes and magnetic field strengths of the acceleration regions, though: as yet we have no direct constraint on the magnetic field strength in FRI jets (although TeV inverse-Compton emission should in principle provide one; Hardcastle \\& Croston, in prep.). This leaves us with the possibility of UHECR acceleration in the lobes, the largest-scale components of both FRI and FRII radio galaxies. At first sight these appear less promising candidates for UHECR acceleration, since there is little direct evidence for {\\it in situ} particle acceleration in the lobes. However, in the case of the 600-kpc giant lobes of Cen A (Hardcastle \\etal\\ 2009, hereafter H09) we showed that the high-frequency radio data from the {\\it Wilkinson Microwave Anisotropy Probe} ({\\it WMAP}) are consistent with the idea that the lobes contain at least some relatively energetic leptons; they do not rule out the idea that particle acceleration is ongoing at some level. Similarly, while we do not as yet have a robust inverse-Compton measurement of the magnetic field in the lobes of any FRI radio galaxy, the available limits in the case of Cen A constrain the field strength to be comparable to or greater than the equipartition value. H09 argued that the known size, and the limits on $B$, for the giant lobes meant that they could {\\it confine} protons of energies of order $10^{20}$ eV, and could therefore {\\it accelerate} protons to such energies, provided that a relatively efficient acceleration process was able to operate. We also showed that, provided that the energy index for the accelerated cosmic rays is relatively flat, the energetic requirements for the acceleration of the PAO UHECR plausibly associated with Cen A are trivially satisfied --- UHECR need only account for a small fraction of the total source energetics. Our preferred acceleration mechanism involved scattering off relativistic turbulence within the lobes, which requires the assumption that the internal energy density is not dominated by thermal particles (see also O'Sullivan \\etal\\ 2009) but is otherwise consistent with observations. We will return to the question of particle content and lobe energetics later in the paper, but in the next section I will show that a model in which UHECR are accelerated in the giant lobes is unique in providing some predictions for the spatial and energetic properties of UHECR which may already be testable using the PAO data. Finally, it should be noted that none of the above mechanisms are mutually exclusive. In fact, it seems highly likely that, in a source like Cen A, hadronic cosmic rays can be accelerated in the nucleus and the kpc-scale jet as well as in the giant lobes. Particles accelerated in the inner few kpc will eventually pass into the giant lobes and will then be confined (and potentially accelerated) there for some time before escaping. Hybrid models of this form potentially reduce the problems of acceleration purely in the lobes, by providing a seed population of cosmic rays at say $10^{17}$ -- $10^{18}$ eV and therefore reducing the required UHECR acceleration time in the lobes. A corollary of this, unfortunately, is that the ability of the giant lobes to {\\it confine} UHECR, irrespective of whether they can accelerate them, implies that the UHECR will be {\\it emitted} by a source like Cen A on scales of the giant lobes, whatever their original acceleration site. Even if all UHECR were generated at the nucleus, we would not expect a source like Cen A to appear `point-like' at the resolution of the PAO, so we cannot use the observed large-scale excess of UHECR around Cen A to argue that acceleration takes place either wholly or even partly in the giant lobes. This limitation should be borne in mind in what follows. The remainder of the paper is structured as follows. In Section \\ref{constraints} I show that the requirement that the lobes can confine high-energy particles gives a potentially interesting constraint on their radio luminosity, and argue that this means that if the PAO UHECR are protons they are likely to originate in a small number of bright nearby radio galaxies, all probably nearby well-studied objects. In Section \\ref{particle} I discuss our best existing constraints on the particle content of FRI lobes and the implications for cosmic ray acceleration and composition. Finally, in Section \\ref{outlook} I discuss the implications of a picture in which particles are accelerated in radio galaxy lobes for the future of cosmic ray astronomy. Throughout the paper I use a cosmology in which $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm m} = 0.3$ and $\\Omega_\\Lambda = 0.7$. The distance to Cen A is taken to be 3.7 Mpc (the mean of 5 distance estimates given in Ferrarese \\etal\\ 2007). ", "conclusions": "The principal results of this paper may be summarized as follows: \\begin{enumerate} \\item Stochastic acceleration of UHECR in the large-scale lobes of radio galaxies may be possible, but there are strong (though model-dependent) constraints on the properties of the radio galaxies that can accelerate them to the highest energies ($10^{20}$ eV). \\item These constraints imply that only a small number of local radio galaxies can be involved in the acceleration of UHECR, if the UHECR are protons, and that UHECR energies will cut off steeply around the energies currently being observed by the PAO; this model is testable in principle using existing radio surveys and up-to-date UHECR arrival positions, and is consistent with much of the available data. \\item However, if UHECR are heavy nuclei with $Z > 1$, as suggested by the latest PAO composition results, then many more radio galaxies can be sites of UHECR acceleration, and it may be that the nearest radio galaxy, Cen A, will be the only identifiable source in the cosmic-ray sky. \\end{enumerate}" }, "1003/1003.3992_arXiv.txt": { "abstract": "We present the results of an analysis of hard X-ray observations of the C2.7 solar flare detected by the $RT$-2 Experiment onboard the $Coronas-Photon$ satellite. We detect hard X-ray pulsations at periods of $\\sim$12 s and $\\sim$15 s. We find a marginal evidence for a decrease in period with time. We have augmented these results using the publicly available data from the $RHESSI$ satellite. We present a spectral analysis and measure the spectral parameters. ", "introduction": "Quasi-Periodic Pulsations (QPPs), a common feature of solar flare emission, have been observed for many years (\\citet{Young61}) in all frequency bands ranging from radio to hard X-rays with periodicities varying from a few milliseconds to several seconds (\\citet{Aschwanden87}; \\citet{Fleishman02}; \\citet{Tan08}). The long period QPPs (periodicity $>$ 10 s) observed in the microwave emission of solar flares are also seen in hard X-rays (\\citet{Kane83}; \\citet{Nakajima83}; \\citet{Asai01}; \\citet{Nakariakov03}) and they could be resulting from some MHD oscillations in the source region or due to modulation of electron acceleration and injection mechanisms. QPPs are generally associated with the emission from the flare accelerated non-thermal electrons, because thermal parameters are not expected to show sudden changes and pulsations. QPPs can be observed in all the stages of a flare, prominently in hard X-rays, microwave and white light emissions. MHD oscillations in the magnetic loop can cause modulation in the magnetic field strength and magnetic mirror ratio (\\citet{Zaitsev82}, \\citet{Zaitsev89}; \\citet{Zimovets09}) resulting in periodic variation towards the flare foot points. Generally, most of the non-thermal hard X-rays are observed at the foot points of the magnetic loops involved in the magnetic reconnection. This emission is produced by the accelerated non-thermal electrons from the reconnection region hitting the foot points of the magnetic loops. The oscillations in the accelerated electrons in turn reflect as QPPs in the hard X-ray emission profile. Similarly, microwave emission produced by the interaction of the accelerated electrons with the magnetic field as a result of gyro-synchrotron process also show QPPs (like hard X-rays, as both the emissions are from the same population of electrons). Observationally, QPPs can be seen in the light curves of the solar flare emission in the respective wave-bands. The profiles of photons or energy fluxes of hard X-rays associated with the non-thermal electrons show QPPs or damping oscillations. Since the basic cause of QPPs have implications for particle acceleration mechanism, it is important to investigate QPPs at diverse source intensities. In this paper, we present the results obtained from the observation of the C2.7 solar flare detected on 2009 July 5 using the RT-2 experiment onboard the Coronas-Photon satellite. Since this is the first result from this experiment, we describe in detail the methodology used in deriving the response matrix and spectral fitting. We augment our results by using the publicly available $RHESSI$ data. We examine the spectral and temporal characteristics of the flare and investigate the implications to the electron acceleration mechanisms. In \\S 2 a brief description of the RT-2 experiment is given. Observations and analysis results (RT-2 and RHESSI data) are given in \\S 3 and finally in \\S 4 a detailed discussion of the results are presented along with relevant conclusions. ", "conclusions": "\\citet{Jakimiec09} have investigated QPPs in about 50 flares using $Yohkov$ and $BATSE$ hard X-ray data and have derived a correlation between the QPP periods (ranging from 10 s to 150 s) and sizes of loop-top sources. From the RHESSI data we derive a size of the X-ray emitting region in the 6 -- 12 keV region of 7\" (5 Mm), corresponding to the 50\\% level of the peak emission, for the 2009 July 5 flare. This size agrees with the correlation derived by \\citet{Jakimiec09}. \\citet{Desai87} detected fast oscillations in several solar hard X-ray flares and observed magneto-hydrodynamic signature of the loop dynamics. \\citet{Jakimiec09} conclude that the hard X-ray oscillations are confined to the loop-top sources and the observations are described with a model of oscillating magnetic traps. \\citet{Fleishman08} have made a detailed analysis of the 2003 June 15 solar flare (GOES X1.3 class) and detected hard X-ray (based on $RHESSI$ data) and microwave oscillations with periods ranging from 10s to 20 s. They, however, conclude that QPPs are associated with quasi-periodic acceleration and injection of electrons. The possible detection of a decreasing trend in the periodicity can put further constraints on the magneto-hydrodynamic models. Several flare observations as well as numerical simulation studies have been reported on the periodic and quasi-periodic oscillations of flare intensity in the radio and X-ray energy bands. Such oscillations show the typical size of reconnection site, configuration of loops formed during the reconnection, and plasmoid or CME launched above the reconnection X-point. As shown in Figure 2, the rising phase of the flare shows the nominal exponential increase. However, as the flare attains the peak level, the intensity goes through moderate quasi-periodic oscillations, which are more prominent in the 20 -- 35 keV energy band. Moreover, the modulation index (i.e., the degree of quasi-periodic oscillation) is higher at the low-energy band (i.e., $\\sim$13.5\\% in the 20 -- 35 keV band) than in the high-energy band ($\\sim$7\\% above 35 keV bands). It shows the production of copious amount of electrons over a limited range of energies. The flare profile observed at 15.4 GHz correlates with the rising phase of the flare, but the oscillations are not clearly seen, which may be a limitation imposed by the sensitivity of the measurement. The white-light images from LASCO, associated with this flare, show a rather slow moving CME (i.e., speed in the range of 50 -- 150 km s$^{-1}$). This is consistent with our finding of the production of particles in a limited energy range. A comparison of profiles shown in Figure 2 with the RHESSI spectrum reveals a gradual steepening of the spectrum from the flare rising phase to the start of the decay phase, although the average intensity of the flare remained nearly same level in this period. Thus, most of the accelerated electrons have been generated and injected from the reconnection site. This is the brightest solar flare detected by the RT-2 experiment in the first ten months of operation. Several other solar flares, particularly during the eruptions that have taken place from 2009 Oct 22 and Nov 2, are also recorded. A detailed investigation on faint flares during this solar minimum are going on and a flare list would be published separately. From 2009 December onwards, communication to the satellite is not responding, though attempts are on to revive the system." }, "1003/1003.1359_arXiv.txt": { "abstract": "Oppenheimer and Snyder found in 1939 that gravitational collapse in vacuum produces a ``frozen star\", i.e., the collapsing matter only asymptotically approaches the gravitational radius (event horizon) of the mass, but never crosses it within a finite time for an external observer. Based upon our recent publication on the problem of gravitational collapse in the physical universe for an external observer, the following results are reported here: (1) Matter can indeed fall across the event horizon within a finite time and thus BHs, rather than ``frozen stars\", are formed in gravitational collapse in the physical universe. (2) Matter fallen into an astrophysical black hole can never arrive at the exact center; the exact interior distribution of matter depends upon the history of the collapse process. Therefore gravitational singularity does not exist in the physical universe. (3) The metric at any radius is determined by the global distribution of matter, i.e., not only by the matter inside the given radius, even in a spherically symmetric and pressureless gravitational system. This is qualitatively different from the Newtonian gravity and the common (mis)understanding of the Birkhoff's Theorem. This result does not contract the ``Lemaitre-Tolman-Bondi\" solution for an external observer. ", "introduction": "Oppenheimer \\& Snyder\\cite{OS39} studied the problem of black hole (BH) formation from gravitational collapse and arrived at two conclusions which have deeply influenced our understanding of astrophysical BH formation ever since. The first conclusion is: ``The total time of collapse for an observer comoving with the stellar matter is finite.\" However it should be realized that the comoving observer is also within the event horizon with the collapsing matter, once a BH is formed. The second and last conclusion of the paper is: ``An external observer [{\\bf O} hereafter] sees the star asymptotically shrinking to its gravitational radius [the radius of the event horizon of the BH of the same mass, $R_{\\rm H}=2GM/c^2$ hereafter.].\" This means that {\\bf O} will never witness the formation of an astrophysical BH. Given the finite age of the universe and the fact that all observers are necessarily external, the last conclusion of Ref.~\\refcite{OS39} seems to indicate that astrophysical BHs cannot be formed in the physical universe through gravitational collapse. Recently, Vachaspati, Stojkovic \\& Krauss\\cite{Krauss} have stressed that ``The process of BH formation is generally discussed from the viewpoint of an infalling observer. However, in all physical settings it is the viewpoint of the asymptotic observer [i.e. {\\bf O}] that is relevant.\" They analyzed the process of the self-collapse of a domain wall (a massive shell with no thickness) and concluded that {\\bf O} sees the domain wall asymptotically shrinking to $R_{\\rm H}$, i.e., a BH is never formed within a finite time to {\\bf O}. This is a further confirmation to the conclusion of Ref.~\\refcite{OS39}. Vachaspati et al.\\cite{Krauss} then went on to study the quantum mechanical effect of the contracting shell and found that the matter accumulating just outside $R_{\\rm H}$ actually produces radiation, which they called pre-Hawking radiation. They concluded that ``Evaporation by pre-Hawking radiation implies that {\\bf O} can never lose objects down a BH.\" Combining the above two works separated by nearly 70 years, a very surprising scenario seems inevitable: Gravitational collapse will not produce BHs, but result in complete conversion of matter into radiation. This scenario, if correct, would have profound implications to our understanding of general relativity which has long been considered to robustly predict the existence of BHs, as well as a vast amount of astronomical observations which can, and perhaps only, be understood by invoking BHs\\cite{Zhang}. However, both of the above works are over-simplified and do not catch all the essence of gravitational collapse in the physical universe, because both investigations only considered gravitational contraction in vacuum and the work of Ref.~\\refcite{Krauss} did not allow a finite thickness of the contracting shell. To overcome the drawbacks of these two works discussed above, Liu \\& Zhang\\cite{Liu} studied the gravitational collapse of a single shell and double-shells onto a pre-existing BH; these shells can have finite thicknesses and the outer shell in the double-shell case mimics the matter outside the collapsing shell in the physical universe. The gravitational contractions studied in the two previous works can be considered as special cases of that studied in Ref.~\\refcite{Liu}. The main conclusion of Liu \\& Zhang\\cite{Liu} is that matter does not accumulate outside $R_{\\rm H}$, but instead falls straight across it, within a finite time of {\\bf O}. In the rest of this paper, we first review briefly the main results in Ref.~\\refcite{Liu}, and then discuss several issues related to the ``frozen star\" paradox, nature of astrophysical BHs, gravitational singularity in the physical universe, and finally applicability of the Birkhoff's theorem. All calculations and discussions in this report are within the framework of Einstein's general relativity. ", "conclusions": "" }, "1003/1003.6115_arXiv.txt": { "abstract": "Marginal likelihoods for the cosmic expansion rates are evaluated using the `Constitution' data of 397 supernovas, thereby updating the results in some previous works. Even when beginning with a very strong prior probability that favors an accelerated expansion, we obtain a marginal likelihood for the deceleration parameter $q_0$ peaked around zero in the spatially flat case. It is also found that the new data significantly constrains the cosmographic expansion rates, when compared to the previous analyses. These results may strongly depend on the Gaussian prior probability distribution chosen for the Hubble parameter represented by $h$, with $h=0.68\\pm 0.06$. This and similar priors for other expansion rates were deduced from previous data. Here again we perform the Bayesian model-independent analysis in which the scale factor is expanded into a Taylor series in time about the present epoch. Unlike such Taylor expansions in terms of redshift, this approach has no convergence problem. ", "introduction": "It is generally accepted that a more appropriate way to measure the acceleration of expansion of the universe is to resort to a cosmographic or model-independent analysis. In the conventional model-based analyses of distance modulus-redshift ($\\mu-z$) data of Type Ia supernova (SN Ia), the accelerated expansion of the universe is an indirect inference based on the best fit values of parameters, such as the density parameters $\\Omega_m$, $\\Omega_{\\Lambda}$, etc. On the other hand, in a model-independent approach, the scale factor $a(t)$ is expanded as a Taylor series in time about the present epoch and the marginal likelihoods of its coefficients are computed using the data. The marginal likelihood for the deceleration parameter gives an estimate of the acceleration of cosmic expansion. Since practically one has to truncate the series to some finite order, the basic assumption here is that $a(t)$ is expressible as a truncated Taylor series or polynomial. Evaluating the deceleration parameter by adopting this method, it was confirmed model-independently that the universe is undergoing an accelerated expansion \\citep{mvj1,mvj2}. In this paper we report the updating of the marginal likelihood for each of the expansion coefficients found in the above work. This is performed for the case of a fifth order polynomial. A notable result in the present Bayesian model-independent analysis is that even when beginning with a very strong prior probability that favors an accelerated expansion, the marginal likelihood for the deceleration parameter $q_0$ is found peaked around $q_0=0$ in the spatially flat case. It is also found that the new data significantly constrains the cosmic expansion rates appearing in the Taylor expansion, when compared to the previous data. We also note that successive terms in the series decreases sufficiently fast, thereby verifying the assumption of a converging Taylor series in time for the cosmic scale factor. Other model-independent approaches, which Taylor expand the distance modulus $\\mu$ in terms of redshift $z$, have also gained attention in recent years [See for eg. \\citep{turner,visser,lima,seikel}]. But a drawback of this method is that, in principle, it converges only for $\\mid z\\mid <1$ \\citep{visser,lima}. The argument behind this assertion is as follows: For an expanding universe, $z<0$ corresponds to the future and $z=-1$ is the redshift when the universe has expanded to infinite size. Since $z=-1$ is a pole, by standard complex variable theory, the radius of convergence of a series about $z=0$ is atmost $\\mid z\\mid =1$, so that it fails to converge for $z>1$. When compared to this, our approach of expanding the scale factor in terms of $t$ about the present epoch $t_0$ is advantageous, for the series converges for all times. Even the lookback time $T\\equiv t-t_0$ is evaluated by numerically solving an equation which involves a Taylor series in time. Hence there is no convergence problem in the present work. However, it should be noted that all analyses which make use of such Taylor expansions, in practice, employ polynomials and hence convergence is not a serious problem. For instance, one can see that there is convergence in certain special cases of the low order polynomial fit by \\cite{lima}. ", "conclusions": "We assumed that a Taylor series form for the scale factor $a(t)$ is valid and attempted to find the coefficients in this expansion using the recent Constitution SN data. The new marginal likelihoods obtained for its coefficients give valuable information regarding the expansion history of the universe. It is found that there is significant constraining of these parameters when compared to previous analyses using the data in \\citep{perl1}. The shift in the computed mean value of the deceleration parameter $q_0$, from that found in the previous analysis is noteworthy. Even when we start with a prior probability distribution that strongly favors an accelerating universe, the marginal likelihood for the deceleration parameter obtained from the present analysis using the Constitution data is found peaked around $q_0=0$. However, we reiterate that the considerable spread still found in the likelihoods of these parameters indicate freedom in the choice of their numerical values. A distinguishing feature of our analysis is that the marginal likelihoods for each parameter obtained in the previous case is chosen as the prior probability distribution in the present one, thereby implementing the Bayesian method in true spirits. The work is also intended as a demonstration of this fundamental requirement in Bayesian analysis. However, we have noted that the results obtained in this paper may heavily depend on the prior chosen for $h$. Thus it is important to evaluate expansion rates using prior for $h$ deduced from other measurements too. It is expected that in future when the SN dataset becomes large enough, the expansion coefficients get sharply peaked marginal likelihoods and become the most basic model-independent description of the expansion history of the universe." }, "1003/1003.1735_arXiv.txt": { "abstract": "We present a new signature by which to one could potentially discriminate between a spectrum of gravitational radiation generated by a self-ordering scalar field vs that of inflation, specifically a comparison of the magnitude of a flat spectrum at frequencies probed by future direct detection experiments to the magnitude of a possible polarization signal in the Cosmic Microwave Background (CMB) radiation. In the process we clarify several issues related to the proper calculation of such modes, focusing on the effect of post-horizon-crossing evolution. ", "introduction": " ", "conclusions": "" }, "1003/1003.4332_arXiv.txt": { "abstract": "The cumulative luminosity distribution functions (CLFs) of radio millisecond pulsars (MSPs) in globular clusters (GCs) and in the Galactic field at a frequency of 1.4 GHz have been examined. Assuming a functional form, $N \\propto L^q$ where $N$ is the number of MSPs and $L$ is the luminosity at 1.4 GHz, it is found that the CLFs significantly differ with a steeper slope, $q=-0.83 \\pm 0.05$, in GCs than in the Galactic field ($q=-0.48 \\pm 0.04$), suggesting a different formation or evolutionary history of MSPs in these two regions of the Galaxy. To probe the production mechanism of MSPs in clusters, a search of the possible relationships between the MSP population and cluster properties was carried out. The results of an investigation of 9 GCs indicate positive correlations between the MSP population and the stellar encounter rate and metallicity. This provides additional evidence suggesting that stellar dynamical interactions are important in the formation of the MSP population in GCs. ", "introduction": "} It is well known that the formation rate per unit mass of low-mass X-ray binaries (LMXBs) is orders of magnitude greater in globular clusters (GCs) than in the Galactic field (Katz 1975; Clark 1975). The high formation rate of LMXBs in GCs is attributed to the frequent dynamical interactions in the dense stellar environment. As an additional formation channel, the binaries in a GC can also be a result of a standard evolutionary path identified for MSP formation in the Galactic field. This has stimulated many theoretical and observational studies to investigate the relative contribution of these two formation processes of compact binaries in the population of GCs (e.g. Fregeau 2008; Pooley et al. 2003; Pooley \\& Hut 2006). With the superior sub-arcsecond spatial resolution of the \\emph{Chandra X-Ray Observatory}, remarkable progress has been made in the understanding of the formation processes of close binaries in GCs. For example, Pooley et al. (2003) found a positive correlation between the number of close X-ray binaries in GCs and the stellar encounter rate, $\\Gamma_{\\rm c}$. Specifically, Pooley et al. (2003) found an approximately linear relationship between the number of LMXBs and $\\Gamma_{\\rm c}$, indicating a dependence on the properties of GCs. A similar relationship has also been reported by Gendre et al. (2003) and taken together with the results of Pooley et al. (2003) provide evidence for the dynamical origin of LMXBs in GCs. Since millisecond pulsars (MSPs) have long been proposed as the descendants of LMXBs, they are also expected to have a dynamical origin in GCs. Due to the existence of extensive pulsar surveys, 140 MSPs have been detected in 26 different clusters and a statistical study of their relationship to cluster parameters is desirable. \\footnote{For updated statistics, please refer to http://www2.naic.edu/$\\sim$pfreire/GCpsr.html} However, previous studies were not successful in finding evidence for the dynamical origin of MSPs in the clusters due to the lack of a relation between the pulsar population and $\\Gamma_{\\rm c}$ in the GCs (e.g. Ransom 2008). This can be ascribed to the observational bias in the pulsar searches. As the distance of the GCs spans a rather wide range (cf. Harris 1996), the sensitivity of the observations can differ in the searches toward different clusters and hence induce selection effects in the observed sample (Ransom 2008). Therefore, the observed number of MSPs is not representative of an unbiased sample for the analysis. In this paper, we present a method to alleviate the aforementioned problem and investigate the possible relationship between the number of MSPs and the cluster properties. In \\S2, an investigation of the cumulative luminosity distribution functions of MSPs in a number of selected GCs is carried out. We subsequently use the obtained results in a correlation analysis in \\S3 and discuss the physical implications of the possible correlation in \\S4. ", "conclusions": "} The CLFs of nine GCs, each containing a population of MSPs has been examined. Upon comparison of the MSP population in GCs with that in the Galactic field, it has been found that the slopes of the CLFs inferred in these two populations significantly differ. It is natural to speculate that the CLF is somehow related to the magnetic field and spin of the MSPs. Wang, Jiang \\& Cheng (2005) have compared the distributions of the spin period and the dipolar surface magnetic field for both cluster and disk populations (cf. Fig.~2 and Fig.~3 in their paper). Despite the broader distribution for the disk population, their mean values are not dissimilar in both populations and therefore cannot solely explain the difference of CLFs. Apart from the radio luminosity functions, the X-ray emission properties of the MSPs in the GCs are also found to be very different from those in the Galactic field. While the MSPs in the Galactic field generally require a hot polar cap component plus a non-thermal power-law tail to model their X-ray spectra (cf. Zavlin 2006), the X-rays from a majority of the MSPs in GCs are purely thermal in nature (see Hui et al. 2009 and the references therein for a recent review). Cheng \\& Taam (2003) suggest the absence of non-thermal X-ray from the cluster MSPs can be possibly related to the complicated magnetic field structure. Since the stellar interaction in GCs is much more frequent than that in the Galactic field, MSPs in the GCs can possibly change their companion several times throughout their lives. Since the orientation of the binary after each exchange can differ, the direction of the angular momentum accreted during the mass transfer phase subsequent to each exchange can vary possibly affecting the magnetic field configuration at the neutron star surface. Such an evolution could lead to a much more complicated magnetic field structure for the MSPs in the GCs than in the case in the Galactic field. In such a complicated magnetic field, Ruderman \\& Cheng (1988) have argued that high energy curvature photons will be emitted and subsequently converted into pairs to quench the accelerating region. This provides an explanation for the absence of non-thermal emission in the cluster MSPs. For the same reason, the complicated magnetic field structure can also possibly alter the coherent radio emission and result in a different radio luminosity of the cluster MSPs in comparison with the disk population. Adopting the best-fit normalization inferred from the CLFs of individual cluster as an unbiased estimate of the number of MSPs, we have further examined the relationships between the pulsar population and the physical properties in GCs. We have found the positive correlations of $N_{0}$ versus $\\Gamma_{\\rm c}$ as well as $N_0$ versus [Fe/H] at a relatively high confidence level. A marginal positive correlation between $N_{0}$ and $v_{\\rm escape}$ is also suggested. Although a high escape speed implies the presence of a deeper gravitational potential well and hence a higher neutron star retention, this correlation is not sufficiently significant to warrant such an interpretation. Hence, we do not discuss this relation any further and focus on the physical implications of the $N_{0}-\\Gamma_{\\rm c}$ and $N_{0}-$[Fe/H] relations. Due to the different selection effects in the pulsar search surveys, it is not feasible to directly use the detected MSP populations in GCs for a statistical analysis. Instead, we alleviate the problem by taking $N_{0}$ as the estimator for the number of pulsars with pseudo radio luminosites at 1.4~GHz larger than 1~mJy~kpc$^{2}$. With this consideration, we have found a correlation between $N_{0}$ and $\\Gamma_{\\rm c}$ at a confidence level $>98\\%$. We have further found that the strength of this correlation is robust and independent of the choice of the luminosity cut-off by repeating the analysis with different thresholds. This provides evidence for the dynamical formation of MSPs in GCs. For a competing scenario that the MSPs have a binary origin similar to the Galactic field, one should expect the number of MSPs to scale with the cluster mass, $M_{GC}$, instead of $\\Gamma_{\\rm c}$. However, we do not find any convincing relationship between $N_{0}$ and $M_{GC}$ (see Table~\\ref{correl}). The absence of correlation with $M_{GC}$ provides additional support for the dynamical formation scenario. Taken together with the difference in the X-ray luminosity functions of LMXBs in the field and in globular clusters (see Voss et al. 2009; Kim et al. 2009), it is likely that the MSPs have different origins/evolutions in globular clusters relative to the Galactic field. We note that the logarithmic slope of the power-law fit in the $N_{0}-\\Gamma_{\\rm c}$ relationship (i.e. $0.69\\pm0.11$) is not dissimilar to that of the number of X-ray sources versus $\\Gamma_{\\rm c}$ ($0.74\\pm0.36$ Pooley et al. 2003). This dependence on the two-body encounter rate suggests a possible relationship between the MSP population and close X-ray binaries in GCs. Apart from the whole X-ray binary population, Pooley et al. (2003) and Gendre et al. (2003) have also examined the relationship for the individual class of LMXBs which has a logarithmic slope of $0.97\\pm0.5$. Although the large uncertainty of this slope resulting from the limited sample of LMXBs precludes a definitive conclusion concerning the link between LMXBs and MSPs, it is consistent with such an interpretation. Theoretical arguments (Verbunt \\& Hut 1987) suggest that the number of LMXBs is linearly proportional to the stellar encounter rate of the cluster, however direct comparison of their relationship with the current two-body encounter rate may be misleading. As the MSPs are long lived and are produced by the previous generations of LMXBs, they can have a different formation rate from the LMXB population currently observed. This point is important since the relaxation time at the cluster core is generally longer than the lifetime of LMXBs (cf. Harris 1996). Therefore, the continuous mass segregation at the cluster center can result in a evolution of the stellar collision frequency and hence a varying formation rate of compact binary systems. Nevertheless, the combination of X-ray and HST observations of Cen A (see Jord\\'{a}n et al. 2007) indicate that globular clusters with LMXBs are characterized by higher stellar encounter rates than those devoid of LMXBs. In addition to the $N_{0}-\\Gamma_{\\rm c}$ relation, we have also found a positive correlation between $N_{0}$ and the metallicity of the GCs. It has been noted that observational evidence suggests that bright LMXBs are preferably formed in metal-rich clusters in our Milky Way as well as other galaxies (e.g Bellazzini et al. 1995; Maccarone et al. 2004; Jord\\'an et al. 2004). Ivanova (2006) proposes that the absence of the outer convective zone in metal-poor main sequence donor stars in the mass range of $0.85\\msun$ - $1.25 \\msun$, in comparison to their metal rich counterparts can be responsible, since the absence of magnetic braking in such stars precludes orbital shrinkage, thereby, significantly reducing the binary parameter space for the production of bright LMXBs. For the conventional scenario that LMXBs are the progenitors of MSPs, the positive correlation between $N_{0}$ and [Fe/H] is not unexpected since the MSP number should scale with that of their progenitors. While the stellar encounter rate has been widely accepted as a parameter to indicate which clusters are likely to host a large MSP population, our study suggests that the metallicity can also be an important parameter. To explore this hypothesis, we suggest that pulsar searches be carried out toward metal-rich GCs, such as Liller~1 which has the highest metallicity ([Fe/H]=0.22) among all 150 GCs in the Milky Way (cf. Harris 1996). Furthermore, its two-body encounter rate is estimated to be comparable with that of 47~Tuc. Therefore, according to these parameters, it is very likely to host a considerable number of MSPs. With a dedicated search, this hidden population may be revealed." }, "1003/1003.1721_arXiv.txt": { "abstract": "The aim of the present study is to test whether the cold accretion of gas through a \"cosmic filament\" \\citep{Mac06} is a possible formation scenario for the polar disk galaxy NGC 4650A. If polar disks form from cold accretion of gas, the abundances of the HII regions may be similar to those of very late-type spiral galaxies, regardless of the presence of a bright central stellar spheroid, with total luminosity of few $10^{9} L_{\\odot}$. We use deep long slit spectra obtained with the FORS2 spectrograph at the VLT in the optical and near-infrared wavelength ranges for the brightest HII regions in the disk polar disk of NGC 4650A. The strongest emission lines ([OII] $H_{\\beta}$, [OIII], $H_{\\alpha}$) were used to derived oxygen abundances, metallicities and the global star formation rates for the disk. The deep spectra available allowed us to measure the Oxygen abundances ($12 + log (O/H)$) using the {\\it Empirical method} based on intensities of the strongest emission lines, and the {\\it Direct method}, based on the determination of the electron temperature from the detection of weak auroral lines, as the [OIII] at 4363 \\AA. The Oxygen abundance measured for the polar disk is then compared with those measured for different galaxy types of similar total luminosities, and then compared against the predictions of different polar ring formation scenarios. The average metallicity values for the polar disk in NGC 4650A is $Z=0.2 Z_{\\odot}$, and it is lower that the values measured for ordinary spirals of similar luminosity. Moreover the gradient of the metallicity is flat along the polar disk major axis, which implies none or negligible metal enrichment from the stars in the older central spheroid. The low metallicity value in the polar disk NGC 4650A and the flat metallicity gradient are both consistent with a later infall of metal-poor gas, as expected in the cold accretion processes. ", "introduction": "\\label{intro} The hierarchical, merger-dominated picture of galaxy formation is based on the Cold Dark Matter (CDM) model \\citep{Col00}, which predicts that the observed galaxies and their dark halo (DH) were formed through a repeated merging process of small systems (\\citealt{Del06}; \\citealt{Gen08}). In this framework, major and minor mergers of disky systems do play a major role in the formation of spheroid and elliptical galaxies (\\citealt{Naa07}; \\citealt{Bou07}), in all environments and from the Local Group to high-redshift universe \\citep{Con03}. The gas fraction is a key parameter in the physics of such gravitational interactions: if it is high enough, an extended and massive disk structure can survive (\\citealt{Spr05}; \\citealt{Rob08}). Galaxies can get their gas through several interacting processes, such as smooth accretion, stripping and accretion of primordial gas, which are equally important in the growth of galaxies. Recent theoretical works have argued that the accretion of external gas from the cosmic web filaments, with inclined angular momentum (\\citealt{Dav01}, \\citealt{Sem05}), might be the most realistic way by which galaxies get their gas. This process may also explain the build-up of high redshift disk galaxies (\\citealt{Ker05}, \\citealt{Ker08}, \\citealt{Bro08}, \\citealt{Dek09}, \\citealt{Bou09}). The relative share of all gravitational interactions depends on the environments and it drives many morphological features observed in galaxies, such as bars and polar rings. Galaxies with polar rings (PRGs) generally contain a central featureless stellar spheroid and an elongated structure, the ``polar ring'' (made up by gas, stars and dust), which orbits in a nearly perpendicular plane to the equatorial one of the central galaxy \\citep{Whi90}. The decoupling of the angular momentum of the polar structure and the central spheroid cannot be explained by the collapse of a single protogalactic cloud: thus a ``second event'' must have happened in the formation history of these systems. This is the reason why studying PRGs promise to yield detailed information about many of the processes at work during galaxy interactions and merging (\\citealt{Iod02}, \\citealt{Iod02A}, \\citealt{Res94}, \\citealt{Res02}, \\citealt{Bou03}). The debate on the origin of PRGs is still open and two main processes have been proposed {\\it i)} a major dissipative merger or {\\it ii)} gas accretion. In the merging scenario, the PRG results from a ``polar'' merger of two disk galaxies with unequal mass, (\\citealt{Bek97}; \\citealt{Bek98}; \\citealt{Bou05}): the morphology of the merger remnants depends on the merging initial orbital parameters and the initial mass ratio of the two galaxies. In the accretion scenario, the polar ring may form by a) the disruption of a dwarf companion galaxy orbitating around an early-type system, or by b) the tidal accretion of gas stripping from a disk galaxy outskirts, captured by an early-type galaxy on a parabolic encounter (\\citealt{Res97}; \\citealt{Bou03}; \\citealt{Han09}). In the latter case, the total amount of accreted gas by the early-type object is about $10\\%$ of the gas in the disk donor galaxy, i.e. up to $10^{9}$ $M_{\\odot}$. Both major merger and accretion scenarios are able to account for many observed PRGs morphologies and kinematics, such as the existence of both wide and narrow rings, helical rings and double rings \\citep{Whi90}. Very recently, a new mechanism has been proposed for the formation of wide disk-like polar rings: a long-lived polar structure may form through cold gas accretion along a filament, extended for $\\sim 1$ Mpc, into the virialized dark matter halo \\citep{Mac06}. In this formation scenario, there is no limits to the mass of the accreted material, thus a very massive polar disk may develop either around a stellar disk or a spheroid. \\citet{Bro08}, by using high-resolution cosmological simulations of galaxy formation, have confirmed and strengthened the formation scenario proposed by \\citet{Mac06}. In this case polar disk galaxies can be considered as extreme examples of angular momentum misalignment that occurs during the hierarchical structure formation. In the merging history, an inner disk formed first after the last major merger of two galaxies with a 1:1 mass ratio, then, due to the high gas fraction, the galaxy rapidly forms a new disk whose angular momentum depends on the merger orbital parameters. At later times, gas continues to be accreted along the disk which may be also in a plane perpendicular to the inner disk. The morphology and kinematics of one simulated object, in both simulations, are similar to those observed for NGC4650A: in particular, \\citet{Bro08} found that such formation mechanism can self-consistently explain both morphology and kinematics of central spheroid and polar structure, and all the observed features (like colors and colors gradient, longevity, spiral arms, HI content and distribution). NGC~4650A is the prototype for PRGs (see Fig. \\ref{slit}). Its luminous components, inner spheroid and polar structure were studied in optical and near-infrared (NIR) photometry, spectroscopy, and in the radio emission, HI 21 cm line and continuum (\\citealt{Arn97}, \\citealt{Gal02}, \\citealt{Iod02}, \\citealt{Swa03}, \\citealt{Iod06}). The polar structure in NGC~4650A is a disk, very similar to that of late-type spirals or LSB galaxies, rather than a ring. The polar disk stars and dust can be reliably traced to $\\sim 1.2$ kpc radius from the galaxy nucleus, and the surface brightness profiles have an exponential decrease (\\citealt{Iod02}, \\citealt{Gal02}). Furthermore, the rotation curves measured from the emission and absorption line optical spectra are consistent with those of a disk in a differential rotation rather than that of a narrow ring \\citep{Swa03}. This is also confirmed by the HI 21 cm observations \\citep{Arn97} which show that the gas is five times more extended than the luminous polar structure, with a position-velocity diagram very similar to those observed for edge-on disks. The polar disk is very massive, since the total HI mass is about $10^{10} M_{\\odot}$, which added to the mass of stars, makes the mass in the polar disk comparable with the total mass in the central spheroid \\citep{Iod02}. The morphology of the central spheroid resembles that of a low-luminosity early-type galaxy: the surface brightness profile is described by an exponential law, with a small exponential nucleus; its integrated optical vs. NIR colors are similar to those of an intermediate age stellar population (\\citealt{Iod02}, \\citealt{Gal02}). New high resolution spectroscopy in NIR on its photometric axes suggests that this component is a nearly-exponential oblate spheroid supported by rotation \\citep{Iod06}. These new kinematic data, together with the previous studied, set important constraints on the possible formation mechanisms for NGC4650A. Because of the extended and massive polar disk, with strong $H_\\alpha$ emissions, NGC 4650A is the ideal candidate to measure the chemical abundances, metallicity and star formation rates (SFR) via spectroscopic measurements of line emissions along the major axis of the polar disk. The goal is to compare the derived values for the metallicity and SFR with those predicted by the different formation scenarios. As we shall detail in the following sections, if the polar structure forms by accretion of primordial cold gas from cosmic web filament, we expect the disk to have lower metallicities of the order of $Z\\sim 1/10 Z_{\\odot}$ \\citep{Age09} with respect to those of same-luminosity spiral disks. We shall adopt for NGC4650A a distance of about 38 Mpc based on $H_{0} = 75 \\ km \\ s^{-1} \\ Mpc^{-1}$ and an heliocentric radial velocity $V = 2880\\ km \\ s^{-1}$, which implies that 1 arcsec = 0.18 kpc. ", "conclusions": "The present study could be considered a step forward both to trace the formation history of NGC4650A and to give hints on the mechanisms at work during the building of a disk by cold accretion process. As mentioned in the Sec.\\ref{intro}, the new kinematic data obtained for the central spheroid \\citep{Iod06}, together with the previous studies, set important constraints on the possible formation mechanisms for NGC4650A. In particular, the merging scenario is ruled out because, according to simulations (e.g. \\citealt{Bou05}), a high mass ratio of the two merging galaxies is required to form a massive and extended polar disk as observed in NGC~4650A: this would convert the intruder into an elliptical-like, not rotationally supported, stellar system. This is in contrast with the high maximum rotation velocity ($\\sim 80 \\div 100$ km/s) observed in the outer regions of the central spheroid \\citep{Iod06}. Both the high baryonic mass (star plus gas) in the polar structure and its large extension cannot reconcile with a polar ring formed via the gradual disruption of a dwarf satellite galaxy (as explained by \\citealt{Iod02}). Differently, a wide polar ring and/or disk (as observed in NGC4650A) may form both around a disk or an elliptical galaxy through the tidal accretion of gas stripped from a gas-rich donor, in the case of large relative velocity and large impact parameter and for a particular orbital configuration (e.g. \\citealt{Bou03}). Therefore, the two formation scenarios which can be really envisioned in the specific case of NGC4650A are the tidal accretion and the accretion of external primordial cold gas from cosmic web filaments (\\citealt{Mac06}; \\citealt{Bro08}). To this aim we have derived the metallicity and SFR for the polar disk in NGC4650A in order to compare them with those predicted by different formation scenarios. The main results of the present work are (see also Sec. \\ref{res}): \\emph{i}) the low value of the metallicity derived for the polar disk $Z = 0.2Z_{\\odot}$, in spite of its high luminosity, $M_{B} = -19.3$ (see Fig. \\ref{conf}), \\emph{ii}) the lack of any metallicity gradient along the polar disk, which suggests that the metal enrichment is not influenced by the stellar evolution of the older central spheroid, \\emph{iii}) the metallicities expected for the present SFR at three different epochs, $1.02 Z_{\\odot} \\le Z \\le 1.4 Z_{\\odot}$, are higher than those measured from the element abundances and this is consistent with a later infall of metal-poor gas. In the following we will address how these results reconcile with the predictions by theoretical models (see Sec. \\ref{theory}) and may discriminate between the two formation mechanisms. If the polar ring/disk formed by the mass transfer from a gas-rich donor galaxy, the accreted material comes only from the outer and more metal-poor regions of the donor: is the observed metallicity for the polar component in NGC4650A consistent with the typical values for the outer disk of a bright spiral galaxy? According to \\citet{Bre09}, the metallicity of very outer regions of a bright spiral is $0.2 Z_{\\odot} \\le Z \\le 1.1 Z_{\\odot}$: the observed value for NGC4650A, $Z=0.2Z_{\\odot}$, is close to the lower limit.\\\\ The cold accretion mechanism for disk formation predicts quite low metallicity ($Z=0.1Z_{\\odot}$) (\\citealt{Dek06}, \\citealt{Ocv08}, \\citealt{Age09}): such value refers to the time just after the accretion of a misaligned material, so it can be considered as initial value for Z before the subsequent enrichment. How this may reconcile with the observed metallicity for NGC4650A? We estimated that the present SFR for the polar disk ($SFR = 0.06 M_{\\odot} yr^{-1}$) is able to increase the metallicity of about 0.2 after 1Gyr (see Sec.\\ref{res}): taking into account that the polar structure is very young, less than 1Gyr \\citep{Iod02}, an initial value of $Z=0.1Z_{\\odot}$, at the time of polar disk formation, could be consistent with the today's observed metallicity.\\\\ This evidence may put some constraints also on the time-scales of the accretion process. The issue that need to be addressed is: how could a cosmic flow form a ring/disk only in the last Gyr, without forming it before? Given that the average age of 0.8 Gyr refers to the last burst of star formation, reasonably the gas accretion along the polar direction could have started much earlier and stars formed only recently, once enough gas mass has been accumulated in the polar disk. Earlier-on, two possible mechanisms may be proposed. One process that may happen is something similar to that suggested by \\citet{Mar09} for the quenching of star formation in early-type galaxies: given that star formation takes place in gravitationally unstable gas disks, the polar structure could have been stable for a while and this could have quenched its star formation activity, until the accumulated mass gas exceeds a stability threshold and star formation resumes. Alternatively, according to the simulations by \\citet{Bro08}, the filament could have been there for several Gyrs with a relative low inclination, providing gas fuel to the star formation in the host galaxy first, about 3 Gyrs ago. Then large-scale tidal fields can let the disk/filament misalignment increase over the time till 80-90 degrees and start forming the polar structure during the last 1-2 Gyrs. This picture might be consistent not only with the relative young age of the polar disk, but also with the estimate of the last burst of star formation in the central spheroid \\citep{Iod02}. One more hint for the cold accretion scenario comes from the fact that the metallicity expected by the present SFR turns to be higher than those directly measured by the chemical abundances. As suggested by \\citet{Dal07}, both infall and outflow of gas can change a galaxy's metallicity: in the case of NGC4650A a possible explanations for this difference could be the infall of pristine gas, as suggested by (\\citealt{Fin07}; \\citealt{Ell08}). The lack of abundance gradient in the polar disk, as typically observed also in LSB galaxies, suggests that the picture of a chemical evolution from inside-out, that well reproduce the observed features in spiral disks (\\citealt{Mat89}, \\citealt{Boi99}), cannot be applied to these systems. In particular, observations for LSB galaxies are consistent with a quiescent and sporadic chemical evolution, but several explanation exist that supports such evidence \\citep{deB98}. Among them, one suggestion is that the disk is still settling in its final configuration and the star formation is triggered by external infall of gas from larger radii: during this process, gas is slowly diffusing inward, causing star formation where conditions are favorable. As a consequence, the star formation is not self-propagating and the building-up of the disk would not give rise to an abundance gradient. The similarities between LSB galaxies and the polar disk in NGC4650A, including colors \\citep{Iod02} and chemical abundances (this work), together with the very young age of the polar disk, the presence of star forming regions towards larger radii, the warping structure of outer arms (\\citealt{Gal02}, \\citealt{Iod02}), and the constant SFR along the disk (this work), suggest that the infall of metal-poor gas, through a similar process described above, may reasonably fit all these observational evidences. Given that, are there any other observational aspects that can help to disentangle in a non ambiguous way the two scenarios? One important feature which characterize NGC4650A is the high baryonic (gas plus stars) mass in the polar structure: it is about $12 \\times 10^{9} M_\\odot$, which is comparable with or even higher than the total luminous mass in the central spheroid (of about $5 \\times 10^{9} M_\\odot$). In the accretion scenario, the total amount of accreted gas by the early-type object is about $10\\%$ of the gas in the disk donor galaxy, i.e. up to $10^{9}$ $M_{\\odot}$, so one would expect the accreted baryonic mass (stellar + gas) to be a fraction of that in the pre-existing galaxy, and not viceversa, as it is observed for NGC4650A. Furthermore, looking at the field around NGC4650A, the close luminous spiral NGC4650 may be considered as possible donor galaxy \\citep{Bou03}, however the available observations on the HI content for this object show that NGC4650 is expected to be gas-poor (\\citealt{Arn97}; \\citealt{Van02}). So, where the high quantity of HI gas in NGC4650A may come from? If the polar disk forms by the cold accretion of gas from filaments there is no limit to the accreted mass. Given all the evidences shown above, we can infer that the cold accretion of gas by cosmic web filaments could well account for both the low metallicity, the lack of gradient and the high HI content in NGC4650A. An independent evidence which seems to support such scenario for the formation of polar disks comes from the discovery and study of an isolated polar disk galaxy, located in a wall between two voids \\citep{Sta09}: the large HI mass (at least comparable to the stellar mass of the central galaxy) and the general underdensity of the environment can be consistent with the cold flow accretion of gas as possible formation mechanism for this object. The present work remarks how the use of the chemical analysis can give strong constraints on the galaxy formation, in particular, it has revealed an independent check of the cold accretion scenario for the formation of polar disk galaxies. This study also confirmed that this class of object needs to be treated differently from the polar ring galaxies, where the polar structure is more metal rich (like UGC5600, see Sec. \\ref{res}) and a tidal accretion or a major merging process can reliable explain the observed properties \\citep{Bou03}. Finally, given the similarities between polar disks and late-type disk galaxies, except for the different plane with respect to the central spheroid, the two classes of systems could share similar formation processes. Therefore, the study of polar disk galaxies assumes an important role in the wide framework of disk formation and evolution, in particular, for what concern the ``rebuilding'' of disks through accretion of gas from cosmic filaments, as predicted by hierarchical models of galaxy formation \\citep{Ste02}." }, "1003/1003.1498_arXiv.txt": { "abstract": "We present a search for long-period variable (LPV) stars among giant branch stars in M15 which, at [Fe/H] $\\sim$ --2.3, is one of the most metal-poor Galactic globular clusters. We use multi-colour optical photometry from the 0.6-m Keele Thornton and 2-m Liverpool Telescopes. Variability of $\\delta$V $\\sim$ 0.15 mag is detected in K757 and K825 over unusually-long timescales of nearly a year, making them the most metal-poor LPVs found in a Galactic globular cluster. K825 is placed on the long secondary period sequence, identified for metal-rich LPVs, though no primary period is detectable. We discuss this variability in the context of dust production and stellar evolution at low metallicity, using additional spectra from the 6.5-m Magellan (Las Campanas) telescope. A lack of dust production, despite the presence of gaseous mass loss raises questions about the production of dust and the intra-cluster medium of this cluster. ", "introduction": "The onset of radial pulsation on the red giant branches is important in stellar evolution: it is one of several linked phenomena (including dust production and substantial mass loss) that control the endpoint of stellar evolution and injection of mass into the interstellar medium. While these phenomena are linked, the relative timing of their appearance is poorly understood \\citep{MvL07}. In particular, it is debatable whether pulsation can provide enough energy to assist mass loss in these evolved stars (e.g.~\\citealt{Bowen88}). Optical photometric variability in highly-evolved stars is known to be less pronounced in metal-poor systems. \\citet{FW98} considered Long-Period Variable stars (LPVs) in both the Galactic Disc and globular clusters in order to prove this dependence, though they only consider stars showing large-amplitude variability as LPVs and do not consider semi-regular variables (SRVs, which are included in the definition of LPVs for the purposes of this work). This raises the question of whether pulsation is capable, or required, to drive mass loss from metal-poor stars. We describe here the search for LPVs in one of the most metal-poor Galactic globular clusters, M15. This cluster is the only one known to harbour a dusty and/or gaseous interstellar medium and has several infrared-excessive giant stars (\\citealt{ESvL+03}; \\citealt{vLSEM06}; \\citealt{BWvL+06}), giving the strong implication that dusty stellar winds are present in the cluster and, by further implication, pulsation-driven winds. No LPVs have so-far been found in M15, though several candidates were identified by \\citet{MW75} (Table \\ref{M15CudworthTable}). These were followed-up by \\citet{Welty85}, who could not find photometric variability in three targets (K169, K288 and K709) and retained K757 and K825 as candidate variables (identifiers from \\citealt{Kustner21}). ", "conclusions": "This study presents a search for long-period variability among giants in the globular cluster M15 and has confirmed the most metal-poor variables known in our Galaxy: K757 and K825. These stars are very close to the RGB-tip and show lightcurves characteristic of pulsation, albeit with very low amplitude. Their periods place them on the `long secondary period' (LSP) sequence (\\citealt{WAA+99}; \\citealt{Wood00}), though no `primary' period has been found. The pulsational velocities are similar to those measured at the base of the wind, but it is not certain that pulsation is required to launch or drive the wind. In any case, these pulsations appear not to lead to dust production, which must be caused by a different mechanism, possibly episodically. Although we do find evidence for acceleration in the singly-ionised calcium line profiles, radiation pressure on grains cannot therefore be held responsible for driving the wind. Instread, Alfv\\'en waves may couple to a weakly-ionised gas and thus drive the wind, something which could be facilitated at lower metallicity \\citep{vLOG+10}." }, "1003/1003.4568_arXiv.txt": { "abstract": "Type Ia supernovae (SNe Ia) play an important role in the study of cosmic evolution, especially in cosmology. There are several progenitor models for SNe Ia proposed in the past years. In this paper, by considering the effect of accretion disk instability on the evolution of white dwarf (WD) binaries, we performed detailed binary evolution calculations for the WD + red-giant (RG) channel of SNe Ia, in which a carbon--oxygen WD accretes material from a RG star to increase its mass to the Chandrasekhar mass limit. According to these calculations, we mapped out the initial and final parameters for SNe Ia in the orbital period--secondary mass ($\\log P^{\\rm i}-M^{\\rm i}_2$) plane for various WD masses for this channel. We discussed the influence of the variation of the duty cycle value on the regions for producing SNe Ia. Similar to previous studies, this work also indicates that the long-period dwarf novae offer a possible ways for producing SNe Ia. Meanwhile, we find that the surviving companion stars from this channel have a low mass after SN explosion, which may provide a means for the formation of the population of single low-mass WDs ($<$0.45$\\,M_{\\odot}$). ", "introduction": "% \\label{sect:intro} Type Ia Supernovae (SNe Ia) are excellent cosmological distance indicators due to their high luminosities and remarkable uniformity. They have been applied successfully in determining cosmological parameters (e.g. $\\Omega$ and $\\Lambda$; Riess et al. 1998; Perlmutter et al. 1999). However, several key issues related to the nature of their progenitors and the physics of the explosion mechanisms are still not well understood (Hillebrandt \\& Niemeyer 2000; Podsiadlowski et al. 2008; Wang et al. 2008a), and no SN Ia progenitor system has been conclusively identified from before the explosion. It is generally believed that SNe Ia are thermonuclear explosions of carbon--oxygen white dwarfs (CO WDs) in binaries (for the review see Nomoto et al. 1997). Over the past few decades, two families of SN Ia progenitor models have been proposed, i.e. the double-degenerate (DD) and single-degenerate (SD) models. Of these two models, the SD model is widely accepted at present. It is suggested that the DD model, which involves the merger of two CO WDs (Iben \\& Tutukov 1984; Webbink 1984; Han 1998), likely leads to an accretion-induced collapse rather than to an SN Ia (Nomoto \\& Iben 1985). For the SD model, the companion is probably a main-sequence (MS) star or a slightly evolved subgiant star (WD + MS channel), a red-giant star (WD + RG channel), or an He star (WD + He star channel) (e.g. Hachisu et al. 1996; Li $\\&$ van den Heuvel 1997; Langer et al. 2000; Han $\\&$ Podsiadlowski 2004, 2006; Chen $\\&$ Li 2007, 2009; Meng et al. 2009; L\\\"{u} et al. 2009; Wang et al. 2009a,b; Wang, Li $\\&$ Han 2010). Note that, some recent observations have indirectly suggested that at least some SNe Ia can be produced by a variety of different progenitor systems (e.g. Patat et al. 2007; Voss \\& Nelemans 2008; Wang et al. 2008b; Justham et al. 2009). An explosion following the merger of two WDs would leave no remnant, while the companion star in the SD model would survive and potentially be identifiable. Tycho's supernova is a Galactic SN Ia. Ruiz-Lapuente et al. (2004) find in the remnant region that Tycho G, a star similar to the sun but with a lower gravity, moves at more than three times the mean velocity of the stars there. They argued that Tycho G could be the surviving companion of the supernova. Note that there has been no conclusive proof yet that any individual object is the surviving companion star of a SN Ia. It will be a promising method to test SN Ia progenitor models by identifying their surviving companions (e.g. Wang \\& Han 2009). The WD + RG channel is a possible ways to produce SNe Ia, and is supported by some observations. It is suggested that, RS Oph and T CrB, both recurrent novae are probable SN Ia progenitors and belong to the WD + RG channel (e.g. Belczy$\\acute{\\rm n}$ski \\& Mikolajewska 1998; Hachisu et al. 1999a, 2007; Sokoloski et al. 2006). Meanwhile, by detecting Na I absorption lines with low expansion velocities, Patat et al. (2007) suggested that the companion of the progenitor of SN 2006X may be an early RG star. Additionally, Voss \\& Nelemans (2008) studied the pre-explosion archival X-ray images at the position of the recent SN 2007on, and they consider its progenitor as a WD + RG system. Xu \\& Li (2009) recently emphasized that the mass-transfer through the Roche lobe overflow (RLOF) in the evolution of WD binaries may become unstable (at least during part of the mass-transfer lifetime). This important feature has been ignored in nearly all of the previous theoretical works on SN Ia progenitors except for King et al. (2003) \\footnote{King et al. (2003) adopted a similar method in Li \\& Wang (1998) to produce SNe Ia with long period dwarf novae in a semi-analytic approach.} and Xu \\& Li (2009), who inferred that the mass-accretion rate onto the WD during dwarf nova outbursts can be sufficiently high to allow steady nuclear burning of the accreted matter and growth of the WD mass. Following the work of Xu \\& Li (2009), Wang, Li \\& Han (2010) studied the WD binaries towards SNe Ia systematically. However, the interest of Wang, Li \\& Han (2010) mainly focused on the WD + MS channel of SNe Ia. The purpose of this paper is to study the WD + RG channel towards SNe Ia in a comprehensive manner, and to show the final parameter spaces of companions at the moment of SN Ia explosion. In Section 2, we simply describe the numerical code for the binary evolution calculations. The binary evolutionary results are shown in Section 3. Finally, a discussion is given in Section 4. ", "conclusions": "\\label{6. DISCUSSION} \\begin{figure} \\begin{center} \\includegraphics[width=6.cm,angle=270]{fig3.ps} \\caption{Regions in the initial orbital period--secondary mass plane ($\\log P^{\\rm i}$, $M^{\\rm i}_2$) for the WD + RG channel that produce SNe Ia with initial WD mass of 1.2$\\,M_{\\odot}$, but for different duty cycle values.} \\end{center} \\end{figure} Compared with most of previous investigations on the WD + RG channel of SNe Ia, the main difference is that we considered the effect of the accretion disk instability on the evolution of WD binaries. In our results, there is no WD + RG system with period as long as $\\sim$$10^{\\rm 3}$ days as indicated in Li \\& van den Heuvel (1997) and Hachisu et al. (1999a). This is because if the initial period of a WD binary is too long, the mass-transfer rate between the RG donor star and the WD will be too high and the optically thick wind will occur and take much H-rich material away from the binary system. Finally, the RG donor star has no enough material to be accumulated onto the WD. Similar to previous studies (e.g. King et al. 2003; Xu \\& Li 2009), our study also indicates that the long-period dwarf novae offer a possible ways for producing SNe Ia. In particular, there is an advantage for this work, i.e. the SN Ia explosion is always occur in a WD binary of small secondary/primary mass ratio ($<$1), and with very little of the H envelope of the secondary remaining in our investigations. This feature will greatly reduce the possibility of H contamination of the SN Ia ejecta (see also King et al. 2003), which is consistent with the defining characteristic of most SNe Ia having no detectable H (Branch et al. 1995). This also has broad implications towards the work that is currently placing very tight limits on the amount of H entrained in the SN Ia ejecta, determined through the analysis of late-time spectra (e.g. Mattila et al. 2005; Leonard 2007). However, the results in this paper depend on many uncertain input parameters, in particular for the duty cycle which is poorly known. The main uncertainties lie in the facts that it varies from one binary system to another and may evolve with the orbital periods and mass-transfer rates (e.g. Lasota 2001; Xu \\& Li 2009). This is the reason why we choose an intermediate value (0.01) rather than other extreme values (e.g. 0.1 or $10^{-3}$). Furthermore, we also did some tests for a higher or lower value of the duty cycle. In Fig. 3, we show the influence of the variation of the duty cycle value on the regions for producing SNe Ia with initial WD mass of 1.2$\\,M_{\\odot}$. We see that, for a high value (0.1), the right boundaries of the regions will be shifted to higher period, while, for a low value ($10^{-3}$), the right boundaries of the regions will be shifted to lower period. For the low value of the duty cycle, it will have a high mass-accretion rate of WDs during outbursts, so that the accreting WDs will lose too much mass via the optically thick wind, preventing them increasing masses to the Ch mass. Thus, a low value of the duty cycle will reduce the regions for producing SNe Ia. In this paper, we set the metallicity $Z=0.02$. For the WD + RG channel, varying the metallicity would have strong influence on the regions for producing SNe Ia (Fig. 2), e.g. high metallicity leads to larger radii of zero-age MS (ZAMS) stars, then the left boundaries of the regions will be shifted to longer period. Meanwhile, stars with high metallicity evolve in a way similar to those with low metallicity but less mass (Umeda et al. 1999; Chen \\& Tout 2007; Meng et al. 2009). Thus, for the WD binary systems with particular orbital periods, the companion mass increases with metallicity. At present, the existence of a population of single low-mass ($<$0.45$\\,M_{\\odot}$) WDs (LMWDs) is supported by some observations (e.g. Maxted et al. 2000; Kilic et al. 2007). The formation of single LMWDs is still unclear. It is suggested that single LMWDs could result from single old metal-rich stars which experiences severe mass loss prior to the He flash (Kalirai et al. 2007; Kilic et al. 2007). However, the study of initial-final mass relation for stars by Han et al. (1994) implies that only LMWDs with masses $\\ga$0.4$\\,M_{\\odot}$ might be produced through such a single-star channel, even at high metallicity (Meng, Chen \\& Han 2008). Thus, it would be difficult to conclude that single stars can produce the LMWDs. The companion in the WD + RG channel would survive and evolve to a WD finally. In Fig. 2, we see that the companion stars have a low mass at the moment of SN explosion. The companion stars will be stripped of some mass due to the impact of SN ejecta. Marietta et al. (2000) presented several high-resolution two-dimensional numerical simulations of the impact of SN Ia explosion with companions. They find that a RG donor star will lose almost its entire envelope (96\\%$-$98\\%) owing to the impact of the SN Ia explosion and leave only the core of the star. Thus, the surviving companion stars from this channel will have a relatively low mass after SN explosion and evolve to a WD finally, which provides a possible pathway for the formation of the population of single LMWDs ($<$0.45$\\,M_{\\odot}$). Meanwhile, we also suggest that the observed, apparently single LMWDs may provide evidence that at least some SN Ia explosions have occurred with non-degenerate donor stars (such as RG donor stars). The CO WD + RG systems can be formed by binary evolution. Wang, Li \\& Han (2010) find that there is one channel which can form CO WD + RG systems and then produce SNe Ia. In the detailed binary evolution procedure, the primordial primary first fills its Roche lobe at the thermal pulsing asymptotic giant branch stage. A common envelope is then easily formed owing to dynamically unstable mass-transfer during the RLOF stage. After the common envelope ejection, the primordial primary becomes a CO WD, then a CO WD + MS system is produced. The MS companion star continues to evolve until the RG stage, i.e. a CO WD + RG system is formed. For the CO WD + RG systems, SN Ia explosions occur for the ranges $M_{\\rm 1,i}\\sim5.0$$-$$6.5\\,M_\\odot$, $M_{\\rm 2,i}\\sim1.0$$-$$1.5\\,M_\\odot$, and $P^{\\rm i}\\ga 1500$\\,days, where $M_{\\rm 1,i}$, $M_{\\rm 2,i}$ and $P^{\\rm i}$ are the initial mass of the primary and the secondary at ZAMS, and the initial orbital period of a binary system. The WD + RG channel has a long delay time from the star formation to SN explosion due to the RG donor star with low initial masses ($\\la$1.5$\\,M_{\\odot}$). Thus, this channel can contribute to the old population of SNe Ia implied by recent observations (Mannucci et al. 2006; Totani et al. 2008; Schawinski 2009). The old population of SNe Ia may have an effect on models of galactic chemical evolution, since they would return large amounts of iron to the interstellar medium much later than previously thought. It may also have an impact on cosmology, as they are used as cosmological distance indicators." }, "1003/1003.0457_arXiv.txt": { "abstract": "We present six new transits of the exoplanet OGLE-TR-111b observed with the Magellan Telescopes in Chile between April 2008 and March 2009. We combine these new transits with five previously published transit epochs for this planet between 2005 and 2006 to extend the analysis of transit timing variations reported for this system. We derive a new planetary radius value of $1.019 \\pm 0.026~R_{J}$, which is intermediate to the previously reported radii of $1.067\\pm0.054~R_J$ \\citep{Winn2007} and $0.922\\pm0.057~R_J$ \\citep{Diaz2008}. We also examine the transit timing variation and duration change claims of \\citet{Diaz2008}. Our analysis of all eleven transit epochs does not reveal any points with deviations larger than $2\\sigma$, and most points are well within $1\\sigma$. Although the transit duration nominally decreases over the four year span of the data, systematic errors in the photometry can account for this result. Therefore, there is no compelling evidence for either a timing or a duration variation in this system. Numerical integrations place an upper limit of about $1~M_{\\oplus}$ on the mass of a potential second planet in a 2:1 mean-motion resonance with OGLE-TR-111b. ", "introduction": "Transiting exoplanets provide a wealth of information for studies of the physical parameters of planets and their environments. For example, the combination of several accurately timed transits of a known transiting exoplanet can be used not only to improve estimates of the planetary radius and orbital parameters of the star-planet system, but also to detect additional objects. Detecting potential variations of parameters such as the inclination and duration of the transits would indicate a precesing planetary orbit, potentially caused by another planet \\citep{MiraldaEscude2002}. We can also use transit timing to search for additional planets or moons, as discussed in several recent theory papers \\citep{Holman2005, Agol2005, Heyl2007, Ford2007, Simon2007, Kipping2009a, Kipping2009b}. The idea is that the presence of additional objects will perturb the orbit of the transiting planet, producing transit timing variations (TTVs) or transit duration variations (TDVs). Those TTVs and TDVs can be detected by monitoring transits over many orbital periods. The absence of such variations can be also used to place limits on the mass and orbital parameters of additional objects in those planetary systems and to gain insight into the systems' architectures. Recent observations show hints of timing variations for some transiting planets, but no definitive detection of additional planets or satellites has been reported using this technique. The most interesting results so far are (1) the absence of TTVs in several systems, which do not host planets more massive than several Earth masses in low-order resonant orbits (see a summary of constraints that can be placed in Table~\\ref{table:ttvlit}); (2) the tentative detection of TDVs in GJ436, roughly 3 minutes per year \\citep{Coughlin2008}, a trend consistent with the presence of a low-mass companion ($<12~M_{\\oplus}$) in a close exterior but non-resonant orbit; this result is consistent with the $8~M_{\\oplus}$ limit placed by transit timing \\citep{Bean2008}; and (3) the preliminary detection of TTVs with a maximum residual of $156 \\pm 48$ sec ($3.3\\sigma$) over a period of 2 years reported by \\citet{Diaz2008} for OGLE-TR-111b, the subject of this paper. OGLE-TR-111b is a $0.5 M_J$ hot Jupiter orbiting its host star, a faint ($I=15.5$) K dwarf, every 4.01 days. This object was first announced as a transiting planet candidate by \\citet{Udalski2002}, and was confirmed to have planetary mass by \\citet{Pont2004}. The physical parameters of the planet were refined over the next two years, with several new radial velocity measurements \\citep{Gallardo2005, Silva2006, Santos2006}. The first high precision transit photometry was provided by \\citet{Winn2007}, with two \\emph{I}-band transits of the planet on 2006 Feb 21 and Mar 5. Shortly after, \\citet{Minniti2007} published a \\emph{V}-band transit from 2005 April 9 and noted that the midtime occurred 5 minutes earlier than expected from the ephemeris in \\citet{Winn2007}, although with only three epochs they could draw no firm conclusions. A follow-up paper by \\citet{Diaz2008} reported two consecutive \\emph{I}-band transits of OGLE-TR-111b on 2006 Dec 19 and 23. Combining all five epochs, they concluded that the previously claimed TTVs were real, with the residuals spanning $-156\\pm48$ to $+98\\pm39$ seconds. Among other scenarios, they noted that if OGLE-TR-111b were in an eccentric orbit with $e\\sim0.3$, the observed TTVs would be consistent with the presence of an Earth-mass planet near an exterior 4:1 resonant orbit. Additionally, \\citet{Diaz2008} noted two parameters with marginally discrepant values across the five transits (see Table~\\ref{table:litparams}). Compared to the results from \\citet{Winn2007}, the \\citet{Diaz2008} values for the planetary radius disagreed at the 10\\% level, or $1.3\\sigma$, and the total transit duration differed by $1.6\\sigma$. The radius ratio discrepancy was suggested to be the result of the parameters chosen for the image subtraction photometry, which focused on precise timing rather than on an accurate transit depth determination. The duration variation, if real, could be due to a perturber decreasing the orbital inclination, which would offer another way of determining the properties of the third body in the system suggested by their TTVs. Here we present six new transits observed during 2008 and 2009, which double the number of high-quality transit light curves available for OGLE-TR-111b. In \\S~\\ref{section:obs} we describe the collection and analysis of the new data. In \\S~\\ref{section:fitting} we describe the transit model fitting, and discuss additional sources of error not included in the formal fit. In \\S~\\ref{section:results} we combine the six new transits with the five previously published observations and provide a new analysis of parameter variation in the OGLE-TR-111 system. In \\S~\\ref{section:conclusions} we discuss the implications of our results. ", "conclusions": "\\label{section:conclusions} We have tested the previously claimed presence of TTVs and TDVs in the OGLE-TR-111 system by adding six new transit epochs, observed between 2008 and 2009, to the five previously published results by \\citet{Winn2007,Minniti2007,Diaz2008}. This new analysis not only doubles the number of available data points, but also extends the TTV baseline from two to four years. In addition, combining the six new transits data allows us to provide a new, more precise value of the radius of this planet. We find a new radius for the planet of $1.019 \\pm 0.026~R_{J}$, which is intermediate to the previously reported radii by \\citet{Winn2007} and \\citet{Diaz2008}, and is more precise. We find a slight variation over time of the duration of the transits of OGLE-TR-111b, as well as variations of other parameters, such as the inclination and semimajor axis of the orbit. Those variations could, in principle, be attributed to perturbations of the orbit of OGLE-TR-111b produced through interaction with additional planet(s) in the system, but we demonstrate that the variations can be instead explained by systematic errors in the data, and therefore should not be attributed to other planets. We have also computed the transit midtimes of our new transits with formal precisions of 20-40 seconds, and more accurate precisions of 35-50 seconds for the full transits (and almost 2 minutes for the half-transit) once systematic errors are considered. The errors on the literature transits similarly increased when the photometry is refit using the same method to account for systematics, to 60-110 seconds depending on the light curve. A longer time baseline and more precise timing data is still necessary to test further for the presence of other planets in the OGLE-TR-111 system, especially in potentially stable non-resonant orbits, but with the present results we conclude that OGLE-TR-111 belongs in the category of systems summarized in Table~\\ref{table:ttvlit} for which there is no sign of additional planets more massive than a few $M_{\\oplus}$ in low-order resonant orbits, including a limit of $1~M_{\\oplus}$ near the 2:1 resonances. The presence of massive (Earth-like) moons around OGLE-TR-111b is still possible, but to detect those we would require timing precision of a few seconds or better, beyond the current capability of ground-based instrumentation for this system." }, "1003/1003.2980_arXiv.txt": { "abstract": "{The quiet-Sun X-ray emission is important for deducing coronal heating mechanisms, but it has not been studied in detail since the emph{Orbiting Solar Observatory} (\\emph{OSO}) spacecraft era. Bragg crystal spectrometer X-ray observations have generally concentrated on flares and active regions. The high sensitivity of the RESIK (REntgenovsky Spectrometer s Izognutymi Kristalami) instrument on the \\emph{CORONAS-F} solar mission has enabled the X-ray emission from the quiet corona to be studied in a systematic way for the first time. } {Our aim is to deduce the physical conditions of the non-flaring corona from RESIK line intensities in several spectral ranges using both isothermal and multithermal assumptions.} {We selected and analyzed spectra in 312 quiet-Sun intervals in January and February 2003, sorting them into 5 groups according to activity level. For each group, the fluxes in selected spectral bands have been used to calculate values parameters for the best-fit that lead to a intensities characteristic of each group. We used both isothermal and multitemperature assumptions, the latter described by differential emission measure (DEM) distributions. RESIK spectra cover the wavelength range ($3.3-6.1$~\\AA). This includes emission lines of highly ionized Si, S, Cl, Ar, and K, which are suitable for evaluating temperature and emission measure, were used. } {The RESIK spectra during these intervals of very low solar activity for the first time provide information on the temperature structure of the quiet corona. Although most of the emission seems to arise from plasma with a temperature between 2~MK and 3~MK, there is also evidence of a hotter plasma ($T \\sim 10$~MK) with an emission measure 3 orders smaller than the cooler component. Neither coronal nor photospheric element abundances appear to describe the observed spectra satisfactorily.} {} ", "introduction": "The RESIK X-ray spectrometer on the Russian {\\it CORONAS-F} solar orbiting mission (circular polar orbit: altitude 550~km, period 96 minutes) obtained numerous high-resolution flare and active-region spectra in the $3.3-6.1$~\\AA\\ range over the period August~2001~--~May~2003. The RESIK instrument (Sylwester et al., 2005) was a bent crystal spectrometer with four spectral channels in which solar X-ray emission was diffracted by crystal wafers made of silicon (Si 111, $2d = 6.27$~\\AA) and quartz (Qu $10\\bar 10$, $2d = 8.51$~\\AA). Although most previous spacecraft crystal spectrometers suffered from the strong instrumental backgrounds caused by fluorescence of the crystal material, RESIK had a system of pulse-height analyzers enabling primary solar photons to be distinguished from secondary photons produced by fluorescence. The background could thus be practically eliminated for much of the period 2003 January~--~March when the non flaring solar X-ray activity was often below C1 class as measured by the {\\it GOES} 1~--~8~\\AA\\ sensor. The sensitivity of RESIK was maximized by not having a collimator placed in front of the crystals. This, like the equivalent Bragg Crystal Spectrometer on the {\\it Yohkoh} spacecraft (operational 1991--2001), produced some spectral confusion when two simultaneous bright sources were present on the Sun, but in practice this rarely occurred. The low-activity corona gives rise to X-ray line profiles in RESIK spectra having large line widths through spatial broadening. Several analyses have been done on spectra of solar flares from RESIK (Sylwester et al., 2006; Chifor et al., 2007; Sylwester, Sylwester and Phillips, 2008), but here we report on spectra obtained during a period of sustained low solar activity, which have high statistical quality because of the relatively high sensitivity of RESIK. The observed spectral shapes including continua are available for analysis of the temperature structure of the emitting regions, specifically the differential emission measure as a function of electron temperature $T$; and from this, the absolute abundances of the elements giving rise to the spectral lines can be assessed. \\begin{figure}[t] \\vspace{3mm} \\hspace{10mm} \\centering \\includegraphics[width=5.5cm]{12907fg1.ps} \\vspace{3mm} \\caption{Histogram showing the respective \\emph{GOES} activity for selected 312 time intervals (January-February 2003).} \\label{FigGam} \\end{figure} Quiet-Sun X-ray spectra have not received nearly as much attention in the past as those from flares, but some properties of the quiet Sun have been widely studied using its ultraviolet emission, which has been measured by the following experiments: Orbiting Solar Observatory \\emph{OSO} series (1962-1975; see for instance Dupree and Reeves, 1971; Dupree et al., 1973), \\emph{Aerobee} rocket spectrometer (1969, see Malinovsky and Heroux, 1973), 9 months of the \\emph{Skylab} mission (May 1973-February 1974, see Vernazza and Reeves, 1978), the series of 9 Solar EUV Rocket Telescope and Spectrograph (SERTS) flights (the first in 1989, see Brosius et al., 1996, 1998), the \\emph{SOHO} mission (starting in December 1995, CDS and SUMER data, see Warren, Mariska and Lean, 1998). The data obtained have been used as well to construct a reference atlas of quiet-Sun ultraviolet radiation (Curdt, Landi and Feldman, 2004) from which differential emission measure (DEM) distributions can be inferred. Brosius et al. (1996) calculated DEM distributions for quiet-Sun conditions during two observing periods near the maximum of Cycle~21 (1991 May 7 and 1993 August 17) based on SERTS data. Their DEM solutions included a power law in the temperature range $6.3\\times 10^4$~K--$5.0\\times 10^5$~K with hot plasma as shown by a localized maximum at 5~MK. A DEM with similar distribution was obtained by Kretzschmar, Lilensten, and Aboudarham (2004) based on \\emph{SOHO} SUMER data, having a power-law shape in the range $2.0\\times 10^4$~K--$2.0\\times 10^5$~K with a high-temperature bump at 1.1~MK. Ralchenko, Feldman, and Doschek (2007) have studied quiet-corona spectra as observed by SUMER during the 2000 June 13--19 period, deducing that the observed line intensities can be satisfactorily described by a model with two Maxwellian electron distributions, a first population corresponding to an isothermal temperature of $\\sim 1.3\\times 10^6$~K, and a second, smaller population ($\\sim 5\\%$) of hot ({300--400}~eV) electrons accounting for the intensities of highly charged Ar and Ca ion lines observed by SUMER. \\begin{figure}[t] \\centering \\includegraphics[width=8cm]{12907fg2.ps} \\vspace{-0mm} \\caption{ Composite of four full-disk EIT/SOHO images obtained in the wavelength bands 304~\\AA, $171$~\\AA, $195$~\\AA, and $284$~\\AA, on 2003~February~24 around 19:00~UT. This was the time of lowest activity in the period of the RESIK quiet-Sun spectra analyzed in this work. } \\label{FigVibStab} \\end{figure} The DEM analysis was also performed using \\emph{SOHO} Coronal Diagnostic Spectrometer (CDS) data for the internetwork, network, and bright network regions of the quiet Sun by O'Shea et al.~(2000). They find that the DEM distributions differ in each region over the temperature range 0.25~--~1~MK. Recently, Young et al. (2007) have published a quiet-Sun extreme ultraviolet (EUV) spectrum in the ranges {170--211}~\\AA\\ and {246--292}~\\AA, as observed on 2006 December 23 by the \\emph{Hinode} EUV Imaging Spectrometer (EIS). \\begin{table}[h] \\caption{Wavelength intervals used for DEM studies} % \\label{table:1} % \\centering % \\begin{tabular}{c c c} % \\hline\\hline % No & Range [\\AA] & Dominant contributor \\\\ % \\hline % \\vspace{-3mm} & & \\\\ 1 & 3.40~-~3.62 & cont. + K~{\\sc xviii }$2p$+sat. \\\\ 2 & 3.62~-~3.80 & cont. + Ar~{\\sc xviii }$2p$, S~{\\sc xvi }$4p, 5p$ \\\\ 3 & 3.85~-~4.10 & cont. + Ar~{\\sc xvii }$2p$+sat, S~{\\sc xv }$4p$ \\\\ 4 & 4.10~-~4.25 & cont. + S~{\\sc xv }sat. \\\\ 5 & 4.35~-~4.43 & cont. + S~{\\sc xv }$3p$+sat \\\\ 6 & 4.43~-~4.52 & cont. + Cl~{\\sc xvi }$2p$+sat \\\\ 7 & 4.68~-~4.74 & cont. + S~{\\sc xvi }$2p$, + Si~{\\sc xiv }$8p$ \\\\ 8 & 4.74~-~4.81 & cont. + Si~{\\sc xiv }$6p$ \\\\ 9 & 5.00~-~5.13 & S~{\\sc xv }$2p$+sat + cont. \\\\ 10 & 5.26~-~5.35 & Si~{\\sc xiii }$5p$+sat. + cont. \\\\ 11 & 5.36~-~5.50 & Si~{\\sc xiii }$4p$+sat. + cont. \\\\ 12 & 5.50~-~5.63 & Si~{\\sc xii } sat. + cont. \\\\ 13 & 5.64~-~5.72 & Si~{\\sc xiii }$3p$ + cont. \\\\ 14 & 5.73~-~5.86 & Si~{\\sc xii } sat. + cont. \\\\ 15 & 5.90~-~6.00 & continuum \\\\ \\hline % \\end{tabular} \\vspace{-5mm} \\end{table} The quiet-Sun hard X-ray emission observed by \\emph{RHESSI} has been examined by Hannah et al. (2007) using a fan-beam modulation technique during seven periods of off-pointing of the \\emph{RHESSI} spacecraft between 2005 June and 2006 October. They established new upper limits on the 3--200~keV X-ray emission for when the \\emph{GOES} level of activity was below A1 class, updating much earlier measurements (Peterson et al., 1966). Schmelz et al. (2009) have investigated the emission of a nonflaring active region based on the ten filters of the \\emph{Hinode} X-ray Telescope data using two independent algorithms to reconstruct the differential emission measure distribution. In addition to the typical low-temperature emission measure ($T<5$~MK), they find a very hot component ($\\sim 30$~MK) with small emission measure. These findings have recently been modified by Schmelz 2009, priv. comm.: the temperature of the hotter component is now much lower, around 10~MK. Reale et al. (2009) have investigated the \\emph{Hinode} XRT data averaged over one hour during a nonflaring period (2006 November~12), finding a hotter component (temperature $\\sim 6.3$~MK) corresponding to a nonflaring active region. These observational results are supported by the theoretical work of Klimchuk et al. (2008) who use hydrodynamic simulations of nanoflares to predict a small amount of hot plasma in addition to the dominant 2--3~MK plasma component. The RESIK spectra recorded during solar minimum can bridge the gap between the results obtained from ultraviolet and X-ray images and spectra and the models of coronal plasma heating. In earlier work, we analyzed RESIK spectra to determine the conditions of flaring plasmas, but here we apply the same techniques of analyzing emission during low-level periods to deduce the properties of the quiet-Sun corona. ", "conclusions": "\\begin{figure} \\centering \\includegraphics[width=8 cm]{12907fg7.ps} \\vspace{2mm} \\caption{ The histogram representation of differential emission measure (DEM) distributions as obtained from RESIK spectra for individual classes of solar activity. Different colors correspond to the following classes: black (dark) represents the B4-B5 class and yellow (lightest) A9-B1 class. } \\label{FigVibStab} \\end{figure} We analyzed 312 individual RESIK spectra {grouped} into 5 different activity classes based on \\emph{GOES} levels of activity recorded at the time of the spectra collection. {All selected spectra were taken during the nonflaring, low activity conditions prevailing in the corona, though some weak active-region emission was always present. The spectral observations (all with a five-minute integration time) cover the period between 2003 January~1 and March~14.} The analysis assumed both isothermal and multithermal distributions, the latter leading to determination of DEM distributions for each activity class. The DEMs were obtained with the Withbroe-Sylwester Bayesian iterative, maximum likelihood procedure. The fluxes integrated over 15 wavelength bands covering the range 3.3~-~6.1~\\AA\\ were used as an input data for the deconvolution. We found it necessary to allow for the emitting plasma composition to be nonstandard, i.e. neither coronal nor photospheric, by varying the abundances of those elements making important contributions to the line emission in the spectra. This analysis has never been done before except for a few flares (to be published, Montreal COSPAR). The main results of the present study follow. \\begin{enumerate} \\item It is necessary to use the multithermal approach in analyzing the spectra. The results obtained in the isothermal approximation provide different values of $T$ and $EM$ parameters depending on the ratio considered. However, the values of so-called thermodynamic measure that is directly related to the total thermal energy plasma content are very close, because they are surprisingly independent of the particular ratio used for its determination. \\item It appears necessary to allow for the plasma composition differences when using the multitemperature approach in the analysis. Neither coronal nor photospheric composition models are able to describe the observed spectra satisfactorily. This result may also bias the outcome of any isothermal analysis performed with an isothermal approach. \\item The two-temperature character of the DEM shape determined here for non-flaring plasmas has also been obtained for flares (Sylwester et al. 2008). The presence of the higher-$T$ component, with $T$ somewhat below 10~MK, is physically important. The emission measure associated with this hotter plasma is $\\sim 3$~orders of magnitude smaller than in the generally accepted $T \\sim 2 - 3$~MK component. This higher-$T$ component is required because it is impossible to reproduce the observed spectra without it. \\end{enumerate} It is worth noting that the presence of a hotter component in active region emission has been recently suggested by Reale et al. (2009) from their analysis of \\emph{Hinode} XRT images of an active region and by Schmelz et al. (2009), although the very high temperature of the hotter component discussed by Schmelz et al. (2009) has been reduced to a value near what is obtained in this work (Schmelz 2009, priv. comm.)." }, "1003/1003.2452_arXiv.txt": { "abstract": "The Fermi Gamma-Ray Space Telescope has more than doubled the number of Gamma-Ray Bursts (GRBs) detected above $100\\;$MeV within its first year of operation. Thanks to the very wide energy range covered by Fermi's Gamma-ray Burst Monitor (GBM; 8$\\;$keV to 40$\\;$MeV) and Large Area Telescope (LAT; 25$\\;$MeV to $>\\,$300$\\;$GeV) it has measured the prompt GRB emission spectrum over an unprecedentedly large energy range (from $\\sim 8\\;$keV to $\\sim 30\\;$GeV). Here I briefly outline some highlights from Fermi GRB observations during its first $\\sim 1.5\\;$yr of operation, focusing on the prompt emission phase. Interesting new observations are discussed along with some of their possible implications, including: (i) What can we learn from the Fermi-LAT GRB detection rate, (ii) A limit on the variation of the speed of light with photon energy (for the first time beyond the Planck scale for a linear energy dependence from direct time of arrival measurements), (iii) Lower-limits on the bulk Lorentz factor of the GRB outflow (of $\\sim 1000$ for the brightest Fermi LAT GRBs), (iv) The detection (or in other cases, lack thereof) of a distinct spectral component at high (and sometimes also at low) energies, and possible implications for the prompt GRB emission mechanism, (v) The later onset (and longer duration) of the high-energy emission ($>\\,$100$\\;$MeV), compared to the low-energy ($\\lesssim\\,$1$\\;$MeV) emission, that is seen in most Fermi-LAT GRBs. ", "introduction": "The Energetic Gamma-Ray Experiment Telescope (EGRET) on-board the Compton Gamma-Ray Observatory (CGRO; 1991$-$2000) was the first to detect high-energy emission from GRBs. EGRET detected only five GRBs with its Spark Chambers (20$\\,$MeV to 30$\\,$GeV) and a few GRBs with its Total Absorption Shower Counter (TASC; $1-200\\;$MeV). Nevertheless, these events already showed diversity. The most prominent examples are GRB~940217, with high-energy emission lasting up to $\\sim 1.5\\;$hr after the GRB including an $18\\;$GeV photon after $\\sim 1.3\\;$hr,~\\cite{Hurley93} and GRB~941017 which had a distinct high-energy spectral component~\\cite{Gonzalez03} detected up to $\\sim 200\\;$MeV with $\\nu F_\\nu \\propto \\nu$. This high-energy spectral component had $\\sim 3$ times more energy and lasted longer ($\\sim 200\\;$s) than the low-energy (hard X-ray to soft gamma-ray) spectral component (which lasted several tens of seconds), and may be naturally explained as inverse-Compton emission from the forward-reverse shock system that is formed as the ultra-relativistic GRB outflow is decelerated by the external medium~\\cite{GrGu03,PW04}. Nevertheless, better data are needed in order to determine the origin of such high-energy spectral components more conclusively. The Italian experiment Astro-rivelatore Gamma a Immagini LEggero (AGILE; launched in 2007) has detected GRB~080514B at energies up to $\\sim 300\\;$MeV, and the high-energy emission lasted longer ($>\\,$13$\\;$s) than the low-energy emission ($\\sim 7\\;$s)~\\cite{Giuliani08}. Below are some highlights of Fermi GRB observations so far and what they have taught us. ", "conclusions": "" }, "1003/1003.5102_arXiv.txt": { "abstract": "{} {The earliest phases of massive star formation are found in cold and dense infrared dark clouds (IRDCs). Since the detection method of IRDCs is very sensitive to the local properties of the background emission, we present here an alternative method to search for high column density in the Galactic plane by using infrared extinction maps. Using this method we find clouds between 1 and 5\\,kpc, of which many were missed by previous surveys. By studying the physical conditions of a subsample of these clouds, we aim at a better understanding of the initial conditions of massive star formation. } {We have made extinction maps of the Galactic plane based on the $3.6-4.5$~$\\mu$m color excess between the two shortest wavelength Spitzer IRAC bands, reaching to visual extinctions of $\\sim$100\\,mag and column densities of $9\\times10^{22}~\\mathrm{cm^{-2}}$. From this we compiled a new sample of cold and compact high extinction clouds. % We used the MAMBO array at the IRAM 30m telescope to study the morphology, masses and densities of the clouds and the dense clumps within them. The latter were followed up by pointed ammonia observations with the 100m Effelsberg telescope, to determine rotational temperatures and kinematic distances.} { Extinction maps of the Galactic plane trace large scale structures such as the spiral arms. The extinction method probes lower column densities, $N_{\\mathrm{H_2}} \\sim 4\\times 10^{22}\\,\\mathrm{cm^{-2}}$, than the 1.2 mm continuum, which reaches up to $N_{\\mathrm{H_2}} \\sim 3\\times 10^{23}\\,\\mathrm{cm^{-2}}$ but is less sensitive to large scale structures. The 1.2 mm emission maps reveal that the high extinction clouds contain extended cold dust emission, from filamentary structures to still diffuse clouds. Most of the clouds are dark in 24\\,$\\mu$m, but several show already signs of star formation via maser emission or bright infrared sources, suggesting that the high extinction clouds contain a variety of evolutionary stages. The observations suggest an evolutionary scheme from dark, cold and diffuse clouds, to clouds with a stronger 1.2\\,mm peak and to finally clouds with many strong 1.2 mm peaks, which are also warmer, more turbulent and already have some star formation signposts.} {} ", "introduction": "Massive stars play a fundamental role in the evolution of galaxies through their strong UV radiation, stellar winds and supernovae explosions, which contribute to the chemical enrichment of the interstellar medium. Massive stars are rare, hence usually found at large distances. They form very rapidly while still deeply embedded in their natal molecular clouds. These characteristics impose several observational obstacles, like the necessity of high resolution and sensitivity in an un-absorbed frequency range, to study their formation. Currently, the earliest stage of massive star formation is thought to take place in the very dense clumps found in Infrared Dark Clouds (IRDCs). The properties of IRDCs are shown by \\citet{carey:1998} to be dense ($n_{\\mathrm{gas}}>10^5\\,\\mathrm{cm^{-3}}$) and cool ($T<20\\,\\mathrm{K}$) aggregations of gas and dust in the Galaxy. They contain clumps with typical masses of $\\geq100\\,\\mathrm{M_\\odot}$ \\citep{rathborne:2006,pillai:2006,simon:2006b}. From IRDC clumps to the next stage, the high-mass protostellar objects \\citep[HMPOs,][]{beuther:2002,sridharan:2002}, the temperatures increase ($30\\,\\mathrm{K}<\\mathrm{T}<60\\,\\mathrm{K}$), the line widths increase, densities and masses rise \\citep{motte:2007}. HMPOs are usually found prior to the formation of ultra compact H{\\sc ii} (UCH{\\sc ii}) regions, before the newly formed star begins to ionize its surrounding medium. \\citet{motte:2007} demonstrated the difficulty of finding massive objects in an early evolutionary phase: in their survey of 3 deg$^2$ in Cygnus X, they found little evidence for dense clumps without any trace of star formation, however dense clumps with already ongoing star formation were found to be abundantly present. Based on these results, the statistical life time of the high-mass protostars and prestellar cores was estimated $\\leq3\\times10^4$\\,yr \\citep{motte:2007}, which is much shorter than what is found in nearby low-mass star-forming regions: $\\sim$$3\\times10^5$\\,yr \\citep{kirk:2005}. It is these very early stages, which provide important clues to construct a theoretical model of massive star formation, since the initial fragmentation of the gas and dust in a clump will be different in the case of monolithical collapse \\citep{krumholz:2009} compared to the competitive accretion model \\citep{bonnell:2001,clark:2008}. Massive stars generally form in clusters \\citep{lada:2003}, of which the precursors are massive clumps or the so-called precluster forming clumps, hereafter just clumps, of a $\\sim$1\\,pc size. For many massive star-forming regions, we do not yet have the capacity to resolve the clumps into prestellar cores and study the fragmentation \\citep{beuther:2007a,rathborne:2007,rathborne:2008,zhang:2009,swift:2009}. In this paper, we report on the physical parameters of the clumps, such as their morphology, density and temperature. Based on this, we hypothesize on an evolutionary sequence of cluster formation. Our understanding of the clumps increased considerably with the discovery of IRDCs. IRDCs are detected by a local absence of infrared (IR) emission against the diffuse mid IR emission of the Galactic plane \\citep{perault:1996,egan:1998} and are observed numerously throughout the Milky Way \\citep{simon:2006b}. At the typical low temperatures of IRDCs ($\\sim20$\\,K), the dust emission peaks in the far-infrared and is optically thin at mm/submm wavelengths. For a majority of IRDCs the mm dust emission coincides with the morphology of the IR absorption \\citep{rathborne:2006,pillai:2006}. Many clumps in IRDCs show signs of star formation via infrared emission at 24\\,$\\mu$m, or SiO emission from shocks driven by outflows \\citep{motte:2007,beuther:2007b,chambers:2009}. Observations tell that clumps in IRDCs span a very wide range of masses, indicating that not all will form clusters with massive stars \\citep{rathborne:2006,pillai:2006}. The detection method of IRDCs is very sensitive to the local properties of the background emission. Also, not all massive dust condensations will be infrared dark if there is enough foreground emission. Hence, to find the high mass end of molecular clouds in an unbiased fashion, new, complementary approaches are needed. We have developed such a new method, well known from studies of low mass star-forming regions, to target more efficiently the most massive clouds: \\citet{lada:1994} pioneered the method of measuring high amounts of extinction through stellar color excess in the infrared. Applied to the 2\\,$\\mu$m data of the 2MASS survey, they covered the range up to 40\\,magnitudes in visual extinction, $A_V$. However, this is not sufficient to probe the dense birthplaces of massive stars. Here the results of the Spitzer Space Telescope GLIMPSE survey \\citep{benjamin:2003} came to help: by applying the extinction curve of \\citet{indebetouw:2005} we have extended the color excess method to reach up to peaks in $A_V$ of $\\sim$100 magnitudes (or column densities $N_\\mathrm{H_2}$ of $9\\times10^{22}$\\,cm$^{-2}$), thus entering the realm where massive star formation becomes possible. The extinction method, however, is limited by the number of available background stars, and will therefore detect mainly nearby clouds (discussed in Sect.~\\ref{sec:hecmap}). In the meanwhile, complementary, unbiased dust continuum surveys were carried out: the ATLASGAL survey of the complete inner Galactic plane at 870\\,$\\mu$m by \\citet{schuller:2009} and the 1.1\\,mm BOLOCAM survey \\citep{rosolowsky:2009} of the Galactic plane accessible from the northern hemisphere. We selected the more compact and high extinction (mean $A_{\\mathrm{V}}>20\\,\\mathrm{mag}$ or $N_\\mathrm{H_2}>2\\times10^{22}$\\,cm$^{-2}$) sources from large scale extinction maps of the inner Galactic plane ($-60\\degr < l < 60\\degr$, $0.9\\degr -42$\\degr. It is thus missing some of the most interesting star formation regions of the Galaxy, like the Carina nebula region. In this context, we undertook a multi-band NIR AO campaign on the main Carina region clusters with the Multi-Conjugate Adaptive Optics (MCAO) Demonstrator \\citep[MAD, ][]{MBD07}. Beside deep NIR photometry of the individual clusters, our survey was designed to provide us with high-resolution imaging of the close environment of a sample of 60 O/WR massive stars in the Carina region. Unfortunately, the bad weather at the end of the second MAD demonstration run in January 2008 prevented the completion of the project. Valuable $H$ and $K_\\mathrm{S}$ photometry of the sole Tr~14 cluster could be obtained. The 2\\arcmin\\ field of view (fov) still provides us with high-quality information of the surrounding of $\\sim$30 early-type stars with masses above 10~\\msol. It also constitutes the most extended AO mosaic ever acquired. \\\\ \\begin{table} \\centering \\caption{Field centering (F.C.) for on-object (Tr~14) and on-sky observations and coordinates of the natural guide stars (NGSs).} \\label{tab: fov} \\begin{tabular}{cccc} \\hline \\hline & RA & DEC & V mag \\\\ \\hline Tr~14 F.C.& 10:43:55.00 & $-$59:33:03.0 & \\dots \\\\ Sky F.C. & 10:43:06.53 & $-$59:37:46.1 & \\dots \\\\ \\\\ NGS1 & 10:43:59.92 & $-$59:32:25.4 & 9.3 \\\\ NGS2 & 10:43:57.69 & $-$59:33:39.2 & 11.2 \\\\ NGS3 & 10:43:48.82 & $-$59:33:24.8 & 10.7 \\\\ \\hline \\end{tabular} \\end{table} \\begin{table} \\centering \\caption{Log of the MAD observations of Tr~14. } \\label{tab: diary} \\begin{tabular}{ccccccc} \\hline \\hline DP & RA & DEC & DIT & NDIT & $N_\\mathrm{IMA}$ & Tot. \\\\ \\hline \\\\ \\multicolumn{7}{c}{Trumpler 14 observations} \\\\ \\\\ \\#1 & 10:43:56.0 & $-$59:32:46 & 2s & 30 & 2$\\times$14 & 28 min \\\\ \\#2 & 10:43:56.5 & $-$59:33:29 & 2s & 15 & 2$\\times$8 & 8 min \\\\ \\#3 & 10:43:53.5 & $-$59:33:28 & 2s & 15 & 2$\\times$8 & 8 min \\\\ \\#4 & 10:43:53.5 & $-$59:32:41 & 2s & 15 & 2$\\times$8 & 8 min \\\\ \\\\ \\multicolumn{7}{c}{Sky field observations (MCAO in open loop)}\\\\ \\\\ \\#1 & 10:43:07.5 & $-$59:37:29 & 2s & 30 & 8 & 8 min \\\\ \\#2 & 10:43:08.0 & $-$59:38:12 & 2s & 15 & 8 & 4 min \\\\ \\#3 & 10:43:05.0 & $-$59:38:11 & 2s & 15 & 8 & 4 min \\\\ \\#4 & 10:43:05.0 & $-$59:37:24 & 2s & 15 & 8 & 4 min \\\\ \\hline \\end{tabular} \\end{table} \\begin{figure} \\centering \\includegraphics[width=6cm]{13688fg1.png} \\caption{2MASS $K$ band image of Tr~14. The cross and the large circle indicate the centre and the size of the 2\\arcmin\\ MAD field of view. The four 57\\arcsec$\\times$57\\arcsec\\ boxes show the position of the CAMCAO camera in the adopted 4-point dither pattern, while the selected NGSs are identified by the smaller circles. } \\label{fig: fov} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=\\columnwidth,angle=-90]{13688fg2a.pdf} \\includegraphics[width=\\columnwidth,angle=-90]{13688fg2b.pdf} \\caption{Averaged FWHM and Strehl ratio maps as computed over the MAD field of view (yellow circle). The three red crosses show the locations of the NGSs. North is to the top and East to the right.} \\label{fig: fwhm} \\end{figure} Located inside Carina at a distance of 1.5-3.0 kpc, Tr~14 is an ideal target to search for multiplicity around massive stars because it contains more than 10 O-type stars and several hundreds of B-type stars \\citep{VBF96}. Large differences in the distance to Trumpler 14 arise from adopting different extinction laws and evolutionary tracks \\citep{CRV04}. Differences in distance can partially account for diverse estimates of mass and structural parameters. Its mass was first estimated to be 2000~\\msun\\ \\citep{VBF96}. However, the photometry used by these authors barely reached the turn-on point of the pre-main sequence (PMS), while they extrapolated the mass assuming a Salpeter initial mass function (IMF). More recently, \\citet{AAV07} used much deeper IR photometry, which revealed the very rich PMS population and provided a more robust mass estimate of 9000~\\msun. \\citet{VBF96} reported a core radius of 4.2~pc, while \\citet{AAV07} revised it to 1.14 pc, and detected for the first time a core-halo structure, which is typical of these young clusters \\citep[e.g.,][]{BVC04}. Tr~14 is indeed very young, not yet relaxed and has been forming stars in the last 4 Myr \\citep{VBF96}. \\begin{figure*} \\centering \\includegraphics[width=1.5\\columnwidth]{13688fg3.jpg} \\caption{False colour image of the 2\\arcmin\\ fov of the MAD observations (blue is $H$ band; red, $K_\\mathrm{S}$ band). North is to the top and East to the left.} \\label{fig: mosaic} \\end{figure*} The layout of the paper is as follows. Sections~\\ref{sect: obs} and \\ref{sect: photom} describe the observations, data reduction and photometric analysis. Section ~\\ref{sect: tr14} presents the NIR properties of Tr~14 and discusses the cluster structure. Section~\\ref{sect: comp} analyses the pairing properties in Tr~14. Section~\\ref{sect: bias} describes an artificial star experiment designed to quantify the detection biases in the vicinity of the bright stars. It also presents two simple models that generalise the results of the artificial star experiment. As such, it provides support to the results of this paper. Finally, Sect.~\\ref{sect: mst} investigates the cluster mass segregation status. and Sect.~\\ref{sect: ccl} summarizes our results. \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{13688fg4.jpg} \\caption{Close-up view of the $K_\\mathrm{S}$ band image around selected massive stars. The image size is 200$\\times$200 pixels, corresponding to approximately 5.5$\\times$5.5\\arcsec\\ on the sky.} \\label{fig: psf} \\end{figure} ", "conclusions": "\\label{sect: ccl} Using the ESO MCAO demonstrator MAD, we have acquired deep $H$ and $K_\\mathrm{S}$ photometry of a 2\\arcmin\\ region around the central part of Tr~14. The average IQ of our campaign is about 0.2\\arcsec\\ and the dynamic range is about 10~mag. The image presented in Fig.~\\ref{fig: fov} is by far the largest AO-corrected mosaic ever acquired. Using PSF photometry, we investigated the sensitivity of faint companions detected in the vicinity of bright sources. We derived several empirical relations that can be used as input for instrumental simulations, to estimate the performance of AO techniques versus seeing-limited techniques or, as done later in this paper, to build first-order analytical models of the impact of some observational biases. In particular, the contrast vs.\\ separation limit has been validated over a 5 magnitude range by an artificial star experiment. Despite a probably significant contamination by field stars, the Tr~14 CMD shows a very clear PMS population. Its location in the CMD can be reproduced by PMS isochrones with contraction ages of 3 to $5\\times10^5$~yr. Interestingly, Tr~14 cannot be significantly further away than the distance obtained by \\citet{CRV04} i.e., 2.5~kpc, as this would result in an even earlier contraction age. We derive the surface density profile of the cluster core and of different subpopulations. For stars brighter than $K_\\mathrm{S}=18$~ mag, the surface density profiles are well reproduced by EFF87 profiles over our full fov, and we provide quantitative constraints on the spatial extent of the cluster and on its stellar contents. Adopting the core-halo description suggested by \\citet{AAV07}, we report that the transition between the core and the halo is not covered by our data, implying that the core is strictly dominating the density profile in a radius of 0.9~pc at least. Using colour criteria to select the most likely cluster members, the density profiles of the more massive MS stars are best described by a power-law (or, equivalently, by an EFF87 profile with a very small core radius). We also investigated the companionship properties in Tr~14. We showed that the number of companions and the pair association process is on average well reproduced by chance alignment from a uniform population randomly distributed across the field. Only stars with a brightness ratio close to unity or with a separation of less less than 0.5\\arcsec\\ cannot be explained by spurious alignment and are thus true binary candidates. This does not imply that large light-ratio and/or wider pairs do not exist, but rather that they cannot be individually disentangled with statistical arguments. Still, 19\\%\\ of our massive star sample have a high probability physical companion. Focusing on the 0.5\\arcsec-2.5\\arcsec\\ separation range, where the observational biases are unable to invalidate our results, we compared the companion distributions of massive stars with those of lower mass stars. In Tr~14, the high-mass stars ($M>10$~\\msun) tend to have more solar-mass companions than lower-mass comparison samples. Those companions are brighter on average, thus more massive. Finally, no difference could be found in the spatial distribution of the companions of low and high-mass stars. Lastly, we employed the MST technique of \\citet{AGP09_mst} to investigate possible mass segregation in Tr~14. Again we found marginally significant results (at the 1.5$\\sigma$ level), suggesting some degree of mass segregation for the more massive stars of the cluster ($M>10$~\\msun) . Although the sensitivity of the method to incompleteness is still not fully quantified, we note that early dynamical evolution can reproduce the observed hints of mass segregation in Tr~14, despite the cluster's young age." }, "1003/1003.4512_arXiv.txt": { "abstract": "We present new radial-velocity measurements of HAT-P-13, a star with two previously known companions: a transiting giant planet ``b'' with an orbital period of 3 days, and a more massive object ``c'' on a 1.2~yr, highly eccentric orbit. For this system, dynamical considerations would lead to constraints on planet b's interior structure, if it could be shown that the orbits are coplanar and apsidally locked. By modeling the Rossiter-McLaughlin effect, we show that planet b's orbital angular momentum vector and the stellar spin vector are well-aligned on the sky ($\\lambda = 1.9\\pm 8.6$~deg). The refined orbital solution favors a slightly eccentric orbit for planet b ($e=0.0133\\pm 0.0041$), although it is not clear whether it is apsidally locked with c's orbit ($\\Delta\\omega = 36_{-36}^{+27}$~deg). We find a long-term trend in the star's radial velocity and interpret it as evidence for an additional body ``d'', which may be another planet or a low-mass star. Predictions are given for the next few inferior conjunctions of c, when transits may happen. ", "introduction": "Precise radial-velocity measurements have revealed more than 30 multiple-planet systems (Wright~2010). However, in only a few cases have transits been detected for any of the planets in those systems. Those cases are potentially valuable because the transit observables---the times of conjunction, orbital inclination, and projected spin-orbit angle, among others---provide a much more complete description of a planetary system, which may in turn give clues about its formation and evolution. The Corot-7 system has two orbiting super-Earths, one of which transits (L{\\'e}ger et al.~2009, Queloz et al.~2009). The HAT-P-7 system has a transiting hot Jupiter in a polar or retrograde orbit, as well as a longer-period companion that could be a planet or a star (P\\'al et al.~2008, Winn et al.~2009, Narita et al.~2010). The HAT-P-13 system, the subject of this paper, features a G4 dwarf star with two previously known orbiting companions (Bakos et al.~2009). The inner companion (HAT-P-13b, or simply ``b'' hereafter) is a transiting hot Jupiter in a 2.9 day orbit. The outer companion (``c'') has an eccentric 1.2~yr orbit and a minimum mass ($M_c\\sin i_c$) of about 15 Jupiter masses, although its true mass ($M_c$) and orbital inclination ($i_c$) are unknown. In particular, transits of companion c have neither been observed nor ruled out. Batygin, Bodenheimer, \\& Laughlin (2009) and Mardling~(2010) showed that it may be possible to use the observed state of the system to determine planet b's Love number $k_2$, a parameter that depends on the planet's interior density distribution. This would be of great interest, as few other methods exist for investigating the interior structure of exoplanets. The method is based on the theoretical expectation that tidal evolution has aligned the apsides of the orbits of b and c. This method has not yet yielded meaningful constraints on $k_2$, partly because of the large uncertainty in the eccentricity of b's orbit. Another relevant parameter is the mutual inclination between the orbits, which is not known at all. Radial-velocity observations are usually powerless to determine mutual inclinations, unless the planets are in a mean-motion resonance (see, e.g., Correia et al.~2010). However, for a transiting planet it is possible to assess the alignment between the orbit and the stellar equator through the Rossiter-McLaughlin (RM) effect. A system with mutually inclined planetary orbits might also be expected to have large angles between the orbits and the stellar equator. In particular, Mardling (2010) presented a formation scenario for HAT-P-13 involving gravitational scattering by a putative third companion, which could have caused large mutual inclinations and a large stellar obliquity. In this paper we present new radial-velocity measurements of HAT-P-13 bearing on all these issues. The new data are presented in \\S~\\ref{sec:rv}. Our analysis is presented \\S~\\ref{sec:analysis}, and includes evidence for a third companion ``d'' (\\S~\\ref{subsec:third}), refined estimates of the eccentricity and apsidal orientation of b's orbit (\\S~\\ref{subsec:ecc}), modeling of the RM effect (\\S~\\ref{subsec:lambda}), and updated predictions for the next inferior conjunction (possible transit window) of companion c (\\S~\\ref{subsec:dt}). In \\S~\\ref{sec:disc} we discuss the implications for further dynamical investigations of HAT-P-13. ", "conclusions": "\\label{sec:disc} HAT-P-13 was already a noteworthy system, as the first known case of a star with a transiting planet and a second close companion. We have presented evidence for a third companion in the form of a long-term radial acceleration of the star. The properties of the newly discovered long-period companion will remain poorly known until additional RV data are gathered over a significant fraction of its orbital period. Our analysis of the Rossiter-McLaughlin effect shows that planet b's orbital axis is aligned with the stellar rotation axis, as projected on the sky. Our new data also agree with the previous finding that the orbit of planet b is slightly eccentric. The latter two findings are relevant to the second reason why HAT-P-13 is noteworthy: its orbital configuration may represent an example of two-planet tidal evolution. In this scenario, first envisioned by Wu \\& Goldreich (2002) and investigated further by Mardling~(2007), tidal circularization of the inner planet's orbit is delayed due to gravitational interactions with the outer planet. The interactions drive the system into a state of apsidal alignment, where it remains as both orbits are slowly circularized. As it turned out, the specific planetary system that inspired Wu \\& Goldreich~(2002) was irrelevant to their theory, because the ``outer planet'' was found to be a spurious detection (Butler et al.~2002). Batygin et al.~(2009) welcomed HAT-P-13 as a genuine system that followed the path predicted by Wu \\& Goldreich~(2002), and with the additional virtue that the inner planet is transiting. For this interpretation to be valid, the apsides of b and c must be aligned, whereas we have found the angle between the apsides to be $36_{-36}^{+27}$~deg, differing from zero by 1$\\sigma$. We do not consider this result to be significant enough to draw a firm conclusion, especially in light of the uncertainties due to the {\\it ad hoc} stellar jitter term and our simplified treatment of the influence of companion d. Further RV monitoring and observations of occultations are needed to make progress. Batygin et al.~(2009) also showed that the existence of transits would allow for an empirical estimate of the tidal Love number $k_2$ of planet b, as mentioned in \\S~1. The requirement that the apsidal precession rates of b and c are equal leads to a condition on $k_2$, because b's precession rate is significantly affected by its tidal bulge. Subsequent work by Mardling~(2010) showed that for a unique determination of $k_2$ it is necessary for the mutual inclination $\\Delta i$ between orbits b and c to be small. If instead the orbits are mutually inclined, then tidal evolution drives the system into a state in which $e_b$ and $\\Delta i$ undergo oscillations: a cycle in parameter space, instead of a fixed point. Furthermore, Mardling~(2010) argued that a large mutual inclination should be considered plausible, or even likely, given c's high eccentricity. She proposed that b and c began with nearly circular and coplanar orbits, but c's orbit was strongly perturbed by an interaction with a hypothetical outer planet. Those same perturbations would likely have tilted c's orbit. The relation, if any, between the newly-discovered HAT-P-13d and Mardling's hypothetical outer planet is unclear. In her scenario, the outer planet is ejected from the system. This seems important to the scenario, as otherwise d would continue interacting with c, and interfere with the tidal evolution of b and c. Thus, unless d's pericenter was somehow raised to avoid further encounters with c, it does not seem likely to have played the role envisioned by Mardling~(2010). Of course the scenario could still be correct even if the third companion d was not the scattering agent; a fourth (ejected) companion may have been responsible. Our study of the Rossiter-McLaughlin effect pertains to the angle $\\psi_{\\star,b}$ between planet b's orbit and the stellar equator, and has no {\\it direct} bearing on the angle $\\Delta i$ between the orbital planes of b and c. However, there is an {\\it indirect} connection, through the nodal precession that would be caused by mutually inclined orbits. As shown by Mardling~(2010), planet b is far enough from the star that its orbital precession rate is likely to be dominated by the torque from c, rather than the quadrupole moment $J_2$ of the star. The critical orbital distance inside which the stellar torque is dominant is $\\sim$$(2 J_2 a_c^2 M_c/M_\\star)^{0.2}$ (Burns 1986), which is 0.020~AU assuming $J_2=2\\times 10^{-7}$ as for the Sun. This is smaller than the actual orbital distance of 0.043~AU. Hence if $\\Delta i$ were large, then b's orbit would nodally precess around c's orbital axis, which would cause periodic variations in $\\psi_{\\star,b}$. Therefore, at any given moment in the system's history, we would be unlikely to observe a small value of $\\psi_{\\star,b}$ unless $\\Delta i$ were small. However, it is impossible to draw firm conclusions about $\\Delta i$ because of the dependence on initial conditions, the possible effects of companion d, and the fact that only the sky-projected angle $\\lambda$ is measured rather than the true obliquity $\\psi_{\\star,b}$. It may be possible to estimate $\\Delta i$ based on transit timing variations of planet b (Nesvorn{\\'y} \\& Beaug{\\'e} 2010; Bakos et al.~2009). An even more direct estimate of $\\Delta i$ could be achieved if transits of c were detected. The existence of transits would show that $i_c$ is nearly $90^\\circ$, as is $i_b$. This would suggest $\\Delta i$ is small, although it would still be possible that the orbits are misaligned and their line of nodes happens to lie along the line of sight. The most definitive result would be obtained by observing the Rossiter-McLaughlin effect during transits of c, and comparing c's value of $\\lambda$ with that of planet b. In effect, the rotation axis of the star would be used as a reference line from which the orientation of each orbit is measured (Fabrycky 2009). This gives additional motivation to observe HAT-P-13 throughout the upcoming conjunctions of companion c." }, "1003/1003.3104_arXiv.txt": { "abstract": "Using deep $\\it Chandra$ ACIS observation data for Cygnus A, we report evidence of non-thermal X-ray emission from radio lobes surrounded by a rich intra-cluster medium (ICM). The diffuse X-ray emission, which are associated with the eastern and western radio lobes, were observed in a 0.7--7 keV $\\it Chandra$ ACIS image. The lobe spectra are reproduced with not only a single-temperature Mekal model, such as that of the surrounding ICM component, but also an additional power-law (PL) model. The X-ray flux densities of PL components for the eastern and western lobes at 1 keV are derived as $77.7^{+28.9}_{-31.9}$ nJy and $52.4^{+42.9}_{-42.4}$ nJy, respectively, and the photon indices are $1.69^{+0.07}_{-0.13}$ and $1.84^{+2.90}_{-0.12}$, respectively. The non-thermal component is considered to be produced via the inverse Compton (IC) process, as is often seen in the X-ray emission from radio lobes. From a re-analysis of radio observation data, the multiwavelength spectra strongly suggest that the seed photon source of the IC X-rays includes both cosmic microwave background radiation and synchrotron radiation from the lobes. The derived parameters indicate significant dominance of the electron energy density over the magnetic field energy density in the Cygnus A lobes under the rich ICM environment. ", "introduction": "Radio lobes, in which jets release a fraction of the kinetic energy originating from active galactic nuclei (AGNs), store enormous amounts of energy as relativistic electrons and magnetic fields. Relativistic electrons in the lobes emit synchrotron radiation (SR) at radio frequencies and boost the seed photons into the X-ray and $\\gamma$-ray ranges via the inverse Compton (IC) process. Candidates for seed photons are cosmic microwave background (CMB) photons (e.g., Harris \\& Grindlay 1979), infrared (IR) photons from the host AGN (Brunetti et al. 1997) and SR photons emitted in the lobes. If the seed photon sources can be identified, the energy densities of relativistic electrons ($u_{\\rm e}$) and magnetic fields ($u_{\\rm m}$) can be determined from a comparison of the SR and IC fluxes, respectively. These energy densities can provide important clues regarding the energy of astrophysical jets and the evolution of radio galaxies. \\begin{figure*}[t] \\centerline{ \\includegraphics[angle=0,width=7.55cm]{fig1-1.eps} \\includegraphics[angle=0,width=7.55cm]{fig1-2.eps}} \\caption{ $\\it Left$: The co-added raw 0.7--7 keV ACIS image of Cygnus A. Color scale of the image shows the photon counts for each pixel; scale bar is shown below. White contours represent the radio strength of the 1.3 GHz band observed by the VLA, the contours levels are 0.053, 0.2, 0.8, 1.2, 2, 5 and 10 $\\rm Jy~beam^{-1}$ for a beam size of 1.19$\\arcmin$$\\times$1.12$\\arcsec$. $\\it Right$: ACIS contours image of Cygnus A at 0.7--7 keV, superposed on the 1.3 GHz VLA image of the $\\rm Jy~beam^{-1}$ for each pixel; scale bar is shown below. Five contours represent X-ray brightness of 3, 6, 15, 30, 50 and 100 counts per pixel. } \\label{X-rayImage} \\end{figure*} Due to the faintness of the IC X-rays emitted from the lobes, the objects located at the edge of the cluster of galaxies have been targeted, owing to the poor radiation from the intra-cluster medium (ICM). In the past 20 years, around 30 objects have been observed with $ASCA$ and $ROSAT$ (e.g., Kaneda et al. 1995; Feigelson et al. 1995; Tashiro et al. 1998, 2001), as well as with the $\\it Chandra~X$-$\\it ray~ observatory$, $\\it XMM$-$\\it Newton$ and $\\it Suzaku$ (e.g., Burunetti et al. 2001; Isobe et al. 2002, 2005; Hardcastle et al. 2002; Grandi et al. 2003; Comastri et al. 2003; Croston et al. 2004; Kataoka \\& Stawarz 2005; Tashiro et al. 2009). In most cases, the seed photons were determined to be CMB photons. The measured IC X-ray flux from the lobes often requires that $u_{\\rm e}$ is considerably greater than $u_{\\rm m}$ (e.g., Tashiro et al. 1998; Isobe et al. 2002), as well as that the magnetic field ($B_{\\rm IC}$) is smaller than the magnetic field estimated under equipartition ($B_{\\rm eq}$), $B_{\\rm IC}/B_{\\rm eq}=0.1$--1 (e.g., Croston et al. 2005). Against this background, it is of great interest to investigate the energy balance between the relativistic electrons and the magnetic fields under ICM pressure in order to argue the energetics from the nuclei to inter-galactic space. In this paper, we present an examination of the diffuse lobe emission from one of the brightest radio lobe objects surrounded by ICM, Cygnus A. Cygnus A is a well-known FR II (Fanaroff \\& Riley 1974) radio galaxy with an elliptical host. The radio images show symmetrical double-lobe morphology with an extremely high radio flux density $S_{\\rm SR}=1598$ Jy at 1.3 GHz (B{\\^i}rzan et al. 2004), which makes it the brightest radio galaxy in the observable sky. Measuring radio fluxes between 151 MHz and 5000 MHz, Carilli et al. (1991) found spectral breaks whose break frequencies vary with position on the lobes. The spectral energy index below and above the break is 0.7 and 2, respectively. X-ray observations show diffuse X-ray emission, which is considered to originate from ICM (e.g., Smith et al. 2002), as well as a cavity corresponding to the radio lobe (e.g., Wilson et al. 2006). In addition to the thermal emissions, we examined non-thermal X-ray emissions from the lobes of Cygnus A utilizing the excellent spatial resolution of {\\it Chandra}. The structure of this paper is as follows. We describe the archived $\\it Chandra$ observation data on Cygnus A and the results of careful X-ray analysis in $\\S~2$ and $\\S~3$. Following the presentation of the results, namely, the suggestion of non-thermal X-ray emissions from the lobes, we report the spectral energy distribution of the Cygnus A lobes, as determined from radio and X-ray data, and estimate the emission of seed photons and the physical parameters of the lobes in $\\S~4$. Finally, we summarize these results in the last section. Throughout this paper, we adopt a cosmology with $H_0=71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm M}=0.27$ and $\\Omega_{\\Lambda}=0.73$ (Komatsu et al. 2009), where $1\\arcmin$ corresponds to 63 kpc at the red shift $\\it z$=0.0562 (Stockton et al. 1994) of Cygnus A. ", "conclusions": "Using $\\it Chandra$ deep observation data (230 ks) for Cygnus A, we carefully analyzed the X-ray spectra of the lobes and the regions surrounding the lobes. Our findings are as follows. \\begin{itemize} \\item In $\\it Chandra$ X-ray images, among emissions originating from ICM, we confirmed extended X-ray emission regions corresponding to the eastern and western lobes. \\item The X-ray spectra of the lobe regions could not be reproduced by a single Mekal model, and we found that the addition of a PL component was more appropriate than the addition of an additional Mekal component in the statistical analysis. The best-fit photon indices of the eastern and western lobe regions were $1.69^{+0.07}_{-0.13}$ and $1.84^{+2.90}_{-0.12}$, and the flux densities at 1 keV were $77.7^{+28.9}_{-31.9}~\\rm nJy$ and $52.4^{+42.9}_{-42.4}~\\rm nJy$, respectively. \\item The obtained X-ray and radio SED of the lobes supported the IC mechanism for X-ray emission. Furthermore, the X-rays are likely produced via both SSC processes below $\\sim 10^{18}$ Hz (4 keV) and CMB/IC processes above $\\sim10^{18}$ Hz (4 keV). This is the first case of a lobe where SSC emission has been found to affect IC emission. \\item The derived physical parameters under the SSC model indicate that the energy density of electrons dominates that of magnetic fields both in the eastern and western lobes, as often reported from other radio lobe objects. \\end{itemize}" }, "1003/1003.3797_arXiv.txt": { "abstract": "We present new {\\it KPNO 0.9-m} optical and {\\it VLA} \\HI\\ spectral line observations of the Orion dwarf galaxy. This nearby (D $\\simeq$ 5.4 Mpc), intermediate-mass (M$_{\\rm dyn} \\simeq$ 1.1\\,$\\times$\\,10$^{10}$ \\msun) dwarf displays a wealth of structure in its neutral ISM, including three prominent ``hole/depression'' features in the inner \\HI\\ disk. We explore the rich gas kinematics, where solid-body rotation dominates and the rotation curve is flat out to the observed edge of the \\HI\\ disk ($\\sim$6.8 kpc). The Orion dwarf contains a substantial fraction of dark matter throughout its disk: comparing the 4.7\\,$\\times$\\,10$^{8}$ \\msun\\ of detected neutral gas with estimates of the stellar mass from optical and near-infrared imaging (3.7\\,$\\times$\\,10$^{8}$ \\msun) implies a mass-to-light ratio $\\simeq$13. New \\halpha\\ observations show only modest-strength current star formation ($\\sim$0.04 \\msun\\,yr$^{-1}$); this star formation rate is consistent with our 1.4 GHz radio continuum non-detection. ", "introduction": "\\label{S1} Dwarf galaxies are excellent laboratories for studying the processes that shape galaxy evolution in canonically ``simple'' conditions. They lack the rotational shear intrinsic to spirals, meaning that star formation (SF) is initiated primarily by local mechanisms. Further, the structures in the ISM are not immediately destroyed, allowing one to study the interplay between SF and the ISM. Targeted observations and larger scale surveys have demonstrated the diagnostic power of \\HI\\ spectral line imaging of nearby dwarf systems (e.g., {Puche \\etal\\ 1992}\\nocite{puche92}; {Puche \\& Westpfahl 1994}\\nocite{puche94}; {Hunter \\etal\\ 2007}\\nocite{hunter07}; {Oh \\etal\\ 2008}\\nocite{oh08}; {Walter \\etal\\ 2008}\\nocite{walter08}). Nearby, gas-rich dwarf galaxies also offer comparatively ``simple'' environments in which to perform detailed kinematic analysis. Many dwarfs display solid-body rotation that is well-suited to precision rotation curve work (e.g., {Weldrake \\etal\\ 2003}). Using simple dynamical arguments, most nearby systems appear to be dark-matter dominated \\citep{mateo98}, making them important laboratories for understanding the missing baryons problem. Numerous investigations of nearby systems have revealed rich kinematic structure in the form of shells and holes in the neutral gas (e.g., {Kim \\etal\\ 1999}\\nocite{kim99}; {Walter \\& Brinks 1999}\\nocite{walter99}, {2001}\\nocite{walter01}; {Ott \\etal\\ 2001}\\nocite{ott01}; to name just a few). It is commonly assumed that these structures result from ``feedback'' processes (i.e., stellar winds and SNe; e.g., {Tenorio-Tagle \\& Bodenheimer 1988}\\nocite{tenoriotagle88}). However, debate continues on whether all such features arise from simple stellar evolution processes (e.g., {Rhode \\etal\\ 1999}\\nocite{rhode99}; {S{\\' a}nchez-Salcedo 2002}\\nocite{sanchezsalcedo02}). Here, we present new \\HI\\ and optical observations of the Orion dwarf galaxy. Originally discovered by \\citet{giovanelli79}, subsequent {\\it Arecibo} observations showed it to be \\HI-rich \\citep{springob05}. However, the low Galactic latitude ($-$12.3\\arcdeg) and high associated foreground reddening (A$_{\\rm R} =$ 1.959 mag; {Schlegel \\etal\\ 1998}\\nocite{schlegel98}) have apparently conspired to keep this nearby system away from detailed observational scrutiny to date. Below we discuss the first spatially resolved study of the ISM and stars within the Orion dwarf; we find modest-level SF in a dark-matter dominated disk that harbors a wealth of small-scale structure. The distance of the Orion dwarf remains uncertain, primarily due to its low galactic latitude and the resulting difficulty of resolved stellar population work. Previous authors have suggested distances based on the brightest stars method after solving for foreground extinction; for example, \\citet{karachentsev96} find D = 6.4 Mpc. More recent infrared work by \\citet{vaduvescu05} lowers this value to 5.4 Mpc; we adopt this latter value in the present work, but note that the distance remains a significant source of uncertainty. We summarize the basic parameters of the Orion dwarf in Table~\\ref{t1}. ", "conclusions": "\\label{S5} New \\HI\\ and optical imaging of the Orion dwarf have been presented. The stellar component occupies the inner region of a mostly well-ordered \\HI\\ disk. Column densities rise above the 10$^{21}$ cm$^{-2}$ level in multiple areas, although the central stellar component is primarily coincident with lower surface density gas in the inner disk. This area contains three regions of lower than average column density, higher than average velocity dispersion, and departures from symmetric isovelocity contours; we term these features ``holes/depressions''. The areas of active SF as traced by \\halpha\\ emission show a strong preference for high surface density gas; some of the active regions are well outside of the central stellar component of the system. The Orion dwarf is well-suited for rotation curve analysis. Through an iterative procedure we extract a rotation curve that is essentially flat at $\\sim$82 \\kms\\ out to the edge of the \\HI\\ disk ($\\simeq$6.8 kpc). At the last rotation curve point, the implied dynamical mass M$_{\\rm dyn}$ $\\simeq$ 10.6$\\times$10$^{9}$ \\msun. Applying the models of \\citet{bell01}, we use optical and infrared images to derive a stellar mass of (3.7\\,$\\pm$\\,1.5)\\,$\\times$\\,10$^{8}$ \\msun. This can be compared with the neutral gas mass of (4.7\\,$\\pm$\\,1.5)\\,$\\times$\\,10$^{8}$ \\msun\\ (which includes a correction for helium and molecular material). The Orion dwarf is thus dark matter dominated throughout the disk. We investigate the kinematics of the ``hole/depression'' regions. While these appear as prominent breaks in position-velocity space, and are evident in individual channel maps, a radius-velocity analysis does not show signs of expansion of these features at the present sensitivity and velocity resolution. Unfortunately, this lack of a measured expansion velocity for these features precludes a direct derivation of the energetic requirements for their formation. While the evidence for ``feedback'' (i.e., the effects of energy released by stellar evolution processes on the surrounding ISM) is tantalizing, we are unable to directly test this hypothesis. In addition to higher velocity resolution \\HI\\ imaging (necessary to measure a potentially slower expansion rate for the ``hole/depression'' structures), we also require deep, single-star photometry to unambiguously connect the past SF with the current ISM features. The high foreground reddening toward the Orion dwarf will make this challenging; however, the curious properties of this system suggest that the investment will be a worthwhile one in the quest to connect SF activity with characteristics of the ISM in dwarf galaxies." }, "1003/1003.3873_arXiv.txt": { "abstract": "Recent observations have gathered a considerable sample of high redshift galaxy candidates and determined the evolution of their luminosity function (LF). To interpret these findings, we use cosmological SPH simulations including, in addition to standard physical processes, a detailed treatment of the Pop III$-$Pop II transition in early objects. The simulated high-$z$ galaxies match remarkably well the amplitude and slope of the observed LF in the redshift range $5 < z < 10$. The LF shifts towards fainter luminosities with increasing redshift, while its faint-end slope keeps an almost constant value, $\\alpha \\approx -2$. The stellar populations of high-$z$ galaxies have ages of 100-300 (40-130) Myr at $z=5$ ($z=7-8$), implying an early ($z>9.4$) start of their star formation activity; the specific star formation rate is almost independent of galactic stellar mass. These objects are enriched rapidly with metals and galaxies identified by HST/WFC3 ($M_{UV}<-18$) show metallicities $ \\approx 0.1 \\Zsun$ even at $z=7-8$. Most of the simulated galaxies at $z\\approx 7$ (noticeably the smallest ones) are virtually dust-free, and none of them has an extinction larger than $E(B-V) = 0.01$. The bulk (50\\%) of the ionizing photons is produced by objects populating the faint-end of the LF ($M_{UV} < -16$), which JWST will resolve up to $z=7.3$. PopIII stars continue to form essentially at all redshifts; however, at $z=6$ ($z=10$) the contribution of Pop III stars to the total galactic luminosity is always less than 5\\% for $M_{UV}<-17$ ($M_{UV}<-16$). The typical high-$z$ galaxies closely resemble the GRB host galaxy population observed at lower redshifts, strongly encouraging the use of GRBs to detect the first galaxies. ", "introduction": "\\label{introduction} The search for the most distant galaxies, located at the beginning of the cosmic dawn, is now entering its maturity. The last few years have witnessed a tremendous increase in the data available, and the number of candidates at redshifts as high as $z=10$, corresponding to only half a billion years after the Big Bang. This has been made possible by a combination of new technologies and refined selection methods. In the first class of triggers, it is easy to acknowledge the role of the Hubble Space Telescope (HST). Thanks to dedicated surveys including the Hubble Ultra Deep Field and its predecessor, the Hubble Deep Field, we have been able to collect information on the luminosities and number counts of galaxies located at the end of the reionization epoch. Immediately after, follow-up experiments performed with the newly installed Wide Field Camera (WFC3), yielding sky images in the F105W (Y-band, $1.05 \\mu$m), F125W (J-band, $1.25 \\mu$m) and F160W bands (H-band, $1.60 \\mu$m), have allowed to push the exploration to very faint (e.g. AB mag = 28.8 in the above bands) galaxies as remote as $z=10$. In addition, the WFC3 crafted filters have considerably alleviated the contamination problem due to interlopers and provided more precise photometric redshift estimates. The standard selection method applied to these survey data sets is based on the dropout technique introduced by Steidel et al. (1996) and later constantly refined and improved by several authors (e.g. Giavalisco et al. 2004, Bouwens et al. 2007). Though this method has proved to be very solid in identifying high-redshift sources, it has the drawback that the exact source redshift cannot be determined with complete confidence. This uncertainty can be partly overcome by also using the longer wavelength infrared data, such as that provided by the {\\it Spitzer} satellite; by building a more complete Spectral Energy Distribution (SED), the stellar mass, age and redshift of a given source can be constrained further. Other complementary techniques to search for distant ($z > 5$) galaxies are also widely used, among which the narrow-band spectroscopy (Malhotra et al. 2005; Shimasaku et al. 2006; Taniguchi et al. 2005; Kashikawa et al. 2006) aimed at detecting the Ly$\\alpha$ line, carrying a large fraction of the bolometric luminosity, is definitely the most established one. Such narrow-band searches have yielded the record-holding most distant galaxy at $z=6.96$ (Iye et al. 2006). Finally, another series of experiments involve searching for remote galaxies behind foreground galaxy clusters acting as magnification lenses (Schaerer \\& Pell\\'o 2005; Richard et al. 2008; Bradley et al. 2008). Although these searches result in deeper magnitudes, their interpretation is hampered by the lens modelling and by the extremely narrow field of views, rendering it difficult to keep cosmic variance under control. As a final remark, we note that the most distant, spectroscopically confirmed, cosmic object is a Gamma Ray Burst (GRB090423 at $z=8.2$, Salvaterra et al 2009b; Tanvir et al. 2009). Although not a galaxy, the presence of this indicator implies that star formation was already well under way at those early epochs, thus further encouraging deeper galaxy searches. In addition, the GRB can be seen as a signpost of the underlying galaxy which could possibly be detected in the future knowing its exact position. Such a finding would be of the utmost importance as GRBs are mostly associated with star forming dwarf galaxies (Savaglio, Glazebrook \\& Le Borgne 2009) which are now considered to be the dominant sources of (re)ionizing photons at high redshifts (Choudhury \\& Ferrara 2007; Choudhury, Ferrara \\& Gallerani 2008). What have we learned from this wealth of experimental results ? The most solid piece of information that can be determined from the data appears to be the luminosity function (LF) and, less robustly, its evolution. It is useful to briefly recap the present observational situation marching towards increasing redshift. Bouwens et al (2007) present a comprehensive view of galaxy candidates from the UDF/ACS/GOODS fields using NICMOS in the redshift range $z=4-6$. They identify 1416 (627) V-dropouts ($i$-dropouts) corresponding to $z\\approx 5$ ($z\\approx 6$) down to an absolute UV magnitude of $M_{UV}\\approx -17$ with a LF described by a Schechter function with characteristic luminosity and faint-end slope given by $M_{UV}^* = -20.64 \\pm 0.13$ and $\\alpha=-1.66\\pm 0.09$ ($M_{UV}^* = -20.24 \\pm 0.19, \\alpha=-1.74\\pm 0.16$) respectively. The same group (Bouwens et al. 2008) has extended the data analysis to include $z\\approx 7$ $z$-dropouts (8 candidates at $z=7.3$) and J-dropouts (no candidates at $z \\approx 9$). More recently, the installation of WFC3 on board the HST has triggered a new series of searches. Oesch et al. (2010) used data collected during the first-epoch WFC3/IR program (60 orbits) in the Y, J, H bands reaching a magnitude limit of AB$\\approx 29 (5\\sigma)$. They identify 16 $z$-dropouts in the redshift range $z=6.5-7.5$ from which they obtained a LF with ($M_{UV}^* = -19.91 \\pm 0.09, \\alpha=-1.77\\pm 0.20$), essentially confirming the previous findings while extending it to fainter luminosities ($M_{UV} \\approx -18$). Bouwens et al. (2010a) pushed the investigation to $z=8.0-8.5$ by using 5 Y-dropouts. Finally, Bouwens et al. (2009) were able to identify three J-dropouts. If confirmed, these sources would be the most distant objects detected so far. Similar studies using the same data has been performed by Bunker et al. (2010), who find a comparable number of $z-$ and $Y-$dropouts. McLure et al. (2010) did not apply specific color cuts as in the previous works, thus finding a larger number of candidates; however, they pointed out that about 75\\% of the candidates at $z>6.3$ (100\\% at $z>7.5$) allow a $z<2$ interloper solution. A recent analysis of the three HUDF and of the deep ($\\sim 27.5$ AB mag), wide-area ($\\sim 40$ arcmin$^2$) WFC3 Early Release Science reveals 66 and 47 candidate galaxies at $z=7$ and $z=8$, respectively (Bouwens et al. 2010b). After carefully modelling the selection volume of each field and of the possible contamination by spurious sources, the LF obtained from these data, while consistent with previous derivations of $M_{UV}^*$ and of the normalization, shows a steeper faint-end slope with $\\alpha=-1.94\\pm 0.24$ and $\\alpha=-2.00\\pm 0.33$ at $z=7$ and $z=8$, respectively. Besides the LF, tentative information on the physical properties of these sources can be extracted from their SED, exploiting available {\\it Spitzer} data (Eyles et al. 2005; Yan et al. 2006; Stark et al. 2009). In a recent study Labb\\'e et al (2010b), based on follow-up {\\it Spitzer/IRAC} observations, analyzed the SED of 12 $z$-dropout and 4 Y-dropout candidates. None of them is detected in the {\\it Spitzer/IRAC} 3.6 $\\mu$m band to a magnitude limit of AB=26.9 $(2\\sigma)$, but a stacking analysis reveals a robust detection for the $z$-dropout sample and a strong upper limit for the Y-dropout one. The stacked SEDs are consistent with a stellar mass of about $10^9 M_\\odot$, no dust reddening, sub-solar metallicity, and best-fit ages of about 300 Myr, implying a formation epoch $z\\approx 10$. These results for the stacked sample should be compared with those obtained by Finkelstein et al (2010) who performed an object-by-object analysis and found similar ages but with a considerable spread, allowing ages as low as a few Myr. One of the major triggers to look for very high-$z$ galaxies is the quest for the reionization sources. The ionizing photon budget provided by the candidate high-$z$ galaxies is often estimated by extrapolating their LF to lower luminosities, a step that introduces a considerable uncertainty in the final determination. Having this in mind, it is still interesting to note that most studies tend to agree on the fact that the integrated UV specific luminosity for the detected galaxies at $z=7-8$ falls short of accounting for the ionizing power required to reionize the intergalactic medium. Of course, this conclusion is subject to at least two major unknown factors, these being the gas clumping factor (affecting its ability to recombine), and the escape fraction of ionizing photons (affected by dust and neutral hydrogen absorption within galaxies). Additionally, the effect of poorly constrained ages and metallicities (including the presence of metal-free, massive Pop III stars), further complicate the calculation. In spite of the large experimental effort, surprisingly little attention has been devoted by modelers to the very high redshift universe. Most of the work has so far concentrated on a semi-analytical approach (Stiavelli, Fall \\& Panagia 2004; Schneider et al. 2006; Bolton \\& Haehnelt 2007; Mao et al. 2007; Samui, Subramanian, Srianand 2009; Trenti et al. 2010) to compute the luminosity function, number counts and emissivity evolution of high-$z$ galaxies. Albeit quite fast and versatile, these methods cannot provide detailed information on the properties of the galaxies, often being based on simplified assumptions. Numerical dedicated simulations attempting to model galaxy populations beyond $z=5-6$ are also very scarce, with the partial exceptions constituted by the works by Nagamine et al. (2006) and Finlator, Dav\\'e \\& Oppenheimer (2007). Our approach is novel and different in spirit from all the previous theoretical ones. As our main aim is to model very high redshift reionization sources, we can afford smaller simulation boxes, thereby reaching the high resolutions required to resolve the dominant reionization sources - dwarf galaxies. Most importantly, though, we have implemented a careful treatment of metal enrichment and of the transition from Pop III to Pop II stars, along with a careful modelling of supernova feedback. Here we are interested in deriving the LF plus other observables from the simulations and to cast them in a form that can be compared directly with the available data or used to make new predictions for the James Webb Space Telescope (JWST). ", "conclusions": "The present results allow to build a coherent and quantitative description of the properties of elusive high-redshift, possibly primordial, galaxies. They are also very useful to interpret the data coming from deep surveys as the HST/WFC3 and future ones. However, there is considerable room for improvement left by our study. In the following we would like to elaborate on the uncertainties and shortcomings of our findings. We first note that resolving the dwarf galaxy population and following the PopIII transition process along with the large variety of physical processes implemented in the simulation limits the size of the cosmic volume that can be simulated. Resolution is certainly an important issue, as it is well known (see, e.g. Governato et al. 2010) to affect the simulated star formation rates and cause the loss of sub-galactic structures. The dependence of the results from resolution has been presented an analyzed for the same set-up of the present simulations in Tornatore et al. (2007) to which we refer the interested reader (see Fig. 1 of that paper). Resolution might also alter the details of the \"Pop III wave\" evolution, since galactic substructure allows star formation to occur at the edges of galaxies as well. The PopIII-PopII transition is also very dependent on the assumed IMF of PopIII stars. Our conclusions are valid under the assumption that PopIII stars were very massive ($M\\ge 100\\;\\Msun$) and the first metal production is driven by the explosion of pair-instability SNe (see Schneider et al. 2006 for alternatives). Our box is also too small to properly describe cosmic reionization, let alone that we are not even attempting to properly treat radiative transfer. These issues have been already addressed in previous works of our group; as already stated we are concerned here with the properties of high-z galaxies, which are presumably more affected by their internal physics rather than by the environment, as we explain below. Comparing the simulated volumes with the observed ones is very challenging as considerable uncertainty exists on the latter (see discussion in Appendix B of Bouwens et al. 2009). However, as it could be induced from the extension of the LF towards the most luminous and rarest objects at z=7, we estimate that the volume sampled by experiments should be about 30 times larger than our simulated one. The next caveat comes from the fact that we have neglected the effects of minihalos (virial temperature $< 10^4$ K). Our resolution does not allow us to track the formation of such objects, whose stellar contribution remains very uncertain due to radiative feedback effects (Haiman \\& Bryan 2006; Susa \\& Umemura 2006; Ahn \\& Shapiro 2007, Okamoto, Gao \\& Theuns 2008, Salvadori \\& Ferrara 2009). The presence of such small collapsed structures, if able to form stars, could alter the reionization history to some extent and increase the number counts of high-redshift galaxies, if detectable. As far as reionization is concerned, it has already been shown by Choudhury \\& Ferrara (2007) that acceptable fits to all relevant reionization data can be obtained without any need for PopIII stars. The bulk of the ionizing photons in those models is produced by halos with virial temperatures just above $10^4$ K, with increasingly better solution if normal, PopII stars are allowed to form in minihalos. As shown in Fig. \\ref{fig:mhalo}, where we plot the UV magnitude as function of the total mass of z=7 galaxies, sources detectable with JWST have halo masses $M_h \\simgt 10^9 \\Msun$. This corresponds to circular velocities $v_c \\simgt 50$ km s$^{-1}$. These objects are large enough that suppression by UVB photoionization filtering is at best marginal, if not negligible at all, as most of the works above agree upon. Internal mechanical feedback might be indeed more important. This process is however is already included at best in the simulations when computing the star formation rate of individual galaxies, modulo the many uncertainties that still plague our understanding of such phenomenon. \\begin{figure} \\center{\\includegraphics[scale=0.42]{highzlf_fig10.ps}} \\caption{\\label{fig:mhalo} {\\it Upper panel:} $M_{UV}$ vs. total (dark+baryonic) mass relation for simulated galaxies at $z=7$; the JWST detection limit is shown. {\\it Lower panel:} differential fractional ionizing photon rate distribution as a function of total galaxy mass. The last bin on the left shows the contribution of larger halos too rare to be caught in our relatively small simulation box, computed by integrating the observed LF down to $M_{UV}=-25$.} \\end{figure}" }, "1003/1003.4986_arXiv.txt": { "abstract": "We present direct imaging observations at wavelengths of 3.3, 3.8 (\\lprime\\,band), and 4.8 (M band) \\micron, for the planetary system surrounding HR 8799. All three planets are detected at \\lprime\\,. The $c$ and $d$ component are detected at 3.3 \\micron, and upper limits are derived from the M band observations. These observations provide useful constraints on warm giant planet atmospheres. We discuss the current age constraints on the HR 8799 system, and show that several potential co-eval objects can be excluded from being co-moving with the star. Comparison of the photometry is made to models for giant planet atmospheres. Models which include non-equilibrium chemistry provide a reasonable match to the colors of $c$ and $d$. From the observed colors in the thermal infrared we estimate \\teff$<$ 960 K for $b$, and \\teff=1300 and 1170 K for $c$ and $d$, respectively. This provides an independent check on the effective temperatures and thus masses of the objects from the \\citet{ Marois08} results. ", "introduction": "Over the past decade, a number of techniques have dramatically expanded our understanding of the nature of exoplanets. Initial detection of systems via radial velocity variations has been followed up by studies of transits, astrometric confirmation, and gravitational microlensing. All of these approaches have been helpful in developing our picture of planetary system architectures, and providing insight into formation. The direct imaging of extrasolar planets is the latest technique to provide useful information, with detection, of objects, first, around low-mass objects \\citep{Chauvin04}, followed, more recently, by the detection of several planets around intermediate mass stars \\citep{Kalas08, Marois08, Lagrange08}. Direct images of extrasolar planets not dominated by insolation have the potential to provide a wealth of information about the size, temperature, composition, and even formation history of these objects. The recently discovered planetary system around HR 8799 \\citep{Marois08}, with three massive planets at large orbital separations, may provide one of the most useful laboratories to constrain the spectral energy distribution of young giant planets. HR 8799 is a young A5V star, thought to be approximately 30-160 Myr old \\citep{Marois06}. It is known to have a bright debris disk at large orbital radii \\citep{Williams06} with approximately 0.1 \\mearth\\, of material at 50\\,K. A tenuous inner disk is thought to exist as well \\citep{Chen06}. Taken together with the planets, this suggests a system with an inner dust disk truncated by the $d$ component at 24 AU, and an outer disk truncated by the $b$ component beyond $\\sim$ 80 AU. Su et al. (2009) also find an outflow of small grains, likely caused by gravitational stirring from the planets. The star also has a number of other interesting properties, including X-ray emission \\citep{Hearty99, Schroeder07} and a deficiency in refractory metals in the stellar atmosphere \\citep[the $\\lambda$ Boo phenomenon;][]{Gray99, Gray03}. Multiple metallicity estimates are consistent it being metal-poor \\citep[Fe/H $\\simeq$\\, -0.5;][]{Gray99}, but having solar abundance lighter elements, such as carbon and oxygen \\citep{Sadakane06, Gerbaldi07}. The planets around HR 8799 are interesting, particularly in the context of planet formation alternatives. If the planets formed in situ, the core accretion hypothesis \\citep{Pollack96} would require formation of 10 \\mearth\\, cores at distances of 40 and 70 AU well before the dispersal of the gas. A disk fragmentation scenario \\citep{Boss97, Nero09, Dodson09, Helled09} may be able to more naturally explain the massive planets and their location. A more detailed look at the planet's environment and their spectral energy distributions (SEDs) can provide clues to these alternatives. Models of the spectral energy distribution for planets \\citep{Burrows97,Baraffe03} are used to estimate the temperature and mass of the objects. However, these models are currently only constrained by field brown dwarfs and objects such as 2M1207, whose formation history is uncertain \\citep{Lodato05, Mamajek07}. Are giant planets formed in a circumstellar disk around a normal star different? Objects such as the planets around HR 8799, as well as the planets around Fomalhaut \\citep{Kalas08} and $\\beta$ Pic \\citep{Lagrange08} may provide our first opportunity to address this topic. With multi-wavelength constraints of the brightness of the planets we can disentangle effects from clouds, composition, and vertical mixing that may affect the measurement of the temperature, and thus the mass of the planets. By comparison of observations with field objects and current models, we may be able to develop a better understanding of the physics leading to the SED's of giant exoplanets. The 3-5 \\micron\\, region provides access to both the 3.4 \\micron\\, CH$_4$ feature and the CO bandhead at 4.7 \\micron. Observations of this portion of the SEDs provides useful constraints on the relative amount of CO to CH$_4$ in the atmospheres of these planets. The relative absorption of these species, in the near-infrared, is used to define the transition in spectral type between the L and T sequence among field brown dwarfs. Observations of brown dwarfs in this region \\citep{Leggett07} indicate that their colors are best reproduced by models that have substantial vertical mixing between the hot lower layers and the cooler upper atmosphere. Comparison of HR 8799 planets to these objects will allow us to better understand how closely exoplanet atmospheres may mimic those of field brown dwarfs. Constraining the SED of giant planets in this region may help guide future direct imaging planet searches. Both Jupiter \\citep{Gillett69} and Gliese 229 B \\citep{Oppenheimer95} have a broad peak in the flux at 4-5 \\micron, suggesting that this may be a robust feature of brown dwarfs and giant planets in the \\teff=100-1000 K range. However, various models indicate different bands as being preferable (c.f. \\citet{Marley96}, and \\citet{Burrows97}). In order to understand whether searching in L' or M band is preferable, it would be helpful to understand giant planet colors over the expected temperature range of these objects. In this paper we present observations at 3.3, 3.8 and 4.8 \\micron\\, of the planets orbiting HR 8799. In Section 2 we describe the MMTAO and Clio camera observations of the HR 8799 system. Section 3 details the data reduction and analysis. Section 4 describes the photometric and astrometric results. Section 5 analyzes the age estimates for HR 8799. Section 6 compares these results to theoretical models and observations of brown dwarfs in the field. Finally, in Sections 7 and 8 we discuss the interpretation of the results and summarize our conclusions. ", "conclusions": "The key feature of the 3-5 \\micron\\, results is the blue color of the photometry, compared to equilibrium models. While the \\lprime-M colors can be matched with the HB07 models, the [3.3]-\\lprime\\, model colors are too red unless effective temperatures of $>$1400 K are used. Vertical mixing can explain both the M band and 3.3 \\micron\\,observations in the S06 models. The difference between these similar model families suggests that the details of how the non-equilibrium chemistry is modeled appears to be an important factor in the resulting SED. \\citet{Janson10} reach a similar conclusion from their analysis of $c$ s spectrum, in the L band. Even the S06 models have a discrepancy between the effective temperature derived from the [3.3]-\\lprime\\, color and the best fit according to the \\lprime\\, photometry, with the \\teff\\, predicted to be $\\sim$ 100-400 K lower from the \\lprime photometry compared to the color-derived temperature. The correspondence of the color-derived temperatures from S06 models with those of field brown dwarfs suggest that the model temperatures are reasonable. Based on the effective temperatures, the derived masses from the cooling models of \\citet{Baraffe03} are 12$\\pm$2 \\mjup\\, for $c$ and 11$\\pm$2 \\mjup\\, for $d$, for an assumed age of 60 Myr. The temperature limit on $b$ results in a mass limit of $<$ 9 \\mjup. These mass estimates are consistent with the estimates from \\citet{Marois08}. The large age range for the HR 8799 system (30-160 Myr) increases the uncertainty in the estimated mass of these objects. This error term is larger than that from the photometry, resulting in a mass range of $\\sim\\pm$3 \\mjup\\, for the objects. Combined with the photometric uncertainty, this indicates a mass of 12$\\pm4$ \\mjup\\, for $c$, 11$\\pm$4 for $d$, and an upper limit on $b$ of $<$ 12 \\mjup. The dimness of $c$ and $d$, given the color-derived temperatures is problematic for any evolutionary model of a giant planet. Even a very old object will not be small enough to match the \\lprime\\, apparent magnitudes. The best fits to \\teff\\, of $c$ and $d$ result in radii for the objects of 0.7 \\rjup\\, and 1.0 respectively. Given the young age of HR8799, a possible explanation of this discrepancy is that dust extinction of approximately 1.3 and 0.4 magnitudes at \\lprime\\, is reducing the observed flux for $c$ and $d$ respectively. Since HR 8799 is relatively nearby an shows no sign of foreground extinction, this dust obscuration would need to be intrinsic to the objects themselves. Alternatively, lower gravity, metallicity, details of the chemical abundance versus height, or even incorrect opacity estimates for methane may be affecting the colors. Models with \\logg=4 for the S06 were also compared to the observations, but were not different from the \\logg=4.5 models shown. However, it is possible that gravity, combined with effects not properly captured in these models (such as metallicity or details of the chemical abundance variation versus height) may be affecting the colors for these objects in a more significant way than for field brown dwarfs. Typical field brown dwarfs have median ages of $\\sim$3 Gyr \\citep{Dahn02}. Hence, the typical gravity of a \\teff\\,= 1000\\,K object in the field is \\logg\\, = 5.0, while these objects are expected to have \\logg\\,= 4.3. If the vertical mixing is similar for both the field brown dwarfs and the planets, it is reasonable to expect that the stronger than expected CO absorption, and the weaker than expected CH$_4$ absorption is an effect of the lower gravity. Interestingly, the NIR colors of these objects, as well as 2M1207b, are redder than field brown dwarfs of the same brightness \\citep{Marois08}. The NIR color difference could well be explained by their youth, and thus their lower gravity. Indeed, the young field L dwarf 2M0141-4633, found by \\citep{Kirkpatrick06}, has nearly identical NIR photometry to $c$ and $d$,and is discrepant with other field L dwarfs. The estimated mass of 6-25 \\mjup for this object indicates it should have similar gravity to the HR 8799 planets. This suggests that the red NIR colors are a low gravity effect. Kirkpatrick et al. attribute the red colors to lower collision-induced absorption of H$_2$, which primarily suppresses the K band. At lower gravities the H$_2$ CIA is reduced, causing the redder H-K color. Thus, while the color trend is opposite in the thermal IR, it appears that both effects might be explained by the youth, and low gravity, of the objects. While in principle, the balance between CO and CH$_4$ in these atmospheres is governed by Le Chatelier's principle \\citep{Burrows01}, the details of how the vertical mixing occurs, appears to be important in predicting the emergent spectra. An additional effect may be the metallicity of the atmospheres \\citep{Fortney08}. In higher metallicity objects, the formation of CO might expected to be enhanced since more oxygen is available relative to hydrogen, compared to lower metallicity atmospheres. This suggests that giant planet atmospheres would be blue in the 3-5 \\micron\\, region for more metal-rich objects. Although HR 8799 is metal-poor star, detailed observations \\citep{Gray99} indicate solar abundance for lighter elements \\citep{Sadakane06}, a consequence of the $\\lambda$ Boo phenomenon. This suggests that models with solar metallicity are probably appropriate for interpreting the colors of the HR 8799 planets. The solar metallicity (or perhaps metal-poor if the $\\lambda$ Boo phenomenon is discounted) cannot explain the blue 3-5 \\micron\\, colors. If the formation of the planets caused them to be significantly metal-enhanced relative to their stars, this may be able to explain their colors. Further observations may be able to disentangle the metallicity of these objects from the other parameters which govern their SEDs. At the same time, detailed models are needed to explore the potentially complex interplay between gravity, vertical mixing, and metallicity. It is interesting to note that mass estimates, derived from model fits to the \\lprime\\, photometry are not significantly different from color-derived estimates. For example, the non-equilibrium fits from the HB07 models, for an age of 60 Myr, result in mass estimates of 9, 10.5, and 11 \\mjup\\, for $b$, $c$, and $d$, respectively. If we use the S06 models, the estimates are 7, 9, and 10 \\mjup respectively. The HB07 are likely a better estimate of the effective temperature from the \\mlprime values, since these models treat the temperature profile in a self-consistent way with the vertical mixing, while the S06 models do not take this into account. The SED's for these objects, shown in Fig \\ref{models}, indicate that the non-equilibrium are consistent with the \\lprime-M colors, while being discrepant with the [3.3]-\\lprime\\, colors. Although the internal inconsistencies point toward a need for better models and further observations, the current results all indicate the objects are massive ($>$ 10 \\mjup), providing a challenge to models for explaining the stability of the system \\citep{Fabrycky09, Reidemeister09}, as well as its formation." }, "1003/1003.5178_arXiv.txt": { "abstract": "{In Helio- and asteroseismology, it is important to have continuous, uninterrupted, data sets. However, seismic observations usually contain gaps and we need to take them into account. In particular, if the gaps are not randomly distributed, they will produce a peak and a series of harmonics in the periodogram that will destroy the stellar information. An interpolation of the data can be good solution for this problem. In this paper we have studied an interpolation method based on the so-called 'inpainting' algorithms. To check the algorithm, we used both VIRGO and CoRoT satellite data to which we applied a realistic artificial window of a real CoRoT observing run to introduce gaps. Next we compared the results with the original, non-windowed data. Therefore, we were able to optimize the algorithm by minimizing the difference between the power spectrum density of the data with gaps and the complete time series. In general, we find that the power spectrum of the inpainted time series is very similar to the original, unperturbed one. Seismic inferences obtained after interpolating the data are the same as in the original case.} ", "introduction": "Helio- and asteroseismology are power tools to accurately determine the structure of the stellar interiors (e.g. Chris\\-ten\\-sen-Dalsgaard et al. 1996; Chaplin et al. 2008), their dynamics (Thompson et al. 1996; Garc\\'\\i a et al. 2008) as well as global parameters as their masses, radius and ages (e.g. Stello et al. 2009). To do so, it is important to have continuous data without regular gaps that would introduce a series of spurious peaks in the power spectrum (e.g. Mosser et al. 2008). For instance, the time series obtained with the observations of the CoRoT (Convection, Rotation and planetary Transits) satellite (Michel et al. 2008) are periodically perturbed by high-energy particles hitting the satellite when it is crossing the South Atlantic Anomaly (SAA) (e.g. Auvergne et al. 2009). The presence of repetitive gaps, which come from this regular perturbation, induces spurious peak in the power spectrum. To reduce the influence of these non-desirable peaks, it is commonly used to interpolate the data. In some cases, a linear interpolation is sufficient to do so (e.g. Appourchaux et al. 2008; Benomar et al 2009; Garc\\'\\i a et al. 2009, Deheuvels et al. 2010) but in other cases a more sophisticated algorithm is necessary (e.g. Mosser et al. 2009). In this paper, we propose a different algorithm based on the so-called {\\it inpainting} techniques (Elad et al. 2005; Pires et al. 2009) that seems to be especially suited for our purposes. All improvements in the gap-filling data are of special importance for the analysis of CoRoT data but also for the forthcoming Kepler observations, for which very long time series (more than 3.5 years) are being expected for thousands different stars covering the HR diagram (e.g. Bedding et al. 2010; Chaplin et al. 2010; Stello et al. 2010). ", "conclusions": "We have shown that the inpainting based on MSDCT is a powerful interpolation algorithm which is well adapted to correct the data gaps in helio and asteroseismic observations. We already applied it to the CoRoT data of the solar-like target HD\\ 170987 (Mathur et al. 2010a). We are planning to integrate it into our asteroseismic automatic pipeline for the analysis of Kepler data (Mathur et al. 2010b), and to use it to correct the GOLF velocity time series (Garc\\'\\i a et al. 2005)." }, "1003/1003.0277_arXiv.txt": { "abstract": "Surprisingly, the question ''Is there Life in the Universe outside Earth?'' has been raised, in rational terms, almost only in the western literature throughout the ages. In a first part I justify this statement. Then I try to develop an explanation of this fact by analyzing the different aspects of the notion of ''decentration''. ", "introduction": "The question ``Are we alone in the Universe?'' is one of the main motivations of this Conference ``Pathways Towards Habitable Planets.'' It is often claimed to be ``as old as Humanity itself'' It indeed looks very natural since Life is spread out over the whole Earth and therefore even a child rising his eyes toward the sky can ask ``Is there also Life out there?'' But, very surprisingly, there is almost no written occurence of this question in ``non western'' ancient cultures. In a first part of this paper I justify this statement. Then I will try to understand why it is so. I will thus be led to first clarify what can characterize and delimit ``Western'' culture. Then I will propose a hypothesis to explain why the question of Life in the Universe has almost never been raised by non-western cultures. Finally I will address the question ``why did this movement start in Greece?''. There are generally two ways to consider the question of extraterrestrial life. First, the point of view of living organisms, leading to the question ``Is there Life elsewhere in the Universe ?'', which is the subject of exobiology and extraterrestrial intelligence, leading to the question ``Are we alone?'' or``Is there anybody out there?'', which is subject of SETI (Search for Extra-Terrestrial Intelligence). And a different, but connected, question is the nature of Life: how different can it be from terrestrial life? This question is symbolized by the word ``Alien'' often found in the literature. Here I will treat these three questions as if they were only one. ", "conclusions": "The personal views presented here are open to debate. Disagreement with them is of course always possible,but any disagreeing opinion should at least offer an alternative explanation of the fact pointed out here that the extraterrestrial life debate seems to be essentially restricted to ``Western\" literature. This first attempt is not the last word and deserves further investigations, in particular the search for the possible occurrence of the extraterrestrial life debate in other parts of the world." }, "1003/1003.2334_arXiv.txt": { "abstract": "{Star clusters are studied widely both as benchmarks for stellar evolution models and in their own right. Cluster age distributions and mass distributions within galaxies are probes of star formation histories, and of cluster formation and disruption processes. The vast majority of clusters in the Universe is small, and it is well known that the integrated fluxes and colours of all but the most massive ones have broad probability distributions, due to small numbers of bright stars.} {This paper goes beyond the description of predicted probability distributions, and presents results of the analysis of cluster energy distributions in an explicitly stochastic context.} {The method developed is Bayesian. It provides posterior probability distributions in the age-mass-extinction space, using multi-wavelength photometric observations and a large collection of Monte-Carlo simulations of clusters of finite stellar masses. The main priors are the assumed intrinsic distributions of current mass and current age for clusters in a galaxy. Both UBVI and UBVIK data sets are considered, and the study conducted in this paper is restricted to the solar metallicity.} {We first use the collection of simulations to reassess and explain errors arising from the use of standard analysis methods, which are based on continuous population synthesis models: systematic errors on ages and random errors on masses are large, while systematic errors on masses tend to be smaller. The age-mass distributions obtained after analysis of a synthetic sample are very similar to those found for real galaxies in the literature. The Bayesian approach on the other hand, is very successful in recovering the input ages and masses over ages ranging between 20 Myr and 1.5 Gyr, with only limited systematics that we explain.} {Taking stochasticity into account is important, more important for instance than the choice of adding or removing near-IR data in many cases. We found no immediately obvious reason to reject priors inspired by previous (standard) analyses of cluster populations in galaxies, i.e. cluster distributions that scale with mass as $M^{-2}$ and are uniform on a logarithmic age scale.} ", "introduction": "In the early decades of astrophysics, star clusters have been our main key to the understanding of stellar evolution. While clusters continue to provide precious constraints on stellar physics, they are today studied in their own right and as tracers of the histories of galaxies. It has become clear that a significant fraction of star formation occurs in clusters, and that events such as interacting galaxies can trigger their formation \\citep{Harris1991,Meurer1995, Barton2000, DiMatteo2007}. Questions have been raised regarding the IMF in clusters in various environments, about the systematic trends in their colour distributions, about their lifetimes as gravitationally bound objects and about the initial and current cluster mass functions. Resolved observations of individual stars remain the most precise way of investigating the nature of clusters and will be possible out to distances of 10 Mpc with future extremely large telescopes. However measurements of the integrated light of unresolved star clusters reach far beyond this scale already today, and will remain the path of choice for the studies of large samples. All our studies of individual clusters and of cluster populations in galaxies rest on our ability to estimate their current ages, masses and metallicities, while accounting for extinction. The standard method of analysis of integrated cluster light is based on the direct comparison of the observed colours with predictions from {\\em continuous population synthesis models}. These models predict fluxes with the assumption that each mass bin along the stellar mass function (SMF) is populated according to the average value given by this SMF. Studies based on continuous population synthesis models have led to results that have a large impact on today's description of cluster ``demographics\". For instance, it is now usually admitted that the current cluster mass function decreases with mass as a power law with an index close to $-2$ \\citep{ Zhang1999, Bik2003, Boutloukos2003} and the debate on the cluster survival rate also rests on distributions obtained using continuous models \\citep{Vesperini1998,Fall2001,Lada2003,Rafelski2005}. The continuous approach has been coupled with statistical data analysis, for instance to provide the impression that including near-IR photometry (K band) solves the age-metallicity degeneracy for clusters \\citep{Goudfrooij2001, Puzia2002, Anders2004, Bridzius2008}. {Still in the context of continuous population synthesis, \\citet{Fernandes2010} followed by \\citet{Delgado2010} developed a Bayesian analysis of the integrated spectra of star clusters.} The continuous population synthesis models are strictly valid only in the limit of a stellar population containing an infinite number of stars. Real clusters, however, count a finite number of stars. Furthermore most of the light is provided by a very small number of bright stars, in particular in the near-IR. The so-called {\\em stochastic fluctuations} in the integrated photometric properties are the result of the random presence of these luminous stars. Some of these can be quantified using selected information provided by continuous population synthesis models \\citep[e.g.][]{Lancon2000, Cervino2002, Cervino2004, Cervino2006}, but others require the use of {\\em discrete population synthesis models} \\citep{Barbaro1977, Girardi1993, Bruzual2002,Deveikis2008,Popescu2009, Piskunov2009}. The predicted luminosity and colour distributions depend strongly on the total mass (or star number) in the cluster, and can be far from Gaussian even when the total mass exceeds $10^5$~M$_{\\odot}$. The most probable colours are offset from those predicted by continuous population synthesis when masses are below $10^4$~M$_{\\odot}$, because the single most luminous star in such clusters will be more often on the main sequence than in the red giant phases of evolution. Attempts to describe the colour distributions analytically have made progress \\citep[e.g.][]{Cervino2006}, but are not yet easily applicable. The present piece of work is based on discrete population synthesis. For the first time, we use the discrete models not only to predict colour distributions but to {\\em analyse} the energy distributions of clusters. We present a Bayesian approach to the probabilistic determination of age, mass and extinction, based on a large library of Monte-Carlo simulations of clusters. {This method is a close analog to the one introduced by \\citet{Kauffmann2003} for the study of star formation histories in the Sloan Digital Sky Survey. However the variety of observable properties has completely different origins in both contexts: stochasticity at a given age, mass and metallicity plays a predominant role here, while different star formation histories provide all the diversity in the model collections used for galaxy studies. } We compare determinations based on the Bayesian approach with traditional estimates, thus providing a new insight into systematic effects and their consequences. In this first paper, we focus on data sets consisting of either UBVI or UBVIK photometry. Future work will extent to other pass-bands and the addition of the metallicity dimension. ", "conclusions": "Studies of star cluster populations in galaxies have been based until now on {\\em continuous} population synthesis models, that provide a very poor approximation of the integrated light of clusters of small and intermediate masses because this light is determined by a very small number of luminous stars. Based on large collections of Monte-Carlo simulations of star clusters that each contain a finite number of stars, this paper explores systematic errors that occur when the integrated fluxes of realistic clusters of small masses are analysed in terms of mass and age. Our main collection is built with the cluster age and mass distributions of \\citet{Fall2009}, extrapolated to masses lower than those observed in the Antennae galaxies. With the standard methods (continuous models), large systematic errors affect estimated ages and large random errors affect masses. If observational uncertainties on cluster fluxes are large and, as a consequence, quality-of-fit criteria fail to reject the numerous poor fits, systematic errors (of a few tenths of a dex) are also present in the estimated masses. Derived age-mass or age-luminosity distributions for samples in which actual ages and masses are distributed as in our main collection display clustered patterns that very closely resemble those found in empirical samples in the current literature. We find no immediately obvious reason to reject the age and mass distributions of our main collection, but clearly this essential point requires detailed study with real observations. A Bayesian method has been described and implemented, in order to account explicitly for the finite nature of clusters in the analysis. It is shown that their age and mass can be recovered with error bars that will be small enough for many purposes. Young small mass clusters will remain difficult to age-date because HR-diagrams of many of them identically look like truncated main sequences with no ionizing or post-main sequence stars. At intermediate ages, the variability of luminous AGB stars is expected to cause difficulties that we have not yet solved. The comparison between the results obtained with UBVI and UBVIK data sets shows that, in the stochastic context, the benefits of adding the K band to optical observations are rather small, except for the mass determination of clusters older than 1\\,Gyr. Clearly, adding near-IR or UV information is secondary, compared to the need to move from continuous to stochastic cluster models. The Bayesian analysis method can now be applied to existing data on cluster samples in nearby galaxies with the aim of constraining the actual age and mass distributions of these clusters. We will also extend the study of systematic errors to the case where metallicity is an unknown parameter." }, "1003/1003.4914_arXiv.txt": { "abstract": "{% Sulphur is a volatile $\\alpha$-element which is not locked into dust grains in the interstellar medium (ISM). Hence, its abundance does not need to be corrected for dust depletion when comparing the ISM to the stellar atmospheres. The abundance of sulphur in the photosphere of metal-poor stars is a matter of debate: according to some authors, [S/Fe] versus [Fe/H] forms a plateau at low metallicity, while, according to other studies, there is a large scatter or perhaps a bimodal distribution. In metal-poor stars sulphur is detectable by its lines of Mult.\\,1 at 920\\,nm, but this range is heavily contaminated by telluric absorptions, and one line of the multiplet is blended by the hydrogen Paschen $\\zeta$ line. We study the possibility of using Mult.\\,3 (at 1045\\,nm) for deriving the sulphur abundance because this range, now observable at the VLT with the infra-red spectrograph CRIRES, is little contaminated by telluric absorption and not affected by blends at least in metal-poor stars. We compare the abundances derived from Multiplets 1 and 3, taking into account NLTE corrections and 3D effects. Here we present the results for a sample of four stars, although the scatter is less pronounced than in previous analysis, we cannot find a plateau in [S/Fe], and confirm the scatter of the sulphur abundance at low metallicity.} ", "introduction": "\\sloppy The light elements between O and Ti of even atomic number are referred to as $\\alpha$-elements because they are mainly produced by successively adding an $\\alpha$-particle from nucleus to nucleus. Among them is sulphur, a volatile element which is not locked into dust in the interstellar medium (ISM), so that the sulphur abundance derived for the ISM can be directly compared to the sulphur abundance derived in stars. The $\\alpha$-elements are crucial probes of the chemical evolution of a stellar population: they are almost exclusively released by Type II Supernovae, while the iron peak elements are produced by Type II SNe, but also, in large amounts, by Type Ia Supernovae. Progenitors of Type II and Type Ia supernovae have very different lifetimes, making the abundance ratio of $\\alpha$-elements to iron-peak elements a powerful diagnostic of the chemical evolution and star formation history of a galaxy. In the Milky Way, stars of lower metallicity are characterised by higher $\\alpha$ to iron abundance ratios than found in the Sun and stars of solar metallicity. This is usually interpreted in terms of the lower contribution of Type Ia SNe. Systems which are characterised by low or bursting star formation, like dwarf Spheroidal galaxies, give time to Type Ia SNe to explode before the enrichment due to Type II SNe has greatly increased. Consequently such systems display rather low $\\alpha$ to iron ratios even at low metallicities. Thus the abundance of $\\alpha$-elements is an important property of any stellar population. For the study of chemical evolution in external galaxies, the more readily available objects are Blue Compact galaxies (BCGs) through analysis of the emission line spectra, and Damped Ly$-\\alpha$ systems (DLAs) through the analysis of resonance absorption lines. In both groups of objects, sulphur is relatively easy to measure in the form of ISM emission (for the BCGs) or absorption (for the DLAs) lines. The investigation of sulphur abundances in the stellar photospheres started with the pioneering work of \\cite{GeorgeW}. They determined the sulphur abundance in six stars out of a sample of nine stars in six Galactic Clusters. Later on, \\cite{clegg81} determined the sulphur content in 20 F- G-type stars with [Fe/H]$\\ge -1$, while \\cite{fran87,fran88} studied 13 and 12 metal-poor stars, respectively. Not many sulphur lines are available in the observed stellar spectra. There is a forbidden line from the ground level of \\ion{S}{i} at 1082.1\\,nm, which is weak, blended, but for which lower and upper level populations are very close to local thermal equilibrium (LTE). This absorption line is measurable in the spectrum of solar-like stars \\citep{ryde06}. However, it becomes undetectable below [Fe/H]$<-0.5$. The permitted Mult.\\,8 at 675\\,nm is weak, but not blended, neither contaminated by telluric absorption. For these transitions the assumption that high and low levels are in LTE is a good approximation, and the same holds for the lines of Mult.\\,6 at 869\\,nm \\citep{takada02}. Both Multiplets 6 and 8 are weak, as a consequence detectable only in stars of solar, or moderately sub-solar metallicity (down to [Fe/H]$\\sim -1.5$). Below such metallicity, usually Mult.\\,1 at 920\\,nm is used, which is, on the other hand, contaminated by telluric absorption. Moreover one of the lines of the triplet (922.8\\,nm) is located in the blue wing of the hydrogen Paschen $\\zeta$ line, and the presence of the hydrogen line makes the abundance determination more difficult. The infra-red lines of Mult.\\,3 at 1045\\,nm are well suited to measure the sulphur abundance. These lines are not as strong as the components of Mult.\\,1, but easily detectable in very metal-poor stars. No blend has to be taken into account in the case of metal-poor stars because the only extra line present in the range, the \\ion{Fe}{i} line at 1045.5\\,nm blending the strongest \\ion{S}{i} line of the triplet, vanishes at sub-solar metallicity. Telluric absorptions in this range are less abundant and less strong than in the case of Mult.\\,1. Observing the 1045\\,nm sulphur lines provides thus a possibility to obtain a reliable sulphur abundance in very metal-poor stars. \\cite{zolfo} analysed the sulphur abundance in a sample of Galactic stars. They investigated, when available, the lines of Multiplets 6, 8, and 1. Because of the blend of the 922.8\\,nm \\ion{S}{i} line of Mult.\\,1 with the hydrogen Paschen $\\zeta$, they performed line profile fitting to derive the sulphur content for all the lines available. The line profile fitting procedure permitted to reproduce the line profile of the hydrogen. In their study they have suggested that in the range $-2.5<$[Fe/H]$<-2.0$ the [S/Fe] ratio shows either a large scatter or a bimodal behaviour. Most of the stars lie on a ``plateau'' at about [S/Fe]=+0.4, while a non negligible number of stars shows a ``high'' value of [S/Fe], around +0.8. This behaviour has no proposed interpretation, but it may be supposed to be due to systematic errors affecting only the analysis of Mult.\\,1. In fact, in the sample of \\cite{zolfo}, the determination of [S/Fe] in this range of metallicity is based mainly on the non contaminated lines of Mult.\\,1. On the other hand, \\cite{nissen04} and \\cite{nissen07} derived the sulphur abundance from equivalent width (EW) measurements of the lines of multiplets 8 and 1, and they find a plateau at low metallicity in the [S/Fe] versus [Fe/H] plot, with no sign of bimodal distribution or scatter. It is therefore of great interest to verify this puzzling finding by the use of an independent and, probably better, diagnostic of the sulphur abundances, as can be afforded by the 1045\\,nm lines. ", "conclusions": "We did an analysis on the \\ion{S}{i} lines of Mult.\\,3 for four metal-poor dwarfs stars of the sample of \\cite{zolfo}. The results of three out of the four candidates are consistent with the previous analysis within errors, and for these three stars the previous analysis relied on Mult.\\,1. We have to notice that for three stars the sulphur abundance we obtain from the lines of Mult.\\,3 is lower with respect to the previous analysis based on Mult.\\,1. For HD\\,181743, whose previous analysis was based on Mult.\\,6, the difference of 0.5\\,dex is above the expected error, and can now be explained with a contamination from the sky in the spectrum analysed by \\cite{zolfo}. From the results of this small sample we cannot confirm the existence of a plateau in the [S/H],[Fe/H] plane. The scatter is larger with respect to the estimated uncertainties, but a similar scatter is present in [Mg/Fe] versus [Fe/H] in the analysis of \\cite{gratton03}. The systematic error related to the oscillator strength of the lines of Mult.\\,1 can be neglected because it acts in the same way for the complete sample of stars. One has to take into account the uncertainties related to the temperature, but these error cannot remove the great scatter in [S/Fe]. In fact for the cooler star a change of 100\\,K in the effective temperature is translated in a change of 0.10\\,dex in the sulphur abundance, while for the hotter star the same difference in temperature would produce only 0.04\\,dex change in A(S). An analysis on Mult.\\,3 in an extended sample of stars would be useful to investigate the behaviour of [S/Fe] at low metallicity." }, "1003/1003.4125_arXiv.txt": { "abstract": "Within the distance of 1 pc from the Galactic center (GC), more than 100 young massive stars have been found. The massive stars at 0.1--1 pc from the GC are located in one or two disks, while those within 0.1 pc from the GC, S-stars, have an isotropic distribution. How these stars are formed is not well understood, especially for S-stars. Here we propose that a young star cluster with an intermediate-mass black hole (IMBH) can form both the disks and S-stars. We performed a fully self-consistent $N$-body simulation of a star cluster near the GC. Stars escaped from the tidally disrupted star cluster were carried to the GC due to a 1:1 mean motion resonance with the IMBH formed in the cluster. In the final phase of the evolution, the eccentricity of the IMBH becomes very high. In this phase, stars carried by the 1:1 resonance with the IMBH were dropped from the resonance and their orbits are randomized by a chaotic Kozai mechanism. The mass function of these carried stars is extremely top-heavy within $10''$. The surface density distribution of young massive stars has a slope of $-1.5$ within $10''$ from the GC. The distribution of stars in the most central region is isotropic. These characteristics agree well with those of stars observed within $10''$ from the GC. ", "introduction": "More than 100 young massive stars have been found in the Galactic centerr (GC) by near-infrared observations \\citep{1995ApJ...447L..95K,2006ApJ...643.1011P,2006JPhCS..54..279L, 2009ApJ...703.1323D,2010ApJ...708..834B}. These stars appear to reside in one or two disks at more than $1''$ from the GC \\citep{2006JPhCS..54..279L,2006ApJ...643.1011P}, while B-type stars within $1''$ (S-stars) have an isotropic and thermal distribution \\citep{2003ApJ...596.1015S,2009ApJ...692.1075G}. Two major scenarios have been proposed for the formation of these stars. One is the in-situ formation in an accretion disk \\citep{2003ApJ...590L..33L} and the other is the migration of a star cluster formed several parsec or more away from the GC \\citep{2001ApJ...546L..39G}. In the 1990s, only OB and Wolf--Rayet (WR) stars could have been observed in the GC. However, recent advance in observation techniques has made it possible to identify many late-type stars, and now we can obtain their spacial distribution. While the OB and WR stars have a power-law distribution with a slope of about $-1.5$, the slope for all stars including late-type giants is nearly 0 \\citep{2009ApJ...703.1323D,2009A&A...499..483B,2010ApJ...708..834B}. The orbits of S-stars appear to be isotropic and thermal \\citep{2003ApJ...596.1015S}, but the most recent observation showed that it is more eccentric with the distribution of $n(e)\\sim e^{2.6}$ \\citep{2009ApJ...692.1075G}. How these distributions have formed is the key to the understanding of the formation process of the young stars near the GC. The mass function of the young stars have also been determined from the observations. It is quite top-heavy \\citep{2009ApJ...703.1323D,2009A&A...499..483B,2010ApJ...708..834B}. Formation scenarios should also explain the mass function. In-situ formation of S-stars is difficult because of the strong tidal field of the central supermassive black hole (SMBH). Therefore, it is necessary to carry young stars to the current location from somewhere outside. Recently, several scenarios have been proposed to carry stars which were born on a gaseous disk at around 0.1--0.5 pc from the GC and to randomize their orbits. The scenarios include migration via the gravitational torques in the stellar disk \\citep{2007MNRAS.374..515L,2009ApJ...702L...1G}, migration as a star cluster core \\citep{2006ApJ...650..901B}, randomization of the orbital elements of S-stars by an IMBH \\citep{2009ApJ...705..361G,2009ApJ...693L..35M} or stellar-mass black holes \\citep{2009ApJ...702..884P}, and the formation of S-stars due to disruptions of binaries \\citep{2008ApJ...683L.151L,2009ApJ...697L..44M}. On the other hand, \\cite{2008Sci...321.1060B} suggested the direct formation of S-stars in a gaseous disk formed from a giant molecular cloud infalling to the GC. However, none of them is well established. We have shown that star clusters can carry young stars to the GC, if an IMBH forms in the clusters. In \\citet{2009ApJ...695.1421F}, we performed a fully self-consistent $N$-body simulation, in which both the internal dynamics of a star cluster and the interaction between the cluster and its parent galaxy are handled correctly. The star cluster migrating into the GC was completely disrupted by the tidal force and the stars escaping from the cluster formed a disk structure. Before the disruption, an IMBH formed through the runaway collisions of stars in the cluster. We found a new migration mechanism of young stars, a 1:1 mean motion resonance with the IMBH. The IMBH carries young stars to the GC by the 1:1 resonance after the disruption of the cluster. In this simulation, however, the spacial resolution around the SMBH was limited to 0.2 pc because of the use of a large softening length for the SMBH. Hence, it was impossible to compare the distributions of the young stars obtained in the simulation with the observed one. In this Letter, we report the result of a new simulation performed using our improved code which does not need softening for the SMBH. In this simulation, we can follow the orbits of stars down to the AU scale, where S-stars reside. We found that the distribution of stars obtained from the simulation agrees very well with the observations in the following three points: the surface density of young massive stars has a slope of $-1.5$, the young stars have an extremely top-heavy mass function, and the orbits of young stars in the inner most region are thermal and isotropic. A sinking star cluster can explain both a young stellar disk and S-stars at the same time. We describe the method of our $N$-body simulation in Section 2. In Section 3, we show the results of simulations. Section 4 is for summary. ", "conclusions": "We performed a fully self-consistent $N$-body simulation of a star cluster near the GC. The star cluster migrated to the GC owing to the dynamical friction and was disrupted by the tidal force. The stars which are initially the members of the star cluster formed a disk structure. Before the disruption, an IMBH formed in the cluster via runaway collisions of stars. After the disruption of the cluster, the IMBH continued to sink to the GC and stars which were caught in the 1:1 mean motion resonance of the IMBH also sank to the GC. Near the GC, the spiral-in of the IMBH slowed because of the depletion of field stars, and the orbit of the IMBH became highly eccentric. At this stage, the stars were kicked out from the resonance and the orbits were efficiently randomized by the non-axisymmetric perturbing potential of the IMBH. We investigated the distributions of the stars carried to the GC by the star cluster and the IMBH. We found that they agree well with the observed ones. The surface density within $10''$ had a slope of $-1.5$. The mass function of the young stars was extremely top-heavy in the inner-most region. The eccentricities and inclinations of the young stars carried near the central SMBH by the resonance were a thermal and isotropic distribution, while young stars in the outer region were distributed in a disk. These distributions agree with that of S-stars and a young stellar disk. Thus, the distributions of ``debris stars'' of the sinking star cluster agree well with the observations. The star cluster scenario with an IMBH can explain the origin of both a young stellar disk and S-stars. Here, we discuss possible conditions for the formation of S-stars. In our simulation, the distances of young stars carried by the resonance were larger than those of S-stars. However, how deeply the IMBH can carry stars to the GC depends on the distance where the orbit of the IMBH becomes eccentric. It occurs when the mass of the IMBH is comparable to the enclosed mass of the field stars \\citep{2007ApJ...656..879M} and this distance strongly depends on the density distribution of the GC and the mass of the IMBH. First, we discuss the density distribution. The enclosed mass of the field stars in our model is only $\\sim10$\\% of that estimated by observed visible stars \\citep{2007A&A...469..125S} at 0.15 pc, where the orbit of the IMBH became eccentric in our simulation. If the density profile is a broken power-law fitted by observed visible stars, the enclosed mass is $\\sim 1000 M_{\\sun}$ even at 0.01 pc \\citep{2007A&A...469..125S,2003ApJ...594..812G}. We also estimate the case of a power-law with $-7/4$, which is theoretically expected \\citep{1976ApJ...209..214B}. In this case, the enclosed mass is 8000$M_{\\sun}$ at 0.01 pc (2000 AU) and 1000 $M_{\\sun}$ at $2\\times10^{-3}$ pc (400 AU). These values are smaller than the upper limit of the enclosed mass from the observation of S2, 3--4 $\\times10^5M_{\\odot}$ at 0.01 pc \\citep{2008ApJ...689.1044G}. Recent observations show a flat density profile of old stars \\citep{2009ApJ...703.1323D,2009A&A...499..483B,2010ApJ...708..834B}. However, it is unlikely that this distribution reflects the true mass distribution. There are certainly dark masses composed of stellar mass black holes, neutron stars, and white and brown dwarfs. On the other hand, the timescale of collisions between main-sequence stars is pretty small, around 0.1--1 Gyr for inner 1 pc, if a stellar cusp developed through thermal relaxation \\citep{1983ApJ...268..565D}. Therefore, the lack of old stars is probably the result of collisional disruptions. Compact objects are not disrupted by collisions, and brown dwarfs are also less likely to be disrupted by collisions due to their high density. Thus, it seems natural to assume that the innermost region of the GC is dominated by dark mass. Next, we discuss the possible range of the IMBH mass. In our simulation, the IMBH is more massive than the observational upper limit, $\\sim 10^4 M_{\\odot}$ \\citep{2004ApJ...616..872R}, but it is possible to form smaller IMBHs from different initial conditions or with higher mass loss rates \\citep{2009ApJ...695.1421F}. Even if we assume an extreme mass loss rate, it is possible to form a star with a few thousand solar masses through runaway collisions because the star cluster near the GC is very compact. For the lower limit, $1500 M_{\\odot}$ is sufficient for the randomization of stars \\citep{2009ApJ...693L..35M}. Less massive IMBHs can carry stars closer to the GC, but it takes a longer time to migrate due to the dynamical friction. Star clusters migrate to around 1pc and are disrupted there. Its typical timescale is 2--10 Myr for a star cluster at 5--10 pc from the GC with around $10^5 M_{\\sun}$\\citep{2009ApJ...695.1421F}. After the disruption, IMBHs migrate due to the dynamical friction. We estimated the timescale for IMBHs migrating from 1 pc to $10^{-3}$ pc using the equation derived from \\citet{2007ApJ...656..879M}. We assumed that the Bahcall--Wolf cusp and obtained 1 Myr for the IMBH with 16000$M_{\\odot}$. This result is consistent with our simulation. We also tried broken power laws \\citep{2003ApJ...594..812G} and obtained similar results. The timescale is inversely proportional to the IMBH mass. Thus, IMBHs with a few thousand solar masses are capable for this scenario." }, "1003/1003.4313_arXiv.txt": { "abstract": "We present a spectroscopic analysis of white dwarfs found in the Kiso survey. Spectroscopic observations at high signal-to-noise ratio have been obtained for all DA and DB stars in the Kiso Schmidt ultraviolet excess survey (KUV stars). These observations led to the reclassification of several KUV objects, including the discovery of three unresolved DA+DB double degenerate binaries. The atmospheric parameters ($\\Te$ and $\\logg$) are obtained from detailed model atmosphere fits to optical spectroscopic data. The mass distribution of our sample is characterized by a mean value of 0.606 \\msun\\ and a dispersion of 0.135 \\msun\\ for DA stars, and 0.758 \\msun\\ and a dispersion of 0.192 \\msun\\ for DB stars. Absolute visual magnitudes obtained from our spectroscopic fits allow us to derive an improved luminosity function for the DA and DB stars identified in the Kiso survey. Our luminosity function is found to be significantly different from earlier estimates based on empirical photometric calibrations of $\\mv$ for the same sample. The results for the DA stars now appear entirely consistent with those obtained for the PG survey using the same spectroscopic approach. The space density for DA stars with $\\mv\\le12.75$ is $2.80\\times10^{-4}$ pc$^{-3}$ in the Kiso survey, which is 9.6\\% smaller than the value found in the PG survey. The completeness of both surveys is briefly discussed. ", "introduction": "White dwarf stars represent the final stage of stellar evolution for main sequence stars whose masses lie between 0.07 and 8 \\msun, which correspond to about 97\\% of the stars in the Galaxy. As the result of the cessation of nuclear reactions, they simply cool off while dissipating the content of their thermal reservoir. Because of these characteristics, the white dwarf luminosity function --- the number of white dwarfs as a function of their intrinsic luminosity --- is a powerful tool that provides an estimate of the contribution of white dwarf stars to the density of matter in the Galaxy. When the luminosity function is derived from a complete sample of white dwarfs, it contains information such as a direct measure of the stellar death rate in the local galactic disk. The comparison with theoretical evolutionary models \\citep[see, e.g.,][]{fon01} then allows us to measure the age of various components of the Galaxy. Likewise, the mass distribution contains information about the amount of mass lost during the evolution of an initial mass distribution \\citep[][hereafter LBH05]{LBH05}. Luminosity functions have been determined for cool white dwarfs discovered in high proper motion surveys, and through UV color-excess surveys for hot white dwarfs. The cool end of the luminosity function was studied by \\citet{ldm88} and \\citet{leggett98}, while the hot end was analyzed by \\citet{fleming86} using white dwarfs identified in the Palomar-Green (PG) survey \\citep{PG} and by Darling (1994; see also \\citealt{wegner94}) using the Kiso Schmidt ultraviolet excess survey (KUV). LBH05 have recently improved upon the analysis of \\citet{fleming86} by applying the spectroscopic method \\citep{bergeron92} to measure the effective temperatures and surface gravities of all the DA stars identified in the PG survey. The improved atmospheric parameters allowed a better determination of the absolute visual magnitude ($M_V$) of each star, and in turn improved the accuracy of the luminosity function calculation. More recently, a similar approach was applied by \\cite{harris06} to a sample of 6000 white dwarfs (or white dwarf candidates) identified in the Sloan Digital Sky Survey (SDSS) Data Release 3. The luminosity function based on the SDSS extends to redder colors than the PG or Kiso surveys, but does not extend to low enough temperatures to cover the end of the white dwarf cooling sequence, although discoveries made through proper motion diagrams may soon change this picture \\citep{kilic06}. In order to measure the luminosity function of white dwarf stars, one needs to carefully define a statistically complete sample. This is a major endeavor, both for common proper motion or UV color excess surveys. \\citet{darling94} attempted to estimate the completeness of the PG and KUV surveys by counting the number of stars discovered in the overlapping fields of both surveys (see also Section 4.4 below). LBH05 discuss at length the completeness of the PG survey (their Section 4) by comparing their improved luminosity function based on spectroscopic $M_V$ values with the results of \\citet{darling94} for the KUV survey (see Figure 10 of LBH05). They evaluate the completeness of the PG survey to 75\\%, while \\citet{darling94} found a lower value of 58\\%. However, this comparison between the PG and the Kiso surveys is fundamentally flawed for at least two reasons. First, in the case of the PG survey, only DA stars are considered while for the Kiso survey, white dwarf stars of {\\it all spectral types} are used in the calculation of the luminosity function. Second, and most importantly, the results of \\citet{darling94} are based on $M_V$ values determined from empirical photometric calibrations; this is similar to the approach used by \\citet{fleming86} for the PG survey. In this paper, we present a study aimed at improving the luminosity function of white dwarf stars by applying the spectroscopic method to the DA and DB stars identified in the Kiso survey. While our primary goal is to improve the comparison of the luminosity functions derived from the PG and Kiso surveys, we also provide an analysis of the global properties of the KUV sample and demonstrate how samples of relatively bright white dwarf stars still hide objects of significant astrophysical interest. Our sample drawn from the Kiso survey is presented in Section 2 and analyzed in Section 3 using the spectroscopic technique. The absolute visual magnitudes obtained from these accurate atmospheric parameter determinations are then used in Section 4 to calculate an improved luminosity function for these stars. A detailed comparison with the results of the PG survey is also presented. Our conclusions follow in Section 5. ", "conclusions": "We presented an analysis of the DA and DB white dwarfs in the KUV survey, and determined the atmospheric parameters for each star from detailed model atmosphere fits to optical spectroscopic data. The $\\mv$ values derived from the atmospheric parameters were compared with those of \\citet{darling94}, which were obtained from photometric empirical calibrations. Our study allowed us to measure directly the impact of the use of state-of-the-art model atmospheres on the determination of absolute magnitudes for white dwarfs. The differences were found to be significant, but had a somewhat smaller impact on the calculation of the luminosity function. We then proceeded to derive the luminosity function of DA and DB stars found in the Kiso survey. We find as a result of our improved $M_V$ values a smaller number of stars in the fainter magnitude bins than estimated by \\citet{darling94}. The comparison of our luminosity function with that of LBH05, for DA stars only, reveals that both functions are similar. We obtained a local space density of white dwarfs of $5.49\\times10^{-4}$ $\\rm{pc}^{-3}$, while this number drops to $2.80\\times10^{-4}$ $\\rm{pc}^{-3}$ for $\\mv\\le12.75$. These results are now entirely consistent with those published in LBH05 for the PG survey and the completeness of both surveys appears comparable. Our spectroscopic survey of white dwarfs in the Kiso survey has also led to an important spectral reclassification of the sample published in \\citet{darling94}. In particular, we have identified three unresolved double degenerate binaries. A two-component fit confirmed that KUV 02196+2816, KUV 03399+0015, and KUV 14197+2514 are unresolved double degenerate binaries composed of a DA and a DB star. These systems were easily identified in our analysis because both components had different spectral types. However, double degenerates composed of two DA stars would go totally unnoticed, as demonstrated by \\citet{liebert91}. One may wonder how many such binaries might be hiding in spectroscopic surveys such as KUV or PG. This in turn could affect our determination of the true space density of white dwarf stars. This era of large scale surveys where the samples can contain up to thousands of stars will certainly help us in the characterization of our Galaxy. Accurate statistical analyses will then provide even more precise determinations of, for example, the white dwarf space densities in the different populations of the Galaxy, the stellar contribution to the mass of the Galaxy, the age of the local galactic disk, the stellar formation and death rates, etc. The issue of completeness is thus of great importance when statistical analyses of these samples are considered. We would like to thank the director and staff of Steward Observatory and Carnegie Observatories for the use of their facilities. We would also like to thank A.~Gianninas for the acquisition of the spectra in the southern hemisphere and for a careful reading of this manuscript. This work was supported in part by the NSERC Canada and by the Fund FQRNT (Qu\\'ebec). P.B. is a Cottrell Scholar of Research Corporation for Science Advancement. \\clearpage" }, "1003/1003.3475_arXiv.txt": { "abstract": "We consider the 1.5 years {\\it Fermi}/Large Area Telescope light curves ($E>100$ MeV) of the flat spectrum radio quasars 3C 454.3 and PKS 1510--089, which show high activity in this period of time. We characterise the {\\it duty cycle} of the source by comparing the time spent by the sources at different flux levels. We consider in detail the light curves covering periods of extreme flux. The large number of high--energy photons collected by LAT in these events allows us to find evidence of variability on timescales of {\\it few hours}. We discuss the implications of significant variability on such short timescales, that challenge the scenario recently advanced in which the bulk of the $\\gamma$--ray luminosity is produced in regions of the jet at large distances (tens of parsec) from the black hole. ", "introduction": "Powerful $\\gamma$--ray emission is a distinctive feature of Flat Spectrum Radio Quasars (FSRQs), radio--loud active galactic nuclei with the relativistic jet closely oriented towards the Earth. Gamma rays with energy above 100 MeV from these sources are widely believed to be produced through the inverse Compton (IC) scattering between highly relativistic electrons in the jet and ambient photons, either the optical--UV photons from the accretion disk (e.g. Dermer \\& Schlickeiser 1993) or reprocessed by the gas in the broad line region (BLR, e.g. Sikora et al. 1994) or the infra--red photons from the dusty torus (e.g. Blazejowski et al. 2000). The sources of target photons dominating the IC emission is basically determined by the location of the emitting region in the jet (e.g. Ghisellini \\& Tavecchio 2009, Sikora et al. 2009). Nuclear optical--UV seed photons dominate if $\\gamma$--rays are produced inside the BLR (at distances $<$0.1--1 pc). On the contrary, if the emission occurs at larger distances, the dominating population of target photons will be that from the torus. In turn, the location of the $\\gamma$--ray production and nature of the target photons determines some of the quantities derived from the radiative models, such as electron energies and cooling times, jet power, magnetic field intensity. \\begin{figure*} \\vskip -2.7 true cm \\centerline{ \\hskip 0.7 truecm \\psfig{file=combo_1yr_noinset.ps,height=17.6cm,width=17.6cm} } \\vskip -3 true cm \\caption{Light curve of PKS 1510--089 (upper panel) and 3C 454.3 (lower panel) from 2008 August 4 (MJD 54682) to 2010 January 31 (MJD 55228) in bins of 1 day. Vertical yellow stripes show the periods of large flux studied in detail in this work.} \\label{1.5anni} \\end{figure*} The overall spectral energy distributions of FSRQs can be well reproduced assuming an emission region located at 300--1000 Schwarzschild radii from the central black hole (e.g. Ghisellini et al. 2010). However, recently, several authors (Sikora et al. 2008; Larionov et al. 2008; Marscher et al. 2008, 2010) argued that the bulk of the emission, especially during large outbursts, is produced at larger distances, even at distances of the order of 10--20 pc from the central black hole, at the expected location of a reconfinement shock (e.g. Sokolov et al. 2004). The main arguments advanced to support this conclusion come from observations in the radio band coupled with the observed peculiar behaviour of the polarisation angle in the optical. This led Marscher et al. (2008) to suggest a general scenario in which blobs (or knots) of material ejected from the central region are forced by the magnetic field to follow an helical path, accounting for the observed rotation of the polarisation angle in the optical. These knots are opaque in the radio band until they reach large distances. The transition to transparency is marked in VLBI maps by the passage of the compact radio core, after which knots become visible and their trajectories can be directly traced. The passage of the knots from the core, interpreted as the location of a standing conical shock, is marked by the huge flares at all wavelengths, triggered by the compression of the plasma in the shock. This scenario has important consequences for the variability of the emission: since the emission region is located at large distances from the central engine, their size is probably large, even if a very small jet opening angle ($\\theta _{\\rm jet}<1$ deg) is assumed. Quantitatively, the light crossing time of the source put a lower limit on the variability timescale expressed as: $t_{\\rm var}>\\theta _{\\rm jet} d (1+z)/c\\delta$, where $d$ is the distance of the emission region from the central engine and $\\delta$ the relativistic Doppler factor. Assuming $d=15$ pc, $\\theta _{\\rm jet}=0.1$ deg and $\\delta=20$ we find $t_{\\rm var}>1.5 (1+z)$ days. Therefore, variability at timescales below 1 day, especially at $\\gamma$--ray energies, would be rather difficult to accommodate in this scheme. The Large Area Telescope (LAT) onboard {\\it Fermi} (Atwood at al. 2009), with its continuous monitoring of the sky, is the ideal instrument to investigate the possible existence of rapid (timescale $t_{\\rm var}<1$ day) $\\gamma$--ray variability in FSRQs. Indeed, variability on timescale of 12 hours has been already reported for the FSRQs PKS 1454--354 ($z=1.424$; Abdo et al. 2009) and PKS 1502+105 ($z=1.83$; Abdo et al. 2010), which underwent major flares in August--September 2008, reaching fluxes of the order of a few times $10^{-6}$ ph cm$^{-2}$ s$^{-1}$. To further probe variability at short timescales we consider here the LAT light curves of two well studied FSRQs, 3C 454.3 ($z=0.859$) and PKS 1510--089 ($z=0.360$), for which the 1--day averaged flux above 100 MeV reached or even exceeded in few occasions $10^{-5}$ ph cm$^{-2}$ s$^{-1}$. One of these events for PKS 1510--089 has been recently interpreted in the framework discussed above (Marscher et al. 2010). Such a large flux allows us to investigate variations occurring on timescales of the order of 3--6 hours, thus providing strong constraints on the theoretical scenarios. ", "conclusions": "We have analysed the LAT light curves of the two FSRQs 3C 454.3 and 1510--089 from August 2008 to January 2010. In this period the two sources displayed several $\\gamma$--ray flares with fluxes approaching or even exceeding $10^{-5}$ ph cm$^{-2}$ s$^{-1}$. We have characterised the variability of the two sources deriving the distribution of the fluxes. We found that in both cases this distribution is well approximated by an increasing power low at low fluxes ($N_F\\propto F^{1.5}$) up to a peak above which the probability rapidly decreases. These cases are therefore different from those well described by log--normal distributions recently found in both accreting systems and blazars in the X--ray band. We have derived light curves around the epochs of the largest flares with bins of 3 and 6 hours, finding that there are several cases in which the flux varies on these timescales, with variations even larger than a factor of 2. This implies a correspondingly very compact emitting region that is difficult to be explained by models that assume that most of the $\\gamma$-ray emission is produced in regions far away from the central engine, at distances of tens of parsecs. The same conclusion, i.e. that the bulk of the emission is produced close to the central engine, was also independently reached by Kovalev at al. (2009), using a different approach." }, "1003/1003.1193_arXiv.txt": { "abstract": "A diversity of resonance configurations may be formed under different migration of two giant planets. And the researchers show that the HD 128311 and HD 73526 planetary systems are involved in a 2:1 mean motion resonance but not in apsidal corotation, because one of the resonance argument circulates over the dynamical evolution. In this paper, we investigate potential mechanisms to form the 2:1 librating-circulating resonance configuration. In the late stage of planetary formation, scattering or colliding among planetesimals and planetary embryos can frequently occur. Hence, in our model, we consider a planetary configuration of two giants together with few terrestrial planets. We find that both colliding or scattering events at very early stage of dynamical evolution can influence the configurations trapped into resonance. A planet-planet scattering of a moderate terrestrial planet, or multiple scattering of smaller planets in a crowded planetary system can change the resonant configuration. In addition, collision or merging can alter the masses and location of the giant planets, which also play an important role in shaping the resonant configuration during the dynamical evolution. In this sense, the librating-circulating resonance configuration is more likely to form by a hybrid mechanism of scattering and collision. ", "introduction": "To date, more than 40 multiple planetary system have been detected beyond our solar system. The commensurability of orbital periods is very ubiquitous in the extrasolar planetary systems. At present, four resonant pair of planets (GJ 876, HD 82943, HD 128311, HD 73526) are reported to be trapped in 2:1 mean motion resonance \\citep{mar01, may04, vog05, tin06}. In recent years, numerous researchers have extensively investigated the dynamics and origin of the 2:1 resonance in the planetary systems \\citep{Goz01, lee02, Haj02, Haj06, Haj03, Ji03a, Ji03b, lee04, Kle04, Bea03, Bea06, Gay08, lee09, voy09}. The two resonance variables for 2:1 resonance, $\\theta_1=\\lambda_1-2 \\lambda_2+\\varpi_1$, $\\theta_2=\\lambda_1-2 \\lambda_2+\\varpi_2$ (where $\\lambda$, $\\varpi$ are the mean longitude and the longitude of periapse respectively, the subscripts 1, 2 denote the inner and outer planets), are categorized to librate about: (1)0$^\\circ$ in symmetric configuration, (2) both 0$^\\circ$ and 180$^\\circ$ respectively, in the so-called antisymmetric configuration, and (3) other degrees different from 0$^\\circ$ or 180$^\\circ$ in asymmetric configuration \\citep{Haj02, Ji03a, Ji03b, Ji03c, Bea03, Bea06, lee04}. The GJ 876 was revealed to be the first 2:1 resonant system \\citep{mar01} near an M dwarf star. The GJ 876 system is in apsidal corotation where the mean motion resonance variables, $\\theta_1$ and $\\theta_2$ librate about $0^\\circ$ with quite slight amplitudes. On the origin of mean motion resonances in the system, a formation scenario is that they were assembled by migration of planets. In the formation of giant planets, if two planets are massive enough to open gaps and not far away from each other in the disk, the material between the region of them can be rapidly cleared off. And then, the dissipation of the stuff outside the outer planet and inside the inner planet may still force two planets approach each other. Any process that makes two bodies approach each other, which are originally separated appropriately, could result in mean motion orbital resonance \\citep{lee02, Kle05, mas06}. Lee and collaborators explored the origin and diversity of the 2:1 mean motion resonance \\citep{lee02, lee04, Kle05, lee06}. They set the two planets on coplanar, circular orbits with, $a_2/a_1\\sim2/1$. The outer planet can migrate toward the center star at a reasonable velocity (\\citealt{war97}; see also \\citealt{lee02} and references therein), \\begin{eqnarray} \\left|\\frac{\\dot{a}}{a}\\right|\\approx\\frac{3\\nu}{2a^2} =\\frac{3}{2}\\alpha\\left(\\frac{H}{a}\\right)^2\\Omega =9.4\\times10^{-5}\\left(\\frac{\\alpha}{4\\times10^{-3}}\\right)\\nonumber\\\\ \\times\\left(\\frac{H/a}{0.05}\\right)^2 \\left(\\frac{M_\\star}{1M_{\\odot}}\\right)^{1/2} \\left(\\frac{a}{1\\rmn{AU}}\\right)^{-3/2}~\\rmn{yr}^{-1} \\end{eqnarray} where $\\nu=\\alpha H^2\\Omega$, is the kinematic viscosity expressed using the Shakura-Sunyaev $\\alpha$-prescription, while $H$, $\\Omega=2\\pi/P$ are the height scale of the protoplanetary disk and the orbital angular velocity at $a$ (where $P\\approx 2\\pi a^{3/2}(GM_\\star)^{-1/2}$), and $M_\\star$ is the mass of the center star. In $(1)$, the values of $\\alpha = 4\\times 10^{-3}$ and $H/a=0.05$ are typical in classical models of protoplanetary disks. If we adopt $M_\\star=1M_\\odot$, the migration velocity of the planet with $a=1\\rmn{AU}$, is $9.4\\times 10^{-5}~\\rmn{yr}^{-1}$ by $(1)$. They find that such migration of the planets can result in resonance capture with the eccentricities of the planets growing quickly. If the eccentricity damping induced by planet-disk interaction is considered, the system may remain stable over secular timescale after the resonance capture has happened. In such circumstances, the eccentricities reach nearly constant values. The investigations imply that the 2:1 symmetric resonant configurations may be easily established under the forced migration due to planet-disk interaction. The numerical explorations and theoretical analysis show that the 2:1 resonance planetary configurations could be quite diverse due to different migration in a slightly eccentric disk \\citep{lee04, Kle05}. For example, it is shown that the orbital solutions of the resonant pairs derived from Keplerian fit, could result in unstable behaviors \\citep{vog05, tin06} over the timescale of several thousand years in the HD 128311 and HD 73526 systems. The best-fit dynamical orbital solutions are given in Table 1, which each system can remain stable over $10^4$ yr. Indeed, both of the systems are stabilized by the 2:1 resonance, however, with a librating-circulating resonance configuration. Planetary configurations such as listed in Table 1 are not observed to lead to the convergent migration scenario mentioned by Lee et al. (2002, 2004, 2006). However, \\citet{tin06} mentioned several possibilities of the origin of librating-circulating resonance configuration. The resonance configurations may be formed either through rapid migration or migration with initial planetary eccentricities or via a dynamical scattering event. \\citet{san06} provide a mixed evolutionary scenario for such resonance configuration by combining an adiabatic migration progress and a sudden gravitational perturbation. In their study, firstly, they show the planets could be captured into mean motion resonances via inward migration. And then, they describe a scenario of a sudden stop of inward migration after the resonance capture happened, which is supported by observations of young protoplanetary disk \\citep{cal05, dal05, mas06}. Finally, a small body ($\\sim 10M_\\oplus$) is scattered during the follow-up dynamical evolution of the system. Therefore, the resonance configuration of the system is turned into a librating-circulating stage, and the eccentricities of the planet oscillate with large amplitudes. The results of the numerical studies are consistent with the dynamical behaviors expressed using solutions listed in Table 1 of HD 128311 and HD 73526 planetary systems \\citep{san06, san07}. \\begin{table} \\centering \\begin{minipage}{100mm} \\caption{Dynamical orbital fits for two systems.} \\begin{tabular}{@{}ccccccc} \\hline & \\multicolumn{2}{c}{HD 128311\\footnote{The orbital parameters reference to \\citet{san06}}} & \\multicolumn{2}{c}{HD 73526\\footnote{\\citet{tin06}}} & \\multicolumn{2}{c}{HD 73526\\footnote{\\citet{san07}}}\\\\ Parameter & b & c & b & c & b & c \\\\ \\hline $\\rmn{Mass}(M_J)$ & 1.56 & 3.08 & 2.9 & 2.5 & 2.415 & 2.55 \\\\ $a(\\rmn{AU})$...... & 1.109 & 1.735 & 0.66 & 1.05 & 0.659 & 1.045 \\\\ $e$................ & 0.38 & 0.21 & 0.19 & 0.14 & 0.26 & 0.1107 \\\\ $\\varpi(deg)$...... & 80.1 & 21.6 & 203 & 13 & 202.9 & 253.7 \\\\ $M(deg)$...... & 257.6 & 166.0 & 86 & 82 & 70.7 & 170.7 \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table} The terrestrial planets are rocky planets, with masses ranging in $1-10M_\\oplus$. The observations show that there are several systems containing several super Earths in close-in orbits, e.g., Gl 581 and HD 40307 (see http://exoplanet.eu/), and recently the nearby solar-like star 61 Virginis was reported to harbor two Neptune-like planets at about 0.22 AU and 0.48 AU \\footnote{The terrestrial planets with masses as small as 2 Earth mass and a period of several months orbiting a K dwarf may be detected in the Habitable zones, by observing 60 groups of 10 nights with HARPS on the ESO 3.6 m telescope (M. Mayor 2009, private communication). Ground-based microlensing is sensitive to Earth-mass planets orbiting at $~1$ AU (D. Bennett 2009, private communication) around late-M stars, while MPF is sensitive to planets down to 0.1 Earth mass in the Habitable zone (near 1 AU) of solar-type, G and K stars. The future E-VLTs and space missions (e.g., GAIA and SIM) will be also hopeful to discover multiple Earth-like planets residing in the terrestrial region in a planetary system.}, and a super Earth \\citep{vog09}. This indicates that terrestrial or Neptune-like planets are very common to survive in the late stage of planet formation. On the other hand, in the numerical studies on the terrestrial planet formation, $2-4$ terrestrial planets are formed with moderate eccentricities \\citep{cha01, ray04, ray05, ray06, Zha09}. It is worthy to pay attention to direct or indirect influence upon the resonance configuration in the very beginning of evolution. At that time, giant planets have been already formed and stopped migrating, while the terrestrial planets may have been created but unstable. If the periods of the giants are commensurable, the terrestrial planets with rather moderate eccentricities could take effect on the dynamical evolution of a system. For example, the earlier works \\citep{san06, san07} show that librating-circulating configurations of the 2:1 mean motion resonance can be formed by combining migration processes and sudden perturbations. In this work, we investigate potential mechanisms of making a librating-circulating resonance configuration. And a novel hybrid mechanism is proposed to explain this issue. In the model, we consider the planetary configuration of the two giants together with few terrestrial planets. We show that a scattering of a moderate mass planet or multiple continuous scattering of the terrestrial bodies can form librating-circulating configuration. Moreover, we also find that the collision and merging can play a vital role in tuning to a librating-circulating mode of the giants. In Section 2, we reproduce a scenario of migration and capture into a 2:1 resonance. In Section 3, we introduce the model and initial setup for our study. We present the main results in Section 4. We conclude the outcomes in Section 5. ", "conclusions": "In this work, we have investigated potential mechanisms to shape 2:1 librating-circulating resonance configuration by considering a planetary system of two giants accompanying with few terrestrial planets with coplanar orbits in the late stage of planet formation. In the model, the system is considered to be gas-free, and two giants with commensurable orbits stop migrating and several terrestrial planets have formed inside the inner giant. The configuration is much closer to real planetary formation scenario according to numerical simulations \\citep{cha01, ray04, Zha09}. In conclusion, we summarize the main results as follows. In the late stage of planetary formation, planet-planet scattering or colliding among planetesimals and embryos can frequently occur. Our results show that not only a single planetary scattering of a terrestrial planet with a moderate mass can result in the librating-circulating configuration, but several continuous planetary scattering with rather smaller terrestrial masses is also at work. Additionally, if two giant planets are initially engaged in a 2:1 symmetric resonance and their eccentricities oscillate with large amplitudes, the collisions arising from the giants and other small bodies may change librating amplitudes of the resonance angles during the evolution. If the apsidal corotation is disintegrated, the configuration may turn into a librating-circulating status. Obviously, the more mass of a perturbing terrestrial body may have much greater influence on the commensurable giant planets. In most simulations, colliding and scattering events can be found and they can increase or decrease the fluctuation in the amplitude of the resonant angles, even dramatically destroy the whole system. In a word, the librating-circulating configuration of mean motion resonance is likely to generate by a mixed mechanism of colliding and scattering. In addition, the librating-circulating configuration could be generated through long-term evolution of planetary formation. A librating-circulating configuration trapped in 3:2 mean motion resonance \\citep{Zha09} is unveiled during the formation of the terrestrial planets, which is similar to the newly-discovered planetary system HD 45346 \\citep{cor09}. They show that the two terrestrial planets were formed within 50 Myr (see their Figure 6 for details), and after that time the frequent orbital crossings of them and their interaction with the inner giant planet may finally lead to a capture of a 3:2 resonance. Such resonant configuration is believed to hold over hundreds of Myrs. However, more resonant configurations for less massive planets are expected to reveal by future higher accuracy space-based projects in search for terrestrial planets (such as TPF, Darwin, SIM). The innovative findings will encourage one to more carefully study their dynamics and origin." }, "1003/1003.5912_arXiv.txt": { "abstract": "{We show that the thermal relic abundance of dark matter can be affected by a new type of reaction: semi-annihilation. Semi-annihilation takes the schematic form $\\psi_i \\psi_j \\rightarrow \\psi_k \\phi$, where $\\psi_i$ are stable dark matter particles and $\\phi$ is an unstable state. Such reactions are generically present when dark matter is composed of more than one species with ``flavor'' and/or ``baryon'' symmetries. We give a complete set of coupled Boltzmann equations in the presence of semi-annihilations, and study two toy models featuring this process. Semi-annihilation leads to non-trivial dark matter dynamics in the early universe, often dominating over ordinary annihilation in determining the relic abundance. This process also has important implications for indirect detection experiments, by enriching the final state spectrum from dark matter (semi-)annihilation in the Milky Way.} \\begin{document} ", "introduction": "\\label{sec:intro} The existence of dark matter is one of the best motivations for physics beyond the standard model (SM). Evidence for dark matter has accumulated at vastly different length scales---from galactic scales and clusters of galaxies to global scales of hundreds of megaparsecs \\cite{Jungman:1995df,Bergstrom:2000pn,Bertone:2004pz}. However, all of these observations infer the existence of dark matter through its gravitational effects alone. In particular, they do not tell us the nature, origin, or composition of this important component of our universe, which is not explained by any SM degree of freedom. A particularly well-motivated class of dark matter candidates are so-called Weakly Interacting Massive Particles, or WIMPs, whose abundance is determined through thermal freeze-out. In the Lee-Weinberg scenario \\cite{Lee:1977ua}, WIMPs $\\psi$ thermalize in the early universe through the annihilation reaction $\\psi \\bar{\\psi} \\rightarrow \\phi \\phi'$, where $\\phi$ and $\\phi'$ are SM degrees of freedom, until their interactions freeze out. A standard relic abundance calculation \\cite{Krauss:1983ik,Scherrer:1985zt,Kolb:1985nn,Srednicki:1988ce,Gondolo:1990dk,Kolb:1990vq,Bottino:1993zx} shows that the dark matter mass density today depends only logarithmically on the relic particle mass and scales inversely with the total annihilation cross section at freeze-out, $\\Omega_{{\\rm DM}} h^2 \\propto \\langle \\sigma v\\rangle^{-1}$. It is remarkable that for a dark matter mass between $10 \\GeV-10 \\TeV$ and an electroweak annihilation cross section, one gets a relic abundance in the ballpark to account for dark matter. This fact is often referred to as the ``WIMP miracle'', and in the context of the gauge hierarchy problem, it is suggestive that the same particles one might introduce to stabilize the Fermi scale could also explain the dark matter in the universe. Since the WIMP paradigm is one of the best explanations for dark matter and since thermal freeze-out is so predictive, it is important to know how to correctly compute the dark matter thermal relic density. This is particularly important for models where the relic computation cannot be reduced to the Lee-Weinberg scenario. As we will argue, the thermal abundance of dark matter can be dramatically affected by the presence of a new dark matter interaction, which we call ``semi-annihilation''. Semi-annihilation occurs when dark matter is stabilized by a larger symmetry than just $Z_2$. In the simplest case with just one dark matter species $\\psi$, there can be an additional allowed reaction \\begin{equation} \\psi \\psi \\rightarrow \\psi \\phi \\end{equation} which preserves a $Z_3$ symmetry. Here, $\\phi$ is a SM state or a new particle which decays to the SM. We see that unlike ordinary annihilation where the total dark matter number changes by two units, in semi-annihilation the total dark matter number changes by only one unit. More generally, dark matter can be composed of more than one stable component $\\psi_i$, and these relic particles can have non-trivial mutual interactions. In this case, a more general semi-annihilation reaction is possible, \\begin{equation} \\psi_i \\psi_j \\rightarrow \\psi_k \\phi , \\end{equation} which often occurs if dark matter is stabilized by ``baryon'' and/or ``flavor'' symmetries, as in QCD-like theories. Such reactions are also allowed in models where dark matter is composed of non-Abelian gauge bosons~\\cite{Hambye:2008bq,Hambye:2009fg,Arina:2009uq}. In this paper, we study explicit examples of such models and find that to correctly compute the relic abundance, semi-annihilation must be included. In fact, semi-annihilation can dominate over standard annihilation for some regions of parameter space. To understand how semi-annihilation fits into the WIMP paradigm, consider a more general framework for dark matter interactions. Start by assuming the existence of a new dark sector composed of $N$ particles $\\psi_i$ which can be either stable or unstable. The possible reactions involving $\\psi_i$ which can take place in the early universe are \\begin{equation} a) \\, \\psi_i \\psi_j \\rightarrow \\phi \\phi^{\\prime},\\;\\;\\; b) \\, \\psi_i \\phi \\rightarrow \\psi_j \\phi^{\\prime},\\;\\;\\; c) \\, \\psi_i \\psi_j \\rightarrow \\psi_k \\phi,\\;\\;\\; d) \\, \\psi_i \\psi_j \\rightarrow \\psi_k \\psi_m,\\;\\;\\; e) \\, \\psi_i \\rightarrow \\psi_j \\phi, \\label{eq:abcde} \\end{equation} where again $\\phi$ and $\\phi^{\\prime}$ are light degrees of freedom in thermal equilibrium with the SM. Different dark matter thermal freeze-out scenarios depend on which of the reactions $a)$--$e)$ are active, and we summarize these main possibilities in \\Tab{tab:reactions}. \\TABLE[t]{ \\begin{tabular}{c|c|c|c|c|} & $\\psi_i \\psi_j \\rightarrow \\phi \\phi^{\\prime}$ & $\\psi_i \\phi \\rightarrow \\psi_j \\phi^{\\prime}$ & $\\psi_i \\psi_j \\rightarrow \\psi_k \\phi$ & $\\psi_i \\rightarrow \\psi_j \\phi$ \\\\ \\hline Lee-Weinberg & $\\checkmark(i=j)$ & $\\checkmark(i=j)$ & $\\times$ & $\\times$ \\\\ Co-annihilation & $\\checkmark$ & $\\checkmark$ & $\\times$ & $\\checkmark$\\\\ Multi-component & $\\checkmark(i=j)$ & $\\checkmark(i=j)$ & $\\times$ & $\\times$\\\\ Semi-annihilation & $\\checkmark(i=j)$ & $\\checkmark(i=j)$ & $\\checkmark$ & $\\times$ \\\\ \\hline \\end{tabular} \\label{tab:reactions} \\caption{Different dark matter freeze-out scenarios and allowed reactions. We do not include the reaction $\\psi_i \\psi_j \\rightarrow \\psi_k \\psi_m$ in this table since it is always present, and we explicitly indicate when only the diagonal $(i= j)$ contribution is allowed.} } \\begin{description} \\item{\\bf Lee-Weinberg}: The simplest case is the Lee-Weinberg scenario \\cite{Lee:1977ua}, where $N=1$ and the only allowed reactions are $a)$, $b)$, and $d)$, with $i=j$. Chemical freeze-out is determined by the annihilation reaction $a)$, and kinetic freeze-out is determined by $b)$. Reaction $d)$ plays no role in the thermal relic computation. \\item{\\bf Co-annihilation}: A slight variation of the standard case is when $N>1$ but there is only one stable dark matter species. Here the relevant reactions are $a)$, $b)$, $d)$, and $e)$. This is the case in the minimal supersymmetric standard model with $R$-parity, where all heavier particles are unstable and eventually decay to the lightest one via reaction $e)$. In principle, this co-annihilation case involves a system of $N$ coupled Boltzmann equations, but as we will review, as long as reactions of type $b)$ with $i \\not= j$ are effective at freeze-out, it is possible to compute the relic density via standard methods \\cite{Griest:1990kh}. \\item{\\bf Decoupled Multi-Component}: These first two examples assume dark matter to be composed of a single particle, but more generally, dark matter could be composed of more than one stable component. Many such multi-component dark matter models have been proposed (see e.g.\\ \\cite{Boehm:2003ha,Ma:2006uv,Hur:2007ur,Hur:2007ur,Feng:2008ya,Fairbairn:2008fb,SungCheon:2008ts,Zurek:2008qg,Morrissey:2009ur}). The standard approach in multi-component models is to assume that each particle thermalizes independently of the others, thus the total dark matter density today is $\\Omega_{\\rm DM} = \\sum_i \\Omega_i$ where the sum runs over all the thermal relics. Said in the language of \\Eq{eq:abcde}, only the diagonal ($i=j$) reactions of type $a)$ and $b)$ are present, and reactions $c)$ and $e)$ are forbidden. Reaction $d)$ may or may not be present in such models, and is relevant for calculating the relic density \\cite{SungCheon:2008ts}. \\item{\\bf Semi-annihilation}: The focus of this paper is on the reaction of type $c)$, which to our knowledge first appeared in \\Ref{Hambye:2008bq,Hambye:2009fg,Arina:2009uq}.\\footnote{In the models considered in \\Ref{Hambye:2008bq,Hambye:2009fg,Arina:2009uq}, a custodial symmetry effectively reduced the relic abundance computation to a single particle system. Here, we consider semi-annihilation in multi-component dark matter models where there is no such simplification.} As long as the triangle inequality $m_k < m_i+m_j$ is satisfied (as well as its crossed versions), the semi-annihilation reaction $\\psi_i \\psi_j \\rightarrow \\psi_k \\phi$ can take place without making any relic particle unstable. Semi-annihilation implies that the relic particles have non-trivial mutual interactions, and typically both reactions of type $c)$ and $d)$ are important in determining the dark matter relic abundance. For simplicity, we focus on the case where each particle in the dark sector is absolutely stable, in which case reaction $e)$ is forbidden. In addition, only the diagonal reactions of type $a)$ and $b)$ are allowed, since off-diagonal contributions would imply that the heavier particle would be unstable by crossing symmetry. As we will see, the absence of off-diagonal reactions of type $b)$ makes these models more difficult to study than the standard cases. \\end{description} The remainder of this paper is structured as follows. In \\Sec{sec:Z3} we present a toy model where dark matter is composed of a single component but semi-annihilations are present and can dominate dark matter production in the early universe. In \\Sec{sec:semimulti} we introduce example multi-component models where semi-annihilation reactions are present, and we give a complete set of Boltzmann equations which can be solved to obtain the relic density today. In particular, we argue that a semi-analytical solution analogous to the Lee-Weinberg scenario or co-annihilation is not in general possible, and thus the equations have to be solved numerically. We present a minimal multi-component toy model with semi-annihilations in \\Sec{sec:bbchi} and numerically study the effects of semi-annihilation on the relic abundance. We explore the effects of semi-annihilations on indirect detection experiments in \\Sec{sec:Indirect}, and conclude in \\Sec{sec:Conclusions}. Computational details, as well as an example supersymmetric QCD model with $N_f = N_c+1$, are given in the appendices. ", "conclusions": "\\label{sec:Conclusions} The presence of dark matter in our universe motivates us to look for physics beyond the SM, and understanding its origin and composition is one of the great open questions in particle physics and cosmology. Particle physics models explaining the origin of the Fermi scale often contain stable massive particles called WIMPs. Assuming WIMPs are thermally produced in the early universe, these particles have the right annihilation cross section to give the observed relic abundance, $\\Omega_{{\\rm DM}} h^2 \\simeq 0.1$, making the WIMP paradigm one of the best solutions to account for dark matter. In this work, we relaxed some assumptions about the symmetry structure of WIMP interactions by allowing relic particles to ``semi-annihilate''. We saw that the semi-annihilation process \\be \\psi_i \\psi_j \\rightarrow \\psi_k \\phi \\ee can have a considerable effect on the dark matter relic abundance. Such processes are present when dark matter has a larger stabilizing symmetry than just $Z_2$, and we showed explicit examples where semi-annihilation is present: a single species dark matter model with a $Z_3$ symmetry and a multiple species dark matter model with ``baryon'' and ``flavor'' symmetries. Along with the standard WIMP scenarios summarized in \\Tab{tab:reactions}, semi-annihilation does occur in realistic particle physics scenarios. Indeed, as sketched in \\App{app:SUSYQCD}, the simple case of a supersymmetric QCD theory with $N_f = N_c+1$ generically involves semi-annihilation. We saw that when dark matter is composed of just one species, the effect of semi-annihilation on the relic density can be derived through a standard freeze-out calculation \\cite{Lee:1977ua}. Semi-annihilation can in fact dominate over ordinary annihilation for some regions of parameter space. However, such single species models are not representation of generic semi-annihilating dark matter models, which usually involve more than one stable dark matter component. In such cases, a semi-analytical solution for the relic density is simply not possible, and one must resort to numerically solving the complete set of coupled Boltzmann equations given in \\Eq{eq:Boltzsystem}. In particular, the simplifying assumptions made when analyzing co-annihilating models \\cite{Griest:1990kh} are not applicable for semi-annihilation. As a consequence, the dark matter dynamics in the presence of semi-annihilation are far more varied than in decoupled multi-component scenarios. While the inclusion of semi-annihilation does not give any new contributions to the direct detection of dark matter, it does have interesting implications for indirect detection experiments. We studied the injection spectrum of light particles in the two toy models we studied, where dark matter (semi-)annihilates through a $\\phi$ ``portal'' which subsequently decays to SM states. The overall integrated flux from (semi-)annihilates in the Milky Way halo is not very affected by semi-annihilation (in the case of one relic particle, is not affected at all). However, the final state spectrum in semi-annihilating models is far richer than in standards scenarios because of the differing kinematics between semi-annihilation and ordinary annihilation. In the language of \\Ref{Griest:1990kh}, semi-annihilation is in some ways the ``fourth exception'' in the calculation of dark matter thermal relic abundances. We find it intriguing that unlike the traditional three exceptions (co-annihilation, annihilation below a mass threshold, and annihilation near a pole), this fourth exception not only affects dark matter interactions in the early universe, but also leaves an imprint today via the indirect detection spectrum. We expect that there are a wide variety of multi-component dark matter models with species changing interactions, motivating further studies of the semi-annihilation process." }, "1003/1003.1133.txt": { "abstract": "%arXiv abstract limit: 24 lines of 80 chars %2345678901234567890123456789012345678901234567890123456789012345678901234567890 Some galaxy clusters show diffuse radio emission in the form of giant halos (GHs) on Mpc scales or minihalos (MHs) on smaller scales. Comparing Very Large Array and \\emph{XMM-Newton} radial profiles of several such clusters, we find a universal linear correlation between radio and X-ray surface brightness, valid in both types of halos. It implies a halo central emissivity $\\nu j_\\nu = 10^{-31.4\\pm0.2} (n/10^{-2}\\cm^{-3})^2 (T/T_0)^{0.2\\pm0.5} \\erg\\se^{-1}\\cm^{-3}$, where $T$ and $T_0$ are the local and central temperatures, and $n$ is the electron number density. We argue that the tight correlation and the scaling of $j_\\nu$, combined with morphological and spectral evidence, indicate that both GHs and MHs arise from secondary electrons and positrons, produced in cosmic-ray ion (CRI) collisions with a strongly magnetized, $B\\gtrsim 3\\muG$ intracluster gas. When the magnetic energy density drops below that of the microwave background, the radio emission weakens considerably, producing halos with a clumpy morphology (\\eg RXC\\,J2003.5\u0096-2323 and A2255) or a distinct radial break. We thus measure a magnetic field $B=3\\muG$ at a radius $r \\simeq 110 \\kpc$ in A2029 and $r\\simeq 50\\kpc$ in Perseus. The spectrum of secondaries, produced from hadronic collisions of $\\sim20\\GeV$ {\\CRIs}, reflects the energy dependence of the collision cross section. We use the observed spectra of halos, in particular where they steepen with increasing radius or frequency, to \\emph{(i)} measure $B\\simeq 10(\\nu/700\\MHz)\\muG$, with $\\nu$ the spectral break frequency; \\emph{(ii)} identify a correlation between the average spectrum and the central magnetic field; and \\emph{(iii)} infer a {\\CRI} spectral index $s\\lesssim -2.7$ and energy fraction $\\xi_p\\sim 10^{-3.6\\pm0.2}$ at particle energies above 10 GeV. Our results favor a model where {\\CRIs} diffuse away from their sources (which are probably supernovae, according to a preliminary correlation with star formation), whereas the magnetic fields are generated by mergers in GHs and by core sloshing in MHs. ", "introduction": "Giant halos (GHs) appear as diffuse radio emission on $\\sim\\mbox{Mpc}$ scales in merging galaxy clusters \\cite[for a review, see][]{FerettiGiovannini08}. GHs were identified in about a quarter of all clusters with X-ray luminosities $L_X>5\\times 10^{44}\\erg \\se^{-1}$ at redshifts $0.2B_{cmb}$ magnetic fields without radio detection at the level given by Eq.~(\\ref{eq:local_correlation}) would rule out the model, unless the {\\CRI} energy fraction is exceptionally low (\\eg due to a low level of star formation; see \\S\\ref{sec:CRI_Origin}). The model can be tested by checking if the different magnetic field estimates it produces in a given halo are self consistent, and agree with independent measurements. In the model, GH clusters are associated with strong, $B\\gtrsim B_{cmb}\\sim 3\\muG$ magnetic fields, whereas weaker, $BE_p) & \\equiv & \\frac{u_{p}(>E_p)}{u_{th}} \\simeq 10^{-3.6\\pm0.2}\\left(\\frac{E_p}{10\\GeV}\\right)^{-0.7} \\fin \\,\\,\\,\\, \\end{eqnarray} Note that with the uncertain and possibly contaminated radio spectra presently available, a steeper {\\CRP} spectrum with $s\\simeq -3$ is possible. The model would be challenged if the uncontaminated spectrum of a substantial halo population turns out to be much steeper than $\\alpha=-1.5$, unless the corresponding steep {\\CRP} spectrum can be explained. The {\\CRP} distribution in Eq.~(\\ref{eq:uCR}) resembles (but has an energy fraction a few $100$ times smaller than) the {\\CRP} distribution found in the solar vicinity above $1\\GeV/$nucleon. This distribution could originate from sources that inject roughly equal energy per decade of {\\CRP} energy ($s_0\\simeq -2.2$), such as supernovae (SNe), if energy-dependent diffusion is significant in the inner halo regions. For example, a simple estimate of {\\CRI} scattering off magnetic irregularities with a Kolmogorov power spectrum yields a diffusion coefficient $D\\simeq 10^{30}(E_p/\\mbox{GeV})^{1/3}(B/\\mu\\mbox{G})^{-1/3}\\cm^2\\se^{-1}$ \\citep{VolkEtAl96}. This implies {\\CRI} diffusion over $\\sim 0.5\\Mpc$ during a Hubble time, and a steepening by $\\Delta s=-1/2$. More substantial steepening is possible if the diffusion function has a stronger energy dependence. For example, the diffusion function is often assumed to scale as $D\\propto E_p^{1/2}$, which could lead to a $\\Delta s=-3/4$ steepening in the {\\CRI} spectrum. The {\\CRI} output of SNe can be crudely estimated \\citep{VolkEtAl96} if we assume that a fraction $f_{\\mbox{\\scriptsize{II}}}$ of the cluster's $Z=0.3Z_{0.3}$ solar metallicity is seeded by Type II SNe, which on average produce $0.1M_\\odot M_{Fe,0.1}$ of iron and deposit a fraction $\\xi=0.3\\xi_{0.3}$ of the $10^{51}E_{51}\\erg$ explosion energy in $E_p>10\\GeV$ {\\CRIs}, \\begin{eqnarray} \\label{Eq:SNeCRI} \\xi_{p}^{SN} \\simeq 0.03 f_{\\mbox{\\scriptsize{II}}} Z_{0.3} E_{51} M_{Fe,0.1}^{-1} \\xi_{0.3} \\fin \\end{eqnarray} This can reproduce Eq.~(\\ref{eq:uCR}) if over the cluster's lifetime, the {\\CRIs} diffuse to distances a few times larger than the radius $R_\\nu$ of the radio halo. Note that if {\\CRI} diffusion is entirely absent, the {\\CRIs} accelerated in SNe would be confined to the cluster, and adiabatic losses could only lower their energy density to the level of Eq.~(\\ref{eq:uCR}). However, they would then retain their flat, $s\\simeq -2.2$ spectrum. An SNe origin of {\\CRIs} can be tested by examining the correlations between $\\myeta$ and (intensive) tracers of SNe activity among different halos. One possible tracer is the local metallicity measured at $r=0.1r_{500}$, tabulated in Table \\ref{tab:halos}. We chose to use $Z(0.1r_{500})$ because \\emph{(i)} it was measured for all the M09 halos with \\emph{XMM-Newton} profiles in \\citet{SnowdenEtAl08}; \\emph{(ii)} the spatially averaged $Z$ is not meaningful when temperature gradients are large; and \\emph{(iii)} $0.1r_{500}$ lies well within the $Z(r>0.02r_{500})\\propto r^{-0.3}$ decline typically found in both cool and non-cool core clusters \\citep{SandersonEtAl09}. While our sample is statistically small, Perseus, which shows a significantly higher $\\myeta$ than in all the other halos in our sample, also shows a slightly higher $Z(0.1r_{500})$ \\citep{SnowdenEtAl08}. However, due to the large uncertainty in abundance measurements, the elevated metallicity in Perseus is not significant ($<1\\sigma$) with respect to some halos. Moreover, at smaller radii $r\\lesssim 0.025r_{500}$, the metallicity in A2029 appears to be higher than in all other halos, and is significantly ($>3\\sigma$) higher than in Perseus \\citep{SnowdenEtAl08}. Although better metallicity statistics may identify a more significant correlation between $\\myeta$ and $Z$, metallicity is probably not the most useful tracer of the SNe contribution to {\\CRIs} in the halo. Metallicity provides a cumulative measure of SNe activity, tracing the metals released from all past SNe in the cluster. In a model where a significant fraction of the {\\CRIs} have already diffused away from the cluster's center, metallicity would not linearly trace the population of {\\CRIs} residing within the halo, especially in the more compact MHs. It is more appropriate to use an intensive tracer of \\emph{recent} SNe activity, such as the star formation rate (SFR) normalized by the cluster's gas mass $M_g$, or the fraction of star forming galaxies. A correlation between an SNe measure and a {\\CRI} tracer, such as $\\myeta$ or the deviation from the luminosity correlation $\\nu P_\\nu/\\bar{L}_{X[0.1,2.4]}^{1.7}$, may be more useful than the metallicity in establishing or ruling out an SNe origin of halo {\\CRIs}. As seen in Fig.~\\ref{fig:MH_RadialTCorrectedProfiles}, the halos with the highest $\\myeta$ in our sample are the MH in Perseus and the GH in A665, while the halos with the lowest $\\myeta$ are the MH in A2029 and the GH in A773. Interestingly, the literature shows evidence for exceptionally high specific star formation in both Perseus and A665, and for a low specific SFR in A2029. (We found no relevant data for A773.) In Perseus, which has the smallest $M_g$ \\citep[by at least a factor of $4$, see][]{FukazawaEtAl04} and one of the most powerful cooling flows within our sample \\citep[\\eg][]{White00,AllenEtAl02}, there is optical-to-UV evidence for a relatively high SFR \\citep[\\eg][]{BregmanEtAl06, RaffertyEtAl08}. In particular, the cD galaxy NGC1275 in Perseus has a high SFR of $\\sim 30M_\\odot \\yr^{-1}$ \\citep{DixonEtAl96} --- the highest in our sample when normalized by $M_g$. Notice that the central galaxy in A1835 has a higher SFR of $\\sim 100M_\\odot \\yr^{-1}$ --- the highest SFR known in such objects \\citep{PetersonFabian06}. However, $M_g$ is $\\sim 10$ times larger in A1835 with respect to Perseus \\citep{FukazawaEtAl04}. Note that regions containing a high density of cosmic-rays are directly observed in Perseus in the form of X-ray cavities, reaching distances $r>100\\kpc$ (M. Markevitch, private communications). A high specific SFR is also inferred in A665, which was found to be the cluster with the highest dispersion in color magnitude relation (a known tracer of star formation) in a sample of 57 X-ray bright clusters \\citep{LopezCruzEtAl04}. In contrast, A2029 has a low specific SFR, as it was shown to have a SFR $\\sim 70$ times lower than in A1835 \\citep{HicksMushotzky05} while its gas mass is only $\\sim 1.8$ times smaller \\citep{FukazawaEtAl04}. While these trends support an SNe origin of {\\CRIs}, more work is needed in order to quantify their significance and compile a comparative statistical analysis. Note that the combination of a strong correlation of $\\myeta$ with the specific SFR and a poor correlation with metallicity, if established, would directly imply that {\\CRI} diffusion is significant. Indeed, it is difficult to explain the steep, $s\\lesssim -2.7$ spectrum without invoking {\\CRI} diffusion. The above estimates of diffusive steepening assume that most of the {\\CRIs} produced by the sources presently dominating the halo have already escaped beyond it. This is consistent with the typical SFR peak at $z\\sim 1$, and with the above estimates of the halo {\\CRI} abundance and the total {\\CRI} output of SNe (\\cf Eqs.~\\ref{eq:uCR} and \\ref{Eq:SNeCRI}). However, such substantial diffusion would introduce some scatter in the radio--X-ray correlation, depending on the {\\CRI} production history of each cluster. Quantitative estimates of the SNe history of GH clusters, needed to compute this scatter, are beyond the scope of this work. \\subsection{Summary and Conclusions} \\label{sec:Summary} We have shown that the radio-X-ray correlation in GH luminosity (Eq.~\\ref{eq:GH_Correlation}) can be generalized (Eq.~\\ref{eq:GeneralizedLx}) to hold for both GHs and MHs (Fig.~\\ref{fig:MH_CorrectedCorrelation}), by correcting for the halo size. A universal, linear relation between the radio and X-ray surface brightness, $\\eta=10^{-4.4\\pm0.2}$, was presented (Eq.~\\ref{eq:local_correlation} and Fig.~\\ref{fig:MH_RadialTCorrectedProfiles}). This, combined with the radial $\\myeta$ and $T$ profiles, implies a universal radio emissivity $\\nu j_\\nu = 10^{-31.4\\pm0.2} n_{-2}^2 (T/T_0)^{0.2\\pm0.5} \\erg\\se^{-1}\\cm^{-3}$ (Eq.~\\ref{eq:nu_j_nu}) near the center of halos. We argued that these results and their applicability to GHs and MHs alike, strongly support one model for all halos, involving secondary {\\CREs} (injected according to Eq.~\\ref{eq:injection}) and strong magnetic fields with $B\\gtrsim B_{cmb}$, while disfavoring other models (Table \\ref{tab:models}). This model makes useful predictions without requiring additional assumptions or fine tuning. Radio emission rapidly fades in regions where $B$ drops below $B_{cmb}$, producing a distinct radial break (\\eg in A2029 and in Perseus; Fig \\ref{fig:MH_RadialTCorrectedProfiles}) or a clumpy/filamentary radio morphology (\\eg in RXC\\,J2003.5\u0096-2323, A2255, and A2319) that can be used to map $B\\simeq B_{cmb}$ contours. Marginally magnetized regions with $B\\lesssim B_{cmb}\\propto (1+z)^2$ are characterized by relatively high polarization and a steeper radio spectrum; their morphology traces the magnetic evolution and can potentially reveal a temporal signal. We expect a higher incidence rate of such transition regions at higher redshift, while no halos should exist at very high redshift. Another direct consequence of the model is radio spectral steepening with increasing $E_e^2\\propto \\nu/B$ (Eq.~\\ref{eq:Ee}), \\ie with increasing $r$ or $\\nu$ or decreasing T, as indeed is observed. Such steepening, and in particular a $B\\simeq 10(\\nu_e/700\\MHz)\\muG$ spectral break (Fig. \\ref{fig:MHSpectra}), gauges the magnetic field, roughly producing an additional $B$ contour for each radio map frequency. The spectral break could be used to accurately map $B$ throughout the halo, using future radio telescopes such as the Murchison Widefield Array (MWA\\footnote{http://www.mwatelescope.org}), the LOw Frequency ARray (LOFAR\\footnote{http://www.lofar.org}), and SKA. A pressure model can be used to extrapolate $B$ throughout the cluster. This indicates central magnetic fields $B_0$ (Eq.~\\ref{eq:B0}) that exceed $10\\muG$ in most halos (see Table \\ref{tab:halos} for lower limits $B_{0,min}$). A correlation between the average radio spectral index $\\langle \\alpha\\rangle$ and $B_0$, implied by the model, was identified in GH data (Fig.~\\ref{fig:MHSpectraAve}). In our model, any source of strong ($B\\gtrsim B_{cmb}$), persistent magnetic fields in the intracluster medium would have similar properties to radio halos, as long as it does not significantly inject additional cosmic rays. This may include some extended radio galaxies, which were recently found to exhibit properties similar to halos \\citep{RudnickLemmerman09}. Conversely, the universal value of $\\myeta$ we predict for any highly magnetized, uncontaminated region in the intracluster medium provides a powerful test of the model. The spectral steepening of the radio signal, the universality of $\\alpha\\simeq-1$ in the centers of halos, and the correlation between $\\langle \\alpha \\rangle$ and $B_{0,min}$ (Fig. \\ref{fig:MHSpectraAve}), indicate a steep {\\CRI} spectrum, $s\\lesssim-2.7$, and thus favor significant {\\CRI} diffusion. In a diffusion model, the most plausible source of the {\\CRIs} (Eq.~\\ref{eq:uCR}) is SNe (\\eg Eq.~\\ref{Eq:SNeCRI}). We show (in \\S\\ref{sec:CRI_Origin}) preliminary evidence for a correlation between $\\myeta$ and the SFR normalized by the gas mass $M_g$, supporting an SNe {\\CRI} origin. None of these properties is expected in an alternative model (K09), in which the secondary {\\CREs} arise from $\\sim 1\\GeV$ {\\CRPs}, which are accelerated in the cluster's virial shock and advected inward with the flow, thus being compressed to $\\sim 10\\GeV$ energies. Note that the data slightly favors a $j_\\nu\\propto n^2T^0$ scaling within each cluster (see \\S\\ref{sec:Theory}), which is natural if {\\CRIs} originate in SNe, rather than the $j_\\nu\\propto n^2T^1$ behavior anticipated if they are accelerated in the virial shock. Also note that adiabatic compression of {\\CRIs} produced at the virial shock and advected with the gas would lead to a radially increasing, $\\myeta\\propto n^{-1/3}$ profile (due the soft equation of state of relativistic particles), which is not observed (see Fig.~\\ref{fig:MH_RadialTCorrectedProfiles}). We stress that although our model and the model of K09 disagree regarding the origin of {\\CRIs}, the {\\CRI} distribution, the spectral properties of halos, and the role of diffusion, we reach the same conclusions regarding the radio mechanism: emission from secondary {\\CREs} in strong magnetic fields. This conclusion is based on \\emph{(i)} the tight radio-X-ray correlation in total GH luminosity and the GH bimodality (K09); \\emph{(ii)} the tight radio-X-ray correlation in both coincident luminosity and surface brightness, in both GHs and MHs, despite their different physical properties; \\emph{(iii)} the strong magnetic fields inferred from Faraday RMs in MHs and possibly (see \\S\\ref{sec:Magnetization}) also in GHs; \\emph{(iv)} the tightening of the brightness correlation at small radii, away from merger shocks, radio relics, and their associated turbulence; \\emph{(v)} the $j_\\nu\\propto n^2 T^\\mypT$ scaling of the radio emissivity within each halo, where $\\mypT\\lesssim 1$; \\emph{(vi)} the coincidence between MH edges and CFs, manifest as a sharp radial cutoff in $\\myeta$ (\\eg in Perseus); \\emph{(vii)} a power-law radial break where $\\myeta(r\\ll r_b)\\simeq\\constant$ and $\\eta(r\\gg r_b)\\propto B^2$, possibly seen in the MHs in A2029, A1835 and RXJ1347, and in the GH in A2319; \\emph{(viii)} the clumpy/filamentary morphology of some halos, where independent evidence for low magnetization is present; \\emph{(ix)} the spectral steepening and the correlation between $\\langle \\alpha \\rangle$ and $B_0$ (this suggests strong magnetization provided that the {\\CRI} spectrum is steep, $s\\lesssim -2.7$). Each aspect of our model can be tested in the near future. The association between the presence of halos and strong, $B\\gtrsim B_{cmb}$ magnetic fields can be directly tested by comparing the magnetization levels independently estimated in halo and in non-halo clusters, as illustrated in \\S\\ref{sec:Magnetization}. The secondary origin of the halos can be tested if the {\\CRIs} are detected directly through their $\\pi^0$ production; for example, such a detection of a {\\CRI} component substantially stronger than in Eq.~(\\ref{eq:uCR}) would rule out our model. The SNe origin of the {\\CRIs} can be tested by carefully examining the correlation between a {\\CRI} measure such as $\\eta$, and an intensive SNe tracer such as the specific SFR, in a sample of halo clusters." }, "1003/1003.3639_arXiv.txt": { "abstract": "K-shell photoabsorption cross sections for the isonuclear \\ion{C}{1} - \\ion{C}{4} ions have been computed using the R-matrix method. Above the K-shell threshold, the present results are in good agreement with the independent-particle results of \\citet{reilman}. Below threshold, we also compute the strong $1s\\rightarrow np$ absorption resonances with the inclusion of important spectator Auger broadening effects. For the lowest $1s\\rightarrow 2p,3p$ resonances, comparisons to available \\ion{C}{2}, \\ion{C}{3}, and \\ion{C}{4} experimental results show good agreement in general for the resonance strengths and positions, but unexplained discrepancies exist. Our results also provide detailed information on the \\ion{C}{1} K-shell photoabsorption cross section including the strong resonance features, since very limited laboratory experimental data exist. The resultant R-matrix cross sections are then used to model the {\\em Chandra} X-ray absorption spectrum of the blazar Mkn 421. ", "introduction": "The inner-shell excitation and ionization features of cosmically abundant elements fall in the spectral ranges covered by the high-resolution X-ray spectrometers onboard the {\\it Chandra} and {\\it XMM-Newton} observatories. The detailed structure and wavelengths of absorption resonances of a given element depend on its ionization and chemical state. High resolution X-ray spectroscopy of these features can in principle be used to probe the physics and chemistry of astrophysical plasmas. Inner-shell photoabsorption resonances have proven particularly useful for investigating the chemical composition of the interstellar gas in the line-of-sight toward bright sources of X-ray continuum radiation, as demonstrated in the pioneering study of \\citet{schattenburg1986} \\citep[see also][]{paerels2001,takei2002,vries2003,juett2004,juett2006,ueda2005,yao2006,yao2009,kaastra2009}. ~\\citet{ness2007} were able to identify several ionization stages of oxygen in the post-outburst circumstellar material of the recurrent nova RS~Oph based on the prominent $1s\\rightarrow 2p$ resonance. These studies have all employed oxygen and higher $Z$ elements. Carbon is the fourth most abundant element in the Galaxy (after H, He, and O), but has not yet been exploited as an X-ray photoabsorption diagnostic. It presents a special challenge for X-ray transmission spectroscopy. X-ray instruments often employ visible/UV light blocking filters based on carbon compounds that are robust to space deployment. These filters imprint strong C K-edge absorption signatures on X-ray spectra, rendering difficult the disentanglement of weaker astrophysical absorption features. The task is hampered further still by a current lack of data describing the expected absorption edge structure and resonances for neutral and ionized C. In order to investigate C absorption features, we perform detailed R-matrix calculations of the K-shell photoabsorption cross-sections of \\ion{C}{1} - \\ion{C}{4}. These cross sections are then used to interpret the x-ray spectra from a high-quality {\\it Chandra} observation of the blazar Mkn~421 and determine relative carbon-ion abundances. ", "conclusions": "Photoabsorption features of carbon ions (\\ion{C}{1}-IV)) are studied using the R-matrix method, yielding, detailed information on the carbon K-shell photoabsorption cross section spectra. Furthermore, we computed photoabsorption cross sections for the additional \\ion{C}{2}, \\ion{C}{3}, and \\ion{C}{4} isonuclear members, for which synchrotron-facility measurements at the ALS and theoretical studies already studied the lowest $1s \\to 2p,3p$ resonance transitions. Our computed resonance positions were within about 0.5 eV of the measured values, and our strengths showed good agreement for some of the resonances but were significantly too large or small, compared to experiment, for certain strong resonances. These computed data are of particular importance for absorption studies of cosmic gas. In turn, a more accurate description of the interstellar absorption near the C K-edge in cosmic sources used as in-flight calibration standards should lead to refinements in the calibration of spectrometers such as the {\\it Chandra} LETGS. Analysis of the LETGS spectrum of Mkn 421 has allowed us to identify interstellar absorption due to \\ion{C}{2} and estimate ion fractions of \\ion{C}{1} and \\ion{C}{2} for the first time using X-rays." }, "1003/1003.1250_arXiv.txt": { "abstract": "We have obtained optical multicolour (UBVRI) linear polarimetric data for 46 of the brightest stars in the area of the open cluster NGC 6124 in order to investigate the properties of the interstellar medium (ISM) that lies along the line of sight toward the cluster. Our data yield a mean polarization efficiency of $P_V/E_{B-V}=$3.1$\\pm$0.62, i.e., a value lower than the polarization produced by the ISM with normal efficiency for an average colour excess of $E_{B-V}=0.80$ as that found for NGC 6124. Besides, the polarization shows an orientation of $\\theta \\sim $8$^\\circ$.1 which is not parallel to the Galactic Disk, an effect that we think may be caused by the Lupus Cloud. Our analysis also indicates that the observed visual extinction in NGC 6124 is caused by the presence of three different absorption sheets located between the Sun and NGC 6124. The values of the internal dispersion of the polarization ($\\Delta P_V\\sim 1.3\\%$) and of the colour excess ($\\Delta E_{B-V}\\sim 0.29$ mag) for the members of NGC 6124 seem to be compatible with the presence of an intra-cluster dust component. Only six stars exhibit some evidence of intrinsic polarization. Our work also shows that polarimetry provides an excellent tool to distinguish between member and non-member stars of a cluster. ", "introduction": "The open cluster NGC 6124 (l= 340.$^\\circ$8, b=+6.$^\\circ$1) is an intermediate-richness, detached cluster of slight concentration with stars in a wide range of brightness located in the Scorpius region. This cluster is seen as projected on the edge of a dust cloud which obscures some of its stars, on the north-western part of cluster area. NGC 6124 contains several red giant stars. This cluster has been studied photoelectrically by Koelbloed (1959), but he only made rough estimates of the distance and of the age of the cluster. In order to get a more accurate determination of those parameters, the study of Koelbloed was extended photoelectrically and photographically to a fainter magnitude limit by Th\\'e (1965). These initial studies of the cluster revealed important discrepancies between the values of the parameters that they obtained. Motivated by some evidence that this cluster is affected by differential reddening, Pedreros (1987) obtained a new set of UBV photoelectric observations for NGC 6124 and analyzed them using a reddening-line slope adequate for this region of the sky. The distance to NGC 6124 that he estimated using the less evolved stars is $563 \\pm 10$ pc. He also found an age of $1.0$ x $10^{8}$ yr and a mean colour excess of $E_{B-V} = 0.80 \\pm 0.06$. A study of the interstellar polarization of NGC 6124 is warranted for three reasons: it provides information on the dust itself, it is a means to trace the galactic magnetic field and it is also useful to establish cluster membership. The comparison between the polarization and extinction values observed along the same line of sight provides tests for extinction models and grain alignment. The latter involves several processes acting simultaneously, but on different time-scales. The rotating dust grains get a substantial magnetic moment that allows them to precess rapidly about magnetic field lines and, as a result, the grains conserve their orientation relative to that of the magnetic field when the latter fluctuates, forcing the axes of the grains to be aligned with respect to the angular momentum of the grains (Lazarian and Cho 2005). Therefore, the observed polarization vectors map the projected direction of the magnetic field on the plane of the sky , and this allows the investigation of the structure of both the macroscopic field in our Galaxy (Mathewson and Ford 1970; Axon and Ellis 1976) and the local field associated with the individual clouds (Goodman et al. 1990). Open Clusters are very good candidates for polarimetric observations, because many of them have been studied through photometric and spectroscopic techniques and detailed information on the colour, luminosity, spectral type, and other properties of their stars is readily available. Thus, the physical parameters of the cluster and its stars can be obtained, and adding polarimetric observations we can study the distribution, size and efficiency of the dust grains which polarize the starlight and the different directions of the galactic magnetic field along the line of sight to the cluster. As some of the open clusters spread over a significant area, it is possible to analyze the evolution of the physical parameters of the dust over the region. Besides, it is also possible to detect the presence (if any) of intra-cluster dust and ,with additional observations of non-member stars in the same region, to investigate the interstellar dust distribution along the line of sight to the cluster. Finally, the polarimetric data can be used to detect the location of any energetic phenomenon that may have occurred in the history of a cluster (Feinstein et al. 2003), and very important by-products of these studies are the detection of stars displaying polarization of non-interstellar origin ,such as stars with extended atmospheres (e.g., Be stars), and of dust associated to possible binary systems, or surrounding the stars (due to evolution or as a formation remnant). Here we present multicolour (UBVRI) measurements of linear polarization vectors for stars observed in the direction to NGC 6124 and we use them to investigate the properties of the dust located along the line of sight to the cluster . In the following sections, we will discuss the observations, the data calibrations and the results obtained both for the individual stars and for the cluster as a whole. ", "conclusions": "\\subsection{Serkowski's Law} To analyze the data, the polarimetric observations in the five filters were fitted in each star of our sample using Serkowski's law of interstellar polarization (Serkowski 1973). That is: $P_{\\lambda}/P_{\\lambda max}=e^{-Kln^2(\\lambda_{max}/\\lambda)} \\ \\ \\ \\ [1]$ We will assume that, if polarization is produced by aligned interstellar dust grains, the observed data (in terms of wavelength, UBVRI) will then follow [1] and each star will have a $P_{\\lambda_{max}}$ and $\\lambda_{max}$ values. To perform the fitting we adopted $K=1.66 \\lambda_{max}+ 0.01$ (Whittet et al. 1992). For all stars in the sample we computed the $\\sigma_{1}$ parameter (the unit weight error of the fit) in order to quantify the departure of our data from the ``theoretical curve'' of Serkowski's law. In our scheme, when a star shows $\\sigma_{1} > 1.5$, it is indicating the presence of intrinsic stellar polarization. Also, $\\lambda_{max}$ values can be used to test the origin of the polarization: those stars which have $\\lambda_{max}$ much lower than the average value of the interstellar medium (0.545 $\\mu m$, Serkowski, Mathewson and Ford 1975) are candidates to have an intrinsic component of polarization as well. Another criterion to detect intrinsic stellar polarization comes from computing the dispersion of the position angle for each star normalized by the average of the position angle errors ($\\bar{\\epsilon}$). The values obtained for $P_{\\lambda_{max}}$, the $\\sigma_{1}$ parameter, $\\lambda_{max}$, and $\\bar{\\epsilon}$ together with the identification of stars, are listed in Table 3. The expression used to calculate the $\\sigma_{1}$ parameter is also shown in the footnote. No star in Table 3 has its $\\lambda_{max}$ much lower than the average for the ISM. Fig. 4 displays $P_{\\lambda}$ and $\\theta_{\\lambda}$ plots for six stars (\\#13,26,27,36,42, and 46) with probable indications of intrinsic polarization. The presence of a non-interstellar component in the measured polarization of the light from a star causes a mismatch between observations and the Serkowski's curve fit, and/or a rotation in the position angle of the polarization vector. This mismatch is clearly shown in the $\\sigma_{1}$ value for star \\#42 ($\\sigma_{1}$= 1.81), and its variations in $P_{\\lambda}$ is a plane curve, indicating more than one component in polarization. In the other five stars, the presence of an intrinsic component of polarization is seen through the variation of $\\theta_ {\\lambda}$ ($\\bar{\\epsilon}$ parameter). In particular, the $P_U$ of star \\#36 is out of the Serkowski's fit and, as it was mentioned before, it is a known binary system. \\subsection{Stokes plane and memberships} To analyze the characteristics of the interstellar medium in the region of NGC 6124, we plot the normalized individual Stokes parameters in the visual filter of the polarization vector {\\bf $\\vec P_V$}, given by $Q_V=P_V cos(2 \\theta_V)$ and $U_V=P_V sin(2 \\theta_V)$, which represent the vector's equatorial components, for each of the observed stars. In this figure the stars are plotted according to the literature with their initial membership status. Filled circles are used for members, open triangles are used for non-members, starred symbols are used for red giant stars and open circles are used for those stars without membership information. The polarimetric technique can help to solve membership problems. This type of plot ($Q_V$ vs. $U_V$) used in combination with photometric information is useful to separate between members and non-members. In Fig. 5 the point of coordinates $Q_V=0$ and $U_V=0$ indicates the dustless solar neighborhood, and any other point represents the direction of the polarization vector with modulus $P_V=\\sqrt{Q_V^2+U_V^2}$ as seen from the Sun. Again, this figure shows a high dispersion in angles. \\subsubsection{Stars close to the Sun} On the left side of the diagram, near point (0,0), there is a group of three stars (\\#6,9,and 31) with very low polarization values (0.11\\%, 0.23\\% and 0.07\\%, respectively). Star 6 was observed by Pedreros (1987) and classified as a non-member of spectral type G5. Star 9, according to Th\\'e (1965), may be a non-member variable star, and star 31 is of spectral type F8V (Houk et al. 1975). The low polarization of these stars is compatible with their small reddening. Therefore, they could be confirmed polarimetrically as non-member stars. Stars \\#3 and 32 (originally considered as members) have their polarization values of 0.89 \\% and 1.02 \\%, respectively. From its photometric data, star 3 (of spectral type G2V) is located near the Sun ($\\sim$ 100 pc), with an $E_{B-V}=0.08$ mag. compatible with its polarization. In Fig. 1 this star is projected on the north-west side of the sky, where a big dust cloud is observed, obscuring several members on this part of the cluster. Therefore, if it were a member, its light would be expected to be highly polarized, but it is not. For Star \\#32, according to its Q parameter (Schmidt-Kaler 1982), we obtain a F0-4V approximate spectral type and a distance of $\\sim$160 pc. From these data, the polarization of the light of both stars is likely caused in a dust cloud near the Sun, and for that we propose them as non-member stars. \\subsubsection{Non-member stars} In Fig. 5 there are two groups of non-member stars with different orientations of polarization vectors, but with similar values in polarization. There is a group made up of 4 non-member stars (\\#34,39,44 and 45) with mean values $\\bar{\\theta}_V$=17$^{o}$ and $\\bar{P}_V$=1.98 \\%. Koelbloed (1959) suggested that \\#44 and \\#45 were members of the cluster with a common spectral type AO, but their reddenings are smaller than those of other A type stars in the main sequence of the cluster; and besides, \\#44 is a B8II star according to Houk et al. (1975). Also, if we take into account that these four stars are very close in projection on the sky (on the west size, see Fig. 1) and in the $Q_V$-$U_V$ plane, indicating similar polarimetric characteristics, probably they are polarized by the same dust cloud. From their photometric data, they are situated nearby the cluster ($\\sim$ 400 pc from the Sun) and in front of it. Therefore, we may confirm them polarimetrically as foreground non-member stars. Also, stars \\#36 (red giant) and \\#42 (member) could form part of this group, but they show some evidence of a non-interstellar component. Other two non-member stars, \\#20 and 28, have very different polarimetric orientations (2$^{o}$.7 and 0$^{\\circ}$.8, respectively) in comparison to the preceding non-members. Both stars are projected on the central part of the cluster (Fig. 1), and to a mean distance of the Sun of approximately 200 pc (from their photometric data). The scatter between the two non-member groups in their position angles ($\\Delta\\theta_V \\sim 16^{\\circ}$) may be showing that the light from these stars is polarized by several overlapped dust components along the line of sight to the cluster, with different orientations of the local magnetic fields. \\subsubsection{Cluster stars} Most stars in our sample have their $Q_V$ values in the range (1.29, 2.6), with a high scatter in polarization angles. In this range we observe a group of stars showing some degree of concentration, considered most of them as members in the literature. To derive mean values, in angle and polarization for the cluster, we use seven member stars (\\#8,18,21,30,37,38, and 40), obtaining $\\bar{P}_V$=2.11\\% and $\\bar{\\theta}_V$=8$^{\\circ}$.2. The wavelength of maximum polarization for the same group of stars amounts to $\\bar\\lambda_{max}=0.57\\pm0.04 \\mu$m, a value very close to that of the ISM. According to their locations on the $Q_V$-$U_V$ plane stars \\#22 and 26 (originally non-members) may be polarimetrically considered now as members. On the upper side of this group, we find three stars \\#5,33, and 43 with polarization values higher than the mean value for the cluster ($P_V$=2.77\\%, 2.78\\%, and 2.7\\%, respectively). Star \\#5, of spectral type B6II/III (Houk et al 1975), is considered as a member by Pedreros (1987), with a 60\\% of membership probability (Baumgardt et al. 2000). Polarimetrically, it shows characteristics in angle similar to the members of the cluster, in spite of its high polarization value. Its estimated distance, according to the spectral type, is approximately of 570 pc. As mentioned before, star \\#33 has a spectral type M0III and probably it is a distant giant not belonging to the cluster. Star \\#43 (spectral type B2III-IV) is considered a member with a membership probability of 77\\% (Baumgardt et al. 2000), but the orientation of the polarization vector (${\\theta_V}$=18$^{o}$.1) is very different from the mean value of the cluster. As this star is an evolved object, an intrinsic polarization component may be expected, but our indicators do not show it. Therefore, these three stars may be possible non-members, due to their position in the $Q_V$-$U_V$ plane. Stars \\#11, 12, 13 and 16 are located on the lower side of the $Q_V$-$U_V$ plane. Star \\#16 is considered as a member by Pedreros (1987), but its polarization value (1.63\\%) and orientation (${\\theta_V}$=176$^{o}$) are both lower than the corresponding values of the rest of members. The estimated spectral type for this star, based on photometric data, gives A8V with a $E_{B-V}$=0.39 mag. and a distance to the Sun smaller than that of the cluster. In Fig. 1, it lies near stars \\#20, and \\#28 on the central part of the cluster, and also in the $Q_V$-$U_V$ plane, which indicates similar polarimetric characteristics in the dust causing the polarization. Polarimetrically, we conclude that star \\#16 may be a non-member in front of NGC 6124. Stars \\#11 and 12 have no membership information, and there is a lack of (U-B) values, therefore it cannot make any prediction about it. Polarimetrically, star\\#11 has a high polarization value to be considered a member, and the orientation of the polarization vector in the star \\#12 is very different from the $\\bar{\\theta}_V$ for the cluster. Star \\#13 displays intrinsic polarization, and its location in the $Q_V$-$U_V$ plane is clearly of a non-member. \\subsubsection{Background Stars} On the right side of the same figure, it can be seen three stars (\\#1,2 and 4) which are the most polarized in our sample with $P_V=$3.36\\%, 3.42\\% and 3.35\\% respectively. On the sky plane (Fig. 1), they are located on the west side of the cluster and to the south of a big dust cloud. Their polarization angles ($\\theta_V$=6$^{\\circ}$.4, 4$^{\\circ}$.8, and 8$^{\\circ}$.7, respectively) are similar to the mean value of the cluster. Star 1 was classified as G8III; therefore, it should have $E_{B-V}=0.87$ mag. Originally, it was considered as a member by Pedreros(1987), but its membership probability is lower than 50\\% (Baumgardt et al. 2000). The other two stars have not membership information. The location of these three stars in the $Q_V$ vs. $U_V$ plot shows characteristics of background stars, their light being polarized for a dust component between them and the cluster. \\subsection{Polarization efficiency and distribution of interstellar matter } The ratio ${P_V}/{E_{B-V}}$ is known as the \"polarization efficiency\" of the interstellar medium and it depends mainly on the alignment efficiency, the magnetic field strength and the depolarization due to radiation traversing more than one dust cloud with different field directions. Fig. 6 displays the relation ($P_V$ vs. $E_{B-V}$) that exists between the reddening and the polarization produced by the dust grains along the line of sight to NGC 6124. Assuming normal interstellar material characterized by R=3.2, the empirical upper limit relation for the polarization efficiency given by $P_{V}=R {A_V} \\sim 9E_{B-V}$ (Serkowski et al. 1975) is depicted by the solid line in this figure. Indeed, this line represents the maximum efficiency of the polarization produced by the interstellar dust. Likewise, the dotted line $P_V=3.5{E_{B-V}}^{0.8}$ represents the most recent estimate of the average efficiency made by Fosalba et al. (2002), valid for $E_{B-V} < 1.0$ mag. The dashed line $P_V=5E_{B-V}$ was drawn as a reference and represents the observed normal efficiency of the polarizing properties of dust given by Serkowski et al. (1975). In this figure the stars are plotted with our polarimetric membership status. Excesses $E_{B-V}$ were obtained from the literature or from the relation between spectral type and colour index following Schmidt-Kaler (1982). In Fig. 6 the majority of stars lie on the right of the interstellar maximum efficiency line (${P_V}/{E_{B-V}} \\sim 9$, Serkowski et al., 1975), indicating that the observed polarization is mostly due to the ISM. Only two stars (\\#3,28) lie on the left of the line; this can be associated with errors in the estimate of their excesses. To derive the mean polarimetric efficiency, we made use of the same 7 stars that we selected to calculate the mean polarization and angle of NGC 6124 (section 4.2), located around the cross symbol in Fig. 6. We obtained a mean efficiency ${P_V}/{E_{B-V}}$=3.1$\\pm $0.62, which is smaller than the value for the interstellar dust given by the Fosalba's estimate (${{P_V}/{E_{B-V}}}\\sim 3.66$) and still much smaller than the value given by the Serkowski's estimate (${{P_V}/{E_{B-V}}}\\sim 5.$) for the average efficiency.. This depolarization in the cluster can be the result of the superposition of several dust clouds on the line of sight with different magnetic field directions, confirming the information given in section 4.1. Also, it shows scattering in colour excess (approximately $0.60$10$^3$\\,M$_{\\sun}$ are the natural habitat of high-mass stars. Whether motivated by the study of the formation, evolution of feedback effects of high-mass stars themselves \\citep{Bik_10,Martins_09,Puga_09} or as probes to investigate the physics and chemistry evolution of the Galaxy \\citep{Messineo_09,Davies_09}, NIR spectroscopic studies have been proven powerful to unveil and characterise the complete high-mass stellar content within massive clusters.\\\\ Although new NIR surveys, such as UKIDSS and VISTA, will surely uncover new obscured massive clusters, some clusters, already known from optical observations, harbour a hidden high-mass stellar component only accessible at NIR wavelengths \\citep[e.g. Cyg~OB2,][]{Knoedl_00}.\\\\ \\object{NGC~7538} (aka Sh~2-158) is a visible \\HII{} region in the Perseus spiral arm and a site of active star formation. The early detection of several luminous NIR and far-IR sources in the vicinity of this region \\citep{Campbell_88}, hinted toward a rather massive nature. In particular, NGC~7538 IRS~1 is a high-mass ($\\sim$30\\,M$_{\\sun}$) protostar with a CO outflow, known to power the ultra-compact (UC)\\HII{} region NGC~7538\\,A. It also has an associated linear methanol maser structure, which might trace a Keplerian-rotating circumstellar disk \\citep{Pestalozzi_04}. Recently, 6.7\\,GHz methanol masers have been detected toward the nearby objects NGC~7538 IRS~9 and NGC~7538\\,S, tracing other young and embedded massive protostars \\citep{Pestalozzi}. \\\\ Although this region has been widely inspected at long wavelengths (mostly in the submillimetre window) and even optical spectroscopy for two stars has been obtained \\citep{Russeil_07}, only a few detailed near-IR photometric studies of the stellar population in NGC~7538 have been conducted \\citep{Balog_04,Ojha_04,McCaughrean_91}. At subarcsecond resolution, the region breaks down into several areas of interest (see Fig.~\\ref{ngc7538}): ({\\sc i}) the presumed powering cluster, centred at the location of IRS~6, ({\\sc ii}) the subcluster sitting amidst the bright NIR reflection nebula (IRS~5), ({\\sc iii}) the cluster located at the southern rim of the optical \\HII{} region (IRS~4), ({\\sc iv}) the active region around IRS~1/2, ({\\sc v}) an embedded stellar cluster at the location of NGC~7538~S, only detected at wavelengths longward of 3.6 $\\mu$m, and ({\\sc vi}) the IRS~9 reflection nebula. Surprisingly, the stellar density detected in the NIR is highest at the southern rim of the dust bubble that bounds the optical \\HII{} region \\citep{Balog_04}, between regions ({\\sc iii}) and ({\\sc iv}). The distance to NGC~7538 has been several times estimated and values between 2.2 and 3.5\\,kpc have been reported \\citep{Moreno_86,Israel_73}. \\begin{figure*}[!ht] \\centering \\includegraphics[bb=53 34 340 340,scale=1.0]{3294fig1.eps} \\caption{ Colour composite image of NGC~7538 obtained with LIRIS ({\\it blue}: $J$, {\\it green}: $H$, {\\it red}: $K_s$). Overlayed contours depict the extension of the \\HII{} region in the DSS/POSSII-F Red map. Stars observed with the LIRIS--MOS mode are labelled in cyan. Object \\#4 was observed with the same instrument in long-slit mode. The green circles and cross are 2MASS point sources of interest.}\\label{ngc7538} \\end{figure*} In this article, we present NIR spectra of a sample of candidate high-mass stars within NGG~7538 together with sub-arcsecond NIR photometry.\\\\ We describe the observations and data reduction in Sect.~\\ref{obs_data_red}. In Sect.~\\ref{results} we analyze the spectro-photometric information of a sample of candidate high-mass stars. Sect.~\\ref{discussion} is dedicated to the estimate the distance, age and mass of the cluster powering the \\HII{} region. Finally, we conclude in Sect.~\\ref{conclusions} with a brief summary. ", "conclusions": "\\label{discussion} \\subsection{Distance, age and mass}\\label{sect_dist} An initial estimate of the distance based on spectro-photometry by \\cite{Blitz} situated NGC~7538 at an heliocentric distance of 2.8$\\pm$0.9\\,kpc. Meanwhile, kinematical estimates delivered a distance for the same region as far as 3.5 kpc \\citep{Israel_73}. \\cite{Russeil_07} determined the spectral type classification of two stars in the \\HII{} region: IRS~6 (O9V), and IRS~5 (O3V), inferring spectro-photometric distances of 1.6$\\pm$0.17 and 4.24$\\pm$0.29\\,kpc, respectively. They estimated an average distance of 2.27$\\pm$0.15\\,kpc for the entire Sh~2-158 region. The most recent estimate of the distance to NGC~7538 has been obtained measuring trigonometric parallaxes of methanol masers, regularly associated with high-mass star forming regions. This study yielded a distance of 2.65$^{+0.12}_{-0.11}$\\,kpc \\cite{Moscadelli}.\\\\ Our revised optical classification of the O-type sources IRS~6 and IRS~5 described in Sect. \\ref{spect_class} renders a new distance estimate to NGC~7538. Considering a classification O3V for the IRS~6 object, we derive a spectro-photometric distance of 2.99$\\pm$0.5\\,kpc to the cluster powering the \\HII{} region. Likewise, the classification of IRS~5 yields a value of 2.39$\\pm$0.4\\,kpc. We report a final spectro-photometric distance of 2.7$\\pm$0.5\\,kpc to this region. Despite the large extinction variations in the region that hampered a robust age determination, \\cite{Balog_04} found an older generation of stars ($\\sim$\\,4 Myr) and a younger population of faint stars closer to $\\sim$\\,1 Myr. A fraction of 30\\% of young stars in the region around NGC~7538 was reported to exhibit a near-IR excess. An upper limit limit to the age of the cluster can be established from the most massive classified star member. According to \\cite{Schaerer_97}, a 60\\,M$_{\\sun}$ star (corresponding to a O3V star) has a main-sequence lifetime $\\sim$2.2\\,Myr before it starts the giant phase. Assuming a coeval star-formation event, the oldest stellar content in NGC~7538 must be, therefore, younger than 2.2\\,Myr. \\\\ At the other end of the mass spectrum, we can use the information derived from the PMS stars to constrain the age of the low-mass population around NGC~7538. To construct a Hertzsprung-Russell diagram (HRD), we use the spectral type to temperature conversion from \\cite{Kenyon&Hartmann}, including the overestimates reported by \\cite{Cohen_79} as temperature error. The absolute $K$ magnitudes are then calculated from the dereddened apparent brightness and assuming a distance of 2.7\\,kpc. The HRD of the PMS stars identified in NGC~7538 is shown in Fig.~\\ref{Age_Mass_pms}. We have depicted in the figure several pre-main-sequence evolutionary tracks and isochrones taken from \\cite{DaRio}. The mass of these PMS stars is estimated in the range 2-4 M$_{\\sun}$. The comparison with the PMS isochrones yields an age range between 0.5 and 2\\,Myr for the PMS identified population. We conclude, thus, that these two estimates result in an age range between 0.5 and 2.2\\,Myr for the powering cluster. \\\\ Clearly, other massive young stellar populations are also present at several locations around NGC~7538. They range from stars for which we observe photospheric spectral features, are detectable in the optical, but also exhibit very strong IR excess (IRS~5 at region {\\sc(ii)}) to objects whose emission is dominated by the presence of a circumstellar disk (region {\\sc(iii)}), and finally to nascent stars that present on-going outflow activity (IRS~1/2 and IRS~9 in regions {\\sc(iv)} and {\\sc(vi)}, respectively).\\\\ These results support the previous idea of a star-forming sequence in the northwest-southeast direction that may be smoother than previously considered. The quantitative analysis of the individual ages of these populations is beyond the scope of this paper. \\begin{figure} \\begin{center} \\resizebox{\\hsize}{!}{ \\includegraphics[bb= 24 0 248 200]{3294fig8.eps} } \\caption{Extinction corrected Hertzsprung-Russell diagram of the identified PMS stars within NGC~7538. The open symbols correspond to PMS stars without an IR excess detected in their LIRIS and IRAC photometry, while the filled ones represent those with an IR excess detected. The solid lines represent the pre-main-sequence isochrones, while the dotted lines correspond to evolutionary tracks by \\cite{DaRio}.}\\label{Age_Mass_pms} \\end{center} \\end{figure} Assuming a Salpeter initial mass function (IMF), and extrapolating it down to 0.8 M$_{\\sun}$, we can calculate the stellar mass of a cluster. Yet, this estimate is only meaningful for a coeval stellar population. This should be the case for the powering cluster of the \\HII{} region, with IRS~6 as the ionizing star. Normalising the stellar mass distribution by the detection of one O3V star with a mass uncertainty between 47 and 64\\,M$_{\\sun}$, we obtain a total mass $\\sim$1.7$\\times$10$^3$\\,M$_{\\sun}$ for the cluster. \\subsection{Spatial distribution} The second most massive identified star is IRS~5 and, although is located amidst the ionised emission, its infrared excess and proximity to the rim of the molecular cloud suggest that it may not be a member of the central cluster, and region {\\sc (ii)} in fact comprises a younger generation of stars. Among the spectroscopically identified B-type stars, only object \\#25 is located amidst the ionised emission, while sources \\#17 and \\#35 are farther away. In the case of \\#35, this source is located in the vicinity of the young region IRS~1/2 or {\\sc (iv)} and exhibits an infrared excess typical of Class~0/I objects.\\\\ In fact, our CMD on the left panel of Fig.~\\ref{CMD_CCD} reveals a deficit of early B-type candidates for sources located within a radius of 30$\\arcsec$ from the O3V ionizing star IRS~6 (objects marked by squares). To further study this sparseness of early B-type stars associated to the cluster, we identify in the CMD B-type candidates of compatible $K_s$-band brightness and $H-K$$\\sim$0.4. These objects appear indicated by diamonds in the left panel of Fig.~\\ref{CMD_CCD}. Analogously to our previous analysis of spectroscopically identified B-type stars, a few of the B-type candidates show a mid-IR excess (Chavarr{\\'{i}}a et al., private communication) suggesting they belong to younger pockets of star formation. Other B-type candidates exhibit only photospheric emission and they are located to the east and shouth of IRS~6, spreading beyond the boundary of the \\HII{} region, possibly due to a mass segregation effect. We sketch a possible picture in which the cavity created by the ionizing stars in the molecular cloud is open in the observer's direction and part of its stellar population is located in the foreground of the neutral molecular cloud, appearing almost aligned with other embedded younger star-forming regions. This hypothesis is supported by the work of \\cite{Balog_04} who found that the stellar density peak of this region at NIR wavelengths is located at the rim of the \\HII{} region. \\\\ This projection effect is particularly important to bear in mind when isolating the stellar populations of the different pockets of star-formation that may be triggered by the expansion of \\HII{} regions." }, "1003/1003.5703_arXiv.txt": { "abstract": "Nonradial pulsations in the primary white dwarfs of cataclysmic variables can now potentially allow us to explore the stellar interior of these accretors using stellar seismology. In this context, we conducted a multi-site campaign on the accreting pulsator \\sdsosf\\ (V386~Ser) using seven observatories located around the world in May 2007 over a duration of 11 days. We report the best fit periodicities here, which were also previously observed in 2004, suggesting their underlying stability. Although we did not uncover a sufficient number of independent pulsation modes for a unique seismological fit, our campaign revealed that the dominant pulsation mode at 609\\,s is an evenly spaced triplet. The even nature of the triplet is suggestive of rotational splitting, implying an enigmatic rotation period of about 4.8 days. There are two viable alternatives assuming the triplet is real: either the period of 4.8 days is representative of the rotation period of the entire star with implications for the angular momentum evolution of these systems, or it is perhaps an indication of differential rotation with a fast rotating exterior and slow rotation deeper in the star. Investigating the possibility that a changing period could mimic a triplet suggests that this scenario is improbable, but not impossible. Using time-series spectra acquired in May 2009, we determine the orbital period of \\sdsosf\\ to be 83.8\\,$\\pm$\\,2.9\\,min. Three of the observed photometric frequencies from our May 2007 campaign appear to be linear combinations of the 609\\,s pulsation mode with the first harmonic of the orbital period at 41.5\\,min. This is the first discovery of a linear combination between nonradial pulsation and orbital motion for a variable white dwarf. ", "introduction": "Cataclysmic variables are interacting binary systems in which a late-type star loses mass to an accreting white dwarf. Photometric variations consistent with nonradial g-mode pulsations were first discovered in the cataclysmic variable GW~Librae in 1998 \\citep{WarneravanZyl98,vanZylet00,vanZylet04}; such pulsations had previously been observed only among the non-interacting white dwarf stars. This discovery has opened a new venue of opportunity to learn about the stellar parameters of accreting variable white dwarfs using asteroseismic techniques \\citep[e.g.][]{Townsleyet04}. A unique model fit to the observed periods of the variable white dwarf can reveal information about the stellar mass, core composition, age, rotation rate, magnetic field strength, and distance \\citep[see the review papers][]{Winget98,WingetaKepler08,FontaineaBrassard08}. There are now thirteen accreting pulsating white dwarfs known \\citep[see][]{vanZylet04,WoudtaWarner04,WarneraWoudt04, Pattersonet05a,Pattersonet05b,Vanlandinghamet05,Araujo-Betancoret05,Gaensickeet06,Nilssonet06,Mukadamet07b,Pavlenko08,Pattersonet08}. \\citet{Szkodyet02a,Szkodyet07,Szkodyet10} are pioneering the effort to empirically establish the pulsational instability strip for accretors and to test the theoretical framework laid down by \\citet{Arraset06}. The instability strip(s) for these pulsators has to be established separately from the \\zzc\\ strip\\footnote{Non-interacting hydrogen atmosphere (DA) white dwarfs are observed to pulsate in a narrow instability strip located within the temperature range 10800--12300\\,K for $\\log~g\\approx 8$ \\citep{Bergeronet95,Bergeronet04, KoesteraAllard00,KoesteraHolberg01,Mukadamet04,Gianninaset05}, and are also known as the \\zzc\\ stars.} because accretion enriches their envelopes with He and metals. This is distinct from the pure H envelope of the non-interacting DA white dwarfs, where H ionization causes them to pulsate as \\zzc\\ stars. \\citet{Arraset06} find a H/HeI instability strip for accreting model white dwarfs with a blue edge near 12000\\,K for a 0.6\\,\\msun star, similar to the \\zzc\\ instability strip. They also find an additional hotter instability strip at $\\approx$15000\\,K due to HeII ionization for accreting model white dwarfs with a high He abundance ($>$\\,0.38). The spectrum of an accreting pulsator includes prominent broad absorption lines from the white dwarf as well as the central emission features from the accretion disk. When the orbital period of a cataclysmic variable is $\\sim$80-90\\,min, it is near the evolutionary orbital period minimum, where the rate of mass transfer is theoretically expected to be the smallest $\\sim10^{-11}$\\,\\msun/yr \\citep{KolbaBaraffe99}. Due to the low rates of mass transfer, the white dwarf is expected to be the source of 90\\% of the optical light observed from these systems \\citep{TownsleyaBildsten02}. This makes it possible to detect white dwarf pulsations in these cataclysmic variables. Accreting pulsators have probably undergone a few billion years of accretion and thousands of thermonuclear runaways. Studying these systems will allow us to address the following questions: to what extent does accretion affect the white dwarf mass, temperature, and composition and how efficiently is angular momentum transferred to the core of the white dwarf. These systems are also crucial in understanding the above effects of accretion on pulsations. Asteroseismology can allow us to obtain meaningful mass constraints for the pulsating primary white dwarfs of cataclysmic variables. Previously any such constraints on the mass of the accreting white dwarf could only be established for eclipsing cataclysmic variables \\citep[e.g.][]{Woodet89,Silberet94,Singet07,Littlefairet08}. Constraining the population, mass distribution, and evolution of accreting white dwarfs is also important for studying supernovae Type Ia systematics. For example, \\citet{Williamset09} show empirically that the maximum mass of white dwarf progenitors has to be at least 7.1\\,\\msun\\, thus constraining the lower mass limit for supernovae progenitors. ", "conclusions": "If even approximately correct, the spin period of 4.8 days is remarkably long for an object accreting matter in a binary with an orbital period of 83.8 minutes. The implied rotational velocity for a period of 4.8\\,days would be $\\leq$1\\,km/s as opposed to the typical rotational velocities of 300--400\\,km/s observed for non-magnetic accreting white dwarfs \\citep{Szkodyet05}. The measured surface velocity (vsin\\,i) for the cataclysmic variables VW Hyi is ~600\\,km/s \\citep{Sionet95}, WZ Sge is ~200--400\\,km/s \\citep{Longet03}, and U Gem is ~50--100\\,km/s \\citep{Sionet94}. \\citet{Kepleret95} find the rotation period of the \\zzc\\ star G\\,226-29 to be 8.9\\,hr, while \\citet{Mukadamet09} compute the rotation period of \\zzc\\ itself to be 1.5 days; all known \\zzc\\ stars are non-magnetic. Whether they pulsate or not, non-interacting white dwarfs in general are known to rotate slowly \\citep{Koesteret98,Bergeret05}. Alternatively, perhaps the splitting of the 609\\,s mode is only indicative of the rotation period of the region of the star that it samples well. In other words, we examine the possibility of differential rotation with a rapidly rotating exterior and a slowly rotating interior. \\citet{PiroaBildsten04} find that the accreted angular momentum is shared with the accumulated envelope on short timescales. Given the small splitting, any differential rotation must be small or constrained to a very localized layer of the star (e.g. the surface). This again provides strong constraints on any diffusive mechanism for angular momentum transport. Evaluating the constraints quantitatively requires defining both the mode eigenfunctions and candidate rotation profiles, as well as how the rotational shear depends on latitude. Such an extensive theoretical analysis is beyond the scope of this observational paper, but will be pursued in the near future by members of our collaboration. Such a long spin period of 4.8 days derived from the triplet spacing cannot be discounted out of hand. While white dwarfs gaining matter are thought to spin up \\citep[e.g.][]{Kinget91,YoonaLanger04}, those undergoing classical novae are thought to eject any accreted mass (see discussion in \\citealt{TownsleyaBildsten04}) or lose small quantities of mass based on the abundance in their ejecta (e.g. \\citealt{Gehrzet98}). Therefore these objects have ample opportunity, depending on how angular momentum is exchanged between the accreted envelope and core, to either gain or shed angular momentum and allow the core to spin up or down over the several Gyr accretion history of old cataclysmic variables like \\sdsos. \\citet{LivioaPringle98} propose a model in which the primary white dwarf loses accreted angular momentum during nova outbursts. The demonstration by \\citet{CharbonnelaTalon05} of how internal gravity waves can extract angular momentum from the solar core during its evolution might provide a model for how the core of an accreting white dwarf could be spun down, with the angular momentum carried away with the envelopes over many nova ejections. The transport of angular momentum in stably stratified layers within stars remains poorly understood. Diffusive prescriptions, even ones which depend on magnetic effects like that of \\citet{Spruit02} used by \\citet{YoonaLanger04}, are contradicted by the observed rotational profile of the Sun \\citep{Thompsonet03}. Despite having been spun down on the main sequence, there is no observed gradient in the rotation profile of the solar core, an essential aspect of angular momentum transport in any diffusive prescription. Additionally such prescriptions predict rotation periods much shorter than those observed for isolated non-magnetic white dwarfs \\citep{Suijset08}. \\citet{Katz75} showed that a radial magnetic field in a partially crystallized white dwarf would be sheared by differential rotation, leading to an increase in the azimuthal component proportional to the cumulative angle of differential rotation. He also indicated that when the field strength reached about $10^5$ Gauss, the star would be locked as a rigid rotator. \\citet{WarneraWoudt02} also demonstrate the need for a magnetic field for rigid body rotation in a white dwarf. Perhaps in the context of \\sdsos\\, this implies that its magnetic field is weak." }, "1003/1003.5229_arXiv.txt": { "abstract": "{ Debris disks are optically thin, almost gas-free dusty disks observed around a significant fraction of main-sequence stars older than about 10 Myr. Since the circumstellar dust is short-lived, the very existence of these disks is considered as evidence that dust-producing planetesimals are still present in mature systems, in which planets have formed~--- or failed to form~--- a long time ago. It is inferred that these planetesimals orbit their host stars at asteroid to Kuiper-belt distances and continually supply fresh dust through mutual collisions. This review outlines observational techniques and results on debris disks, summarizes their essential physics and theoretical models, and then places them into the general context of planetary systems, uncovering interrelations between the disks, dust parent bodies, and planets. It will be shown that debris disks can serve as tracers of planetesimals and planets and shed light on the planetesimal and planet formation processes that operated in these systems in the past. \\keywords { planetary systems: formation - circumstellar matter } } \\authorrunning{A. V. Krivov} % \\titlerunning{Debris disks in planetary systems} % ", "introduction": "\\label{sect:intro} An inventory of our own planetary system uncovers its complex architecture. Eight known planets are arranged in two groups, four terrestrial ones and four giants. The main asteroid belt between two groups of planets, terrestrial and giant ones, comprises planetesimals that failed to grow to planets because of the strong perturbations by the nearby Jupiter \\citep[e.g.,][]{safronov-1969,wetherill-1980}. The Edgeworth-Kuiper Belt (EKB) exterior to the Neptune orbit is built up by planetesimals that did not form planets because the density of the outer solar nebula was too low \\citep[e.g.,][]{safronov-1969,lissauer-1987,kenyon-bromley-2008}. Both the asteroid belt and the Kuiper belt are heavily sculptured by planets, predominantly by Jupiter and Neptune respectively. Accordingly, they include non-resonant and resonant families, as well as various objects in transient orbits ranging from detached and scattered Kuiper-belt objects through Centaurs to Sun-grazers. Short-period comets, another tangible population of small bodies in the inner solar system, must be genetically related to the Kuiper belt that act as their reservoir \\citep{Quinn-et-al-1990}. Asteroids and short-period comets together are sources of interplanetary dust, observed in the planetary region, although their relative contribution to the dust production remains uncertain \\citep{Gruen-et-al-2001}. And this complex system structure was likely quite different in the past. It is argued that the giant planets and the Kuiper belt have originally formed in a more compact configuration (the ``Nice model'', \\citeauthor{Gomes-et-al-2005} \\citeyear{Gomes-et-al-2005}), and that it went through a short-lasting period of dynamical instability, likely explaining the geologically recorded event of the Late Heavy Bombardment (LHB). Similar to the solar system, planetary systems around other stars are more than a star itself and one or several planets. As \\citet{wyatt-2008} pointed out, at the end of the protoplanetary phase a star is expected to be surrounded by one or all of the following components: various planets from sub-Earth to super-Jupiter size; remnant of the protoplanetary disk, both dust and gas; planetesimal belts in which solids continue to grow; planetesimal belts, which are being ground down to dust. And indeed, many systems comprise numerous smaller objects, ranging in size presumably from dwarf planets (like Pluto or Ceres) down to dust. Evidence for this comes from observations of the thermal emission and stellar light scattered by that dust. A common umbrella term for all these ``sub-planetary'' solids is a ``debris disk''. Debris disks are aftermath of planet formation in the past, and they formed in the same process as the planets did. Even in mature systems where the planet formation has long been completed, they continue to evolve collisionally and dynamically, are gravitationally sculptured by planets, and the dust they produce through ongoing collisional cascades responds sensitively to electromagnetic and corpuscular radiation of the central star. Hence debris disks can serve as indicators of directly invisible small bodies, planets, and even stars, are tracers of their formation and evolution, and represent an important component of planetary systems. We start with outlining observational techniques and results on debris disks (Section~\\ref{sect:observations}) and summarizing their essential physics (Section~\\ref{sect:physics}) and models (Section~\\ref{sect:models}). In the subsequent sections, the observational data and their theoretical interpretation are used to draw conclusions about the ``dust end'' of the size distribution, dust itself (Section~\\ref{sect:dust}), its parent planetesimals (Section~\\ref{sect:planetesimals}), planets expected to be present in debris disks systems (Section~\\ref{sect:planets}), and the entire planetary systems (Section~\\ref{sect:planetary_systems}). Section~\\ref{sect:conclusions} provides a summary and lists open questions of the debris disk research. ", "conclusions": "" }, "1003/1003.2409_arXiv.txt": { "abstract": "In this paper we present a path integral formulation of stochastic inflation, in which volume weighting can easily be implemented. With an in-depth study of inflation in a quartic potential, we investigate how the inflaton evolves and how inflation typically ends both with and without volume weighting. Perhaps unexpectedly, complex histories sometimes emerge with volume weighting. The reward for this excursion into the complex plane is an insight into how volume-weighted inflation both loses memory of initial conditions and ends via slow-roll. The slow-roll end of inflation mitigates certain ``Youngness Paradox''-type criticisms of the volume-weighted paradigm. Thus it is perhaps time to rehabilitate proper time volume weighting as a viable measure for answering at least some interesting cosmological questions. ", "introduction": "Introduction} Inflation driven by the potential energy of some effective scalar field \\cite{Guth:1980zm,Linde:1981mu,Albrecht:1982wi} has become a common explanation of the starting state of the radiation-dominated hot big bang model. A key reason for its acceptance is that small quantum fluctuations during the last 60 or so efolds of inflation can develop into almost scale invariant curvature perturbations \\cite{Hawking:1982cz,Guth:1982ec,Bardeen:1983qw} like those that we see today in the cosmic microwave background fluctuations \\cite{Jarosik:2010iu}. Couplings in the inflaton's potential have to be chosen to be very small in order to get the amplitude of the fluctuations suitably low. However, fluctuations in the scalar field increase as the background energy density increases, so in certain circumstances the fluctuations might have a significant effect on the evolution of large patches, leading to ``stochastic inflation'' \\cite{Vilenkin:1983xq}. Such fluctuations might lead to a situation in which part or even in some sense the majority of the universe continues to inflate for all time, i.e.\\ ``chaotic eternal inflation'' \\cite{Linde:1986fd}. The advent of the ``string landscape'' \\cite{Bousso:2000xa,Susskind:2003kw} with its complicated vacuum structure has reinvigorated the search for a suitable measure on inflationary histories in situations where more than one possible history can be conceived of. Much of the debate revolves around the extent to which predictions should be conditioned on observations and, if more inflation leads to more observers, whether and how any ``volume-weighting'' should be implemented. For technical reasons much of this recent work has focussed on models where the inflaton is expected to ``tunnel'' from one vacuum state to another via bubble nucleation \\cite{Garriga:2005av,Easther:2005wi,Vanchurin:2006qp,Bousso:2006ev,Aguirre:2006ak,Aguirre:2006na,Linde:2007nm,DeSimone:2008if,Linde:2008xf}. \\cite{Tegmark:2004qd} is an exception, considering random initial conditions in random potentials, and the ``reheating-volume'' approach has been applied to both stochastic and bubble nucleation models~\\cite{Winitzki:2008yb,Winitzki:2008ph,Winitzki:2008jp}. Quantum cosmological studies \\cite{Hartle:1983ai,Hawking:1998bn,Gratton:2000fj,Hawking:2002af,Hawking:2003bf,Hawking:2007vf,Hartle:2007gi,Hartle:2008ng,Hartle:2009ig,Hartle:2010vi} provide complementary perspectives. The approach discussed in this paper illuminates and expands the approach to stochastic inflation and volume-weighting presented in \\cite{Gratton:2005bi}, in which one follows the evolution of the inflaton in a $\\lambda \\phi^4$ potential in proper time with a Langevin noise term approximating the quantum fluctuations. There expectation values were calculated for the field history and perturbatively corrected for the effects of volume weighting. By allowing for final-time constraints on the field value and considering weighting field values at some time by either the volume at that time or the volume at the final time, \\cite{Gratton:2005bi} began to directly attack the two issues in the debate mentioned above. The current paper addresses more general inflationary models than $\\lambda \\phi^4$ and in some sense corrects the perturbative conclusions of \\cite{Gratton:2005bi} via a non-perturbative treatment of volume weighting by way of a path integral. The change in viewpoint is similar to that in going from the Heisenberg approach to the Feynman approach in quantum mechanics when trying to address a question about the history of the system. An early approach to a Langevin model of stochastic inflation was presented by Hodges in~\\cite{Hodges:1989zz}; more recent work includes \\cite{Martin:2005ir,Martin:2005hb}. Refs.\\ \\cite{Creminelli:2008es,Dubovsky:2008rf} also address eternal inflation in a related manner. In contrast, much of the early work on stochastic eternal inflation \\cite{Goncharov:1987ir,Nakao:1988yi,Sasaki:1988df,Sasaki:1987gy,Linde:1993xx} attempted to follow in time the evolution of a probability distribution for the inflaton with a Fokker-Planck equation (analogous to the Schrodinger approach to quantum mechanics). Such approaches typically broke down after a finite time, when the probability became unnormalizable rising rapidly with field value, leading to the suspicion that Planck-scale effects might be vital in controlling the theory and restoring predictivity. This led in part to proper time volume weighting falling out of favour as a measure on eternal inflation. In addition, puzzles such as the ``Youngness Paradox'' \\cite{Guth:2007ng} (\\textemdash if a fraction more inflation produces exponentially more volume, aren't the most common observers at a given time the youngest ones conceivable?\\textemdash) seem particularly acute with proper time volume weighting. We will see the surprising way in which a constrained path integral approach mitigates all these issues and so it may be suggested that proper time volume weighting should be reinstalled as a useful measure for at least some calculations in stochatic inflation. This paper is organised as follows. First, a measure on slow-roll inflationary histories is presented. Saddle points of histories are discussed, and then volume-weighting is introduced. The $\\lambda \\phi^4$ model is studied in depth, and the way inflation typically ends is investigated. Finally there is a discussion and conclusions. ", "conclusions": "We can assemble what we have learned above into a somewhat cogent picture of volume-weighted stochastic eternal inflation. The field must start off above the eternal inflationary threshold, and then we soon see the volume-averaged field stop decreasing and turn around and begin increasing, indicating that volume effects are outweighing classical drift. From this we may hope that a late-time ``steady state'' situation will arise with late-time results dominating any averaging. After a finite proper time, the volume-averaged field ceases to exist; the system is dominated by strong fluctuations and a ``global'' picture breaks down. Nontheless, we may choose to focus on the observationally relevant but rare regions of the universe where inflation happens to end. Then we find that inflation ends in practically the same slow-roll manner on all proper time slices and hence some level of predictability is restored. The saddle point histories deviate into the complex plane rather than continue to values far above the eternal inflation threshold, indicating that the conventional view of the inflaton as jumping up and down on its potential in the eternal inflation regime might be too simplistic. Because when inflation does end it basically ends in slow-roll, conventional density perturbation calculations should still apply, preserving the successful predictions of conventional treatments of inflation. The general techniques and insights of this paper should apply to many large-field models of inflation. It would be interesting to investigate potentials with multiple vacua. Indeed, for ``Mexican-hat'' type potentials, $V =\\lambda (\\phi^2-\\phi_0^2)^2$, one can analytically obtain an expression for $q$ in terms of $\\phi$ and so obtain an explicit formula for the effective potential for $q$. Thus one could investigate volume-weighting for small-field models of inflation where would might expect its effect to be less pronounced than for the large-field case studied here. (Note that an early work~\\cite{Rey:1986zk} discusses an approximate path integral treatment of the behaviour of the inflaton in a ``new'' inflationary potential.) Numerical investigation of the determinant prefactor would be helpful in getting an idea in how ``classical'' the histories really are. The author has checked numerically that there are no negative modes satisfying the relevant boundary conditions for a sample of representative (real) histories, as one would hope. It would also be possible to go beyond slow-roll, obtaining fourth-order equations for the saddle point histories, though the precise way in which the quantum fluctuations are modelled might need to be thought through more carefully. As in quantum mechanics, we have seen that a path integral approach is particularly useful when asking time-dependent questions and looking for semi-classical histories. It has given us a technique for calculating in volume-weighted eternal inflation that is relevant for observations. We have been able to demonstrate how inflation typically ends normally even with volume-weighting in a manner insensitive to the precise initial conditions. By retreating from demanding a global picture of the universe at all times and rather adopting a more ``top-down'' observationally relevant approach \\cite{Hawking:2002af,Hawking:2003bf,Gratton:2005bi} the path integral has allowed us to push much further than in the Fokker-Planck approach without having to worry about Planck density issues. We have also obtained a different result about the behaviour of the volume-weighted average than in the Langevin treatment of \\cite{Gratton:2005bi}. This is possibly because that work only perturbatively expanded around the classical solution, implicitly forcing one to consider only the subset of histories in which inflation has to end. Finally, let us return to the question of proper-time volume weighting itself. Rather than any intrinsic flaw in the scheme, perhaps it was the gauge-dependence of the questions that proper time volume weighting encouraged one to ask that led to the weighting getting a bad reputation (a canonical example of such a gauge-dependent question being ``which value of the inflaton is most likely at a given time?''). A question that we have addressed in this paper is ``how does inflation end at a given proper time?''. Seeing that the answer only depends very weakly on what that time actually is, we have been able to obtain a satisfactory answer to the more general reasonable question ``how does inflation end?'' even using proper-time volume weighting. So for at least some physically relevant questions perhaps proper-time volume weighting is not so bad after all." }, "1003/1003.0418_arXiv.txt": { "abstract": "{ We have timed four millisecond pulses, PSRs\\,J1721$-$2457, J1745$-$0952, J1810$-$2005, and J1918$-$0642, for up to a total of 10.5 years each using multiple telescopes in the European Pulsar Timing Array network: the Westerbork Synthesis Radio Telescope in The Netherlands, the Nan\\c cay Radio Telescope in France and the Lovell telescope at Jodrell Bank in the UK. The long time span has enabled us to measure the proper motions of J1745$-$0952 and J1918$-$0642, indicating that they have transverse velocities of 200(50) and 54(7)\\,\\kms respectively. We have obtained upper limits on the proper motion of J1721$-$2457 and J1810$-$2005, which imply that they have transverse velocities less than 140 and 400\\,\\kms respectively. In all cases, the velocities lie in the range typical of millisecond pulsars. We present pulse profiles for each pulsar taken from observations at multiple frequencies in the range of 350 to 2600 MHz, and show that J1810$-$2005 shows significant profile evolution in this range. Using our multi-frequency observations, we measured the spectral indices for all four pulsars, and for J1810$-$2005 it appears to be very flat. The flux density of J1918$-$0642 shows extensive modulation which we attribute to the combined effects of refractive and diffractive scintillation. We discuss the possible use of including J1721$-$2457 or J1918$-$0642 in a pulsar timing array, and find that J1918$-$0642 will be useful to include when the timing precision of this pulsar is improved over the next few years. We have searched archival optical observations to detect companions of the binary pulsars, but none were detected. However, we provide lower limits on the masses of the white dwarf companions of PSRs J1745$-$0952 and J1918$-$0642. ", "introduction": "In this paper we present improved timing solutions for four millisecond pulsars (MSPs): PSRs J1721$-$2457, J1745$-$0952, and J1918$-$0642, discovered by \\cite{eb01b}, and J1810$-$2005, that was discovered by \\cite{clm+01}. Long-term timing of MSPs is an important tool to determine masses of the individual stars in binary systems, and by constraining secular effects like proper motion, it can be used to improve statistics on the evolution of these systems. In general, masses of pulsars are not easy to determine. In some cases, when the pulsar is in a binary with another compact object, high precision pulsar timing observations on extended timescales can allow for detecting post-Keplerian parameters of the systems, which can be used to separately measure the individual masses of the stars. Space velocities derived from proper motion measurements of radio pulsars give clues about the evolution of these systems and their birth supernovae (e.g.\\,\\cite{tsb+99,hllk05,lor08}). It is believed that recycled pulsars have lower space velocities than normal, slowly rotating pulsars. However as MSPs are generally the most stable rotators, the effects of their space velocities on their rotational and orbital parameters will be easier to derive from timing measurements of those systems and could be used to deduce the intrinsic properties of the stars. Another interesting and important use of timing % MSPs to high precision, is the formation of a pulsar timing array (PTA, e.g. \\citealt{hjl+08a,jsk+08a}). An instrument like this will use an array of MSPs as the endpoints of a Galaxy-scale gravitational wave (GW) detector. Current estimates predict that to detect a GW background, long-term high precision timing of about 20\\,MSPs is required \\citep{jhlm05,vlml08}. Increasing the number of stable MSPs in the array will improve the detection significance. In order to better understand these systems in general, and to determine their suitability for inclusion in a timing array, their long-term timing behaviour needs to be determined. Moreover, the best frequency at which they should be timed needs to be ascertained. This is a combination of the pulse shape at these frequencies and their intensity. We observed all four pulsars at additional frequecies to find their best possible observing frequency for timing purposes, and discuss their suitability for inclusion in a PTA. Three of the pulsars that are presented in this paper are in low-eccentricity binary systems with white dwarf companions (see Table\\,\\ref{tab:solution}). These systems are usually classified as low-mass binary pulsars (LMBPs). The LMBPs distinguish themselves from intermediate-mass binary pulsars (IMBPs) in having shorter periods ($<10$~ms), very low eccentricities ($<10^{-5}$), and they follow relationships between their orbital period and the eccentricity of their orbit, and their orbital period and the companion mass \\citep{tc99}. The origin of the difference between these two classes of binary MSPs is believed to be an evolutionary effect, and mainly due to the progenitor masses of the companion stars. There are now $\\sim$65 MSP binary pulsars with probable white dwarf companions known (e.g.\\,\\citealt{lor08}) and about 16 of these can be regarded as IMBP candidates \\citep{jcv+06}. Detecting an optical counterpart to one of the binary pulsars allows to derive properties of their white dwarf companions. Because of differences in the cooling properties of white dwarfs in LMBPs and IMBPs, optical observations can in some cases be used to distinguish between the two classes (e.g.\\,\\citealt{vk95}). So far, for all three binary pulsars described in this paper, the most recently published limit on optical magnitude of the companion is $R>24$ \\citep{vbjj05}. ", "conclusions": "Using the WSRT, NRT and Lovell telescopes, we have timed four MSPs for 7.5 to 10.5 years. We have presented updated timing solutions, pulse profiles at multiple frequencies for each pulsar, and scintillation parameters for PSR\\,J1918$-$0642. We have measured transverse velocities for PSRs\\,J1745$-$0952 and J1918$-$0642, and set limits on the velocities of PSRs\\,J1721$-$2457 and J1810$-$2005. All velocities are consistent with previously published distributions for solitary and binary MSPs. We have analysed archival optical observations for the binary MSPs and found no companions to the pulsars. From the magnitde limits we deduce for the companions, we can exclude white dwarfs with thick atmospheres. This indicates that the companions must be heavier than about 0.2\\,\\msun. At this point, the mass restrictions as well as the optical magnitude limits give no conclusive information to classify PSRs\\,J1745$-$0952 or J1810$-$2005 as either LMBPs or IMBPs. For low-eccentricity binary pulsars, the only post-Keplerian parameter that is likely to be measurable is the Shapiro delay. The upper limit for inclination of PSR\\,J1745$-$0952, i$<$34 , suggests that the effect of Shapiro delay in timing will be very low and therefore we can not expect to use pulsar timing analyis to disentangle the individual masses of this system. However, the expected luminosity of $M_R>11.3$ indicates that using a dedicated optical observation, the companion of this pulsar may be detectable, and this will therefore be the most promising way of deducing the pulsar and companion masses." }, "1003/1003.3077_arXiv.txt": { "abstract": "We explored the motion of test particles near slowly rotating relativistic star having a uniform luminosity. In order to derive the test particle's equations of motion, we made use of the radiation stress-energy tensor first constructed by Miller and Lamb \\cite{ML96}. From the particle's trajectory obtained through the numerical integration of the equations of motion, it is found that for sufficiently high luminosity, ``suspension orbit\" exists, where the test particle hovers around at uniform angular velocity in the same direction as the star's spin. Interestingly, it turned out that the radial position of the ``suspension orbit\" was determined by the luminosity and the angular momentum of the star alone and was independent of the initial positions and the specific angular momentum of the particle. Also found is that there exist not only the radiation drag but also ``radiation counter-drag'' which depends on the stellar radius and the angular momentum and it is this radiation counter-drag that makes the test particle in the ``suspension orbit\" to hover around at uniform angular velocity which is greater than that induced by the Lense-Thirring effect (i.e., general relativistic dragging of inertial frame). ", "introduction": " ", "conclusions": "" }, "1003/1003.3900_arXiv.txt": { "abstract": "Recent coronagraphic imaging of the AB Aurigae disk has revealed a region of low polarized scattered light suggestive of perturbations from a planet at a radius of $\\sim 100$ AU. We model this darkened region using our fully non-plane-parallel radiative-transfer code combined with a simple hydrostatic equilibirum approximation to self-consistently solve for the structure of the disk surface as seen in scattered light. By comparing the observations to our models, we find that the observations are consistent with the absence of a planet, with an upper limit of 1 Jupiter mass. ", "introduction": "The Herbig Ae star, AB Aurigae (AB Aur), hosts a much-studied disk of gas and dust thought to be representative of protoplanetary disks during the Jovian planet formation phase. The circumstellar material around AB Aur has been imaged in scattered light by several observatories, including STIS on HST \\citep{1999Grady_etal}, Subaru \\citep{2004Fukagawa_etal}, and the Lyot Project \\citep{Oppe08}. These observations show that AB Aur is surrounded by what appears to be a nearly face-on disk, with structures that look like spiral arms. Using adaptive-optics coronagraphy and polarimetry, \\citet{Oppe08} imaged the AB Aur disk between radii of $43-302$ AU. Their images revealed an azimuthal gap in polarized light at a radius of about 100 AU, along with a 2-$\\sigma$ bright ``spot.'' They interpreted this to indicate the presence of a possible massive planet of $5-37$ Jupiter masses ($M_J$). The position angle (PA) of the putative planet is coincident with the minor axis of the disk, as a well as a gap between two spiral arms observed in other scattered light images \\citep{2004Fukagawa_etal}, suggesting that the darkened region results from anisotropies in the overall disk structure, rather than locally confined to a single perturber in the disk. On the other hand, the spiral arms are not seen by \\citet{Oppe08}. Another alternative is that there is no true gap in the disk at all, but rather that the lack of scattered polarized light in this region is purely a geometrical effect caused by the inclination of the disk \\citep{2009Perrin_etal}. In this scenario, the disk is inclined enough that the far edge of the disk is back scattering rather than forward sacttering, creating a region of lower total polarized intensity in the back scattering region. This is verified by a lower polarization fraction in the supposed gap, but no corresponding decrease in the total scattered light. This is likely to be the correct interpretation of the AB Aur observation. The objective of this paper is to put firm limits on the mass of any planet embedded in AB Aur, complemetary to the interpretation by \\citet{2009Perrin_etal}. Planets embedded in optically thick accretion disks, like the disk around AB Aurigae, are expected to produce perturbations in the density and temperature structure of the disk. \\citet{paper1} and \\citet{paper2} calculated the magnitudes of these perturbations for a range of planet masses and distances. They predicted the formation of a shadow at the position of the planet paired with a brightening just beyond the shadow. \\citet[][henceforth Paper I]{HJC_model} improved upon these calculations, by self-consistently calculating the temperature and density structures under the assumption of hydrostatic equilibrium and taking the full three-dimensional shape of the disk into account rather than assuming a plane-parallel disk. \\citet{Oppe08} suggested that the observed structure resembled models of dust trapped in mean motion resonances with a planet \\citep[e.g.][]{kuch03}. However, AB Aur hosts a gas-dominated disk, as confirmed by detections of numerous molecular species \\citep{1997ApJ...490..792M,2000ApJ...529..391M,2005MNRAS.359..663D,2005Pietu_etal,2005Corder_etal,2008A&A...491..821S} In such a disk, the dynamics of the gas dominate the system and the dust traces the gas on orbital timescales. Moreover, the disk is optically thick, so the scattered light image traces only features in the upper surface of the disk. Thus, the darkened region and bright ``spot'' imaged by \\citet{Oppe08} are unlikely to trace dust concentrations, as might be the case were the disk optically thin, but could perhaps be scattering off structure in the disk created in the wake of a planet. Here we model the observed structure as a shadow caused by the tug of a planet in approximate hydrostatic equilibrium with the gas disk, rather than resonant trapping of ballistic grains. This interpretation is consistent with short stopping time of the dust grains in the gas, $\\sim10^{-3}$ orbits \\citep{YoudinChiang}. We use the algorithms and code described in Paper I and \\citet[][henceforth Paper II]{2009HJC} to model the AB Aurigae system \\citep[stellar mass $2.4\\pm0.2\\,M_{\\odot}$, luminosity 48 $L_\\odot$, accretion rate $10^{-7}\\,M_{\\sun}\\mbox{yr}^{-1}$;][]{garc06,2001Rodgers}. We synthesize images of the disk in scattered light and compare them with the coronagraphic images of the AB Aurigae disk to constrain the mass of the perturbing planet. ", "conclusions": "We find that scattered polarized light images of AB Aur do not indicate the presence of a massive planet. We put an upper limit of 1 $M_J$ on any planet at this location, well below the suggested mass of $5-37\\,M_{J}$ postulated by \\citet{Oppe08}. Indeed, the best models are consistent with the absence of any planet, and support the hypothesis of \\citet{2009Perrin_etal}, that the ``gap'' is simply an inclination effect rather than an indication of a planet. On the other hand, if a massive planet several times Jupiter's mass did exist in the disk, our models suggest that it would have been detectable. In other words, our work suggests that massive planets on 100 AU orbits can now be detected via scattered light imaging of nearby Herbig Ae/Be disks. As a Herbig Ae star, AB Aur will eventually evolve into an A-type main sequence star. We can compare AB Aur to two main sequence A-type stars with directly imaged planets: Fomalhaut's planet is at 120 AU separation from the star \\citep{2008Fomalhaut}, and HR 8799's planets are at 24, 38, and 68 AU separation \\citep{2008HR8799}. Assuming no significant planetary migration, our upper limits indicate that AB Aur is not a younger analog of Fomalhaut. The coronographic spot covers the inner 40 AU of AB Aur's disk, so it may yet prove to be an analog of HR 8799. If planets do migrate significantly after dissipation of the gas disk, and there is some mechanism for expelling planets out to large distances, then AB Aur may evolve into a Formalhaut analog. If systems like Fomalhaut and HR 8799 are common and planet formation and migration occur early, then there may be many planets waiting to be detected in known Herbig Ae/Be disks. The requirements include high angular resolution ($\\lesssim0.1\\arcsec$), good star light suppression, and small inner working angle ($\\lesssim0.3\\arcsec$). Ongoing coronographic surveys, like the SEEDS survey on Subaru and future surveys on the Keck, VLT, and Gemini telescopes, will have the power to probe this regime." }, "1003/1003.4527_arXiv.txt": { "abstract": "{ We report near-infrared spectroscopic observations of the Eta Carinae massive binary system during 2008--2009 using the CRIRES spectrograph mounted on the 8\\,m UT\\,1 Very Large Telescope (VLT Antu). We detect a strong, broad absorption wing in \\ion{He}{i} $\\lambda$10833 extending up to $-1900~\\kms$ across the 2009.0 spectroscopic event. Analysis of archival {\\it Hubble Space Telescope}/Space Telescope Imaging Spectrograph ultraviolet and optical data reveals the presence of a similar high-velocity absorption (up to $-2100~\\kms$) in the ultraviolet resonance lines of \\ion{Si}{iv} $\\lambda\\lambda$1394, 1403 across the 2003.5 event. Ultraviolet resonance lines from low-ionization species, such as \\ion{Si}{ii} $\\lambda\\lambda$1527, 1533 and \\ion{C}{ii} $\\lambda\\lambda$1334, 1335, show absorption only up to $-1200~\\kms$, indicating that the absorption with velocities $-1200$ to $-2100~\\kms$ originates in a region markedly faster and more ionized than the nominal wind of the primary star. Seeing-limited observations obtained at the 1.6\\,m OPD/LNA telescope during the last four spectroscopic cycles of Eta Carinae (1989--2009) also show high-velocity absorption in \\ion{He}{i} $\\lambda$10833 during periastron. Based on the large OPD/LNA dataset, we determine that material with velocities more negative than $-900~\\kms$ is present in the phase range $0.976 \\leq \\phi \\leq 1.023$ of the spectroscopic cycle, but absent in spectra taken at $\\phi \\leq 0.947$ and $\\phi \\geq 1.049$. Therefore, we constrain the duration of the high-velocity absorption to be 95 to $206~\\mathrm{days}$ (or 0.047 to 0.102 in phase). We suggest that the high-velocity absorption component originates from shocked gas in the wind-wind collision zone, at distances of 15 to 45~AU in the line-of-sight to the primary star. With the aid of three-dimensional hydrodynamical simulations of the wind-wind collision zone, we find that the dense high-velocity gas is in the line-of-sight to the primary star only if the binary system is oriented in the sky so that the companion is behind the primary star during periastron, corresponding to a longitude of periastron of $\\omega \\sim 240\\degr-270\\degr$. We study a possible tilt of the orbital plane relative to the Homunculus equatorial plane and conclude that our data are broadly consistent with orbital inclinations in the range $i=40\\degr-60\\degr$.} ", "introduction": " ", "conclusions": "VLT/CRIRES observations of Eta Car provide definitive evidence that high-velocity material, up to $\\sim-1900~\\kms$, was present in the system during the 2009.0 periastron passage. The broad, high-velocity absorption is seen in \\ion{He}{i} $\\lambda$10833 in the VLT/CRIRES dataset only in the spectrum obtained at phase $\\phi=11.991$ to 11.998, showing that it is connected to the spectroscopic event. Near-infrared observations obtained at OPD/LNA from 1992 to 2009 indeed show that the high-velocity absorption in \\ion{He}{i} $\\lambda$10833 is periodic, and tightly connected with phase zero of the spectroscopic cycle as well. Based on the OPD/LNA dataset, we constrained the timescale of detection of the high-velocity gas from 95 to $206~\\mathrm{d}$ (0.047 to 0.102 in phase) around phase zero. We analyzed archival {\\it HST}/STIS ultraviolet data, showing that the \\ion{Si}{iv} $\\lambda$1393, 1402 resonance line also presented a high-velocity absorption component up to $-2100~\\kms$. We presented several reasons why the high-velocity absorption is unlikely to be due to a transitory high-velocity wind of Eta Car~A, or due to a wind eclipse of Eta Car~B. We suggest that our observations provide direct detection of shocked, high-velocity material flowing from the wind-wind collision zone around the binary system. Using detailed 3-dimensional hydrodynamical simulations of the wind-wind collision zone of Eta Car, we found that dense high-velocity gas is in the line-of-sight to the primary star only if the binary system is oriented in the sky such that the companion is behind the primary star during periastron, corresponding to a longitude of periastron of $\\omega \\sim 240\\degr-270\\degr$. Our data is broadly consistent with an orbital inclination in the range $i=40\\degr-60\\degr$. We derived that the high-velocity gas is located at distances of 15 to 45~AU in the line-of-sight to Eta Car~A. More importantly, we can rule out orbital orientations in the range $\\omega \\sim 0\\degr-180\\degr$ for all inclination angles, since these do not produce a significant column density of high-velocity gas in our line-of-sight to match our observations of the high-velocity absorption component. The current 3-D SPH simulations used in this paper do not account for radiative cooling, which makes it difficult to estimate the ionization stage of the high-velocity material in the wind-wind collision zone. In addition to the increase in column density of the high-velocity gas, ionization effects due to the close presence of Eta Car B likely play an important role in explaining the amount of high-velocity absorption seen during periastron. Time-dependent, multi-dimensional radiative transfer modeling of the outflowing gas from the wind-wind collision zone of the Eta Car binary system is urgently warranted, and will allow us to better understand the influence of Eta Car~B on the wind of Eta Car~A across periastron. Ultimately, this will provide constraints on the masses of the stars and on the wind parameters of the Eta Car binary system." }, "1003/1003.5906_arXiv.txt": { "abstract": "{Observations in the cosmological domain are heavily dependent on the validity of the cosmic distance-duality (DD) relation, $D_L(z) (1 + z)^{2}/D_{A}(z) = 1$, an exact result required by the Etherington reciprocity theorem where $D_L(z)$ and $D_A(z)$ are, respectively, the luminosity and angular diameter distances. In the limit of very small redshifts $D_A(z) = D_L(z)$ and this ratio is trivially satisfied. Measurements of Sunyaev-Zeldovich effect (SZE) and X-rays combined with the DD relation have been used to determine $D_A(z)$ from galaxy clusters. This combination offers the possibility of testing the validity of the DD relation, as well as determining which physical processes occur in galaxy clusters via their shapes.} { { We use WMAP (7 years) results by fixing the conventional $\\Lambda$CDM model} to verify the consistence between the validity of DD relation and different assumptions about galaxy cluster geometries usually adopted in the literature.} {We assume that $\\eta$ is a function of the redshift parametrized by two different relations: $\\eta(z) = 1 + \\eta_{0}z$, and $\\eta(z)=1 + \\eta_{0}z/(1+z)$, where $\\eta_0$ is a constant parameter quantifying the possible departure from the strict validity of the DD relation. In order to determine the probability density function (PDF) of $\\eta_{0}$, we consider the angular diameter distances from galaxy clusters recently studied by two different groups by assuming elliptical (isothermal) and spherical (non-isothermal) $\\beta$ models. The strict validity of the DD relation will occur only if the maximum value of $\\eta_{0}$ PDF is centered on $\\eta_{0}=0$.} {It was found that the elliptical $\\beta$ model is in good agreement with the data, showing no violation of the DD relation (PDF peaked close to $\\eta_0=0$ at $1\\sigma$), while the spherical (non-isothermal) one is only marginally compatible at $3\\sigma$.}{The present results derived by combining the SZE and X-ray surface brightness data from galaxy clusters with the latest WMAP results (7-years) favors the elliptical geometry for galaxy clusters. It is remarkable that a local property like the geometry of galaxy clusters might be constrained by a global argument provided by the cosmic DD relation.} \\keywords {X-ray: galaxy clusters, distance scale, cosmic microwave background} \\authorrunning \\titlerunning ", "introduction": "The most useful distances in cosmology are the luminosity distance, $D_L(z)$, and the angular-diameter distance, $D_A(z)$. The expressions of both distances depend on the world models, but the relationship between them, namely \\begin{equation} \\frac{D_{\\scriptstyle L}}{D_{\\scriptstyle A}}{(1+z)}^{-2}=1 \\label{rec0} \\end{equation} is valid for arbitrary spacetimes, a result usually referred to as distance-duality (DD) relation. The above expression can easily be deduced in the context of Friedmann-Robertson-Walker (FRW) cosmologies (Weinberg 1972). However, as originally proven by Etherington (1933), it depends neither on Einstein field equations nor the nature of matter content filling the spacetime. The proof depends crucially on photon conservation (transparency of the cosmic medium) and that sources and observers are linked by null geodesics in a Riemannian spacetime. The DD relation plays an essential role ranging from gravitational lensing studies to analyses of the cosmic microwave blackbody radiation (CMBR) observations, as well as for galaxy and galaxy cluster observations (Schneidder, Ehlers \\& Falco 1999; Komatsu et al.\\ 2011; Lima, Cunha \\& Alcaniz 2003; Cunha, Marassi \\& Lima 2007; Ribeiro 1992, 2005; Ribeiro \\& Stoeger 2003). Indeed, any observational deviation from Eq. (\\ref{rec0}) would be a theoretical catastrophe thereby igniting a major crises in observational cosmology (Ellis 1971, 2007). Although taken for granted in virtually all analyses in cosmology, the DD relation is in principle testable by means of astronomical observations. One may assume a redshift dependence of the form \\begin{equation} \\frac{D_{\\scriptstyle L}}{D_{\\scriptstyle A}}{(1+z)}^{-2}= \\eta(z), \\label{rec} \\end{equation} where $\\eta(z)$ quantifies a possible epoch-dependent departure from the standard photon conserving scenario ($\\eta=1$). Basset \\& Kuns (2004) used both supernovae Ia data as measurements of the luminosity distance $D_{\\scriptstyle L}$ and the estimated $D_{\\scriptstyle A}$ of FRIIb radio galaxies (Daly \\& Djorgovski 2003) and ultra compact radio sources (Gurvitz 1994, 1999; Lima \\& Alcaniz 2000, 2002; Santos \\& Lima 2008) in adopting this kind of approach to test possible new physics. Any source of attenuation (``gray\"' intergalactic dust) or exotic photon interaction must violate the DD relation (More et al. 2009, Avgoustidis et al. 2010), thereby providing new consistency checks of cosmological models. On the other hand, observations of the Sunyaev-Zeldovich effect (SZE) from galaxy clusters are becoming a powerful tool in cosmology (Sunyaev \\& Zeldovich 1972; Cavaliere \\& Fusco-Fermiano 1978, De Filippis et al. 2005, Cunha et al. 2007; Nord et al. 2009; Basu et al. 2010). The combination of SZE and X-ray provides the $D_A(z)$ of galaxy clusters, hence can be used to constrain some cosmological parameters. However, Uzan, Aghanim \\& Mellier (2004) argued that this technique is strongly dependent on the validity of the DD relation. When the DD relation does not hold ($\\eta \\neq 1$), the observationally determined angular distance must be replaced by the value (in the Uzan et al.\\ (2004) notation the correcting term is $\\eta^{-2}$) \\begin{equation} D^{\\: data}_{A}(z)=D_{A}(z)\\eta^{2}, \\label{4} \\end{equation} this quantity reduces to the conventional angular diameter distance (assuming cosmic transparency) only when the DD relation is strictly valid ($\\eta=1$). To quantify the $\\eta$ parameter, Uzan et al.\\ (2004) obtained $D_A(z)$ as given by the cosmic concordance model (Spergel et al. 2003), whereas for $D^{\\: data}_{A}(z)$ they considered 18 angular diameters from the Reese {\\it et al.\\ } (2002) galaxy cluster sample for which a spherically symmetric cluster geometry was assumed. By assuming $\\eta$ to be constant, their statistical analysis provided $\\eta = 0.91^{+ 0.04}_{-0.04}$ (1$\\sigma$) and is therefore only marginally consistent with the standard result, $\\eta=1$. De Bernardis, Giusarma \\& Melchiorri (2006) also searched for deviations from the DD relation by using the {\\bf $D^{\\: data}_{A}(z)$} from galaxy clusters provided by the sample of Bonamente {\\it et al.\\ } (2006). They found a non violation of the DD relation in the framework of the cosmic concordance $\\Lambda$CDM model. { Avgoustidis {et al.\\ }(2010) adopted a DD relation of the form $D_L=D_A(1+z)^{2+\\epsilon}$ to constrain the cosmic opacity by combining the SN Type Ia data compilation of Kowalski { et al.} (2008) with the latest measurements of the Hubble expansion at redshifts in the range $0 < z < 2$ (Stern { et al.} 2009). By working in the context of a flat $\\Lambda$CDM model, they found $\\epsilon=-0.04_{-0.07}^{+0.08}$ (2$\\sigma$).} In the past few years, many studies based on {\\it Chandra } and {\\it XMM} observations have shown that in general galaxy clusters exhibit elliptical surface brightness maps. Simulations have also predicted that dark matter halos show axis ratios typically of the order of $\\approx 0.8$ (Wang \\& White 2009), thereby disproving the spherical geometry assumption usually adopted (Reiprich \\& Boringer 2002; Bonamente et al.\\ 2006, Shang, Haiman \\& Verdi 2009). In this line, the first determination of the intrinsic three-dimensional (3D) shapes of galaxy clusters was presented by Morandi, Pedersen \\& Limousin (2010) by combining X-ray, weak-lensing, and strong-lensing observations. They studied the galaxy cluster MACS J1423.8+2404 and found a tri-axial galaxy cluster geometry with DM halo axial ratios $1.53 \\pm 0.15$ and $1.44 \\pm 0.07$ on the plane of the sky and along the line of sight, respectively. In this letter, we take the validity of the DD relation for granted to access the galaxy cluster morphology. The values of $D_A(z)$ are obtained from the WMAP (7 years) results by fixing the conventional flat $\\Lambda$CDM model whereas the observational measurements $D_A^{\\: data}(z)$ are the angular diameter distances from galaxy clusters obtained via SZE plus X-ray techniques. These samples differ in terms of the assumptions concerning the possible cluster geometries of elliptical and spherical models. Our analysis is based on two parametric representations of $\\eta(z)$ defined by Eq. \\ref{rec} (or Eq. \\ref{4}), namely \\\\ \\\\ \\hspace{1.0cm} I. $\\eta (z) = 1 + \\eta_{0} z$ \\, \\, and \\, \\, II. $\\eta (z) = 1 + \\eta_{0}z/(1+z)$. \\\\ \\\\ \\noindent The first expression is a continuous and smooth one-parameter linear expansion, whereas the second one includes a possible epoch-dependent correction, which avoids the divergence at extremely high z. These deformations of the DD relation effectively parametrize our ignorance of the underlying process responsible for its possible violation. However, we emphasize that these expressions are very simple and have several advantages such as a manageable one-dimensional phase space and a good sensitivity to observational data. Clearly, the second parametrization can also be rewritten as $\\eta (z) = 1 + \\eta_{0}(1-a)$, where $a(z)= (1 + z)^{-1}$ is the cosmic scale factor. This represents an improvement with respect to the linear parametrization, since the DD relation becomes bounded regardless of the redshift values. It will become more useful once higher redshift clusters data become available. The above parametrizations are clearly inspired by similar expressions for the $\\omega(z)$-equation of state parameter of dark energy models (Padmanabhan \\& Choudury 2003; Linder 2003; Cunha, Marassi \\& Santos 2007; Silva, Alcaniz \\& Lima 2007). In the limit of extremely low redshifts ($z<<1$), we have $\\eta = 1$ and $D_{L} = D_{A}$ as should be expected, and, more important for our subsequent analysis, the value $\\eta_0=0$ must be favored by the Etherington result. In other words, for a given data set, the likelihood of $\\eta_0$ must peak at $\\eta_0=0$ to satisfy the cosmic relation. As we shall see, for those accepting the strict validity of the standard DD relation, our analysis suggests that galaxy clusters have an elliptical geometry. In principle, this kind of result is an interesting example of how a cosmological (global) condition correlates with the local physics. \\begin{figure} \\centering \\includegraphics[width=0.55\\linewidth]{15547fg1.eps} \\caption{Galaxy clusters data. The open (blue) and filled (red) circles with the associated error bars represent, respectively, the De Filippis et al.\\ (2005) and Bonamente et al.\\ (2006) samples.} \\label{fig1} \\end{figure} ", "conclusions": "\\label{sec:Conclusions} We have explored some consequences of a deformed distance duality relation, $\\eta(z) = D_{L}(1+z)^{-2}/D_{A}$, based on observations of Sunyaev-Zeldovich effect and X-ray from galaxy clusters. The consistency between the strict validity of the standard relation ($\\eta(z) \\equiv 1$) and the assumptions regarding the geometry used to describe the galaxy clusters (elliptical and spherical $\\beta$ models) has been discussed. The $\\eta(z)$ function was parametrized in two distinct forms, $\\eta = 1 + \\eta_{0}z$ and $\\eta = 1 + \\eta_{0}z/(1+z)$, where $\\eta_0$ is a constant parameter quantifying a possible departure from the strict validity of the duality relation. The basic idea pursued in this work is a simple one. The likelihood of the free parameter appearing in the proposed expressions for $\\eta(z)$ should peak around $\\eta_0 =0$ when the distance duality relation is strictly obeyed. By comparing the De Filippis et al.\\ (2005) (elliptical isothermal $\\beta$ model) and Bonamente et al.\\ (2006) (spherical non-isothermal $\\beta$ model) samples with $D_{A}(z)$ obtained from $\\Lambda$CDM (WMAP7), we show that the elliptical geometry is more consistent with no violation of the duality relation. The uncertainties in $ \\eta_0$ included the systematic plus statistical errors from cluster data. In the case of an elliptical sample (see Fig.\\ \\ref{fig:Analysis}), we found that $\\eta_{0} = -0.056^{+0.1}_{- 0.1}$ and $\\eta_{0} = -0.088^{+ 0.14}_{- 0.14}$ for the linear and non-linear parametrization, respectively. However, the spherical sample (see Fig. 2) is only marginally compatible with $\\eta_{0} = -0.12^{+0.055}_{-0.055}$ and $\\eta_{0} = -0.175^{+ 0.083}_{-0.083}$ for linear and non-linear parametrization, respectively. Our analysis reveals that the elliptical model is compatible with the duality relation validity at 1$\\sigma$, whereas the spherical model is only marginally compatible at 3$\\sigma$. At this point, it is interesting to compare our results with those obtained by following a complementary approach (Holanda, Lima \\& Ribeiro 2010). The $\\eta(z)$ function there was also parametrized as in the present work. However, the overall discussion was based on a model-independent cosmological test by considering $D_A(z)$ from galaxy clusters and the luminosity distances given by two sub-samples of SNe Ia taken from the constitution data (Hicken et al. 2009). Both analyse are consistent with each other and suggest that the elliptical model is more compatible with the validity of the standard duality relation than the spherical case. Summarizing, the statistical analysis presented here provides additional evidence that the true geometry of clusters has an elliptical form. In principle, it is remarkable that a local property (the geometry of galaxy clusters) might be constrained by a global argument such as the one provided by the cosmological distance duality relation. In the near future, as more and larger data sets with smaller statistical and systematic uncertainties become available, the method proposed here (based on the validity of the distance duality relation) can improve the limits on the measurements of cluster geometries. \\centerline{\\bf Acknowledgments} The authors are grateful to an anonymous referee for helpful comments and suggestions that improved the original version of the work. We also thank Antonio Guimar\\~aes for helpful discussions. RFLH is supported by FAPESP (No. 07/52912-2), JASL is partially supported by CNPq (No. 304792/2003-9) and FAPESP (No. 04/13668-0) and MBR is partially supported by FAPERJ." }, "1003/1003.5889_arXiv.txt": { "abstract": "We report on new sensitive CO J=6--5 line observations of several luminous infrared Galaxies (LIRGs: L$_{\\rm IR}$(8--1000$\\mu $m$)\\ga $10$^{11}$\\,L$_{\\odot}$), 36\\% (8/22) of them ULIRGs (L$_{\\rm IR}$$>$10$^{12}$\\,L$_{\\odot}$), and two powerful local AGN: the optically luminous QSO PG\\,1119+120, and the powerful radio galaxy 3C\\,293 using the James Clerk Maxwell Telescope (JCMT) on Mauna Kea in Hawaii. We combine these observations with existing low-J CO data and dust emission Spectral Energy Distributions (SEDs) in the far-infrared - submillimetre from the literature to constrain the properties of the star-forming ISM in these systems. We then build the first {\\it local} CO Spectral Line Energy Distributions (SLEDs) for the {\\it global} molecular gas reservoirs that reach up to high J-levels. These CO SLEDs are neither biased by strong lensing (which affects many of those constructed for high-redshift galaxies), nor suffer from undersampling of CO-bright regions (as most current high-J CO observations of nearby extended systems do). We find: 1) a significant influence of dust optical depths on the high-J CO lines, suppressing the J=6--5 line emission in some of the most IR-luminous LIRGs, 2) low global CO line excitation possible even in vigorously star-forming systems, 3) the first case of a shocked-powered high-excitation CO SLED in the radio galaxy 3C\\,293 where a powerful jet-ISM interaction occurs, and 4) unusually highly excitated gas in the optically powerful QSO PG\\,1119+120. In Arp\\,220 and possibly other (U)LIRGs very faint CO J=6--5 lines can be attributed to {\\it significant dust optical depths at short submm wavelengths} immersing those lines in a strong dust continuum, and also causing the C$^+$ line luminosity deficit often observed in such extreme starbursts. Re-analysis of the CO line ratios available for submillimeter galaxies (SMGs) suggests that similar dust opacities may be also present in these high-redshift starbursts, with genuinely low-excitation of large amounts of {\\it SF-quiescent} gas the only other possibility for their often low CO (high-J)/(low-J) line ratios. We then present a statistical method of separating these two almost degenerate possibilities, and show that high dust optical depths at submm wavelengths can impede the diagnostic potential of submm/IR lines (e.g. starburst versus AGN as gas excitation agents), which is of particular importance for the upcoming observations of the Herschel Space Observatory and the era of~ALMA. ", "introduction": "Obtaining unbiased Spectral Energy Distributions (SEDs) for local and distant galaxy populations provides a crucial yardstick by which to compare their properties and eventually relate such populations along evolutionary paths within any given galaxy evolution framework. Unbiased Spectral Line Energy Distributions (hereafter SLEDs) of the rotational transitions of molecules such as CO, HCN, HCO$^+$ are of particular importance since: \\begin{itemize} \\item they can trace the mass distribution of the star formation fuel, the molecular gas, across its considerable range of properties ($\\rm n$$\\sim $($10^2$--$10^7$)\\,cm$^{-3}$, $\\rm T_k$$\\sim $ (10--200)\\,K) \\item relative molecular line strengths are in principle extinction-free probes of molecular gas properties and AGN-versus-starburst as power sources of IR luminosities of galaxies (Meijerink \\& Spaans 2005; Meijerink, Spaans, \\& Israel 2006) \\item imaging their emission distribution and velocity fields at mm/submm wavelengths yields unique dynamical mass probes of deeply dust-enshrouded star-forming galaxies and QSO host galaxies across the Universe (e.g. Walter et al. 2004; Greve et al. 2005; Tacconi et al. 2006) \\item cooling and thus the thermodynamic state of the molecular gas (heated by stellar far-UV light, cosmic rays, and turbulent motions) is regulated by [C\\,II], [O\\,II] and high-J CO line emission, which in turn may efficiently regulate the local Jeans mass and the stellar IMF (Elmegreen, Klessen, \\& Wilson 2008). \\end{itemize} \\noindent Much of the star formation in the distant Universe occurs in heavily dust-enshrouded IR-luminous systems (e.g. Smail, Ivison, \\& Blain 1997), challenging to image at optical wavelengths even with the current 8-10 meter-class telescopes. Thus it may well be that in the upcoming era of ALMA molecular lines will replace optical and even IR lines as the most potent probes of galaxy structure and of the power sources that drive galaxy evolution across cosmic epoch. The recent discovery of significant dust optical depths {\\it even at submm wavelengths} in Arp\\,220 (a local prototype of dust-enshrouded extreme starbursts) by Sakamoto et al. 2008 further highlights how deeply obscured such extreme star-forming systems can~be. Local templates are key in understanding not only the properties of local galaxy populations, but also of those in the distant Universe. To this end we have undertaken a large multi-J CO and HCN line survey of local Luminous Infrared Galaxies (LIRGs) -- dust-obscured, star-forming systems with SFR$\\sim $(10--100)\\,$\\rm M_{\\odot}\\, yr^{-1}$ (eg. Sanders \\& Ishida 2004), for which the dominant fraction of the bolometric luminosity is in the rest-frame IR with $\\rm L_{IR}$(8-1000)\\,$\\mu $m)$\\ga $$\\rm 10^{11}\\,L_{\\odot}$. The prodigious star formation events in LIRGs/ULIRGs, most often due to dissipative galaxy interactions/mergers of gas-rich progenitors, make them the best local analogues of the submillimeter galaxies (SMGs) (Tacconi et al. 2006; Iono et al. 2009), dust-enshrouded starbursts at high redshifts with even higher star formation rates ($\\rm SFR \\sim$$ \\rm 10^3\\,M_{\\odot}\\,yr^{-1}$) and the sites of a significant part of cosmic star formation history (e.g. Hughes et al. 1998; LeFloch et al. 2009). Finally we observed the hosts of two local AGN: the powerful radio galaxy 3C\\,293, and the optically luminous but radio quiet QSO PG\\,1119+120, as part of a pilot study of the molecular gas excitation in the presence of a bona fide AGN. The latter remain as the most effective beacons of the most distant galaxies where molecular lines have been detected (Walter et al. 2004), and can be the cause of distinct molecular gas excitation conditions in high redshift quasars (Schleicher, Spaans, \\& Klessen~2010). In this paper we report on sensitive CO J=6--5 line measurements for galaxies in our survey. This allows: a) a first systematic glimpse of their high-excitation molecular gas phase, b) constraints on their CO SLEDs, and c) direct comparisons with starburst and AGN-dominated systems at high redshifts where predominantly only high-J CO lines are currently available (e.g. Solomon \\& Vanden Bout 2005; Omont 2007). Throughout this work we adopt a flat $\\Lambda $-dominated cosmology with $\\rm H_0=71\\,$km\\,s$^{-1}$\\,Mpc$^{-1}$ and $\\Omega_{\\rm m}=0.27$, and calculate luminosity distances using the NED calculator developed by Wright (2006). ", "conclusions": "We report on our new sensitive CO J=6--5 observations of LIRGs and two powerful AGN with the JCMT, part of a now completed multi-J CO and HCN molecular line survey. Our findings are as follows: \\noindent 1. Large dust optical depths at short submm wavelengths can suppress the CO J=6--5 line in some extreme starbursts by immersing it in strong continuum dust emission. This can yield faint CO J+1$\\rightarrow $J lines for J+1$\\geq $6, even in ULIRGs whose large amounts of dense star-forming molecular are intrinsically luminous in such lines. Such high optical depths can easily account for the so-called [CII] line luminosity deficity known to exists in such systems. \\noindent 2. Similar conditions may be present in high redshift galaxies, yielding deceptively ``cool'' CO SLEDs at high frequencies even in extreme starbursts such as the submillimeter~galaxies. \\noindent 3. Global dust emission SEDs cannot unambigiously distinguish between high dust optical depths at far-IR/submm wavelengths or large amounts of cold dust, as they can both contribute to the far-IR/submm part of the dust emission SED in a similar fashion. \\noindent 4. The very low CO {\\it and} HCN line excitation in Arp\\,193 demonstrates that low global gas excitation remains possible even in vigorously star-forming LIRGs. \\noindent 5. The suppressing effects of high dust optical depths on the high frequency part of CO SLEDs are difficult to distinguish from genuine low gas excitation. Low-frequency ($<$350\\,GHz) line observations of highly dipolar heavy rotor molecules (e.g. HCN, CS), with their high critical densities but unhindered by potentially large dust extinctions at submm wavelengths, can ``break'' this degeneracy. In the simplest such application we propose that LIRGs with large HCN/CO J=1--0 but very low CO (6--5)/(3--2) line ratios are likely to have dust-affected rather than low-excitation CO SLEDs. \\noindent 6. Remarkably high CO line excitation, above that typical for star-forming gas, is found for the hosts of two prominent AGN, an optically bright QSO (PG\\,1119+120) and a radio galaxy (3C\\,293). The latter could be the first known case of a shock-excited CO SLED, likely powered by a strong jet-ISM interaction. As a result much of its large molecular gas reservoir is hot ($\\rm T_{kin}$$>$100\\,K) and dense ($\\rm n(H_2)$$\\geq $10$^4$\\,cm$^{-3}$) while most of its dust mass remains ``cool'' ($\\rm T_{dust}$$\\sim $15\\,K) typical of quiescent ISM with low star formation rates. In summary, the emerging picture of CO SLEDs in LIRGs and AGN seems diverse, with high dust optical depths at short submm wavelengths capable of ``quenching'' high-J CO line emission in extreme starbursts while low global gas excitation remaining a possibility in such systems. Finally, shocks and possibly X-rays seem capable of surpassing far-UV photons as the main excitation contributors in AGN hosts. The now spaceborne Herschel Space Observatory is ideally suited for fully characterizing these environments via high-J molecular line observations of local IR-bright galaxies and AGN, paving the way for ALMA and the study of such phenomena across cosmic epoch." }, "1003/1003.2467_arXiv.txt": { "abstract": "We compare the near-infrared (NIR) $H$ band photometric and morphological properties of low-$z$ ($z<0.3$) 3CR radio galaxies with samples of BL~Lac object and quasar host galaxies, merger remnants, quiescent elliptical galaxies, and brightest cluster galaxies drawn from the literature. In general the 3CR host galaxies are consistent with luminous ($\\sim L^\\star$) elliptical galaxies. The vast majority of FR~II's ($\\sim80$\\%) occupy the most massive ellipticals and form a homogeneous population that is comparable to the population of radio-loud quasar (RLQ) host galaxies in the literature. However, a significant minority ($\\sim20$\\%) of the 3CR FR~II's appears under-luminous with respect to quasar host galaxies. All FR~II objects in this faint tail are either unusually red, or appear to be the brightest objects within a group. We discuss the apparent differences between the radio galaxy and RLQ host galaxy populations. RLQs appear to require $\\gtsim10^{11}~M_\\odot$ host galaxies (and $\\sim10^{9}~M_\\odot$ black holes), whereas radio galaxies and RQQs can exist in galaxies down to $\\sim 3 \\times10^{10}~M_\\odot$. This may be due to biases in the measured quasar host galaxy luminosities or populations studied, or due to a genuine difference in host galaxy. If due to a genuine difference, it would support the idea that radio and optical active galactic nucleii are two separate populations with a significant overlap. ", "introduction": "\\label{sec-intro} Much effort has been invested over the last decade and a half to characterize the host galaxies of quasars using the {\\it Hubble Space Telescope} ({\\em HST}). Unfortunately by their very nature quasars remain difficult to study, and results remain ambiguous due to selection effects. Detailed isophotal analysis of quasar host galaxies is impossible due to the strongly anisotropic contamination by the nucleus and point-spread function (PSF) of the instrument, which also hinder accurate spectroscopy and thus measurement of host galaxy dynamics. In the absence of any forthcoming space-based coronographic mission (e.g., the Terrestrial Planet Finder Coronograph -- TPF-C~\\footnote{http://planetquest.jpl.nasa.gov/TPF-C/tpf-C\\_index.cfm}), researchers have begun to explore the nature of type 2 (heavily obscured) active galactic nucleii (AGN) to study their environments in greater detail (e.g.,~\\citealt{zakamska+06}). Type 2 AGNs have been identified from the Sloan Digital Sky Survey (SDSS) out to $z=0.83$~\\citep{zakamska+03}, but the sample remains incomplete and biases difficult to measure. Radio galaxies offer us an alternative sample of type 2 radio-loud AGN, detectable to far higher redshifts in a statistically complete manner. In this paper, we compare the overall near-infrared (NIR; $H$ band) galaxy properties of the $z<0.3$ 3CR radio galaxies with those of quasars and BL~Lac objects at similar redshift, and to quiescent early-type galaxies and mergers available in the literature. For the reasons outlined above, we intend to use the radio galaxies as a link or bridge between the quasars and quiescent galaxies which can be studied in much greater structural, dynamical and population detail. Nearby radio galaxies have been statistically studied in recent papers~\\citep{best+05b,mauchsadler07} and are well-known to occupy luminous galaxies in general. Indeed it is by now accepted that the primary requirement for radio-loud AGN activity is a supermassive black hole $\\gtsim 10^{9}~M_\\odot$, which generally requires a very luminous elliptical host galaxy. The principal aim of this paper is to explore in detail the subset of galaxies drawn from the nearby universe that are capable of hosting a radio-loud AGN. We wish to determine what additional effects the host galaxy plays in determining the radio loudness of a source, and to identify any sources that are hosted by ``unusual'' galaxies. A secondary aim is to determine whether there are any systematic biases introduced in the study of quasar host galaxies. Our approach is to analyze the results of isophotal and parametric modeling of the host galaxies of 3CR radio sources from~\\citet{madrid+06} -- hereafter Paper I, and~\\citet{floyd+08} -- hereafter Paper~II, and to compare the results to those for similarly studied samples of galaxies and AGN. In Paper~I we presented the imaging and photometry for the first part of our survey. In Paper~II we presented the modeling used here, and showed that two different techniques (widely adopted by the galaxy morphology and quasar host galaxy communities, respectively) reliably recovered the properties of the NIR host galaxies of the low-redshift ($z<0.3$) 3CR radio sources without strong optical nuclei. Furthermore, for synthetic quasars (created by taking a true galaxy and adding an artificial central point source), the quasar modeling technique was shown to recover the correct underlying morphological parameters to within 10\\%, and galaxy luminosity within to 2\\%. The galaxies associated with the 3CR radio sources have long been known to be typically elliptical~\\citep{matthews+64}, with a number of them associated with the brightest galaxies in a cluster or group~\\citep{burns90,best+07}. Recent surveys with {\\em HST} have revealed a wealth of information on the environments of these uniquely powerful sources at a range of wavelengths~\\citep{dekoff+96,martel+99,dekoff+00,allen+02,madrid+06,privon+08,floyd+08}, and uncovered numerous new jets, dust disks, etc. But until recently~\\citep{donzelli+07,floyd+08}, the 3CR lacked a detailed and systematic classification of their host galaxies in such a way that can be straightforwardly compared with existing samples of quiescent galaxies in the literature. Observing in the infrared offers important advantages in the study of both galaxies and AGNs providing a direct measure of dynamically-dominant old stellar population, an almost-constant mass-to-light ratio~\\citep{zibetti+02}, and a peak in the ratio of galaxy to AGN continuum, giving the best chance of detecting the host galaxy in cases where AGN is unobscured. Thus, it seems a natural application of the NICMOS $H$ band data set to make a direct comparison to similar samples in the literature. We need to be certain which radio galaxies correspond to which RLQs in order to be able to use radio galaxies as a proxy for quasars in detailed spectroscopic, dynamical and structural studies of the effect of AGN on environment and vice versa. This paper is structured as follows. In Section~\\ref{sec-samp} we describe the samples used in this paper, and refer briefly to the observations and data reduction used. In Section~\\ref{sec-res} we describe the results of sample-wide comparisons according to different observational properties of the galaxies. In Section~\\ref{sec-disc} we discuss those comparisons, property-by-property across all the samples studied, and look at the $R_e-\\mu_e$ Kormendy relation, host-to-nuclear luminosity distribution and host luminosity versus extended radio power distributions. We pick out several anomalously faint 3CR host galaxies for discussion in further detail in Section~\\ref{sec-faint}. In Section~\\ref{sec-cont} we attempt to draw this discussion together and look at the issues remaining in the study of quasar host galaxies. Finally we conclude in Section~\\ref{sec-conc} with a summary of our findings and suggestions for future study. ", "conclusions": "\\label{sec-conc} Here, we again summarize the main findings. \\begin{itemize} \\item{\\bf In terms of ellipticity,} the low-$z$ 3CR are found to be in excellent agreement with~\\citet{pahre99} sample of ellipticals drawn from across a range of environments, and with Floyd and Dunlop quasar samples. The mergers within the 3CR (identified as such in Paper~II) fit in well with the general merger population of~\\citet{rothberg+04}. \\item {\\bf In terms of \\sersic index,} there is good agreement with the Floyd et al. quasar and Urry et al. BL~Lac object samples, and very poor agreement with the merger population, even when considering just the 3CR mergers in isolation from the remainder of the 3CR sample. \\item {\\bf In terms of host galaxy luminosity,} the 3CR are generally well matched to the~\\citet{BBF92,pahre99} and ~\\citet{mobasher+99} elliptical galaxy samples, but exhibit far more faint objects than the quasar samples and the BCG's. \\item The Virgo cluster early-type galaxy sample of Ferrarese et al. offers an interesting contrast to the 3CR, clearly illustrating that the latter is drawn from only the most luminous section of a cluster's population. \\item The 3CR and its merging subsample have a similar luminosity range to the mergers of~\\citet{rothberg+04}, but the merging radio galaxies are generally found to be larger, with the Rothberg mergers on the low side of the Kormendy relation. \\item In terms of radio--optical properties, the FR~I's in the 3CR unify well with the properties expected of BL~Lac objects, and FR~II's with RLQs. However, a larger spread is seen in both the morphology and the host galaxy luminosity of the 3CR sources than would be expected from samples of BL~Lac objects and RLQs studied so far. \\end{itemize} We confirm findings from earlier work (e.g.,~\\citealt{ledlowowen96,dunlop+03}) that an elliptical host galaxy is a prerequisite for radio-loud AGN activity. However, we identify several radio galaxies ($\\sim20$\\%) that fall below the observed RLQ host galaxy luminosity cutoff at $\\sim L^\\star$. These objects have a luminosity distribution closer to that of the normal elliptical galaxy population, and were missed in the NVSS-6dFGS survey due to the $K$ band flux limit. The same finding is echoed in the morphological comparison -- the 3CR host galaxies exhibit a range of \\sersic index that is entirely consistent with the general giant elliptical galaxy population, whereas quasar host galaxies show a narrower range of morphology. We conclude that the morphological difference is due to quasar selection effects and / or contamination of quasar host galaxy flux by scattered nuclear flux. It would seem highly probable that given accurate coronographic observations, or a higher resolution image of the central regions of quasars, we would observe the same spread in morphologies in the host galaxies of these objects as well. However, we do not believe that the luminosity discrepancy is so easily explained. The fainter FR~II host galaxies identified here clearly merit more detailed study and form an interesting subset of the 3CR, containing dusty sources and brightest apparent members small groups. While we now have a reasonable view of the quasars of host galaxies, (and an impressively deep view in a small number of cases -- see~\\citealt{bennert+08}), several problems remain. One is the technical issue of separating any scattered nuclear flux from the host galaxy flux. Note that this is distinct from the problem of simply separating out the host galaxy from the PSF. The problem is reviewed by~\\citet{young+09} who argue that the effect is likely to affect the measured luminosities of quasar host galaxies significantly. The extent of the problem can be easily tested using space-based optical polarimetry of a bright quasar. The second issue is one of sample bias. Existing quasar host galaxy studies focus on small samples with likely selection biases toward bright host galaxies. We need a statistical survey of low-redshift quasar host galaxies in at least two bands in order to be able to place meaningful constraints on the masses of these objects to provide a baseline against which to compare higher redshift samples." }, "1003/1003.0654_arXiv.txt": { "abstract": "We present data from the archival plates at Harvard College Observatory and Sonneberg Observatory showing the field of the solar type pre-main sequence star GM Cep. A total of 186 magnitudes of GM Cep have been measured on these archival plates, with 176 in blue sensitivity, 6 in visible, and 4 in red. We combine our data with data from the literature and from the American Association of Variable Star Observers to depict the long-term light curves of GM Cep in both B and V wavelengths. The light curves span from 1895 until now, with two densely sampled regions (1935 to 1945 in B band, and 2006 until now in V band). The long-term light curves do not show any fast rise behavior as predicted by an accretion mechanism. Both the light curves and the magnitude histograms confirm the conclusion that the light curves are dominated by dips (possibly from extinction) superposed on some quiescence state, instead of outbursts caused by accretion flares. Our result excludes the possibility of GM Cep being a FUor, EXor, or McNeil's Nebula type star. Several special cases of T Tauri stars were checked, but none of these light curves are compatible with that of GM Cep. The lack of periodicity in the light curve excludes the possibility of GM Cep being a KH 15D system. ", "introduction": "GM Cep is a solar type variable star in the $\\sim$ 4 Myr-old open cluster Tr 37 \\citep{sic04, sic05}, which is located at a distance of 900 pc \\citep{con02}. The coordinates of GM Cep are $21^h 38^m 16^s.48$ and $+57^{\\circ} 32' 47''.6$. It has a late-type spectral classification of G7V-K0V, with a mass of $\\sim$2.1 $M_\\sun$ and radius estimate of 3 - 6 $R_\\sun$ \\citep{sic08}. A companion star has been hypothesized as part of the physical mechanism for the variability in the GM Cep system, but it has not been seen. The first recorded photometric data for GM Cep was taken at Sonneberg Observatory \\citep{mor39} and showed with the visual magnitude varying from 13.5 to 15.5 mag. \\citet{suy75} showed that GM Cep had a stable period of up to $\\sim$100 days, and it was experiencing rapid variation between 14.2 and 16.4 mag. \\citet{sic08} listed and summarized the available data in the literature and depicted a long term light curve in multiple wavelengths. This list contains 16 magnitude, in V band and 5 in B band, most of which were taken in 2006 and later. The only one B band magnitude before 2006 was taken from \\citet{kun86}, with B = 17.31 mag. It is much fainter than any other available B magnitude values, and the simultaneous V magnitude is not significantly high. \\citet{sic08} took it as an outlier and did not include it in their analysis. Although the data for GM Cep in the literature span from 1939 until 2007, the time history is rather spotty, and there are few magnitudes before 2000. \\citet{sic08} invoked several possible mechanisms to explain the large rapid variability of GM Cep's optical magnitude, the fast rotation rate, and the strong mid-IR excesses. The rapid variability can be explained by the strong outbursts of FUor systems (which brighten by $\\geqslant$ 4 mag), in which the mass accretion rate through the circumstellar disk of a young star increases by orders of magnitude \\citep{har96}. Another proto-stellar system, EXor (with outbursts $\\geqslant$ 2 mag), was also interpreted as a mass accretion event \\citep{leh95} and proposed to be similar to GM Cep. \\citet{sic08} also give comparisons between the observational features of GM Cep and several better-known systems. For example, RW Aur, which is often quoted as a triple system, shares the features of a strong and variable P Cygni H$\\alpha$ profile, a powerful disk, a large accretion rate, and a strong double-peaked $OI$ emission line with GM Cep \\citep{ghe93, ale05, suy75}. Another similar system is GW Ori, a 1-Myr old G5 star with a fast rotation rate of $V\\sin i$ = 43 km s$^{-1}$ \\citep{bou86}, variability up to 1 mag in JHK\\footnote{VizieR Online Data Catalog, II/250 \\citep{sam04}}, and strong IR excess \\citep{mat91, mat95}. CW Tau, a K3 star, has large magnitude variations of 2 mag\\footnotemark[\\value{footnote}] , a rapid rotation rate of $V\\sin i$ = 28 km s$^{-1}$ \\citep{muz98}, a P Cygni H$\\alpha$ profile, and a deep, broad $OI$ absorption at 7773 \\AA. McNeil's Nebula \\citep{mcn04} has its emission line spectrum at optical wavelengths similar to the spectrum of GM Cep. KH 15D, a pre-main sequence binary system with a precessing disk or ring \\citep{ham05}, is another system that provides an example of a possible explanation for the mechanism of GM Cep. However, without a long-term light curve of GM Cep, these physical explanations cannot be properly tested, and the observational comparisons cannot be made. A long-term light curve can be used to search for outbursts, periodicities, repetitive features, and other observational features that these mechanisms predict. To obtain a long-term light curve, we visited Harvard College Observatory and Sonneberg Observatory, searched through the archival plates showing this field, and obtained 186 magnitude estimates from 1895 until 1993. We also collected the 75 visual observations from the database of the American Association of Variable Star Observations (AAVSO) from 2006 to present. A long-term light curve for GM Cep was plotted from these data. ", "conclusions": "In this paper, we present data of GM Cep from all available series archival plates and some patrol plates from Harvard College Observatory and Sonneberg Observatory. We obtained 186 new magnitudes for GM Cep (176 in blue, 6 in visible, and 4 in red) ranging from 1895 until 1993. Another 75 V band magnitudes were drawn from the AAVSO database. By combining our data from archival plates, AAVSO data, and previously-published data collected by \\citet{sic08}, long term B and V band light curves for GM Cep were constructed. The B band light curve shows a generally constant magnitude ($\\sim$14-14.5) with occasional dips to $\\sim$16.5. Fast variations are found in both the B and V band light curves. The magnitude histograms in both B and V bands show cut-offs at the low magnitude (bright) end and long extended tails at the high magnitude (faint) end, which implies that the light curve is composed of dips caused by varying extinction instead of outbursts caused by accretion superposed on quiescence state. The lack of large outbursts in the past century implies that it is not a FUor or EXor star, or a McNeil's Nebula type star. The lack of periodicity in the light curve also excludes the possibility of GM Cep being a KH 15D type star. Several special cases of T Tauri stars (RW Aur, GW Ori and CW Tau) were checked, but none of these light curves are compatible with that of GM Cep." }, "1003/1003.4947_arXiv.txt": { "abstract": "{\\large The variations of gravity were measured with a high precision LaCoste-Romberg D gravimeter during a total solar eclipse to investigate the effect of solar eclipse on the gravitational field. The observed anomaly ($7.0 \\pm 2.7$) \\mss during the eclipse implies that there may be a shielding property of gravitation. } ", "introduction": "Although gravitation may has the property of shielding in theories, it is very difficulty to test the possible effect experimentally. If gravitation were carried by particles, a mass between two bodies could partially shield each of them from the gravity of the other. Anomalies can be expected in the motions of certain artificial Earth satellites during eclipse seasons that behave like shielding of the Sun's gravity as suggested by VanFlandern [1]. The possible existence of gravitional shielding and gravitational-wave absorption [2] and some theoretical analysis of a weak shielding of the gravitational interaction by a disk of high temperature superconducting materials have been investigated [3,4,5]. An experiment of electrically charged pendulum [6] was carried out during an eclipse to test the Saxl's effect [7] although there was no noticeable effect observed. Some related work were reviewed by Gillies [8]. If there were gravitational shielding, it would expect that the effect shall be only significant during an eclipse when gravity of the Sun may be shieldly slightly by the moon so that the gravity on the Earth may fluctuate accordingly, however such effect may be extremely small even if it would exist. The present work was thus motivated to test the possible effect of gravitational shielding during the total solar eclipse with a high precision modern gravimeter. ", "conclusions": "The vertical gravitational acceleration measured consists of several components: 1) gravitional forces due to the Earth, the Sun and the Moon, and 2) the earth's rotation. The former includes the static gravity by the Earth and the tidal force by the Sun and the Moon due to changes of moving positions. The tidal component can be calculated theoretically with a precision of 1 \\ug \\s or 1 \\mss, which is a routine practice in geophysics. After making all these corrections, the difference left shown Figure 1) is the variation of vertical gravity during the eclipse due to some unknown effect, which may be a possible shielding effect of gravitation. The solid curve is the averaged values with a 10 minute window and the variation can be more clearly identified. The variation around zero has an amplitude of $\\pm 3 \\sim 4$ \\ug. The important and interesting anomaly is that there exists two regions with significant gravity decrease. One of such region occurred within about 30 mins around 07:30am with a maximum significant decrease of $6.0 \\pm 2.5$ \\ug, and another took place within 30 mins around 10:20am with a maximum change of $7.0 \\pm 2.7$ \\ug. The deviation is calculated by using the standard formulae in measurement data processing. If the solid curve is used for the calculation, the maximum changes shall be $5.3 \\pm 1.4$ \\ug at fist contact and $6.8 \\pm 1.4$ \\ug at fourth contact, respectively. These two changes took place between first contact and fourth contact, and quite closely related to the timing of eclipse phases of first contact and fourth (last) contact. Figure 2 shows the measured gravity variation in the week of the eclipse from 5 March 1997 to 12 March 1997. The significant variation during the eclipse on 9 March 1997 is also shown (detail see Figure 1). In plotting this figure, the data was averaged with a 10 minute moving window so that the curve is more smooth than the actual measured data and the signal looks more significant. We can see that the reading was quite stable before the eclipse and after the eclipse. The change during the eclipse is remarkable. Table I shows the number of data deviated from the average value with a total of 10,080 data. Please note that the actual number of data during the eclipse is much more than those listed this table (with a resampling rate of 1 reading per minute) because the sampling rate during the eclipse is much higher (1 reading per second). The changes are quite significant and they are not the effect of temperature and pressure changes. According to the calibration precision of the LaCoste-Romberg gravimeter provided by the manufacturer, the variation of $8^{o}$C in temperature would lead to 5 \\ug change in gravity reading. The actual temperature change in controlled room temperature during the eclipse is within $\\pm 1^{o}$C, so the actual effect of temperature change is less than 1 \\ug. The actual change in pressure during eclipse from 07:00am to 11:00am is about 1 mmH and the change is less than 3 mmH in that whole day. According to the manufacturer, the effect of actual pressure change on gravity reading shall much less than 1 \\ug. Therefore, the actual noticeable changes of gravity during the eclipse may imply some extra-ordinary phenomenon associated with gravity such as the possible shielding effect of moon on the gravitational force of the Sun. In addition, another puzzle is that the anomalies of the gravity variations occurred at the first and last contact but not during the totality. This certainly requires more precise measurements in the future during totality of a solar eclipse. \\\\ \\def\\tab{\\hspace{2.5in}} Table I: Measured Data Distribution \\\\ --------------------------------------------------------------- \\\\ Data deviation range (\\ug) \\,\\,\\,\\,\\ Number of Data \\\\ --------------------------------------------------------------- \\\\ $< $ 2 \\tab 9948 \\\\ $\\ge$ 2 \\,\\,\\, \\tab 87 \\\\ $\\ge$ 4 \\,\\,\\, \\tab 45 \\\\ --------------------------------------------------------------- \\\\ In summary, we have used the best available gravimeter, with a high precision of 2 $\\sim$ 3 \\ug, to measure the variation of vertical gravity during the total eclipse on 9 March 1997. Although there was no noticable changes around the totality during the solar eclipse, we have observed quite significant decrease in vertical gravity during the first contact and the last contact. The may imply the new property of gravitation, which certainly needs more high precision experiments to be conducted in the future especially during solar eclipse. Although the purpose of this short paper and the present work is not intended to prove the shielding effect of gravitation, however, we would be delighted if the present work can initiate more work on the possible new property of gravitation. {\\bf Acknowledgement}: We would thank the referee(s) for their insightful comments which has greatly improved the manuscript, especially for the kind suggestion of averaging the data over the 10 minute interval. The work was supported by the National Natural Science Foundation of China. The authors are grateful to the help from the Moho geophysical station of Chinese Academy of Sciences." }, "1003/1003.1753_arXiv.txt": { "abstract": "We report observations of the linear polarization of a sample of 49 nearby bright stars measured to sensitivities of between $\\sim$1 and $\\sim$4 $\\times 10^{-6}$. The majority of stars in the sample show measurable polarization, but most polarizations are small with 75\\% of the stars having P $<$ 2 $\\times10^{-5}$. Correlations of the polarization with distance and position, indicate that most of the polarization is of interstellar origin. Polarizations are small near the galactic pole and larger at low galactic latitudes, and the polarization increases with distance. However, the interstellar polarization is very much less than would be expected based on polarization-distance relations for distant stars showing that the solar neighbourhood has little interstellar dust. BS 3982 (Regulus) has a polarization of $\\sim$ 37 $\\times 10^{-6}$, which is most likely due to electron scattering in its rotationally flattened atmosphere. BS 7001 (Vega) has polarization at a level of $\\sim$ 17 $\\times 10^{-6}$ which could be due to scattering in its dust disk, but is also consistent with interstellar polarization in this direction. The highest polarization observed is that of BS 7405 ($\\alpha$ Vul) with a polarization of 0.13\\% ", "introduction": "Linear polarization of starlight provides a powerful technique for investigating the nature of the interstellar medium. Interstellar dust particles aligned to the galactic magnetic field produce interstellar polarization, which is one of the main sources of stellar linear polarization. Studies of this polarization provide information on the dust distribution and magnetic field structure \\citep[e.g.][]{heiles96} and on the nature and size of the dust particles \\citep{whittet92,kim94}. Studies of the polarization of stars close to the Sun have been made by \\citet{piirola77}, \\citet{tinbergen82} and \\citet{leroy93a,leroy93b,leroy99}. They found very little polarization in neraby stars. \\citet{leroy93b} found only 25 stars with definite polarization in a survey of 1000 stars within 50pc. Subsequent analysis showed that almost all of these polarized stars were actually at greater distances when more accurate Hipparcos parallaxes became available \\citep{leroy99}, and that significant interstellar polarization became detectable at distances of about 70pc in some directions and at 150pc in others. \\citet{andersson06} have also reported observations of polarization of southern hemisphere stars, that they attribute to the wall of the Local Bubble at $\\sim$100 pc distance. All of these studies used polarization measurements with accuracies of, at best, $\\sim$10$^{-4}$ in fractional polarization. Recently we have built and tested a new polarimeter, PlanetPol, \\citep{hough06} capable of measuring stellar linear polarization at the parts per million level. Here we report observations of a sample of nearby bright stars measured to sensitivies of generally better than 3 $\\times$ 10$^{-6}$ and in some cases to better than 1 $\\times 10^{-6}$ in fractional polarization. This represents an improvement of a factor of 20 to 100 on previous measurements. In addition to the use of polarization to probe the interstellar medium, it is of interest to know at what level normal stars show intrinsic polarization. There is currently considerable interest in using polarization to study extrasolar planets. Polarization can be used to directly detect unresolved hot-Jupiter type planets \\citep{seager00,lucas06,lucas09}, as a differential technique to detect planets in imaging observations \\citep{schmid05,keller06}, and as a means of characterizing extrasolar planet atmospheres \\citep{bailey07}. Significant polarization from the host star could complicate such observations. In the case of the quiet Sun direct observations by \\citet{kemp87} show linear polarization of $<3 \\times 10^{-7}$. However, more active stars could show higher polarizations, and polarization might also result from exozodiacal disks around the stars. \\begin{table*} \\caption{Properties of Sample Stars - The horizontal line marks the 17 hour division used in figure 4} \\begin{flushleft} \\begin{tabular}{llrlrrrrrr} \\hline BS & Other Names & V & Spectral & Dist & RA & Dec & \\multicolumn{2}{c}{Galactic} & $v \\sin{i}$ \\\\ & & mag & Type & (pc) & hh:mm & dd:mm & Long & Lat & km$s^{-1}$ \\\\ \\hline 3982 & Regulus, $\\alpha$ Leo & 1.35 & B7V & 23.8 & 10 08 & +11 58 & 226.4 & 48.9 & 353 \\\\ 4031 & & 3.44 & F0III & 79.6 & 10 16 & +23 25 & 210.2 & 55.0 & 83 \\\\ 4069 & & 3.07 & M0III & 76.3 & 10 22 & +41 30 & 177.9 & 56.4 & \\\\ 4295 & Merak, $\\beta$ Uma & 2.35 & A1V & 24.3 & 11 01 & +56 23 & 149.2 & 54.8 & 32 \\\\ 4301 & Dubhe, $\\alpha$ Uma & 1.79 & K0Iab & 37.9 & 11 03 & +61 45 & 142.8 & 51.0 & $<$17 \\\\ 4335 & & 3.01 & K1III & 45.0 & 11 09 & +44 30 & 165.8 & 63.2 & 10 \\\\ 4357 & & 2.56 & A4V & 17.7 & 11 14 & +20 31 & 224.2 & 66.8 & 180 \\\\ 4359 & & 3.32 & A2V & 54.5 & 11 14 & +15 26 & 235.4 & 64.6 & 5 \\\\ 4518 & & 3.71 & K0.5IIIb & 60.1 & 11 46 & +47 47 & 150.3 & 28.4 & 10 \\\\ 4527 & & 4.54 & A7V & 69.4 & 11 48 & +20 13 & 235.0 & 73.9 & \\\\ 4534 & $\\beta$ Leo & 2.14 & A3V & 11.1 & 11 49 & +14 34 & 250.6 & 70.8 & 110 \\\\ 4540 & $\\beta$ Vir & 3.61 & F9V & 10.9 & 11 51 & +01 46 & 270.5 & 60.8 & 3 \\\\ 4905 & & 1.76 & A0p & 24.8 & 12 54 & +55 58 & 122.2 & 61.2 & 33 \\\\ 4910 & & 3.38 & M3III & 62.1 & 12 56 & +03 24 & 305.5 & 66.2 & \\\\ 4915 & $\\alpha^2$ CVn & 2.90 & A0p & 33.8 & 12 56 & +38 19 & 118.3 & 78.8 & \\\\ 4932 & & 2.83 & G8III & 31.3 & 13 02 & +10 58 & 312.3 & 73.6 & 8 \\\\ 5054 & Mizar & 2.27 & A2V & 24.0 & 13 24 & +54 56 & 113.1 & 61.6 & 13 \\\\ 5191 & & 1.85 & B3V & 30.9 & 13 47 & +49 19 & 100.7 & 65.3 & 226 \\\\ 5235 & & 2.68 & G0IV & 11.3 & 13 55 & +18 24 & 5.3 & 73.0 & 18 \\\\ 5340 & Arcturus, $\\alpha$ Boo & $-$0.04 & K1.5III & 11.3 & 14 16 & +19 11 & 15.1 & 69.1 & 8 \\\\ 5429 & & 3.58 & K3III & 45.6 & 14 32 & +30 22 & 47.3 & 67.8 & 8 \\\\ 5435 & & 3.00 & A7III & 26.1 & 14 32 & +38 18 & 67.3 & 66.2 & 135 \\\\ 5563 & & 2.08 & K4III & 38.8 & 14 51 & +74 09 & 112.6 & 40.5 & 8 \\\\ 5793 & $\\alpha$ CrB & 2.21 & A0V & 22.9 & 15 35 & +26 43 & 41.9 & 53.8 & 132 \\\\ 5849 & & 3.84 & B9IV & 44.5 & 15 43 & +26 18 & 41.7 & 51.9 & 100 \\\\ 5854 & & 2.64 & K2IIIb & 22.5 & 15 44 & +06 26 & 14.2 & 44.1 & 8 \\\\ 6092 & & 3.74 & B5IV & 96.4 & 16 20 & +46 18 & 72.5 & 45.0 & 20 \\\\ 6095 & & 3.74 & A9III & 59.9 & 16 22 & +19 09 & 35.3 & 41.3 & 135 \\\\ 6148 & & 2.79 & G7IIIa & 45.3 & 16 30 & +21 29 & 39.0 & 40.2 & 10 \\\\ 6149 & & 3.90 & A0V & 50.9 & 16 31 & +01 59 & 17.1 & 31.8 & 142 \\\\ 6212 & & 2.89 & G0IV & 10.8 & 16 41 & +31 36 & 52.7 & 40.3 & 5 \\\\ 6299 & & 3.20 & K2III & 26.3 & 16 58 & +09 23 & 28.4 & 29.5 & 8 \\\\ \\hline 6324 & & 3.91 & A0V & 49.9 & 17 00 & +30 56 & 52.9 & 36.2 & 60 \\\\ 6410 & & 3.13 & A3IV & 24.1 & 17 15 & +24 50 & 46.8 & 31.4 & 305 \\\\ 6556 & $\\alpha$ Oph & 2.10 & A5III & 14.3 & 17 35 & +12 34 & 35.9 & 22.6 & 240 \\\\ 6603 & & 2.77 & K2III & 25.1 & 17 43 & +04 34 & 29.2 & 17.2 & 8 \\\\ 6623 & & 3.42 & G5IV & 8.4 & 17 46 & +27 43 & 52.4 & 25.6 & 8 \\\\ 6629 & & 3.75 & A0V & 29.1 & 17 48 & +02 42 & 28.0 & 15.4 & 212 \\\\ 6688 & & 3.74 & K2III & 34.2 & 17 54 & +56 52 & 85.2 & 30.2 & 8 \\\\ 6703 & & 3.71 & G8III & 41.5 & 17 58 & +29 15 & 54.9 & 23.8 & 10 \\\\ 6705 & & 2.23 & K5III & 45.2 & 17 57 & +51 29 & 79.1 & 29.2 & 8 \\\\ 6872 & & 4.32 & K2III & 70.4 & 18 20 & +36 04 & 63.5 & 21.5 & 8 \\\\ 7001 & Vega, $\\alpha$ Lyr & 0.03 & A0V & 7.8 & 18 37 & +38 47 & 67.4 & 19.2 & 5 \\\\ 7235 & & 3.00 & A0V & 25.5 & 19 05 & +13 52 & 46.9 & 3.2 & 360 \\\\ 7405 & $\\alpha$ Vul & 4.45 & M0III & 90.9 & 19 29 & +24 40 & 59.0 & 3.4 & \\\\ 7528 & & 2.90 & B9.5IV & 52.4 & 19 45 & +45 08 & 78.7 & 10.2 & 128 \\\\ 7557 & Altair, $\\alpha$ Aql & 0.77 & A7V & 5.1 & 19 51 & +08 52 & 47.7 & $-$8.9 & 245 \\\\ 7582 & & 3.83 & G8III & 44.6 & 19 48 & +70 16 & 102.4 & 20.8 & 10 \\\\ 7635 & & 3.53 & M0III & 84.0 & 19 59 & +19 30 & 58.0 & $-$5.2 & 8 \\\\ \\hline \\end{tabular} \\end{flushleft} \\label{tab_stars} \\end{table*} \\begin{table*} \\caption{Previous Polarization Measurements} \\begin{flushleft} \\begin{tabular}{lcrrrr} \\hline BS & Polarization (Heiles) & \\multicolumn{2}{c}{Polarization (Tinbergen)$^a$} & \\multicolumn{2}{c}{Polarization (Piirola)$^a$}\\\\ & \\multicolumn{1}{c}{P(\\%)} & Q/I & U/I & Q/I & U/I \\\\ \\hline 3982 & 0.060$\\pm$0.120 & 0$\\pm$7 & -4$\\pm$7 & 7$\\pm$11 & 3$\\pm$ 11\\\\ 4031 & 0.050$\\pm$0.120 & & & & \\\\ 4295 & 0.000$\\pm$0.120 & 4$\\pm$7 & 6$\\pm$7 & 13$\\pm$6 & 3$\\pm$6 \\\\ 4301 & 0.040$\\pm$0.120 & 3$\\pm$7 & 22$\\pm$7 & $-$3$\\pm$13 & $-$2$\\pm$13 \\\\ 4335 & 0.030$\\pm$0.120 & 22$\\pm$7 & $-$14$\\pm$7 & &\\\\ 4357 & 0.000$\\pm$0.120 & & & 4$\\pm$7 & $-$6$\\pm$7 \\\\ 4359 & 0.010$\\pm$0.120 & $-$1$\\pm$7 & $-$11$\\pm$7 & & \\\\ 4518 & 0.060$\\pm$0.120 & $-$33$\\pm$7 & 1$\\pm$7 & & \\\\ 4534 & 0.030$\\pm$0.120 & 12$\\pm$7 & 3$\\pm$7 & 0$\\pm$14 & $-$12$\\pm$14 \\\\ 4540 & 0.042$\\pm$0.026 & 4$\\pm$7 & 3$\\pm$7 & & \\\\ 4905 & 0.010$\\pm$0.120 & & & & \\\\ 4910 & 0.020$\\pm$0.120 & & & & \\\\ 4915 & 0.020$\\pm$0.120 & & & & \\\\ 4932 & 0.010$\\pm$0.120 & 6$\\pm$7 & $-$6$\\pm$7 & & \\\\ 5191 & 0.060$\\pm$0.000 & & & & \\\\ 5235 & 0.007$\\pm$0.012 & $-$6$\\pm$7 & $-$12$\\pm$7 & $-$6$\\pm$9 & 3$\\pm$9 \\\\ 5340 & 0.030$\\pm$0.120 & $-$1$\\pm$7 & $-$6$\\pm$7 & $-$10$\\pm$8 & $-$11$\\pm$8 \\\\ 5429 & 0.030$\\pm$0.120 & & & & \\\\ 5435 & 0.000$\\pm$0.120 & & & & \\\\ 5563 & 0.100$\\pm$0.120 & 6$\\pm$7 & $-$1$\\pm$7 & & \\\\ 5793 & 0.060$\\pm$0.120 & & & 14$\\pm$6 & $-$1$\\pm$6 \\\\ 5849 & 0.030$\\pm$0.120 & & & & \\\\ 5854 & 0.030$\\pm$0.120 & 1$\\pm$7 & 0$\\pm$7 & & \\\\ 6092 & 0.010$\\pm$0.000 & & & & \\\\ 6095 & 0.057$\\pm$0.012 & & & & \\\\ 6149 & 0.010$\\pm$0.120 & & & & \\\\ 6212 & 0.000$\\pm$0.200 & & & & \\\\ 6410 & 0.020$\\pm$0.120 & & & & \\\\ 6556 & 0.010$\\pm$0.100 & $-$4$\\pm$7 & $-$6$\\pm$6 & & \\\\ 6603 & 0.150$\\pm$0.120 & & & & \\\\ 6629 & 0.008$\\pm$0.001 & 12$\\pm$7 & 15$\\pm$7 & & \\\\ 6688 & & $-$36$\\pm$7 & $-$41$\\pm$7 & & \\\\ 7001 & 0.020$\\pm$0.120 & 11$\\pm$7 & 7$\\pm$7 & 4$\\pm$4 & 6$\\pm$4 \\\\ 7528 & 0.030$\\pm$0.120 & & & & \\\\ 7557 & 0.016$\\pm$0.002 & 15$\\pm$7 & 0$\\pm$7 & 2$\\pm$6 & $-$7$\\pm$6 \\\\ \\hline \\end{tabular} a - Polarizations in units of 10$^{-5}$ \\end{flushleft} \\label{tab_ppol} \\end{table*} ", "conclusions": "Polarization measurements of a sample of 49 nearby bright stars have been measured to accuracies about 20 to 100 times better than those of any previous measurements. In contrast to previous observations which have generally been unable to detect many polarized stars at these distances, we find significant polarization in many of the stars. The polarization increases with distance and shows much higher values at low galactic latitudes than towards the galactic pole. The distribution of polarization strongly suggests that the high polarization stars and probably most of the lower polarization stars, are showing interstellar polarization. The results indicate that polarization measured at the parts per million level provides a very sensitive probe of the interstellar medium in the solar vicinity. The polarization observed near the Sun is much less than would be expected based on the polarization of distant stars, thus confirming the presence of the local cavity or bubble seen in absorption line measurements and in the soft X-ray background. Polarization shows litle correlation with CaII absorption due to warm interstellar gas. The data is not consistent with the hypothesis of \\cite{frisch05} that polarization in nearby stars is due to interstellar dust entrained in the heliospere. Regulus shows a larger polarization than expected for its position. Regulus is known to be a rapidly rotating star, and the polarization direction agrees with the minor axis of the rotational flattening as measured by interferometry. The polarization is reasonably consistent with that expected due to electron scattering in the atmosphere of the flattened star." }, "1003/1003.4801_arXiv.txt": { "abstract": "% The Virtual Observatory is a new technology of the astronomical research allowing the seamless processing and analysis of a heterogeneous data obtained from a number of distributed data archives. It may also provide astronomical community with powerful computational and data processing on-line services replacing the custom scientific code run on user's computers. Despite its benefits the VO technology has been still little exploited in stellar spectroscopy. As an example of possible evolution in this field we present an experimental web-based service for disentangling of spectra based on code KOREL. This code developed by P. Hadrava enables Fourier disentangling and line-strength photometry, i.e. simultaneous decomposition of spectra of multiple stars and solving for orbital parameters, line-profile variability or other physical parameters of observed objects. We discuss the benefits of the service-oriented approach from the point of view of both developers and users and give examples of possible user-friendly implementation of spectra disentangling methods as a standard tools of Virtual Observatory. ", "introduction": "% The astronomical spectroscopy uses many special techniques to analyse stellar spectra and estimate physical properties of targets studied. Basically they consist in comparison of the observed spectra with theoretical models which, however, may be of very different level of sophistication. For instance, a simple comparison of suitably defined effective centres of spectral lines with their laboratory wavelengths gives Doppler shifts, which in the case of spectroscopic binaries enables one to determine their orbital parameters. Detailed comparison of equivalent widths and shapes of line profiles with synthetic spectra may reveal effective temperatures, gravity acceleration, abundances and other physical parameters of stellar atmospheres. In practice, however, the spectra of components of the binary are blended and the information on orbital and atmospheric parameters are entangled. Several techniques for separation of component spectra from a series of spectra has been proposed which enable also to develop the so called spectra disentangling, i.e. a method of simultaneous separation of the spectra and determination of physical parameters governing their variability. In particular, the method of Fourier disentangling introduced and implemented in program KOREL by \\citet{h95} proved to be efficient and viable for a further generalisation. To allow the application of such a powerful method on a number of different objects in a scalable way, we attempted to embed the KOREL in the infrastructure of Virtual Observatory. ", "conclusions": "The Fourier disentangling is already well-established method of stellar spectra analysis with the wide range of applications. The KOREL web service is probably one of the first attempts to adapt the legacy stellar spectra analysis code for the Virtual Observatory service. The advantages of solution adopted are evident, although some level of user conservatism has to be expected." }, "1003/1003.2420_arXiv.txt": { "abstract": "We develop an optimized technique to extract density--density and velocity--velocity spectra out of observed spectra in redshift space. The measured spectra of the distribution of halos from redshift distorted mock map are binned into 2--dimensional coordinates in Fourier space so as to be decomposed into both spectra using angular projection dependence. With the threshold limit introduced to minimize nonlinear suppression, the decomposed velocity--velocity spectra are reasonably well measured up to scale $k=0.07\\hompc$, and the measured variances using our method are consistent with errors predicted from a Fisher matrix analysis. The detectability is extendable to $k\\sim 0.1\\hompc$ with more conservative bounds at the cost of weakened constraint. ", "introduction": "The evolution of large scale structure, as revealed in the clustering of galaxies observed in wide--deep redshift surveys has been one of key cosmological probes. Structure formation is driven by a competition between gravitational attraction and the expansion of space-time, which enables us to test our model of gravity at cosmological scales and the expansion of history of the Universe~\\citep{Wang:2007ht,Linder:2007nu,Guzzo:2008ac,2009JCAP...10..004S,Simpson:2009zj,Guzik:2009cm,McDonald:2008sh,Stril:2009ey,Bean:2010zq}. Maps of galaxies where distances have been measured from redshifts show anisotropic deviations from the true galaxy distribution \\citep{2000AJ....120.1579Y,2001Natur.410..169P,2003astro.ph..6581C,2003MNRAS.346...78H,2004MNRAS.353.1201P,2005ApJ...630....1Z,2005AA...439..877L,2006PhRvD..74l3507T,2008ApJ...676..889O,2008arXiv0812.2480G,2008AA...486..683G,Guzzo:2008ac}, because galaxy recession velocities include components from both the Hubble flow and peculiar velocities. In linear theory, a distant observer should expect a multiplicative enhancement of the overdensity field of tracers due to the peculiar motion along the line of sight \\citep{Davis:1982gc,1987MNRAS.227....1K,1989MNRAS.236..851L,1990MNRAS.242..428M,1991MNRAS.251..128L,1992ApJ...385L...5H,1994MNRAS.266..219F,1995ApJ...448..494F}. In principle, the observed spectra in redshift space can be decomposed into both density--density and velocity--velocity spectra using angular projection dependence~\\citep{2009JCAP...10..004S,Percival:2008sh,2009MNRAS.397.1348W,Song:2010vh}. With a local linear bias, the real-space galaxy density field is affected, while the peculiar velocity term is not. In this paper, we attempt to extract velocity--velocity spectra as an unbiased tool to trace the history of structure formation. A theoretical formalism~\\citep{2009MNRAS.397.1348W} was derived for forecasting errors when extracting velocity--velocity spectra out of the observed redshift space distortion maps. However, it is not yet fully understood what the optimal technique is to practically decompose the spectra as theory predicts. We propose a statistical technique to extract it up to the limit of theoretical estimation. Our method utilizes the distinct angular dependence of density--density and velocity--veclocity spectra to decompose them from two--dimensional redshift power spectra, and is consistent with the theoretical estimate from Fisher matrix analysis. We present the detailed formalism in the next section. The Fisher matrix analysis to decompose spectra is briefly reviewed, then we present the method to decompose spectra in an optimal way with mock data. We discuss statistical method to minimize the effect by nonlinear suppression. ", "conclusions": "We propose a statistical tool to decompose $P_{gg}(k)$ and $P_{\\Theta\\Theta}(k)$ practically out of redshift distortion maps, with a few assumptions: 1) perfect correlation between density and velocity fluctuations, 2) confidence on theoretical prediction of velocity dispersion effect within threshold limit. The results show that the true value of velocity--velocity spectra up to $k=0.07\\hompc$ are successfully recovered using theoretical dispersion effect. The detectability is extendable up to $k\\sim 0.1\\hompc$ with more conservative threshold limit at the cost of weakened constraint. We find that the theoretical dispersion effect can be estimated from $P_{\\Theta\\Theta}(k)$ parameters using weighted average at $k<0.1\\hompc$. In linear regime, $P_{\\Theta\\Theta}(k)$ is well--measured with this estimated $\\sigma_v$ as much as with the true fixed $\\sigma_v$ of the simulation. We find that the biased measurement of $P_{\\Theta\\Theta}(k)$ is mainly caused by the unpredictable non--linear supression effect at $k>0.1\\hompc$. The detectability limit in scale can be extended by parameterizing this effect \\citep{Tang10}, but we scope our range of interest in linear regime in this paper." }, "1003/1003.0049_arXiv.txt": { "abstract": "Recent spectro-polarimetric observations of a sunspot showed the formation of bipolar magnetic patches in the mid penumbra and their propagation toward the outer penumbral boundary. The observations were interpreted as being caused by sea-serpent magnetic fields near the solar surface \\citep{dalda08}. In this Letter, we develop a 3D radiative MHD numerical model to explain the sea-serpent structure and the wave-like behavior of the penumbral magnetic field lines. The simulations reproduce the observed behavior, suggesting that the sea-serpent phenomenon is a consequence of magnetoconvection in a strongly inclined magnetic field. It involves several physical processes: filamentary structurization, high-speed overturning convective motions in strong, almost horizontal magnetic fields with partially frozen field lines, and traveling convective waves. The results demonstrate a correlation of the bipolar magnetic patches with high-speed Evershed downflows in the penumbra. This is the first time that a 3D numerical model of the penumbra results in downward directed magnetic fields, an essential ingredient of sunspot penumbrae that has eluded explanation until now. ", "introduction": "It is well-known that the sunspot penumbra (the outer part of sunspots) has a very complicate filamentary structure and a strong non-stationary outflow. This outflow is responsible for the so-called Evershed effect, a Doppler shift of the spectral lines emerging from sunspots \\citep{evershed1909}. The magnetic structure of the penumbra can be represented as a mixture of two magnetic field components with different inclinations and strengths \\citep[e.g.,][]{degenhardt91,schmidt92,title93,lites93, stanchfield97,bellot04,sanchez05,borrero05,beck08}. However, the radial Evershed flow (in particular, the high-speed 'Evershed coluds') is associated with the more strongly inclined, almost horizontal field \\citep[e.g.][]{title93,shine1994,bellot03}. Recently, significant progress in our understanding of the Evershed effect was made by numerical simulations \\citep{heinemann2007, scharmer2008,rempel2009,rempel_sci09,kiti09a}. These studies suggested that the Evershed flow is a consequence of overturning magnetoconvection in the presence of inclined magnetic fields, and that the driving mechanism is associated with traveling convective waves that propagate in the direction of the magnetic field inclination \\citep{hurlburt2000,kiti09a}. The issue is not resolved and other interpretations are possible \\citep[see][]{schliche09}. Thus, it is important to confront the simulations with as many observations as possible. In this Letter, we examine the idea that the sea-serpent penumbral field lines detected by \\cite{dalda08} in high-resolution Hinode measurements are related to the same mechanism of overturning convection and traveling convective waves in a strong inclined magnetic field. Our analysis is based on the radiative MHD simulations of \\citet{kiti09a}. ", "conclusions": "In this Letter, we have used the results of numerical simulations of magnetoconvection in strong inclined magnetic field to interpret polarimetric observations of a sunspot penumbra. The results reproduce the moving bipolar magnetic elements observed in high-resolution SOHO/MDI and Hinode/SOT data and also their properties, supporting the sea-serpent model proposed by \\cite{dalda08}. The simulations explain the sea-serpent structure and dynamics of the penumbral field as a consequence of solar magnetoconvection in a highly inclined, strong magnetic field, which forms filamentary structures and has properties of traveling convective wave. The physical picture schematically illustrated in Figure~\\ref{scheme} is the following. Convective cells in sunspot penumbrae are deformed under the action of the inclined magnetic field, forming filamentary structures and producing high-speed Evershed flows \\citep{kiti09a}. The magnetic field lines are stretched by the downward flows and dragged under the surface. The points where the magnetic field lines cross the solar surface are observed as magnetic patches of positive and negative polarities. Note that the negative patch is closer to the umbra, in agreement with the observations. The convective cells move in the direction of the magnetic field inclination because of the traveling convective wave behavior. Therefore, the bipolar magnetic patches also move in the same direction. Thus, the numerical simulations connect the sea-serpent structure of the moving bipolar magnetic pathes observed in the penumbra with the process of overturning magnetoconvection, traveling convective waves, and the Evershed flow." }, "1003/1003.0878_arXiv.txt": { "abstract": "In the study of Planck-scale (``quantum-gravity induced\") violations of Lorentz symmetry, an important role was played by the deformed-electrodynamics model introduced by Myers and Pospelov. Its reliance on conventional effective quantum field theory, and its description of symmetry-violation effects simply in terms of a four-vector with nonzero component only in the time-direction, rendered it an ideal target for experimentalists and a natural concept-testing ground for many theorists. At this point however the experimental limits on the single Myers-Pospelov parameter, after improving steadily over these past few years, are ``super-Planckian\", {\\it i.e.} they take the model out of actual interest from a conventional quantum-gravity perspective. In light of this we here argue that it may be appropriate to move on to the next level of complexity, still with vectorial symmetry violation but adopting a generic four-vector. We also offer a preliminary characterization of the phenomenology of this more general framework, sufficient to expose a rather significant increase in complexity with respect to the original Myers-Pospelov setup. Most of these novel features are linked to the presence of spatial anisotropy, which is particularly pronounced when the symmetry-breaking vector is space-like, and they are such that they reduce the bound-setting power of certain types of observations in astrophysics. ", "introduction": "A large effort has been devoted over the last decade (see, {\\it e.g.}, Refs.~\\cite{grbgac,LQGDispRel,astroSchaefer,astroBiller,astroKifune,gacNature1999,urrutiaPRL,gacPIRANprd,jaconature,PiranNeutriNat,ellisPLB2009,gacSMOLINprd2009,fermiNATURE,gacPRL2009} and references therein) toward establishing that it is possible to actually study experimentally some minute effects introduced at the ultra-high ``Planck scale\" $M_P (\\simeq 1.2 \\cdot 10^{28} eV$), the scale expected to characterize quantum-gravity effects. At this point the scopes of this ``quantum-gravity phenomenology\"~\\cite{gacLRR} extend over a rather large ensemble of candidate quantum-gravity effects, inspired by (and/or formalized within) several models that are believed to be relevant for the understanding of the quantum-gravity problem. We here focus on one of these research programmes which has been driven by a model first introduced by Myers and Pospelov~\\cite{Myers2003}, as a candidate description of the Lorentz-symmetry-violation effects that are expected in some approaches to the quantum-gravity problem~\\cite{grbgac,LQGDispRel,urrutiaPRL,gacSMOLINprd2009}. This model adopts effective field theory for the description of Lorentz-symmetry-violation effects that are suppressed by a single power of the Planck scale (linear in $1/M_P$) and its proposal was primarily grounded on the observation~\\cite{Myers2003} that there is a unique such correction term which could be added to Maxwell theory, \\begin{equation} \\Delta \\mathcal L_{QG}= \\frac{1}{2M_P} n^\\alpha F_{\\alpha\\delta}n^\\sigma \\partial_\\sigma(n_\\beta\\varepsilon^{\\beta\\delta\\gamma\\lambda}F_{\\gamma\\lambda}) ~, \\label{eq:lagrangian} \\end{equation} if one enforces some relatively weak assumptions, including gauge invariance and the characterization of the symmetry-breaking structure in terms of an external four-vector $n^\\alpha$. Myers and Pospelov provided an even more definite and manageable framework by restricting their attention~\\cite{Myers2003} to the case in which the four-vector $n_\\alpha$ only has a time component, $n_\\alpha = (n_0,0,0,0)$. Then, upon introducing the convenient notation $\\xi \\equiv (n_0)^3$, one arrives at the following modified Maxwell Lagrangian density: \\begin{equation} \\mathcal L_{MP}=-\\frac{1}{4}F_{\\mu\\nu}F^{\\mu\\nu}+\\frac{\\xi}{2M_{P} } \\varepsilon^{jkl} F_{0 j} \\partial_0F_{k l}\\, , \\label{eq:MP} \\end{equation} and in particular it is then possible to exploit the simplifications provided by spatial isotropy. This Myers-Pospelov effective-field-theory model of Planck-scale modified electromagnetism has attracted much attention over the last few years. For phenomenologists it provided an ideal target (see {\\it e.g.} Refs~\\cite{gacLRR,mattinglyLRR,liberati0805,Galaverni:2007tq, Maccione:2008iw} and references therein), because of the presence of a single parameter and because (unlike most other fashionable proposals for the study of the quantum-gravity problem~\\cite{gacLRR}) its reliance on standard effective field theory poses no challenges at the level of ``physical interpretation\" of the formalism. This vigorous effort of investigation of the Myers-Pospelov model has produced a quick pace of improvement of experimental bounds, and, while the rough estimate invited by a quantum-gravity intuition~\\cite{gacLRR,mattinglyLRR,liberati0805} would be $\\xi \\sim 1$, the Myers-Pospelov parameter $\\xi$ is now constrained to be much smaller than 1, with some analyses~\\cite{Galaverni:2007tq, Maccione:2008iw} even suggesting a bound at the level $\\xi < 10^{- 15}$. We here observe that however these bounds are not applicable to the general correction term $ \\Delta \\mathcal L_{QG}$ of Eq.~(\\ref{eq:lagrangian}), since they exploit significantly the spatial isotropy regained by the {\\it ad hoc} choice $n_\\alpha = (n_0,0,0,0)$. And actually this {\\it ad hoc} choice is only available for a restricted class of frames of reference: even imposing ``by brute force\" $n_\\alpha = (n_0,0,0,0)$ in some desired frame of reference, then the four-vector $n_\\alpha$ will of course still acquire a spatial component in other (boosted) frames. Since the main strategy for constraining the Myers-Pospelov parameter has relied on various astrophysics observations, conducted in different ``laboratory frames\", these are concerns that necessarily must be investigated, at least in order to establish to which extent those limits are vulnerable to the presence of a (perhaps small, but necessarily nonzero) spatial component in frames other than the ``preferred frame\". In the next section we therefore propose a phenomenology centered on the more general form of the $ \\Delta \\mathcal L_{QG}$ of Eq.~(\\ref{eq:lagrangian}), involving therefore an arbitrary (four-parameter) four-vector $n_\\alpha$, and we describe the resulting equations of motion for the electromagnetic field. Since the types of data that are most useful and are likely to still be most useful to set bounds on this framework concern regimes that involve classical electromagnetic waves, we shall here be satisfied with an analysis confined at the level of some modified Maxwell equation for classical electromagnetic waves. In this respect we adopt the same perspective of the original analysis by Myers and Pospelov~\\cite{Myers2003}, but for our purposes it is valuable to provide, as we shall, a more detailed description of the Planck-scale modifications of classical electromagnetic waves, whereas Ref.~\\cite{Myers2003} focused exclusively on the form of the dispersion (/``on-shell\") relation. In Section~3 we investigate the features that are likely to be most relevant from the phenomenology perspective, which concern dispersion, birefringence and a possible longitudinal component. In Section~4 we provide a rough quantitative characterization of the effects introduced by the spatial components of $n_\\alpha$, focusing mainly on cases with space-like symmetry-breaking vector and stressing that the magnitude of the effects is not exclusively governed by the magnitude of the spatial components of $n_\\alpha$: there are direction-dependent (anisotropic) effects, and even small values of the spatial components of $n_\\alpha$ produce large effects within a certain range of directions. Section~5 offers some closing remarks. ", "conclusions": "" }, "1003/1003.3340_arXiv.txt": { "abstract": "HR\\,8799 is a $\\lambda$ Bootis, $\\gamma$ Doradus star hosting a planetary system and a debris disk with two rings. This makes this system a very interesting target for asteroseismic studies. This work is devoted to the determination of the internal metallicity of this star, linked with its $\\lambda$ Bootis nature (i.e., solar surface abundances of light elements, and subsolar surface abundances of heavy elements), taking advantage of its $\\gamma$ Doradus pulsations. This is the most accurate way to obtain this information, and this is the first time such a study is performed for a planetary-system-host star. We have used the equilibrium code CESAM and the non-adiabatic pulsational code GraCo. We have applied the Frequency Ratio Method (FRM) and the Time Dependent Convection theory (TDC) to estimate the mode identification, the Brunt-Va\\\"is\\\"al\\\"a frequency integral and the mode instability, making the selection of the possible models. When the non-seismological constraints (i.e its position in the HR diagram) are used, the solar abundance models are discarded. This result contradicts one of the main hypothesis for explaining the $\\lambda$ Bootis nature, namely the accretion/diffusion of gas by a star with solar abundance. Therefore, according to these results, a revision of this hypothesis is needed. The inclusion of accurate internal chemical mixing processes seems to be necessary to explain the peculiar abundances observed in the surface of stars with internal subsolar metallicities. The use of the asteroseismological constraints, like those provided by the FRM or the instability analysis, provides a very accurate determination of the physical characteristics of HR\\,8799. However, a dependence of the results on the inclination angle $i$ still remains. The determination of this angle, more accurate multicolour photometric observations, and high resolution spectroscopy can definitively fix the mass and metallicity of this star. ", "introduction": "The A5 V star HR\\,8799 (V342 Peg, HD\\,218396, HIP\\,114189) has been extensively studied in the last years. The first studies of this star were done in the context of asteroseismology. This technique is being developed as an efficient instrument for the study of stellar interiors and evolution \\citep{libroastro}. Through the comparison of the observed pulsational modes with theoretical calculations, the general and internal characteristics of the star can be determined. \\cite{schuster} firstly reported HR\\,8799 as a possible SX Phoenicis type. \\cite{zerbi} observed this star in a multisite multicolour photometric campaign, with Str\\\"omgren filters, and found three independent pulsational frequencies ($f_1=1.9791\\;\\rm{c}\\,\\rm{d}^{-1}$ ($\\equiv$ 22.906 $\\mu$Hz), $f_2=1.7268\\;\\rm{c}\\,\\rm{d}^{-1}$ ($\\equiv$ 19.986 $\\mu$Hz), and $f_3=1.6498\\;\\rm{c}\\,\\rm{d}^{-1}$ ($\\equiv$ 19.095 $\\mu$Hz), units are cycles per day), making it one of the 12 first $\\gamma$ Doradus pulsators known \\citep{kaye1}. This pulsating stellar group is composed of Main Sequence (MS) stars in the lower part of the classical instability strip \\citep{tdcma}. Their pulsation modes have periods in the range [0.5,3] days, that is, they are asymptotic g-mode pulsators. \\cite{gray} obtained an optical spectrum of HR\\,8799, and assigned an spectral type of kA5 hF0 mA5 V $\\lambda$ Bootis, reporting an atmospheric metallicity of [M/H]=$-0.47$. They also noted that HR\\,8799 may be also a Vega-type star, characterized by a far IR excess due to a debris disk. \\cite{sadakane} developed a deep study of the metal abundances of this star, confirming its $\\lambda$ Bootis nature (with surface chemical peculiarities). Two years later, \\cite{Marois} reported the presence of a planetary system around this star. It was the first detection of such a system carried out by direct imaging. This was the starting point of a set of approximately ten studies about this system during 2009, none of them from the asteroseismic point of view, even considering that the three detected pulsational frequencies can be used to better understand the host star, in particular its $\\lambda$ Bootis nature and evolutionary status. Discovered by \\citet{Morgan43}, the $\\lambda$ Bootis-type stars are non-magnetic, moderately-rotating, Population I stars with spectral types from late B - early A to F (dwarfs)\\footnote{Two percent of the A-type stars belong to this class \\citep{GrayCorbally02}}, which show peculiarities in the morphology and abundance of the Fe-peak element lines. In particular, these lines are unusually weak considering their spectral types. Significant deficiencies in their abundances (up to 2\\,dex) are found, whereas C, N, O and S have solar abundance \\citep{Paunzeniliev02}. Different theories have attempted to explain the $\\lambda$ Bootis nature from both observational (photometry, spectroscopy) and theoretical investigations. It is not our aim to discuss all of them here \\citep[for a interesting review see][]{Paunzen03}. Nevertheless, it may be worth describing the most accepted scenario relying on the accretion of inter-stellar medium gas by the star \\citep{VennLambert90}. The accretion/diffusion scenario would explain the abundances found at the base of the outer convective zone of these stars, since convective layers are assumed to remain chemically homogeneous. This accretion/diffusion scenario is based on the accretion of inter-stellar medium by the star, and the mixing of these elements with those of the star due to diffusion and rotationally mixing processes. The accretion rate required to maintain this situation is of the order of $10^{-10}$-$10^{-14}\\,{\\rm M}_{\\odot}$ per year \\citep{turcotte}, and, once the accretion has ceased, the metal deficiencies should disappear in approximately 1 Myr due to diffusion and internal mixing processes. Thus, a possible interpretation is that $\\lambda$ Bootis stars are young A-type stars (in a pre-main sequence or zero-age main sequence evolutionary stage), still interacting with their primordial clouds of gas and dust. Interestingly, \\citet{Paunzeniliev02} found that most of the known $\\lambda$ Bootis stars lie between the ZAMS (zero-age main sequence) and the TAMS (terminal-age main sequence, with ages of several hundreds million years). In this case the most likely scenario would be a MS star with solar abundance passing through an interstellar cloud \\citep{kamppaunzen}. However, the chemical mixing due to internal processes, such as rotationally-induced mixing, cannot be discarded as a possible explanation of the observed abundances. $\\lambda$ Bootis stars and other types of objects in the same region of the HR diagram, such as $\\delta$ Scuti and $\\gamma$ Doradus stars, are considered as particularly suitable for the asteroseismological study of poorly known hydrodynamical processes occurring in stellar interiors, like the extent of the convective core, mixing of chemical elements, redistribution of angular momentum \\citep{Zahn92,steph1,steph2}, etc. $\\lambda$ Bootis-type stars are also pulsating stars. Therefore, asteroseismology can be used to obtain information about the internal structure of these objects. Several works have been devoted to study $\\lambda$ Bootis stars with $\\delta$ Scuti pulsations, for instance \\cite{Paunzen98,Casas09}. In addition, the combined use of space and ground-based observations improves the potential of the modelling of the star \\citep{bruntt}. The present work aims at a comprehensive asteroseismic modelling of HR\\,8799, focusing on the discussion of the $\\lambda$ Bootis nature of the star. More precisely, we want to answer whether the observed abundance is intrinsic to the star or an effect of surface processes. This answer drives the search for possible mechanisms explaining the observed abundances, that is, if there is accretion or not, together with internal chemical transport (rotationally-induced mixing, gravitational settling, radiative levitation, etc.). Up to now, only three $\\lambda$ Bootis stars have been reported to be $\\gamma$ Doradus pulsators: HD\\,218427 \\citep{rodriguez06a}, HD\\,239276 \\citep{rodriguez06b}, and HR\\,8799. As $\\gamma$ Doradus stars are asymptotic g-mode pulsators, they are very good candidates for testing the internal structure of the star, in particular its internal metallicity. Some of the most updated tools adapted for this purpose are used: 1) the evolutionary code CESAM \\citep{Morel08}, and 2) the pulsation code GraCo \\citep{graco1,graco2}. Using these tools we performed a massive numerical study of HR\\,8799 in an attempt of constraining physical and theoretical parameters. In this work we will follow the same scheme used for the study of RV Arietis, 29 Cygnis and 9 Aurigae \\citep{Casas06,Casas09,Moya06au}. ", "conclusions": "In this work, the first comprehensive asteroseismologic study of the planetary system host HR\\,8799, a $\\lambda$ Bootis star presenting $\\gamma$ Doradus pulsations has been carried out. This asteroseismic work is specially important for the determination of the internal abundances of this kind of stars, a previous step to understand the physical mechanism responsible for the surface chemical peculiarities of the $\\lambda$ Bootis group. The asymptotic g modes observed in HR\\,8799 make it possible to deeply test its internal metallicity, something hard to do with $\\delta$ Scuti pulsators, for example. For the asteroseismologic study, we have used part of the most updated codes and physics for describing $\\gamma$ Doradus stars: CESAM and GraCo codes, TDC theory, FRM, and mode identification using Str\\\"omgren multicolour photometry. We have constructed a dense grid of equilibrium models, and carried out a first selection using the available data ($T_{\\rm{eff}}$, $\\log\\,\\rm{g}$, and luminosity), leaving the metallicity and mass as free parameters in order to test the internal metallicity of the star. The first selection shows that there are no models with solar metallicity fulfilling the observations, since the stellar luminosity derived from the observations is smaller than any of the possible luminosities of models with solar metallicity. This contradicts the main assumption of the theories explaining the $\\lambda$ Bootis nature, i.e. that these stars have solar metallicity, whereas the observed abundances are due to surface phenomena. The asteroseismic study to fix the metallicity of the star is done using 1) the FRM, to estimate possible mode identifications and Brunt-V\\\"ais\\\"al\\\"a integrals, 2) energy balance to determine the stability of the modes, selecting models in a certain temperature range and $\\alpha_{\\rm{MLT}}$ and, 3) mode identification with multicolour photometric observables, selecting only models for which FRM predicts the mode with the highest amplitude to have $\\ell=2$. This is the first time that such a study has been done for a planet-hosting star, but the general procedure was proposed in \\cite{completo}. All these steps provide a small group of acceptable models depending on the rotation velocity. For a moderate rotation velocity (implying $i$ around $50^\\circ$) the mass and the metallicity are very precise determined ([M/H]=$-0.32\\pm0.1$ and M=$1.32\\pm 0.01\\rm{M}_{\\odot}$). For a rotation velocity at the limit of the validity of the FRM (around $i\\approx 36^\\circ$), the method is less selective and provides some possible values of the metallicity and mass ([M/H]=$-0.32$ and $-0.12\\pm0.1$, and M=[1.32,1.33] and [1.44,1.45] $\\rm{M}_{\\odot}$ respectively). We point out that inclination angles lower than $18^\\circ$ cannot be possible to preserve the $\\lambda$ Bootis nature of the star. Nevertheless, this limit must be also verified with more accurate observations. These results have some consequences. The first one is related to the most accepted theory of the $\\lambda$ Bootis nature, namely the accretion/diffusion scenario given by \\cite{VennLambert90}. The result of this work, discarding solar metallicity as the internal metallicity of the star, together with the fact that some $\\lambda$ Bootis stars have debris disks not connected with the star \\citep{chen}, makes this explanation unlikely, but not negligible \\citep{su}. In any case, the equation developed by \\cite{VennLambert90} to describe the individual abundances as a sum of internal abundance plus that of accreted material, must be corrected to include a possible non-solar internal abundance as following \\begin{equation} \\epsilon(m)=(1-f)\\epsilon_*(m)+f\\epsilon_{\\rm{ISM}}(m) \\end{equation} \\noindent where $f$ is a mixture factor, and $\\epsilon$, $\\epsilon_*$ and $\\epsilon_{\\rm{ISM}}$ are the abundances observed, internal and coming from the interstellar medium of each chemical element, respectively. Therefore, the study of internal chemical mixing processes seems to be the key to explain the $\\lambda$ Bootis nature, at least for HR\\,8799, as the solar abundances for C, N, O and S observed on its surface have still to be explained. On the other hand, we have seen that the asteroseismological study needs an accurate determination of the rotation velocity, not just the projected rotation velocity, of the star, for a very accurate determination of the stellar characteristics. Therefore, a consequence of this study is the need for a precise determination of the inclination angle $i$, of the multicolour photometric amplitudes and phases of $f_2$, and some information of $m$ values through time-series if high resolution spectroscopy. These determinations would help to carry out a definitive selection of the models." }, "1003/1003.3030_arXiv.txt": { "abstract": "The physical sizes of supernova remnants (SNRs) in a number of nearby galaxies follow an approximately linear cumulative distribution, contrary to what is expected for decelerating shock fronts. This phenomenon has been variously attributed to observational selection effects, or to a majority of SNRs being in ``free expansion'', with shocks propagating at a constant velocity into a tenuous ambient medium. We compile multi-wavelength observations of the 77 known SNRs in the Magellanic Clouds, and argue that they provide a fairly complete record of the SNe that have exploded over the last $\\sim 20$~kyr, with most of them now in the adiabatic, Sedov phase of their expansions. The roughly linear cumulative distribution of sizes (roughly uniform in a differential distribution) can be understood to result from the combination of the deceleration during this phase, a transition to a radiation-loss-dominated phase at a radius that depends on the local gas density, and a probability distribution of densities in the interstellar medium varying approximately as $\\rho^{-1}$. This explanation is supported by the observed powerlaw distributions, with index $\\sim -1$, of three independent tracers of density: neutral hydrogen column density, H$\\alpha$ surface brightness, and star-formation rate based on resolved stellar populations. In this picture, the observed cutoff at a radius of 30~pc in the SNR size distribution is due to a minimum in the mean ambient gas density in the regions where supernovae (SNe) explode. We show that M33 has a SNR size distribution very similar to that of the Magellanic Clouds, suggesting these features, and their explanation, may be universal. In a companion paper (Maoz \\& Badenes 2010), we use our sample of SNRs as an effective ``SN survey'' to calculate the SN rate and delay time distribution in the Magellanic Clouds. The hypothesis that most SNRs are in free expansion, rather than in the Sedov phase of their evolution, would result in SN rates that are in strong conflict with independent measurements, and with basic stellar evolution theory. ", "introduction": "The ecology of galaxies is dominated by supernova (SN) explosions, which inject energy and enriched material into the interstellar medium and trigger the formation of the next generations of stars. Many fundamental aspects of SNe are still poorly understood, both for the core-collapse (CC) SN explosions that are thought to end the lives of massive ($\\ga 8M_\\odot$) stars and for the type-Ia SNe (SNe~Ia) that are believed to be the thermonuclear combustions of CO white dwarfs (WDs) that approach the Chandrasekhar mass. Nevertheless, the two flavors of SNe deposit a similar amount ($\\sim10^{51}$ erg) of kinetic energy into the surrounding medium, leaving behind supernova remnants (SNRs) that remain visible for thousands of years. In recent times, the study of young SNRs at X-ray wavelengths has emerged as a new way to explore the physics of CC and Type Ia SN explosions \\citep[see][and references therein]{badenes10:PNAS_review}, but most known SNRs are too old to provide much information about the specific events that originated them \\citep[see discussions in][]{rakowski05:G337,badenes09:SNRs_LMC}. In spite of this, much can be learned by studying the entire population of SNRs, young and old, in nearby galaxies. Because the timescales for SNR evolution are short compared to most of the processes that affect the structure of galaxies, SNR catalogues provide a clean record of the environments where SNe explode, which can be used to put constraints on the properties of their progenitors \\citep[e.g.][]{badenes09:SNRs_LMC}. Moreover, by considering the properties of the entire population of SNRs in a galaxy together with the bulk properties of the gas they are expanding into, we can gain insights into the evolutionary phases of SNRs \\citep{woltjer72:SNR-review}, the structure of galaxies on scales comparable to the average SNR size \\citep{cox05:ISM}, and the cycles of matter and energy in the interstellar medium \\citep{ferriere01:ISM}. In this paper and in a companion publication (Maoz \\& Badenes 2010, henceforth Paper~II), we use the SNR population in the Magellanic Clouds to explore some of these issues. The Magellanic Clouds (MCs) have the advantage of being close enough to study key aspects of their global structure in great detail, and they also harbor a large and extensively observed population of SNRs. Thus, they are the optimal setting to study the interplay between local density, star formation, SN explosions, and SNR evolution on a galactic scale. Our ultimate goal, and the focus of Paper~II, is to derive the SN rate and delay time distribution (i.e., the SN rate as a function of time following a brief burst of star formation) in the Magellanic Clouds. However, this cannot be done without understanding first the relationship between the lifetime of SNRs and the properties of their local environments. This is the subject of the present work. The evolution of SNRs has been the subject of many theoretical studies \\citep[e.g.,][]{woltjer72:SNR-review,chevalier82:selfsimilar,cioffi88:Radiative_SNRs,blondin98:Radiative_SNRs,truelove99:adiabatic-SNRs}. Because accurate ages are only known for a handful of young, often historical objects, any observational tests of these theoretical models must rely on SNR size as a proxy for age. Given that SNRs of equal ages will have different sizes if they expand in different media, this necessarily brings the role of local density into the picture. Previous works on the distribution of SNR sizes initially focused on the Milky Way and the MCs \\citep[e.g.][]{mathewson84:SNR-Magellanic-Clouds,green84:SNR_Statistics,hughes84:SNR_N_D,berkhuijsen87:SNRs_N_D,chu88:LMC_SNRs_environments}, but more recent efforts have also explored other galaxies in the Local Group, including M31 \\citep{magnier97:M31_SNRs_ROSAT}, M33 \\citep{long10:M33_SNRs}, and M83 \\citep{dopita10:M83_SNRs}. With few exceptions, these studies gave little consideration to the bulk properties of the gas in the galaxies hosting the SNRs. In many cases, their samples were also affected by issues of completeness and biases from working at a single wavelength. Not surprisingly, these efforts have failed to produce a unified, physically motivated picture of the evolution of SNRs in the interstellar medium. Here, we propose a first approximation to the problem in the context of the Magellanic Clouds. This paper is organized as follows. In \\S~\\ref{catalog}, we present a compilation of multi-wavelength observations for the 77 known SNRs in the Magellanic Clouds, and we argue that it provides a fairly complete record of all the SNe that have exploded over the last $\\sim 20$~kyr. In \\S~\\ref{distribution}, we examine the size distribution of SNRs in both galaxies, and we find that the cumulative distribution is close to linear (i.e., the differential is close to uniform), within the uncertainties associated with the relatively small number of objects, up to a marked cutoff at a physical radius of $\\sim$ 30 pc. We also show that the SNR size distribution in M33 has very similar properties, suggesting that these features might be widespread. In \\S~\\ref{physics}, we propose a physical model to explain this distribution, based on the assumption that most objects are in the Sedov (adiabatic) stage of their evolution, and that they rapidly fade away once they transition to the radiative stage, at an age that depends on the local density. Under these conditions, the uniform distribution of SNR sizes requires that the gas density in the Clouds have a probability distribution described by a power law with an index of $-$1 (i.e., $\\delta P / \\delta \\rho \\sim \\rho^{-1}$). In \\S~\\ref{densityestimates}, we test this requirement by examining the distribution of three independent density tracers in the Clouds: neutral hydrogen column density, H$\\alpha$ surface brightness, and star-formation rate based on resolved stellar populations. We find that these tracers are indeed well described by powerlaws with a $-$1 index. This lends credence to our model, and provides us with the crucial means to estimate the visibility time of SNRs in different locations, which we review briefly in \\S~\\ref{sec:disc}, and more extensively in Paper~II, where we use it to derive the SN rate and delay time distribution in the Clouds. We conclude by summarizing our main results and outlining avenues for future work in \\S~\\ref{sec:summary}. \\begin{figure*} \\includegraphics[width=0.7\\textwidth]{lmcsnrp1fig1a.eps} \\includegraphics[width=0.7\\textwidth]{lmcsnrp1fig1b.eps} \\caption{Maps of the LMC \\citep[top, from][]{kim03:LMC_HI_Parkes_ATCA} and the SMC \\citep[bottom, from][]{stanimirovic99MNRAS.302..417S} in the HI 21~cm line. The positions of the SNRs from Table 1 are indicated by crosses: 54 objects in the LMC, 23 objects in the SMC. The contours are at column densities of 5, 10, 30, 50, and $100\\times10^{20}\\,\\mathrm{cm^{-2}}$. \\label{fig-HIMaps}} \\end{figure*} ", "conclusions": "\\label{sec:summary} In this paper, we have proposed a physically motivated scenario for the evolution of SNRs in the interstellar medium of galaxies, and we have applied it to the sample of SNRs in the Magellanic Clouds. We compiled multi-wavelength observations of the 77 known SNRs in the MCs from the existing literature (Table 1 and Figure \\ref{fig-HIMaps}). We verified that this compilation is fairly complete, and that the size distribution of SNRs is approximately flat, within the allowed uncertainties, up to a cutoff at $r\\sim30$ pc, as discussed by other authors before. We noted that these features are also present in the larger SNR sample assembled by \\cite{long10:M33_SNRs} for M33. In our model, the flat size distribution can be explained if most SNRs are in the Sedov, decelerating, stage of their expansion, quickly fading below detection as soon as they reach the radiative stage at an age (size) that depends on the local density. Under these circumstances, a flat distribution of SNR sizes arises naturally if the probability distribution of densities in the gas follows a power law with index $-1$. To test this hypothesis, we have examined the distributions of three different density tracers in the Clouds: HI column density, H$\\alpha$ flux, and SFR obtained from resolved stellar populations. We have verified that these tracers do indeed indicate a density distribution that follows a power law with index $-1$, over a wide range of densities, and that this agrees with our theoretical understanding of the dynamics of the interstellar medium. Our scenario implies that SNRs will remain visible for different times at different locations in the Magellanic Clouds, depending on the local density. This visibility time is a key ingredient in the calculation of SN rates and SN delay time distributions, which we derive for the Magellanic Clouds in Paper~II. The absolute value of the visibility time cannot be determined from theoretical arguments alone, but we have used the results from Paper~II to verify that the range of values we obtain ($13$ to $23$ kyr for regions of mean density, depending on the tracer) is consistent with the existing limits for the maximum ages of SNRs, and with basic stellar evolution theory. It would be interesting to extend our analysis of SNR sizes and interstellar medium densities to other galaxies in the Local Group. On the one hand, this would allow us to test the validity of our SNR evolution scenario in different settings. On the other hand, the techniques developed here and in Paper~II would yield more accurate SN rates and more detailed delay time distributions with larger SNR samples, provided the SFHs in the host galaxies could be determined with enough spatial and temporal resolution. A growing number of nearby galaxies have published SNR catalogues, but samples of the quality of the one we have assembled here for the Magellanic Clouds or the one published by \\cite{long10:M33_SNRs} for M33 are hard to obtain. Interstellar extinction and distance uncertainties plague the Milky Way SNRs, and the radio detection limits become comparable to the fluxes of the fainter Magellanic Cloud SNRs for distances beyond several hundred kpc \\citep[see][]{chomiuk09:SNR_Luminosity_Function}. With a judicious investment of observing time, however, reasonably complete multi-wavelength catalogues of SNRs in a few well-suited nearby galaxies could become available in the near future." }, "1003/1003.4426_arXiv.txt": { "abstract": "We present multiwavelength investigation of morphology, physical-environment, stellar contents and star formation activity in the vicinity of star-forming region Sh 2-100. It is found that the Sh 2-100 region contains seven \\hii regions of ultracompact and compact nature. The present estimation of distance for three \\hii regions, along with the kinematic distance for others, suggests that all of them belong to the same molecular cloud complex. Using near-infrared photometry, we identified the most probable ionizing sources of six \\hii regions. Their approximate photometric spectral type estimates suggest that they are massive early-B to mid-O zero-age-main-sequence stars and agree well with radio continuum observations at 1280 MHz, for sources whose emissions are optically thin at this frequency. The morphology of the complex shows a non-uniform distribution of warm and hot dust, well mixed with the ionized gas, which correlates well with the variation of average visual extinction ($\\sim$ 4.2 - 97 mag) across the region. We estimated the physical parameters of ionized gas with the help of radio continuum observations. We detected an optically visible compact nebula located to the south of the 850 $\\mu$m emission associated with one of the \\hii regions and the diagnostic of the optical emission line ratios gives electron density and electron temperature of $\\sim$ 0.67 $\\times$ 10$^{3}$ cm$^{-3}$ and $\\sim$ 10$^{4}$ K, respectively. The physical parameters suggest that all the \\hii regions are in different stages of evolution, which correlate well with the probable ages in the range $\\sim$ 0.01 - 2 Myr of the ionizing sources. The spatial distribution of infrared excess stars, selected from near-infrared and IRAC color-color diagrams, correlates well with the association of gas and dust. The positions of infrared excess stars, ultracompact and compact \\hii regions at the periphery of an \\hi shell, possibly created by a WR star, indicate that star formation in Sh 2-100 region might have been induced by an expanding \\hi shell. ", "introduction": "Massive star formation is poorly understood as compared to low-mass stars, because their formation takes place in the dense core of a molecular cloud of high visual extinction, usually observable at far-infrared (FIR) to millimeter wavelengths and their evolutionary time scales are much shorter ($\\leq$ 10$^{5}$ yr) (e.g., Bernasconi \\& Maeder 1996). The radio free-free emission in terms of ultracompact \\hii (\\uchii) regions, represents a later evolutionary sequence of a massive protostar outlined by Beuther et al. (2007), where the associated high mass protostar may still be in the accretion phase (e.g, Shepherd \\& Churchwell 1996) or has already ceased the accretion (Garay \\& Lizano 1999). The \\uchii region can further evolve into the less obscure stage of compact \\hii (\\chii) and classical \\hii regions (Garay \\& Lizano 1999) by the disruption of associated gas and dust, revealing both, the embedded high-mass and lower mass stellar population at shorter wavelengths. Most of the high-mass star-forming regions in the Galaxy lie at a distance of more than 1 kpc. As a result, the problem in interpreting observations of massive star formation originates in source confusion at the core of distant molecular clouds due to their cluster mode formation. Statistically, it has been shown that the most luminous protostars form in molecular clouds associated with evolved \\hii regions (Dobashi et al. 2001), where the interplay between the stars' radiation field and associated gas and dust makes the environment even more complex. Therefore, a detailed study of the star-forming region hosting young massive stars, using optical, near-infrared (NIR), mid-infrared (MIR) and radio bands, is necessary to understand the high-mass star formation scenario in these complexes. In this paper, we have studied a young star-forming region (SFR) K3-50 shown in Fig. 1, which contains a group of \\hii regions, namely A, B, C, D, E, and F (cf. Israel 1976). The \\hii region C itself consists of two \\uchii regions, C1 and C2 (Harris 1975). The kinematic distance towards K3-50 ranges from 7.9 to 9.3 kpc (Harris 1975; Bronfman et al. 1996; Araya et al. 2002). In the present study we have adopted a widely used value of distance of $\\sim$ 8.7 kpc for our analysis. We also derived distances of three \\hii regions that come out to be close to the value of 8.7 kpc. The group of \\hii regions is located within an area of radius $\\sim$ 3\\arcm.5 centered on $\\alpha_{2000} = 20^{h}01^{m}42^{s}$, $\\delta_{2000} = +33^{\\circ}33^{\\prime}49^{\\prime\\prime}$, in the proximity of an evolved diffuse nebula Sh 2-100 of $\\sim$ 4$^{\\prime}$ in size (Sharpless 1959). Sh 2-100 is itself a part of the large molecular cloud complex W58, which consists of several classes of \\hii regions with widely varying physical parameters. The observed radio luminosity of the complex can be explained by the presence of OB association (cf. Israel 1976). In Fig. 1 one can notice that K3-50D, K3-50E, and K3-50F are associated with nebulosity. A faint nebulosity is also seen $\\sim$ 15.5$^{\\prime\\prime}$ away in the southern direction of K3-50C2, whereas K3-50A appears as a stellar point-like source and is marked with a circle. The regions K3-50B and K3-50C1 are optically invisible radio sources, implying that the extinction is high enough to obscure the regions. By comparing the expected infrared Br$\\alpha$ line fluxes as predicted from radio continuum fluxes, to the observed Br$\\alpha$ line fluxes, Roelfsema et al. (1988) estimated visual extinction of 15, 26, 97 and 32 mag towards A (over 25$^{\\prime\\prime}$ aperture), B (over 5$^{\\prime\\prime}$ aperture), C1 and C2 (over 15$^{\\prime\\prime}$ aperture) respectively, while the extinction towards D was found to be 2 mag (5$^{\\prime\\prime}$ aperture). Observations at radio wavelengths suggest that these regions harbor at least one massive OB star (Harris 1975; Israel 1976; DePree et al. 1994). The \\uchii regions, astronomical masers, outflows, detection of high density tracer molecules and infrared (IR) sources with IR-excess are the trademarks of young star-forming regions. The evidence of weak bipolar molecular outflow by Phillips \\& Mampaso (1991) with CO $J$ = 2 - 1 transition, ionized outflow by DePree et al. (1994) with H76 alpha radio recombination line, detection of molecular core mapped in the \\hco and CS (2-1) lines by Vogel \\& Welch (1983) and Bronfman et al. (1996) suggest the youthfulness of the Sh 2-100 region. These observations along with the presence of astronomical masers (H$_2$O by Kurtz \\& Hofner 2005; OH by Baudry \\& Desmurs 2002) imply active star formation is going on around Sh 2-100. The current star formation activity in the vicinity of Sh 2-100 region is probably the effect of feedback from the earlier generation stars of the W58 complex (Israel 1980). Though there are several studies in the radio and infrared domain, most of them are concentrated on K3-50A and K3-50C. To continue our multiwavelength investigations of massive SFRs (cf. Ojha et al. 2004a; Ojha et al. 2004b; Ojha et al. 2004c; Tej et al. 2006; Samal et al. 2007; Pandey et al. 2008), we have studied an area of 8$^{\\prime}$ \\into 8$^{\\prime}$ of Sh 2-100 region to understand the star formation activity in this region. This paper presents new results from optical and infrared photometry, optical spectroscopy and low frequency radio continuum observations. Based on these observations, we have carried out a detailed study of the ionizing sources of individual \\hii regions, physical characteristics and the nature of associated lower mass population. With the following layout of the paper we tried to interpret the possible star formation scenario of W58 cloud complex. In Sect. 2, we describe our observations and the reduction procedures. In Sect. 3, we discuss other available datasets used in the present study. Sect. 4 describes the general morphology of the region. Sect. 5 describes the infrared and radio components associated with the region. In Sect. 6, we discuss the nature of individual regions. Sect. 7 is devoted to general discussion and star formation scenario in W58 complex. We present the main conclusions of our results in Sect. 8. ", "conclusions": "In this paper we have presented a multi-wavelength study of the stellar contents and physical environment of star-forming region (K3-50) in the proximity of diffuse \\hii region Sh 2-100. K3-50 consists of seven \\hii regions (A, B, C1, C2, D, E and F) of different classes and are thought to be a part of W58 molecular cloud complex located at a distance of $\\sim$ 8.7 kpc. Our main conclusions are as follows: 1. We have identified the ionizing star of the optical visible nebula K3-50D. Optical spectroscopy suggests that an O4V star is responsible for the ionization. We have estimated a distance of $\\sim$ 8.6 kpc for the K3-50D region. The approximate distances for the K3-50E and K3-50F regions have been estimated from the NIR photometry and radio continuum flux densities following the methodology by Comer\\'{o}n \\& Torra (2001). The distance estimation suggests that both the regions are likely the part of the W58 complex. 2. We identified the probable exciting sources for six \\hii regions using the $J$ vs. $J-H$ CM diagram. The photometric spectral types of the ionizing sources agree well within a subclass with that derived from radio observations based on the number of Lyc photons. 3. We derived physical parameters for all the five \\hii regions using 1280 MHz radio continuum observation. The \\hii regions cover a wide range in size ($\\sim$ 0.2 to 1.6 pc), optical depth ($\\sim$ 0.18 to 1.92), rms electron density ($\\sim$ 740 to 6390 cm$^{-3}$) and rate of Lyc photons ($\\sim$ 5.5 $\\times$ 10$^{47}$ s$^{-1}$ to $\\sim$ 1.1 $\\times$ 10$^{49}$ s$^{-1}$). These physical parameters reflect different evolutionary stages of \\hii regions. 4. We estimated the total mass of the cloud associated with one of the \\hii regions (K3-50A) which was found to be $\\sim$ 7800 $\\msun$. At the center of the cloud core, we detected a possible massive YSO with $M_{\\star}$ $\\sim$ 42 $\\pm$ 4.4 $M_\\odot$, accreting with an effective envelope accretion rate of Log($\\dot{M}_{env}$) $\\sim$ -3.8$\\pm$ 0.2 $M_\\odot$ yr$^{-1}$. 5. Combining IRAC photometry with ground-based NIR observations within an area of $\\sim$ 8$^{\\prime}$ $\\times$ 8$^{\\prime}$, we found a total of 150 objects with circumstellar disks (i.e., YSOs), suggesting a site of active star formation. The distribution of these YSOs (Class II, Class I and $H-K$ $>$ 1.8 sources) in Sh 2-100 SFR correlates well with the association of gas and dust within the region. 6. The distribution of YSOs at the periphery of \\hi shell on a large scale supports the speculation by Israel (1980) that the star formation activity around K3-50 region might have been induced by an expanding spherical bubble which was created by strong stellar winds from older generation OB stars of the complex in the south-eastern direction. The Sh 2-100 SFR region represents a broad sample of different stages of massive star formation all in the same cloud, from the cold, dense core at source C1 to the diffuse H II region at source F. The conclusions drawn in this paper on this cloud should be relevant and applicable to further deep study of the region with high resolution and high sensitivity instruments." }, "1003/1003.3206_arXiv.txt": { "abstract": "In this paper we consider a unique model of inflation where the universe undergoes rapid asymmetric oscillations, each cycle lasting $\\sim 10^{6}$ Planck times. Over many-many cycles the space-time expands to mimic the standard inflationary scenario. Moreover, these rapid oscillations leave a distinctive periodic signature in $\\ln k$ in the primordial power spectrum, where $k$ denotes the comoving scale. Although the cyclic-inflation model contains additional parameters as compared to the standard power-law spectrum, the improvement to the fit of the 7-year WMAP data is significant. ", "introduction": "Primordial inflation has been very successful in explaining the perturbations in the cosmic microwave background (CMB) radiation and the large scale structures of the universe~\\cite{wmap7}, for a recent review see~\\cite{RM}. However inflation has some outstanding problems, in particular it doesn't encode a non-singular geodesically complete evolution~\\cite{Linde1}. We will present a unique singularity free geodesically complete realization of inflation in the context of cyclic cosmologies. In cyclic cosmological models, rather than having a singular {\\it beginning of time}, our universe can be made {\\it eternal} in both past and future~\\cite{tolman,narlikar,Starobinsky,ekcyclic,barrow,phantom,peter,emergent,BM,saridakis}. However, in most cyclic scenarios the effort has been to produce primodial fluctuations within a single long cycle~\\footnote{Recently, there have been attempts to calculate how perturbations can evolve through various cycles~\\cite{robert}.}. Although there have been progress~\\cite{justin,BMS,matter}, this has proved rather challenging in comparison to the success which inflation enjoys in explaining the observed near scale invariant perturbations in the CMB. In this paper we provide a simple alternative to the standard inflation and cyclic universe scenarios in the form of a cyclic-inflation model which tries to incorporate the successes of both; our model includes a non-singular cyclic phase of evolution where in every cycle the universe contracts a little less than it expands leading to an overall growth. In fact, over many-many cycles the space-time resembles that of inflation. Thus, in close analogy with inflation, we can explain how the seed perturbations generated at much higher energy densities can be stretched to the observable scales, and why the spectrum is nearly scale free. Additionally, the model leaves distinct signatures of the rapid oscillations the universe undergoes by modifying the power spectrum with periodic signatures. Last but not the least, it turns out that the cyclic-inflationary phase requires a negative potential energy, but the universe can gracefully exit to a positive potential region marking the end of the inflationary phase and the onset of standard radiation dominated era. Thus our model may provide a way of reaching a positive energy vacuum from a plethora of negative energy vacua in the string landscape~\\cite{landscape}. Finally, Tolman's entropy problem (which is equivalent to the problem of geodesic incompleteness in our model) can be naturally addressed by including a pre-inflationary emergent phase where the scale-factor starts to oscillate periodically as we approach past infinity, the size of the universe never becomes vanishingly small~\\cite{emergent}. Let us consider a simple cyclic inflation model, where the universe is mostly dominated by radiation, and the cycles are (approximately) periodic in energy densities. This follows if firstly, we assume that quantum gravitational effects trigger a bounce whenever some critical Planckian energy density is reached. Secondly, we need a $-$ve cosmological constant (CC), $-\\La$, which ensures that the universe turns around and starts to re-collapse once the radiation energy density dilutes and becomes equal to $\\La$. Thus contrary to common expectations, in the presence of matter the universe does not get stuck in an AdS vacua~\\cite{linde,BM}, but rather starts to cycle. These cycles are typically short, the time period, $\\tau=\\al{M_p/ \\sqrt{\\La}}$, where $\\al\\sim \\cO(1)$ and $M_p=2.4\\times 10^{18}$~GeV~\\cite{BM}. We shall show that in order to obtain the correct amplitude of CMB fluctuations we will require $\\La$ to be close to the conventional string/GUT scale, $\\La^{1/4} \\sim 10^{-3}M_p$, so that $\\tau\\sim 10^6 M_p^{-1}$. Now, provided there is exchange of energy between radiation and some other forms of matter, then, as a natural consequence of the second law of thermodynamics, one expects the cycles to be asymmetric. The total entropy in the universe increases monotonically, and by the same factor in every cycle: $S_{n+1}/S_n\\equiv 1+3\\ka$. Since entropy is proportional to the volume this means that if, for instance, we track the scale factor at the bounce point of consecutive cycles then \\be \\label{kappa} {a_{n+1}}/{a_n}= 1+\\ka\\for \\ka\\ll1 \\ee While the above scenario can be realized in many different ways, here we will consider a simple toy model with two species, massless radiation and some massive particle which interact with each other. It is clear from (\\ref{kappa}) that, over many asymmetric cycles the evolution of the universe looks very similar to that of ordinary inflation with an average Hubble expansion rate $ H_{\\mt{av}}={\\ka/\\tau}$. We will see later that this ``cyclic inflationary'' phase can indeed address the usual cosmological puzzles such as isotropy, horizon, flatness and homogeneity. \\begin{figure}[htbp] \\begin{center} \\includegraphics[width=0.35\\textwidth,angle=0,scale=0.8]{winfig1.eps} \\end{center} \\caption{Typical potentials: In the positive energy side three different possibilities are consistent with our model. \\label{fig:potential}} \\end{figure} What about the spectrum of the primordial fluctuations? To match the COBE normalization, the power-spectrum associated with metric fluctuations must be given by: $\\cP_{\\Phi}\\sim 10^{-10}$. Now in general, since matter couples very weekly to gravity, in the sub-Hubble phase (when the wavelength of a given comoving mode is smaller than the cosmological time-scale), when the metric fluctuations are generated from the matter fluctuations the amplitude is suppressed by the Planck scale. Typically we have \\be \\cP_{\\Phi}\\propto k^3\\Phi_k^2\\sim {\\rho/ M_p^4}\\sim 10^{-10}\\,. \\label{amplitude} \\ee Once the wavelength becomes larger than the cosmological time scale, the metric fluctuations effectively freeze at the value of $\\rho$ when the particular mode crosses the Hubble radius. This intuitive picture will essentially let us argue why the perturbations in our model will have a near scale invariant spectrum with a distinctive periodic feature (in $\\ln k$) that, in fact, provides a significant improvement in the WMAP 7-yr fit, and hence may be detectable in the future experiments. The two most important parameters governing the physics are $\\La$ and $\\ka$. While the former determines the amplitude of fluctuations, the latter characterizes the wiggles on top of the near scale-invariant spectrum. Finally, let us discuss the graceful exit problem in this inflationary scenario. If we are stuck in a $-$ve CC, then the above inflationary phase persists forever and one can never obtain an universe like ours. Fortunately, one can exit the inflationary phase if instead of a negative cosmological constant we have a dynamical scalar field whose potential {\\it interpolates} between a negative and a positive cosmological constant as depicted in Fig.~\\ref{fig:potential}. Since $V(\\phi)\\ra -\\La$ as $\\phi\\ra \\infty$, we can realize the inflationary phase, but the scalar field keeps rolling towards smaller $\\phi$ and eventually there comes a (last) cycle, in this last contraction phase the scalar field gains enough energy to zoom through the minimum and reach the $+$ve CC phase. The paper is organized as follows: In the following section~\\ref{sec:background}, we describe our toy model and the background evolution which mimics the inflationary space-time. Next, in section~\\ref{sec:perturbations}, we explain why we expect to get a nearly scale-invariant spectrum with characteristic small wiggles in our model. We also try to fit our model with the WMAP data and report our findings. In~\\ref{sec:conclusions}, we end with a discussion of the standard cosmological puzzles in the context of our model and an outlook towards future research directions. ", "conclusions": "" }, "1003/1003.5616_arXiv.txt": { "abstract": "By numerically solving the relativistic Boltzmann equations, we compute the time scale for relaxation to thermal equilibrium for an optically thick electron-positron plasma with baryon loading. We focus on the time scales of electromagnetic interactions. The collisional integrals are obtained directly from the corresponding QED matrix elements. Thermalization time scales are computed for a wide range of values of both the total energy density (over 10 orders of magnitude) and of the baryonic loading parameter (over 6 orders of magnitude). This also allows us to study such interesting limiting cases as the almost purely electron-positron plasma or electron-proton plasma as well as intermediate cases. These results appear to be important both for laboratory experiments aimed at generating optically thick pair plasmas as well as for astrophysical models in which electron-positron pair plasmas play a relevant role. ", "introduction": " ", "conclusions": "" }, "1003/1003.0033_arXiv.txt": { "abstract": "{ When clusters of galaxies are viewed in projection, one cannot avoid picking up a fraction of foreground/background interlopers, that lie within the virial cone, but outside the virial sphere. Structural and kinematic deprojection equations are known for the academic case of a static Universe, but not for the real case of an expanding Universe, where the Hubble flow (HF) stretches the line-of-sight distribution of velocities. Using 93 mock relaxed clusters, built from the dark matter (DM) particles of a hydrodynamical cosmological simulation, we quantify the distribution of interlopers in projected phase space (PPS), as well as the biases in the radial and kinematical structure of clusters produced by the HF. {The stacked mock clusters are well fit by an $m$=5 Einasto DM density profile (but only out to 1.5 virial radii), with velocity anisotropy (VA) close to the Mamon-{\\L}okas model with characteristic radius equal to that of density slope $-2$. The surface density of interlopers is nearly flat out to the virial radius, while their velocity distribution shows a dominant gaussian cluster-outskirts component and a flat field component. This distribution of interlopers in PPS is nearly universal, showing only small trends with cluster mass, and is quantified. A local $\\kappa$=2.7 sigma velocity cut is found to return the line-of-sight velocity dispersion profile (LOSVDP) expected from the NFW density and VA profiles measured in three dimensions. The HF causes a shallower outer LOSVDP that cannot be well matched by the Einasto model for any value of $\\kappa$. After this velocity cut, which removes 1 interloper out of 6, interlopers still account for 23$\\pm1$\\% of all DM particles with projected radii within the virial radius (surprisingly very similar to the observed fraction of cluster galaxies lying off the Red Sequence) and over 60\\% between 0.8 and 1 virial radius. The HF causes the best-fit projected NFW or $m$=5 Einasto model to the stacked cluster to underestimate the true concentration measured in 3D by 6$\\pm$6\\% (16$\\pm$7\\%) after (before) the velocity cut. These biases in concentration are reduced by over a factor two once a constant background is included in the fit. The VA profile recovered from the measured LOSVDP by assuming the correct mass profile recovers fairly well the VA measured in 3D, with a slight, marginally significant, bias towards more radial orbits in the outer regions. } {These small biases in the concentration and VA of the galaxy system are overshadowed by important cluster-to-cluster fluctuations caused by cosmic variance and by the strong inefficiency caused by the limited numbers of observed galaxies in clusters.} An appendix provides an analytical approximation to the surface density, projected mass and tangential shear profiles of the Einasto model. Another derives the expressions for the surface density and mass profiles of the NFW model projected on the sphere (for future kinematic modeling). } ", "introduction": "\\label{intro} The galaxy number density profiles of groups and clusters of galaxies falls off slowly enough at large radii that material beyond the virial radius (within which these structures are thought to be in dynamical equilibrium) contribute non-negligibly to the projected view of cluster, i.e. to the radial profiles of surface density, line-of-sight velocity dispersion and higher velocity moments. In principle, this contamination of observables by \\emph{interlopers}, defined here as particles that lie within the virial cone but outside the virial sphere, is not a problem, since we know how to express deprojection equations when interlopers extend to infinity along the line-of-sight. Consider the projection equation \\begin{equation} \\Sigma(R) = \\int_{-\\infty}^{+\\infty} \\nu(r)\\, \\d s = 2\\,\\int_R^\\infty \\nu(r) {r\\,\\d r\\over \\sqrt{r^2-R^2}} \\ , \\label{projec} \\end{equation} where $\\Sigma$ and $\\nu$ are the projected and space number densities, respectively, while $R$ and $r$ are the projected and space radial distances (hereafter, radii), respectively. Equation~(\\ref{projec}) can be deprojected through Abel inversion\\footnote{Alternatively, the projection equation~(\\ref{projec}) corresponds to a convolution and can therefore be deprojected with Fourier methods (see Discussion in \\citealp{MB10} and references therein).} to yield \\begin{equation} \\nu(r) = -{1\\over \\pi}\\int_r^\\infty {\\d \\Sigma/\\d R \\over \\sqrt{R^2-r^2}}\\,\\d R \\ . \\label{deprojec} \\end{equation} The projection to infinity is explicit in equations~(\\ref{projec}) and (\\ref{deprojec}). However, the Hubble expansion complicates the picture, as the Hubble flow moves background (foreground) objects to high positive (negative) line-of-sight velocities. \\begin{figure}[ht] \\includegraphics[width=\\hsize]{vzvsDlosH0.ps} \\includegraphics[width=\\hsize]{vzvsDlosH100.ps} \\caption{Line-of-sight velocity as a function of real-space line-of-sight distance (see Fig.~\\ref{scheme}) for particles inside the virial cone obtained by stacking 93 cluster-mass halos in the cosmological simulation described in Sect.~\\ref{data} without (\\emph{top}) and with (\\emph{bottom}) the Hubble flow (1 particle in 5 is shown for clarity). The \\emph{red dashed horizontal lines} roughly indicate the effects of a radius-independent $3\\,\\sigma$ clipping. Note that the velocity-distance relation without Hubble flow shown here is not entirely realistic, because the simulation was run in the context of an expanding Universe (and cannot be run in a static Universe, for lack of knowledge of realistic initial conditions), but should be accurate enough to illustrate our point. } \\label{hubflow} \\end{figure} This is illustrated in Figure~\\ref{hubflow} which shows how the line-of-sight velocity vs. real-space distance relation is affected by the Hubble flow. The line-of-sight distances are computed as the segment length QP in Figure~\\ref{scheme}. \\begin{figure}[ht] \\includegraphics[height=\\hsize,angle=-90]{Schema_zoom.ps} \\caption{Representation of the virial cone with halo particles inside the inscribed virial sphere and interlopers outside. Also shown is our definition of projected radius (CQ) and line-of-sight distance (OP and QP respectively in the observer and halo reference frames) for a random point P. For illustrative purposes, the distance to the cone is taken to be very small, so that the cone opening angle is much larger than in reality. \\label{scheme}} \\end{figure} Now, clipping the velocity differences to, say, $\\kappa = 3$ times the cluster velocity dispersion (averaged over a circular aperture, hereafter aperture velocity dispersion), $\\sigma_v$, gets rid of all the distant interlopers. More precisely, the radius, $r_{\\rm max}$, where the Hubble flow matches $\\kappa\\,\\sigma_v$ is found by solving $H_0\\,r_{\\rm max} = \\kappa\\,\\sigma_v$ (where $H_0$ is the Hubble constant) yielding \\begin{equation} {r_{\\rm max} \\over r_{\\rm v}} = \\kappa\\,\\sqrt{\\Delta\\over 2}\\,\\left ({\\sigma_v\\over v_{\\rm v}}\\right ) \\ , \\label{rmax} \\end{equation} where $r_{\\rm v}$ is the virial radius where the mean density is $\\Delta$ times the critical density of the Universe, $\\rho_{\\rm c} = 3 H_0^2/(8\\pi G)$ (where $G$ is the gravitational constant), and where $v_{\\rm v} = \\sqrt{\\Delta/2}\\, H_0\\,r_{\\rm v}$ is the circular velocity at the virial radius. Clusters are thought to have density profiles consistent with the \\citeauthor*{NFW96} profile (\\citeyear{NFW96}, hereafter NFW) profiles, \\begin{equation} \\nu(r) = {1/\\left(\\ln2-1/2\\right)\\,\\over \\left(r/r_{-2}\\right)\\,\\left(r/r_{-2}+1\\right)^2} \\,\\left[{M\\left(r_{-2}\\right)\\over 4\\,\\pi\\,r_{-2}^3}\\right] \\ , \\label{rhoNFW} \\end{equation} where $r_{-2}$ is the radius of density slope $-2$, in number \\citep*{LMS04}, luminosity \\citep{LM03} and mass (\\citeauthor{LM03}; \\citealp*{BG03, KBM04}), with a concentration, $c=r_{\\rm v}/r_{-2}$, of 3 to 5. For isotropic NFW models, the aperture velocity dispersion is (Appendix A of \\citealp{MM07}) $\\sigma_v = \\eta \\,v_{\\rm v}$, where $\\eta\\simeq 0.62$ (weakly dependent on concentration), and equation~(\\ref{rmax}) then becomes \\begin{equation} {r_{\\rm max}\\over r_{\\rm v}} \\simeq 13.2\\,\\left ({\\kappa\\over 3}\\right )\\,\\sqrt{\\Delta\\over 100} \\ . \\label{rmax2} \\end{equation} Equation~(\\ref{rmax2}) indicates that a 3-sigma clipping will remove all material beyond 13 ($\\Delta = 100$) or 19 ($\\Delta = 200$) virial radii.\\footnote{With our chosen cosmology, the overdensity at the virial radius is $\\Delta \\simeq 100$, but many authors prefer to work with $\\Delta=200$, and we will do so too.} So, in the deprojection equation~(\\ref{deprojec}), the upper integration limit must be set to this value of $r_{\\rm max}$. Although this effective cutoff in line-of-sight distances is quite far removed from the cluster, it is not clear whether there may still be a measurable bias in the concentration of clusters that one measures by comparing the surface density distribution of galaxies in clusters with NFW models projected out to infinity. Moreover, it is not clear how accurate are such measures given the finite number of galaxies observed within clusters. Finally, it is not clear how the stretching of the velocities affects the kinematic analyses of clusters, especially in the case of nearby clusters where the opening angle of the cone is non-negligible, leading to an asymmetry between the foreground and background absolute velocity distributions. For example, is the anisotropy of the 3D velocity distribution (hereafter velocity anisotropy or simply anisotropy) \\begin{equation} \\beta(r) = 1 - {1\\over 2}\\,{\\left\\langle v_\\theta^2(r)+v_\\phi^2(r)\\right\\rangle \\over \\left\\langle v_r^2(r)\\right\\rangle} \\label{betadef} \\end{equation} or equivalently\\footnote{$\\beta(r)$ enters the Jeans equation of local hydrostatic equilibrium, while ${\\cal A}(r)$ is a more physical definition of velocity anisotropy.} \\begin{equation} {\\cal A}(r) = [1-\\beta(r)]^{-1/2} = \\left [{2\\,\\left\\langle v_r^2(r)\\right\\rangle \\over \\left\\langle v_\\theta^2(r)+v_\\phi^2(r)\\right\\rangle} \\right]^{1/2} \\label{betapdef} \\end{equation} affected by the Hubble flow? One can add interlopers beyond the virial radius as a separate component to the kinematical modeling \\citep{vanderMarel+00,Wojtak+07}. Unfortunately, we have no knowledge of the distribution of interlopers in projected phase space (projected distance to the halo center and line-of-sight velocity). This paper provides the distribution of interlopers in projected phase space (projected distance to the halo center and line-of-sight velocity) as measured on nearly 100 stacked halos from a well-resolved cosmological simulation. We additionally measure the bias in the measured surface density and line-of-sight velocity dispersion and kurtosis profiles compared to those obtained in a Universe with no Hubble flow, and estimate how this bias affects the recovered concentration and velocity anisotropy of the cluster. In this paper, we use interchangeably the terms `clusters' and `halos'. We present in Sect.~\\ref{data} the cosmological simulations we use and how the individual halos were built. In Sect.~\\ref{stack} we explain how we stack these halos. In Sect.~\\ref{stats}, we present the statistics on the halo members and the interlopers in projected phase space. Then in Sect.~\\ref{removal}, we explain how we remove the outer interlopers and show analogous statistics on the cleaned stacked halo in Sect.~\\ref{statswohivilop}. We proceed in Sect.~\\ref{conc} to measure the biases induced by the Hubble flow and the imperfect interloper removal on the estimated concentration parameter and anisotropy profile. We discuss our results in Sect.~\\ref{discuss}. ", "conclusions": "\\label{discuss} This work analyzes the distribution of particles in projected phase space around dark matter halos in cosmological simulations. The particles are split among halo particles within the virial sphere and interlopers within the virial cone but outside the virial sphere (Fig.~\\ref{scheme}). The reader should be careful that the analyses presented here cannot be directly applied to observations of clusters of galaxies, as they work with halo particles instead of galaxies within clusters, and assume the halo centers to be determined quite precisely (from real space measurements). We find a universal distribution of interlopers in projected phase space, i.e. with little dependence on halo mass (Figs.~\\ref{vhistint1}c and \\ref{univvsmass}). In particular, we note that velocity cuts cannot distinguish the quarter of particles that are interlopers from those in the virial sphere (Fig.~\\ref{ilopfracs}), as was previously noted by \\cite{Cen97}. We find that the distribution of interlopers in projected phase space displays a roughly constant surface density (Figs.~\\ref{Rhists} and \\ref{univvsrad}) and a distribution of line-of-sight velocities that is the sum of a quasi-gaussian component, caused by the halo outskirts (out to typically 8 virial radii, Fig.~\\ref{vhistint1}b) and a uniform component caused by particles at further distances from the halo (Figs.~\\ref{vhists} and \\ref{vhistint1}). The cosmological simulations allow us to optimize the ratio of maximum velocity to line-of-sight velocity dispersion that recovers the latter quantity. Although this may seem to be a circular argument (since $\\sigma_{\\rm los}(R)$ depends on the velocity cut), it has been widely used in the past, usually in iterative form, with a $3\\,\\sigma$ cutoff. We find that this cutoff is not restrictive enough and causes an overestimate of the line-of-sight velocity dispersion profile (based upon mass and velocity anisotropy models derived from the cosmological simulations): up to 10\\% for the isotropic NFW velocity cut, which is reduced to 5\\% for the ML anisotropy velocity cut (Fig.~\\ref{siglos3}). We recommend instead a velocity cut at $2.7 \\,\\sigma_{\\rm los}(R)$ on the best iterative fit to the line-of-sight velocity dispersion for the NFW model with $r_{\\rm a}=r_{-2}$ ML anisotropy. Alternatively, one can use a velocity cut at $\\kappa = 2.6$ for the $m$=5 Einasto model, modeled (again) with $r_{\\rm a} = r_{-2}$ ML anisotropy, but this underestimates the line-of-sight velocity dispersion near the virial radius (Fig.~\\ref{siglos3}). We illustrate (Figs.~\\ref{phasespacewcuts}, \\ref{Rvzdist}, \\ref{vhists}, \\ref{ilopfracs}, and \\ref{sdens}) how the distribution of particles in projected phase space is altered once the high velocity interlopers are rejected with this new velocity filter (besides limiting the line-of-sight to typically $\\pm 17\\,r_{200}$, the main effect is to remove the flat velocity component). The fraction of interlopers within the virial cone drops from 27\\% (with an observer at distance $D=90 \\, h^{-1} \\, \\rm Mpc$) to $23.1$$\\pm$$0.6\\%$ (independent of $D$ for $D\\ga 17\\,\\left\\langle r_{200}\\right\\rangle$) when the velocity cut is applied (where the uncertainty is taken from the end of Sect.~\\ref{ilopstatsaftervcut}). This fraction of interlopers can be directly inferred from the NFW or Einasto model \\begin{equation} f_{\\rm i} = {\\hat M_{\\rm p}(r_{200})-1 + \\pi\\,\\hat\\Sigma_{\\rm bg} \\over \\hat M_{\\rm p}(r_{200}) + \\pi\\,\\hat\\Sigma_{\\rm bg}} \\, \\label{fofi} \\end{equation} where $\\hat M_{\\rm p} = M_p/M_{200}$ is the projected virial mass in virial units (i.e. in units of the mass within the virial sphere), while $\\hat\\Sigma_{\\rm bg}$ is given in equation~(\\ref{Sigmabgvir}). For the NFW model, one then obtains $f_{\\rm i}$ = 26.8\\% and 24.0\\%, respectively before ($\\hat\\Sigma_{\\rm bg}=0.0286$) and after ($\\hat\\Sigma_{\\rm bg}=0.0126$) the velocity cut, while with the $m$=5 Einasto model with $c=4$, the corresponding percentages of interlopers are 25.7\\% and 22.8\\%. These theoretical predictions are in excellent agreement with the fractions obtained from the simulations. Note that the omission of the background ($\\hat\\Sigma_{\\rm bg}$) term in equation~(\\ref{fofi}) reduces $f_{\\rm i}$ by typically 10\\% in relative terms, relative to the fractions after the velocity cut. % Applying equation~(\\ref{fofi}) to models of different concentrations leads to roughly a power-law variation of $f_{\\rm i}$ with slope --0.32 (NFW) or --0.49 ($m$=5 Einasto). Therefore, \\emph{the fraction of interlopers should be (slightly) more important in the more massive halos}, since they have (slightly) lower concentrations \\citep{NFW97,MDvdB08}. In comparison, using 62 clusters from the same simulation as the one we have analyzed (we have 53 clusters in common), \\cite{Biviano+06} found that among particles selected in cones of projected radius $1.5 \\, h^{-1} \\, \\rm Mpc$ around cluster-mass halos (after their velocity cut), $18$$\\pm$$1.4$\\% of them lie outside the sphere of the same radius.\\footnote{The error is taken as their dispersion over the square root of their number of halos.\\label{error}} Their halos have a median virial radius of $r_{200} = 0.93 \\, h^{-1} \\, \\rm Mpc$ (1.08 times our median) and hence a virial mass of $M_{200} = 1.9\\times 10^{14} h^{-1} M_\\odot$. Their Figure~7 indicates that their velocity cut is roughly 1180, 1105, and $780\\,\\, \\rm km \\, s^{-1}$, at projected radii 0.6, 1.0 and $1.5 \\, h^{-1} \\, \\rm Mpc$, respectively. Since NFW concentration scales as $M^{-0.1}$ \\citep{NFW97,MDvdB08}, their median concentration should be 3.9, hence their scale radius should be $930/3.9 = 238 \\, h^{-1} \\, \\rm kpc$. Assuming an NFW model, we deduce that their median circular velocity at the scale radius is $914 \\, \\rm km \\, s^{-1}$, and find that their velocity cuts correspond to $\\kappa = 2.1, 2.2$, and 1.8, at the three projected radii chosen above. These fractions are consistent with the values of $\\hat\\kappa$ one can read off of Figure~3 of \\cite{Wojtak+07} that illustrates the same velocity cut model \\citep{dHK96}. We then considered a cone of projected size $1.5/0.93 = 1.6\\,r_{200}$. Adopting their typical $\\kappa = 2$, we then found that after a $2\\,\\sigma$ velocity cut, the fraction of particles with $r > 1.6\\,r_{200}$ is now 21.3\\%. This fraction is still marginally significantly larger than \\citeauthor{Biviano+06}'s fraction of 18\\% (assuming the same errors as above). We attribute this discrepancy to their variable $\\kappa$ velocity cut, which differs from our fixed $\\kappa$ one. \\cite{Wojtak+07} tried several interloper removal schemes and definitions (using a different $\\Lambda$CDM cosmological simulation). Their local $3\\,\\sigma$ cut leads to $20.4$$\\pm$$1.7$\\% of interlopers remaining within the virial cone. Given the quoted errors, the lower fraction of interlopers found by \\citeauthor{Wojtak+07} is marginally consistent with ours. \\emph{This fraction of 23\\% of interlopers after the velocity cut is surprisingly close to the fraction of blue galaxies} (i.e. galaxies off the Red Sequence) observed within SDSS clusters, as \\cite{YMvdB08} find roughly 22\\% of blue galaxies within SDSS clusters of masses $>10^{14} h^{-1} M_\\odot$. Admittedly, it is dangerous to match the dark matter distribution with the galaxy distribution, since galaxies are biased tracers of the matter distribution. In fact, galaxies are biased relative to dark matter halos (e.g. \\citealp{CWK06}), which in turn are biased relative to the dark matter particle distribution (e.g. \\citealp{MW96,CLMP98}). If, in the end, the SDSS galaxies analyzed by \\citeauthor{YMvdB08} are unbiased tracers of the dark matter distribution, then this close agreement would be expected if all blue galaxies are caused by projection effects. But if projections also pick up red galaxies in groups, then some blue galaxies would need to survive within the virial sphere for the match to hold. However, the \\citeauthor{YMvdB08} group finder is fairly efficient in separating groups along the line-of-sight, so we conclude that the fraction of blue galaxies within the virial sphere should be small. In other words, star formation appears to be strongly quenched when galaxies penetrate the virial spheres of clusters. When no velocity cut is performed, a maximum likelihood fit of the concentration of the projected NFW model to the projected radii of a stacked cluster of nearly $300\\,000$ particles out to $r_{200}$ ($r_{100}$) leads to a 14$\\pm$7\\% (22$\\pm$6\\%) underestimate of the true concentration parameter (Table~\\ref{cbiastab}, where most of the uncertainty comes from cosmic variance). Similar biases occur with the $m$=5 Einasto model. But after the velocity cut, these biases decrease by a factor two, and are no longer statistically significant (Table~\\ref{cbiastab}). Moreover, the inclusion in such fits of a constant background as an extra parameter also strongly decreases the bias, even when the maximum projected radius is as low as $r_{200}$ (Table~\\ref{cbiastab}). In fact, inspection of Table~\\ref{cbiastab} indicates that, for $R_{\\rm max} = r_{100}$ or $3\\,r_{200}$, the background (fixed or free) has a greater influence than the velocity filter in removing the bias on measured concentration. Surprisingly, for $R_{\\rm max} = 3\\,r_{200}$, a physically motivated fixed background added to the NFW model is slightly less effective in reducing the concentration bias than is a free background. When the maximum radius is $3\\,r_{200}$ and no velocity cut is performed, the NFW model with a free (respectively fixed) background underestimates the concentration (Table~\\ref{cbiastab}) by $2$$\\pm$$11\\%$ ($15$$\\pm$$9\\%$). This insignificant (marginally significant) bias is caused by the strong decrease of the surface density profile once the Hubble flow is added to the peculiar velocities (Fig.~\\ref{hbias}). These small biases suggest that the fairly low concentration ($c_{200}=2.9$$\\pm$$0.2$) for the galaxy distribution in clusters found by \\cite{LMS04}, who fit an NFW model with a free constant background to the distribution of projected radii in the range $0.02 < R/r_{200} < 2.5$, but who did not make a velocity cut for lack of velocity data, is incompatible with true cluster concentrations of $c=4.0$ at the $2\\,\\sigma$ level. The lower concentration bias with the two-component model is expected, because the single component NFW or Einasto models cannot capture the flat surface density at large radii (Fig.~\\ref{sdens}), because other halos are projected along the line-of-sight. While a two-component model of halo (to infinity) + constant background is better able to recover the halo concentration than a single-component model (Table~\\ref{cbiastab}), it is not wise to estimate the halo concentration from a two-component model with a halo term whose line-of-sight is limited to the sphere (Appendix~\\ref{appnfwsph}) plus a near constant background term arising from our universal interloper surface density model (eq.~[\\ref{SigmaivsA}] or simply [\\ref{SigmaiNFW}]): the single-component NFW captures better the \\emph{total} surface density profile than this halo+background model, especially if the maximum projected radius is beyond the virial radius, as the interloper surface density has a discontinuous slope at the virial radius (Fig.~\\ref{sdens}). On the other hand, the universal distribution of interlopers in projected phase space might be useful to model the internal kinematics (hence total mass profile) of clusters of galaxies, where the full distribution of galaxies in projected phase space is the sum of these interlopers and an NFW-like model projected onto the virial sphere. We are preparing tests of the mass/anisotropy modeling of clusters, groups, and galaxies (through their satellites) using this interloper model. We also performed 2D fits to individual halos of typically 700 particles (not shown here). The dispersion of the concentrations were much larger (typically 0.16 dex) than the biases obtained from the stacked virial cone (typically 0.05 dex, i.e. 10\\% errors, see Table~\\ref{cbiastab}), which means that shot noise and cosmic variance dominate the bias caused by the Hubble flow. The line-of-sight velocity dispersion profile shows a concavity (in log-log) near the virial radius (Fig.~\\ref{siglos3}), which is caused by the Hubble flow (Fig,~\\ref{hbias}). The velocity anisotropy profile recovered from this velocity dispersion profile, assuming the correct mass distribution, is close to the true anisotropy profile, with a slight, marginally significant, radial bias in the envelopes of clusters in comparison with the anisotropy profile recovered in 3D (Fig.~\\ref{betainv}), as was previously noted by \\cite{Biviano07}. In summary, the density profile of $\\Lambda$CDM halos falls fast enough that the effects of the Hubble flow perturbing the standard projection equations produce only small biases in comparison with the shot noise of clusters with less than 1000 galaxies, as well as the large cosmic variance of the halos. These results have been obtained with the dark matter particles of a cosmological $N$-body simulation (with additional gas and galaxy components). They will need to be confirmed with future more realistic simulations of the \\emph{galaxy} distribution." }, "1003/1003.5085_arXiv.txt": { "abstract": "Using cosmological MHD simulations of the magnetic field in galaxy clusters and filaments we evaluate the possibility to infer the magnetic field strength in filaments by measuring cross-correlation functions between Faraday Rotation Measures (RM) and the galaxy density field. We also test the reliability of recent estimates considering the problem of data quality and Galactic foreground (GF) removal in current datasets. Besides the two self-consistent simulations of cosmological magnetic fields based on primordial seed fields and galactic outflows analyzed here, we also explore a larger range of models scaling up the resulting magnetic fields of one of the simulations. We find that, if an unnormalized estimator for the cross-correlation functions and a GF removal procedure is used, the detectability of the cosmological signal is only possible for future instruments (e.g. SKA and ASKAP). However, mapping of the observed RM signal to the underlying magnetization of the Universe (both in space and time) is an extremely challenging task which is limited by the ambiguities of our model parameters, as well as to the weak response of the RM signal in low density environments. Therefore, we conclude that current data cannot constrain the amplitude and distribution of magnetic fields within the large scale structure and a detailed theoretical understanding of the build up and distribution of magnetic fields within the Universe will be needed for the interpretation of future observations. ", "introduction": "\\label{sec:intro} Magnetic fields in the Universe are found in almost all studied environments. In particular, their presence in the inter-galactic medium \\citep[IGM; see][ for a recent review]{2009ASTRA...5...43B} and in the intra-cluster medium (ICM) is confirmed by diffuse radio emission as well as by observations of Faraday Rotation Measures (RM) towards polarized radio sources within or behind the magnetized medium \\citep[e.g.][]{2006AN....327..539G}. On the largest scales, like those of filaments, magnetic fields are notoriously difficult to measure and available data is still incomplete. This is especially difficult because these measurements require either a high thermal density (for RMs) or the presence of relativistic particles (for the synchrotron emission). Therefore, measurements of the magnetic field strength have been successfull for high density regions of collapsed objects (e.g. galaxies and galaxy clusters), and thus, fields significantly below the $\\mu$G level can hardly be detected. Recently an interesting attempt to constrain the value of large scale cosmic magnetic fields was done by \\citet{2009arXiv0906.1631L}. These authors detected a positive cross-correlation signal between the galaxy distribution in the SDSS Sixth Data Release \\citep{2008ApJS..175..297A} and the RM values extracted from the \\citet{2009ApJ...702.1230T} catalog. Using the amplitude of this signal, together with a simplified model for the magnetic fields configuration in the Universe (estimated from its mean electron density), and computing the RM typical values expected from this coherent field in a given length scale, they were able to derive limits for the corresponding cosmic magnetic fields. In this work, we want to investigate: ({\\it i}) to what extent a self-consistent treatment of the cosmological RM signal based on magneto-hydrodynamical (MHD) simulations of structure formation changes the expected shape and amplitude of such a correlation signal, and ({\\it ii}) how such an approach is affected by the presence of the Galactic foreground (GF) and noise in the final RM signal. Both points are of extreme importance, if robust field properties are to be derived from any observed signal. Furthermore, the appearance of magnetic field reversals (as observed in galaxy clusters at various length scales) will alter the cosmological signal magnitude and shape, whereas the residuals of any foreground and measurement errors will bias the relation between the amplitude of the correlation function and the underlying cosmological field. In order to self-consistently treat the cosmic magnetic fields, we make use of several cosmological MHD simulations which compute the resulting magnetization of the cosmological structures (e.g. amplitude and structure) following different models for the origin and seeding process of such magnetic fields. We also construct magnetic field models with much higher magnetization amplitude in the low density regions to test how the resulting signatures of more extreme models affect our results. Here we scale up the predicted amplitude of the magnetic field in filaments by several orders of magnitude to test if such strong magnetic fields in low density regions significantly effect the expected correlation signal. By introducing GF and adding noise to the signal on top of the underlying cosmological signal, we can study how the shape and amplitude of the cross-correlation function would be modified when considering actual observations. To avoid further complications we ignore the cosmological evolution of magnetic fields, which, in principle, would be consistently treated within our cosmological MHD simulations. Hence, we neglect the evolution of the cosmic magnetic field seen in the simulation as a result of the structure formation process, and assume the present day magnetization of the simulated universe to be present up to the redshift of the sources. The paper is organized as follows. In Section \\ref{sec:simul} we describe the cosmological MHD simulations used and how we compute the synthetic RM catalogs. In Section \\ref{sec:cross} we discuss the cross-correlation estimators used, the estimation of the intrinsic uncertainties due to the limited number of lines of sight probing the magnetization of the cosmological structures, the different signals expected for the various magnetization of the universe, as well as the uncertainties induced by the redshift distribution of the sources. In Section \\ref{sec:obs} we show how the shape and amplitude of the signal is affected by the recipe normally used to remove the foreground signal, due to observational noise and to the Galaxy itself. In Section \\ref{sec:res} we summarize the combination of all the effects, and present the resulting observable signal of the different magnetic field models. Finally, our conclusions are given in Section \\ref{sec:conc}. ", "conclusions": "\\label{sec:conc} Using cosmological MHD simulations of the magnetic field in galaxy clusters and filaments we evaluated the possibility to infer the magnetic field strength in filaments by measuring cross-correlation functions between RMs and the galaxy density field. We find that the shape of the cross-correlation function using the normalized estimator $\\omega_{RM}$ (in absence of any noise or foreground signal) nicely reflects the underlying distribution of magnetic field within the large scale structure. However, a very large number of lines of sight probed by RM measurements (much more than the 3072 used in this investigation) are needed to overcome the statistical noise induced by the particular magnetic field realization within the cosmic structures, in order to distinguish between the wide range of models we used here. In general, the RM signal is strongly dominated by the denser regions (e.g. those populated by galaxy clusters and groups) and not by the low density ones, like filaments. On this point, the magnetic field associated with filaments already changes by several orders of magnitudes within the different models used here. Aditionally, the normalized estimator $\\omega_{RM}$ is extremely sensitive to measurement errors and to the presence of the GF (despite attempts to remove it by subtracting a smoothed map). It is fair to say that given the current measurement errors in the available RMs and our knowledge of the GF, present studies cannot determine the magnetization magnitude of the Universe based only on the cross-correlation $\\omega_{RM}$, whatever the significance of the measured signal is. On the contrary, the shape of the unnormalized estimator $\\xi_{RM}$ (the same as used by Lee et al. 2009) is relatively insensitive against the presence of measurement errors for the RMs and for the presence of the GF (as long as the described removal technique is used). Its amplitude, however, is quite strongly affected by measurement uncertainties. Current measurement errors (as for example those inherited by the Taylor's published sample) suppress the signal by a significant amount in such a way that it is impossible to relate the amplitude of the cross-correlation function to the underlying magnetization of the the large scale structure. However, we expect that future radio telescopes will be able of reaching error magnitudes of order of 1 rad m$^{-2}$ that could make the correction of the signal possible. Unfortunately, this estimator does not nicely encode in its shape the details of the magnetization of the large scale structure and, especially, its amplitude is extremely sensitive to the {\\it magnetic depth} of the Universe. Therefore, any interpretation of an observed signal is limited by our knowledge of the redshift distribution of the sources (towards the RM signals measured), as well as by our knowledge of the distribution and evolution of the cosmic universal magnetization. Future observational data will help to put better constraints on theoretical models for the origin of cosmological magnetic fields which, in return, can be implemented in next generation of MHD cosmological simulations in order to draw a self-consistent picture that can be compared against observations. In summary, we conclude that current RM observations cannot constrain the amplitude and distribution of magnetic fields within the large scale structure. On the other hand, future datasets, based on a larger number of observations with more accurate RMs, might be able to shed light on the magnetic field distribution and evolution within these structures. However, very detailed model predictions are needed in order to compare with any observed cross-correlation signal. It will be a quite demanding task for future cosmological simulations to provide detailed enough information of the large scale structure magnetization process within a large enough volume to produce useful templates of such correlation functions which can then be compared directly to the observations." }, "1003/1003.0496_arXiv.txt": { "abstract": "The Sunyaev-Zel'dovich Effect (SZE) has been observed toward six massive galaxy clusters, at redshifts $0.091{\\leq}z{\\leq}0.322$ in the 86-102 GHz band with the Y. T. Lee Array for Microwave Background Anisotropy (AMiBA). We modify an iterative method, based on the isothermal $\\beta$-models, to derive the electron temperature $T_{\\rm e}$, total mass $M_{\\rm t}$, gas mass $M_{\\rm g}$, and integrated Compton $Y$ within $r_{2500}$, from the AMiBA SZE data. Non-isothermal universal temperature profile (UTP) $\\beta$ models are also considered in this paper. These results are in good agreement with those deduced from other observations. We also investigate the embedded scaling relations, due to the assumptions that have been made in the method we adopted, between these purely SZE-deduced $T_{\\rm e}$, $M_{\\rm t}$, $M_{\\rm g}$ and $Y$. Our results suggest that cluster properties may be measurable with SZE observations alone. However, the assumptions built into the pure-SZE method bias the results of scaling relation estimations and need further study. ", "introduction": "\\label{sec:intro} The Sunyaev-Zel'dovich Effect (SZE) is an useful tool for studies of galaxy clusters. This distortion of the Cosmic Microwave Background (CMB) is caused by the inverse Compton scattering by high energy electrons as the CMB propagates through the hot plasma of galaxy clusters \\citep{SZ1972}. The SZE signal is essentially redshift independent, making it particularly useful for determining the evolution of large-scale structure. For upcoming SZE cluster surveys \\citep{Ruhl2004,Fowler2004,Kaneko2006,Ho2008}, it is important to investigate the relations between SZE flux density and other cluster properties, such as mass, temperature, and gas fraction. By assuming that the evolution of clusters is dominated by self-similar gravitational processes, we can predict simple power law relations between integrated Compton $Y$ and other cluster properties \\citep{Kasier1986}. Strong correlations between integrated SZE flux and the mass of clusters are also suggested by numerical simulations \\citep{dasilva2004,Motl2005,Nagai2006}. These relations imply the possibility of determining the masses and temperatures of clusters, and investigating cluster evolution at high redshift, with SZE observation data alone. \\citet{Joy2001} and \\citet{Bonamente2007} demonstrated an iterative approach based on the isothermal $\\beta$ model to estimate the values of electron temperature $T_e$, total mass $M_t$, gas mass $M_g$, and Compton-$Y$ from SZE data alone. In this paper, we seek to derive the same cluster properties from the AMiBA SZE measurements of six clusters. Due to the limited $u-v$ space sampling, the AMiBA data do not provide useful constraints on the structural parameters, $\\beta$ and $r_c$, in a full iterative model fitting. Instead, we adopt $\\beta$ and $r_c$ from published X-ray fits and use a Markov Chain Monte-Carlo (MCMC) method to determine the cluster properties ($T_e, M_t, M_g$ and $Y$). We also estimate these cluster properties from AMiBA data with structural constraints from X-ray data using the non-isothermal universal temperature profile model \\citep{Hallman2007}. All quantities are integrated to spherical radius $r_{2500}$ within which the mean over-density of the cluster is $2500$~times the critical density at the cluster's redshift. We then investigate the scaling relations between these cluster properties derived from the SZE data, and identify correlations between those properties that are induced by the iterative method. We note that \\citet{Locutus2008} investigate the scaling relations between the values of Compton $Y$ from AMiBA SZE data and other cluster properties from X-ray and other data. All results are in good agreement. However, we are concerned that there are embedded relations between the properties we derived using this method. Therefore, we also investigate the embedded scaling relations between SZE-derived properties as well. We assume the large-scale structure of the Universe to be described by a flat $\\Lambda$CDM model with $\\Omega_{\\rm m} = 0.26$, $\\Omega_{\\rm \\Lambda} = 0.74$, and Hubble constant $H_{\\rm 0} = 72 \\ \\rm km \\, s^{-1} \\, Mpc^{-1}$, corresponding to the values obtained using the WMAP 5-year data \\citep{WMAP5}. All uncertainties quoted are at the 68\\% confidence level. ", "conclusions": "\\label{sec:discuss} We derived the cluster properties, including $T_{\\rm e}$, $r_{2500}$ ,$M_{\\rm t}$, $M_{\\rm g}$ and $Y$, for six massive galaxy clusters ($M_{\\rm t}(r_{2500})>2\\times 10^{14}M_{\\odot}$) mainly based on the AMiBA SZE data. These results are in good agreement with those obtained solely from the OVRO/BIMA SZE data, and those from the joint SZE-X-ray analysis of Chandra-OVRO/BIMA data. In the comparison, the SZE-X-ray joint analysis gives smaller error bars than the pure SZE results, because currently the uncertainty in the measurement of the SZE flux is still large. On the other hand, in our current SZE-based analysis, due to the insufficient $u$-$v$ coverage of the 7-element AMiBA we still need to use X-ray parameters for the cluster model i.e., the $\\beta$ and $\\theta_{\\rm c}$ for the $\\beta$-model. However, \\citet{apex2009} have deduced $\\beta$ and $\\theta_{\\rm c}$ from an APEX SZE observation alone recently. For AMiBA, the situation will be improved when it expands to its 13-element configuration with 1.2m antennas \\citep[AMiBA13;][]{Ho2008}, and thus much stronger constraints on the cluster properties than current AMiBA results are expected. Furthermore, with about three times higher angular solution, we should be able to estimate $\\beta$ and $\\theta_{\\rm c}$ from our SZE data with AMiBA13 and make our analysis purely SZE based \\citep{Ho2008,Sandor2008}. Nevertheless, the techniques of using SZE data solely to estimate cluster properties are still important, because many upcoming SZE surveys will observe SZE clusters for which no X-ray data are available \\citep{Ruhl2004,Fowler2004,Kaneko2006,Ho2008}, especially for those at high redshifts. \\citet{Hallman2007} suggested that adopting the UTP $\\beta$ model for SZE data on galaxy clusters will reduce the overestimation of the integrated Compton $Y_{500}$ and gas mass. However, the $Y_{2500}$ values we obtained with the UTP model are not smaller than those obtained with the isothermal model. The $M_{g}(r_{2500})$ values deduced using the UTP model are even larger than those deduced using the isothermal model. For the case of integrated Compton Y, when we compare $Y_{500}$ deduced using the UTP model $Y_{500,\\rm UTP}$, and those deduced using the isothermal model $Y_{500,\\rm iso}$, we found that the $Y_{500,\\rm UTP}$ are smaller than $Y_{500,\\rm iso}$, as predicted by \\citet{Hallman2007}. The reason is that the Compton $y$ profile predicted using the UTP $\\beta$ model will decrease more quickly than the profile predicted by the isothermal $\\beta$ model, with increasing radius. Therefore, the ratio $Y_{\\Delta,\\rm UTP}/Y_{\\Delta,\\rm iso}$ will decrease as $\\Delta$ decreases. We also noticed that the electron temperature values obtained with the isothermal model are significantly higher than the temperatures deduced from X-ray data for most clusters we considered. The temperatures of clusters obtained using the UTP model are lower than those obtained with the isothermal model and thus are in better agreement with those deduced from X-ray data. Therefore, in the UTP model, with similar $Y_{2500}$ and lower temperature, we should get larger $M_{g}$. The electron temperatures derived using the UTP $\\beta$ model are in better agreement with X-ray observation results than those derived using the isothermal $\\beta$ model. This result implies that the UTP $\\beta$ model may provide better estimates of the electron temperature when we can use only the $\\beta$ model parameters from X-ray observation. However, we noticed that the UTP $\\beta$ model produced larger errorbars than the isothermal $\\beta$ model did. These increased errors are based on the uncertainties of $\\beta$ and $r_{c}$ which we insert by hand. On the other hand, because we treat $\\beta$ and $r_{c}$ as independent parameters in this work, the uncertainty could be over estimated due to the degeneracy between these two parameters. If we can access to the likelihood distributions of $\\beta$ and $r_{c}$ of the UTP $\\beta$ model derived from observation, the error-bars might be reduced significantly. There is a concern that the scaling relations among the purely SZE-derived cluster properties may be implicitly embedded in the formalism we used here. In this paper, we also investigate for the first time the embedded scaling relations between the SZE-derived cluster properties. Our analytical and numerical analyses both suggest that there are embedded scaling relations between SZE-derived cluster properties, with both the isothermal model and the UTP model, while we fix $\\beta$. The embedded $Y$-$T$ and $M$-$T$ scaling relations are close to the predictions of self-similar model. The results imply that the assumptions built in the pure-SZE method significantly affect the scaling relation between the SZE-derived properties. Therefore, we should treat those scaling relations carefully. Our results suggest the possibility of measuring cluster parameters with SZE observation alone. The agreement between our results and those from the literature provides not only confidence for our project but also supports our understanding of galaxy clusters. The upcoming expanded AMiBA with higher sensitivity and better resolution will significantly improve the constraints on these cluster properties. In addition, an improved determination of the $u$-$v$ space structure of the clusters directly from AMiBA will make it possible to measure the properties of clusters which currently do not have good X-ray data. The ability to estimate cluster properties based on SZE data will improve the study of mass distribution at high redshifts. On the other hand, the fact that the assumptions of cluster mass and temperature profiles significantly bias the estimations of scaling relations should be also noticed and treated carefully." }, "1003/1003.2214_arXiv.txt": { "abstract": "Neutral hydrogen is ubiquitous, absorbing and emitting 21~cm radiation throughout much of the Universe's history. Active sources of perturbations, such as cosmic strings, would generate simultaneous perturbations in the distribution of neutral hydrogen and in the Cosmic Microwave Background (CMB) radiation from recombination. Moving strings would create wakes leading to 21~cm brightness fluctuations, while also perturbing CMB light via the Gott-Kaiser-Stebbins effect. This would lead to spatial correlations between the 21~cm and CMB anisotropies. Passive sources, like inflationary perturbations, predict no cross correlations prior to the onset of reionization. Thus, observation of any cross correlation between CMB and 21~cm radiation from dark ages would constitute evidence for new physics. We calculate the cosmic string induced correlations between CMB and 21~cm and evaluate their observability. ", "introduction": "\\label{sec:intro} Neutral hydrogen emits and absorbs radiation at 21~cm in its rest frame throughout the history of the Universe. The brightness of the 21~cm line observed today from various redshifts is determined by the spatial distribution and peculiar velocity of the neutral hydrogen that emitted it. The radiation from any particular redshift has small angular variations, similar to those in the Cosmic Microwave Background (CMB) radiation. There is also variation along the line of sight, with a correlation depth of approximately $10$~Mpc~\\cite{Eisenstein:2005su,Furlanetto:2006jb}. In the absence of physics that violates free streaming, one does not expect a correlation among slices in the line-of-sight direction that are separated by more than this correlation length. Since observations can measure the redshift of 21~cm emission precisely, we should be able to extract information about this line of sight variation from 21~cm observations. It is this third dimension of information encoded in 21~cm radiation that can make it a powerful new tool to learn about fundamental physics. One path forward is to look for physics that violates free streaming and generates correlations among distinct redshift shells~\\cite{Furlanetto:2006jb}. Several sources of such a signal have been identified, though they are limited to the epoch during and after reionization~\\cite{Alvarez:2005sa,21cm-xcorr,Adshead:2007ij,Sarkar:2008sz}. Though passive perturbations, such as those seeded by inflation, cannot generate such correlations, active sources of perturbation could. Active sources of perturbation include cosmic strings and other topological defects, such as textures \\cite{Turok:1989ai}, as well as the field gradients that are left behind after global phase transitions \\cite{JonesSmith:2007ne}. Of these, cosmic strings are the most thoroughly studied and have perhaps the richest phenomenology \\cite{Myers:1900zz}, so they will be our focus. Moving strings generate density wakes in the material through which they pass \\cite{wakes}, including neutral hydrogen. As a result, the brightness of the 21~cm emission, which depends on the density and the velocity of the emitters, will be directly perturbed in the string's wake~\\cite{Khatri:2008zw}. The same moving strings also perturb the background CMB light, monotonically shifting the spectra and creating line discontinuities in the temperature screen of the sky---a phenomenon known as the Gott-Kaiser-Stebbins (GKS) \\cite{gott,kaiser} effect. GKS is a special case of the Integrated Sachs-Wolfe (ISW) effect \\cite{Sachs:1967er} caused by time-varying gravitational potentials sourced by moving strings. Thus, there will be some correlation between string-sourced 21~cm brightness fluctuations at any redshift and a part of the CMB temperature anisotropy. Strings will also induce some cross-correlation in 21~cm fluctuations coming from different redshifts. However, as we discuss in Section~\\ref{sec:physics}, we expect this effect to be relatively small. Precision measurements of the CMB temperature anisotropy limit the contribution of cosmic string perturbations to be less than about 10\\% of the total over the range of scales covered by WMAP~\\cite{Wyman:2005tu}. As such, the bulk of the Universe's anisotropy comes from the primordial inflationary fluctuations, for which the cross-correlation between sufficiently-separated shells is zero. The fact that any string contribution to primordial anisotropy must be small means that the intrinsic signal for which we are searching will be faint. However, since there is nothing in the standard $\\Lambda$CDM model that could generate such cross correlations before the epoch of reionization, we can be sure that discovery of such a cross correlation would be a sign of new physics. To examine the prospects for detecting these string-induced correlations, we will evaluate the signal-to-noise density of this effect and the volume over which the signal can be found. Our aim is to identify the optimal way of looking for stringy phenomena in the 21~cm radiation from the dark ages. This effect is present over a vast volume---from $z\\lse100$ until the onset of reionization---which may make a detection of the CMB/21cm correlation feasible under optimistic observing scenarios. However, because of uncertainties in reionization physics, we exclude a large region ($z<20$) from our present study. We will also leave the calculation of 21cm $\\times$ 21cm cross-correlation from different redshifts for future work. In this paper we adopt the string network model of \\cite{vectorstring,ABR97}, which was previously implemented in CMBFAST \\cite{cmbfast} and made available publicly as CMBACT~\\cite{cmbact}. For this work, we have implemented this string model in CAMB~\\cite{Lewis:2007kz}. Our code will be made publicly available soon \\cite{ustobe}. ", "conclusions": "By incorporating the segment model of strings into the CAMB\\_sources Einstein-Boltzmann code, we have calculated the cross correlation between the 21~cm brightness temperature and the CMB. Our chief result is plotted in Fig.~\\ref{fig:s2n-gmu}, which represents the signal-to-noise available in CMB-21~cm correlation studies given an all-sky survey, a string tension $G\\mu$, and our observing strategy, which we characterize by a single integration time and telescope properties parameter, $x$. Unfortunately, noise from the sky temperature overwhelms the signal we are calculating for observationally-allowed cosmic string tensions. However, there is room for more work. We have thus far only included the signal from the CMB-21~cm correlations for $z>20$, where we do not need to know the physics of reionization. Since the signal should extend until the epoch of reionization $z\\simeq10$, we may eventually be able to include $\\mathcal{O}(10^2)$ more redshift bins and $\\mathcal{O}(10^3)$ data points in the $z-\\ell$ plane, once we can accurately model reionization. Another contribution to our signal that was neglected in this study can be found from a cross-correlation study between different redshift bins, and this signal will be addressed in future work~\\cite{ustobe}. It is possible that we will find a detectable signal when these extra sources of data are included in the analysis. Though our calculations have not predicted a detectable signal, we reiterate that any detection of a cross correlation between the 21~cm radiation from $z>20$ and the CMB would be a clear sign of new physics. Active sources, such as cosmic strings, are a promising candidate for the sort of new physics that could generate this correlation. However, the fact that inflation very likely provides the dominant source of perturbations in the Universe makes detectability of any signal difficult: the inflationary perturbations must necessarily enter as ``noise'' for any study like this, and are typically much larger than any alternate sources of perturbation, at least prior to recombination where measurements of the CMB limit any non-inflationary perturbations. One way around this would be to contrive some active source that was suppressed prior to recombination which subsequently sourced a larger fraction of perturbations in the dark ages. On the other hand, a non-detection of such a cross correlation would constitute a strong constraint on active sources with such unusual properties." }, "1003/1003.4604_arXiv.txt": { "abstract": "We present the results of deep radio observations with the Australia Telescope Compact Array (\\textit{ATCA}) of the globular cluster NGC 6388. We show that there is no radio source detected (with a r.m.s. noise level of 27 $\\mu$Jy) at the cluster centre of gravity or at the locations of the any of the \\textit{Chandra} X-ray sources in the cluster. Based on the fundamental plane of accreting black holes which is a relationship between X-ray luminosity, radio luminosity and black hole mass, we place an upper limit of $\\sim$1500~M$_{\\odot}$ on the mass of the putative intermediate-mass black hole located at the centre of NGC 6388. We discuss the uncertainties of this upper limit and the previously suggested black hole mass of 5700~M$_{\\odot}$ based on surface density profile analysis. ", "introduction": "Following the early discoveries of X-ray sources in globular clusters in the mid-1970s (Clark 1975; Clark, Markert \\& Li 1975), it was proposed that the X-ray emission of these clusters was due to accretion of intracluster material released by stellar mass loss onto central black holes (Bahcall \\& Ostriker 1975; Silk \\& Arons 1975). This started a debate about whether globular clusters contain black holes of intermediate masses (i.e. greater than the $\\sim$ 30 M$_{\\odot}$ limit for black holes formed through normal single star evolution, but less than the $10^{5}$ M$_{\\odot}$ seen in the smallest galactic nuclei). The difficulties of stellar dynamics has prompted a search for accretion constraints on the presence of intermediate-mass black holes. As pointed out by Maccarone (2004) and Maccarone et al. (2005), deep radio searches may be a very effective way to detect intermediate-mass black holes in globular clusters and related objects. Indeed, for a given X-ray luminosity, supermassive mass black holes produce far more radio luminosity than stellar-mass black holes. The relation between black hole mass and X-ray and radio luminosity empirically appears to follow a \"fundamental plane\", in which the ratio of radio to X-ray luminosity increases as the $\\sim$0.8 power of the black hole mass (Falcke \\& Biermann 1996, 1999; Merloni, Heinz \\& Di Matteo 2003; Falcke, Kording \\& Markoff 2004). Also, as the luminosity of accretion onto a black hole decreases, the ratio of radio to X-ray power increases \\citep{Corbel03,gallo}. Accretion theory suggests that the Bondi-Hoyle rate \\citep{bondi} overestimates the actual accretion rate by 2-3 orders of magnitude, e.g. \\citep{perna}. Thus, the X-ray luminosities from accretion of the interstellar medium by intermediate-mass black holes in globular clusters are likely to be well below detection limits of current X-ray observatories. Considering the prediction of \\citet{miller}, that the black holes should have about 0.1 per cent of the total cluster mass, the radio luminosities of the brightest cluster central black holes may be detectable with existing instrumentation \\citep{mak3}. Several methods have been considered for proving the existence of these intermediate-mass black holes (in the $10^{2}-10^{4}$ M$_{\\odot}$ range), but to date there is no conclusive evidence for their existence. Searches for radio emission from globular clusters have mostly yielded only upper limits (Maccarone, Fender \\& Tzioumis 2005; De Rijcke, Buyle \\& Dejonghe 2006; Bash et al. 2008). Although, the cluster G1 in M31 seems to have evidence for harbouring an intermediate-mass black hole \\citep{ulve}, including radio detection. However, beyond globular clusters, there are other possibilities for intermediate-mass black holes. They may be produced in the core collapses of $\\sim 100-1000$-M$_{\\odot}$ Population III stars, see e.g. \\citep{fryer}. Other good candidates for hosting intermediate-mass black holes are thought to be young dense star clusters (Portegies Zwart \\& McMillian 2002; Portegies Zwart et al. 2004; G\\\"{u}rkan, Freitag \\& Rasio 2004) and ultraluminous X-ray sources, whose X-ray luminosities well exceed the Eddington luminosity of a ten solar-mass compact object; for more details see \\citep{Kaaret01,zampi}. The main evidence favouring an intermediate-mass black hole in NGC 6388 is that the observed surface density profile has a power-law shape with a slope $\\alpha = -0.2$ in the inner one arcsecond of the cluster. This slope is shallower than expected for a post core collapse cluster and is consistent with the presence of an intermediate-mass black hole \\citep{surf,r}. The surface density profile provided an estimated mass of 5700 $\\pm$ 500 M$_{\\odot}$ \\citep{Lanzoni07} for the central black hole in NGC 6388 and it motivated us to propose radio observations of the source. Here, we report on radio observations with the Australia Telescope Compact Array (\\textit{ATCA}) of NGC 6388 that led to an upper limit on the mass of the putative intermediate-mass black hole located at the centre of NGC 6388. In Sec.~2, we describe the analysis of an archival \\textit{Chandra} observation and our new \\textit{ATCA} radio observations of NGC 6388. In Sec.~3, we describe the results of the X-ray and radio observations. Then we discuss, in Sec.~4, the methodology for setting an upper limit on the mass of a central black hole in NGC 6388 and the uncertainties. ", "conclusions": "\\citet{Lanzoni07} reported the possible presence of a black hole with a mass of 5700 $\\pm$ 500 M$_{\\odot}$ at the centre of the globular cluster NGC 6388. \\textit{Chandra} and \\textit{XMM-Newton} observational data analysis were carried out by \\citet{nucita}. Removing the pixel randomisation allowed us to identify a unique source coincident with the cluster centre of gravity with properties consistent with those expected for a black hole accreting at a low rate. With the X-ray detection and optical surface density fit, the only missing piece of the puzzle was a radio detection. On the basis of the X-ray luminosity, \\citet{nucita} predicted an upper limit of $<$ 3 mJy radio flux on NGC 6388. Deep radio observations with the Australia Telescope Compact Array allowed us to reach a sensitivity of 27$\\mu$Jy/beam, but did not reveal any radio sources within the cluster. We interpreted the radio non-detection by using the fundamental plane relating the radio and X-ray properties of black holes accreting at low rates, assuming that the X-ray flux is related to a black hole accretion luminosity. We obtained an upper limit on the black hole mass of $M < $~735~M$_{\\odot}\\pm$~244~M$_{\\odot}$ (1$\\sigma$). Taking into account the uncertainties on the radiative efficiency of accretion and on the fraction of Bondi-Hoyle accretion rate reaching the black hole, we concluded that the centre of NGC 6388 can not host a black hole with a mass in excess of $1500$ M$_{\\odot}$ at a 3$\\sigma$ confidence level." }, "1003/1003.1217_arXiv.txt": { "abstract": "A two-dimensional electromagnetic particle-in-cell simulation with the realistic ion-to-electron mass ratio of 1836 is carried out to investigate the electrostatic collisionless shocks in relatively high-speed ($\\sim 3000$ km s$^{-1}$) plasma flows and also the influence of both electrostatic and electromagnetic instabilities, which can develop around the shocks, on the shock dynamics. It is shown that the electrostatic ion-ion instability can develop in front of the shocks, where the plasma is under counter-streaming condition, with highly oblique wave vectors as was shown previously. % The electrostatic potential generated by the electrostatic ion-ion instability propagating obliquely to the shock surface becomes comparable with the shock potential and finally the shock structure is destroyed. It is also shown that in front of the shock the beam-Weibel instability gradually grows as well, consequently suggesting that the magnetic field generated by the beam-Weibel instability becomes important in long-term evolution of the shock and the Weibel-mediated shock forms long after the electrostatic shock vanished. It is also observed that the secondary electrostatic shock forms in the reflected ions in front of the primary electrostatic shock. ", "introduction": "Collisionless shock is one of the most interesting phenomena in plasma physics; it dissipates the kinetic energy of the plasma flow into the thermal energy and the electromagnetic energy not by the Coulomb collision but by the collective effect associated with the electric and the magnetic fields. For example, the universe is filled with hot, tenuous collisionless plasmas and a variety of collisionless shocks are produced due to violent phenomena, such as supernova explosions. These shocks are believed to accelerate charged particles to high energies to generate cosmic rays. The electrostatic shock \\citep{Moiseev63}, or the ion-acoustic shock \\citep{Chen}, is also one of the collisionless shocks and it forms in unmagnetized collisionless electron-ion plasmas if the Mach number is not so large and the temperature ratio of electrons to ions is relatively large \\citep{Mason72}. The electrostatic shocks were observed in various experiments with double-plasma devices \\citep{Taylor70, Ikezi73, Bailung08}, with Q-machines \\citep{Takeuchi98}, with photo-ionized plasmas \\citep{Cohn72}, and with laser plasmas \\citep{Koopman67,Romagnani08}. Recent experiments with intense lasers also showed a possible formation of the electrostatic shock at a high shock speed of $\\sim 1000$ km s$^{-1}$ \\citep{Morita09}. In space, these shocks are observed, for example, in the auroral zone of the Earth \\citep{Mozer81} as well. In the astrophysical context, there are no clear observations of the electrostatic shocks so far. However, they can also be driven in the universe at a wide range of the plasma flow speed. From the theoretical point of view, the electrostatic shock has been investigated with % hybrid or particle-in-cell (PIC) simulations extensively \\citep[e.g.,][]{Forslund70b,Mason71,Mason72}. Recent simulations also showed a possible formation of very high Mach number electrostatic shocks \\citep{Sorasio06}. The requirement of the relatively high electron-to-ion temperature ratio for the electrostatic shocks is also one of the conditions for the electrostatic ion-ion instability \\citep{Stringer64,Ohnuma65,Fried66,Gary87}, which is also called the ion/ion acoustic instability, to develop. Since there exist the reflected ions in front of the electrostatic shocks, this instability can grow to generate electrostatic waves there. The wave vector of the instability is oblique or even almost perpendicular to the streaming direction \\citep{Forslund70a,Gresillon75} and therefore multi-dimensional simulations are necessary to investigate the effect of the instability. Although most of the simulations were carried out in one dimension and therefore could not deal with the instability, \\citet{Karimabadi91} carried out a two-dimensional electrostatic hybrid simulation and showed that the electrostatic ion-ion instability indeed develops in front of the electrostatic shock and affects the structure of the shock significantly. (It should be noted that recently \\citet{Ohira08} showed that this instability can also develop in the foot region of the magnetized collisionless shocks in supernova remnants.) In astrophysical plasmas and laboratory plasmas with recent laser facilities, flows of collisionless plasma whose velocity is even faster than 1000 km s$^{-1}$ can be generated. In such high-speed flows, the electromagnetic instabilities, e.g., the Weibel-type instabilities \\citep{Weibel59, Fried59}, appears to be important as well as the electrostatic instabilities and can affect the evolution of the electrostatic shocks. The simulation by \\citet{Karimabadi91} was however an electrostatic one and the effects of such electromagnetic instabilities were therefore not included. In addition, since the simulation box size of their simulation was not large compared with the structures generated by the electrostatic instability, larger scale simulations are also desirable. In this paper, we investigate the electrostatic shocks propagating at a relatively high speed with a two-dimensional electromagnetic PIC simulation. In particular, we focus on the influence of the electrostatic ion-ion instability and the Weibel-type instabilities on the shock formation. ", "conclusions": "We have carried out a two-dimensional electromagnetic PIC simulation for the case of counter-streaming plasmas at a relatively high flow velocity with a large electron-to-ion temperature ratio of 9. At first we have confirmed that the electrostatic shock forms in the early time evolution as was shown in the previous works \\citep{Forslund70b,Mason72}. We also confirmed that the electrostatic ion-ion instability develops in front of the shock due to the counter-streams of the ions \\citep{Karimabadi91}. Then, we found that the electric field generated by this instability results in the strong fluctuation in the ion density and finally leads to destroy the shock itself. It was also found that the electromagnetic beam-Weibel instability develops much slower than the electrostatic instability but it becomes predominant in the later time. This suggests the possibility that the Weibel-mediated collisionless shock is formed according to the scenario shown in \\citet{Kato08} long after the electrostatic shock disappears. It was also observed that the secondary electrostatic shock forms in the reflected ions in front of the primary electrostatic shock." }, "1003/1003.3815_arXiv.txt": { "abstract": "Lynds dark cloud LDN1622 represents one of the best examples of anomalous dust emission, possibly originating from small spinning dust grains. We present Cosmic Background Imager (CBI) 31~GHz data of LDN1621, a diffuse dark cloud to the north of LDN1622 in a region known as Orion East. A broken ring with diameter $\\approx 20$~arcmin of diffuse emission is detected at 31~GHz, at $\\approx 20-30$~mJy~beam$^{-1}$ with an angular resolution of $\\approx 5$~arcmin. The ring-like structure is highly correlated with Far Infra-Red emission at $12-100~\\mu$m with correlation coefficients of $r \\approx 0.7-0.8$, significant at $\\sim10\\sigma$. Multi-frequency data are used to place constraints on other components of emission that could be contributing to the 31~GHz flux. An analysis of the GB6 survey maps at 4.85~GHz yields a $3\\sigma$ upper limit on free-free emission of 7.2~mJy~beam$^{-1}$ ($\\la 30$~per cent of the observed flux) at the CBI resolution. The bulk of the 31~GHz flux therefore appears to be mostly due to dust radiation. Aperture photometry, at an angular resolution of $13$~arcmin and with an aperture of diameter $30$~arcmin, allowed the use of IRAS maps and the {\\it WMAP} 5-year W-band map at 93.5~GHz. A single modified blackbody model was fitted to the data to estimate the contribution from thermal dust, which amounts to $\\sim 10$ per cent at 31~GHz. In this model, an excess of $1.52\\pm 0.66$~Jy ($2.3\\sigma$) is seen at 31~GHz. Future high frequency $\\sim 100-1000$~GHz data, such as those from the {\\it Planck} satellite, are required to accurately determine the thermal dust contribution at 31~GHz. Correlations with the IRAS $100~\\mu$m gave a coupling coefficient of $18.1\\pm4.4~\\mu$K~(MJy/sr)$^{-1}$, consistent with the values found for LDN1622. ", "introduction": "\\label{sec:introduction} Anomalous Microwave Emission (AME) is the name given to excess microwave emission, observed at frequencies in the range $\\sim10-60$~GHz, that is strongly correlated with far infrared (FIR) dust emission \\citep{Leitch97,Kogut96,Finkbeiner02,Banday03,deOliveira-Costa04,Finkbeiner04,Watson05,Casassus08,Davies06,Dickinson09a,Scaife09,Tibbs10}. This dust-correlated emission is known to be a significant source of contamination for Cosmic Microwave Background (CMB) data, that must be separated accurately from the CMB signal \\citep{Bennett03,Bonaldi07,Miville-Deschenes08,Gold09,Dickinson09b}. Although there is still some debate about the physical mechanism that is responsible for the emission, the most favoured explanation is in terms of small, rapidly spinning dust grains \\citep{Draine98a,Draine98b}. Assuming this is the case, microwave observations of spinning dust represent a new way of studying the properties of interstellar dust grains and its environment within the interstellar medium \\citep{Ali-Hamoud09,Dobler09,Ysard10}. Accurate observations, at frequencies in the range $5-100$~GHz, are now required to confirm spinning dust grains as the source of the anomalous emission and to study the spectrum to infer information about the grains and their environs. One of the best examples of spinning dust emission, comes from observations of the dark cloud LDN1622 \\citep{Finkbeiner02,Casassus06}. LDN1622 is a small ($\\approx 10$~arcmin) dark cloud at the low Galactic latitude edge of the giant molecular cloud Orion B, on the outer edge of Barnard's loop \\citep{Maddalena86}. The spectrum between 1~GHz and 3000~GHz is well-fitted by a superposition of optically thin free-free emission, thermal dust at a temperature of $T \\sim 15$~K and a spinning dust component with a peak at $\\sim 30$~GHz \\citep{Casassus06}; at 30~GHz, and on angular scales $\\la 20$~arcmin, the spectrum is dominated by spinning dust emission. Furthermore, correlations with IRAS infra-red maps indicate a better correlation with the emission from Very Small Grains (VSGs), as expected if the origin of the cm emission is spinning dust. About $25$~arcmin to the north-east of LDN1622 lies another dark cloud, LDN1621, which is more diffuse than LDN1622 and forms a broken ring-like structure. These two clouds together were termed as {\\it Orion East} by \\cite{Herbig72} and are thought to lie at a distance of $\\sim 400$~pc, although there is some debate about whether it lies much closer to us at $\\sim140$~pc \\citep{Wilson05,Kun08}. LDN1622 is a higher density dust cloud in which low mass star formation is beginning, as indicated by the presence of T-Tauri stars \\citep{Herbig72} but there are no bright OB stars in the vicinity; the UV radiation is predominately from the Orion OB1 association. By contrast, LDN1621 is a lower density, ring-like cloud, with no obvious pre-main sequence stars within \\citep{Lee05}. Here we present 31~GHz data of the area around LDN1621 and compare it with multi-frequency data, to estimate the contributions of free-free, thermal dust and anomalous dust components. This is compared with previous observations of the LDN1622 dark cloud to the south-west. Section~\\ref{sec:data_reduction} describes the observations and data reduction while Section~\\ref{sec:discussion} discusses multi-frequency maps of the LDN1621 region. Section~\\ref{sec:analysis} gives a quantitative analysis of the contributions from free-free and thermal dust emissions and the possible origin of the 31~GHz excess. Conclusions are given in Section~\\ref{sec:conclusions}. ", "conclusions": "\\label{sec:conclusions} LDN1621 is a region of diffuse emission $\\approx 25$~arcmin to the north of LDN1622. Observations with the CBI at 31~GHz show a broken ring of emission, that is strongly correlated with FIR emission at $12-100~\\mu$m, with Pearson correlation coefficients in the range $\\approx 0.6-0.8$. Optical and H$\\alpha$ data show absorption of a strong background of emission from warm ionized gas in the Eastern arm of Orion. This suggests that LDN1621 and LDN1622 are in the foreground of Orion (at a distance of $\\sim 500$~parsec), possibly as close as 120~pc \\citep{Wilson05}. No H$\\alpha$ emission, associated directly with LDN1621, is seen. This suggests that LDN1621 itself is not emitting significant free-free emission, although the effects of dust extinction do not allow a strong constraint to be placed. Low frequency radio data also do not show evidence of diffuse emission associated with LDN1621. The 31~GHz emission is at $\\approx 20-30$~mJy~beam$^{-1}$ while an analysis of the GB6 map at 4.85~GHz provides a strong ($3\\sigma$) upper limit of 7.2~mJy~beam$^{-1}$ at 31~GHz for free-free emission. The FIR-correlated emission at 31~GHz therefore appears to be mostly due to radiation associated with dust. IRAS data alone do not allow a reliable extrapolation of the Rayleigh-Jeans thermal dust tail to 31~GHz. {\\it WMAP} data at 93.5~GHz combined with IRAS data allowed the flux density to be estimated in an aperture of diameter 30~arcmin at an angular resolution of 13~arcmin. A single modified blackbody indicates that the thermal dust is $\\sim 10$~per cent of the total 31~GHz flux, corresponding to an excess of $1.52\\pm0.66$~Jy ($2.3\\sigma$). The dust-correlated emission has a coupling coefficient, relative to $100~\\mu$m, of $18.1\\pm4.4~\\mu$K~(MJy/sr)$^{-1}$, consistent with that observed from LDN1622. Orion East (consisting of both LDN1621 and LDN1622) appear to be part of the same system of dust clouds, emitting significant anomalous emission at frequencies $\\sim 30$~GHz. Spinning dust is an obvious candidate for the physical mechanism responsible for the bulk of the emission. High sensitivity data, covering a wide range of frequencies ($\\sim5-300$~GHz), is required to study such clouds in more detail. Data from the {\\it Planck} satellite will be particularly useful in constraining the Rayleigh-Jeans dust tail, which may be responsible for a significant fraction of the 31~GHz if the emissivity index flattens at longer wavelengths." }, "1003/1003.4268_arXiv.txt": { "abstract": "We investigate whether the subhalos of $\\Lambda$CDM galaxy halos have potentials consistent with the observed properties of Milky Way satellites, particularly those with high-quality photometric and kinematic data: Fornax, Leo I, Sculptor, Sextans, and Carina. We compare spherical models with isotropic velocity dispersion tensors to the observed, circularly averaged star counts, line-of-sight velocity dispersion profiles and line-of-sight velocity distributions. We identify subhalos within the six high-resolution dark matter halos of the Aquarius Project for which the spherically averaged potentials result in excellent fits to each of the five galaxies. In particular, our simple one-integral models reproduce the observations in the inner regions, proving that these data are fully consistent with $\\Lambda$CDM expectations and do not require cored dark matter distributions. For four of the five satellites the fits require moderately cusped {\\it stellar} density profiles. The star count data for Leo~I, however, do require a cored distribution of star counts. Current data suggest that these five satellites may be hosted by $\\Lambda$CDM subhalos with maximum circular velocities in the range 10 to 30~km~s$^{-1}$. ", "introduction": "The internal dynamics of the dwarf spheroidal (dSph) satellite galaxies of the Milky Way (MW) offer perhaps the best prospects for investigating the properties of the dark matter in the nearby universe. These galaxies are dominated by dark matter and the brightest of them are sufficiently close that high-precision line-of-sight velocities for large samples of their stars can be measured using high-resolution, multi-object spectroscopy ~\\citep{Walker:2008fc}. The current datasets represent an improvement upon the earliest such observations ~\\citep{Aaronson1983,Mateo1993} by factors of tens to thousands. Analysis of the kinematic data, in combination with improved photometric measurements, have not only confirmed earlier indications that the classical dSphs are dark matter dominated, but have also revealed the surprising property that their mean central densities are similar even though their luminosities span a very wide range of values~\\citep{Mateo:1998wg,Walker2007,Gilmore2007,Strigari:2008ib}. Current cosmogonic theory makes strong predictions for the internal structure of dark matter halos. N-body simulations of halo formation in hierarchical clustering cosmogonies have shown that halos develop strongly cusped density profiles which are almost independent of halo mass and cosmological parameters \\citep{nfw96,nfw97}. Subsequent simulations have confirmed this result \\citep[e.g.][]{Navarro:2008kc,Stadel:2009}, showing, in addition, that cuspy profiles are retained even after halos fall into larger ones and suffer extensive tidal stripping \\citep{Kazantzidis:2005su,springel2008}. Thus, to the extent that the dark matter distributions in the inner parts of halos have not been significantly disturbed by the galaxies forming within them, halo profiles offer a strong and direct test of the $\\Lambda$CDM cosmogony in a regime not probed by microwave background and large-scale structure data. An in-depth analysis of the central regions of Milky Way satellites is particularly important given that a number of recent studies of the structure of these galaxies have claimed that shallow central density profiles provide a better description of their dark matter halos than the cuspy profiles characteristic of $\\Lambda$CDM \\citep{Goerdt:2006rw,sanchez2006,Gilmore2007}. Models with dark matter core radii of $\\sim 100$~pc have been shown to provide good fits to the kinematic data sets of the classical dSphs \\citep{Angus:2009jh}. If confirmed, the shallow cores suggested by these studies might indicate a lower central phase-space density than expected if the dark matter is a cold collisionless particle \\citep{Tremaine:1979we,Hogan:2000bv}. Kinematic studies typically treat dSph galaxies as spherical, dynamically equilibrated systems. With these simple assumptions there is a strong degeneracy between the statistics of stellar orbits (i.e. whether velocity dispersions are isotropic, or are radially or tangentially biased) and the shape of the stellar and dark matter density profiles \\citep[e.g.][]{Evans:2008ik}. This ambiguity is reflected in the broad range of models used in recent attempts to constrain the dark matter density profiles of the dSphs \\citep{Strigari:2008ib,Lokas:2009cp,Walker:2009zp,Wolf:2009tu}. The intrinsic parameter degeneracies of the models cast doubt on the robustness of inferences favoring cored or cuspy central density profiles, even given the high-quality data that are now available \\citep{Strigari:2006ue,Walker:2009zp}. Breaking these degeneracies may only be possible by exploiting additional observational constraints, for example, through measurement of internal stellar proper motions \\citep{Wilkinson:2001ut,Strigari:2007vn}. Motivated by the simple question of whether the observed dSphs are kinematically consistent with the $\\Lambda$CDM theory of structure formation, we here use six high-resolution halo simulations performed as part of the Aquarius Project \\citep{springel2008} to search for subhalos whose properties would allow them, in principle, to host the well observed dSph satellites of the Milky Way. We restrict our attention to the five satellites with abundant, high-quality stellar kinematic data: Fornax, Sculptor, Leo I, Carina, and Sextans \\citep{Mateo:2007xh,Walker:2008fc}. Assuming isotropic velocity dispersions, we identify an Aquarius subhalo with a dark matter potential which results in a good simultaneous fit to each satellite's photometric and kinematic data. In addition to focusing specifically on $\\Lambda$CDM and making direct use of realistic potentials from the Aquarius simulations, our modeling differs from previous work in that we simultaneously fit both photometry and kinematics. We allow for mildly cusped {\\em stellar} profiles with $\\rho_\\star \\sim r^{-a}$ near the centre, where $a$ is in the range 0 to 1. Projections of such profiles are fitted to the photometric data, and with the results in hand we make predictions for the kinematic data, both traditional second moment (line-of-sight velocity dispersion) profiles and full line-of-sight velocity distributions. The latter comparison allows us to test whether simple, spherical, isotropic, $\\Lambda$CDM-based models are consistent with higher moments of the observed velocity distribution. For each of the five satellites we study, the best-fitting Aquarius subhalo provides an excellent statistical fit to the data. For four of the five this requires a mildly cusped stellar distribution similar to those found in brighter early-type galaxies. The only galaxy that requires a true core in the star distribution is Leo~I, but with such a profile its kinematics are still consistent with a $\\Lambda$CDM subhalo. We present circular velocity curves for the best-fitting subhalo hosts for each of the MW satellites. Not surprisingly, for a given satellite, the circular velocity curves of ``good'' subhalos are very similar at the radii that are well sampled by the stellar tracers. This is a consequence of our assumptions of spherical symmetry and isotropy. Finally, we determine the mass of the best-fitting subhalos, both at the time of accretion onto the host halo and at high redshift, and we show that these quantities have more scatter than the present-day central potentials or maximum circular velocities. ", "conclusions": "We have investigated whether the gravitational potentials of subhalos in N-body simulations of $\\Lambda$CDM halo formation are consistent with the high-quality photometric and kinematic data available for five of the brighter satellites of the Milky Way. We find that a direct mapping is, in fact, possible between each of these satellites and a subset of the dark matter subhalos in the six high resolution simulations of the Aquarius Project. Star count profiles with inner cusps scaling as $r^{-a}$ with $0\\leq a \\leq 1$ can provide good fits to the observed counts. Placed in the measured Einasto-like potentials of appropriately selected subhalos, they also fit the observed, nearly flat line-of-sight velocity dispersion profiles very well, even under the restrictive assumption of negligible velocity anisotropy. Such isotropic models fit the {\\it shapes} of the observed line-of-sight velocity distributions well, in addition to their second moments. We have measured the present-day maximum circular velocities of the ``best-fit'' subhalos for each of these five satellites. These range from 10 to 30~km~s$^{-1}$. Subhalos consistent with hosting the observed systems at $z=0$ have peak circular velocities (i.e. the largest maximum circular velocity they {\\it ever} had) ranging from 12 to 50~km~s$^{-1}$. The maximum past masses of the main progenitors of these subhalos range up to $\\sim 5\\times 10^9$ M$_\\odot$, while their masses at $z=7$ (the approximate lower bound on the redshift of reionization~\\citep{Dunkley:2008ie}) range up to $\\sim 10^9$ M$_\\odot$. At $z=0$ their Galactocentric distances range from 40 to 400~kpc. Our results indicate that current data on faint Milky Way satellites are consistent with these galaxies living in $\\Lambda$CDM halos. They do not, however, explain {\\it why} galaxies living in such subhalos should have the observed properties. Exploring this issue is an important task for future work \\citep[e.g.][]{Li:2009kv,Cooper:2009kx,Sawala:2009,Okamoto:2009rw,Busha2010}." }, "1003/1003.4097_arXiv.txt": { "abstract": "An abundance analysis has been conducted for a sample of nine post-AGB candidate stars; eight of them have not been explored before. We find four very promising objects like HD~105262, HD~53300 and CpD$-62^o5428$ among them. We find strong evidence of dust-gas separation through selective depletion of refractive elements in HD~105262. The same effect is also observed in HD~53300, CpD$-62^o5428$ and HD~114855 although abundance peculiarities are relatively smaller for the last two stars. We find strong enrichment of nitrogen for HD~725, HD~842, HD~1457, HD~9233 and HD~61227 but no further evidence to support their post-AGB nature. We have compared the observed [N/C] ratios of these stars with the predictions of evolutionary models which include the rotation induced mixing. ", "introduction": "The post-AGB stars (hereinafter PAGB) are the late stage of evolution of low and intermediate mass stars (1 to 8M$\\odot$) when they transit from AGB to Planetary Nebulae (PN). Since at the end of AGB evolution, most of the outer envelope is lost, circumstellar shells (detectable in infrared sub-millimeter to radio wavelengths) are commonly observed. However, for less massive progenitors, the longer transition times would result in the dissipation of circumstellar material for them. The atmospheric chemical compositions of PAGB stars and their circumstellar envelopes are those inherited from the local interstellar medium (ISM) but strongly modulated by the products of nucleosynthesis being dredge-up at different stages of evolution through successive mixing events. They enrich ISM with the products of nucleosynthesis through strong stellar winds. They are important contributors of C, N and s-process elements to the ISM. AGB evolution has been described in Herwig (2005) and post-AGB evolution in van Winckel (2003) and Garc\\1a-Lario(2006). Among intermediate mass stars, those in mass range 2-4M$_{\\odot}$ would experience Third Dredge Up (TDU) where the the product of helium burning as well as s-process elements are transported to the outer envelope and will become carbon stars (C/O $>$ 1). On the other hand, low mass stars (M$<$1.8M$_{\\odot}$) may not undergo sufficient thermal pulses and subsequent dredge up to reach C/O $>$ 1 stage. In higher mass stars the carbon would be quickly converted to nitrogen due to hot bottom burning (HBB) thereby preventing them from becoming carbon stars (Lattanzio et al. 1996, Groenewegen \\& Marigo 2004, Herwig 2005). However, the metallicity also has strong influence on the mass limits which determines the chemical dichotomy. The minimum mass needed to activate the HBB, number of thermal pulses needed to produce carbon stars and the efficiency of the dredge up are strongly affected by metallicity. This effect can be seen through the higher proportion of C-rich PN found in the metal-poor systems like Magellanic Clouds. Although the basic scheme of the post-AGB evolution as presented by Iben \\& Renzini (1983) is generally accepted, these AGB models and the calculations of yields from AGB were affected by the number of quantities e.g mass-loss, mixing length being kept as free parameters. Further development by Groenewegen \\& de Jong (1993), Boothroyd \\& Sackmann (1999) made better approximations for these parameters thereby getting better agreement with observations. But still these models are called synthetic models since they use analytical expression for thermal pulse phase. More complete models by Karakas et al. (2002), Herwig (2004) follow all the pulses in detail. A careful testing of these model prediction is warranted since the yield they produce still depends upon the adopted treatment of convection and mass-loss. \\subsection{Post-AGB detections} Observationally, a range of objects with diverse characteristics are found under this class, hence different strategies to identify them. The IRAS two color diagrams have been very useful in detecting these objects (Kwok et al. 1987, van der Veen et al. 1988; 1989, Van de Steene \\& Pottasch 1993, Garc\\1a-Lario et al. 1997, Van de Steene et al. 2000 and more recently Su\\'arez et al. 2006). These authors have studied candidates with PN like colors with supplementary data in longer wavelengths to confirm their advanced evolutionary status. These objects were found to be very faint in optical wavelengths. The investigation of optically bright IRAS sources with IR fluxes pointing to the existence of dusty shells (see e.g. Hrivnak et al. 1989, Pottasch \\& Parthsarathy 1988, Oudmaijer et al. 1992, van Winckel 1997) has also led to the detection of many post-AGB stars. However, these stars occupy different parts of the IRAS two color diagram. The systematic studies of high galactic latitude supergiants e.g. by Luck, Lambert \\& Bond 1990 from the candidate list of Bidelman (1951) have also resulted in more detections. Although many display double-peaked energy distributions, objects like HR 6144 and BD$+39^o4926$ are exceptions. Among these high galactic latitude supergiants, a small subgroup called UU Her characterized by high radial velocities small amplitude pulsation and large IR excess also contains sizeable fraction of post-AGB stars. Hot post-AGB stars are found from the studies of B stars found in the Galactic Halo (McCausland et al. 1992, Conlon et al. 1993, Moehler \\& Heber 1998). UV bright objects in globular clusters also contain objects like Bernard 29 in M13 (Conlon, Dufton \\& Keenan 1994), No 1412 in M4 (Brown, Wallerstein \\& Oke 1990) which show chemical compositions similar to those of halo B stars. The variable stars like RV Tau and population II Cepheids contain a noticeable fraction of post-AGB stars (Giridhar et al. 1994, Maas et al. 2002, Giridhar et al. 2005, Maas, Giridhar \\& Lambert 2007). It is a consequence of post-AGB evolutionary track intersecting the high luminosity end of instability strip. The above mentioned detections were made on samples strongly biased towards candidates showing optically brighter counterparts, generally stars located at high galactic latitude and belonging to low mass populations. More recent color selected and flux limited samples have led to the detection of rapidly evolving heavily obscured post-AGB stars. These objects included in GLMP catalog (Garc\\1a-Lario 1992) do not show preference to F and G spectral type and high galactic latitude but have flatter distribution in spectral type and follow galactic distribution which corresponds to more massive population (Garc\\1a-Lario 2006). The center stars of these highly obscured post-AGB stars is usually of B type suggesting a fast post-AGB evolution. They possess circumstellar molecular shells which are detectable in CO or OH at sub-millimeter or radio wavelengths. Most of these objects are O-rich which is expected for stars developing the HBB at the AGB phase. \\subsection {The observed chemical compositions} The chemical composition studies indicate a diverse behaviour too. The classical post-AGB stars (displaying double peaked SED) with enhancement of carbon (carbon -rich) and s-process elements caused by Third Dredge-Up make only a relatively smaller fraction of the known post-AGB stars. The well-known examples are HD 56126, IRAS 065330-0213, HD 158616, HD 187785 (Klochkova 1995, Bakker \\& Lambert 1998, Reyniers 2002 and Reddy et al. 2002). Their IR spectrum contains 21 micron features. The other subgroup (O-rich) also show double peaked SED but have no signature of TDU. The typical examples are HD 161796, 89 Her, HD 133656, SAO 239853, HR 4912 (Luck, Lambert \\& Bond 1983, van Winckel 1997, Giridhar, Arellano Ferro \\& Parrao 1997). A subgroup with C/O nearly one but showing selective depletion of easily condensable elements like Fe, Cr, Ca and Sc has been identified. HR 4049, HD 46703, BD$+39^o4926$, HD 52961 (Lambert, Hinkle \\& Luck 1988, Bond \\& Luck 1987, Waelkens et al. 1991) are well known post-AGB stars showing this effect. The same effect is also seen in RV Tau stars; e.g. IW Car, AC Her, EP Lyr, AD Aql and many others (Giridhar et al. 1994, Giridhar, Lambert \\& Gonzalez 2000, Giridhar et al. 2005, Maas, van Winckel \\& Waelkens 2002, van Winckel et al. 1998, Gonzalez et al. 1997) and also in population II Cepheids like ST Pup, CC Lyr (Gonzalez \\& Wallerstein 1996, Maas et al. 2007). Scenarios based upon a single star with dust-gas separation occurring in the stellar wind and binary stars with dust-gas separation occurring into circumbinary disks have been discussed in several papers. Keplerian circumbinary disk has been observed for HD 44179 van Winckel et al. (1995), Waters et al. (1998) and more recently by Bujarrabal et al. (2005). More such detection of circumbinary disks by de Ruyter et al. (2006), van Winckel (2007) and interferometric studies by Derroo et al. (2006; 2007) have lent support to the existence of compact passive disc around many post-AGB binaries. A mixed chemistry has been observed in the infrared spectra of some evolved objects where features of both O-rich and C-rich dust species are present e.g. J-type carbon stars exhibiting silicate dust emission (Lloyd Evans 1990), PAGB HD 44179 showing O-rich circumbinary disk (van Winckel et al. 1995, Waters et al. 1998) and EP Lyr shows emission features of C-PAH emission as well as those of crystalline silicates Gielen et al. (2009). The analysis of 350 ISO spectra of post-AGB and PN by Garc\\1a-Lario \\& Perea Calder\\'on (2003) has made a strong impact on our understanding of PAGB - PN transition based upon the shape of the infrared spectrum and evolution of gas phase molecular bands and that of the solid state features detected in the 1 to 60 $\\mu$ spectral range. These authors have identified two main chemical evolutionary sequences (for C-rich and O-rich stars) through which the process of condensation and growth of the dust-grain produced in circumstellar envelopes till the star becomes PN can be followed. Given the importance of post-AGB objects in understanding the late stages of evolution of low and intermediate mass stars and their contribution to enrichment of the interstellar medium (ISM), it is hardly surprising that the detection of new post-AGB stars still continue to be one of the most important scientific objectives of many surveys. The recent Torun catalog of post-AGB objects (Szczerba et al. 2007) contains 326 likely post-AGB stars and 107 possible candidates. We continue our exploration of optically bright post-AGB candidates, based upon the criteria of their PN like colors in two color diagrams of van der Veen \\& Habing (1988), high-galactic latitude and also Str\\\"omgren $c_{1}$ index (as listed in Bidelman 1993) with modest facilities available. Our sustained effort in the past has been rewarded with the detection of many interesting objects such as HD 725, HD 27381, HD 137569, HD 172481, HD 21853 and HD 331319 (Arellano Ferro, Giridhar \\& Mathias 2001 and Giridhar \\& Arellano Ferro 2005). This paper is the third in a series devoted to the search of more post-AGB stars among high-galactic latitude A-G supergiants (see Table \\ref{table1}). In section 2 we present description of the observational material used and reduction technique. The abundance analysis approach and error analysis is presented in section 3. The results of individual stars are presented in section 4. In section 5 we discuss large nitrogen enhancements observed in some sample stars. Section 6 contains the summary and conclusions. \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{Basic data for sample stars.} \\label{table1} \\begin{tabular}{llclrrllll} \\hline \\multicolumn{1}{l}{Star}& \\multicolumn{1}{l}{SpT.}& \\multicolumn{1}{c}{V}& \\multicolumn{1}{c}{l}& \\multicolumn{1}{c}{b}& \\multicolumn{1}{c}{IRAS}& \\multicolumn{1}{l}{12 $\\mu$}& \\multicolumn{1}{l}{25 $\\mu$}& \\multicolumn{1}{l}{60 $\\mu$}& \\multicolumn{1}{l}{100 $\\mu$}\\\\ \\multicolumn{1}{l}{}& \\multicolumn{1}{l}{}& \\multicolumn{1}{l}{(mag.)}& \\multicolumn{1}{c}{($^{o}$)}& \\multicolumn{1}{c}{($^{o}$)}& \\multicolumn{1}{l}{}& \\multicolumn{1}{l}{(Jy)}& \\multicolumn{1}{l}{(Jy)}& \\multicolumn{1}{l}{(Jy)}& \\multicolumn{1}{l}{(Jy)}\\\\ \\hline HD 725 & F5Ib-II & 7.08 & 117.56 & $-5.1$ & 00091+5659 & 0.36 & 0.25L & 0.40L & 14.43L\\\\ HD 842 & A9Iab & 7.96 & 117.52 & $-6.6$ & & & & & \\\\ HD 1457 & F0Iab & 7.85 & 118.92 & $-2.2$ & & & & & \\\\ HD 9233 & A4Iab & 8.10 & 128.15 & $-3.3$ & 01289+5853 & 0.31 & 0.37L & 0.40L & 9.27L \\\\ HD 53300 & A0Ib & 7.00 & 219.12 & $+0.4$ & 07018-0513 & 0.75 & 0.31 & 0.41: & 2.80: \\\\ HD 61227 & F0Ib & 6.38 & 239.15 & $-1.2$ & 07351-2339 & 0.62 & 0.25L & 0.40L & 6.36L \\\\ HD 105262 & B9 Ib & 7.09 & 264.5 &+72.4& &&&&\\\\ HD 114855 & F5Ia/Iab& 8.39 & 306.24 & $+8.0$ & 13110-5425 & 0.31L & 2.11 & 7.60 & 5.33 \\\\ CpD $-62^o5428$ & A7Iab & 9.94 & 326.20 & $-11.1$ & 16399-6247 & 0.25L & 1.28 & 1.94 & 2.17L \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*} \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{Derived physical parameters and radial velocities for program stars.} \\label {table2} \\begin{tabular}{lccccrcl} \\hline \\multicolumn{1}{l}{Star}& \\multicolumn{1}{c}{$T_{\\rm eff}$}& \\multicolumn{1}{c}{log~g}& \\multicolumn{1}{c}{$\\xi$}& \\multicolumn{1}{l}{[Fe/H]}& \\multicolumn{1}{c}{V$_{r}$(Hel)}& \\multicolumn{1}{l}{Observatory}& \\multicolumn{1}{l}{Date Obs.}\\\\ \\multicolumn{1}{l}{}& \\multicolumn{1}{c}{(K)}& \\multicolumn{1}{c}{}& \\multicolumn{1}{c}{(km s$^{-1}$)}& \\multicolumn{1}{l}{}& \\multicolumn{1}{c}{(km s$^{-1}$)}& \\multicolumn{1}{l}{}& \\multicolumn{1}{l}{}\\\\ \\hline HD 725 & 7000 & 1.0 & 4.65& $-0.20$ &$-56.9$ & OHP &1999 Jul 6 \\\\ & & & & &$-58.9$ & VBO &2006 Dec 8 \\\\ HD 842 & 7000 & 1.0 & 2.3 & $-0.25$ &$-31.2$ & OHP &2000 Oct 7 \\\\ & 7000 & 1.0 & 3.1 & &$-30.4$ & VBO &2008 Sep 22 \\\\ HD 1457 & 7300 & 0.75& 3.4 & $-0.20$ &$-38.5$ & OHP &2000 Oct 10 \\\\ & & & & &$-38.2$ & VBO &2008 Oct 9 \\\\ HD 9233 & 7750 & 1.0 & 4.2 & $-0.21$ &$-29.9$ & OHP &1999 Jul 11 \\\\ & & & & &$-31.9$ & VBO &2008 Sep 22 \\\\ HD 53300 & 7500 & 0.5 & 2.4 & $-0.62$ &$+58.4$ & McD &2007 Nov 2 \\\\ & & & & &$+58.3$ & VBO &2006 Feb 18 \\\\ HD 61227 & 7000 & 1.0 & 4.0 & $-0.38$ &$+18.5$ & OHP &2000 Oct 7 \\\\ & & & & &$+18.3$ & VBO &2005 Jan 22 \\\\ & & & & &$+17.9$ & VBO &2005 Mar 28 \\\\ & & & & &$+18.4$ & VBO &2006 Feb 10 \\\\ & 7250 & 1.0 & 3.2 & &$+18.8$ & VBO &2006 Feb 12 \\\\ & & & & &$+18.2$ & VBO &2006 Feb 14 \\\\ & 7000 & 1.0 & 4.9 & &$+18.0$ & VBO&2006 Feb 18 \\\\ HD 105262 & 8500 & 1.5 & 2.8 & $-1.87$ &$+18.5$ & McD&2009 May 19 \\\\ HD 114855 & 6000 & 0.5 & 4.7 & $-0.11$ &$+73.0$ & LCO &2008 Feb 14 \\\\ & & & & &$ -3.3$ & VBO &2006 Feb 16 \\\\ & & & & &$ -2.6$ & VBO&2006 Feb 17 \\\\ & & & & &$-13.6$ & VBO &2009 Jan 2 \\\\ & & & & &$ +2.5$ & VBO &2009 Jul 11 \\\\ CpD $-62^o5428$ &7250 & 0.5 & 4.6& $-0.45$&$-29.9$& LCO &2008 Feb 14 \\\\ \\hline \\end{tabular} \\end{minipage} \\flushleft{Sources of spectra: Haute Provence Observatory (OHP), Vainu Bappu Observatory (VBO), Las Campanas Observatory (LCO), McDonald Observatory (McD).} \\end{table*} \\section {Observations and Data reductions} The spectra for this work were obtained largely using the ELODIE spectrograph at 1.93m telescope of the Haute-Provence Observatory (OHP) giving 42,000 resolution (Baranne et al. 1996) and the echelle spectrometer at 2.3m telescope of the Vainu Bappu Observatory (VBO), Kavalur giving 28,000 resolution in the slitless mode (Kameswara Rao et al. 2005). A few, but very important spectra were also obtained with MIKE spectrograph on the 6.5m Magellan telescope at the Las Campanas Observatory giving about 24,000 resolution and a spectral coverage from 3350 to 9400 \\AA, and the 2D Coud\\'e echelle spectrograph (Tull et al. 1995) on the 2.7m telescope at the McDonald Observatory giving 60,000 resolution. The spectroscopic reductions were carried out using the tasks contained in IRAF software of NOAO. Our program stars were generally in the temperature range 7000K to 7500K hence the spectra were not crowded enabling us to measure the line strengths with an accuracy of 8 to 10\\% at the resolution employed. ", "conclusions": "" }, "1003/1003.3448_arXiv.txt": { "abstract": "Local Group dwarf spheroidal satellite galaxies are the faintest extragalactic stellar systems known. We examine recent data for these objects in the plane of the Baryonic Tully-Fisher Relation (BTFR). While some dwarf spheroidals adhere to the BTFR, others deviate substantially. We examine the residuals from the BTFR and find that they are not random. The residuals correlate with luminosity, size, metallicity, ellipticity, and susceptibility of the dwarfs to tidal disruption in the sense that fainter, more elliptical, and tidally more susceptible dwarfs deviate farther from the BTFR. These correlations disfavor stochastic processes and suggest a role for tidal effects. We identify a test to distinguish between \\LCDM\\ and MOND based on the orbits of the dwarf satellites of the Milky Way and how stars are lost from them. ", "introduction": "Recent years have seen enormous progress in the discovery and measurement of the tiny dwarf satellite galaxies of the Local Group. These include the long known, ``classical'' dwarf spheroidal galaxies \\citep{mateo} as well as the more recently discovered ``ultrafaint'' dwarfs and the satellites of M31 \\citep{willman,zucker,grill06,maj07,belokurov,grill09,martin}. In addition to identifying these systems, kinematic data from measuring the velocities of individual stars has become available for many systems. These now consist of thousands of individual stars for the classical dwarfs \\citep{walker07}, with rapidly improving data for the other types of systems \\citep{SG,jason}. The Local Group dwarfs appear to be the most dark matter dominated objects in the universe \\citep{mateo,wilk02,gerry,SG,koch,strigari,walker,boom}, consistent with the trend of dark matter domination increasing with decreasing surface brightness \\citep{dBMH96,MdB98a}. As such, they provide a unique probe of structure formation at the smallest accessible scales. They presumably reside in the sub-halos thought to inhabit the large dark matter halos of galaxies like the Milky Way (MW) and M31. If so, physical processes specific to their environment, like tidal disruption and ram pressure stripping, may play a role in the evolution of the luminous content of the dwarfs as they orbit the primary structure. It is therefore interesting to investigate whether, and the extent to which, the dwarfs obey the scaling relations established for brighter galaxies. In this paper we investigate how the dwarfs behave in the Baryonic Tully-Fisher plane. Rotating disk galaxies define a tight relation between baryonic mass and rotation velocity \\citep{btforig,verhTF}. This has recently been extended \\citep{stark,trach} to low rotation velocities ($\\sim 20\\;\\mathrm{km}\\,\\mathrm{s}^{-1}$) comparable to the Local Group dwarf spheroidals. An obvious question is whether the Local Group dwarfs that are satellites of the MW and M31 continue this relation. ", "conclusions": "We have examined the adherence of the Local Group dwarf satellites of the Milky Way and M31 to the Baryonic Tully-Fisher Relation. We find that most of the brighter dwarfs are largely consistent with the extrapolation of the BTFR fitted to isolated, late type, gas rich, rotating disk galaxies. The fainter dwarfs, especially the ultrafaint dwarfs, are not. More importantly, we find that residuals from the BTFR are not random, correlating well with luminosity and ellipticity. The amount of deviation from the BTFR also correlates with metallicity, size, and the susceptibility of the dwarfs to tidal perturbation. We have considered a number of possible interpretations for the observed behavior, as we summarize below. \\paragraph{Insufficient Kinematic Accuracy:} Heroic efforts have been made to find new dwarfs and to measure their velocity dispersions. This is a challenging endeavor. While the deviations of some dwarfs from the BTFR are formally significant, that significance is not overwhelming (typically 2 to $4 \\sigma$). Moreover, the analysis assumes that the dwarfs are spherical and in stable equilibrium. The assumption of sphericity at least is violated for the most deviant dwarfs. We therefore consider one possibility to be that there are no genuine deviations from the BTFR. This hypothesis predicts that as the data improve, so too will agreement with the BTFR. \\paragraph{Gas Removal:} Dwarfs that deviate from the BTFR do so in the sense that they seem to be lacking luminosity for their velocity dispersion. This may be explained if baryons are removed before they form stars. Several possible mechanisms to accomplish this removal include the suppression of star formation by cosmic reionization, ionization from Pop.~III stars, removal of cold gas by supernova feedback, and ram pressure stripping. Cosmic reionization is often invoked in the context of the dwarf spheroidals, and is an attractive solution if only these objects are considered. If we simultaneously consider slightly larger gas dominated disk galaxies, it becomes clear that reionization is not in itself an adequate explanation for the observed trends in the data. Some other mechanism must be acting to suppress the accumulation of cold baryons in a manner that becomes more severe with decreasing $V_c$. Whatever this mechanism is, it presumably affects the dwarf spheroidals as well. Cosmic reionization may be an additional factor acting only at scales $V_c < 20\\;\\mathrm{km}\\,\\mathrm{s}^{-1}$. Feedback from supernovae is a candidate mechanism for affecting star formation across all halo masses. This provides a qualitatively appealing explanation for the trend in the detected baryon fraction with halo mass. Supernovae provide the kinetic energy to drive gas beyond escape velocity, but the escape velocity increases with increasing halo mass so progressively more baryons are retained. This mechanism may provide a natural explanation for the residual correlation with luminosity and metallicity. However, quantitative tests remain wanting, as does an explanation for the correlations of the residuals with ellipticity and tidal susceptibility. In the ram pressure stripping hypothesis, galaxies deviate from the BTFR when they fall into the halo of the current host galaxy and their cold gas is ablated by the ram pressure of the hot gas in the halo. This hypothesis requires that a sufficient amount of hot gas be present in the halos of the host galaxies, which is not obviously the case. It predicts that the ages of stars in the dwarfs is related to the time of infall, with the dwarfs that fall in first losing the most gas, deviating by the largest amount from the BTFR, and having the oldest stars. Strictly speaking, the star formation history of the pre-infall dwarf is not constrained, so there should be some scatter in these predictions. However, no star formation can occur after infall and gas stripping, so the ages of the youngest stars should correspond well to the time of infall with a sharp truncation in star formation after that time. This would provide a natural explanation for why some dwarfs appear to have had relatively recent star forming events but now contain no cold gas. One would also expect the metallicities of the stars to reflect the star formation history. The first dwarfs to fall in would have had the least time for enrichment and have the lowest [Fe/H]. This predicts that [$\\alpha$/Fe] should also correlate with the amplitude of deviation $F_b$, to the extent to which we expect the objects with the briefest enrichment time to have the highest [$\\alpha$/Fe]. The gas removal hypotheses provide a potential explanation for the correlation of BTFR residuals with luminosity and metallicity. However, they provide no obvious explanation for the correlation with ellipticity and tidal susceptibility. This occurs more naturally in the following two scenarios. The removal of gas and the concomitant truncation of star formation and its consequences for metal enrichment may also occur as a result of tidal stripping in the following hypotheses. \\paragraph{Stellar Stripping:} The correlation of the residuals with ellipticity in addition to luminosity suggests a role for tidal effects. In this hypothesis, the dwarfs become progressively more distorted due to tidal disruption as they orbit the Milky Way. Stars are lost in the process, reducing the luminosity of the dwarfs. At the present time, the computed tidal radii of the dwarfs greatly exceed their luminous extent. This leads us to infer that the orbits of the dwarfs must be highly eccentric in this scenario, with most of the stripping occurring during pericenter passage. This hypothesis predicts that the mass required to reconcile each dwarf with the BTFR may exist in a tidal stream. It further predicts that the age and metallicities of stars in the predicted streams should be consistent with those of the parent body. Examples exist where this may already be observed. \\paragraph{MOND:} Low surface brightness dwarf spheroidals provide a strong test of an alternative to dark matter, MOND. They should very nearly follow the BTFR, which is a consequence of the specific form of the modified force law in MOND. While some dwarfs do indeed adhere to the BTFR, others deviate substantially. The ultrafaint dwarfs of the Milky Way have MONDian mass-to-light ratios in the tens to hundreds. This is fatal for MOND \\textit{if} these dwarfs are in a stable equilibrium, the spherical approximation used in the analysis is adequate, and the kinematic data are to be trusted. Intriguingly, the sizes of the dwarfs relative to their MONDian tidal radii correlate strongly with the degree of deviation from the BTFR. Indeed, the discrepancy for MOND sets in precisely where the theory predicts that non-equilibrium effects become strong. It therefore appears that the unacceptably high mass-to-light ratios may be a result of the dwarfs being out of equilibrium. This should be testable, in the sense that the deviant dwarfs should show evidence of tidal disruption while the dwarfs that adhere to the BTFR should not. Notably, stripping of the deviant dwarfs should be ongoing and not restricted to pericenter passage as in the stellar stripping hypothesis. It is of course possible that some combination of these effects is at work. As a dwarf satellite approaches its host on its orbit, it is subject to both tidal forces and ram pressure effects. Perhaps gas is lost first due one or both of these effects, with stars being tidally liberated later. This makes it somewhat difficult to distinguish between the various hypotheses, but it should be possible." }, "1003/1003.5164_arXiv.txt": { "abstract": "We present a detailed study of the effect of internal bremsstrahlung photons in the context of the minimal supersymmetric standard models and their impact on $\\gamma$-ray dark matter annihilation searches. We find that although this effect has to be included for the correct evaluation of fluxes of high energy photons from neutralino annihilation, its contribution is relevant only in models and at energies where the lines contribution is dominant over the secondary photons. Therefore, we find that the most optimistic supersymmetric scenarios for dark matter detection do not change significantly when including the internal bremsstrahlung. As an example, we review the $\\gamma$-ray dark matter detection prospects of the Draco dwarf spheroidal galaxy for the MAGIC stereoscopic system and the CTA project. Though the flux of high energy photons is enhanced by an order of magnitude in some regions of the parameter space, the expected fluxes are still much below the sensitivity of the instruments. ", "introduction": " ", "conclusions": "" }, "1003/1003.2869_arXiv.txt": { "abstract": "Mechanisms for the generation of primordial non-Gaussian metric fluctuations in the context of multiple-field inflation are reviewed. As long as kinetic terms remain canonical, it appears that nonlinear couplings inducing non-gaussianities can be split into two types. The extension of the one-field results to multiple degrees of freedom leads to \\textsl{gravity} mediated couplings that are ubiquitous but generally modest. Multiple-field inflation offers however the possibility of generating \\textsl{non-gravity} mediated coupling in isocurvature directions that can eventually induce large non-Gaussianities in the metric fluctuations. The robustness of the predictions of such models is eventually examined in view of a case study derived from a high-energy physics construction. ", "introduction": "It is now clearly understood that standard single field inflation cannot produce significant non-Gaussianities (NG) during or immediately after the inflationary phase. The result obtained by Maldacena in Ref. \\cite{2003JHEP...05..013M} explicitly shows that standard single field inflation leads to no or very little primordial non-Gaussianities. And this result appears to be very robust, independent on the details of the model. This point is best illustrated by the expression of the bispectrum in the squeezed limit. Defining the time dependent curvature modes $\\zeta(t,\\vk)$ and taking advantage of the statistical isotropy of the universe, the power spectrum $\\mP_{\\zeta}$ of the field $\\zeta$ can be defined as \\begin{equation} \\langle\\zeta(t,\\vk_{1})\\zeta(t,\\vk_{2})\\rangle=(2\\pi)^3\\delta_{\\rm Dirac}(\\vk_{1}+\\vk_{2})\\,\\mP_{\\zeta}(k_{1},t) \\end{equation} and its bispectrum\\footnote{In this context, both the power spectrum and the bispectrum will eventually be time independent at super-Hubble scales.} $\\mB_{\\zeta}$ as \\begin{equation} \\langle\\zeta(t,\\vk_{1})\\zeta(t,\\vk_{2})\\zeta(t,\\vk_{3})\\rangle=(2\\pi)^3\\delta_{\\rm Dirac}(\\vk_{1}+\\vk_{2}+\\vk_{3})\\,\\mB_{\\zeta}(\\vk_{1},\\vk_{2},\\vk_{3},t). \\end{equation} In the squeezed limit, i.e. when $k_{1}\\ll k_{2}\\approx k_{3}$, the bispectrum scales like $(n_{s}-1) \\mP_{\\zeta}(k_{1}) \\mP_{\\zeta}(k_{2})$ where $n_{s}$ is the spectral index \\cite{2004JCAP...10..006C}. Not only are the nonlinear couplings naturally small -- say of order unity\\footnote{Note however that the amount of NGs determined by a dimensionless quantity such as $\\mB/\\mP^{3/2}$ is of the order of $\\mP_{*}(n_{s}-1)$ where $\\mP_{*}\\approx 10^{-5}$ is the amplitude of the metric fluctuations.} -- they are even suppressed by the slow-roll parameters (that ensures that $n_{s}$ is close to unity). There are then two possible strategies to escape the limits set by Maldacena's results. One can modify the kinetic term by introducing higher order terms in the action that are not due to the potential shape. An example is provided by the Dirac-Born-Infeld action \\cite{2004PhRvD..70j3505S}. Such models will succeed in producing large NGs if precisely the kinetic term is, at the time of horizon crossing, at a non standard running point. That does not change however the squeezed limit case but allows large NG couplings for more equilateral type configurations of modes. This has been put forward as a powerful way for discriminating models \\cite{2004JCAP...08..009B}. Another way of evading the constraints of standard single field inflation is to introduce multiple scalar degrees of freedom. It can actually be argued that this is a natural hypothesis since it is unlikely that only one fundamental degree of freedom will be light (e.g. compared to the Hubble energy scale) during the epoch of inflation. What is more hypothetical is whether those extra degrees of freedom can have observational consequences. By definition, degrees of freedom that do not participate in the metric fluctuation, at a given time, are called isocurvature modes. There is no reason why the isocurvature modes should remain so all along the history of the universe and various mechanisms have been put forward that can lead to a transfer of modes, from isocurvature to adiabatic modes. For instance the curvaton model is based on the survival of (massive) isocurvature modes until late after the end of inflation that can alter the subsequent expansion history of the universe \\cite{2002PhLB..524....5L}. This is a particular case of modulated inflation \\cite{2003astro.ph..3614K,PhysRevD.69.023505,2004PhRvD..70h3004B}. Other mechanisms assume that isocurvature modes can change the end-point of inflation or alter the (p)-reheating effects (see \\cite{2009PhRvL.103g1301B} and contribution by A. Frolov, this volume). Such mechanisms can also happen in the context of hybrid inflation. It does not mean yet that it induces non-Gaussian metric fluctuations. That would happen only if there are nonlinearities in the isocurvature-curvature transfer or if isocurvature modes are intrinsically NG at the time of transfer. This latter situation is in particular advocated in Refs. \\cite{1990PhRvD..42.3936S,2002PhRvD..65j3505B,2002PhRvD..66j3506B,2003PhRvD..67l1301B} where isocurvature modes are shown to be able to develop large NG after horizon crossing. The aim of this paper is to show how different models that have been put forward in the literature differ in their mechanisms for producing NGs and how they differ in the amplitude and/or shapes of NGs they produce. This will be described in section 2 where it is argued that one can distinguished between gravity and non-gravity mediated contributions. The mere construction of working mechanisms cannot however be fully satisfactory. It is now clear that models can lead to a variety of observational signatures. However, whether there exist natural realizations for those models from high-energy physics point of view is largely open. That will be tentatively addressed in section 3. ", "conclusions": "The exploration of the various types of coupling terms that appear in the action leads to distinguish gravity from non-gravity mediated couplings. In one-field inflation, because the field fluctuations and the metric fluctuation are locked together, only the former can be found. In multiple-field inflation however this is not necessarily so and it opens the possibility of having a richer phenomenology. Whereas the gravity mediated couplings are ubiquitous but induce only modest effects, the non-gravity mediated couplings can be very efficient, although nothing ensures that they are generically at play. Giving the freedom one has in building potentials it is however certainly possible to design models exhibiting any kind of scale and geometrical dependences in the bispectrum (this has been tentatively explored in Ref. \\cite{2009arXiv0911.2780B}), but actual constructions motivated by high-energy physics point to hybrid inflation type models where couplings are eventually of local type. The extended $D$-term susy model presented here is an example of such a construction. A few lessons can be drawn from its analysis, \\begin{itemize} \\item in such a susy framework masses and coupling constants are naturally protected. This is at the expense of the doubling of the number of scalar degrees of freedom; \\item this extension of $D$-term inflation leads to hybrid models where the end line, the critical line where the inflationary period terminates, is linear in all field directions. This is at variance with the construction proposed in \\cite{2008PThPh.120..159S} where NGs originate from a nonlinear critical line; \\item in such a model there is no need for specific fine tunings, neither in the initial conditions nor in the parameter space. It is also to be noted that such models induce a dumping of the rare event tails - on both sides. \\item the transfer of modes, from isocurvature to adiabatic direction, takes place at the time the inflation stops. It leads to local types for bispectrum, and trispectrum, shapes. \\item calculations were carried here at leading order in perturbation theory. The conditions for such calculations to be valid during the sub-Hubble evolution is that the coupling constant $\\nu^2$ is less than unity. Once the evolution is super-Hubble however the amplitude of the fluctuation couplings is driven by $\\nu^2 (N-N_{*})$ which ought then to be small. There is therefore a regime where perturbation theory can be used during sub-Hubble evolution but not during the whole super-Hubble evolution. That could lead to a new phenomenology; e.g. to nonlocal effects in the shape of the high order spectra. \\end{itemize} The model at hand is therefore rather sound with not so much freedom in its predictions. But this is by no means exhaustive. Other scenarios are certainly possible." }, "1003/1003.0577_arXiv.txt": { "abstract": "The problems of using the spectral index of radio galaxies in various tests, in particular, in selecting distant radio sources are considered. The history of the question of choosing a criterion of searching for distant radio galaxies based on the spectral index is presented. For a new catalog of 2442 radio galaxies constructed from NED, SDSS, and CATS data, an analytical form of the spectral index.redshift relation has been determined for the first time. The spectral index.angular size and spectral index.flux density diagrams have also been constructed. Peculiarities of the distribution of sources on these diagrams are discussed. \\mbox{\\hspace{1cm}}\\\\ \\noindent PACS: 98.54.-h, 98.54.Gr, 98.62.Ve, 98.70.Dk, 98.80.Es Key words: radio galaxies: observations, radio continuum, spectral index. ", "introduction": "Radio galaxies are among the most powerful observed space objects, making it possible to use such radio sources to investigate the properties of the Universe at various cosmological epochs. Therefore, constructing samples of radio galaxies from various red- shift ranges is one of the most important tasks in observational cosmology. The sharp fall in the number of observed radio sources with increasing redshift forces us to seek for ways of rapidly selecting distant objects. One of the ways is selection by the spectral index $\\alpha< -1.0$ ($S\\sim\\nu^{\\alpha}$, where S is the flux density and $\\nu$ is the frequency). It is based on the fact that the farther the object, the more probable that it will have a steep spectrum. This is one of the first discovered and strongest criteria in searching for distant galaxies based on radio-astronomical data. It was established independently in several works devoted to the identification of radio sources and the analysis of radio- spectra statistics. Among the first papers, we will note the paper by Whitfield (1957), who pointed out a correlation of the spectral index for a radio source with its distance, and the papers by Dagkesamanskii (1969, 1970), who found that there were no distant objects with spectral indices $\\alpha>-0.7$ for the 3C sources. Subsequently, Tielens et al. (1979) determined that the fraction of optically identified sources with ultrasteep spectra (spectral indices $\\alpha^{5000}_{178}<-1$) decreased with decreasing $\\alpha$. Following this work, Blumenthal and Miley (1979) showed the spectral index for 3C and 4C radio sources of various populations to depend on the properties of the objects: their apparent magnitude, redshift, radio luminosity, and angular size. The detected correlation suggested that the steep-spectrum sources were, on average, further away and more luminous than the sources with less steep spectra of the same populations (radio galaxies and quasars). Laing and Peacock (1980) investigated the relation between radio spectrum and radio luminosity for samples of extragalactic sources at 178 and 2700 MHz. The spectra were measured for the extended regions of the radio sources, which were classified by morphological types. At low frequencies, the degree of spectral curvature was found to correlate with the luminosity for sources with hot spots. At high frequencies, the correlation between spectral index and luminosity was confirmed. Laing and Peacock (1980) also confirmed the existence of a spectral index.redshift relation at 178 and 2700 MHz for FR II sources (Fanaroff and Riley 1974). At present, many groups and authors use such a selection of candidates for distant radio galaxies (see, e.g., Chambers et al. 1988; Wieringa and Katgert 1991; Soboleva and Temirova 1991; R$\\ddot{o}$ttgering et al. 1997; de Breuck et al. 2000; Pedani 2003; Verkhodanov et al. 2003; Gopal-Krishna et al. 2005; Klamer et al. 2006; Bornancini et al. 2006; Kopylov et al. 2006). Although the criterion based on the spectral index is very efficient, its explanation is not yet completely clear (De Young 2002). Three main, widely used ideas that explain this dependence can be highlighted: -- if the spectral steepness for many radio galaxies increases with frequency, then the spectra of distant objects must be steeper than those of near ones because of the factor (1+z); -- the losses due to the Compton scattering of cosmic microwave background (CMB) photons by relativistic electrons grow, since the CMB radiation density increases as $(1+z)^4$; these growing losses lead to an aging of the population of electrons, causing a cutoff or an increase in the steepness of the high-energy part of the spectrum where the losses are greatest; -- the selection effect: only bright sources are seen at high redshifts, and it can then be said that precisely these sources exhibit the greatest ``depletion'' in the high-energy part of the spectrum. Klamer et al. (2006) give yet another possible explanation for the $\\alpha-z$ correlation. They noted that the steep-spectrum sources are rather rare among the nearest FR I radio galaxies, while those with such spectra are located in regions with a high baryon density. It can then be assumed that if there is evolution of the environment of powerful radio galaxies reflected in the richness of the surrounding cluster, then, on average, the radio galaxies are more likely located in regions with a higher ambient density than in less dense regions. Hence follows the observed redshift dependence of the spectral steepness. This can play its role when the gas density as one goes into the past increases as $(1+z)^3$, while the injection of electrons with steeper spectra occurs naturally and is described as a function of the redshift in the first- order Fermi acceleration processes attributable to the decreasing expansion velocity of the hot spots in the denser and hotter intergalactic medium (Athreya and Kapahi 1998). Other authors (Kharb et al. 2008), who studied a sample of 13 ``powerful classical double'' FR II radio galaxies, also analyzed this $\\alpha-z$ correlation and pointed out the existence of a redshift dependence of the spectral index for both the hot spots and the core. From the presence of a correlation between the spectral indices for the hot spots and a core with a flatter spectrum, Kharb et al. (2007) drew the conclusion in favor of choosing the explanation of the effect by the electron aging model. Although as yet there is no general, common consensus on clarifying the nature of the $\\alpha-z$ correlation, the very fact that this empirical dependence exists is important for studies. Almost all of the distant radio galaxies found has passed through the stage of such a selection. An example of an object selected in this way is the radio galaxy that is the red- shift record-holder with z=5.19 and spectral index $\\alpha=-1.63$ (van Breugel et al. 1999). Another example is the object RC 0311+0507 investigated in the ``Big Trio'' Program (Kopylov et al. 2006). Its redshift is z=4.514, it has the record radio luminosity for radio galaxies with $z>4$, and its spectral index is $\\alpha=-1.33$. We prepared a sample of distant ($z>0.3$)radio galaxies (Khabibullina and Verkhodanov 2009a, 2009b, 2009c) using the NED\\footnote{\\tt http://nedwww.ipac.caltech.edu} CATS\\footnote{\\tt http://cats.sao.ru} (Verkhodanov et al. 2005), and SDSS\\footnote{\\tt http://www.sdss.org} (Schneider et al. 2007) databases for the subsequent application in various statistical and cosmological tests (Verkhodanov and Parijskij 2003, 2009; Miley and De Breuck 2008), in which a large number of objects of the same nature is required to carry out a study. To compile the primary list, we used the NED database, from which we selected objects with the following parameters: the redshift ($z>0.3$) and morphological properties.radio galaxies. The initial list contained 3364 object. This sample of objects is contaminated by objects with incomplete information or by objects with different properties. Therefore, the next stage was to clean the initial sample of superfluous sources. For this purpose, we selected the objects that were removed from the primary list (Khabibullina and Verkhodanov 2009a): (1) with photometrically determined redshifts and (2) with quasar properties based on available published data. The final catalog contains 2442 sources with spectroscopic redshifts, photometric magnitudes, radio flux densities, sizes of the radio sources, and radio spectral indices calculated from the results of the cross-identification of the list of selected radio galaxies with the radio catalogs stored in CATS in the frequency range from 30 GHz to 325 MHz. By default, the flux densities are given on the Baars scale. The data collected in the catalog can be used for cosmological luminosity-redshift, angular size-redshift, and age-redshift tests. In addition, the statistical diagrams for the parameters of radio galaxies and their evolutionary properties can be investigated using the cataloged data. The goal of this paper is to construct and analyze the spectral index-redshift, spectral index-angular size, and spectral index-flux density diagrams for a pure sample of radio galaxies, which can be used to estimate the redshifts of radio galaxies and to calculate their luminosity function. ", "conclusions": "We presented the results of our investigation of the $\\alpha -z$, $\\alpha - \\theta$, and $\\alpha -S$ diagrams for a sample of radio galaxies containing 2442 objects (Khabibullina and Verkhodanov 2009a, 2009b, 2009c) and constructed from the spectroscopic data collected in NED, SDSS, and archives of the Big Trio program (Parijskij et al. 1996, 1999). For the objects of our list, we found the existence of a regression $\\alpha(z)$ and established its analytical form: $\\alpha(z)=-0.73- 0.15z$. The estimate of the regression parameters is stable, since is was made using the median values of large subsamples of radio galaxies. This dependence can also be used to preselect objects at given z,to estimate the distances to radio galaxies in the first approximation, and to study the luminosity function. Nevertheless, when using the derived relation $\\alpha(z)$, we should take into account the fact that the detected regression was obtained from an incomplete sample. For the correlation $\\alpha(S_{1400})$, the regression parameters were found to evolve with increasing z. Since correlations (Disney et al. 2008) suggesting the existence of a multiparameter fundamental plane of radio galaxies are found for a number of their physical parameters (total mass, baryon fraction, luminosity, etc.), it may be concluded that the evolution of the dependence $\\alpha(S_{1400})$ is also indicative of the possible evolution of the fundamental plane. In turn, the existence of a correlation between the variations in parameters of the fundamental plane and spectral indices argues for the explanation of the dependence $\\alpha(z)$ by the electron aging model, which complements the conclusions by Kharb et al. (2007). \\noindent {\\small {\\bf Acknowledgments}. We wish to thank the referees for useful remarks that allowed the content of our paper to be improved. In our study, we used the NASA/IPAC Extragalactic Database which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. We also used the CATS database of radio-astronomical catalogs (Verkhodanov et al. 1997, 2005) and the radio-astronomical data processing system FADPS\\footnote{\\tt http://sed.sao.ru/$\\sim$vo/fadps\\_e.html} (Verkhodanov et al. 1993; Verkhodanov 1997b). This work was supported by ``Leading Scientific Schools of Russia'' grants (S.M. Khaikin.s school) and the Russian Foundation for Basic Research (project nos. 09-02-00298 and 09-02-92659-IND). One of us (O.V.V.) also thanks the Russian Foundation for Basic Research (project no. 07-02-01417) and the Foundation for Support of Russian Science (the ``Young Doctors of the Russian Academy of Sciences'' Program) for partial support. }" }, "1003/1003.5214_arXiv.txt": { "abstract": "Stellar population models of absorption line indices are an important tool for the analysis of stellar population spectra. They are most accurately modelled through empirical calibrations of absorption line indices with the stellar parameters effective temperature, metallicity, and surface gravity, the so-called fitting functions. Here we present new empirical fitting functions for the 25 optical Lick absorption line indices based on the new stellar library MILES. The major improvements with respect to the Lick/IDS library are the better sampling of stellar parameter space, a generally higher signal- to-noise, and a careful flux calibration. In fact we find that errors on individual index measurements in MILES are considerably smaller than in Lick/IDS. Instead we find the rms of the residuals between the final fitting functions and the data to be dominated by errors in the stellar parameters. We provide fitting functions for both Lick/IDS and MILES spectral resolutions, and compare our results with other fitting functions in the literature. A Fortran 90 code is available online in order to simplify the implementation in stellar population models. We further calculate the offsets in index measurements between the Lick/IDS system to a flux calibrated system. For this purpose we use the three libraries MILES, ELODIE, and STELIB. We find that offsets are negligible in some cases, most notably for the widely used indices H$\\beta$, Mg$b$, Fe5270, and Fe5335. In a number of cases, however, the difference between flux calibrated library and Lick/IDS is significant with the offsets depending on index strengths. Interestingly, there is no general agreement between the three libraries for a large number of indices, which hampers the derivation of a universal offset between the Lick/IDS and flux calibrated systems. ", "introduction": "\\label{intro} Stellar population models of absorption line indices are a key tool for the analysis of star cluster and galaxy absorption spectra. They are used to derive the fundamental stellar population properties such as age, metallicity and element abundance ratios. In particular, optical absorption line diagnostics in the spectra of evolved stellar populations have successfully been adopted in the past in studies on galaxy evolution \\citep*[e.g.][]{worthey92,davies93,vazdekis97,kuntschner98,worthey98,trager98,henry99, kuntschner00,trager00,thomas05} and globular cluster formation \\citep[e.g.][]{kpatig98,forbes01,kuntschner02,brodie05,puzia05}. The Lick/IDS system \\citep{burstein84,faber85,gorgas93,worthey94,worthey97,trager98} is the standard set of absorption line indices that has been used extensively during the last two decades for studying absorption features of stellar populations. This system consists of index definitions for 25 prominent absorption features between 4000 and $6500\\;$\\AA\\ present in the spectra of evolved stellar populations. For studies of galaxy and star cluster evolution, absorption lines need to be modelled for stellar populations \\citep*[e.g.][]{maraston98,maraston05,bc03,vazdekis99,leitherer99,worthey94,worthey97,rose94,trager00,TMB03,thomas04}. A convenient way goes through the use of empirical calibrations. This is motivated by the fact that theoretical model atmospheres are known to suffer from incomplete line lists and continuum uncertainties. \\citep*[e.g.][]{korn05,coelho07,merino05,lee09,walcher09}. Empirical calibrations on the other hand have the disadvantage to be hardwired to the chemical abundance pattern of the Milky Way, which can be overcome in a semi-empirical approach as in the models by \\citet{trager00}, \\citet{TMB03,thomas04} and \\citet{schiavon07} Empirical calibrations can be inserted in the models in two ways. In the first and most widely used approach, absorption line indices enter stellar population modelling through calibrations of the empirical relationship between the indices and the stellar atmospheric parameters T$_{\\rm eff}$, $\\log g$ and [Fe/H] as provided by stellar libraries. As these calibrations are usually obtained through polynomial fitting procedures they are commonly referred to as 'fitting functions'. The quality of the final stellar population model critically depends on the accuracy with which these relationships can be inferred from stellar libraries, i.e.\\ the coverage of stellar parameter space and the reliability of the index measurements. The computational procedure with which the fitting functions are determined is a further crucial step in producing accurate models. A number of studies in the literature are devoted to such empirical calibrations for various stellar libraries, either for the Lick indices, parts of the Lick indices or other prominent absorption features \\citep*{buzzoni92,buzzoni94,worthey94,borges95,gorgas99,cenarro02,schiavon07,maraston09}. Alternatively to the use of fitting functions, absorption line indices can be measured directly on the synthetic spectral energy distribution (SED) from stellar population models that are based on empirical stellar libraries. The benefit of this method is that the full SED can be compared pixel-by-pixel to observations \\citep[e.g.][]{panter07,tojeiro07}. The major strength of fitting functions, instead, lies in the fact that they allow for interpolation between well populated regions of stellar parameter space which increases the accuracy of the model in stellar parameter space that is only sparsely sampled by empirical stellar libraries. Moreover, each absorption index or spectral feature is represented by an individual fitting function, which is optimised to best reproduce its behaviour in stellar parameter space. Fitting functions are also easier to implement in a stellar population synthesis code, and models based on fitting functions are better comparable. The widely used fitting functions of \\citet{worthey94} and \\citet{worthey97} are based on the Lick/IDS stellar library \\citep{burstein84,faber85}. They are adopted in most stellar population models \\citep{worthey94A,vazdekis96,trager00,TMB03,thomas04,thomas05,annibali07} in the literature. Other fitting functions based on the same stellar library exist \\citep{buzzoni92,buzzoni94,borges95} and lead to overall consistent results in the final stellar population model \\citep{maraston03}. Major progress has been made with the advent of a new generation of stellar libraries \\citep{jones99,prugniel01,leBorgne03,miles} that have led to considerable improvements regarding coverage of stellar parameter space, spectral resolution, signal-to-noise ratio, and flux calibration. In particular the latter is a critical step forward. As the Lick/IDS system is not flux calibrated, observations have to be re-calibrated onto the Lick/IDS system through comparison with Lick standard stars. This requirement hampers the analysis of data samples for which spectra of such calibration stars are either not available at sufficient quality or do not cover the appropriate rest-frame wavelength range. This problem is most imminent in high redshift observations and in galaxy redshift surveys such as the Sloan Digital Sky Survey \\citep{york00}. The new flux calibrated libraries allow the analysis of flux calibrated spectra at any redshift without spectroscopic standard stars. Flux-calibrated stellar libraries in the literature that are suitable for stellar population modelling include the \\emph{Jones} \\citep{jones99}, \\emph{ELODIE} \\citep{prugniel01}, \\emph{STELIB} \\citep{leBorgne03} and \\emph{MILES} \\citep{miles} libraries. The \\emph{MILES} library is particularly well suited for stellar population modelling of absorption line indices owing to its favourable combination of spectral resolution, wavelength range, stellar parameter coverage, and quality of flux calibration. In this paper we present new Lick index fitting functions based on the \\emph{MILES} stellar library. To take advantage of the full spectral resolution of the \\emph{MILES} library we have produced fitting functions for both the lower Lick/IDS resolution ($8-11\\;$\\AA\\ FWHM) and the higher resolution of the \\emph{MILES} library (2.3 \\AA\\ FWHM). A new version of the TMB stellar population model of absorption line indices based on these new fitting functions will be presented in a subsequent paper. This paper is organized as follows. In Section~\\ref{stlib} we present the Lick indices measured on the \\emph{MILES} library and a quality evaluation of the index measurements. We discuss offsets between the flux calibrated MILES and the Lick/IDS systems. The empirical fitting method is presented in Section~\\ref{fitf} along with the resulting fitting functions. In Section~\\ref{comps} we compare the fitting functions of this work with fitting functions from the literature. We summarise in Section~\\ref{concs}. \\begin{table*} \\centering \\caption{Typical Lick index errors and offsets to the Lick/IDS library. M-$\\sigma$ and L-$\\sigma$ corresponds to index errors at the resolution of the \\emph{MILES} and Lick/IDS libraries, respectively. T98-$\\sigma$ are the index errors presented in \\citet{trager98} for the Lick/IDS library. $I_{lib}$ are indices measured on the libraries (\\emph{MILES}, \\emph{ELODIE} and \\emph{STELIB}) for which offsets to the Lick/IDS library are presented. $I_{L}$ are indices measured on the Lick/IDS library.} \\label{offtable} \\begin{tabular}{rlllllllr@{0}llr@{0}llr@{0}l} \\hline \\multicolumn{2}{|c|}{\\bf INDEX} & & \\multicolumn{3}{|c|}{\\bf Error} & & \\multicolumn{9}{|c|}{\\bf Offset $I_{lib}=a\\cdot I_{Lick}+b$}\\\\ & & & & & & & \\multicolumn{3}{|c|}{MILES} & \\multicolumn{3}{|c|}{ELODIE} & \\multicolumn{3}{|c|}{STELIB} \\\\ i & NAME & & \\multicolumn{1}{|c|}{M-$\\sigma$} & \\multicolumn{1}{|c|}{L-$\\sigma$} & \\multicolumn{1}{|c|}{T98-$\\sigma$} & & \\multicolumn{1}{|c|}{$a$} & \\multicolumn{2}{|c|}{$b$} & \\multicolumn{1}{|c|}{$a$} & \\multicolumn{2}{|c|}{$b$} & \\multicolumn{1}{|c|}{$a$} & \\multicolumn{2}{|c|}{$b$} \\\\ \\hline 1 & H$\\delta_{A}$ & & 0.164 & 0.125 & 0.64 & & 0.960 & -&.054 & 0.955 & &.721 & 0.940 & &.823 \\\\ 2 & H$\\delta_{F}$ & & 0.093 & 0.075 & 0.40 & & 0.965 & &.049 & 0.936 & &.397 & 0.956 & &.242 \\\\ 3 & CN$_{1}$ & & 0.0042 & 0.0038 & 0.018 & & 0.912 & &.008 & 0.897 & -&.012 & 0.986 & -&.010 \\\\ 4 & CN$_{2}$ & & 0.0050 & 0.0042 & 0.019 & & 0.907 & &.006 & 0.900 & -&.008 & 0.985 & -&.013 \\\\ 5 & Ca4227 & & 0.063 & 0.047 & 0.25 & & 0.904 & &.074 & 0.771 & &.163 & 0.918 & -&.057 \\\\ 6 & G4300 & & 0.112 & 0.093 & 0.33 & & 0.858 & &.625 & 0.870 & &.646 & 0.924 & &.565 \\\\ 7 & H$\\gamma_{A}$ & & 0.142 & 0.107 & 0.48 & & 0.976 & -&.148 & 0.967 & -&.057 & 1.022 & -&.735 \\\\ 8 & H$\\gamma_{F}$ & & 0.069 & 0.059 & 0.33 & & 0.963 & -&.038 & 0.962 & &.016 & 0.999 & -&.238 \\\\ 9 & Fe4383 & & 0.155 & 0.127 & 0.46 & & 0.932 & -&.220 & 0.929 & -&.184 & 0.915 & &.796 \\\\ 10 & Ca4455 & & 0.073 & 0.056 & 0.22 & & 0.747 & -&.067 & 0.785 & -&.105 & 0.891 & -&.228 \\\\ 11 & Fe4531 & & 0.122 & 0.096 & 0.37 & & 0.857 & &.290 & 0.838 & &.390 & 0.877 & -&.002 \\\\ 12 & C$_{2}$4668 & & 0.179 & 0.156 & 0.57 & & 0.903 & &.484 & 0.913 & &.295 & 0.992 & &.512 \\\\ 13 & H$\\beta$ & & 0.063 & 0.051 & 0.19 & & 0.981 & &.126 & 0.996 & &.015 & 1.004 & &.032 \\\\ 14 & Fe5015 & & 0.139 & 0.115 & 0.41 & & 0.902 & &.084 & 0.926 & &.178 & 0.989 & &.168 \\\\ 15 & Mg$_{1}$ & & 0.0017 & 0.0013 & 0.006 & & 0.911 & &.0004 & 0.923 & &.005 & 0.903 & -&.009 \\\\ 16 & Mg$_{2}$ & & 0.0023 & 0.0014 & 0.007 & & 0.918 & -&.003 & 0.940 & &.0006 & 0.960 & -&.013 \\\\ 17 & Mg$b$ & & 0.053 & 0.045 & 0.20 & & 0.964 & &.108 & 0.935 & &.247 & 1.003 & -&.026 \\\\ 18 & Fe5270 & & 0.058 & 0.047 & 0.24 & & 0.923 & &.101 & 0.919 & &.180 & 0.932 & &.173 \\\\ 19 & Fe5335 & & 0.063 & 0.044 & 0.22 & & 0.960 & &.135 & 0.963 & &.032 & 0.946 & &.110 \\\\ 20 & Fe5406 & & 0.044 & 0.031 & 0.18 & & 0.874 & &.269 & 0.913 & &.165 & 0.853 & &.264 \\\\ 21 & Fe5709 & & 0.060 & 0.050 & 0.16 & & 0.979 & -&.026 & 0.907 & &.015 & 1.019 & -&.046 \\\\ 22 & Fe5782 & & 0.057 & 0.043 & 0.19 & & 0.920 & &.037 & 0.879 & -&.004 & 0.906 & &.088 \\\\ 23 & Na D & & 0.082 & 0.064 & 0.21 & & 0.990 & -&.162 & 0.979 & -&.069 & 0.993 & -&.071 \\\\ 24 & TiO$_{1}$ & & 0.0021 & 0.0017 & 0.006 & & 0.918 & -&.005 & 0.895 & -&.006 & 0.918 & &.0003 \\\\ 25 & TiO$_{2}$ & & 0.0022 & 0.0016 & 0.006 & & 0.904 & &.0007 & 0.912 & &.005 & 0.940 & &.009 \\\\ \\hline \\end{tabular} \\end{table*} \\begin{figure*} \\begin{minipage}{17cm} \\centering \\includegraphics[scale=0.2]{Hdelta_A.ps}\\includegraphics[scale=0.2]{Hdelta_F.ps}\\includegraphics[scale=0.2]{CN_1.ps}\\includegraphics[scale=0.2]{CN_2.ps}\\\\ \\includegraphics[scale=0.2]{Ca4227.ps}\\includegraphics[scale=0.2]{G4300.ps}\\includegraphics[scale=0.2]{Hgamma_A.ps}\\includegraphics[scale=0.2]{Hgamma_F.ps}\\\\ \\includegraphics[scale=0.2]{Fe4383.ps}\\includegraphics[scale=0.2]{Ca4455.ps}\\includegraphics[scale=0.2]{Fe4531.ps}\\includegraphics[scale=0.2]{C_24668.ps}\\\\ \\caption{Index by index comparison between index strengths measured on the MILES library and the Lick/IDS library. Each panel shows the residual as a function of index strength. Included are also least-square fits of the residuals (black lines, coefficients presented in Table~\\ref{offtable}) that show clear index strength dependent offsets between the two libraries. Red crosses are sigma clipped data points in the least-square fitting routine. Typical index errors are indicated in the right top corners (see text for more details). Included are also offsets to the Lick/IDS library derived in this work for the \\emph{ELODIE} (magenta lines) and \\emph{STELIB} (green lines) libraries, as well as the offsets derived in \\citet{sanchez09} for the \\emph{MILES} library (dotted lines).} \\label{offig} \\end{minipage} \\end{figure*} \\begin{figure*} \\begin{minipage}{17cm} \\centering \\includegraphics[scale=0.2]{Hbeta.ps}\\includegraphics[scale=0.2]{Fe5015.ps}\\includegraphics[scale=0.2]{Mg_1.ps}\\includegraphics[scale=0.2]{Mg_2.ps}\\\\ \\includegraphics[scale=0.2]{Mgb.ps}\\includegraphics[scale=0.2]{Fe5270.ps}\\includegraphics[scale=0.2]{Fe5335.ps}\\includegraphics[scale=0.2]{Fe5406.ps}\\\\ \\includegraphics[scale=0.2]{Fe5709.ps}\\includegraphics[scale=0.2]{Fe5782.ps}\\includegraphics[scale=0.2]{NaD.ps}\\includegraphics[scale=0.2]{TiO_1.ps}\\\\ \\includegraphics[scale=0.2]{TiO_2.ps}\\\\ \\contcaption{} \\end{minipage} \\end{figure*} ", "conclusions": "\\label{concs} We have derived new empirical fitting functions for the relationship between Lick absorption indices and stellar atmospheric parameters (T$_{\\rm eff}$, [Fe/H] and $\\log g$) described by the \\emph{MILES} library of stellar spectra, both for the resolution of the \\emph{MILES} library and for the resolution of the Lick/IDS library. The \\emph{MILES} library consists of 985 stars selected to produce a sample with extensive stellar parameter coverage. The \\emph{MILES} library was also chosen because it has been carefully flux-calibrated, making standard star derived offsets unnecessary. This becomes important when comparing stellar population models to high redshift data where no resolved individual stars are available. We find the index measurements of the \\emph{MILES} spectra to have very high quality in terms of observational index errors. These errors are also found to be significantly smaller than for the Lick/IDS library. This was expected since the \\emph{MILES} library was observed nearly thirty years after the Lick/IDS library. Given the high quality of the index measurements, index errors should not be the major error sources for the final fitting functions. We instead find indications that the stellar parameter estimates are significant error sources. Lick Index offsets between the \\emph{MILES} library and the classic Lick/IDS library are derived in order to be able to compare stellar population models based on this work with models in the literature. We find these offsets to be dependent on index strength and have therefore derived least-square fits for the residual between the two libraries. Offset to the Lick/IDS library are also derived for the flux-calibrated \\emph{ELODIE} and \\emph{STELIB} libraries. We find clear offset deviations between the libraries. The largest deviations are found for the \\emph{STELIB} library compared to the other two libraries, which is also the library having least stars in common with the Lick/IDS library. The deviations in offsets found between the three libraries undermine the derivation of universal offsets between the Lick/IDS and these flux-calibrated systems. We compare the fitting functions of this work to fitting functions in the literature, namely the fitting functions of \\citet{worthey94}, \\citet{worthey97} and \\citet{schiavon07}. Generally we find good agreement within the rms of the residuals between the data and the fitting functions of this work. The differences found in the comparisons vary significantly from index to index and especially from one stellar parameter region to another for individual indices. However, the major differences are found in the outskirts of stellar parameter space, i.e. at the temperature and metallicity ends. This is probably due to a low number of data points in these regimes for the stellar libraries, inducing uncertainties which result in the major differences found. In a forthcoming paper (Thomas et al. in prep.) the fitting functions of this work will be implemented in stellar population models following the techniques of \\citet{maraston05} and \\citet{TMB03}. A user friendly Fortran 90 code is available online at www.icg.port.ac.uk/$\\sim$johanssj to easy the implementation of our fitting functions in population synthesis codes." }, "1003/1003.2386.txt": { "abstract": "{}{}{}{}{} % 5 {} token are mandatory \\abstract % context heading (optional) % {} leave it empty if necessary {Despite all the studies, the geometry of the wind at the origin of the blueshifted broad absorption lines (BAL) observed in nearly 20\\% of quasars still remains a matter of debate.} % aims heading (mandatory) {We want to see if a two-component polar+equatorial wind geometry can reproduce the typical BAL profiles observed in these objects.} % methods heading (mandatory) {We built a Monte Carlo radiative transfer code (called MCRT) to simulate the line profiles formed in a polar+equatorial wind in which the photons, emitted from a spherically symmetric core are resonantly scattered. Our goal is to reproduce typical \\ion{C}{iv} line profiles observed in BAL quasars and to identify the parameters governing the line profiles.} % results heading (mandatory) {The two-component wind model appears to be efficient in reproducing the BAL profiles from the P Cygni-type profiles to the more complex ones. Some profiles can also be reproduced with a pole-on view. Our simulations provide evidence of a high-velocity rotation of the wind around the polar axis in BAL quasars with non P Cygni-type line profiles.} % conclusions heading (optional), leave it empty if necessary {} ", "introduction": "\\label{lintro} Depending on the selection technique and the definition used, about 20\\% to 30\\% of the quasars detected in recent surveys show the presence of the broad absorption line (BAL) troughs associated with the emission lines in their rest frame UV spectrum (e.g. Knigge et al. \\cite{kn08}, Ganguly et al. \\cite{ga08}). These BALs, reminiscent of the P Cygni-type profiles seen in the spectra of massive stars, are mainly observed in high ionization lines like \\ion{C}{iv} and \\ion{Si}{iv} and are sometimes detected in lower ionization species like \\ion{Mg}{ii}. They reveal strong outflows from quasars (Scargle \\cite{sc72}), which can reach velocities up to 0.2 c (Foltz et al. \\cite{fo83}). Despite the large number of observations, the physical and geometrical properties of the wind at the origin of the BALs remain largely unknown (e.g. Brotherton \\cite{br07}). Moreover, the distance at which those objects are found ($z \\geq 1.5$, so that \\ion{C}{iv} is shifted in the optical domain) hampers direct observation of the regions at the origin of the BALs even with the best telescopes presently available. Thus, all the information we can get about the inner regions of BAL quasars comes mainly from indirect observations. The first attempts to model the BAL profiles considered the resonant scattering of photons emitted by a continuum source in a spherically symmetric stellar-like wind (e.g. Scargle et al. \\cite{sc72}, Surdej \\& Hutsem\\'ekers \\cite{su87}). However, the growing number of observed spectra displaying a huge variety of line profiles (Korista et al. \\cite{ko93}) revealed the need for other wind models. Facing the diversity of line profiles, Turnshek (\\cite{tu84a}) proposed that BAL quasars could be broadly divided into two samples: those quasars that exhibit smooth P Cygni-type profiles, and other ones that display an absorption trough that is detached in velocity from the associated weaker and wider emission peak. These observations indicate that the properties of the wind are more complex than the simple spherically symmetric outflow inferred for stellar winds (Lee \\& Blandford \\cite{le97}). However, as emphasized by Turnshek (\\cite{tu84b}), it is very likely that distinct types of BAL QSOs do not exist but are instead different manifestations of the same phenomenon. The similarities of the emission line, optical continuum, and infrared properties of BAL and non-BAL QSOs (e.g. Weymann et al. \\cite{we91}, Gallagher et al. \\cite{ga99}, Reichard et al. \\cite{re03}, Gallagher et al. \\cite{ga07}), as well as the spectropolarimetric observations (e.g. Schmidt \\& Hines \\cite{sc99}, Ogle et al. \\cite{og99}, Lamy \\& Hutsem\\'ekers \\cite{la04}), favor a unification by orientation scheme for the BAL QSOs over the evolutionary scheme (Hazard et al. \\cite{ha84}, Becker et al. \\cite{be00}). In the unification by orientation scheme, only a fraction (roughly corresponding to the observed fraction of BAL QSOs) of the continuum source is covered by optically thick material producing the broad absorption lines, which suggests a disk-like equatorial geometry for the BAL region (e.g. Turnshek \\cite{tu84a}, Hamann et al. \\cite{ha93}, Murray et al. \\cite{mu95}, Elvis et al. \\cite{el00}, Yamamoto \\cite{ya02}). Such a geometry is supported by theoretical studies and commonly accepted, since the QSOs are thought to be powered by accretion of matter onto a supermassive black hole in the form of a disk, from which the wind could be launched. However, the recent discovery of radio loud BAL QSOs (e.g. Becker et al. \\cite{be00}) and subsequent radio variability studies reveal polar outflows in at least some of them (Brotherton et al. \\cite{br06}, Zhou et al. \\cite{zh06}, Ghosh \\& Punsly \\cite{gh07}). Models combining polar and equatorial components have also been suggested (e.g. Lamy \\& Hutsem\\'ekers \\cite{la04}) and evaluated from a theoretical point of view (Pereyra et al. \\cite{pe04}, Proga et al. \\cite{pr00}, Proga \\cite{pr03}, Proga \\& Kallman \\cite{pr04}) In this context and given the similarities between typical BAL profiles (e.g. Korista et al. \\cite{ko93}) and the line profiles produced by a two component polar+equatorial wind like the one presented by Bjorkman et al. (\\cite{bj94}), our goal in this first paper is to determine whether such a simple two-component wind can qualitatively reproduce the various types of line profiles observed among the BAL QSOs. We also try to identify the key ingredients needed to reproduce BAL profiles. In a second paper, we will investigate the effect of microlensing on these profiles, aiming at a realistic interpretation of the spectral differences observed in gravitationally lensed BAL QSOs like H1413+117 (cf. Hutsem\\'ekers et al. \\cite{hu09}). In Sect.\\ref{sec1}, we present MCRT, the Monte Carlo radiative transfer code we implemented in order to simulate resonance line profiles in a two-component axisymmetric wind. In Sect.\\ref{parmstud} we briefly identify the influence of the wind model parameters on the line profiles we computed. In Sect.\\ref{fitbal}, we show how MCRT is able to reproduce typical \\ion{C}{iv} BAL QSOs line profiles. We discuss the results of the line profile fitting and summarize our conclusions in the last two sections of the paper. %__________________________________________________________________ ", "conclusions": "In this study, we used a combination of a Monte Carlo radiative transfer code and a simple two-component polar+equatorial wind model in which photons are emitted from a central spherically symmetric source and resonantly scattered in the wind to reproduce typical \\ion{C}{iv} resonance line profiles selected from a homogeneous sample of BAL QSO spectra. Although the lack of uniqueness of the line profile fitting does not allow us to strongly constrain the geometry of the wind, we can summarize our main findings as follows \\begin{enumerate} \\item The diversity of BAL profiles produced by the adopted polar+equatorial model ranges from the typical P Cygni-type profiles to the detached absorption ones, reproducing those observed in a homogeneous sample of BAL QSOs. \\item While in some cases the line profiles can be reproduced by a single equatorial wind, we find it necessary to use a two-component polar+equatorial wind in a majority of objects. \\item The viewing angle to the wind is generally large (disk seen near edge-on); however in some cases, the line profiles can also be reproduced when assuming a pole-on view, in accordance with the results of recent radio surveys of BAL QSOs. In this context, it would be interesting to obtain good quality spectra of bona-fide polar BAL quasars and try to fit their line profiles by assuming the pole-on view. \\item The equatorial wind is rotating, and the rotational velocity at the base of the wind can reach a significant fraction of the polar terminal speed. \\end{enumerate} A possible way to break the degeneracy between the various parameter combinations of the two-component model that can reproduce the observed BAL profiles is to use gravitational microlensing. Indeed, a microlens moving across the quasar inner regions can differentially magnify the different line-forming regions, inducing line profile variations from which the geometry of the outflow can in principle be retrieved (e.g. Hutsem\\'ekers \\cite{hu93}, Hutsem\\'ekers et al. \\cite{hu94}, Lewis \\& Belle \\cite{le98}, Chelouche \\cite{ch03,ch05}). Our code MCRT has been explicitly built to integrate these microlensing effects. The effect of microlensing on BAL profiles, their use for deriving the physical properties of the outflow, and application to a known lensed system will be presented in a second paper. % %______________________________________________________________ %" }, "1003/1003.4354_arXiv.txt": { "abstract": "% We study the response of the low-degree solar p-mode frequencies to the unusual extension of the minimum of solar surface activity since 2007. Helioseismic observations collected by the space-based, Sun-as-a-star GOLF instrument and by the ground-based, multi-site network GONG (integrated signal) are analyzed. Temporal variations of the low-degree ($l=0,1,2$), p-mode frequencies are obtained. Although the known correlation of the frequency changes with the solar surface activity is recovered for the period 1996--2007, since the second half of 2007 and until July 2009 (latest period analyzed) we notice a peculiar behavior amongst modes of different angular degrees. In particular, a clear increase of the $l=0$ and $l=2$ p-mode frequencies is consistently obtained since late 2007, while the $l=1$ frequencies follow the general decreasing trend of surface activity. We interpret these differences in the frequency shifts of individual low-degree modes as indicative of variations at high latitudes in the magnetic flux beneath the surface of the Sun related to the onset of solar cycle~24. ", "introduction": "Temporal variations of the low-degree (low-$l$), solar p-mode frequencies with solar activity were first reported by \\citet{wood85}, who found that the $l = 0$ and $l=1$ mode frequencies in the 5-min band decreased by $\\sim 0.42$~$\\mu$Hz between 1980 (near solar maximum) and 1984 (near solar minimum). These early observations were later on confirmed by \\citet{palle89} using data spanning the entire solar cycle 21 (1977-1988). As higher quality and continuous helioseismic data became available, more detailed analysis were carried out revealing the highly-correlated sensitivity of the solar oscillation acoustic frequencies to the solar surface activity at low- \\citep{chaplin01,gelly02,salabert04} and high-angular degrees \\citep{chano01,howe02,salabert06a}. Angular-degree dependence of the frequency shifts at intermediate and high $l$ was also observed \\citep{chano01}. Moreover, \\citet{howe02} showed close temporal and spatial correlation of the latitude distribution of the high-degree shifts with the surface magnetic field. Marginal but still significant degree dependence of the low-degree frequency shifts was uncovered as well \\citep{chano04b,chaplin04}. However, the origin of the frequency shifts is far from being properly understood. The form and the degree dependence of the shifts would favor near-surface phenomenon but however they cannot be purely explained by structural changes: the magnetic field is somehow involved in the mechanism. As the solar oscillation frequencies, the p-mode amplitudes and linewidths, for instance, were also proven to be sensitive to the solar activity cycle in both Sun-as-a-star \\citep{chaplin00,salabert03,chano04a} and spatially-resolved \\citep{komm00,salabert06b} observations. The frequency shifts being closely correlated with solar surface activity proxies during the past solar cycles \\citep{broom09}, the response of the solar oscillations to the current unusually long and deep solar activity minimum is of particular interest. By analyzing 4768 days collected by the space-based helioseismic instrument Global Oscillations at Low Frequency (GOLF) instrument \\citep{gabriel95} onboard the {\\it Solar and Heliospheric Observatory (SOHO)} spacecraft, \\citet{salabert09} observed that the frequency shifts of the $l=0$ and $l=2$ modes show a sharp rise from the end of 2007, while no significant surface activity is visible on the Sun. On the other hand, the $l = 1$ modes follow the general decreasing trend of solar surface activity. The differences between individual angular degrees can be interpreted as different geometrical responses to the spatial distribution of the solar magnetic field beneath the surface of the Sun, indicating variations in the magnetic flux at high latitudes related to the onset of solar cycle~24. Significant variations of the p-mode frequencies during the current minimum in contrast to the surface activity observations over the same period were also reported by \\citet{broom09}. However, they did not analyzed the variations in term of individual angular degree. Furthermore, \\citet{howe09} showed that the lack of sunspots and the low-activity levels during the current minimum can be explained by a slower than usual jet stream associated with the production of sunspots. These streams originating from the poles every 11 years migrate slowly below the surface towards the equator. We present here an updated analysis of the GOLF observations until July 2009, allowing us to include an extra measurement with a 100\\% filling factor in comparison with \\citet{salabert09}. We also present results obtained by performing the same analysis on the integrated time series of the Global Oscillation Network Group \\citep[GONG;][]{harvey96}. We discuss the temporal dependence of the frequency shifts at individual angular degree and their behaviors during the current extended minimum of activity. ", "conclusions": "We analyzed integrated helioseismic observations collected by the space-based GOLF and ground-based, multi-site network GONG (integrated signal) instruments, and studied the response of the low-degree p-mode frequencies to the unusually extended and deep surface activity minimum of solar cycle~23. While closely correlated with the surface activity proxies during the past solar cycles, the temporal variations of the individual $l=0$ and $l=2$ mode frequencies (full-disk), which are more sensitive to higher latitudes, show an upturn from the end of 2007, when no significant activity is observable on the surface of the Sun. On the other hand, the variations of the $l=1$ mode frequencies (full-disk), which are less sensitive to higher latitudes, show no evidence of an upturn and follow the decreasing trend of surface activity. Similar results were obtained with the GOLF and the integrated GONG data sets. These differences between individual angular degrees may indicate that the magnetic effects related to the new solar cycle~24 happening beneath the surface and responsible for the frequency shifts have started in late 2007 at high latitudes." }, "1003/1003.2659_arXiv.txt": { "abstract": "{The single degenerate model is the most widely accepted progenitor model of type Ia supernovae (SNe Ia), in which a carbon-oxygen white dwarf (CO WD) accretes hydrogen-rich material from a main sequence or a slightly evolved star (WD +MS) or from a red giant star (WD + RG), to increase its mass and explodes when approaching the Chandrasekhar mass. The explosion ejecta may impact the envelope of and strip off some hydrogen-rich material from the companion. The stripped-off hydrogen-rich material may manifest itself by means of a hydrogen line in the nebular spectra of SNe Ia. However, no hydrogen line is detected in the nebular spectra.} {We compute the remaining amounts of hydrogen in red giant donors to see whether the conflict between theory and observations can be overcome. } {By considering the mass-stripping effect from an optically thick wind and the effect of thermally unstable disk, we systematically carried out binary evolution calculation for WD + MS and WD + RG systems. } {Here, we focus on the evolution of WD + RG systems. We found that some donor stars at the time of the supernova explosion contain little hydrogen-rich material on top of the helium core (as low as 0.017 $M_{\\odot}$), which is smaller than the upper limit to the amount derived from observations of material stripped-off by explosion ejecta. Thus, no hydrogen line is expected in the nebular spectra of these SN Ia. We also derive the distributions of the envelope mass and the core mass of the companions from WD + RG channel at the moment of a supernova explosion by adopting a binary population synthesis approach. We rarely find a RG companion with a very low-mass envelope. Furthermore, our models imply that the remnant of the WD + RG channel emerging after the supernova explosion is a single low-mass white dwarf (0.15 $M_{\\odot}$ - 0.30 $M_{\\odot}$).} {The absence of a hydrogen line in nebular spectra of SNe Ia provides support to the proposal that the WD + RG system is the progenitor of SNe Ia. } ", "introduction": "\\label{sect:1} Although type Ia supernovae (SNe Ia) are clearly important in determining cosmological parameters, e.g., $\\Omega_{\\rm M}$ and $\\Omega_{\\Lambda}$ (Riess et al. \\cite{REI98}; Perlmutter et al. \\cite{PER99}), the progenitor systems of SNe Ia have not yet been confidently identified (Hillebrandt \\& Niemeyer \\cite{HN00}; Leibundgut \\cite{LEI00}). It is widely believed that a SN Ia is produced by the thermonuclear runaway of a carbon-oxygen white dwarf (CO WD) in a binary system. The CO WD accretes material from its companion to increase its mass. When its mass reaches its maximum stable mass, it explodes as a thermonuclear runaway and almost half of the WD mass is converted into radioactive nickel-56 (Branch \\cite{BRA04}). Two basic scenarios have been presented. One is the single degenerate (SD) model, which is widely accepted (Whelan \\& Iben \\cite{WI73}; Nomoto et al. \\cite{NTY84}). In this model, a CO WD increases its mass by accreting hydrogen- or helium-rich matter from its companion, and explodes when its mass approaches the Chandrasekhar mass limit. The companion may be a main-sequence star (WD+MS) or a red-giant star (WD+RG) (Yungelson et al. \\cite{YUN95}; Li et al. \\cite{LI97}; Hachisu et al. \\cite{HAC99a}, \\cite{HAC99b}; Nomoto et al. \\cite{NOM99, NOM03}; Langer et al. \\cite{LAN00}; Han \\& Podsiadlowski \\cite{HAN04}; Chen \\& Li \\cite{CHENWC07, CHENWC09}; Han \\cite{HAN08}; Meng, Chen \\& Han \\cite{MENGXC09}; Meng \\& Yang \\cite{MENGXC09a}; L\\\"{u} \\cite{LGL09}; Wang et al. \\cite{WANGB09a,WANGB09b,WANGB09c}). The SD model has also been verified by many observations (see Meng \\& Yang \\cite{MENGXC09b}). An alternative is the theoretically less probable double degenerate (DD) model (Iben \\& Tutukov \\cite{IBE84}; Webbink \\cite{WEB84}), in which a system of two CO WDs loses orbital angular momentum by means of gravitational wave radiation and finally merges. The merger remnant may explode if the total mass of the system exceeds the Chandrasekhar mass limit (see the reviews by Hillebrandt \\& Niemeyer \\cite{HN00} and Leibundgut \\cite{LEI00}). In the single degenerate model, the companion persists after the supernova explosion. The supernova ejecta collides with the envelope of and strips some hydrogen-rich material from the surface of the companion (Marietta et al. \\cite{MAR00}; Meng, Chen \\& Han \\cite{MENGXC07}; Pakmor et al. \\cite{PAK08}). The stripped-off hydrogen-rich material may manifest itself by means of narrow H$_{\\rm \\alpha}$ emission or absorption lines in later-time spectra of SNe Ia (Chugai \\cite{CHU86}; Filippenko \\cite{FIL97}). The amount of the stripped-off material determines whether or not the narrow hydrogen line can be observed. Marietta et al. (\\cite{MAR00}) ran several high-resolution two-dimensional numerical simulations of the collision between the ejecta and the companion. They claimed that about $0.15-0.17$ $M_{\\odot}$ of hydrogen-rich material is stripped from a MS or a subgiant (SG) companion and 0.5 $M_{\\odot}$ from red giant (RG) star. Meng, Chen \\& Han (\\cite{MENGXC07}) used a simple analytic method but a more physical companion model than that used in Marietta et al. (\\cite{MAR00}) to simulate the interaction between SNe Ia ejecta and companions, and found that the minimum mass of the stripped-off material from a MS or SG star is $0.035$ $M_{\\odot}$. However, the results of Marietta et al. (\\cite{MAR00}) and Meng, Chen \\& Han (\\cite{MENGXC07}) did not include confirmation by observations, i.e., no hydrogen line was detected in nebular spectra of some SNe Ia and the upper mass limit to the stripped-off material was set to be 0.02 $M_{\\odot}$ (Mattila et al. \\cite{MAT05}; Leonard \\cite{LEO07}). Pakmor et al. (\\cite{PAK08}) used a more physical companion model than and a similar numerical simulation to that of Marietta et al. (\\cite{MAR00}) to recalculate the interaction between the supernova ejecta and companion. They found results similar to those of Marietta et al. (\\cite{MAR00}). In certain circumstances, they claimed that these results agree with observations, and hence that theory does not conflict with observations. However, the special conditions envisaged by Pakmor et al. (\\cite{PAK08}) appear to be unrealistic according to detailed binary population synthesis results (Meng \\& Yang \\cite{MENGXC09a}). Based on the results of Pakmor et al. (\\cite{PAK08}), the amount of stripped-off material is between 0.06 $M_{\\odot}$ and 0.16 $M_{\\odot}$, which is consistent with the discovery of Marietta et al. (\\cite{MAR00}) and Meng, Chen \\& Han (\\cite{MENGXC07}) (see also Meng \\& Yang \\cite{MENGXC09a}). The results of Pakmor et al. (\\cite{PAK08}) therefore do not resolve the confliction between theory and observations. Justham et al. (\\cite{JUSTHAM09}) proposed that the rotational effect of a CO WD may prevent its thermonuclear runaway until the accretion phase has ended, which could produce a RG companion with a low-mass envelope and reconcile theory and observations. We also suggest that failure to detect a hydrogen line in nebular spectra of some SNe Ia may imply that the WD + RG channel is a means of producing SNe Ia. The amount of hydrogen-rich material obtained by Marietta et al. (\\cite{MAR00}) by means of WD + RG channel is higher than observed (0.5 $M_{\\odot}$), which may be due to the simplistic RG model used by Marietta et al. (\\cite{MAR00}). In Sect. \\ref{sect:2}, we describe our binary evolution model. We present our evolutionary and binary population synthesis results in Sect. \\ref{sect:3}, and our discussions and conclusions in Sect. \\ref{sect:4}. ", "conclusions": "\\label{sect:4} \\subsection{Age}\\label{sect:4.1} We have found that if a SN Ia originates from the WD + RG channel, the donor star may almost become a helium WD and only some hydrogen-rich material remains on top of the helium WD (as little as 0.017 $M_{\\odot}$), which means that the upper limit mass to the stripped-off from the companion by supernova ejecta is 0.017 $M_{\\odot}$. No hydrogen line is then expected in the nebular spectra of some SNe Ia and our results may explain the conflict between theory and observation, i.e, theory predicts that the stripped-off material should be greater than 0.035 $M_{\\odot}$ (Marietta et al. \\cite{MAR00}; Meng, Chen \\& Han \\cite{MENGXC07}; Meng \\& Yang \\cite{MENGXC09b}), while observations indicate that the upper mass limit of the stripped-off material is 0.02 $M_{\\odot}$ (Mattila et al. \\cite{MAT05}; Leonard \\cite{LEO07}). No hydrogen line has been detected in nebular spectra of some SNe Ia may indicate that the progenitors of the observed SNe Ia are from WD + RG systems. If the observed SNe Ia (SN 2001el, 2005am, and 2005cl) were produced by the WD + RG channel, they should originate in an old population. Unfortunately, there is no constraint on the age of the three SNe Ia. We checked the types of their host galaxies and found that apart from the host galaxy of SN 2005cl (MCG-01-39-003), which is a S0 galaxy (Wang \\cite{WANGXF09}; Bufano et al. \\cite{BUFANO09}), they are both spiral galaxies, i.e. the host galaxy of SN 2001el (NGC 1448) is a Scd galaxy (Wang et al. \\cite{WANGLF03}; Wang et al. \\cite{WANGXF06}) and the host galaxy of SN 2001am (NGC 2811) is a Sa galaxy (Bufano et al. \\cite{BUFANO09}). The progenitor of SN 2005cl may therefore belong to an old population, but we are unable to infer any information about the population of the other SNe Ia. However, we note that all three SNe Ia are located at the edge of their host galaxy and that SN 2005cl is even located in a tail extending from MCG-01-39-003. Is this phenomenon evidence of an old population? It is possible because halo stars in general belong to an old population. All three SNe Ia are located at the edge of their host galaxy may be an observational select effect because we are more likely to observe a SN Ia at the outskirts of a host galaxy rather than its inner part. This selection effect might increase the probability that a SN Ia is produced by the WD +RG channel with a low-mass envelope being observed. \\subsection{WD + RG system}\\label{sect:4.2} Relative to that of the WD + MS system, the Galactic birth rate of the WD + RG channel is low (see Meng \\& Yang \\cite{MENGXC09b} and Wang, Li \\& Han (\\cite{WANGB09c})). However, since the Galactic birth rate of SNe Ia predicted by the model in Meng \\& Yang \\cite{MENGXC09b} is lower than that inferred from observations, the WD + RG channel should be carefully investigated because the progenitors of some SNe Ia (e.g. SN 2006X and SN 2007on) are possible WD + RG systems (Patat et al. \\cite{PATAT07}; Voss \\& Nelemans \\cite{VOSS08}). In addition, some recurrent nova (belonging to WD + RG) are suggested to be the candidates of SNe Ia progenitors (Hachisu et al. \\cite{HAC99b}; Hachisu \\& Kato \\cite{HK06}; Hachisu et al. \\cite{HKL07}). In this paper, we even proposed that the prevalence of the WD + RG channel is why no hydrogen line was detected in nebular spectra of some SNe Ia, although the probability of its occurrence is low. \\subsection{Interaction between supernova ejecta and companion}\\label{sect:4.3} Marietta et al. (\\cite{MAR00}) performed several high-resolution two-dimensional numerical simulations of the collision between the supernova ejecta and companion, and found that a red-giant donor will lose almost its entire envelope (96\\% - 98\\%) due to the impact leaving only the core of the star ($\\simeq0.42M_{\\odot}$). The RG star used in Marietta et al. (\\cite{MAR00}) consists of a helium core of $0.42 M_{\\odot}$ and an envelope of $0.56 M_{\\odot}$, which are not comparable to those obtained by our simulations (see Figs \\ref{Fig3} and \\ref{Fig4}). In addition, the radius of their RG model is 180 $R_{\\odot}$, which corresponds to an orbital period of $\\sim900$ days. Too long to compare with our simulation (see Fig. \\ref{Fig1}), the orbital period leads to a lower envelope binding energy than produced by the model developed in this paper since the binding energy of the envelope is determined mainly by the radius of the RG star (Meng, Chen \\& Han \\cite{MENG08}). The envelope of the RG model used by Marietta et al. (\\cite{MAR00}) is then more likely to be stripped off and the amount of material stripped-off by the RG companion in Marietta et al. (\\cite{MAR00}) might be overestimated. A more detailed numerical simulation of the interaction between supernova ejecta and an RG companion should therefore be performed by a more physical companion model than that in Marietta et al. (\\cite{MAR00}). \\subsection{Alternative explanation of the lack of hydrogen}\\label{sect:4.4} The absence of hydrogen lines in the nebular spectra of SNe Ia may have other explanations. The RG companion with a small hydrogen-rich envelope may be the results of either a fine-tuning effect as suggested in this paper, or a physical process. For example, Justham et al. (\\cite{JUSTHAM09}) suggested that a rotational effect of WD could prevent the thermonuclear runaway occurring until the accretion phase has ended, which could also produce a RG companion with a low-mass envelope. Rotation may also increase the probability of a SN Ia from WD +RG channel with low-mass envelope being observed. An alternative explanation of the lack of hydrogen is that the amount of stripped-off material might has been dramatically overestimated as discussed above. \\\\ Based on the discussions above, further attempts to observe hydrogen lines in nebular spectra of SNe Ia are encouraged to check our suggestion. The WD + RG system may also be an origin of single low-mass white dwarfs." }, "1003/1003.2690_arXiv.txt": { "abstract": "The inter-relation of clouds, solar irradiance and surface temperature is complex and subject to different interpretations. Here, we continue our recent work, which related mainly to the period from $1960$ to the present, back to $1900$ with further, but less detailed, analysis of the last thousand years. The last $20$ years is examined especially. Attention is given to the mean surface temperature, solar irradiance correlation, which appears to be present (with decadal smoothing) with a $22$-year period; it is stronger than the 11-year cycle correlation with one year resolution. UV in the solar radiation is a likely cause. Cloud data are taken from synoptic observations back to 1952 and, again, there appears to be a correlation - with opposite phase for high and low clouds - at the $20$-$30$ y level. Particular attention is devoted to answering the question, `what fraction of the observed increase in mean Global temperature ($\\sim 0.7^{\\circ}$C) can be attributed to solar, as distinct from man-made, effects?' We conclude that a best estimate is 'essentially' all from $1900$ to $1956$ and less than 14\\% from $1956$ to the present. ", "introduction": "The undoubted `global warming' over the past half century or so has focused attention on the role of changes in solar irradiance (and the solar wind) on a variety of timescales and the relevance of cloud cover (CC). As is well known, the effect of `solar forcing' on the Earth's climate is not fully understood (eg Foukal et al, $2004$; $2006$). In particular, the observed temperature changes are greater than would have been expected, so this is one reason for yet another examination of the problem. Another is our analysis of the maps of Voiculescu et al ($2006$) in which we were unable to find a good meteorological reason for the observed geographical pattern of the regions having strong cloud cover (CC), solar irradiance (denoted SI) correlations. As is well known, (eg Kristjansson and Kristiansen, $2000$ and Erlykin et al $2009$a), the observed CC, Sunspot Number (SSN) - which can be taken in first order as a proxy for SI - correlation is most unlikely to be due to cosmic ray variations, as proposed by a number of authors, and SI variations are favoured. In what follows, we refer to `SI' but are mindful that the closely related solar wind may be the operative agent, instead. The possible distinction is taken up later. Here, we examine the variations over the $20$th Century of both temperature and SI (via SSN) and CC. Data on `temperature versus time' are available for most of the $20$th Century for many points on the Earth's surface (eg Hegerl et al, $2007$) and sunspot records are readily available over an even longer period. A complementary study is made, albeit with less rigour, of the last thousand years and, in view of the contemporary significance of the extended period of low (zero) sunspot numbers, the last $20$ years is examined in some detail. The biggest problem relates to cloud cover. Clouds are inevitably connected with temperature change and their relationship to surface temperature is cloud-height dependent. Satellite data have been available since $1983$, only, and even here there are calibration uncertainties (eg Norris, 2000). Cloud data over a longer period, post-$1952$, are, however, available from the synoptic reports summarised by Norris ($1998$, $2004$). These relate to `upper' and `low' clouds and cover specific latitude ranges : $30^{\\circ}$ S - $30^{\\circ}$ N (ocean) and $30^{\\circ}$ N - $60^{\\circ}$N (ocean), and these data are used here. The main thrust of the paper is to study the correlation of changes in solar irradiance with `climate' (temperature and clouds) for the various time periods from the standpoint of both the $11$-year and $22$-year cycles. We aim to check that the temperature variations are in fact excessive and, in particular, to study the reason for the $22$-year cycle being so much stronger than that for the $11$-year cycle (as observed already by Miyahara et al, $2008$). We then go on to determine a best estimate of the contribution of `natural' (SI) effects to the well known increase of Global temperature since $1900$. It is appreciated that others have also examined this topic but an independent study is clearly desirable. ", "conclusions": "The relation between changes in solar irradiance and changes in mean surface temperature and cloud cover at various levels has been examined. The comparatively high temperature changes associated with changes in the solar irradiance (as evinced by change in sunspot numbers) are confirmed; changes in UV as distinct from longer wavelengths are a strong candidate. Positive feedback may also be a contributory process, such as occurs in the Arctic where an increased surface temperature melts ice and reduces the albedo so that the temperature rise is enhanced. A similar situation pertains in the stratosphere. The increase in $\\Delta$T with respect to expectation as a function of `integration time' - 1y, 11y and 22y - points to an inertial effect, ie the time-constant of the atmosphere/Earth's surface temperature system, but it must be said that the difference between the 22-year and 11-year temperature variations - a factor $\\sim 5$ - is rather dramatic. Our result is bigger than the factor 1.7 found by Miyahara et al (2008) (for 26-year and 12-year cycles) from tree ring data over the last 500 years. It might be thought that the 22-year cycle temperature change was due to a solar wind effect, in view of the well-known 22-year cycle of the solar magnetic field direction. Indeed, a cosmic ray origin might even be postulated. However, inspection of the CR induced atmospheric ionization versus time (eg Bazilevskaya et al, 2008) does not show an adequate 22-year modulation. Instead, we prefer an intrinsic-to-the sun origin, involving ultra-violet radiation together with positive feedback. The changes in cloud cover correlate appropriately with the temperature changes on the 20-30 y scale but the mechanism is still unclear. The longer term slow increase in low cloud cover (and small reduction in high cloud cover) is of the opposite `sign' to expectation. An explanation in terms of anthropogenic causes for the temperature rise seems likely, although it cannot be ruled out that the slow cloud cover changes are an artifact. The lack of an explanation for the actual geographical pattern of the strong correlation of low cloud cover with solar irradiance (negative)(see Section 1), is still present. An explanation in terms of the changes to the atmospheric circulation `geography' not being as predicted by the model of Haigh (2007) is a distinct possibility, although the magnitude of the effect may well be as has been estimated. The mechanism responsible for the temperature and Cloud Cover changes is clearly `solar' but whether the initiating energy is supplied by radiation (UV, as described) or whether it is the solar wind is not yet clear. However, in view of the energy in the solar wind being only of order one millionth of that in sunlight, the solar wind hypothesis has difficulties: a `positive feedback' of the magnitude required would appear to be very unlikely. Our estimate of `less than 14\\% for the period 1956-2002 is confirmed. For the previous 50 years, changes in solar irradiance appear to be responsible. Thus, extra forcing (presumably anthropogenic) started to become important only in the 1950s. This conclusion confirms that of Lean et al (2005), and others (notably that of the IPCC)." }, "1003/1003.3233_arXiv.txt": { "abstract": "We explore the novel possibility that the inflaton responsible for cosmological inflation is a gauge {\\it non}-singlet in supersymmetric (SUSY) Grand Unified Theories (GUTs). For definiteness we consider SUSY hybrid inflation where we show that the scalar components of gauge non-singlet superfields, together with fields in conjugate representations, may form a D-flat direction suitable for inflation. We apply these ideas to SUSY models with an Abelian gauge group, a Pati-Salam gauge group and finally Grand Unified Theories based on $SO(10)$ where the scalar components of the matter superfields in the $\\sixteen$s may combine with a single $\\csixteen$ to form the inflaton, with the right-handed sneutrino direction providing a possible viable trajectory for inflation. Assuming sneutrino inflation, we calculate the one-loop Coleman-Weinberg corrections and the two-loop corrections from gauge interactions giving rise to the ``gauge $\\eta$-problem'' and show that both corrections do not spoil inflation, and the monopole problem can be resolved. The usual $\\eta$-problem arising from supergravity may also be resolved using a Heisenberg symmetry. ", "introduction": "The inflationary paradigm remains extremely successful in solving the horizon and flatness problems of the standard Big Bang cosmology, and at the same time in explaining the origin of structure of the observable Universe~\\cite{Guth:1980zm,Liddle:2000cg}. Several schemes for inflation have been proposed including chaotic inflation~\\cite{Linde:1983gd}, which predicts large tensor perturbations~\\cite{Lyth:1996im}, in contrast to hybrid inflation~\\cite{Linde:1991km, Linde:1993cn, Linde:1997sj, Jeannerot:1997is, BasteroGil:1999fz} which predicts small ones. The main advantage of hybrid inflation is that, since it involves small field values below the Planck scale, it allows a small field expansion of the K\\\"ahler potential in the effective supergravity (SUGRA) theory, facilitating the connection with effective low energy particle physics models such as SUSY extensions of the Standard Model (SM) and GUTs~\\cite{Dvali:1994ms}. A long standing question in inflation models is: Who is the inflaton? We are still far from answering this question. Indeed it is still unclear whether the inflaton, the (presumed) scalar field responsible for inflation, should originate from the observable (matter) sector or the hidden (e.g.~moduli) sector of the theory. However the connection between inflation and particle physics is rather difficult to achieve in the observable sector due to the lack of understanding of physics beyond the Standard Model (SM) and in the hidden sector due to the lack of understanding of the string vacuum. However over the past dozen years there has been a revolution in particle physics due to the experimental discovery of neutrino mass and mixing~\\cite{King:2003jb}, and this improves the prospects for finding the inflaton in the observable sector. Indeed, if the SM is extended to include the seesaw mechanism~\\cite{seesaw} and SUSY~\\cite{Chung:2003fi}, the right-handed sneutrinos, the superpartners of the right-handed neutrinos, become excellent inflaton candidates. Motivated by such considerations, the possibility of chaotic (large field) inflation with a sneutrino inflaton~\\cite{Murayama:1992ua} was revisited~\\cite{Ellis:2003sq}. Subsequently three of us with Shafi suggested that one (or more) of the singlet sneutrinos could be the inflaton of hybrid inflation~\\cite{Antusch:2004hd}. Despite the unknown identity of the inflaton, conventional wisdom dictates that it must be a gauge singlet since otherwise radiative corrections would spoil the required flatness of the inflaton potential. For example in SUSY models scalar components of gauge non-singlet superfields have quartic terms in their potential, due to the D-terms, leading to violations of the slow-roll conditions which are inconsistent with recent observations by WMAP. In addition, gauge non-singlet inflatons would be subject to one-loop Coleman Weinberg corrections from loops with gauge fields which could easily lead to large radiative corrections that induce an unacceptably large slope of the inflaton potential. Furthermore a charged inflaton is in general also subject to two-loop corrections to its mass which can easily be larger than the Hubble scale~\\cite{Dvali:1995fb}. Such a contribution is in principle large enough to spoil inflation for any gauge non-singlet scalar field, leading to a sort of ``gauge $\\eta$-problem''. In this paper we shall argue that, contrary to conventional wisdom, the inflaton may in fact be a gauge {\\it non}-singlet (GNS). For definiteness we shall confine ourselves here to examples of SUSY hybrid inflation \\footnote{We note that GNS inflation may be applied to other types of inflation other than SUSY hybrid inflation.} and show that the scalar components of gauge non-singlet superfields, together with fields in conjugate representations, may form a D-flat direction suitable for inflation. Along this D-flat trajectory the usual F-term contributes the large vacuum energy. We apply these ideas first to a simple Abelian gauge group $G=U(1)$, then to a realistic SUSY Pati-Salam model, then to $SO(10)$ SUSY GUTs, where the scalar components of the matter superfields in the $\\sixteen$s may combine with a single $\\csixteen$ to form the inflaton, with the right-handed sneutrino direction providing a possible viable trajectory for inflation. We emphasize that, in sneutrino inflation models, the right-handed sneutrino has previously been taken to be a gauge singlet, as for example in SUSY GUTs based on $SU(5)$ rather than $SO(10)$. However, one of the attractive features of $SO(10)$ SUSY GUTs is that it {\\it predicts} right-handed neutrinos which carry a charge under a gauged $B-L$ symmetry. The right-handed sneutrinos of SUSY $SO(10)$, being charged under a gauged $B-L$ symmetry, have not previously been considered as suitable inflaton candidates, but here they may be. Indeed, assuming the sneutrino inflationary trajectory, we calculate the one-loop Coleman-Weinberg corrections and the two-loop corrections usually giving rise to the ``gauge $\\eta$-problem'' and show that both corrections do not spoil inflation. In addition we show that the monopole problem~\\cite{Kibble:1976sj} of $SO(10)$ GUTs can be resolved. We shall also show that the usual $\\eta$-problem arising from SUGRA~\\cite{Copeland:1994vg} may be resolved using a Heisenberg symmetry~\\cite{Binetruy: 1987} with stabilized modulus~\\cite{Antusch:2008pn}. The layout of the remainder of the paper is as follows. In Section~2 we introduce the idea of SUSY hybrid inflation with a GNS inflaton, focusing on the example of an Abelian gauge group $G=U(1)$. In Section~3 we discuss a realistic model of this kind based on the SUSY Pati-Salam gauge group, specializing to the case of the right-handed sneutrino inflationary trajectory. In Section~4 we embed the preceding Pati-Salam model into $SO(10)$ SUSY GUTs. Section~5 confronts the issues associated with radiative corrections for a GNS inflaton at one and two loops. Section~6 shows how the $\\eta$\\,-\\,problem in SUGRA may be resolved in this class of models using a Heisenberg symmetry with stabilized modulus. Section~7 summarizes and concludes the paper. ", "conclusions": "In this paper we have explored the novel possibility that, within SUSY GUTs, the inflaton responsible for cosmological inflation is a gauge {\\it non}-singlet under some gauge group $G$. For definiteness we have considered SUSY hybrid inflation where we have shown that the scalar components of gauge non-singlet superfields, together with fields in conjugate representations, may form a D-flat direction suitable for inflation. We have first sketched an explicit example of this scenario based on the Abelian gauge group $G=U(1)$. We then presented a realistic model of this kind based on the relevant part of the SUSY Pati-Salam gauge group under which the inflaton transforms, namely $G=SU(4)_C\\times SU(2)_R$. In such a framework we have shown how it is possible for the inflaton to consist of the scalar components of a pair of gauge non-singlet matter superfields $R^c$ and $\\oR^c$, in conjugate representations under the Pati-Salam gauge group, rolling along a D-flat valley and coupled to a pair of gauge non-singlet Higgs superfields $H^c$ and $\\oH^c$, also in conjugate representations of the same gauge group, whose VEVs end inflation, according to the hybrid inflation scenario, breaking the Pati-Salam gauge group. We emphasize that it is the components of the matter superfields which form the inflaton. We have then extended the model to $SO(10)$ SUSY GUTs. Such a scenario is perfectly suited to sneutrino inflation in SUSY GUTs, allowing the inflaton to be a conjugate pair of right-handed sneutrinos, and the pair of Higgs superfields to break the GUT gauge group at the end of inflation. We have shown that in this case, if the inflaton and Higgs directions relevant for inflation lie along the right-handed neutrino direction, then this mechanism for inflation solves the monopole problem. Assuming the sneutrino trajectory for simplicity, we have then systematically examined the obvious objections to having a charged inflaton, namely the one- and two-loop gauge corrections to the potential which might be thought to threaten the flatness of the potential and so violate the slow-roll conditions, and have shown that such corrections do not pose a threat to this scheme. The key to the success of this mechanism is that the inflaton VEV during the inflationary epoch breaks the GUT gauge group but preserves D-flatness so only the F-term breaks SUSY. The inflaton therefore only couples to gauge bosons and gauginos which are heavy (and degenerate) which effectively suppresses the one- and two-loop gauge corrections. With the inclusion of SUGRA, the $\\eta$-problem may be resolved by appealing to a Heisenberg symmetry which involves a modulus field stabilized during inflation. We remark that the conjugate matter representations naturally arise from string theory constructions where generically several copies of matter of the $SO(10)$ $\\sixteen$ and $\\csixteen$, for example, appear as massless modes, where there are three more $\\sixteen$s than $\\csixteen$s which accounts for the three chiral families. In such a framework we are suggesting that one or more pairs of the extra $\\sixteen$s and $\\csixteen$s could be responsible for inflation, and their coupling to Higgs fields might trigger part of the GUT symmetry breaking at the end of inflation, without leading to excessive monopole abundance. The components of the extra $\\sixteen$s and $\\csixteen$s which develop VEVs during inflation lie along the right-handed neutrino directions, and mixing of these components with the physical right-handed neutrinos (in the three chiral $\\sixteen$s) could lead to interesting consequences associated with reheating and non-thermal leptogenesis at the end of inflation which should be explored in future work. In conclusion, we find that the idea that the inflaton is a gauge non-singlet is viable in the framework of SUSY hybrid inflation, and this opens up the possibility of having right-handed sneutrino inflation in Pati-Salam or $SO(10)$ SUSY GUTs." }, "1003/1003.4763_arXiv.txt": { "abstract": "Inflation predicts primordial scalar perturbations with a nearly scale-invariant spectrum and a spectral index \\textit{approximately} unity (the Harrison--Zel'dovich (HZ) spectrum). The first important step for inflationary cosmology is to check the consistency of the HZ primordial spectrum with current observations. Recent analyses have claimed that a HZ primordial spectrum is excluded at more than $99\\%$ c.l. Here we show that the HZ spectrum is only marginally disfavored if one considers a more general reionization scenario. Data from the Planck mission will settle the issue. ", "introduction": " ", "conclusions": "" }, "1003/1003.0283_arXiv.txt": { "abstract": "The thermal dominant state in black hole binaries (BHBs) is well understood but rarely seen in ultraluminous X-ray sources (ULXs). Using simultaneous observations of M82 with Chandra and XMM-Newton, we report the first likely identification of the thermal dominant state in a ULX based on the disappearance of X-ray oscillations, low timing noise, and a spectrum dominated by multicolor disk emission with luminosity varying to the 4th power of the disk temperature. This indicates that ULXs are similar to Galactic BHBs. The brightest X-ray spectrum can be fitted with a relativistic disk model with either a highly super-Eddington ($L_{\\rm disk}/L_{\\rm Edd} = 160$) non-rotating black hole or a close to Eddington ($L_{\\rm disk}/L_{\\rm Edd} \\sim 2$) rapidly rotating black hole. The latter interpretation is preferred, due to the absence of such highly super-Eddington states in Galactic black holes and active galactic nuclei, and suggests that the ULX in M82 contains a black hole of 200-800 solar masses with nearly maximal spin. On long timescales, the source normally stays at a relatively low flux level with a regular 62-day orbital modulation and occasionally exhibits irregular flaring activity. The thermal dominant states are all found during outbursts. ", "introduction": "Ultraluminous X-ray sources (ULXs) are non-nuclear X-ray sources with apparent luminosities above the Eddington limit of stellar-mass black holes (BHs). Variable ULXs are black hole binaries (BHBs) and may harbor intermediate-mass BHs \\citep{col99,mak00}. The emission of BHBs has been classified into four states based on spectral and timing properties: the quiescent, hard, thermal dominant (TD), and steep power-law states \\citep{mr06,rem06}. The TD state is the best understood and is well described by the standard accretion disk model \\citep{sha73}. The emergent X-ray spectrum is described by a multicolor disk (MCD) model with two parameters: the disk inner radius ($R_{\\rm in}$) and the temperature ($T_{\\rm in}$) at that radius. In the TD state, $R_{\\rm in}$ is constant and the accretion disk is thought to extend all the way to the ``innermost stable circular orbit'' (ISCO) around the BH \\citep{tak97}. The ISCO radius depends solely on the mass and spin of the BH. Therefore, a TD spectrum from an accreting BH can be used to shed light on its mass and spin. The TD state has been found during outbursts of many BHBs and exhibits specific properties. The disk bolometric luminosity is $L_{\\rm disk} = 4 \\pi \\sigma R_{\\rm in}^2 T_{\\rm in}^4$. For constant $R_{\\rm in}$, a 4th power relation between the luminosity and the disk inner temperature, $L_{\\rm disk} \\propto T_{\\rm in}^4$, is observed \\citep{kub04}. Due to up scattering of disk photons in the corona, sometimes the observed exponent is lower than 4, but can be recovered by applying hardening correction \\citep{mcc09}. Moreover, sources in the TD state have low levels of short term variability, with very weak or absent narrow band timing noise (quasi-periodic oscillations; QPOs) and weak broad band power continuum \\citep{mr06,rem06}. Spectral surveys of ULXs in nearby galaxies show they are rarely in the TD state \\citep{fen05,sto06,win06,sor09}. The spectra of a few ULXs are super-soft and can be modeled by an MCD with an inner temperature close to 0.1 keV. However, repeated observations revealed an evolution pattern inconsistent with the $L_{\\rm disk} \\propto T_{\\rm in}^4$ relation ruling out interpretation as thermal dominant disk emission \\citep{liu08}. Some other ULXs show soft excesses in the spectra that could be fitted by cool disk emission \\citep{kaa03,mil04}. However, the cool disk is not the dominant component in their spectra, thus, the sources are not in the TD state. Finding a ULX in the TD state would demonstrate that the ULX has emission states similar to Galactic BHBs and allow inference of the BH mass and spin. The starburst galaxy M82 contains one of the most luminous ULXs, CXOM82 J095550+694047 \\citep[=\\xa;][]{kaa01,mat01}, in nearby galaxies. The ULX was first identified with Chandra at a luminosity higher than the Eddington limit of a $500 M_{\\odot}$ BH \\citep{kaa01,mat01}. On the sky, it lies near a super star cluster, which was speculated to be the birth place of a massive BH \\citep{por04}. Low frequency QPOs and broadband timing noise, detected in the central region of M82 \\citep{str03,dew06,muc06} and later confirmed to originate from this ULX \\citep{fen07}, suggest that the ULX harbors a massive BH. \\citet{kin05} also suggests that the ULX contains a massive BH which is the nucleus of a satellite galaxy merging with M82. Positive identification of the emission states requires both timing and spectral information. Here, we describe simultaneous observations exploiting the high angular resolution of Chandra to isolate the ULX spectrum from diffuse emission and nearby sources and the large collecting area of XMM-Newton to obtain timing information. We use the joint timing and spectral information to identify the ULX emission states and unveil the nature of the source. ", "conclusions": "Joint Chandra and XMM-Newton observations have enabled us to precisely measure the spectral and timing behavior of this ULX in M82. Disappearance of QPOs and low timing noise, an X-ray spectrum best-fitted with an MCD model, and an $L_{\\rm disk} \\propto T_{\\rm in}^4$ pattern in the spectral evolution provide strong motivation to interpret this behavior in terms of the TD state seen in stellar-mass BHBs. The MCD model indicates that the source spectrum shows spectral curvature at a few keV and thus could also be fitted by more complicated models like a cool optically thick corona which has been proposed recently as an indicator of a new `ultraluminous state' \\citep{gla09}. However, a transition from the hard spectrum (incompatible with the steep power-law state) seen at lower luminosities to the ultraluminous state would be difficult to explain and we, therefore, do not discuss this proposed state in detail. Simply scaling the $L_{\\rm disk} \\propto T_{\\rm in}^4$ pattern to those from Galactic BHBs suggests that the source contains a more massive BH, because the compact object mass is proportional to the square root of the disk luminosity if two BHs have the same disk inner temperature and spin \\citep{mak00}. Fitting with a fully relativistic MCD model leads to a consistent result. A conservative estimate of the BH mass by allowing an Eddington ratio of as high as 10 suggests a BH of at least 200 solar masses but less than 800 solar masses. This is coincident with the theoretical estimate of the mass of this BH if it is created in a nearby star cluster by runaway collisions \\citep{por04}. The fast spin of the source makes it efficient in extracting gravitational energy. The spectral fitting also indicates a relatively high inclination angle, 59-79 degrees, of the accretion disk. For a maximally spinning BH of a few hundred solar masses, the relativistic effect has largely eliminated the limb-darkening effect; viewing the disk at a high inclination angle receives almost the same flux as at a low angle. Therefore, the observed high flux is not a problem for a nearly edge-on disk. The X-ray flux from M82 is modulated at a period of 62 days, interpreted as due to orbital motion of this ULX binary, which must contain a giant or super-giant star inferred from the period assuming Roche-lobe overflow \\citep{kaa06,kaa07}. With knowledge of the BH mass (assuming 200-800 solar masses) and the 62-day binary orbital period, the binary separation of the ULX can be calculated with the assumption of Roche lobe overflow as $(3-4) \\times 10^{13}$~cm, which is insensitive to the mass of the companion star. This is larger than all separations of low-mass BHBs with a dynamical measurement of the mass, and larger than that of GRS 1915+105, which has the largest separation known so far, by a factor of a few. This ULX and GRS 1915+105 share a similarity that they have both been active for many years without quenching to the quiescent state. The large separation and consequently a huge reservoir of accretion mass could be the factor that determines their long-term activity \\citep{dee09}." }, "1003/1003.5659_arXiv.txt": { "abstract": "It is well known that the classical gravitational two body problem can be transformed into a spherical harmonic oscillator by regularization. We find that a modification of the regularization transformation has a similar result to leading order in general relativity. In the resulting harmonic oscillator, the leading-order relativistic perturbation is formally a negative centrifugal force. The net centrifugal force changes sign at three Schwarzschild radii, which interestingly mimics the innermost stable circular orbit (ISCO) of the full Schwarzschild problem. Transforming the harmonic-oscillator solution back to spatial coordinates yields, for both timelike and null weak-field Schwarzschild geodesics, a solution for $t,r,\\phi$ in terms of elementary functions of a variable that can be interpreted as a generalized eccentric anomaly. The textbook expressions for relativistic precession and light deflection are easily recovered. We suggest how this solution could be combined with additional perturbations into numerical methods suitable for applications such as relativistic accretion or dynamics of the Galactic-centre stars. ", "introduction": "When Schwarzschild geodesics appear in classic tests of general relativity, the important result is an integral over the geodesic: orbital precession or deflection of light. Similary, in modern tests of relativity involving binary pulsars \\citep[for a review, see][]{lrr-2006-3} the observable effects are also cumulative over many orbits. In the case of the recently-discovered S-stars near the Galactic Centre, the cumulative effects of relativity are no longer the principal quantity of interest. The highly eccentric examples from the S-stars \\citep{2008ApJ...689.1044G, 2009ApJ...692.1075G}, which experience a range of gravitational regimes, motivate an interest in tracking relativistic effects as they vary along an orbit. In particular, some recent work has drawn attention to relativistic effects on redshifts near pericentre passage \\citep{2006ApJ...639L..21Z,2009ApJ...690.1553K,2010ApJ...711..157A}. These effects can be calculated numerically, and some of them also by post-Newtonian perturbation theory, but a simpler method is desirable. Such a method is suggested by Levi-Civita or Kustaanheimo-Stiefel (LC or KS) regularization, which are transformations of the classical gravitational two-body problem to an equivalent harmonic oscillator. This type of regularization was originally introduced in two dimensions \\citep{1920LC} and much later extended to three dimensions \\citep{1964Kustaanheimo,1965KS}. KS regularization has an extensive literature, including applications to $N$-body simulations \\citep{1974CeMec..10..185A,1974CeMec..10..516A,1989ApJS...71..871J,1993CeMDA..57..439M}. The classical result suggests that the LC or KS regularization could be used to transform the general relativistic problem into a perturbed harmonic-oscillator. We find even better: a modification of the LC/KS transformation acting on the geodesics of the leading order Schwarzschild metric in the isotropic or harmonic gauge \\citep[cf.][\\unskip~Section~8.2]{1972gcpa.book.....W}, \\begin{equation} ds^2 = - \\left( 1 - {2M\\over r} + {2M^2\\over r^2} + O\\left(M^3\\over r^3\\right) \\right) dt^2 + \\left( 1 + {2M\\over r} + O\\left(M^2\\over r^2\\right) \\right) d{\\bf x}^2, \\label{metric} \\end{equation} yields an unperturbed circular/spherical harmonic oscillator. As a result, the solution is analytic. The difference from the classical case is a negative centrifugal-force term in the transformed space. This terms encodes the leading-order effects of precession, deflection of light, and the innermost stable orbit. Because orbits in the Schwarzschild spacetime do not leave the orbital plane, in this paper we mainly consider the two-dimensional or LC case. The three-dimensional or KS case is similar, but algebraically more complicated, as it involves introducing a fourth spatial dimension. ", "conclusions": "We have derived timelike and null geodesics in the leading-order Schwarzschild metric in terms of elementary functions. The expressions (\\ref{solnb}) for bound orbits and (\\ref{solnu}) for unbound orbits, together with (\\ref{solnull}) for light rays, are all simple generalizations of well-known expressions in classical celestial mechanics. The usual formulas for relativistic orbital precession and light deflection are easily recovered. A feature resembling the innermost stable circular orbit in the full Schwarzschild metric is also present. The technique we have used is a modification of the Levi-Civita or Kustaanheimo-Stiefel regularization transformation and transforms the geodesic equation into a spherical harmonic oscillator. The simplicity of the result, notwithstanding the non-trivial route used to derive it, hints at some underlying symmetry in the Schwarzschild problem. We speculate that it is somehow related to the separability of the Hamilton-Jacobi and other equations in the Schwarzschild and Kerr metrics \\citep[cf.][]{1983mtbh.book.....C} but have not attempted to investigate this. As mentioned in the Introduction, the original motivation for this work was to find useful formulas applicable to the highly-eccentric Galactic-centre stars, whose orbits pass through a large range of gravitational regimes. Future observations of these stars aiming to detect relativistic effects will require computation of relativistic effects on both stellar orbits and light rays at many points along an orbit, for many orbits, in order to fit the orbital parameters. The solutions in this paper allow a simpler, more efficient method for carrying out those computations. The analytic solutions will not be sufficient on their own because the Galactic-centre stars also experience additional Newtonian perturbations due to local matter \\citep{2008AJ....135.2398M}, but they can be incorporated into numerical methods, specifically, generalized leapfrog integrators. Such algorithms evolve alternately under two Hamiltonians, which are integrable separately. The idea goes back to \\cite{1991AJ....102.1528W} and \\cite{1991CeMDA..50...59K}. Some recent developments on adaptive stepsizes appear in \\cite{2007CeMDA..98..191E} and are applied to the specific problem of Galactic-centre stars in \\cite{2009ApJ...703.1743P}. We note, however, that the present work is limited to test particles, and hence will not be applicable for binary orbits or self-gravitating disc simulations unless a generalization is found. Another potential application may be the use of the solutions in relativistic disc simulations as an alternative to the widely used pseudo-Newtonian potentials \\citep[see especially][]{1980A&A....88...23P,1996ApJ...461..565A,2009A&A...500..213A}, an advantage being that the solutions in this paper are well-defined approximations and include a more complete repertoire of general-relativistic effects for the same computational budget." }, "1003/1003.0556_arXiv.txt": { "abstract": "% Variability is a main property of active galactic nuclei (AGN) and it was adopted as a selection criterion using multi epoch surveys conducted for the detection of supernovae (SNe). We have used two SN datasets. First we selected the AXAF field of the STRESS project, centered in the Chandra Deep Field South where, besides the deep X-ray surveys also various optical catalogs exist. Our method yielded 132 variable AGN candidates. We then extended our method including the dataset of the ESSENCE project that has been active for 6 years, producing high quality light curves in the R and I bands. We obtained a sample of $\\sim$4800 variable sources, down to R=22, in the whole 12 deg$^{2}$ ESSENCE field. Among them, a subsample of $\\sim$500 high priority AGN candidates was created using as secondary criterion the shape of the structure function. In a pilot spectroscopic run we have confirmed the AGN nature for nearly all of our candidates. ", "introduction": "% Since the discovery of AGN, variability was established as a main property of the population and it was among the first ones to be explored \\citep{KBSmith63}. The luminosities of AGN have been observed to vary in the whole electromagnetic range and the majority of the objects are exhibiting continuum variations of about 20\\% on timescales of months to years \\citep{KBHook94}. From a physical point of view, variations can set limits on the size of the central emitting region and the differences in the variability properties in the X-ray, optical and radio bands provide important information on the underlying structure. The mechanism of variability itself is still unknown and a variety of models have been proposed \\citep{KBTerl92,KBHawk93,KBKawa98}. Thus, the study of the AGN variability is very important and can put constraints on the models describing the AGN energy source and the AGN structure. On the other hand, supernovae (SN) are very powerful cosmological probes and their systematic discovery outside the local Universe has led to major scientific results, like the confirmation that the Universe is accelerating \\citep{KBRiess98,KBPerl99}. Such studies require well sampled light curves and large statistical samples which can be achieved by monitoring wide areas of the sky to very faint limiting magnitudes. This kind of surveys produce huge amounts of data that can be suitable also for other scientific studies. For example, given that the time sampling is adequate, data gathered during SN searches can be used to detect AGN through variability. One of the main purposes of this work is to explore this possibility and create suitable tools for the efficient selection of AGN in such databases. The two projects that have provided us with their data are: the Southern inTermediate Redshift ESO Supernova Search \\citep[STRESS,][]{KBBott08} and the ESSENCE (Equation of State: SupErNovae trace Cosmic Expansion) survey \\citep{KBmikn07}. ", "conclusions": "We have applied a variability selection method to data collected by SN searches in order to detect AGN through variability. We have been very successful in detecting new AGNs, which had escaped traditional selection techniques and have confirmed a large number of already known AGNs in these fields. This proves that the AGN field can benefit from such synergic AGN-SN surveys. After a spectroscopic follow-up, a considerable fraction of our variable candidates turned out to be ``variable galaxies'' with narrow emission lines and properties consistent with LLAGNs diluted by the host galaxy. By combining the criterion for variability with a secondary criterion concerning the shape of the SF, we created highly reliable AGN samples, since $\\sim$97\\% of our candidates belonging to such a sample was confirmed as BLAGN. In an era that large survey telescopes are being developed, such variability studies can give valuable feed-back both for determining the strategy of the observations as well as for the development of software and pipelines that will allow the scientific community to fully exploit the huge datasets that will be produced." }, "1003/1003.2881_arXiv.txt": { "abstract": "We analyze the dynamics of a Dirac-Born-Infeld (DBI) field in a cosmological set-up which includes a perfect fluid. Introducing convenient dynamical variables, we show the evolution equations form an autonomous system when the potential and the brane tension of the DBI field are arbitrary power-law or exponential functions of the DBI field. In particular we find scaling solutions can exist when powers of the field in the potential and warp-factor satisfy specific relations. A new class of fixed-point solutions are obtained corresponding to points which initially appear singular in the evolution equations, but on closer inspection are actually well defined. In all cases, we perform a phase-space analysis and obtain the late-time attractor structure of the system. Of particular note when considering cosmological perturbations in DBI inflation is a fixed-point solution where the Lorentz factor is a finite large constant and the equation of state parameter of the DBI field is $w=-1$. Since in this case the speed of sound $c_s$ becomes constant, the solution can be thought to serve as a good background to perturb about. ", "introduction": "The inflationary paradigm remains to date the most successful explanation for the origin of the observed temperature fluctuations of the cosmic microwave background (CMB) (see, e.g., \\cite{Linde,Lyth:1998xn} for reviews). However, establishing its origin in fundamental theory has not been quite as successful and so for many it remains a fascinating paradigm in search of an underlying theory. The favourite candidate for this is String theory and for a nice review of the construction of inflation models in string theory, see \\cite{Baumann:2009ni}. One interesting model, recently proposed from string theory is Dirac-Born-Infeld (DBI) inflation \\cite{Silverstein:2003hf,Alishahiha:2004eh,Chen:2004gc,Chen:2005ad,Shandera:2006ax}, where inflation is driven by the motion of a D3-brane in a warped throat region of a compact internal space. In this model, since the inflaton is the position of a D-brane with a DBI action, its kinetic term is inevitably non-canonical. In addition to this kinetic term, its effective action includes a potential arising from the quantum interaction between D-branes, with the brane tension encoding geometrical information about the throat region of the compact space. Because of these novel ingredients, the predictions of DBI inflation are quite different from the ones from the standard slow-roll inflation models and it has led to an intense period of research into the scenario, including work into the background dynamics and linear perturbations \\cite{Spalinski:2007dv,Spalinski:2007qy,Chimento:2007es,Ward:2007gs,Spalinski:2007un,Kinney:2007ag,Tzirakis:2008qy,Czuchry:2008km}. From the phenomenological viewpoint, one of the main reasons why this model has attracted attention is because of its sizable equilateral type of primordial non-Gaussinaity, first pointed by \\cite{Silverstein:2003hf} and explored further in \\cite{Alishahiha:2004eh,Chen:2004gc,Chen:2005ad,Chen:2005fe,Chen:2006nt,Huang:2006eh,Arroja:2008ga,Langlois:2008wt,Langlois:2008qf,Arroja:2008yy,Langlois:2009ej,Gao:2009gd,Chen:2009bc,Us,Mizuno:2009cv,Gao:2009at,Mizuno:2009mv,RenauxPetel:2009sj,Chen:2009zp,Koyama:2010xj,Chen:2010xk}. Such a large signal can not be obtained in standard slow-roll inflation, hence it opens up the possibility of distinguishing this model from that of the standard slow-roll inflation-- although not necessarily distinguishing it from more general slow-roll models going beyond single field inflation as they may also lead to large signals. Furthermore, because of the dependence of the effective four dimensional string tension on the nature of the compact internal space, it is possible to place further constraints on the parameters related with compactifications such as flux numbers. For the most up to date observational constraints on and consequences of DBI inflation, see \\cite{Kecskemeti:2006cg,Lidsey:2006ia,Baumann:2006cd,Bean:2007hc,Lidsey:2007gq,Peiris:2007gz,Kobayashi:2007hm,Gmeiner:2007uw,Lorenz:2007ze,Bean:2007eh,Bird:2009pq,Bessada:2009pe,Fuzfa:2005qn,Fuzfa:2006pn}. Recently the idea of low scale inflation arising from the DBI action has been invoked to explain the late time acceleration we associate with dark energy \\cite{Martin:2008xw,Ahn:2009hu,Chiba:2009nh,Ahn:2009xd}. Given the results mentioned above and inspired by the possibility of inflation being an attractor solution in DBI models, we believe there is a need to understand the late-time attractor structure for as general a DBI set-up as possible. It is well known that for a canonical scalar field with a potential, scaling solutions can exist where the ratio of the kinetic and potential energy of the scalar field maintain the same ratio, and understanding the stability of these solutions is important in determining the nature of the late-time solutions \\cite{Lucchin:1984yf,Ratra_Peebles,trac,Ferreira:1997hj,Copeland:1997et,vdHCW,Heard:2002dr,Padmanabhan:2002cp,Tsujikawa:2004dp,Calcagni:2004wu,Copeland:2004qe,Copeland:2009be}. Among the earlier work, the phase-space analysis proposed by \\cite{Copeland:1997et} is particularly powerful because it allows us to make use of suitable dimensionless dynamical variables, in order to establish the global stability of such scaling solutions. For related work which analyzes the dynamics in DBI models with general inflationary potentials, see \\cite{Meng:2004ap,Underwood:2008dh,Franche:2009gk,Franche:2010yj}. Now, this method for analysing the stability of the solutions has only been applied to the case where the DBI field has a quadratic potential and the associated D-brane is in the anti-de Sitter throat \\cite{Guo:2008sz}. The late-time attractor structure for the case without a potential, known as tachyon cosmology, has been studied extensively, \\cite{Aguirregabiria:2004xd,Piazza:2004df,Copeland:2004hq,Gumjudpai:2006hg,Tsujikawa:2006mw,Gong:2006sp,Sen:2008yt,Quiros:2009mz,Li:2010eu} and the case with another degree of freedom, has been studied in \\cite{Gumjudpai:2009uy,Saridakis:2009uk}. In this paper, in order to make the cosmological application of scaling solutions in DBI models more complete and as a natural extension of \\cite{Guo:2008sz}, we obtain the late-time attractor structure of the system including a perfect fluid plus a DBI field whose potential and brane tension are arbitrary power-law or exponential functions of the DBI field. The rest of the paper is arranged as follows. In section \\ref{basic-eqn} we present the model and basic equations. Then, in section \\ref{section_powerlaw} we consider the models where the potential and brane tension are arbitrary power-law functions of the DBI field. Special emphasis is given to a new set of fixed-point solutions which at first site appear singular in the equations of motion, but on closer inspection are well defined. This is followed in section \\ref{section_exp} with an analysis of the models where the potential and brane tension are arbitrary exponential functions of the DBI field. Finally, we summarise in section \\ref{sec_summary}. ", "conclusions": "} Successful models of inflation arising within string theory from the DBI action have generated a great deal of interest recently. With their non-canonical kinetic terms, non-trivial potentials and brane tensions, they have led to a number of fascinating results including the prediction of distinctive non-Gaussian fluctuations in the CMB. Although most models investigated to date have had specific functional forms for the potential and brane tension, in this paper we have decided to broaden the class of models being discussed and so have analysed the dynamics associated with more general forms for these potential and brane tension functions, including in the analysis the presence of a background perfect fluid in a flat FRW universe. Following the approach developed in \\cite{Copeland:1997et}, we have introduced a suitable set of dynamical variables $x$, $y$ and $\\tilde{\\gamma}$ in Eq.~(\\ref{defn_xy}) which has allowed us to determine the phase-space portrait of the system. In particular, we have established the late time behaviour of these systems, demonstrating where appropriate the attractor nature of the solutions. In Sec.~\\ref{section_powerlaw}, we have considered the models where the potential and brane tension are given by power-law functions of the DBI field ($V(\\phi) = \\sigma |\\phi|^p$, $f(\\phi) = \\nu |\\phi|^r$). The standard fixed-point solutions of this system are summarised in TABLE~\\ref{fixed-points-Summary_power}, where we see that the late-time attractor nature of the solutions depends on $q \\equiv -1/(p+r)$. The interesting cases of scaling where the ratio of the kinetic to potential energies of the DBI field is a constant are found to exist only for $q=1/2$ ($(a4)$, $(a5)$) and $q=0$ ($(b2)$, $(b3)$). This is because if we require $x$ and $y$ to be nonzero constants, there are only two possibilities, that is, $\\tilde{\\gamma}=0$ (leading to scaling with $q=1/2$) and $\\tilde{\\gamma}=1$ (leading to scaling with $q=0$). These scaling solutions are then shown to be stable for certain regions of the parameter space. In addition to these, we have also explicitly demonstrated the existence and stability of an interesting inflationary solution $(c3)$ specific to the DBI field with constant $\\tilde{\\gamma}$ which differs from $\\tilde{\\gamma}=1$ and $\\tilde{\\gamma}=0$. The DBI system is rich. For example the evolution equations (\\ref{x_evol_eq_mod})-(\\ref{gamma_evol_eq_mod}) can appear singular when some of $x$, $y$ or $\\tilde{\\gamma}$ either tend to zero or unity, which one is singular depends on the value of $q$. On the face of it, the equations appear to be ill-defined, but in practice it turns out that these points can actually be late-time attractor solutions for the system. In section~\\ref{fixedpoints0/0} we obtain these fixed points and determine their stability. These are summarised in Table~\\ref{fixed-points-Summary_0/0}. Having established all the fixed-points and their stability for a given set of parameters, we have gone on to determine which of these solutions will be the late-time attractor. Particular care is required when considering the cases where the eigenvalues associated with the perturbations vanish, implying that the solution is marginally stable. Whether these are the late time attractors for the system depends upon the stability of the other allowed fixed-points. In particular those fixed points with three negative eigenvalues, means that these marginally stable fixed-point solutions can not be the late-time attractor as their stability is weaker than that of the fixed-points. However, if all other fixed-points have positive eigenvalues, then the marginally stable fixed-point solution can turn out to be the late-time attractor. We have summarised the possible late-time attractor solutions for a given set of $q$, $p$, $r$ in Table~III. In Sec.~\\ref{section_exp}, we have considered the models where the potential and brane tension are exponential functions of the DBI field ($V(\\phi) = C e^{-\\lambda \\phi}$, $f(\\phi) = D e^{-\\mu \\phi})$. This system has a similar dynamical structure to that of the power-law models with $q=0$. However, there are additional scaling solutions present in the exponential case because we can construct an autonomous system for the case with $\\lambda = -\\mu$, leading to solutions where $\\tilde{\\gamma}$ is a constant between $0$ and $1$ ($(c4)$,$(c5)$). In this case, we find that once we specify the values of the parameters, $\\lambda$, $\\sigma$, $\\nu$ and $w_m$ we can judge which of these two fixed-points will be the late-time attractor. The stability of these two fixed-points are summarised in Table~IV. We have also shown that the special case ($V(\\phi) = \\sigma |\\phi|^p$, $f(\\phi) = \\nu |\\phi|^{-p}$) can be discussed in terms of the limiting behaviour $\\lambda \\to 0$ or $\\lambda \\to \\infty$ in the models with an exponential potential and brane tension satisfying $\\lambda=-\\mu$. There is an overlap between elements of this work and other published material. For example, we are able to reproduce a number of results obtained earlier in \\cite{Guo:2008sz} where the authors considered the case with a massive potential and AdS throat. It turns out to be a special case of $q=-1/2$ discussed in Sec.~III in this paper. The existence of the scaling solution ($c5$) was originally pointed out by Martin and Yamaguchi in \\cite{Martin:2008xw}. The authors stability analysis followed that of \\cite{Ratra_Peebles}, which although providing a proof that the fixed-point solution is attractive, it did not explain how the the solution could be realised from all initial values. In our approach we have gone into more details, following the analysis of \\cite{Copeland:1997et}. In particular by making the phase space compact, we have been able to establish all the fixed-points in the compact phase space, allowing us to then properly discuss the late-time attractor structure of the solution. Of most interest to us though is the existence and stability of a class of cosmologically relevant solutions. We have found that a fixed point solution $(c3)$ (relativistic potential dominated solutions) where $\\tilde{\\gamma}$ is a constant satisfying $ 0 < \\sqrt{3}/\\sqrt{\\tilde{\\lambda}^2 + 3} < 1$ can be the late-time attractor for $q=-1/2$ with $r/p < -1$ (see also \\cite{Ahn:2009xd}). Although it can not be applied to the AdS throat ($r=-4$) as $r/p > -1$, we believe this solution is very interesting and important. In calculations of primordial perturbations of the DBI inflation models, properly incorporating the time dependence of the sound speed $c_s$ is a complicated issue \\cite{Garriga:1999vw,Alishahiha:2004eh}, and in fact it is usually assumed to be constant. (For a recent approach to relaxing this assumption, see \\cite{Lorenz:2008et}.) Since in our case $c_s=\\tilde{\\gamma}$ is constant for the fixed-point solution $(c3)$, this serves as a good background to use when we consider cosmological perturbations in DBI inflation." }, "1003/1003.0884_arXiv.txt": { "abstract": "{}{To investigate the positions and source sizes of X-ray sources taking into account Compton backscattering (albedo).} {Using a Monte Carlo simulation of X-ray photon transport including photo-electric absorption and Compton scattering, we calculate the apparent source sizes and positions of X-ray sources at the solar disk for various source sizes, spectral indices and directivities of the primary source.} {We show that the albedo effect will alter the true source positions and substantially increase the measured source sizes. The source positions are shifted up to $\\sim 0.5''$ radially towards the disk centre and 5 arcsecond source sizes can be two times larger even for an isotropic source (minimum albedo effect) at 1~Mm above the photosphere. X-ray sources therefore should have minimum observed sizes, thus FWHM source size (2.35 times second-moment) will be as large as $\\sim 7''$ in the 20-50 keV range for a disk-centered point source at a height of 1~Mm ($\\sim 1.4''$) above the photosphere. The source size and position change is the largest for flatter primary X-ray spectra, stronger downward anisotropy, for sources closer to the solar disk centre, and between the energies of 30 and 50 keV.} {Albedo should be taken into account when X-ray footpoint positions, footpoint motions or source sizes from e.g. RHESSI or Yohkoh data are interpreted, and suggest that footpoint sources should be larger in X-rays than in optical or EUV ranges.} ", "introduction": " ", "conclusions": "" }, "1003/1003.2553_arXiv.txt": { "abstract": "The origin of the long secondary periods (LSPs) in red variables remains a mystery up to now, although there exist many models. The light curves of some LSPs stars mimic an eclipsing binary with a pulsating red giant component. To test this hypothesis, the observational data of two LSP variable red giants, 77.7795.29 and 77.8031.42, discovered by the MACHO project from the LMC, are collected and analyzed. The probable eclipsing features of the light curves are simulated by the Wilson-Devinney (W-D) method. The simulation yields a contact and a semidetached geometry for the two systems, respectively. In addition, the pulsation constant of the main pulsating component in each binary system is derived. By combining the results of the binary model and the pulsation component, we investigate the feasibility of the pulsating binary model. It is found that the radial velocity curve expected from the binary model has a much larger amplitude than the observed one and a period double the observed one. Furthermore, the masses of the components based on the density derived from the binary orbit solution are too low to be compatible with both the evolutionary stage and the high luminosity. Although the pulsation mode identified by the pulsation constant which is dependent on the density from the binary-model is consistent with the first or second overtone radial pulsation, we conclude that the pulsating binary model is a defective model for the LSP. ", "introduction": "Among the variable red giant stars, one sub-type exhibits long secondary periods (LSPs). The light curves of these stars exhibit not only a short primary period but also a long secondary period, which is approximately nine times longer than the short one. This phenomenon has been known for several decades \\citep{pay54,hou63}. Some samples of these LSPs variables are shown in \\citet{kis99}. An interest in the stars with the LSPs has been renewed by the study of \\citet{woo99}. This paper shows that in the LMC, $\\sim$25~\\% of all variable asymptotic giant branch (AGB) stars show LSP roughly nine times longer than the short primary period which is typically $\\sim$ 30--200 days. Meanwhile, a study of the bright local pulsating red giants indicates that at least one-third of these stars exhibit LSPs \\citep{per04}. Soszynski also gives $\\sim$30\\% of pulsating red giants in the LMC with LSPs \\citep{sos07b}. In the period-luminosity (P-L) diagram, it is interesting to see that the LSP variables follow a distinct sequence (sequence D), which is roughly parallel to the radial pulsation sequences A, B, and C for variable red giants. The LSPs present in variable red giants have attracted a lot of attention since their discovery, but their origin still remains mysterious. Since the LSPs are several times longer than the fundamental radial periods, they could not be caused by normal radial pulsations. Moreover, \\citet{woo99} and \\citet{woo04} note that the LSPs can not be explained by $g^{+}$ mode for the oscillatory $g ^{+}$ mode is evanescent in convective region and it is unlikely to be observable in a red giant because of its convective envelope. The $g^{-}$ mode is dynamically unstable in the convection envelope, and unable to lead to any oscillation. Regarding the nonradial $p$ modes, their periods are rather shorter than those of the fundamental radial modes, so they can not explain the LSPs either. Alternatively, if the LSPs are caused by some strange modes, they would be extremely damped and should not be seen \\citep{woo00}. One more possible explanation is the rotating spheroid model, which can explain the shape of the velocity curve, but there is no reason for the rotating period to bring about the observed P-L relation. It seems that variable red giants with long secondary periods (hereinafter referred to as LSPVs) can not be easily interpreted as pulsating red giants. In this situation, a hypothesis of binarity arose. It has been suggested that the sequence D stars could be components of close binary systems, and the LSPs could be interpreted as the light variations caused by ellipsoidal binary motions or eclipses. \\citet{sos04,sos07b} find that the sequence D stars overlap with and have a direct continuation of the sequence E stars that are mostly confirmed to be binaries. Radial velocity variations on the time-scale of LSPs have also been measured in a number of sequence D stars, such radial velocity variations being a requirement if the sequence D stars are binaries \\citep{hin02,oli03,woo04,nic09}. Some observational arguments supporting the binary hypothesis have already been reviewed by \\citet{sos07a}. After examining the light curves of LSPVs collected from the MACHO database, we note that many of them show some nearly regular and stable, eclipse-like light variations at the long secondary periods with a large amplitude in comparison with that of the primary pulsation. It seems that the LSPs can be easily interpreted as eclipses by an orbiting component. If this can be shown to be the case, it would be direct evidence to support the binary hypothesis for the sequence D stars. To test this idea, we propose a model -- pulsating binary, i.e. a binary system with a pulsating component. We begin with some very probable eclipsing LSPs candidates whose light curves look like eclipsing binary with primary and secondary minima. The light curves of two LSP stars are collected and analyzed by using the Wilson-Devinney ( W-D, hereafter) code and the power-spectrum method. Afterwards, the proposed eclipsing and pulsating nature as well as their evolutionary properties are discussed. ", "conclusions": "The origin of the LSPs is unknown and there exist many explanations. The light curves of some LSPVs are eclipse-like and it seems that this phenomenon is due to an invisible component orbiting around the pulsating red giant. To test this hypothesis, we propose a model -- pulsating binary, and select two LSP stars to analyze their orbit motions and pulsation nature. On the assumption of binarity, we simulate the photometric light curves of the systems by using the W-D method. The photometric solutions give us a configuration for the binary system: a contact system for 77.7795.29 and a semidetached system for 77.8031.42 with its Roche lobe fully filled. It means that the LSP star may have strong interaction and very probably mass transfer with the other component via Roche lobe overflow. However, Roche lobe overflow will rule out the binary hypothesis if we accept the view of \\citet{woo04}, which argues that the mass transfer of Roche lobe will result in a short merger timescale of about 1000 yrs. Moreover, the calculated effective temperatures and the ``$L_{1}/(L_{1}+L_{2})$'' values in Table 1 suggest that the secondary star is also a red giant with a similar temperature and luminosity to that of the primary red giant. This would lead to a double-lined spectroscopic binary. However, no observer of radial velocities in these systems has reported seeing spectral lines from the secondary star \\citep{hin02,oli03,woo04,nic09}. There are also some serious problems about the synthesized velocity curves. The simulation produces one cycle of velocity curve for two cycles of the light curve, while the observed velocities show one cycle of the radial velocity curve for one cycle of the light curve. The synthesized full velocity amplitudes for the star 77.7797.29 and 77.8031.42 are much greater than the typical values of other LSPVs. The full amplitudes of the secondary components are even larger and would lead to great velocity separations between the primary and the secondary stars, which have never been found by spectral observations. Fourier analysis is applied to investigate the intrinsic oscillation of the LSPVs over the raw data sets. Using the parameters ($R_{1}/A$ and $q$) obtained from the W-D code, the mean densities of the LSPVs are deduced and they are consistent with the red giant phase. Then the pulsation constants are obtained by using the classical equation $Q=P_{\\rm{pul}}(\\rho_{1}/\\rho_{\\sun})^{1/2}$ to identify the pulsation mode. For star 77.7795.29, one pulsating frequency is detected and it is caused by the first overtone radial pulsation, and for star 77.8031.42, the only pulsating frequency is caused by the second overtone radial pulsation. These agree with the theoretical value for red giants and the conclusion of \\citet{woo99}. The stellar properties of LSPVs are derived by using the information from the pulsating binary model. We calculate the bolometric magnitude, luminosity, radius and mass of the primary star, and find some of them conflict very seriously with the evolutionary properties of red giants. In particular, the masses for the star 77.7795.29 and 77.8031.42 are both less than 0.4$M_{\\sun}$, and their luminosities are too high for such low masses. It is difficult to imagine how such low mass stars could have such large luminosities. The situation gets even worse if we apply the mass ratio and the value of ``$L_{1}/(L_{1}+L_{2})$'' to calculate the mass and luminosity of the secondary star. Therefore, the radial velocities and the masses computed from the pulsating binary model do not agree with some observations and facts about red giants. We conclude that the model ``pulsating binary'' has some deficiencies in dealing with the observed properties of LSPVs and that the binary hypothesis for explaining the LSPs seems unreasonable." }, "1003/1003.0786.txt": { "abstract": "{ %About 40 years after the identification of the optical counterpart of the Crab pulsar, only a few rotation-powered pulsars have been identified in the optical domain. Of these, only four have spin-down ages of $\\la10\\,000$ years: the Crab pulsar, PSR\\, B1509$-$58, PSR\\, B0540$-$69 in the Large Magellanic Cloud (LMC), and the Vela pulsar. The study of the younger, and brighter, pulsars is important to understand the optical emission properties of isolated neutron stars through observations which, even in the 10m-class telescope era, are much more challenging for older and fainter objects. PSR\\, B0540$-$69, the second brightest ($V\\sim 22$) optical pulsar, is obviously a very interesting target for these investigations.} % {The aim of this work is threefold: %. Firstly, we aim at constraining the pulsar proper motion and its velocity on the plane of the sky, %and at improving the determination of the pulsar coordinates through optical astrometry. Secondly, we aim at obtaining a more precise characterisation of the pulsar optical spectral energy distribution (SED) through a consistent set of multi-band, high-resolution, imaging photometry observations %and at studying the relation with the X-ray spectrum, including the presence of a spectral turnover between the two bands. Last, we aim at measuring the pulsar optical phase-averaged linear polarisation, for which only a preliminary and uncertain measurement was obtained so far from ground-based observations %, and at testing the predictions of different neutron star magnetosphere models .} % {We performed high-resolution observations of \\psr\\ with the \\wfpcn\\ (\\wfpc) aboard the \\hstn\\ (\\hst), in both direct imaging and polarimetry modes.} % {From multi-epoch astrometry we set a $3 \\sigma$ upper limit of 1 mas yr$^{-1}$ on the pulsar proper motion, implying a transverse velocity $<250$ km s$^{-1}$ at the 50 kpc LMC distance. Moreover, we determined the pulsar absolute position with an unprecedented accuracy of 70 mas. From multi-band photometry we characterised the pulsar power-law spectrum and we derived the most accurate measurement of the spectral index ($\\alpha_{O} = 0.70 \\pm 0.07$) which indicates a spectral turnover between the optical and X-ray bands. Finally, from polarimetry we obtained a new measurement of the pulsar phase-averaged polarisation degree ($PD =16\\%\\pm4\\%$), consistent with magnetosphere models depending on the actual intrinsic polarisation degree and depolarisation factor, and we found that the polarisation vector ($22^{\\circ} \\pm 12^{\\circ}$ position angle) is possibly aligned with the semi-major axis of the pulsar-wind nebula and with the apparent proper motion direction of its bright emission knot.} % {By using the \\wfpc\\ on the \\hst, we performed a comprehensive optical study (astrometry, photometry, and polarimetry) of \\psr. Deeper studies with the \\hst\\ can only be possible with the refurbished \\acsn\\ (\\acs) and with the new {\\em Wide Field Camera 3} ({\\em WFC3}).} ", "introduction": "The pulsar \\psr\\ in the Large Magellanic Cloud (LMC) is often referred to as the Crab pulsar ``twin'' because of it is very similar in age ($\\sim 1700$ years), spin period ($P= 50$ ms), and rotational energy loss ($\\dot {E} \\sim 10^{38}$ erg~s$^{-1}$). \\psr\\ is the second pulsar discovered in X-rays by the {\\em Einstein} Observatory (Seward et al. 1984) and the first extragalactic pulsar detected at any wavelength. Like the Crab, \\psr\\ is also embedded in a bright pulsar wind nebula (PWN) visible at wavelengths from the optical to the soft/hard X-rays (e.g. De Luca et al. 2007; Petre et al. 2007; S\\l{}owikowska et al. 2007). After its discovery, \\psr\\ has been observed by nearly all X-ray satellites, {\\em EXOSAT} (\\\"Ogelman \\& Hasinger 1990), {\\em Ginga} (Nagase et al. 1990), {\\em ROSAT} (Finley et al. 1993), {\\em BeppoSax} (Mineo et al. 1999), {\\em Chandra} (Kaaret et al. 2001), {\\em Rossi-XTE} (de Plaa et al. 2003), {\\em ASCA} (Hirayama et al. 2002), {\\em Integral} (S\\l{}owikowska et al. 2007), and {\\em Swift} (Campana et al. 2008). Its X-ray light curve is characterised by a single, broad peak, very much at variance with, e.g. that of the Crab pulsar. In radio, the distance to the LMC made \\psr\\ undetectable for a long time until pulsations were finally detected from Parkes (Manchester et al. 1993). Giant radio pulses were discovered by Johnston \\& Romani (2003), aligned in phase with the peak of the X-ray pulse (Johnston et al. 2004), making \\psr\\ the second youngest radio pulsar to feature this phenomenon. \\psr\\ is also one of the handful of rotation-powered pulsars with a measured braking index, obtained from X-ray observations (e.g., Zhang et al. 2001; Cusumano et al. 2003; Livingstone et al. 2005). In the optical, \\psr\\ is the second brightest isolated neutron star ($V\\sim 22$) identified so far (see Mignani 2009a,b for recent reviews). Optical pulsations were detected by Middleditch \\& Pennypacker (1985) soon after the X-ray discovery, making \\psr\\ the third optical pulsar after the Crab (Cocke et al. 1969) and Vela pulsars (Wallace et al. 1977). However, it was only through high-resolution imaging observations with the ESO {\\em New Technology Telescope} (\\ntt) that its optical counterpart was actually identified (Caraveo et al. 1992; Shearer et al. 1994). The \\psr\\ optical light curve (Middleditch et al. 1987; Gouiffes et al. 1992; Boyd et al. 1995; Gradari et al. 2009) is characterized by a single broad peak, very similar to the X-ray and radio light curves, with a significant dip on top. The optical spectral energy distribution (SED) of \\psr\\ was first measured by Middleditch et al. (1987) from high-speed multi-band photometry which suggested the presence of a possible excess in the U band with respect to an otherwise monotonic power-law continuum ($F_{\\nu} \\propto \\nu^{-\\alpha_{O}}$). A re-analysis of the same multi-band photometry measurements, however, did not yield evidence of the claimed U-band excess (Nasuti et al. 1997). The power-law spectrum demonstrates that the optical emission from \\psr\\ is of magnetospheric origin, with an optical luminosity consistent with the expectations of the Pacini \\& Salvati (1987) model. Low-resolution spectra of \\psr\\ were obtained by Hill et al. (1997), with the \\hstn\\ (\\hst), and by Serafimovich et al. (2004), with the \\vltn\\ (\\vlt), although the latter was affected by the contamination from the supernova remnant. By using multi-band imaging photometry of \\psr\\ from archival \\hst\\ observations, Serafimovich et al. (2004) confirmed that the optical spectrum is dominated by a power-law continuum, although with a spectral index different from the values previously published. They also reported a tentative proper motion measurement for \\psr, which, however, was not confirmed by De Luca et al. (2007) using a large time baseline of \\hst\\ observations. Phase-resolved polarimetry observations of \\psr\\ were performed by Middleditch et al. (1987) but only yielded an upper limit on the phase-averaged polarisation degree, while image polarimetry observations of Chanan \\& Helfand (1990) focused on the PWN only. More recently, image polarimetry observations of \\psr\\ were performed by Wagner \\& Seifert (2000) using the \\vlt. They reported a phase-averaged polarisation degree of $\\approx$ 5\\% (with no associated error) which was probably affected by the contribution of the unresolved PWN. Thus, only preliminary, or uncertain, phase-averaged optical polarisation measurements exist for \\psr. In this paper, we report the results of new astrometry, photometry, and polarimetry measurements of \\psr\\ performed with the \\hst\\ as a part of a dedicated programme aimed at the comprehensive study of the pulsar and of its PWN (Mignani et al., in preparation). From the pulsar astrometry we will constrain its proper motion and its velocity on the plane of the sky, and we will provide an update reference position for future observations. From the pulsar photometry we will measure anew the optical SED, we will study the relation with the X-ray spectrum, and we will verify the presence of a spectral turnover between the two energy bands. From the pulsar polarimetry we will measure its phase-averaged polarisation and we will test the predictions of different neutron star magnetosphere models. The paper is divided as follows: observations and data analysis are described in Sect. 2, while the results are presented and discussed in Sect. 3 and 4, respectively. ", "conclusions": "\\subsection{The pulsar astrometry} Unfortunately, the lack of a measurable proper motion for \\psr\\ does not allow to search for possible connections between the pulsar kinematics, its polarisation properties, and the PWN morphology, which have been found, e.g. for the Crab and Vela pulsars (see Mignani 2009c and references therein). The computed $3 \\sigma$ upper limit of 1 mas yr$^{-1}$ (see Sect. 3.1) on the \\psr\\ proper motion sets a corresponding upper limit of $\\sim 250$ km s$^{-1}$ on its transverse velocity. Although this value is lower than the peak of the pulsar transverse velocity distribution (e.g., Hobbs et al. 2005), it does not allow to claim a peculiarly low velocity for \\psr\\ unless one lower the constraints on the measured upper limit to $1 \\sigma$. Indeed, several pulsars feature transverse velocities lower than $100$ km s$^{-1}$ including, e.g. the Vela pulsar for which a velocity of $\\sim 65$ km~s$^{-1}$ has been inferred from \\hst\\ astrometry (Caraveo et al. 2001). \\begin{table} \\begin{center} \\caption{Compilation of the \\psr\\ coordinates available from the literature and associated errors $\\Delta \\alpha$ and $\\Delta \\delta$. } \\begin{tabular}{llllr} \\\\ \\hline $\\alpha_{J2000} ^{(hms)}$ & $\\Delta \\alpha ^{(s)}$ & $\\delta_{J2000}^{(\\circ ~'~\")}$ & $\\Delta \\delta ^{(\")}$ & Source \\\\ \\\\\\hline 05 40 11.03 & 0.40 & -69 19 57.5 & 2.0 & X-rays (1) \\\\ 05 40 11.16 & 0.15 & -69 19 57.79 & 0.90 & Optical (2) \\\\ 05 40 10.980 & 0.090 & -69 19 55.17 & 0.50 & Optical (3) \\\\ 05 40 11.221 & 0.090 & -69 19 54.98 & 0.50 & X-rays (4) \\\\ 05 40 11.173 & 0.120 & -69 19 54.41 & 0.70 & Optical (5) \\\\ 05 40 11.160 & 0.040 & -69 19 53.90 & 0.20 & X-rays (6) \\\\ \\hline 05 40 11.202 & 0.009 & -69 19 54.17 & 0.05 & Optical (7) \\\\ \\hline \\label{coo} \\end{tabular} \\end{center} (1) Seward et al. (1984); (2) Manchester \\& Peterson (1989); (3) Caraveo et al. (1992); (4) Kaaret et al. (2001); (5) Serafimovich et al. (2004); (6) Livingstone et al. (2005); (7) this work \\end{table} The computed average coordinates of \\psr, together with their associated errors, are listed in Table \\ref{coo}, for comparison with those obtained from previous works. After its discovery by {\\em Einstein} (Seward et al. 1984), the coordinates of \\psr\\ were firstly revised by Manchester \\& Peterson (1989), from optical timing observations, and later by Caraveo et al. (1992), from the astrometry of the optical counterpart identified in their \\ntt\\ images. We note that the discovery of the pulsar in radio (Manchester et al. 1993) did not allow to obtain more precise coordinates. Indeed, its very low radio flux density at the LMC distance, $S_{400}= 0.7$ $\\mu$Jy, makes it difficult to obtain an accurate radio timing position. Thus, the optical coordinates of Caraveo et al. (1992) were thereafter assumed as a reference. In particular, they were used to compute the pulsar X-ray and radio timing solution by, e.g. Zhang et al. (2001), Cusumano et al. (2003) but, surprisingly enough, not by Finley et al. (1993) and Manchester et al. (1993). \\chan\\ observations (Kaaret et al. 2001) yielded an updated X-ray position. The former, and actually not the latter as the authors claim, was used as a reference by Johnston et al. (2004) for the pulsar X-ray/radio timing solution, and by de Plaa et al. (2003). A new optical position was obtained by Serafimovich et al. (2004) from \\hst\\ data, while a new X-ray position was derived by Livingstone et al. (2005) based on {\\em Rossi-XTE} timing observations. \\begin{figure} \\centering \\includegraphics[height=8.5cm,angle=0,clip]{13870fig1_small.ps} \\caption{$15\\arcsec \\times 15\\arcsec$ region of the \\psr\\ field observed with the \\hst/\\wfpc\\ in the F555W filter (October 1995 observation). North to the top, East to the left. The crosses mark the different positions of the pulsar obtained from previous X-ray and optical observations (black) and from our \\wfpc\\ observation (white). The cross arms correspond to the $1 \\sigma$ error in each coordinate. Different positions are numbered according to the reference publications (see Table \\ref{coo}). The pulsar is indicated by the arrow for clarity. } \\label{astro} % Give a unique label \\end{figure} To visualise the difference and the improvement on the determination of the \\psr\\ position, we plotted all coordinates from Table \\ref{coo} on the October 1995 \\wfpc\\ F555W image of the field (Fig. \\ref{astro}). We see that the original pulsar coordinates from Seward et al. (1984) and from Manchester \\& Peterson (1989) fall $\\approx 3\\farcs5$ away from the pulsar position and, actually, out of the nebula shell. On the other hand, the coordinates of Caraveo et al. (1992) and Kaaret et al. (2001) are indeed within the nebula but they are clearly offset from the actual pulsar position, while those of Serafimovich et al. (2004) and Livingstone et al. (2005) are more consistent with the pulsar position, although the former have much larger errors. Thus, thanks to the sharp angular resolution of the \\wfpc\\ and to the high astrometric accuracy of \\tmass, our coordinates are the most accurate obtained so far, providing a factor of $\\ga 4$ improvement with respect to the most recently published values. In particular, our \\wfpc\\ coordinates supersede both the optical coordinates of Caraveo et al. (1992), which were obtained using images with lower spatial resolution (0\\farcs13/pixel) and the {\\em GSC 1.0} ($\\sigma_{s}\\sim1\\farcs0$; Lasker et al. 1990) as a reference catalogue, and those of Serafimovich et al. (2004), which were obtained using an early release of the {\\em GSC 2.0}, as well as the X-ray coordinates of Kaaret et al. (2001) and Livingstone et al. (2005), obtained through \\chan\\ imaging and {\\em Rossi-XTE} timing observations, respectively. The better determination of the \\psr\\ absolute position is important for follow-up non-imaging observations. In the case of, e.g. spectroscopy, it will make it possible to accurately centre the pulsar in the slit when blind presets are made necessary, e.g. for ground-based observations due to the difficulties in resolving the pulsar through lower spatial resolution and seeing-dominated acquisition images. As a consequence, this will allow to use narrower slits without causing any loss of signal from the pulsar. This is important to minimise the background contamination from the supernova remnant which hampered ground-based spectroscopic observations performed so far (see, e.g. Serafimovich et al. 2004). \\subsection{The pulsar spectrum} \\begin{figure} \\centering \\includegraphics[height=8.5cm,bb=20 180 560 720, angle=0,clip]{13870fig2.ps} \\caption{Optical spectral energy distribution of \\psr\\ derived from the available multi-band \\wfpc\\ photometry (Table \\ref{data}). Points are labelled according to the filter names. The dashed line is to the best fit power-law spectrum. } \\label{spec} % Give a unique label \\end{figure} \\begin{figure} \\centering \\includegraphics[height=9cm,angle=270,clip]{13870fig3.ps} \\caption{Optical spectral energy distribution of \\psr\\ (points) compared with the power-law model (Kaaret et al. 2001) best-fitting the \\chan\\ X-ray spectrum (solid line) and its extrapolation in the optical domain. Dashed lines correspond to a $1 \\sigma$ uncertainty on the model parameters. } \\label{multispec} % Give a unique label \\end{figure} We used our multi-band \\wfpc\\ photometry of \\psr\\ to better investigate its optical SED. For consistency with previous works, we computed the interstellar extinction towards the pulsar using as a reference an $E(B-V)=0.2$ and $R=3.1$, which have been verified by several independent measurements (see Serafimovich et al. 2004 and references therein). We derived the extinction coefficients in the \\wfpc\\ filters from the extinction curves of Fitzpatrick (1999). The pulsar SED is shown in Fig. \\ref{spec}, after correction for the interstellar extinction. The plotted errors on the spectral fluxes are purely statistical and do not account for the systematic 0.02-0.04 magnitude error on the \\wfpc\\ zero points (Heyer et al. 2004). A linear fit to the data points yields a spectral index $\\alpha_{O}=0.70\\pm 0.07$. As seen, our fit clearly confirms that there is no evidence for the \"U-band excess\" originally claimed by Middleditch et al. (1985). We compared our best-fit power-law spectrum with that measured by Serafimovich et al. (2004). They derived a spectral index $\\alpha_{O} = 1.07^{+0.20}_{-0.19}$ which is somewhat steeper than that measured by us, although not formally incompatible. As we stated above, the power-law slope is only marginally affected by the difference in the assumed interstellar extinction correction. Indeed, applying the extinction coefficients of Cardelli et al. (1989) to our spectral fluxes, as done by Serafimovich et al. (2004), yields a power law slope which is substantially identical to that obtained using the extinction coefficients of Fitzpatrick (1999). Thus, the difference in the power-law slope is most likely intrinsic to the photometry measurements. First of all, Serafimovich et al. (2004) fitted a lower number of points (see Table \\ref{data}) and used the flux measured in the medium/narrow-band F547M and F658N filters, while we fitted a larger number of points and we only used the flux measured in wide-band filters. Moreover, while we used similar photometry apertures, we applied aperture corrections specific for the \\wfpc, which are wavelength-dependent (Holtzman et al. 1995), and we applied the CTE correction which, apparently, was not accounted for by Serafimovich et al. (2004). The updated value of the \\psr\\ optical spectral index is now more consistent with those of the other optically identified pulsars which all feature relatively steep spectral indexes ($\\alpha_{O} \\ga 0.4$), while only the Crab and Vela pulsar feature a nearly flat power-law spectrum (see Mignani et al. 2007a for a comparison). In particular, \\psr\\ may be the pulsar with the steepest optical spectral index. We also note that the new value of $\\alpha_{O}=0.70 \\pm 0.07$ is very close to that of the X-ray spectral index $\\alpha_{X}=0.83 \\pm 0.13$ measured with \\chan\\ (e.g., Kaaret et al. 2001). Interestingly, however, the optical spectrum is not the continuation of the X-ray one. Indeed, the \\hst\\ spectral fluxes lie about a factor of 100 below the low-energy extrapolation of the \\chan\\ 0.1-10 keV X-ray spectrum (see Fig. \\ref{multispec}). This difference can not be accounted for by possible problems in the subtraction of the PWN background or by other kinds of systematic effects, and confirms the presence of a spectral turnover between the optical and the X-ray bands (see also Fig. 14 of Serafimovich et al. 2004). Spectral turnovers between these two energy bands are a common feature of the magnetospheric emission of many other rotation-powered pulsars, as indicated by the differences in their optical and X-ray spectral indexes (e.g., Mignani et al. 2007a). For comparison, we plotted in Fig. \\ref{allspec} the optical, near ultraviolet (NUV), and near infrared (NIR) spectral fluxes for all rotation-powered pulsars with an optical counterpart, together with the model 0.1-10 keV X-ray spectra. In the case of the Crab and Vela pulsars, the spectral turnover is in the form of a single break in the optical-to-X-ray power-law spectrum. In the case of \\psr, instead, a double break is required to join the optical and the X-ray SEDs, unless the actual interstellar extinction is larger by a factor of 2 with respect to current best estimates, which would make the spectrum flatten in the blue. This apparent double break in the optical-to-X-ray power-law spectrum might be also present in other, much closer, rotation-powered pulsars. Of course, for some pulsars (especially for the fainter ones) the comparison is hampered by the paucity of spectral flux measurements at NIR/optical/NUV wavelengths, which makes it difficult to establish the presence of a power-law component. Bearing this caveat in mind, a possible double break can be recognised in the optical-to-X-ray power-law spectrum of the middle-aged pulsars PSR\\, B0656+14 ($\\approx$100 kyrs) and, to a lesser extent, Geminga ($\\approx$ 340 kyrs), for both of which a power-law tail is hinted in the NIR. On the other hand, a single break is probably required for the old pulsars PSR\\, B1929+10 ($\\approx 3$ Myrs) and PSR\\, B0950+08 ($\\approx$ 18 Myrs), while only for the young PSR\\, B1509$-$58 the optical/NIR power-law seems to nicely fit the extrapolation of the X-ray one. So, three different trends are recognised in our sample which spans four age decades. This means that the breaks in the power-law spectrum do not correlate with the neutron star age and, thus, are not indicative of any evolution in the emission processes in the neutron star magnetosphere. This is consistent with the lack of evidence for a power-law spectrum evolution, both in the optical (Mignani et al. 2007a) and in the X-rays (Becker 2009). The shape of the power-law optical-to-X-ray spectrum does not correlate with the dipole magnetic field either with, e.g. the Crab and PSR\\, B0950+08 both featuring a single spectral break but having magnetic fields different by two orders of magnitudes. It is possible that the number of observed spectral breaks is related to the geometry of the emission regions and to the particle distribution in the neutron star magnetosphere which, however, are difficult to reconstruct without having information on both the X-ray and optical light curve profiles. At the same time, is not possible to determine whether such double breaks are simply associated with a change in the power-law spectrum or they are associated with absorption processes in the neutron star magnetosphere, a possibility proposed by Serafimovich et al. (2004) for \\psr. Observations in the far UV (FUV) would be fundamental to clarify this issue. Unfortunately, only few pulsars have been observed in this spectral region by the {\\em EUVE} satellite (Korpela \\& Bowyer 1998) but only Geminga (Bignami et al. 1996) and PSR\\, B0656+14 (Edelstein et al. 2000) have been detected amongst the pulsars shown in Fig. \\ref{allspec}, with their FUV fluxes being more or less compatible with the extrapolations of the black body components fitting the optical spectra. \\begin{figure*}[ht] \\centering \\includegraphics[height=6.5cm,angle=270,clip]{13870fig4.ps} \\includegraphics[height=6.5cm,angle=270,clip]{13870fig5.ps} \\includegraphics[height=6.5cm,angle=270,clip]{13870fig6.ps} \\includegraphics[height=6.5cm,angle=270,clip]{13870fig7.ps} \\includegraphics[height=6.5cm,angle=270,clip]{13870fig8.ps} \\includegraphics[height=6.5cm,angle=270,clip]{13870fig9.ps} \\includegraphics[height=6.5cm,angle=270,clip]{13870fig10.ps} \\includegraphics[height=6.5cm,angle=270,clip]{13870fig11.ps} \\caption{Same as Fig. \\ref{multispec} but for all rotation-powered pulsars with an optical counterpart and flux measurements in at least two bands. Optical flux values are taken from the compilation in Mignani et al. (2007a). X-ray spectral models are taken from the source publications: Crab (Willingale et al. 2001), PSR\\, B1509$-$58 (Gaensler et al. 2002), Vela (Manzali et al. 2007), PSR\\, B0656+14 (De Luca et al. 2005), Geminga (Caraveo et al. 2004), PSR\\, B1929+10 (Becker et al. 2006), PSR\\, B0950+08 (Becker et al. 2004). Left panels: young pulsars ($\\tau < 10\\,000$ years). Right panels: middle-aged and old pulsars ($\\tau > 100\\,000$ years). In each column objects are sorted according to their spin-down age and labelled. Different X-ray spectral components are shown by the dotted red, green (black body), and blue (power-law) lines while the solid blue lines show the composite spectra. Only best-fits are plotted for clarity. } \\label{allspec} % Give a unique label \\end{figure*} \\subsection{The pulsar polarisation} Our value of the phase-averaged polarisation ($PD = 16\\%\\pm 4\\%$) is consistent with the upper limit of $15 \\%$ inferred by Middleditch et al. (1987) from phase-resolved polarimetry of the pulsar. However, the measured $PD$ is higher than $PD \\approx 5\\%$ obtained from \\vlt\\ phase-averaged polarisation observations by Wagner \\& Seifert (2000), for which, however, no error estimate is given, so that its significance can not be assessed. Thus, ours is the only significant measurement ($\\sim 4 \\sigma$ level) of the \\psr\\ phase-averaged polarisation obtained so far. Measurements of the phase-averaged polarisation degree have been obtained so far for all the young ($\\la 10\\,000$ years old) rotation-powered pulsars with an optical counterpart, albeit with a different degree of confidence (see Mignani et al. 2007b and S\\l{}owikowska et al. 2009 for a summary). For the Crab, a value of $PD=9.8\\% \\pm 0.1\\%$ was inferred from phase-resolved polarimetry observations (S\\l{}owikowska et al. 2009), while image polarimetry observations with the \\vlt\\ yielded $PD=9.4\\%\\pm 4\\%$ for the Vela pulsar (Mignani et al. 2007b) and $PD=10.4\\%$ for PSR\\, B1509$-$58 (Wagner \\& Seifert 2000), with the latter measurement being admittedly very uncertain and quoted with no error. The phase-averaged polarisation of \\psr, $PD = 16 \\%\\pm 4\\%$, is thus consistent with the measurements obtained for the other young pulsars, as one might expect from their similar parameters, like spin down luminosity and/or magnetic field strength. Although the current data base is still extremely limited, with the measurement obtained for PSR\\, B1509$-$58 in wait for confirmation and with the phase-resolved polarimetry observations of the middle-aged ($\\sim 100$ kyrs) PSR\\, B0656+14 (Kern et al. 2003) covering only one third of the period, it nonetheless suggests that the phase-averaged optical polarisation degree of rotation-powered pulsars is typically around $10\\%$, with our value of $PD$ for PSR\\, B0540$-$69 being possibly somewhat larger. On the theoretical side, the comparison of the observed phase-averaged polarisation degrees with the predictions of different pulsar magnetosphere models is complicated both by the degree of complexity of such models and by the limits in model simulations. However, for both the Crab and Vela pulsars, a detailed comparison with different pulsar magnetosphere models, showed that the measured phase-averaged polarisation degree of the emerging optical radiation requires rather low degree of polarisation intrinsic to the source and/or strong depolarisation factors (S\\l{}owikowska et al. 2009; Mignani et al. 2007b). Recent 3D outer-gap model calculations of optical and/or high-energy radiation with detailed incorporation of electron-positron gyration have been carried out so far for the Crab pulsar only (Takata et al. 2007). In general, the results seem to fit the measured optical polarization level (see S\\l{}owikowska et al. 2009) for some particular values of the viewing angle (the angle between the line of sight and the spin axis). Other models addressing the problem of polarisation characteristics are relatively more simplified in terms of the model assumptions in use (Dyks et al. 2004; Petri \\& Kirk 2005). All models mentioned above - relevant for young energetic pulsars - rely on spatially extended sources of emission (either within the magnetosphere - Dyks et al 2004; Takata et al. 2007) or in the pulsar wind zone (Petri \\& Kirk 2005). As a consequence, the depolarisation effects due to rotation, photon finite time of flight, and magnetic field structure are significant in all three models. This might explain the somewhat larger polarisation degree of PSR\\, B0540$-$69. However, one should keep in mind that the shape of its optical light curve makes it more difficult to infer the geometry of the emission region with respect to the Crab and Vela pulsars, which provides crucial input parameters for model simulations. Interestingly, for both the Crab and Vela pulsars the polarisation position angle features a remarkable alignment with the axis of symmetry of the X-ray structures (torus and jets) observed by \\chan, with the pulsar spin axis, and with the proper motion vector (S\\l{}owikowska et al. 2009; Mignani et al. 2007b). Unfortunately, for \\psr\\ the scenario is less clear since no proper motion has been measured so far and \\chan\\ observations (Gotthelf \\& Wang 2000; Petre et al. 2007) could only partially resolve the morphology of the PWN, although its clear asymmetry might hint at the existence of a torus and, possibly, of a jet. A similar PWN morphology was also observed in the optical from \\hst/\\wfpc\\ observations (Caraveo et al. 2000). Most noticeably, the \\wfpc\\ images shows the existence of a bright emission knot south west of the pulsar %, likely identified with the PWN \"southern maximum\" of Caraveo et al. (1992), which is apparently moving at a speed of $\\sim 0.04~c$ (De Luca et al. 2007) and which can be interpreted as the head of a possible jet emitted by the pulsar. We note that the orientation of the pulsar polarisation vector ($PA =22^{\\circ} \\pm 12^{\\circ}$) is consistent, accounting for possible projection effects, with that of the semi-major axis of the PWN (see Fig. \\ref{pol}). Moreover, the pulsar phase-averaged polarisation vector is also possibly aligned with the apparent direction of motion of the knot ($\\approx 230^{\\circ}$ east from north)\\footnote{We note that accounting for the proper motion of the LMC (Costa et al. 2009) and for the galactic rotation does not significantly affect this apparent alignment.}. We estimated the chance alignment probability between the two vectors to be $\\sim 0.04$, small but still not negligible. If real, however, this peculiar alignment would indicate a possible physical connection between the pulsar and the knot, as already proposed in De Luca et al. (2007), and would support the pulsar/jet scenario. The measurement of both the knot and of the PWN polarisation structure (Mignani et al. in preparation) would further reinforce this scenario. Assuming for \\psr\\ the same PWN geometry as the Crab and Vela pulsars, the direction of the jet would suggest that a putative torus would indeed extend along the semi-minor axis of the nebula, and not along the semi-major one, as previously hypothesised from \\chan\\ and \\hst\\ observations (Gotthelf \\& Wang 2000; Caraveo et al. 2000). \\begin{figure}[h] \\centering \\includegraphics[height=8.5cm,angle=0,clip]{13870fig12_small.ps} \\caption{Same as in Fig. \\ref{astro} but zoomed over a $8\\arcsec \\times 8\\arcsec$ region. \\psr\\ is labelled. The image has been smoothed over $3\\times3$ pixel cells for a better visualisation of the extended PWN emission. Isophotal contours (linear spacing) are drawn as solid lines. The thick lines show the measured $1 \\sigma$ uncertainty ($\\pm 12^{\\circ}$) on the pulsar polarisation position angle, while the thin arrows show the assumed uncertainty ($\\pm 30^{\\circ}$) on the apparent proper motion direction of the bright emission knot in the PWN (De Luca et al. 2007). } \\label{pol} % Give a unique label \\end{figure}" }, "1003/1003.0626_arXiv.txt": { "abstract": "{The precise and accurate modelling of a terrestrial planet like Venus is an exciting and challenging topic, all the more interesting since it can be compared with that of the Earth for which such a modelling has already been achieved at the milliarcsecond level} {We want to complete a previous study (Cottereau and Souchay, 2009), by determining at the milliarcsecond level the polhody, i.e. the torque-free motion of the axis of angular momentum of a rigid Venus in a body-fixed frame, as well as the nutation of its third axis of figure in space, which is fundamental from an observational point of view. } {We use the same theoretical framework as Kinoshita (1977) did to determine the precession-nutation motion of a rigid Earth. It is based on a representation of the rotation of rigid Venus, with the help of Andoyer variables and a set of canonical equations in Hamiltonian formalism} {In a first part we have computed the polhody, i.e. the respective free rotational motion of the axis of angular momentum of Venus with respect to a body-fixed frame. We have shown that this motion is highly elliptical, with a very long period of 525 cy to be compared with 430 d for the Earth. This is due to the very small dynamical flattening of Venus in comparison with our planet. In a second part we have computed precisely the Oppolzer terms which allow to represent the motion in space of the third Venus figure axis with respect to Venus angular momentum axis, under the influence of the solar gravitational torque. We have determined the corresponding tables of coefficients of nutation of the third figure axis both in longitude and in obliquity due to the Sun, which are of the same order of amplitude as for the Earth. We have shown that the coefficients of nutation for the third figure axis are significantly different from those of the angular momentum axis on the contrary of the Earth. Our analytical results have been validated by a numerical integration which revealed the indirect planetary effects.} {This paper is a complementary study of Cottereau and Souchay (2009).It gives a precise determination both of the torque free motion of Venus, as well as the nutation of the third Venus figure axis in space for a short time scale, when considering the planet as a rigid body.} ", "introduction": "Venus which can be considered as the twin sister of the Earth, in view of its global characteristics (size, mass, density), has been the subject of a good amount of investigations on very long time scales, to understand its slow retrograde rotation (243 d) and its rather small obliquity ($2^{\\circ}.63$) (Goldstein 1964; Carpenter 1964; Goldreich and Peale 1970; Lago and Cazenave 1979; Dobrovoskis 1980; Yoder 1995; Correia and Laskar 2001, 2003). Habibullin (1995) made an analytical study on the rotation of a rigid Venus. In Cottereau and Souchay (2009) we presented an alternative study, from a theoretical framework already used by Kinoshita (1977) for the rigid Earth. We made an accurate description of the motion of rotation of Venus at short time scale. We calculated the ecliptic coordinates of Venus orbital pole and the reference point $\\gamma_{0V}$ which is the equivalent of the vernal equinox for Venus. Our value for the precession in longitude was $\\dot{\\Psi}=4474\".35$t/cy$\\pm 66.5$. We have performed a full calculation of the coefficients of nutation of Venus and presented the complete tables of nutation in longitude $\\Delta \\Psi$ and obliquity $\\Delta \\epsilon$, for the axis of angular momentum due to both the dynamical flattening and triaxiality of the planet. In this paper, the study begun in Cottereau and Souchay (2009) is completed. First in section \\ref{torque} we consider the torque free rotational motion of a rigid Venus. We recall the parametrization of Kinoshita (1977) and the important equations of Kinoshita (1972) which are used to solve this torque free motion. The important characteristics (amplitude, period, trajectory) of the free motion are given . Cottereau and Souchay (2009) supposed that the relative angular distances between the three poles (of angular momentum, figure and rotation) are very small as it is the case for the Earth. In this paper we want to determine accurately the motion of the third Venus figure axis which is the fundamental one in an observational point of view. To do that we reject the hypothesis of coincidence of the poles. Thus we determine the Oppolzer terms depending on the dynamical flattening and the trixiality of Venus. Then we compare these terms with the corresponding terms of nutation for the axis of angular momentum, as determined by Cottereau and Souchay (2009) and the Oppolzer terms determined by Kinoshita (1977) for the Earth in section \\ref{section3}. We give the complete tables of the coefficients of nutation of the third figure axis of Venus. We compare them, with the coefficients of nutation of the angular momentum axis, taken in Cottereau and Souchay (2009) (section \\ref{section4}). Finally in section \\ref{section6} we determine the nutation of the angular momentum axis by numerical integration using the ephemeris DE405. We validate the analytical results of Cottereau and Souchay (2009) down to a precision of the order of a relative $10^{-5}$. We show that the discrepancies between the numerical integration and analytical results (Cottereau and Souchay, 2009) are caused by the indirect planetary effect. i.e the small contribution to the nutation, which is due to the periodic oscillations of the orbital motion of Venus. In this paper as it was the case in Cottereau and Souchay (2009), our domain of validity is roughly 3000 years. ", "conclusions": "In this paper we achieved the accurate study of the rotation of Venus, for a rigid model and on a short time scale begun by Cottereau and Souchay (2009), by applying analytical formalisms already used for the rigid Earth (Kinoshita, 1972, 1977). The differences between the rotational characteristics of Venus and our planet, due to the slow rotation of Venus and its small obliquity, have been highlighted. Firstly we have precisely determined the polhody, i.e the torque free rotational motion for a rigid Venus. We have adopted the theory used by Kinoshita (1972, 1992) and we have given the parametrization and the equations of motion to solve the motion. We have shown that the polhody is significantly elliptic, on the contrary of the Earth for which it can be considered as circular in first approximation. Moreover it is considerably slower. Indeed, the period of the torque free motion is $525.81$ cy for Venus whereas it is $303$ d for our planet, when considered as rigid. Then we have determined the motion of the third figure axis, which is fundamental from an observational point of view. We have calculated the Oppolzer terms due to the gravitational action of the Sun using the equation of motion of Kinoshita (1977) as well as the corresponding development of the disturbing functions. We have compared them with the coefficients of nutation for the angular momentum axis taken in Cottereau and Souchay (2009). One of the important results is that these Oppolzer terms depending on the dynamical flattening are of the same order of amplitude as the coefficients of nutation themselves, whereas for the Earth (Woolard, 1953; Kinoshita, 1977) these Oppolzer terms are very small with respect to the coefficients of nutation of the angular momentum axis. Moreover we have computed the Oppolzer terms depending on the triaxiality, which is not done in Kinoshita (1977) in the case of the Earth for which they are negligible. In the case of Venus,these Oppolzer terms are significant. Even larger than the corresponding coefficients of nutation of the angular momentum axis. Thanks to our Oppolzer terms we have also been able to give the tables of nutation of the third figure axis from which we computed the nutation for a 4000 d time span. The comparison with the nutation of the angular momentum axis, calculated from Cottereau and Souchay (2009), is also given. The nutation of the third figure axis is significantly smaller peak to peak than the nutation of the angular momentum axis in longitude, and less important in obliquity. The amplitude of the largest nutation coefficient in longitude of the third figure axis (1.5\") is half the one of the angular momentum axis (3\"). The amplitude of the nutation in obliquity of the third figure axis is 0\".18 peak to peak whereas the amplitude of the angular momentum axis is 0\".25. The nutations of the third figure axis, in obliquity and in longitude, are dominated by three sinusoids associated with the arguments $2L_{s}$, $M$ and $2\\Phi$, with respective periods 112.35 d, 224.70 d and 121.51 d. The nutation of the angular momentum axis is dominated by two sinusoids with the argument $2L_{s}$ and $2\\Phi$. Our results have shown that although the axis of angular momentum and the third figure axis can be considered identical in the case of the Earth (Kinoshita, 1977), this approximation does not hold in the case of a slowly rotating planet as Venus. At last we have validated our analytical results down to a relative accuracy of $10^{-5}$ with a numerical integration. Moreover, we have confirmed, by using results in Souchay and Kinoshita (1996) for the nutation of a rigid Earth, that the differences between our analytical computation and our numerical integration are essentially due to the indirect planetary effects, which was not taken into account by Cottereau and Souchay (2009). We think that this study is fundamental to understand the behavior of Venus rotation in a very accurate and exhaustive way for short time scales, and should be a necessary starting point to another similar study including non rigid effects (elasticity, atmospheric forcing etc...)." }, "1003/1003.2415_arXiv.txt": { "abstract": "This guide describes the functionality and use of the \\aipcls{} by explaining its extensions and restrictions compared to the \\texttt{article} class of standard \\LaTeX. It is not a manual to be used on its own but should be used together with an introductory manual on \\LaTeX{} such as \\cite{A-W:LLa94}. ", "introduction": "The \\aipcls{} is a \\LaTeXe{} document class for conference proceedings of the American Institute of Physics and other documents with similar layout requirements. Your file will be used to reproduce your paper as is, the only modifications done by the publisher are adding appropriate page numbers and the copyright line. It is therefore essential that you embed all fonts when saving your file. This version of the guide explains all features of the class some of which are only applicable to certain proceeding layouts. The class provides essentially the same markup as implemented by \\LaTeX's standard \\texttt{article} class. In addition to this it implements the following: \\begin{itemize} \\item extended set of front matter commands, \\item automatic placement of floats into column or page areas including turning of table floats by 90\\textdegree{} if necessary, \\item allows mixing column and page-wide floats without getting the numbering out of sync, \\item footnotes will appear below bottom floats, \\item extended set of citation commands if the \\texttt{natbib} system is installed, \\item support for table notes, \\item support for textual page references like ``on the next page''. \\end{itemize} Due to the extended functionality an article written for \\LaTeX{}'s standard article class might need adjustments in the following places before it can be used with the \\aipcls{} (a more detailed description is given in later sections): \\begin{itemize} \\item In the preamble, since the \\aipcls{} requires a |\\layoutstyle| declaration. \\item In the front matter, since the \\aipcls{} uses an extended set of title/author declarations. \\item In the body of floats, since the \\aipcls{} only allows a single |\\caption| command and processes the body in horizontal mode. \\end{itemize} ", "conclusions": "" }, "1003/1003.5003_arXiv.txt": { "abstract": "Owing to their more extensive sky coverage and tighter control on systematic errors, future deep weak lensing surveys should provide a better statistical picture of the dark matter clustering beyond the level of the power spectrum. In this context, the study of non-Gaussianity induced by gravity can help tighten constraints on the background cosmology by breaking parameter degeneracies, as well as throwing light on the nature of dark matter, dark energy or alternative gravity theories. Analysis of the shear or flexion properties of such maps is more complicated than the simpler case of the convergence due to the spinorial nature of the fields involved. Here we develop analytical tools for the study of higher-order statistics such as the bispectrum (or trispectrum) directly using such maps at different source redshift. The statistics we introduce can be constructed from cumulants of the shear or flexions, involving the cross-correlation of squared and cubic maps at different redshifts. Typically, the low signal-to-noise ratio prevents recovery of the bispectrum or trispectrum mode by mode. We define power spectra associated with each multispectra which compresses some of the available information of higher order multispectra. We show how these can be recovered from a noisy observational data even in the presence of arbitrary mask, which introduces mixing between {\\em Electric} (${\\rm E}$-type) and {\\em Magnetic} (${\\rm B}$-type) polarization, in an unbiased way. We also introduce higher order cross-correlators which can cross-correlate lensing shear with different tracers of large scale structures. ", "introduction": "Weak lensing surveys play an important role as cosmological probes complementary to Cosmic Microwave Background (CMB) surveys and large-scale galaxy surveys. In principle, they can probe evolution of the dark matter power spectrum at a moderate redshifts in an unbiased way; for a recent review, see \\cite{MuPhysRep08}. However, the first weak lensing measurements were published within the last decade \\citep{BRE00,Wittman00,KWL00,Waerbeke00} so this is a young field. There has been a tremendous progress in the last decade since the first measurements, on various fronts, including analytical modeling, technical specification and the control of systematics. Weak lensing probes cosmological perturbations on smaller angular scales where the perturbations are both nonlinear and non-Gaussian. The two-point correlation function (or, equivalently, the power spectrum) of weak lensing is weakly sensitive to the cosmological constant $\\Omega_\\Lambda$. It also depends only on a degenerate combination of amplitude of matter power spectrum $\\sigma_8$ and the matter density parameter $\\Omega_M$. This degeneracy can be broken by use of the three-point correlation function which is diagnostic of the non-Gaussianity induced by gravity \\citep{Vil96,JainSeljak97}. The higher order statistics needed to analyze weak lensing in detail are however prone to the effects of noise due to the intrinsic ellipticity distribution, finite number of galaxies (shot noise) and sample (or cosmic) variance owing to the partial sky coverage. Ongoing and planned weak lensing surveys, such as the CFHT legacy survey{\\footnote{http://www.cfht.hawai.edu/Sciences/CFHLS/}}, Pan-STARRS {\\footnote{http://pan-starrs.ifa.hawai.edu/}}, the Dark Energy Survey, and further in the future, the Large Synoptic Survey Telescope {\\footnote{http://www.lsst.org/llst\\_home.shtml}}, JDEM and Euclid will provide a wealth of information in terms of mapping the distribution of mass and energy in the universe. A large fractional sky coverage as well as accurate photometric redshift determination will help to probe higher order correlation functions at a higher signal-to-noise(S/N) then is possible at present (see e.g. \\cite{Pen03}). Almost all previous studies concerning weak lensing have used a real space description and employed correlation function of various orders to probe underlying mass clustering. Future surveys will cover a good fraction of the sky, if not the entire sky. This motivates us to use a harmonic description and employ the mutispectra - which are simply the Fourier representations of the higher order correlation functions. Weak lensing observations require masks with complicated topology to deal with presence of foreground objects. Our formalism in harmonic space can deal with arbitrary mask. The formalism described here is completely general and can deal with fields with arbitrary spins including shear and flexions. In the absence of information concerning the redshifts of the source galaxies, traditional weak lensing studies have tended to operate in projection or in 2D slices. Indeed, before the advent of 3D weak lensing, almost all studies were carried out in such a way \\citep{JSW00}. Due to lack of all-sky coverage most studies also used a ``flat sky'' approach (\\citet{MuJai01,Mu00,MuJa00}) which has led to study of lower order non-Gaussianity using convergence $\\kappa$ as well as shear $\\gamma$ \\citet{MuJai01,Valageas00, MuVa05,VaMuBa04,VaMuBa05}. As the next stage of development tomographic approaches went beyond projected survey by binning galaxies in redshift slices which can further tighten the constraints \\citet{TakadaWhite03,TakadaJain04}. More recently, the use of photometric redshifts to study weak lensing in three dimensions was introduced by \\citet{Heav03}. It was later developed by many authors \\citet{HRH00, HKT06, HKV07, Castro05}, and was shown to be a vital tool in constraining dark energy equation of state \\citep{HKT06}, neutrino mass \\citep{Kit08} and many other possibilities. In our present study we propose estimators for non-Gaussianity using projected data. In most cosmological studies, the power spectrum remains the most commonly used statistical probe and its evolution and characterization remain the best studied. As has already been pointed out, however, higher order statistics can help to break degeneracies e.g. in $\\Omega_M$ and $\\sigma_8$, but are more difficult to probe observationally (see e.g. \\citet{BerVanMell97,JainSeljak97, Hui99,Schneider98,TakadaJain03}) given the low signal to noise associated with them. As a result most commonly used higher order probes compress the available information to a simple number (e.g. ``skewness'' or ``kurtosis'') which are one-point statistics \\citep{Pen03}. The higher order correlation functions have been detected observationally \\citep{BerVanMell97,BerVanMell02} and future larger samples are expected to improve the statistical situation still further. A recent study by \\citet{Mu3D10} has underlined the usefulness of working with two-point correlations (or their Fourier analog the power spectrum) at each order. At the level of the bispectrum, a single power spectrum tends to compress some of the available information; there is more than one degenerate power spectrum associated with the higher order multispectra. We extend the recent study by \\citet{Mu3D10} and provide systematic results at each order for spin$-2$ fields, thus generalizing our previous results for spin-$0$ fields such as the convergence. Understanding higher order statistics also provide an indication of the scatter associated with estimation of lower order statistics \\citet{TakadaJain09}. Their results extend the formalism developed in \\citet{Heav03} and \\citet{Castro05} and we follow the all-sky expansion introduced in weak lensing studies by \\citet{Stebbins96}. Statistics of the shear $\\gamma$ or convergence $\\kappa$ seen in a lensing map are sensitive to the statistics of underlying density contrast $\\delta = {\\delta \\rho_b / \\rho_b}$ of background density $\\rho_b$. An accurate understanding of the gravitational clustering is therefore essential for modeling the weak lensing statistics. At present we lack a detailed analytical model of the non-linear gravitational clustering. In the absence of a clear analytical picture, beyond the perturbative regime the modeling of the higher order statistics are done using a hierarchical ansatz (see e.g. \\citet{Fry84,Schaeffer84, BerSch92,SzaSza93, SzaSza97, MuBaMeSch99, MuCoMe99a, MuCoMe99b, MuMeCo99, MuCo00, MuCo02, MuCo03}). The hierarchical ansatz assumes a factorisable model for the higher order correlation functions. Different hierarchical ansatze differs the way they assign amplitudes to the various tree amplitudes. We will employ a very generic form of the hierarchical ansatz which has already been tested in modeling of projected statistics. The method is flexible enough to allow for a more specific prescription. Other approaches to model nonlinear gravity include the halo models \\citep{CooSeth02} which has also been employed in weak lensing studies. In a recent paper, \\citet{Mu3D10} advocated the use of higher order cumulants and their correlators of convergence $\\kappa$ along with their associated power spectra for studying dark matter clustering beyond power spectra in 3D. The aim of the present paper is to extend that study to fields with non-zero spin (e.g. shear). Shear can directly be used from observed data to study the dark matter clustering beyond the usual approximation of Gaussianity. Although in this study we primarily focus on shear, the analytical results are directly applicable for higher order spin fields sometime used in weak lensing known as {\\it flexions}. Previous studies have also focused on the cumulant correlators of the shear $\\gamma$ and $\\kappa$ \\citep{BerVanMell97,BerVanMell02,BerMellWaer03}. While cumulant correlators for convergence fields can be probed in a relatively straight forward manner, even to arbitrary order \\citep{Mu00}, the analogous studies are more involved for shear fields. There have been some studies involving smoothed $M_{ap}$ statistics to study cumulant correlators from shear map by using compensated filters \\citep{MuVaCross05}, but in this present paper we develop a more systematic approach to probe higher order statistics directly through spinorial objects. The focus of the present study is to develop tools in the Fourier domain and to use the power spectra associated with the multispectra as opposed to cumulant correlators. Given the non-trivial topology of weak lensing surveys correlation functions indeed are a good option to work with, but, as we will show, the power spectra we develop can also be estimated in the presence of arbitrary mask and noisy data. The paper is arranged as follows. In \\textsection2 we discuss the basic formalism and introduce some terminology and notation. In \\textsection3, we introduce the power spectra related to bispectra and trispectra for various combinations of the convergence, shear as well as flexion fields. In \\textsection4 we focus on statistical description of underlying matter distribution and relate them to power spectra defined in previous sections for weak lensing observables. Finally \\textsection5 is reserved for discussions and future prospects. \\begin{figure} \\begin{center} {\\epsfxsize=6. cm \\epsfysize=6. cm {\\epsfbox[301 425 590 713]{cl.eps}}} \\end{center} \\caption{The plots correspond to the power spectrum of convergence $C_l^{\\kappa\\kappa}$. The source redshift is $z_1=z_2=1$. See text for more details.} \\label{fig:ps} \\end{figure} ", "conclusions": "Future weak lensing surveys will play a crucial role in cosmological research, particularly in further reducing the uncertainty in fundamental properties of the standard cosmological model, including those that describe the evolution of equation of state of dark energy \\cite{Euclid}. It is known that, by exploiting both the angular diameter distance and the growth of structure, weak lensing surveys are able to constrain cosmological parameters, as well as testing the accuracy that general relativity provides as a description of gravity \\cite{HKV07,Amendola08,Benyon09, Schrabback09, Kilbinger09}. Constraints from weak lensing surveys are complementary to those obtained from cosmic microwave background studies and from galaxy surveys as they probe structure formation in the dark sector at a relatively low redshift range. Initial studies in weak lensing were restricted to studying two-point functions in projection for the entire source distribution. It was, however, found that binning sources in a few photometric redshift bins can improve the constraints \\cite{Hu99}. More recently a full 3D formalism has been developed which uses photometric redshifts of all sources without any binning\\cite{Heav03,Castro05,HKT06}. These studies have demonstrated that 3D lensing can provide more powerful and tighter constraints on the dark energy equation of state parameter, on neutrino masses \\cite{deBernardis09}, as well as testing braneworld and other alternative gravity models. Most of these 3D works have primarily focused on power spectrum analysis, but in future accurate higher-order statistic measurement should be possible (e.g. \\cite{TakadaJain04,Semboloni09}). In a recent work \\citep{Mu3D10} extended previous studies using convergence maps to probe gravity-induced non-Gaussiniaty. These studies go beyond previous analysis of power-spectrum estimation from convergence data. Recovery of cosmological information however is even more straightforward from shear maps. In this paper we have generalized the results previously obtained for convergence to shear. We have gone beyond conventional power spectrum analysis by defining power spectra associated with higher-order multispectra in such a way as to compress information available in the multispectra to a single derived power spectrum. Generic results which do not depend on detailed modeling of multispetcra have been obtained and later specialized with concrete models using specific forms for the gravity-induced clustering as encoded in the higher-order multispectra. The generic results we have presented do not depend on the specific nature of non-Gaussianity and are able to handle primordial as well as secondary non-Gaussianity. Though the higher-order multispectra contain a wealth of information in through their shape dependence, they are partly degenerate. It would be desirable to exploit the entire information content encoded in higher order multispectra for constraining structure formation scenarios. However, the direct determination of the multispectra and their complete shape dependence from noisy data can be a very difficult task. In this paper we have employed a set of statistics called ``cumulant correlators'' which were first used in real space in the context of galaxy surveys \\citep{SzaSza93,SzaSza97} and later extended to CMB studies \\citep{MuPol10,MuPol09}. We have presented a general formalism for the study of the power spectra or the Fourier transforms of these correlators. We present a 3D analysis which takes into account the radial as well as on the surface of the sky decomposition. Previous studies in weak lensing mainly concentrated on convergence maps containing spin-0 objects that can be treated relatively easily. However, from the point of view of observational data reduction, it is much more natural to focus on galaxy shear which is a natural byproduct of weak lensing surveys. In this paper we have extended the concept of cumulant correlators to higher spin objects to {\\it shear} as well as {\\it flexions}. We consider the correlations of cumulants which are constructed from such spinorial objects. The analytical results we obtain are valid for an {\\it arbitrary} mask and our formalism allows us to correct for the effect of mask at arbitrary order for general spinorial field. The analytical results are sufficiently generic to include a tracer field for the large scale structure and probe mixed bispectrum involving underlying mass distribution as well as the large scale tracer field. We have restricted this study to the third and fourth order, as data from observations gets increasingly noise dominated as the order increases. However, our analysis can be generalized to higher order and some of our results are indeed valid at arbitrary order. At third order, we define a power spectrum which compresses information associated with a bispectrum to a power spectrum. This power spectrum $C_l^{\\sg\\sgp,\\sgpp}(r_2,r_1)$ is the cross-power spectrum associated with squared convergence maps $\\sg(r_1,\\oh)\\sgp(r_1,\\oh)$ constructed at a specific radial distance $r_1$ against $\\sgpp(r_2,\\oh)$ at $r_2$. In a similar manner we also associate power spectra $C_l^{\\sg\\sgp,\\sgpp\\sgppp}$ and $C_l^{\\sg\\sgp\\sgpp,\\sgppp}$ with associated trispectra $T^{l_1l_2}_{l_3l_4}(L;r_i)$. There are two different power spectra at the level of trispectra which are related to the respective real-space correlation functions $\\langle \\sg(r_1,\\oh)\\sgpp(r_1,\\oh)\\sgpp(r_2,\\oh')\\sgppp(r_2,\\oh)\\rangle$ and $\\langle \\sg(r_1,\\oh)\\sgp(r_1,\\oh)\\sgpp(r_1,\\oh)\\sgppp(r_2,\\oh')\\rangle$. We expressed these real-space correlators in terms of their Fourier space analogue which take the form of $C_l^{\\sg\\sgp\\sgpp,\\sgppp}(r_2,r_1)$ and $C_l^{\\sg\\sgp,\\sgpp\\sgppp}(r_2,r_1)$. The statistical descriptors we have presented here will provide particularly useful tools for the study of non-Gaussianity in alternative theories of gravity \\citep{Ber04}; which is one of the important science drivers for the future generations of weak lensing surveys. We plan to present detailed results elsewhere in future. The various analytical approximations, such as the Born approximations linking the observed shear or convergence bispectrum (or trispectrum) have been studied with some rigor \\citep{Ham02,ShapiCoo06}. The effect of source clustering, as well as lens and source overlap - which also introduce corrections - have also been studied with some detail. We plan to extend these studies for the power spectra presented here elsewhere. The estimators we have worked with can be generalized further to take into account optimum weighting. This will involve inverse covariance weighting the observed field harmonics. This will then be useful in constructing a {\\it matched filtering} version of the estimator used here to fine tune the study of primordial as well as gravity induced non-Gaussianity. While previous ray tracing simulation of weak lensing used a flat-sky approach \\citep{JSW00}, simulations are now available for the entire observed sky \\citep{Tey09}. This will provide an opportunity to test the analytical results presented here. A detailed discussion will be presented elsewhere (Munshi et al. 2010, in preparation)." }, "1003/1003.3887_arXiv.txt": { "abstract": "{The ``radiation inner edge'' of an accretion disk is defined as the inner boundary of the region from which most of the luminosity emerges. Similarly, the ``reflection edge'' is the smallest radius capable of producing a significant X-ray reflection of the fluorescent iron line. For black hole accretion disks with very sub-Eddington luminosities these and all other ``inner edges'' locate at ISCO. Thus, in this case, one may rightly consider ISCO as the unique inner edge of the black hole accretion disk. However, even for moderate luminosities, there is no such unique inner edge as differently defined edges locate at different places. Several of them are significantly closer to the black hole than ISCO. The differences grow with the increasing luminosity. For nearly Eddington luminosities, they are so huge that the notion of the inner edge losses all practical significance.} \\authorrunning{M.A. Abramowicz, M. Jaroszy{\\'n}ski, S. Kato, J.-P. Lasota, A. R{\\'o}{\\.z}a{\\'n}ska and A. S{\\k a}dowski} \\titlerunning{Leaving the ISCO} ", "introduction": "\\label{section-introduction} Accretion flows on to black holes must change character before matter crosses the event horizon. Two reasons account for this fundamental property of such flows. First, matter must cross the black-hole surface at the speed of light as measured by a local inertial observer \\citep[see e.g.][]{gj-06}, so that if the flow is sub-sonic far away from the black-hole (in practice it is always the case) it will have to cross the sound barrier (well) before reaching the horizon. This is the property of all realistic flows independent of their angular momentum. The sonic surface in question can be considered as the inner edge of the accretion flow. The second reason is related to angular momentum. Far from the hole many (most probably most) rotating accretion flows adapt the Keplerian angular momentum profile. Because of the existence of the Inner-Most Stable Circular Orbit (ISCO) such flow must stop to be Keplerian there. At high accretion rates when pressure gradients become important the flow may extend below the ISCO but the presence of the Inner-Most Bound Circular Orbit (IBCO) defines another limit for a circular flow (the absolute limit being given by the Circular Photon Orbit; the CPO). These critical circular orbits provide another possible definition of the inner edge of the flow, in this case of an accretion disk. The question is: what is the relation between the accretion flow edges? In the case of geometrically thin disks the sonic and Keplerian edges coincide and one can define the ISCO as the inner edge of such disks. \\cite{paczynski-2000} showed rigorously that, independent of viscosity mechanism, presence of magnetic fields etc. the ISCO is the universal inner disk's edge for not too-high viscosities. The case of thin disks is therefore settled\\footnote{In a recent paper \\cite{Pennaetal-10} studied the effects of magnetic fields on thin accretion disk (the disk thickness $H/r \\lta 0.07$, which corresponds to $L \\lta 0.2\\,L_{\\rm Edd}$). They found that to within a few percent the magnetized disks are consistent with the \\cite{nov-tho-1973} model, in which the inner edge coincides with the ISCO.}. However, this is not the case of non-thin accretion disks, i.e. the case of medium and high luminosities. The problem of defining the inner edge of an accretion disk is not just a formal exercise. \\cite{afshordi-2003} explored several reasons which made discussing the location of inner edge $r = r_{\\rm in}$ of the black hole accretion disks an interesting and important issue. One of them was, \\par \\hangindent=0.5cm {\\it Theory of accretion disks is several decades old. With time ever more sophisticated and more diverse models of accretion onto black holes have been introduced. However, when it comes to modeling disk spectra, conventional steady state, geometrically thin-disk models are still used, adopting the classical ``no torque'' inner boundary condition at the marginally stable orbit.} \\noindent The best illustration of this fact is the case of the state-of-art works on measuring the black hole spin $a$ in the microquasar GRS 1915+105 by fitting its observed ``thermal state'' spectra to these calculated \\citep[e.g.][]{sha-2008, mid-2009}. These works use general relativistic version of the classical Shakura-Sunyaev thin accretion disk model worked out by \\cite{nov-tho-1973}. The Novikov-Thorne model assumes that the inner edge of the the disk $r_{\\rm in} \\equiv r_{\\rm ISCO}$ is also the innermost boundary of the radiating region. Because the black hole mass of GRS 1915+105 is known and therefore fixed ($M_0 = 14\\,M_{\\odot}\\pm 4\\,M_{\\odot}$), the surface area $A$ of the radiating region, calculated in the model, depends only on the black hole unknown spin, $a^*$ ($a^* = Jc/GM^2$ with $J$ being the total angular momentum of the black hole). In the thermal state, the disk spectrum is close to that of a sum of black body contributions from different radial locations. Its shape is determined by the radial distribution of temperature, which in the Novikov-Thorne model depends on the spin, $T = T(r, a^*)$. The total radiation power $L$ is determined by the ``averaged'' temperature $T_0 = T_0(a^*)$ and the surface area $A = A(a^*)$ of the radiating region, $L = \\sigma T_0^4 A$. By calculating the spectral shape and power for different $a^*$ in the Novikov-Thorne model, one may find the best-fit estimates for the spin-dependent temperature and area. This is just the main idea of the spin estimate; details of the fitting are far more complex \\citep[see][]{sha-2008, davi-2005, str-2010} and include, for example, a heuristic way of treating a contribution of scattering in accretion disks atmosphere (i.e. the ``hardening factor''). Results obtained this way by \\cite{sha-2008} for GRS 1915+105 showed that $a^* = 0.99$ for the whole luminosity range $L < 0.2\\,L_{\\rm Edd}$. However, for $L > 0.2\\,L_{\\rm Edd}$, the spin estimated by \\cite{sha-2008} was much lower, $a^* \\approx 0.8$. The inconsistent spin estimates at different luminosities indicate that some assumptions adopted by the Novikov-Thorne model are wrong at high luminosities. This is not a surprise, because there are several physical effects known to be important at high luminosities, but ignored in the classical Shakura-Sunyaev and Novikov-Thorne {\\it thin} accretion disk models. These effects are properly included in the {\\it slim}\\footnote{The names {\\it thin} and {\\it slim} refer to the dimensionless vertical geometrical thickness, $h = H/r$. For thin disks it must be $h \\ll 1$, while for slim disks a weaker condition $h < 1$ holds.} accretion disks models, introduced by \\cite{abr-1988}. Advection is perhaps the best known of these ``slim disk effects'', but in the present context equally important is a significant stress due to the radial pressure gradient (for thin disks $dP/dr \\approx 0$). The stress firmly holds matter well inside ISCO and as a result of this, at high luminosities the edge of the plunge-in region may be considerably closer to the black hole than ISCO\\footnote{Matter may be hold well inside ISCO also by magnetic stresses, as pointed out by many authors; see e.g. a semi-analytic model by \\cite{narayan-mad-2003}, or MHD numerical simulations by \\cite{noble-2010}, and references quoted in these papers.}. \\begin{figure}[h] \\centering \\includegraphics[width=9cm]{009-figure-branches.eps} \\caption{The Figure illustrates a few best-known analytic and semi-analytic solutions of the stationary black hole accretion disks. Their location in the parameter space approximately correspond to viscosity $\\alpha = 0.1$ and radius $r = 20~{\\rm M}$. For detailed reviews of these solutions see, {\\tt \\tiny http://www.scholarpedia.org/article/Accretion\\_discs} or \\cite{kat-2008}.} \\label{fig:branches} \\end{figure} Slim disks are assumed to be stationary and axially symmetric. They are described by vertically integrated Navier-Stokes hydrodynamical equations; no magnetic fields are considered. The effective viscosity, believed to be generated by the MHD turbulence \\citep{balb-1991} is described by the ``$\\alpha P$'' Shakura-Sunyaev ansatz. Figure~\\ref{fig:branches} shows the slim disk location with respect to other analytic and semi-analytic disk models, in the parameter space $[\\tau, h, {\\dot m}]$ described by the vertical optical depth $\\tau$, dimensionless vertical thickness $h = H/r$ and dimensionless accretion rate ${\\dot m}= {\\dot M}/{\\dot M}_{\\rm Edd}$, where ${\\dot M}_{\\rm Edd} = 16L_{\\rm Edd}/c^2$ is the critical accretion rate approximately corresponding to the Eddington luminosity ($L_{\\rm Edd} = 10^{38}M/M_{\\odot}\\,$erg/s) in case of a disk around a non-rotating black hole\\footnote{Two warnings about notation. (i) Many authors use a different definition, ${\\dot M}_{\\rm Edd} = L_{\\rm Edd}/c^2$. (ii) We often use the $c = 1 = G$ convention in which $M = r_G = GM/c^2$.}. In this paper, we discuss properties of the inner edge of slim accretion disks around rotating black holes, using models similar to those calculated recently by \\cite{sad-2009}\\footnote{At {\\tt \\tiny http://users.camk.edu.pl/as/slimdisks} a very detailed data base for these solutions is given. It covers the whole parameter space relevant for microquasars and AGN.}. For convenience, we shortly remind the slim disk basic equations in the Appendix~A. In the following Section \\ref{section-definitions}, we list six possible definitions of the inner edge. These definitions reflect different (but partially overlapping) physical meanings and different practical astrophysical applications. In the following six Sections \\ref{section-potential-spout}-\\ref{section-reflection} we calculate the slim disk locations of these six inner edges, and discuss their astrophysical relevance. Some of the results presented here have been anticipated previously by us and other authors in a different context of the Polish doughnuts \\citep[i.e. thick accretion disks; see e.g. a short review by][]{paczynski-1998}; see also \\cite{paczynski-2000} and \\cite{afshordi-2003}. ", "conclusions": "\\label{section-conclusions} We addressed the inner edge issue by discussing behavior of six differently defined ``inner edges'' of slim accretion disks around the Kerr black hole. We found that the slim disk inner edges behave very differently than the corresponding Shakura-Sunyaev and Novikov-Thorne ones. The differences are qualitative. Even for moderate luminosities, ${\\dot M} \\gtrsim 0.3\\,{\\dot M}_{\\rm Edd}$, there is no unique inner edge. Differently defined edges locate at different places. For nearly Eddington luminosities, the differences are huge and the notion of the inner edge losses all practical significance. We summarize the properties and locations of the six inner edges in Table \\ref{table:1}. It refers to $a^*=0$, but the qualitative behavior is similar for $a^*\\not =0$. \\begin{table*}[t] \\caption{Summary of the results (specific numbers refer to the case $a^*=0$).} % \\label{table:1} % \\centering % \\begin{tabular}[width=1.0\\textwidth]{l m{3.cm} m{1.7cm} m{1.25cm} m{2.cm} m{2.cm} m{2.5cm} m{.0001cm}} % \\hline\\hline % & \\centering$r_{\\rm pot}$ & \\centering$r_{\\rm son}$ & \\centering$r_{\\rm var}$ & \\centering$r_{\\rm str}$ & \\centering$r_{\\rm rad}$ & \\centering $r_{\\rm ref}$ &\\\\ \\hline % &&&&&&\\\\ $\\dot m\\lesssim 0.3$ & \\multicolumn{6}{c}{ $ r_{\\rm in} \\approx r_{\\rm pot} \\approx r_{\\rm son} \\approx r_{\\rm var} \\approx r_{\\rm str} \\approx r_{\\rm rad} \\approx r_{\\rm ref} \\approx r_{\\rm ISCO} $ } \\\\&&&&&&\\\\ \\hline % $\\dot m\\gtrsim 0.3$ & \\begin{center} for $\\alpha\\lesssim0.1$ moves inward with increasing $\\dot m$ down to $\\sim r_{mb}$; ~ for $\\alpha\\gtrsim0.1$ and sufficiently high $\\dot m$ disk enters the Bondi regime --- undefined \\end{center}& \\begin{center} departs from ISCO; ~ for $\\alpha\\ll0.1$ $r_{\\rm son}\\approx r_{mb}$; ~ for $\\alpha\\gtrsim0.2$ $r_{\\rm son}>r_{ISCO}$ \\end{center} & \\begin{center} undefined\\end{center} & \\begin{center} moves inward with increasing $\\dot m$ down to BH horizon.\\end{center} & \\begin{center} moves inward with increasing $\\dot m$ down to BH horizon.\\end{center} & \\begin{center} for $0.3\\lesssim\\dot m\\lesssim1.0$ $r_{\\rm ref} 5\\times 10^{-2}$ pc from the central engine. } {} } ", "introduction": "It is commonly accepted that the center of active galaxies (Active Galactic Nuclei -AGNs) hosts a massive black hole (with a mass in the range $10^6-10^9$ M$_{\\odot}$) accreting the surrounding material via the formation of a disk. How the energy released from the central engine interacts with the local environment and contributes to the history of the host galaxy is one of the crucial question of present astrophysical research. In this respect, while the mechanisms of energy output in the form of radiation and relativistic jets are quite well understood, it also seems that the outflowing winds have an important role in the overall energy budget. Although the origin of these winds is still controversial, at our present level of understanding the narrow-line regions, the inner part of an obscuring torus (\\citealt{blustin2005}) and the black hole accretion disk \\citep{elvis2000} are all possible locations. X-ray obscured AGNs (with an intrinsic column density $N_H\\ut> 10^{22}$ cm$^{-2}$) are not completely dark in the soft X-ray band. High resolution {\\it XMM}-Newton and Chandra observations revealed a complex spectrum dominated by emission lines from He-and H-like transitions of elements from carbon to neon as well as by L-shell transitions of {Fe\\,\\textsc{xvii}} to {Fe\\,\\textsc{xxi}} ions (\\citealt{sakoa}, \\citealt{ali}, \\citealt{sambruna}, \\citealt{armentrout2007}). This gas, which shows the signature of a photoionization process (\\citealt{ali}, \\citealt{bianchi2007}), is sometimes referred to as a warm mirror. In unobscured AGNs a modification of the output energy spectrum may also occur as a consequence of absorption by a warm ionized gas along the line of sight. The properties of these so called warm absorbers can be summarized as follows: i) average ionization parameter in the range $\\log \\xi =0-3$, ii) total column density in the range $\\log N_H = 21-22$ cm$^{-2}$, iii) outflow velocities of hundred of km s$^{-1}$ (see e.g. \\citealt{blustin2005}, but also \\citealt{steenbrugge}). Evidence of a multi-phase warm absorber gas was also recently reported for Mrk~841 (\\citealt{longinotti2009}). In general, detecting warm mirror signatures is easier in sources in low flux states, because the emission features are not outshone by the continuum radiation. This was the case for the Seyfert 1 galaxy Mrk 335, whose soft X-ray spectrum resembled the spectra of obscured AGNs when the source was observed at low state (\\citealt{longinotti2008}), but does not show any evidence of a warm absorber in the high flux state (\\citealt{longinotti2007}). The overall properties of the warm mirror (even if it is poorly constrained) and the warm absorber (as described above) are similar so that there is the possibility that they represent the same physical system. Conversely, the interplay between the warm absorber and warm mirror regimes is best studied in sources that display both components. The source NGC 4051, a narrow-line Seyfert galaxy at the redshift of $0.00234$, was at the center of many past investigations in the X-ray band because it offers a unique laboratory where to test present theories and models about the physics of AGNs. The X-ray emission is characterized by rapid variations \\citep{lamer2003,ponti2006} sometimes showing periods of low activity (see \\citealt{lawrence1987} and \\citealt{uttley}). Its power spectral distribution (PSD) in high state resembles the behavior of a galactic black hole system (\\citealt{mchardy}). At high X-ray flux, the spectrum of the galaxy is characterized by a power law with photon index $\\Gamma\\sim 1.8-2$ which becomes harder above 7 keV where a reflection component from cold matter has been observed. On long time-scales, the X-ray light curve of NGC 4051 shows low state flux periods of several months during which the spectrum in the energy range 2-10 keV becomes harder ($\\Gamma \\simeq 1$) and shows a strong iron $K\\alpha$ line (as found by \\citealt{guainazzi98} in Beppo-SAX data). A soft X-ray excess is also evident. As reported by \\citet{warmabsorber}, the high state X-ray spectrum of NGC 4051 in the soft band is a combination of continuum and emission line components. Curvature in the spectrum cannot be explained with simple models, i.e. a single power law or a black body, because an ionized absorber-emitter has to be taken into account as well. In this context, \\citet{krongold} showed that the evolution in time of the properties of the warm absorber can constrain the physical parameters of the absorbing gas. In particular they find that at least two different ionization components are required with matter densities of $\\simeq 10^6$ cm$^{-3}$ and $\\ut> 10^7$ cm$^{-3}$, thus placing the warm absorber in the vicinity of the accretion-disk. Dynamical arguments permit us to infer that the warm absorber gas originates in a radiation-driven high-velocity outflow in accretion disk instabilities (\\citealt{krongold}). On the other hand, as shown by \\cite{pounds}, the low state flux spectrum of NGC 4051 is dominated by narrow emission lines and radiative recombination continua (RRC) from hydrogenic and He-like carbon, oxygen, neon and nitrogen. To be specific, a fit to the identified RRCs yields a mean temperature for the emitting gas of $T\\simeq 4\\times 10^{4}$ K, which favors a scenario invoking a photoionization process. In this case, the soft X-ray spectrum of NFG 4051 in low state is similar to that observed for the prototype Seyfert 2 galaxy NGC 1068 (see \\citealt{ali}). Below we do not repeat the analysis of the EPIC data but refer to \\citet{pounds} for more details on the main results obtained in the energy band 0.3-10 keV. We only say that a comparison between the EPIC PN data for the 2001 and 2002 observations shows that the high state observation flux level is a factor $\\sim 5$ greater with respect to the low state. Furthermore, the spectrum shows a gradual flattening of the continuum slope from 3 keV up to 6.4 keV. It was also noted that when the fit to the 0.3-10 keV band continuum is extrapolated down in the soft X-ray (0.3-3 keV) a strong excess appears in both the two observations, and as is clear from the RGS spectrum, it can be explained by a blending effect of fine structures (emission lines). Here we first conducted a phenomenological study of the emission lines identified in the spectrum of NGC 4051 and compare our results with those known in literature. We further compared the RGS emission line spectrum with synthetic spectra generated with the photoionization code Cloudy 8 (\\citealt{ferland}). For this purpose we followed a similar approach as in \\cite{armentrout2007} (to which we refer for more details) on NGC 4151. The paper is structured as follows: in Sect. 2 we briefly describe the reduction of the {\\it XMM}-Newton data set and describe our phenomenological analysis of the soft $X$-ray spectrum of NGC 4051. In Sects. 3 and 4 we give details on the Cloudy model developed and address some conclusions. \\begin{figure*}[htbp] \\vspace{9.5cm} \\special{psfile=f1.ps vscale=60. hscale=60. voffset=-10 hoffset=480 angle=90.0} \\caption{Fluxed RGS spectrum of NGC 4051 (low state). The first order spectra of the two RGS cameras were combined and the resulting spectrum smoothed with a triangular kernel. The identified lines are labeled with the corresponding ion transition name and vertical dashed lines (the big dips at $\\simeq 13$ \\AA and $\\simeq 21$ \\AA in this plot are due to CCD gaps).} \\label{spectrum} \\end{figure*} ", "conclusions": "Most of the information on the physics and geometry of gas in AGNs is inferred by means of optical spectroscopy and imaging techniques with which it was shown that the AGN central high energy emission is the main source of ionizing photons with an occasional contribution from collisionally ionized plasma. In the last years X-rays observations acquired an important role in AGN studies particularly since Chandra showed the existence, at least for Seyfert 2 galaxies, of extended (a few kpc) X-ray emission \\citep{bianchi2006} similarly to what was observed in the optical band. High resolution spectroscopy in the soft X-ray band ($0.2-2$ keV) confirms the overall scenario, and photoionization seems to be the dominant ionization mechanism which results in a spectrum characterized by recombination lines from He- and H-like transitions of C to Si elements and by Fe-L transitions. In this respect, X-ray high-resolution spectroscopy offers a powerful diagnostic tool because the observed spectral features strongly depend on the physical properties of matter (ionization parameter $U$, electron density $n_e$, hydrogen column density $N_H$ as well as size and location of the emitting clouds). The Seyfert 1 object NGC 4051 shows a very rich emission line $X$-ray spectrum when observed in low-flux state. According to the analysis we conducted on the {\\it XMM}-Newton RGS data, the observed soft $X$-ray features originate in a low-density photoionized gas. In order to constrain the physical properties of the photoionized gas, we simulated synthetic spectra via the Cloudy software (\\citealt{ferland}) and compared them to the RGS data with standard minimization techniques. We found that to describe the overall soft $X$-ray spectrum, at least a three-phase gas is required (two emission components and one warm absorbing component). Referring to the emission components respectively as {\\it low} and {\\it high} ionization components, our fit procedure gave us their physical properties. For the {\\it low} component we have $\\log U\\simeq 0.63$, and $\\log (N_H/{\\rm cm^{2}})\\simeq 22.10$ and for the {\\it high} component we have $\\log U\\simeq 1.90$, and $\\log (N_H/{\\rm cm^{2}})\\simeq 22.20$. Using Cloudy we get for the electron density $n_e$ an upper limit of $\\log (n_e/{\\rm cm^{3}}) \\simeq 9$, which reduces to $\\log (n_e/{\\rm cm^{3}}) \\simeq 5$ when the recombination time scale of {O\\,\\textsc{vii}} is taken into account. Even if the warm absorber gas seems to be required by our fit procedure, its parameters are poorly constrained. Thus it is characterized by $\\log U\\simeq 0.85$, $\\log (N_H/{\\rm cm^{2}})\\simeq 23.36$, and $\\log (n_e/{\\rm cm^{3}}) \\ut< 7$. This technique was successfully applied before to the Seyfert 1 Mrk~335 and the Seyfert 2/starburst galaxy NGC~1365 (\\citealt{longinotti2008,guainazzi2009}). The main difference is that NGC~4051 is characterized by a strong warm absorber component in the high flux state that is still affecting the spectrum even when the nuclear flux is attenuated. Indeed, we found out in our analysis of the low flux state data that the effect of the line of sight medium is not negligible, particularly not in the modeling of the resonance line of the {O\\,\\textsc{vii}} triplet (see Fig. \\ref{fig6}) which is close to several absorption features. For example, the resonance line could be weakened by the same line in absorption\\footnote{As noted by \\citet{sakob} and \\citet{ali}, the resonance line of the {O\\,\\textsc{vii}} triplet could be also enhanced by photoexcitation. Note however that this would also result in a boost of all the higher order resonance transitions of the H-like and He-like ions (Ly-$\\beta$, Ly-$\\delta$, He-$\\beta$ and He-$\\delta$), but this enhancement is not currently observed.} (see e.g. \\citealt{krongold}). Nonetheless, the physical parameters of the warm absorber cannot be well-constrained by the analysis of the low flux state data (see Table 4). The average distance $r$ of each of the photoionized plasma-emitting components from the nuclear source can be estimated by the definition of the ionization parameter $U$ after normalizing to the ionization luminosity $L_{ion}$. However, our results are insensitive to values of the electron density $n_e$ lower than $10^5$ cm$^{-3}$. In this limit, we can only determine a lower limit of the $X$-ray-emitting gas location (Table 4). The analysis carried out in this paper allowed us to identify two ionization states for the line emitting gas and one warm absorber medium. It is interesting to note that \\\\ - The $X$-ray emitting region can be placed at a distance of $r\\ut> 0.05$ pc. Indeed, \\citet{warmabsorber} found that the NGC 4051 $X$-ray narrow-line regions can be placed at a distance of the same order of magnitude. This was also confirmed by the Chandra ACIS-S images of the same galaxy (\\citealt{uttley}), which showed a size of the diffuse emission smaller than that of the optical narrow-line regions (30 -220 pc, \\citealt{christopoulou1997}), thus implying a clear separation between the $X$-ray and optical emissions. This is also naturally expected as a consequence of projection effects: as shown by \\citet{schmitt2003}, who studied a sample of 60 Seyfert galaxies with the Hubble Space Telescope, the Seyfert 1 narrow-line regions objects are more circular and compact than those in the Seyfert 2 galaxies, with the Seyfert 2 subsample characterized by more elongated shapes. This agrees well with the unified picture according to which the conical narrow-line region of a Seyfert 1 galaxy is observed close to the axis of symmetry, while that of a Seyfert 2 galaxy is observed from an orthogonal line of sight. Furthermore, the scale-length found in this paper is consistent with the inner radius of the torus in NGC~4051 as determined by \\citet{blustin2005}, i.e. $r\\simeq 0.15$~pc.\\\\ - The NGC 4051 low state warm absorber is poorly constrained but its existence is nevertheless required by the fit. In particular, we found a lower limit of the warm absorber distance $\\simeq 0.02$ pc, i.e. at least a factor $10$ larger than that measured in the high state flux (\\citealt{krongold}). Indeed, by using the long $XMM$-Newton exposure of NGC 4051 in its high flux state and studying the time evolution of the ionization states of the $X$-ray absorbers, \\citet{krongold} were able to put severe constraints on the physical and geometrical properties of the warm absorber medium. They specifically found that the warm absorber consists of two different ionization components which are located within $3.5$ lt-days (or $0.0029$ pc) from the central massive black hole. This result allowed the authors to exclude an origin in the dusty obscuring torus because the expected dust sublimation radius\\footnote{The torus inner edge has to be at a distance larger than the dust sublimation radius $r_{sub}$. In the particular case of NGC 4051, \\citet{krongold} found $r_{sub}\\simeq 0.01$ pc.} is at least one order of magnitude larger. Hence the authors suggested a model in which the black hole accretion disk is at the origin of a $X$-ray absorber wind, which forms a conical structure moving upward. \\begin{figure*}[t] \\vspace{0.4cm} \\begin{center} $\\begin{array}{c@{\\hspace{0.05in}}c@{\\hspace{0.05in}}c} \\epsfxsize=3.55in \\epsfysize=3.55in \\epsffile{f8a.eps} & \\epsfxsize=3.55in \\epsfysize=3.55in \\epsffile{f8b.eps} \\end{array}$ \\end{center} \\caption{The cartoon shows qualitatively the location of the $X$-ray emitting and absorbing material (see text for details).} \\label{figcartoon} \\end{figure*} If this is the correct picture, when the continuum source is switched off, the compact warm absorber might not be observed anymore during the low state flux of NGC 4051. Our analysis showed instead the existence of a more exterior $X$-ray absorber, which absorbs the soft $X$-ray photons emitted from sources (as for example the inner surface of the conical structure proposed by \\citealt{krongold}) located (in projection) at scales larger than the torus and/or the narrow-line regions. Remarkably, this could indicate the existence of a diffuse warm material filling the wind-generated cone. Figure \\ref{figcartoon} gives a qualitative representation of the model. During the high state flux (left panel) a two ionization component warm absorber (here labeled as II) lying within a few l-days ($\\simeq 0.003$ pc) from the accreting black hole was identified by \\citet{krongold}. \\citet{warmabsorber} found that the NGC 4051 $X$-ray narrow-line regions can be placed at a distance of $r>0.02$ pc, while the optical narrow-line regions are on the scale of tenth of parsec (\\citealt{christopoulou1997}). During the low state flux (right panel), the interior warm absorber might not be observed anymore since the central engine is switched off. A more exterior warm absorber (labeled as I) could now absorb the $X$-ray photons emitted from sources located on the scale larger than the torus and/or the narrow-line regions." }, "1003/1003.4682_arXiv.txt": { "abstract": "{} {Synthetic spectra are needed to determine fundamental stellar and wind parameters of all types of stars. They are also used for the construction of theoretical spectral libraries helpful for stellar population synthesis. Therefore, a database of theoretical spectra is required to allow rapid and quantitative comparisons to spectroscopic data. We provide such a database offering an unprecedented coverage of the entire Hertzsprung-Russell diagram.} {We present the POLLUX database of synthetic stellar spectra. For objects with T$_{\\rm eff} \\le$ 6\\,000 K, MARCS atmosphere models are computed and the program TURBOSPECTRUM provides the synthetic spectra. ATLAS12 models are computed for stars with 7\\,000 K $\\le T_{\\rm eff} \\le$ 15\\,000 K. SYNSPEC gives the corresponding spectra. Finally, the code CMFGEN provides atmosphere models for the hottest stars ($T_{\\rm eff} >$ 25\\,000 K). Their spectra are computed with CMF\\_FLUX. Both high resolution (R$>$150\\,000) optical spectra in the range 3\\,000 to 12\\,000 \\AA\\ and spectral energy distributions extending from the UV to near--IR ranges are presented. These spectra cover the HR diagram at solar metallicity. } {We propose a wide variety of synthetic spectra for various types of stars in a format that is compliant with th Virtual Observatory standards. A user--friendly web interface allows an easy selection of spectra and data retrieval. Upcoming developments will include an extension to a large range of metallicities and to the near--IR high resolution spectra, as well as a better coverage of the HR diagram, with the inclusion of models for Wolf-Rayet stars and large datasets for cool stars. The POLLUX database is accessible at http://pollux.graal.univ-montp2.fr/ and through the Virtual Observatory.} {} ", "introduction": " ", "conclusions": "" }, "1003/1003.2637_arXiv.txt": { "abstract": "Our numerical simulations show that the reconnection of magnetic field becomes fast in the presence of weak turbulence in the way consistent with the Lazarian \\& Vishniac (1999) model of fast reconnection. We trace particles within our numerical simulations and show that the particles can be efficiently accelerated via the first order Fermi acceleration. We discuss the acceleration arising from reconnection as a possible origin of the anomalous cosmic rays measured by Voyagers. ", "introduction": "A magnetic field embedded in a perfectly conducting fluid preserves its topology for all time (Parker 79). Although ionized astrophysical objects, like stars and galactic disks, are almost perfectly conducting, they show indications of changes in topology, ``magnetic reconnection'', on dynamical time scales (Parker 1970, Lovelace 1976, Priest \\& Forbes 2002). Reconnection can be observed directly in the solar corona ( Innes et al 1997, Yokoyama \\& Shibata 1995, Masuda et al. 1994), but can also be inferred from the existence of large scale dynamo activity inside stellar interiors (Parker 1993, Ossendrijver 2003). Solar flares (Sturrock 1966) and $\\gamma$-ray busts (Fox et al. 2005, Galama et al. 1998) are usually associated with magnetic reconnection. Previous work has concentrated on showing how reconnection can be rapid in plasmas with very small collisional rates (Shay et al. 1998, Drake 2001, Drake et al. 2006, Daughton et al. 2006), which substantially constrains astrophysical applications of the corresponding reconnection models. We note that if magnetic reconnection is slow in some astrophysical environments, this automatically means that the results of present day numerical simulations in which the reconnection is enevitably fast due to numerical diffusivity do not correctly represent magnetic field dynamics in these environments. If, for instance, the reconnection were slow in collisional media this would entail the conclusion that the entire crop of interstellar, protostellar and stellar MHD calculations would be astrophysically irrelevant. Here we present numerical evidence, based on three dimensional simulations, that reconnection in a turbulent fluid occurs at a speed comparable to the rms velocity of the turbulence, regardless of either the value of the resistivity or degree of collisionality. In particular, this is true for turbulent pressures much weaker than the magnetic field pressure so that the magnetic field lines are only slightly bent by the turbulence. These results are consistent with the proposal by Lazarian \\& Vishniac (1999, henceforth LV99) that reconnection is controlled by the stochastic diffusion of magnetic field lines, which produces a broad outflow of plasma from the reconnection zone. This work implies that reconnection in a turbulent fluid typically takes place in approximately a single eddy turnover time, with broad implications for dynamo activity (Parker 1970, 1993, Stix 2000) and particle acceleration throughout the universe (de Gouveia dal Pino \\& Lazarian 2003, 2005, Lazarian 2005, Drake et al. 2006). Astrophysical plasmas are often highly ionized and highly magnetized (Parker 1970). The evolution of the magnetic field in a highly conducting fluid can be described by a simple version of the induction equation \\begin{equation} \\frac{\\partial \\vec{B}}{\\partial t} = \\nabla \\times \\left( \\vec{v} \\times \\vec{B} - \\eta \\nabla \\times \\vec{B} \\right) , \\end{equation} where $\\vec{B}$ is the magnetic field, $\\vec{v}$ is the velocity field, and $\\eta$ is the resistivity coefficient. Under most circumstances this is adequate for discussing the evolution of magnetic field in an astrophysical plasma. When the dissipative term on the right hand side is small, as is implied by simple dimensional estimates, the magnetic flux through any fluid element is constant in time and the field topology is an invariant of motion. On the other hand, reconnection is observed in the solar corona and chromosphere (Innes et al. 1997, Yokoyama \\& Shibata 1995, Masuda et al. 1994, Ciaravella \\& Raymond 2008), its presence is required to explain dynamo action in stars and galactic disks (Parker 1970, 1993), and the violent relaxation of magnetic fields following a change in topology is a prime candidate for the acceleration of high energy particles (de Gouveia Dal Pino \\& Lazarian 2003, henceforth GL03, 2005, Lazarian 2005, Drake et al. 2006, Lazarian \\& Opher 2009, Drake et al. 2010) in the universe. Quantitative general estimates for the speed of reconnection start with two adjacent volumes with different large scale magnetic field directions (Sweet 1958, Parker 1957). The speed of reconnection, i.e. the speed at which inflowing magnetic field is annihilated by ohmic dissipation, is roughly $\\eta/\\Delta$, where $\\Delta$ is the width of the transition zone (see Figure 1). Since the entrained plasma follows the local field lines, and exits through the edges of the current sheet at roughly the Alfven speed, $V_A$, the resulting reconnection speed is a tiny fraction of the Alfven speed, $V_A\\equiv B/(4\\pi \\rho)^{1/2}$ where $L$ is the length of the current sheet. When the current sheet is long and the reconnection speed is slow this is referred to as Sweet-Parker reconnection. Observations require a speed close to $V_A$, so this expression implies that $L\\sim \\Delta$, i.e. that the magnetic field lines reconnect in an ``X point''. The first model with a stable X point was proposed by Petschek (1964). In this case the reconnection speed may have little or no dependence on the resistivity. The X point configuration is known to be unstable to collapse into a sheet in the MHD regime (see Biskamp 1996), but in a collisionless plasma it can be maintained through coupling to a dispersive plasma mode (Sturrock 1966). This leads to fast reconnection, but with important limitations. This process has a limited astrophysical applicability as it cannot be important for most phases of the interstellar medium (see Draine \\& Lazarian 1998 for a list of the idealized phases), not to speak about dense plasmas, such as stellar interiors and the denser parts of accretion disks. In addition, it can only work if the magnetic fields are not wound around each other, producing a saddle shaped current sheet. In that case the energy required to open up an X point is prohibitive. The saddle current sheet is generic for not parallel flux tubes trying to pass through each other. If such a passage is seriously constrained, the magnetized highly conducting astrophysical fluids should behave more like Jello rather than normal fluids. Finally, the traditional reconnection setup does not include ubiquitous astrophysical turbulence\\footnote{The set ups where instabilities play important role include Simizu et al. (2009a,b). For sufficiently large resolution of simulations those set-ups are expected to demonstrate turbulence. Turbulence initiation is also expected in the presence of plasmoid ejection (Shibata \\& Tanuma 2001). Numerical viscosity constrains our ability to sustain turbulence via reconnection, however.} (see Armstrong, Rickett \\& Spangler 1994, Elmegreen \\& Scalo 2004, McKee \\& Ostriker 2007, Haverkorn, Lazarian 2009, Chepurnov \\& Lazarian 2010). Fortunately, this approach provides another way of accelerating reconnection. Indeed, an alternative approach is to consider ways to decouple the width of the plasma outflow region from $\\Delta$. The plasma is constrained to move along magnetic field lines, but not necessarily in the direction of the mean magnetic field. In a turbulent medium the two are decoupled, and fluid elements that have some small initial separation will be separated by a large eddy scale or more after moving the length of the current sheet. As long as this separation is larger than the width of the current sheet, the result will not depend on $\\eta$. \\begin{figure}[!t] \\begin{center} \\includegraphics[width=1.0 \\columnwidth]{fig1.eps} \\caption{{\\it Upper plot}: Sweet-Parker model of reconnection. The outflow is limited by a thin slot $\\Delta$, which is determined by Ohmic diffusivity. The other scale is an astrophysical scale $L\\gg \\Delta$. {\\it Middle plot}: Reconnection of weakly stochastic magnetic field according to LV99. The model that accounts for the stochasticity of magnetic field lines. The outflow is limited by the diffusion of magnetic field lines, which depends on field line stochasticity. {\\it Low plot}: An individual small scale reconnection region. The reconnection over small patches of magnetic field determines the local reconnection rate. The global reconnection rate is substantially larger as many independent patches come together. From Lazarian et al. 2004.} \\label{fig_rec} \\end{center} \\end{figure} LV99 we introduced a model that included the effects of magnetic field line wandering (see Figure 1). The model relies on the nature of three-dimentional magnetic field wandering in turbulence. This nature is different in three and two dimensions, which provides the major difference between the LV99 model and the earlier attempts to solve the problem of magnetic reconnection appealing to turbulence (Matthaeus \\& Lamkin 1985). The effects of compressibility and heating which were thought to be important in the earlier studies (Matthaeus \\& Lamkin 1985, 1986) do not play the role for the LV99 model either. The model is applicable to any weakly perturbed magnetized fluid, irrespectively, of the degree of plasma being collisional or collisionless (cf. Shay et al. 1998). Two effects are the most important for understanding of the nature of reconnection in LV99. First of all, in three dimensions bundles of magnetic field lines can enter the reconnection region and reconnection there independently (see Figure~1), which is in contrast to two dimensional picture where in Sweet-Parker reconnection the process is artificially constrained. Then, the nature of magnetic field stochasticity and therefore magnetic field wandering (which determines the outflow thickness, as illustrated in Figure~1) is very different in 2D and the real 3D world (LV99). In other words, removing artificial constraints on the dimensionality of the reconnection region and the magnetic field being absolutely straight, LV99 explores the real-world astrophysical reconnection. Our calculations in LV99 showed that the resulting reconnection rate is limited only by the width of the outflow region. This proposal, called ``stochastic reconnection'', leads to reconnection speeds close to the turbulent velocity in the fluid. More precisely, assuming isotropically driven turbulence characterized by an injection scale, $l$, smaller than the current sheet length, we find \\begin{equation} V_{rec}\\approx \\frac{u_l^2}{V_A}\\left(l/L\\right)^{1/2}\\approx u_{turb}\\left(l/L\\right)^{1/2}, \\label{recon1} \\end{equation} where $u_l$ is the velocity at the driving scale and $u_{turb}$ is the velocity of the largest eddies of the strong turbulent cascade. Note, that here \"strong\" means only that the eddies decay through nonlinear interactions in an eddy turn over time (see more discussion of the LV99). All the motions are weak in the sense that the magnetic field lines are only weakly perturbed. It is useful to rewrite this in terms of the power injection rate $P$. As the perturbations on the injection scale of turbulence are assumed to have velocities $u_l0$) eliminates this singularity; \\item all curvature invariants on the null surface in this case are exactly equal to their values on the horizon of a Schwarzschild-de Sitter black hole of mass $m$ and Hubble constant $H_{0}>0$; \\item therefore, at least in the case $H_{0}>0$ the McVittie metric describes a regular (on the horizon) black hole embedded in an FRW spacetime. \\end{itemize} The items marked $\\diamond$ are new results of the work reported in this paper, correcting existing claims in the literature. ", "conclusions": "\\label{sec:conclusions} In this paper we analyzed the McVittie solutions of Einstein's field equations, which have been debated for nearly 80 years as candidates for describing the gravitational fields of spherically symmetric mass distributions in expanding FRW universes. For simplicity we focused on spatially flat McVittie geometries. Our results show that the McVittie solutions that asymptote to FRW universes dominated by a positive cosmological constant at late times are black holes with regular event horizons. While we were not able to construct the explicit form of the extension of the geometry past the black hole event horizon $r=r_-, t = \\infty$, we showed that such analytic extensions {\\em exist} when $H \\rightarrow H_0 > 0$, being defined by the ingoing null geodesic family. Further, we demonstrated that in this limit the near horizon geometry is asymptotically Schwarzschild-de Sitter, which is essentially a consequence of the cosmic and black hole no-hair dynamics. Thus at least in this case, the shaded area past the surface $r=r_-, t = \\infty$ in Figure \\ref{sdsmcvitt} is a real black hole interior of some sort. When the exterior geometry asymptotes to an FRW cosmology with a power law scale factor, so that $H = p/t \\rightarrow 0$ in the future, the McVittie solutions are singular on the null surface $r=r_- = 2m, t = \\infty$. This surface can still be reached by null geodesics in finite affine parameter, and so is a real singularity. However it is very soft since it takes invariants involving at least two derivatives of the curvature to detect it. While it is naked, it respects the spirit of the strong cosmic censorship \\cite{robm} in that it cannot be observed directly by external observers. Our results correct the claims made in \\cite{nolan} which have propagated through the literature.\\footnote{See for example \\cite{ers}. There are many conflicting incorrect claims, ranging from assertions that McVittie solutions are geodesically complete, over the suggestions that they cannot be black holes because of exotic spacelike singularities, to the statement that McVittie solutions in any accelerating universes are black holes. Our results establish a systematic method for determining what the McVittie solutions {\\em are} and dispel those incorrect claims.} Our most interesting conclusion is that the McVittie solutions in spacetimes which are asymptotically dominated by a positive cosmological constant are black holes. It seems rather remarkable that these black holes were not identified until nearly 80 years after the metric was written down. One might be puzzled that what seems to be an infrared quantity---the cosmological constant---regulates a short-distance property of the geometry, the would-be null singularity. This can be understood as a consequence of the special and restrictive form of the McVittie solutions, which requires the validity of a continuous fluid description down to the shortest scales, prevents the hole from accreting, and enforces that its energy density be exactly homogeneous on some set of spatial slices. As we have discussed, these conditions are not particularly well-motivated physically. One could give up spatial homogeneity of the energy density, as in {\\it e.g.} \\cite{sussman} and its generalizations, or allow the hole to accrete. Another strategy for describing black holes in cosmological backgrounds is to go beyond the simple fluid description of cosmological matter and replace it with a microscopic description. Such ideas motivated the work of \\cite{kastra}, and were behind the proposal of the levitating dark matter in \\cite{kalpad}. A concrete example which realizes this has been found very recently in \\cite{gibma}, based on the papers \\cite{mae1,mae2}, and generalized in \\cite{mae3}. The solutions found in those papers are very interesting, in that they encode large distance properties of the McVittie family. The matter content is a scalar field with an asymptotic runaway potential and a system of gauge fields with nonzero charges. Far from the hole the scalar rolling down the potential provides the source that drives the cosmic expansion. Near the hole the charge contributions correct the effective potential for the scalar and give it a large mass, employing essentially the same dynamics as the supersymmetric attractor mechanism in asymptotically flat black holes \\cite{fixscalars}. Thus close to the black hole horizon the scalar decouples and there is no medium that could create divergences. This method of regulating the null McVittie singularity is local, but of course it requires abandoning both the fluid description near the black hole and the insistence of homogeneity of energy density outside of it, and is therefore completely consistent with our findings. It suggests that there might be other mechanisms for regulating the near horizon geometries of solutions which reduce to McVittie far away that might warrant further study. \\vskip.5cm \\smallskip {\\bf \\noindent Acknowledgements} \\smallskip We are grateful to B.~Carr, H.~C.~Cheng, A.~Gruzinov, M.~Luty, A.~Padilla, M.~Porrati and L.~Susskind for useful discussions, and to A.~Flachi for the help in obtaining a copy of the reference \\cite{vaidya}. We would also like to thank T.~Harada for questions on the previous version of the proof of finiteness of the $\\tau$-time of Eq. (\\ref{newtime}). NK is grateful to RESCEU, University of Tokyo, and in particular to J. Yokoyama, for a kind hospitality during the course of this work. The work of NK was supported in part by the DOE Grant DE-FG03-91ER40674. The work of MK is supported by NSF CAREER grant PHY-0645435. \\appendix" }, "1003/1003.1222_arXiv.txt": { "abstract": "{Ultraviolet radiation is a double-edged sword to life. If it is too strong, the terrestrial biological systems will be damaged. And if it is too weak, the synthesis of many biochemical compounds can not go along. We try to obtain the continuous ultraviolet habitable zones, and compare the ultraviolet habitable zones with the habitable zones of host stars. Using the boundary ultraviolet radiation of ultraviolet habitable zone, we calculate the ultraviolet habitable zones of host stars with masses from 0.08 to 4.00 \\mo. For the host stars with effective temperatures lower than 4,600 K, the ultraviolet habitable zones are closer than the habitable zones. For the host stars with effective temperatures higher than 7,137 K, the ultraviolet habitable zones are farther than the habitable zones. For hot subdwarf as a host star, the distance of the ultraviolet habitable zone is about ten times more than that of the habitable zone, which is not suitable for life existence. ", "introduction": "Typically, stellar habitable zone (HZ) is defined as a region near the host star where water at the surface of a terrestrial planet is in liquid phase, which has been widely researched (eg., Hart, 1978; Kasting et al., 1993; Franck et al., 2000; Noble et al., 2002). The boundary flux of HZ not only depends on luminosity, but also depends on effective temperature ($T_{\\rm eff}$) (eg., Forget \\& Pierrehumbert, 1997; Williams \\& Kasting, 1997; Mischna et al., 2000; Jones, 2004; Jones et al., 2006). As the higher $T_{\\rm eff}$, the less the infrared fraction in luminosity, and the less this fraction, the less the greenhouse effect for a given stellar flux (Jones et al., 2006). Thus, the distances at both the inner and the outer HZ boundaries are closer to host star, with higher $T_{\\rm eff}$, than they would have been if the $T_{\\rm eff}$ effect is not taken into consideration. However, others pointed out that life existence not only needs clement temperature, but also appropriate ultraviolet radiation (eg., Setlow \\& Doyle, 1954; Lindberg \\& Horneck, 1991; Cockell, 1998; Hoyle \\& Wickrasinghe, 2003; Sequra et al., 2003). UV radiation can induce DNA destruction and make life inactivate (Buccino et al., 2006; Tepfer \\& Leach, 2006). And UV radiation is also one of the most important energy source for the synthesis of many biochemical compounds on the primitive Earth (Buccino et al., 2006). The ``Principle of Mediocrity\" is in the ``hard core\" of all the research programs that search for life in the universe (Lakatos, 1974). In the points of this hypothesis, life and intelligence will develop with the same rules of natural selection wherever the proper conditions and the needed time are given (von Hoerner, 1961, 1973). In other words, the conditions that give place to the origin and evolution of life on Earth are average, in comparison to other worlds in the universe (Buccino et al., 2006). Using the ``Principle of Mediocrity\", Buccino et al. (2006) gave the boundary UV radiation of ultraviolet habitable zone (UV-HZ), but not the continuous UV-ZHs. Previously, it was pay attention to the HZs of host stars at main sequence (MS) phase, as the evolution from biochemical compounds to primary life needs very long time. However, life-seeds may migrate from one planet to another (Buccino et al., 2007). One hundred Myr may be too short for the pre-biological evolution, but the features of biology can change greatly in the same period, based on the ``Principle of Mediocrity\". Hence, it is meaningful to study the HZs and the UV-HZs of host stars at post-MS phase (eg., Franck et al., 2000; Noble et al., 2002). Hot subdwarfs are known as Extreme Horizontal Branch stars and believed to be core He-burning objects with extremely thin hydrogen envelopes (less than 0.02 \\mo). They are an important source of far-UV light in the galaxy and successfully used to explain the UV-upturn in elliptical galaxies (Kilkenny et al., 1997; Han et al. 2007). The typical core mass of a hot subdwarf is 0.475 \\mo, which can stably burn more than 160 Myr. Using the boundary UV radiation of UV-HZ (Buccino et al., 2006), we achieve the the distances at both the inner and the outer UV-HZ boundaries, as a function of stellar radius and $T_{\\rm eff}$. Using the data of stellar radius and $T_{\\rm eff}$ (Guo et al., 2009, hereafter Paper I), we calculate the UV-HZs around host stars with masses less than 4.00 \\mo, at zero age main sequence (ZAMS) and at the terminal of main sequence (TMS). Comparing the UV-HZs with the HZs (calculated in Paper I) of the same host stars, we find that the UV-HZs are near to the HZs for the host stars with $T_{\\rm eff}$ from 4,600 to 7,137 K. For the host stars with $T_{\\rm eff}$ lower than 4,600 K, the UV-HZs are closer than the HZs, which means that there is inadequate UV radiation in the HZs. For the host stars with $T_{\\rm eff}$ higher than 7,137 K, the UV-HZs are farther than the HZs, which means that there is too strong UV radiation and DNA may be destructed in the HZs. Using the evolutionary data of a hot subdwarf calculated by Zhang et al. (2009), we obtain both the HZ and the UV-HZ of the hot subdwarf with core mass 0.475 \\mob and envelope mass 0.001 \\mo. It is found that the UV-HZ is about ten times farther than the HZ for the hot subdwarf. This means that the UV radiation in the HZ is about one hundred times stronger than the UV radiation suited to life existence. Therefore, there is no chance that life survive in the HZs of hot subdwarfs, for the damaging UV radiation. The outline of the paper is as follows: we describe our methods in Section 2, show our results in Section 3, present some discussions in Section 4, and then finally in Section 5 we give our conclusions. ", "conclusions": "Firstly, we give the boundary distances of UV-HZ, as a function of stellar radius and $T_{\\rm eff}$. we obtain the UV-HZs around host stars with masses from 0.08 to 4.00 \\mo, and compare the UV-HZs with the HZs. The UV-HZs are closer than the HZs for the host stars with $T_{\\rm eff}$ lower than 4,600 K, and farther than the HZs for the host stars with $T_{\\rm eff}$ higher than 7,137 K. Secondly, we make out the impacts of the differences between HZs and UV-HZs on biological evolution. The UV radiation is very strong in the HZs for upper MS stars, which is negative to live existence. Thirdly, we give the HZ and the UV-HZ of a hot subdwarf, with core mass 0.475 \\mob and envelope mass 0.001 \\mo. The UV radiation in the HZ is from 22.257 to 135.696 times of the UV radiation dissociating DNA, which can easily kill lives in the HZ. Finally, we present discussions about the boundary UV radiation and X-ray and EUV radiation. One may also send any special request to \\it{guojianpo1982@hotmail.com} \\normalfont{or} \\it{guojianpo16@163.com}\\normalfont{.}" }, "1003/1003.3011_arXiv.txt": { "abstract": "{ We analyze the evolution of the perturbations in the inflaton field and metric following the end of inflation. We present accurate analytic approximations for the perturbations, showing that the coherent oscillations of the post-inflationary condensate necessarily break down long before any current phenomenological constraints require the universe to become radiation dominated. Further, the breakdown occurs on length-scales equivalent to the comoving post-inflationary horizon size. This work has implications for both the inflationary ``matching'' problem, and the possible generation of a stochastic gravitational wave background in the post-inflationary universe.\\vskip-.8cm} ", "introduction": "A complete inflationary scenario must provide a ``graceful exit'' from the phase of accelerated expansion and (re)-thermalize the universe, setting the stage for a hot big bang. In many models this process is efficiently driven by parametric resonance sourced by the oscillating inflaton field~\\cite{Traschen:1990sw,Kofman:1994rk,Kofman:1997yn}. The rich phenomenology of preheating is often contrasted with ``standard'' reheating, where the universe is rethermalized via the decay of the inflaton into standard matter. Given that the couplings responsible for the decay must be small in order to protect the inflaton potential against loop corrections, reheating is assumed to be a slow and gradual process (see e.g.~\\cite{Albrecht:1982mp}). Without parametric resonance, it might appear that reheating could take place almost arbitrarily slowly: we must only {\\em insist\\/} that the universe reheat to MeV scale temperatures, a scale about eighteen orders of magnitude below the characteristic energy at the end of GUT scale inflation, in order to be rethermalized in time for Big Bang Nucleosynthesis and to generate a cosmological neutrino background~\\cite{Komatsu:2010fb}. In standard single field scenarios the inflaton rolls slowly through the inflationary portion of the potential, and then performs rapid, damped oscillations about a minimum of the potential after inflation is over. The frequency of these oscillations is much higher than the expansion rate of the universe, so $\\rho$ is almost constant during a single oscillation. Conversely, the pressure $p$ fluctuates rapidly: at maximum amplitude $\\dot\\varphi =0$ and $\\rho= -p = V(\\varphi_{max})$ while at the bottom of the potential $V(\\varphi)=0$ and $\\rho = p = \\dot\\varphi^2/2$. The equation of state parameter, $w$, thus varies between $\\pm1$ during the oscillatory phase. Eventually, the amplitude of these oscillations will be small enough so that the potential is well approximated by the first term in its Taylor expansion around the minimum, $i.e$ a quadratic potential. The {\\em average\\/} equation of state can then be shown to be zero. This leads to the conclusion that in the post-inflationary universe the energy density decays like $1/a^3$, and the scale factor grows like $a(t) \\propto t^{2/3}$, mimicking a matter dominated universe. Since the energy density can decrease by 72 orders of magnitude before reheating must occur, the universe can expand by a factor of up to $10^{24}$ during this effectively matter dominated period.\\footnote{This is a conservative limit but the precise scale one adopts here has no impact on our conclusions.} In this paper, we carefully analyze the evolution of the perturbations in both the scalar field and the metric during this period of coherent oscillations. Given that there are two well-separated time scales in this system (the Hubble time and the oscillation period), we can obtain accurate analytic approximations for the field evolution by expanding in the ratio of the time scales. In the limit that the inflaton is coupled {\\em only\\/} to gravity and has a purely quadratic potential, we show that the perturbations grow and the system becomes nonlinear after the scale factor $a(t)$ has increased by a factor of about $10^6$. We expect that keeping higher order terms in the Taylor series expansion of the potential about its minimum lead to an even earlier onset of non-linearities, making this a conservative estimate. Consequently, while we are free to assume that the universe is matter dominated from the end of inflation to the MeV scale, the phase of coherent oscillations is necessarily much shorter. We see that the perturbations that become non-linear first are those for which the physical momentum becomes comparable to the expansion rate of the universe at the end of inflation. The coherent oscillations thus fragment on length scales roughly equal to the (comoving) horizon size at the end of inflation. We can understand our results heuristically by noting that in a universe dominated by pressureless dust, sub-horizon perturbations grow linearly with $a(t)$. Given the apparent red tilt of the primordial perturbation spectrum, perturbations on scales near the post-inflationary horizon volume have a lower initial amplitude than those at longer scales, but they also spend more time inside the horizon during this effective matter dominated era. The growth of these modes during this period more than compensates for the diminished initial amplitude, and shorter modes become nonlinear before longer modes. However, modes whose frequency is higher than the oscillation frequency of the coherent field ``resolve'' the oscillations. At this point our the analogy with a simple $p=0$ fluid breaks down, and modes with higher frequencies will be suppressed. Given the age and the simplicity of this model, these issues have been touched upon a number of times in the past. Nambu and Sasaki consider a related problem in the context of the invisible axion \\cite{Nambu:1989kh}, while Nambu and Taruya \\cite{Nambu:1996gf} write down the equations of motion for the inflaton perturbations and scalar metric fluctuations, showing that they are described by a Mathieu-like equation. Leach and Liddle \\cite{Leach:2000yw} compute the spectrum of perturbations for modes that leave the horizon near the end of inflation, while Assadullahi and Wands \\cite{Assadullahi:2009nf,Assadullahi:2009jc} show that nonlinearities formed during a generic matter dominated phase likely leads to gravitational wave production, and a high-frequency, stochastic background of gravitational waves in the present-day universe. Recently, Jedamzik, Lemoine and Martin \\cite{Jedamzik:2010dq} (see also \\cite{Jedamzik:2010hq}) discussed the gravitational growth of perturbations in this system, highlighting the parametric resonance ``hidden'' in the scalar field perturbations (see also \\cite{Nambu:1996gf}). The purpose of this analysis is to understand and clarify the connection between the detailed evolution of inflaton and metric perturbations. In particular, we carefully explore the basis of the simple analogy between the post-inflationary coherent oscillations and a dust-dominated universe. We show that the universe necessarily becomes inhomogeneous (and the perturbations nonlinear) if reheating is delayed long enough, and thus compute the maximum duration of a phase of coherent oscillations. We are careful to perform our post-inflationary analysis using quantities that are well-defined at all times: some perturbation variables which are well-defined during inflation have the background field velocity $\\dot{\\varphi}$ in the denominator, and thus become briefly singular during each oscillation of the field. The use of variables that remain smooth is particularly helpful for the numerical calculations we use to verify our analytic results. The layout of this paper is as follows. In Section 2, we briefly review linear perturbation theory for a single scalar field with canonical kinetic term minimally coupled to gravity. We focus on the appropriate choice of variables for this analysis, since the standard Mukhanov-Sasaki variable exhibits singularities for modes inside the horizon during the oscillatory phase. In Section 3, we study the evolution of perturbations that are near-horizon sized at the end of inflation for a quadratic potential. We calculate an appropriately defined power spectrum, recovering the results by Leach and Liddle \\cite{Leach:2000yw} where our calculation overlaps with theirs. Even though the underlying equation of state is rapidly oscillating, we find that the {\\em average\\/} growth of the density contrast of the modes is linear in the scale factor once they re-enter the horizon just like one would expect for a matter dominated universe. Finally, we compute when the first nonlinearities appear, verifying the estimate given above. We see that the first nonlinearities are found on scales comparable to the horizon size at the end of inflation. In Section 4, we discuss our results and highlight a number of open questions identified by our analysis. ", "conclusions": "We have given a full and careful account of the evolution of the universe during a phase of ``coherent oscillations'' following inflation. We have worked in variables that are manifestly finite, and given reliable analytic approximations for the individual modes' evolution. As is well known, the background solution is pressure-free when averaged on time scales long compared to the oscillation time of the fields. All modes grow linearly with the scale factor, matching our expectations for a dust-dominated universe with vanishing pressure. Modes whose natural frequency is much lower than the frequency of the scalar field oscillations during the end of inflation possess a nearly scale invariant power spectrum. For modes whose frequency during the end of inflation exceeds that of the background field oscillations, the power spectrum decreases rapidly with increasing frequency. The breakdown of the analogy between the coherent oscillations of the inflaton and pressureless dust reflects the presence of two intrinsic time scales in our dynamical system, namely the Hubble time and the oscillation period for the scalar field condensate. Conversely, a dust-dominated universe has only the Hubble time. Consequently, the perturbations of the scalar field match those of dust only when we can ``integrate out'' the more rapid scalar field oscillations. Interestingly, the ``long modes'' for which the pressureless dust approximation is reliable can be shown to undergo parametric resonance while they are outside the horizon~\\cite{Jedamzik:2010dq}, and this resonance is the microphysical origin of the constancy of $\\mathcal{R}_k$, as we show via an explicit calculation. Phenomenologically, this distinction between ``long'' and ``short'' modes picks out a unique scale in the post-inflationary universe. At the end of inflation, modes which are approximately horizon-sized have a frequency similar to the oscillations of the coherent inflaton background. Solutions for modes well inside the horizon at the end of inflation decay in amplitude until they are redshifted to lower frequencies by the expansion of the universe. Conversely, modes that left the horizon well before the end of inflation are frozen until they re-enter the horizon. Consequently, modes that are roughly horizon sized at the end of inflation undergo significantly more growth than other modes and will be the first for which nonlinear effects become important, and the breakdown of the coherent oscillations takes place on length scales fixed by the post-inflationary horizon size. Physically, the breakdown of coherence in the oscillating field implies that while we can assume a very long period of matter domination following inflation, the universe can expand by at most a factor $\\sim e^{13}$ or so during the coherent oscillation phase. Beyond this point, $\\delta \\rho/\\rho \\sim 1$, and the field is no longer homogeneous. Once this breakdown occurs, we need to move beyond first order perturbation theory, raising the prospect that even the simplest post-inflationary universe can have highly nontrivial phenomenology. This could include gravitational wave production \\cite{Assadullahi:2009nf,Assadullahi:2009jc,Jedamzik:2010hq}, or even the formation of primordial black holes (e.g. \\cite{1985MNRAS.215..575K}) if the overdensities grow without limit. In the latter case, the primordial black holes will be very small, since their mass will be set by the mass-energy inside a region whose size is equal to the comoving length scale at the end of inflation. These black holes would thus decay quickly via Hawking radiation, reheating the universe (e.g. \\cite{GarciaBellido:1996qt}) via ``non-perturbative'' gravitational effects. Moreover, in addition to any gravitational waves generated via the nonlinear evolution of the field, gravitons Hawking radiated as the black holes decay would generate a background of very high-frequency gravitational waves \\cite{Anantua:2008am}. More generally, this analysis reinforces the importance of carefully exploring the physical processes which govern the universe between inflationary and nucleosynthesis scales. Models with parametric resonance are well known for their rich phenomenology and possible gravitational wave background (e.g. \\cite{Easther:2006gt,Easther:2006vd,Easther:2007vj,GarciaBellido:2007af,Dufaux:2007pt}, but it is becoming increasingly clear that even the simplest potentials can have nontrivial post-inflationary dynamics. We also wish to understand the overall thermal history of the post-inflationary epoch in order to accurately match present-day astrophysical scales to their counterparts during the inflationary era \\cite{Liddle:2003as}. Consequently, while this model is almost as old as inflation itself \\cite{Linde:1983gd}, its phenomenology contains a number of important open questions. The analysis here provides a thorough and complete understanding of the post-inflationary oscillations, and we intend to pursue the issues raised here in future work." }, "1003/1003.0965_arXiv.txt": { "abstract": "One of the aims of the Low Frequency Array (LOFAR) Epoch of Reionization (EoR) project is to measure the power spectrum of variations in the intensity of redshifted 21-cm radiation from the EoR. The sensitivity with which this power spectrum can be estimated depends on the level of thermal noise and sample variance, and also on the systematic errors arising from the extraction process, in particular from the subtraction of foreground contamination. We model the extraction process using realistic simulations of the cosmological signal, the foregrounds and noise, and so estimate the sensitivity of the LOFAR EoR experiment to the redshifted 21-cm power spectrum. Detection of emission from the EoR should be possible within 360 hours of observation with a single station beam. Integrating for longer, and synthesizing multiple station beams within the primary (tile) beam, then enables us to extract progressively more accurate estimates of the power at a greater range of scales and redshifts. We discuss different observational strategies which compromise between depth of observation, sky coverage and frequency coverage. A plan in which lower frequencies receive a larger fraction of the time appears to be promising. We also study the nature of the bias which foreground fitting errors induce on the inferred power spectrum, and discuss how to reduce and correct for this bias. The angular and line-of-sight power spectra have different merits in this respect, and we suggest considering them separately in the analysis of LOFAR data. ", "introduction": "\\label{sec:intro} Studying 21-cm radiation from hydrogen at high redshifts \\citep*{FIE58,FIE59,HOG79,SCO90,KUM95,MAD97} promises to be interesting for several reasons. Fluctuations in intensity are sourced partly by density fluctuations, measurements of which may allow rather tight constraints on cosmological parameters \\citep{MAO08}. They are also sourced by variations in the temperature and ionized fraction of the gas, which means that 21-cm studies may provide information on early sources of ionization and heating, such as stars or mini-QSOs. The period during which the gas undergoes the transition from being largely neutral to largely ionized is known as the Epoch of Reionization \\citep[EoR; e.g.][]{LOE01,BEN06,FOB06}, while the period beforehand is sometimes known as the cosmic dark ages. While the latter has perhaps the best potential to give clean constraints on cosmology, the instruments becoming available in the near future are not expected to be sensitive enough at the appropriate frequencies to study this epoch interferometrically. Several, though, are hoped to be able to study the EoR (e.g.\\ GMRT,\\footnotemark\\ MWA,\\footnotemark\\ LOFAR,\\footnotemark\\ 21CMA,\\footnotemark\\ PAPER,\\footnotemark\\ SKA\\footnotemark), but even so, their sensitivity is not expected to be sufficient to make high signal-to-noise images of the 21-cm emission in the very near future. We seek, instead, a statistical detection of a cosmological 21-cm signal, with the most widely studied statistic being the power spectrum (e.g. \\citealt{MOR04}; \\citealt{BAR05}; \\citealt{MCQ06}; \\nocite{BMH06,BOW07}Bowman, Morales \\& Hewitt 2006, 2007; \\citealt{PRI07}; \\citealt{BAR09}; \\citealt{LID08}; \\citealt{PRI08}; \\citealt{SET08}). Our aim in this paper is to test how well the 21-cm power spectrum can be extracted from data collected with the Low Frequency Array (LOFAR), which is currently under construction. While this is a general-purpose observatory, the EoR project, being one of LOFAR's Key Science Projects, has helped to drive the design of the instrument. We give some details on parameters of the instrument which are relevant to EoR observations in Section~\\ref{subsec:instresp}. The quality of extraction is affected by several factors: the observational strategy and the length of observations, which affect the volume being studied and the level of thermal noise; the array design and layout; the foregrounds from Galactic and extragalactic sources, and the methods used to remove their influence from the data (presumably by exploiting their assumed smoothness as a function of frequency; see e.g. \\citealt{SHA99}; \\citealt{DIM02}; \\citealt{OH03}; \\citealt*{ZAL04}); excision of radio-frequency interference (RFI) and radio recombination lines; and, for example, the quality of polarization and total intensity calibration for instrumental and ionospheric effects. We will not study RFI or calibration here. We will, however, use simulations of the cosmological signal (CS), the foregrounds, the instrumental response and the noise to generate synthetic data cubes -- i.e.\\ the intensity of 21-cm emission as a function of position on the sky and observing frequency -- and then attempt to extract the 21-cm power spectrum from these cubes. We generate data cubes realistic enough so that we can test different observing strategies and methods of subtracting the foregrounds, and look at the effect on the inferred power spectrum. \\footnotetext[1]{Giant Metrewave Telescope, http://www.gmrt.ncra.tifr.res.in/}\\footnotetext[2]{Murchison Widefield Array, http://www.haystack.mit.edu/ast/arrays/mwa/}\\footnotetext[3]{Low Frequency Array, http://www.lofar.org/}\\footnotetext[4]{21 Centimeter Array, http://web.phys.cmu.edu/\\~{}past/}\\footnotetext[5]{Precision Array to Probe the EoR, http://astro.berkeley.edu/\\~{}dbacker/eor/}\\footnotetext[6]{Square Kilometre Array, http://www.skatelescope.org/} We devote the following section to describing the construction of the data cubes and giving a brief description of their constituent parts. Then, in Section~\\ref{sec:ext} we discuss the extraction of the 21-cm power spectrum from the cubes, including our method for subtracting the foregrounds. In Section~\\ref{sec:sens} we present our estimates of the sensitivity of LOFAR to the 21-cm power spectrum, and discuss the character of the statistical and systematic errors on these estimates. We conclude in Section~\\ref{sec:conc} by offering some thoughts on what these results suggest about the merits of different observing strategies and extraction techniques. ", "conclusions": "\\label{sec:conc} In this paper we have studied the extraction of the 21-cm EoR power spectrum from simulated LOFAR data. The simulations allow us to compute the statistical errors on the power spectrum due to thermal noise and sample variance, and these are small enough to raise the possibility of a significant detection of emission from the EoR using only a modest amount of observing time. If we wish to estimate the power spectrum accurately, however, this becomes more challenging once we take into account the presence of fitting errors from the subtraction of astrophysical foregrounds. These errors are correlated (positively or negatively) with the signal and the noise in general, and introduce a scale-dependent bias into our estimate of the power spectrum. We anticipate that simulations such as the ones studied here could be used to estimate and correct for the bias; this would induce a further statistical error which can be straightforwardly computed by using multiple realizations of a simulated observation. Making this sort of correction will always be uncertain, though, so it is desirable to minimize its size. We have looked at the extent to which the size of the correction, as well as the size of the statistical errors, can be reduced by observing for longer or using alternative observational strategies. Before that, though, we tested that extraction is still possible if we do not make the assumption that the {\\it uv} coverage is independent of frequency. We find that this necessitates fitting the foregrounds in the $(u,v,\\nu)$ cube rather than the image cube, as noted by \\citet{LIU09b}. The Wp smoothing method, which we have used previously to fit the foregrounds in the image cube, can be adapted to work in the $(u,v,\\nu)$ cube by fitting the real and imaginary parts independently for each {\\it uv} cell and by varying the regularization parameter, $\\lambda$, across the {\\it uv} plane. This yields results comparable to (in fact, even better than) those we obtain if we assume frequency-independent {\\it uv} coverage and then fit in the image cube. We have also tried using a third-order polynomial to fit the foregrounds in the $(u,v,\\nu)$ cube: this yields results which are acceptable, but not as good as those obtained using Wp smoothing. The main drawback of Wp smoothing in this case is its speed, especially for `lines of sight' near the centre of the {\\it uv} plane where it is best to choose a small value for $\\lambda$ (implying little smoothing). Because Wp smoothing in the image cube is faster, because the polynomial fitting gives worse results than Wp smoothing in the $(u,v,\\nu)$ cube, and because Wp smoothing produces extraction of similar quality in the image and $(u,v,\\nu)$ cubes, we have concentrated on results using frequency-independent {\\it uv} coverage to explore the different scenarios in this paper. We have found that a year's observations (of, say, 600 hours, of which perhaps 360 could be of a single window) should be sufficient to detect cosmological 21-cm emission from towards the end of the EoR. We caution, however, that the approximations employed in this paper prevent us from treating these numbers as more than rough estimates. If we wish to study the power spectrum at small or large scales -- away from the `sweet spot' at intermediate $k$ -- it will be important to be able to synthesize multiple station beams. This allows us to reduce the statistical errors from sample variance and noise. Unfortunately, however, there appears to be no substitute for extending the integration time, especially to probe high redshifts and very small scales. This is because only deep observations can improve the quality of the foreground fitting, and hence reduce the systematic offset between the true signal and the recovered signal. Under the optimistic assumptions that we can synthesize six beams, and that the useful frequency range can be covered using just two frequency bands (the instantaneous frequency coverage is limited), 600 hours of observation of a single window should be enough to yield quite precise and accurate power spectra up to $z\\approx 9$, for $k$ between approximately 0.03 and $0.6\\ h\\ \\mathrm{Mpc}^{-1}$. Pushing to the very highest redshifts accessible with the frequency coverage of LOFAR's high band antennas requires somewhat longer: perhaps 900 hours per frequency band, which corresponds to 1800 hours of observation if there are two frequency bands. With observations of this depth, the limiting factor in the statistical errors comes from sample variance on large scales, which can only be reduced by observing a larger area of sky. This is one of several reasons why the LOFAR EoR project plans to observe multiple -- perhaps five -- independent windows. We have already seen that approximately 600 hours per window is required for the thermal noise errors to be small and the bias to be under control for redshifts less than about 9. For five windows, this corresponds to 3000 hours of observation. Comparing the independent windows will also allow important cross-checks, in particular that systematics are under control. To really push towards precise constraints on the power spectrum towards the start of reionization, the 1800 hours per window that we find yelds high quality extraction at $z>10$ corresponds to 9000 total hours for five windows. This figure may be reduced if a hybrid strategy, in which we integrate for a longer time in lower frequency bands, turns out to be feasible. From the point of view of foreground fitting and power spectrum extraction, ignoring constraints that may be imposed by calibration etc., a hybrid strategy does indeed seem to be feasible. Of course, we have considered this strategy only from the point of view of the power spectrum. If deeper observations at all frequencies would allow us to push beyond the power spectrum, perhaps into a regime where we can observe individual features in the distribution of 21-cm emission towards the end of reionization with reasonable signal to noise, then this would surely be valuable too. Other hybrid strategies are also possible, for example ones in which different windows are observed for different amounts of time. We have not studied them here since they do not really impact the fitting and extraction, which is independent for each window. None the less, they may allow us to obtain high redshift constraints by observing one window deeply, while simultaneously allowing us to beat down sample variance errors on large scales at low redshifts by observing several other windows at reduced depth. In any case, our study suggests that as the amount of time spent observing the EoR with LOFAR is increased, this allows us to make qualitative improvements to the fitting, and to the range of scales and redshifts we can probe accurately. Deeper integration does more than simply allow us to shrink our statistical error bars. This all depends, however, on the robustness of our fitting techniques, and more generally on the level of control we are able to exercise over systematic errors. The Wp smoothing method we have introduced previously appears to work well when it comes to extracting the power spectrum. This holds whether we apply it to an idealized case in which the {\\it uv} coverage of the instrument is constant with frequency, or to a more realistic case in which it varies. We confirm a suspicion we have expressed previously \\citep{NONPAR_09} that the power spectrum may be easier to extract than an apparently simpler statistic such as the rms of the 21-cm signal: the fitting errors are scale-dependent, and a power spectrum analysis allows us to pick out the scales where our method works best without being swamped by small-scale noise. Splitting the power spectrum into angular and line-of-sight components may help us to test the robustness of our conclusions, and perhaps extend the spatial dynamic range we can probe. We have assumed here that the power spectrum of the noise is known to reasonable accuracy, an assumption which will be examined in future work. We will also study in a future paper how different strategies alter our ability to constrain the parameters of reionization models. Finally, we note that foreground fitting and power spectrum extraction are late steps in the collection and analysis of LOFAR EoR data. They depend on earlier and probably more difficult steps, such as instrumental calibration (including polarization, which we have neglected here), correcting for the ionosphere, and the excision of RFI. The results of this paper only reassure us that the later stages are unlikely to be the limiting ones." }, "1003/1003.0681_arXiv.txt": { "abstract": "We report the discovery of kinematic shock signatures associated with a localized radio jet interaction in the merging Seyfert galaxy NGC 5929. We explore the velocity-dependent ionization structure of the gas and find that low ionization gas at the interaction site is significantly more disturbed than high ionization gas, which we attribute to a local enhancement of shock ionization due to the influence of the jet. The characteristic width of the broad low-ionization emission is consistent with shock velocities predicted from the ionization conditions of the gas. We interpret the relative prominence of shocks to the high density of gas in nuclear environment of the galaxy and place some constraints of their importance as feedback mechanisms in Seyferts. ", "introduction": "\\label{sec1} Most Seyfert AGN are associated with weak nuclear radio sources which show radio spectral indices and morphologies (if resolved) consistent with synchrotron-emitting jets and lobes \\citep{ulv84, nagar99}, though their physical properties are poorly constrained. However, they are effective avenues of kinetic and thermal feedback from the active nucleus and may play an important role in determining the evolution of the central spheroid. In the small fraction of Seyferts with kpc-scale radio jets, several studies have uncovered clear signatures of interactions between the jet and surrounding gas, in the form of disturbed emission line profiles in the inner Narrow-Line Region (NLR), as well as close associations between resolved jet structures and NLR gas \\citep[e.g.][]{whittle88,capetti96,fws98,cooke00,cecil02,whittle04}. Depending on the physical make-up of the jet, relativistic or ram pressure can drive fast shocks, compressing, sweeping up and altering the appearance of the NLR. Postshock gas, with a temperature of several $10^7$ K, is a source of ionizing photons which couple the shock properties to the ionization state of the surrounding gas \\citep[e.g.][]{ds96}. Shocks and winds can also change the distribution of ISM phases in the NLR by destroying and ablating clouds \\citep{fragile05}. While the active nucleus is usually the dominant source of ionization even in strongly jetted Seyferts \\citep[e.g.][]{whittle05}, widespread shocks driven by a jet can alter the ionization of the NLR by affecting the properties of the ISM. Studies of shock structure and energetics \\citep{ds96, allen08} predict strong differences between the emission line spectrum of dense post-shock gas and gas that is ionized by, but not in direct dynamical contact with, the shock (the precursor). Therefore, a clear signature of shock ionized gas is a difference between the line profiles of low and high ionization lines, which are preferentially produced by post-shock and precursor gas, respectively \\citep{whittle05}. In this work, we present an HST/STIS spectroscopic study of NGC 5929, a local Seyfert galaxy with a well-studied bi-polar radio jet. Previous ground-based spectroscopic studies find evidence of a localized interaction between the jet and the near-nuclear emission line gas \\citep{whittle86,wakamatsu88,taylor89,wilson89,ferruit97}. The datasets and analysis methods used are briefly reviewed in $\\S2$. Direct shock features and a picture of the interaction are developed in $\\S3$ and $\\S4$. We discuss the role of shocks in Seyferts and AGN feedback in $\\S5$. NGC 5929 has a systemic heliocentric velocity of $cz\\,=\\,2492$ km s$^{-1}$ based on the stellar aborption line redshift of \\citet{nelsonnwhittle95}, which corresponds to $161$ pc arcsec$^{-1}$ (H$_0 = 75$ km s$^{-1}$ Mpc$^{-1}$). \\begin{figure*}[ht] \\label{corrplot} \\centering \\includegraphics[width=0.7\\columnwidth,angle=270]{ngc5929_corrplot.eps} \\caption[NGC 5929: Combined datasets] { Panels of images and spectra of the nuclear region of NGC 5929. The 2D spectra in Panels 2 and 3 correspond to STIS long-slit apertures A and B respectively, as indicated on the images plotted in Panels 1 and 4. The radio map (gaussian smoothed by $0\\farcs3$ to bring out its structure) is plotted in Panel 1 in contours. } \\end{figure*} ", "conclusions": "\\label{sec4} Why do we see such obvious signatures of shocks in this object and not in others? This may be due to the weakness of the AGN: NGC 5929 has an absorption-corrected hard X-ray luminosity of $1.8\\times 10^{41}$ erg s$^{-1}$ -- low compared to average local Seyferts \\citep{cardamone07}. Or perhaps the nuclear environment of the galaxy is dense and gas-rich due to its ongoing merger. This will lower the ionization parameter of nuclear radiation and produce a more compact emission line region. In both scenarios, the influence of the nucleus at larger radii will be relatively unimportant, making shock ionized line emission visible against the general background of centrally ionized gas. If this is indeed the case, it implies that radiative shock ionization from nuclear outflows is widespread in Seyfert NLRs, but is usually secondary to nuclear photoionization processes and only visible in low-luminosity Seyferts or those with dense nuclear environments. A rich nuclear environment raises another possibility that the jet is impacting a dense molecular cloud, enhancing the shock luminosity. This may explain why no clear shock signatures are seen around the NE radio hotspot, though the obscuration of the main dust lane prevents a direct view of this region. The luminosity of the broad H$\\beta$ component is $8\\times 10^{38}$ erg s$^{-1}$. Following \\citet{ds96}, the H$\\beta$ flux scales with the total radiative flux from the shock, with a weak dependence on $V_{sh}$, giving shock luminosities of $5 \\times 10^{41}$ erg s$^{-1}$ -- comparable or slightly less than the total luminous output of the AGN (taking a X-ray bolometric correction of $\\sim 10$). Using standard relationships from \\citet{osterbrock89}, the H$\\beta$ luminosity can be used to derive the mass of ionized gas in the broad component: $3 \\times 10^{4}$ M$_{\\odot}$. If this mass of gas was accelerated to $V_{sh}$, it would have a total kinetic energy of $5 \\times 10^{52}$ ergs. Taking the approximate acceleration timescale to be the crossing time of the region of size 0\\farcs3 (48 pc) at $V_{sh}$ (around $10^{5}$ yr), a lower limit on the `kinetic luminosity' of the jet is estimated to be $1.5 \\times 10^{40}$ erg s$^{-1}$. This is few to several percent of the total AGN energy output. Given that jet outflows are a relatively common feature of Seyfert activity and couple strongly to the ISM through shocks, jet driven feedback can effectively carry as much as energy as the AGN photon luminosity to kpc scales, and transfer a tenth or less of this energy in the form of kinetic energy to the NLR. This can have important consequences for the suppression of bulge star-formation and the energy budget and dynamics of circum-nuclear gas." }, "1003/1003.5549.txt": { "abstract": "A Bayesian analysis of 47 Ursae Majoris (47 UMa) radial velocity data confirms and refines the properties of two previously reported planets with periods of 1079 and 2325 days. The analysis also provides orbital constraints on an additional long period planet with a period $\\sim 10000$ days. The three planet model is found to be $~ 10^{5}$ times more probable than the next most probable model which is a two planet model. The nonlinear model fitting is accomplished with a new hybrid Markov chain Monte Carlo (HMCMC) algorithm which incorporates parallel tempering, simulated annealing and genetic crossover operations. Each of these features facilitate the detection of a global minimum in $\\chi^2$. By combining all three, the HMCMC greatly increases the probability of realizing this goal. When applied to the Kepler problem it acts as a powerful multi-planet Kepler periodogram. The measured periods are $1078\\pm2$, $2391_{-87}^{+100}$, and $14002_{-5095}^{+4018}$d, and the corresponding eccentricities are $0.032\\pm 0.014$, $0.098_{-.096}^{+.047}$, and $0.16_{-.16}^{+.09}$. The results favor low eccentricity orbits for all three. Assuming the three signals (each one consistent with a Keplerian orbit) are caused by planets, the corresponding limits on planetary mass ($M \\sin i$) and semi-major axis are \\\\ ($2.53_{-.06}^{+.07} M_J$, $2.10\\pm 0.02\\rm{au}$), ($0.54\\pm 0.07 M_J$, $3.6\\pm 0.1\\rm{au}$), and ($1.6_{-0.5}^{+0.3} M_J$, $11.6_{-2.9}^{+2.1}\\rm{au}$), respectively. Based on a three planet model, the remaining unaccounted for noise (stellar jitter) is $5.7$m s$^{-1}$. The velocities of model fit residuals were randomized in multiple trials and processed using a one planet version of the HMCMC Kepler periodogram. In this situation periodogram peaks are purely the result of the effective noise. The orbits corresponding to these noise induced periodogram peaks exhibit a well defined strong statistical bias towards high eccentricity. We have characterized this eccentricity bias and designed a correction filter that can be used as an alternate prior for eccentricity, to enhance the detection of planetary orbits of low or moderate eccentricity. ", "introduction": "\\label{sec:introduction} Improvements in precision radial velocity (RV) measurements and continued monitoring are permitting the detection of lower amplitude planetary signatures. One example of the fruits of this work is the detection of a super earth in the habitable zone surrounding Gliese 581 \\citep{Udry2007}. This and other remarkable successes on the part of the observers is motivating a significant effort to improve the statistical tools for analyzing radial velocity data (e.g., \\citealt{LoredoChernoff2003}, \\citealt{Loredo2004}, \\citealt{Cumming2004}, Gregory 2005a \\& b, Ford 2005 \\& 2006, \\citealt{FordGregory2006}, \\citealt{CummingDragomir2010}). Much of the recent work has highlighted a Bayesian MCMC approach as a way to better understand parameter uncertainties and degeneracies and to compute model probabilities. Gregory (2005a, b, c and 2007a, b, c) presented a Bayesian MCMC algorithm that makes use of parallel tempering (PT) to efficiently explore a large model parameter space starting from a random location. It is able to identify any significant periodic signal component in the data that satisfies Kepler's laws and thus functions as a Kepler periodogram~\\footnote{Following on from the pioneering work on Bayesian periodograms by \\citet{Jaynes1987} and \\citet{Brett1988}}. This eliminates the need for a separate periodogram search for trial orbital periods which typically assume a sinusoidal model for the signal that is only correct for a circular orbit. In addition, the Bayesian MCMC algorithm provides full marginal parameters distributions for all the orbital elements that can be determined from radial velocity data. The algorithm includes an innovative two stage adaptive control system that automates the selection of efficient Gaussian parameter proposal distributions. The latest version of the algorithm, \\citet{Gregory2009}, incorporates a genetic crossover operation into the MCMC algorithm. The new adaptive hybrid MCMC algorithm (HMCMC) incorporates parallel tempering, simulated annealing and genetic crossover operations. Each of these techniques was designed to facilitate the detection of a global minimum in $\\chi^2$. Combining all three in an adaptive hybrid MCMC greatly increase the probability of realizing this goal. \\citet{ButlerMarcy1996} first reported a 1090 day companion to 47 UMa using data from Lick Observatory. With additional velocity measurements over 13 yr, \\citet{Fischer2002} announced a long-period second planet, 47 UMa c, with a period of $2594\\pm90$ days and a mass of $0.76 M_J$. \\citet{Naef2004} reported observations from the fiber fed echelle spectrograph ELODIE of 47 UMa, and noted that the second planet was not evident in their data. \\citet{Wittenmyer2007} reported that there is still substantial ambiguity as to the orbital parameters of the proposed planetary companion 47 UMa c. They gave a period of 7586 day for one orbital solution, and 2594 day for two others. In their latest work \\citet{Wittenmyer2009}, their best-fit 2-planet model now calls for $P_2 = 9660$ days. In this paper we present a Bayesian analysis of the latest Lick observatory measurements and a combined Lick plus McDonald Observatory (\\citealt{Wittenmyer2009}) data set. We also report on an investigation of the behavior of the Bayesian HMCMC Kepler periodogram to noise. The noise data sets were formed by randomly interchanging velocity measurements. ", "conclusions": "In this paper, we have demonstrated that a Bayesian adaptive hybrid MCMC (HMCMC) analysis of a challenging data set has helped clarify the number of planets present in 47UMa. HMCMC integrates the advantages of parallel tempering, simulated annealing and the genetic algorithm. Each of these techniques was designed to facilitate the detection of a global minimum in $\\chi^2$. Combining all three in an adaptive hybrid MCMC greatly increase the probability of realizing this goal. The adaptive Bayesian hybrid MCMC is very general and can be applied to many different nonlinear modeling problems. It has been implemented in gridMathematica on an 8 core PC. The increase in a speed for the parallel implementation is a factor 6.6. When applied to the Kepler problem it corresponds to a multi-planet Kepler periodogram which is ideally suited for detecting signals that are consistent with Kepler's laws. However, it is more than a periodogram because it also provides full marginal posterior distributions for all the orbital parameters that can be extracted from radial velocity data. The execution time for a 1 planet blind fit (7 parameters) is $10^6$ iterations per hr. The program scales linearly with the number of parameters being fit. The 47UMa data has been analyzed using 1, 2, 3, 4, and 5 planet models. On the basis of the model selection results we can conclude there is strong evidence for three planets based on a Bayesian false alarm probability of $5.0 \\times 10^{-6}$, however, the longest period orbital parameters are still not well defined. The measured periods (based on the combined data set) are $1078\\pm 2$, $2391_{-87}^{+100}$, and $14002_{-5095}^{+4018}$d, and the corresponding eccentricities are $0.032\\pm 0.014$, $0.098_{-.096}^{+.047}$, and $0.16_{-.16}^{+.09}$. The results favor low eccentricity orbits for all three. Note: the longer time base of the full Lick data set favors a value for $P_3$ at the lower end of the 68\\% credible region of $\\sim 10,000$ days. Assuming the three signals (each one consistent with a Keplerian orbit) are caused by planets, the corresponding limits on planetary mass ($M \\sin i$) and semi-major axis are ($2.53_{-.06}^{+.07} M_J$, $2.10\\pm 0.02\\rm{au}$),\\\\ ($0.54\\pm 0.07 M_J$, $3.6\\pm 0.1\\rm{au}$), and ($1.6_{-0.5}^{+0.3} M_J$, $11.6_{-2.9}^{+2.1}\\rm{au}$), respectively. Based on our three planet model results, the remaining unaccounted for noise (stellar jitter) is $5.7$m s$^{-1}$. A four planet model fit to the Lick data yielded a well defined fourth period of $370.8_{-2.0}^{+2.4}$ days and eccentricity of $0.57_{-0.15}^{+0.22}$, but the combined data set did not yield a well defined fourth period. Even though this period was well defined in the Lick only data, the Bayesian false alarm probability of $\\approx 0.5$ is much too high to warrant any claim of significance. The period is also suspiciously close to one year and might be an artefact of the data reduction. The velocities of model fit residuals were randomized in multiple trials and processed using a one planet version of the HMCMC Kepler periodogram. In this situation periodogram probability peaks are purely the result of the effective noise. The orbits corresponding to these noise induced periodogram peaks exhibit a well defined statistical bias towards high eccentricity. We have characterized this eccentricity bias and designed a correction filter that can be used as an alternate prior for eccentricity to enhance the detection of planetary orbits of low or moderate eccentricity. On the basis of our understanding of the mechanism underlying the eccentricity bias, we expect the eccentricity prior filter to be generally applicable to searches for low amplitude orbital signals in other precision radial velocity data sets." }, "1003/1003.0758_arXiv.txt": { "abstract": "{} {We develop a method for deriving distances from spectroscopic data and obtaining full 6D phase-space coordinates for the RAVE survey's second data release.} {We used stellar models combined with atmospheric properties from RAVE (effective temperature, surface gravity and metallicity) and $\\jk$ photometry from archival sources to derive absolute magnitudes. In combination with apparent magnitudes, sky coordinates, proper motions from a variety of sources and radial velocities from RAVE, we are able to derive the full 6D phase-space coordinates for a large sample of RAVE stars. This method is tested with artificial data, Hipparcos trigonometric parallaxes and observations of the open cluster M67.} {When we applied our method to a set of $\\raveApply$ stars, we found that 25\\% ($\\raveMagc$) of the stars have relative (statistical) distance errors of $< \\raveMagcperc$\\%, while 50\\% ($\\raveMagb$) and 75\\% ($\\raveMaga$) have relative (statistical) errors smaller than \\raveMagbperc{}\\% and \\raveMagaperc{}\\%, respectively. Our various tests show that we can reliably estimate distances for main-sequence stars, but there is an indication of potential systematic problems with giant stars owing to uncertainties in the underlying stellar models. For the main-sequence star sample (defined as those with $\\logg{} > 4$), 25\\% ($\\raveMagcMS$) have relative distance errors $< \\raveMagcMSperc\\%$, while 50\\% ($\\raveMagbMS$) and 75\\% ($\\raveMagaMS$) have relative errors smaller than \\raveMagbMSperc{}\\% and \\raveMagaMSperc{}\\%, respectively. Our full dataset shows the expected decrease in the metallicity of stars as a function of distance from the Galactic plane. The known kinematic substructures in the $U$ and $V$ velocity components of nearby dwarf stars are apparent in our dataset, confirming the accuracy of our data and the reliability of our technique. We provide independent measurements of the orientation of the $UV$ velocity ellipsoid and of the solar motion, and they are in very good agreement with previous work.}{The distance catalogue for the RAVE second data release is available at {\\sf http://www.astro.rug.nl/$\\sim$rave}, and will be updated in the future to include new data releases.} ", "introduction": "The spatial and kinematic distributions of stars in our Galaxy contain a wealth of information about its current properties, its history and evolution. This phase-space distribution is a crucial ingredient if we are to build and test dynamical models of the Milky Way \\citep[e.g.][and references therein]{Bi2005}. More directly, the kinematics of halo stars can be used to trace the Galaxy's accretion history \\citep{HelmiWhite1999}, as has been shown to good effect in many subsequent studies \\citep[e.g.][]{Helmi1999,Kepley2007,Smith2009}. There is also much to learn from the phase-space structure of the disk, where it is possible to identify substructures due to both accretion events and dynamical resonances \\citep[e.g.][]{Dehnen2000,Famaey2005,Helmi2006} or learn about the mixing processes that influence the chemical evolution of the disk \\citep[e.g.][]{Ro2008,Sch2008}. To fully exploit this rich resource, we need to analyse the full six-dimensional phase-space distribution, which clearly cannot be done without a reliable estimate of the distances to the stars under consideration. Therefore obtaining accurate distances and velocities for a representative sample of stars in our Galaxy will be essential if we are to understand both the structure of our own Galaxy and galaxy formation in general. The most dramatic recent development in this field was the Hipparcos satellite mission \\citep{ESA1997,2000A&A...355L..27H}, which carried out an astrometric survey of stars down to $V \\sim 12$ mag with accuracies of up to 1 mas. This catalogue enabled the distances of $\\sim 10,000$ stars to be measured using the trigonometric parallax technique, with parallax errors of less than 5\\% \\citep{vanLeeuwen2007a,vanLeeuwen2007b}. However, in general the resulting parallaxes only probe out to a couple of hundred parsec and are limited to the brightest stars. This limitation of the trigonometric parallax method led researchers to attempt other techniques for calculating distances. One promising avenue is the study of pulsating variable stars, such as RR Lyraes or Cepheids, for which it is possible to accurately determine distances using period-luminosity relations \\citep[see, for example, the reviews of][]{Ga1995,Ga1996}. These have been used effectively to probe the structure of our Galaxy, in particular the study of the old and relatively metal-poor RR Lyrae stars \\citep{Vi2001,Ku2008,Wa2009}. Although pulsating variables can provide accurate tracer populations, the numbers of such stars is clearly limited; ideally we would like to determine distances for large numbers of stars and not just specific populations. As a consequence there have been numerous studies utilising photometric distance determinations, where one estimates the absolute magnitude of a star from its colour. The efficacy of this method can be seen from the work of \\citet{Siegel2002} and \\citet{Juric2008}, who both used this technique to model the stellar density distribution of the Galaxy. Another striking example of the power of this technique was presented by \\citet{Belokurov2006}, where halo turn-off stars were used to illuminate a host of substructures in the Galactic halo. The strength of photometric distances is that they can be constructed for a wide range of stellar populations. An important recent study was carried out by \\citet{Ivezic2008}. In this work they took high-precision multi-band optical photometry from the Sloan Digital Sky Survey \\citep[SDSS;][]{Ab2008} and constructed a photometric distance relation for F- and G-type dwarfs, using colours to identify main-sequence stars and estimate metallicity. Globular clusters were used to calibrate their photometric relation, resulting in distance estimates accurate to $\\sim15$ per cent. This is only possible due to the extremely well-calibrated SDSS photometry and, in any case, is only applicable to F- and G-type dwarfs. To determine distances for entire surveys (with a wide range of different stellar classes and populations) requires complex multi-dimensional algorithms. In this paper we develop such a technique to estimate distances for stars using photometry in combination with stellar atmosphere parameters derived from spectra. One of the motivations behind our study is so that we can complement the Radial Velocity Experiment \\citep[RAVE][]{Steinmetz2006,Zwitter2008}. This project, which started in 2003, is currently measuring radial velocities and stellar atmosphere parameters (temperature, metallicity and surface gravity) for stars in the magnitude range 9~$<$~$I$~$<$~12. By the time it reaches completion in $\\sim2011$ it is hoped that RAVE will have observed up to one million stars, providing a dataset that will be of great importance for Galaxy structure studies. A number of publications have already made use of this dataset \\citep[e.g.][]{Smith2007, Klement+2008,Munari2008,Siebert+2008,Veltz2008}, but to fully utilise the kinematic information we crucially need to know the distances to the stars. Unfortunately, most of the stars in the RAVE catalogue are too faint to have accurate trigonometric parallaxes, hence the importance of a reliable and well-tested photometric/spectroscopic parallax algorithm. When distances are combined with archival proper motions and high precision radial velocities from RAVE, this dataset will provide the full 6D phase-space coordinates for each star. Clearly such an algorithm for estimating distances will be a vital tool when carrying out kinematic analyses of large samples of Galactic stars, not just for the RAVE survey but for any similar study. The future prospects for distance determinations are very promising. In the next decade the Gaia satellite \\citep{Perryman2001} will observe up to $10^{9}$ stars with exquisite astrometric precision. The mission is due to start in 2012, but a final data release will not arrive until near the end of the decade at the earliest. Furthermore, as with any such magnitude limited survey, there will be a significant proportion of stars for which their distances are too great for accurate trigonometric parallaxes to be determined. Therefore, although Gaia will revolutionise this field, it will not close the chapter on distance determinations for stars in the Milky Way and so photometric parallax techniques will remain of crucial importance. In this paper we present our algorithm for determining distances, which we construct using stellar models. When we apply this method to the RAVE dataset we are able to reproduce several known characteristics of the kinematics of stars in the solar neighbourhood. In \\S \\ref{sec:method}, we present a general introduction. We discuss the connection between stellar evolution theory, stellar tracks and isochrones to gain insight in these topics before presenting our statistical methods for the distance determination and testing the method using synthetic data. In \\S \\ref{sec:rave} we apply the method to the RAVE dataset and compare the distances to external determinations, namely stars in the open cluster M67 and nearby stars with trigonometric parallaxes from Hipparcos. Results obtained from the phase-space distribution are presented in \\S \\ref{sec:results} to check whether the data reflect known properties of our Galaxy. We present a discussion of the uncertainties and limitations of the method in \\S \\ref{sec:discussion} and conclude with \\S \\ref{sec:conc}. ", "conclusions": "\\label{sec:conc} We have presented a method to derive absolute magnitudes, and therefore distances, for RAVE stars using stellar models. It is based on the use of stellar model fitting in metallicity, $\\logg$, $\\T$ and colour space. We find that our method reliably estimates distances for main-sequence stars, but there is an indication of potential systematic problems with giant stars owing to issues with the underlying stellar models. The uncertainties in the estimated absolute magnitudes for RGB stars are found to depend mainly on the uncertainties in $\\logg$, while for main-sequence stars the accuracy of $\\T$ is also important (\\S \\ref{sec:method:test}). For the RAVE data the uncertainties in $\\logg$ and $\\T$ give rise to relative distance uncertainties in the range 30\\%-50\\%, although from cross-matching with Hipparcos (\\S \\ref{sec:app:testing:hipp}) it appears that our uncertainties may be overestimated for the brighter stars (with higher signal-to-noise spectra). It is important to note that that some 10\\% of the RAVE stars may be on the red clump, but these are treated as RGB by our pipeline, and hence their distances may be systematically biased. As can be seen in the results section (\\S \\ref{sec:results}), the data accurately reflect the known properties of halo and disk stars of the Milky Way. A variation in metallicity and $v_\\phi$ was found away from the Galactic plane, corresponding to an increase in the fraction of metal-poor halo stars. Existing substructure in the $UV$ velocity plane was recovered, as was the vertex deviation. Upon completion the RAVE survey will have observed a factor of up to $\\sim$20 times more stars than analysed here. Clearly this will be a hugely valuable resource for studies of the Galaxy. In future the Gaia satellite mission \\citep{Perryman2001} will revolutionise this field, recording distances to millions of stars with unprecedented accuracy. However, for large numbers of Gaia stars it will not be possible to accurately constrain the distance due to them being too far away or too faint, which implies that it is crucial to develop techniques such as ours for reliably estimating distances. In the near term it will be possible to improve the accuracy of our pipeline by calibrating it through observations of clusters; a technique which has been used with great success by the Sloan Digital Sky Survey \\citep{Ivezic2008}. Within the RAVE collaboration a project is underway to obtain data for cluster stars \\citep[e.g.][]{Kiss2007} and we aim to incorporate this into future analyses. This may allow us to reduce or remove the reliance on stellar models, which will lessen one of the major sources of uncertainty in our work. Our pipeline will allow us to fully utilise current surveys such as RAVE, and also places us in an ideal position exploit future large-scale spectroscopic surveys that will be enabled by upcoming instruments such as LAMOST." }, "1003/1003.0272_arXiv.txt": { "abstract": "We use the statistics of regions above or below a temperature threshold (excursion sets) to study the cosmic microwave background (CMB) anisotropy in models with primordial non-Gaussianity of the local type. By computing the full-sky spatial distribution and clustering of pixels above/below threshold from a large set of simulated maps with different levels of non-Gaussianity, we find that a positive value of the dimensionless non-linearity parameter $f_{\\rm NL}$ enhances the number density of the cold CMB excursion sets along with their clustering strength, and reduces that of the hot ones. We quantify the robustness of this effect, which may be important to discriminate between the simpler Gaussian hypothesis and non-Gaussian scenarios, arising either from non-standard inflation or alternative early-universe models. The clustering of hot and cold pixels exhibits distinct non-Gaussian signatures, particularly at angular scales of about 75 arcmin (i.e. around the Doppler peak), which increase linearly with $f_{\\rm NL}$. Moreover, the clustering changes strongly as a function of the smoothing angle. We propose several statistical tests to maximize the detection of a local primordial non-Gaussian signal, and provide some theoretical insights within this framework, including an optimal selection of the threshold level. We also describe a procedure which aims at minimizing the cosmic variance effect, the main limit within this statistical framework. ", "introduction": "Since some level of non-Gaussianity is generically expected in all inflation models, due to interactions of the inflaton with gravity and/or from inflaton self-interactions, seeking for deviations from the Gaussian paradigm has recently become a major effort -- and a minor industry -- in cosmology. Properties of the primordial perturbations are uniquely imprinted in the cosmic microwave background (CMB) anisotropy distribution; hence, its analysis is a powerful way of looking at the specifics of the inflationary models (or alternatives to inflation). At the present time, the main challenge is either to detect or to constrain mild or weak departures from primordial Gaussian initial conditions, as the level of non-Gaussianity predicted in the simplest single-field slow-roll inflation is slightly below the minimum value detectable by the Planck satellite, and not within reach of future galaxy surveys. This is essentially why primordial non-Gaussianity is regarded as one of the most promising probes of the inflationary universe (Komatsu et al. 2009b), and it has received a recent boost, both theoretically and observationally, mainly because of the \\textit{W}ilkinson \\textit{M}icrowave \\textit{A}nisotropy \\textit{P}robe (WMAP) data which seems to favor a slightly positive value of the dimensionless non-linearity parameter $f_{\\rm NL}$ (Yadav \\& Wandelt 2008; Komatsu et al. 2009a, 2010; Smith et al. 2009). From the theoretical side, much effort has been directed towards the development of competing scenarios for perturbation generation which go beyond the single-field slow-roll paradigm, for instance by the inclusion in the Lagrangian of non-trivial kinetic terms, the presence of more than one light field during inflation, the temporary violation of slow-roll, or a non-adiabatic initial vacuum state for the inflaton. Examples are the curvaton model, the modulated reheating, DBI or ghost inflation, or multi-field scenarios, some of which imply large departures from Gaussianity (see, for instance, among the plethora of papers on this subject, Linde \\& Mukhanov 1997; Lyth \\& Wands 2002; Acquaviva et al. 2003; Lyth, Ungarelli \\& Wands 2003; Maldacena 2003; Alishahiha et al. 2004; Arkani-Hamed et al. 2004; Bartolo et al. 2004; Dvali, Gruzinov \\& Zaldarriaga 2004; Chen 2005; Seery \\& Lidsey 2005; Bartolo, Matarrese \\& Riotto 2006; Lyth \\& Riotto 2006; Sasaki et al. 2006; Creminelli et al. 2007; Creminelli \\& Senatore 2007; Koyama et al. 2007; Buchbinder et al. 2008; Chen et al. 2008, 2009; Lehners \\& Steinhardt 2008; Matarrese \\& Verde 2008; Sasaki 2008; Bartolo \\& Riotto 2009; Brandenberger 2009; Naruko \\& Sasaki 2009; Senatore, Tassev \\& Zaldarriaga 2009; Silvestri \\& Trodden 2009; Bartolo, Matarrese \\& Riotto 2010). From the observational point of view, the main goal is to constrain the level of primordial non-Gaussianity directly from a real data set, and this is usually achieved by constructing and applying a variety of non-Gaussian estimators such as the 3-point function (Hinshaw et al. 1994; Gangui et al 1994), the genus statistics or the topological genus density (Coles 1988; Gott et al. 1990; Smoot et al. 1994; Colley \\& Gott 2003; Park 2004; Gott et al. 2007), the other Minkowski functionals (Schmalzing \\& Gorski 1998; Winitzki \\& Kosowsky 1998; Banday, Zaroubi \\& Gorski 2000; Hikage et al. 2006, 2008b; Matsubara 2010), the bispectrum and trispectrum (Spergel et al. 2007; Komatsu et al. 2009; Rudjord et al. 2009; Liguori et al. 2010), tensor modes (Coulson, Crittenden \\& Turok 1994), wavelets (Cabella et al. 2005; Curto et al. 2009; Vielva \\& Sanz 2009), pixel and peak statistics (Adler 1981; Bond \\& Efstathiou 1987; Coles and Barrow 1987; Kogut et al. 1995, 1996; Barreiro et al. 1997, 1998; Heavens 1998; Heavens \\& Sheth 1999; Heavens \\& Gupta 2001; Hern{\\'a}ndez-Monteagudo et al. 2004; Rossi et al. 2009; Hou et al. 2010), phase correlations, multifractals, and so forth (see also Komatsu, Spergel \\& Wandelt 2005; Chen \\& Szapudi 2006; Munshi \\& Heavens 2010). In this process, many observational challenges and experimental artifacts come into play; therefore, it is perhaps not surprising that controversial results and a long list of anomalies have been reported so far, ranging from a low value of the quadrupole till North-South or parity asymmetries, strange alignments in the data, and much more (see for example Chiang et al. 2003, 2007; Tegmark et al. 2003; de Oliveira-Costa et al. 2004; Eriksen et al. 2004, 2007; Schwarz et al. 2004; Cruz et al. 2005, 2006, 2007, 2008; Land \\& Magueijo 2005, 2007; Naselsky et al. 2005; Copi et al. 2006, 2007; Vielva et al. 2007; Gurzadyan et al. 2008; Pietrobon et al. 2009; R{\\\"a}th et al 2009; Kim \\& Naselsky 2010). Deviations from Gaussian initial conditions (if any) also carry important consequences on many aspects of the large-scale structure (LSS) of the Universe, and galaxy surveys can provide constraints on non-Gaussianity competitive with those from the CMB alone. There are in fact modifications in the statistics of voids (Kamionkowski, Verde \\& Jimenez 2009), in the distribution of neutral hydrogen and in the intergalactic medium (Viel et. al 2009), in the high-mass tail of the halo distribution (Chiu et al. 1998; Matarrese, Verde \\& Jimenez 2000; Sefusatti \\& Komatsu 2007; LoVerde et al. 2008), in the large-scale skewness of the galaxy distribution (Chodorowski \\& Bouchet 1996), in the number counts of clusters and of density peaks (Desjacques et al. 2009; Jeong \\& Komatsu 2009), in the measurement of the scale dependence of the bias of LSS tracers (Carbone et al. 2008; Dalal et al. 2008; Verde \\& Matarrese 2009; Desjacques \\& Seljak 2010), in the reionization history (Crociani et al. 2009), in the galaxy power spectrum and bispectrum (Scoccimarro 2000; Scoccimarro et al. 2004; Mangilli \\& Verde 2009), in the topology (Park et al. 1998, 2005; Gott et al. 2008; Hikage et al. 2008a), and in the abundance and clustering of galaxies and dark matter halos (Verde et al. 2001; Afshordi \\& Tolley 2008; Grossi et al. 2008; LoVerde et al. 2008; Matarrese \\& Verde 2008; McDonald 2008; Slosar et al. 2008; Taruya et al. 2008; Pillepich et al. 2010). Despite all these remarkable theoretical and observational efforts, till date the experimental detection of a significant deviation from the Gaussian paradigm remains still challenging and not convincing. In this respect, we need to explore alternative statistics more sensitive to deviations from Gaussianity, and to search for unique features which may allow one to distinguish among the myriad of inflation models available in the literature. It is important to adopt different and complementary statistical approaches, and not just a single view, because non-Gaussianity can take innumerable forms. In fact, while Gaussian random processes are theoretically desirable since they are the only ones for which the knowledge of all spectral parameters completely determines all the statistical properties, as soon as we introduce departures from Gaussianity a more complicated scenario emerges, and there is no such statistics which describes fully and uniquely the non-Gaussian nature of a sample. In particular, moving away from standard estimators like the bispectrum, trispectrum, three and four-point functions, skewness, etc, we are interested here in rare events, which can often maximize deviations from what is predicted by a Gaussian distribution. The main goal of the present work is to extend and apply the statistics of the excursion sets, regions above or below a temperature threshold, to models with primordial non-Gaussianity. Specifically, we focus on the \\textit{local} parametrization of non-Gaussianity (Salopek \\& Bond 1990), by including quadratic corrections to the curvature perturbation. We simulate a large set of full-sky maps with different $f_{\\rm NL}$ values, and compute the number density and the spatial clustering of the CMB excursion set regions. We also provide the theoretical formalism to interpret our results. The excursion set statistics is fully characterized in the context of Gaussian random fields (Kaiser 1984; Bardeen et al. 1986), and it has been used in a variety of studies (see for example Jensen \\& Szalay 1986; Bond \\& Efstathiou 1987; Barreiro et al. 2001; Kashlinsky et al. 2001 and references therein). There are also some extensions to non-Gaussian conditions in the literature (i.e. Coles \\& Barrow 1987; Coles 1988; Barreiro et al. 1998). Our analysis differs from those of the previous authors primarily because we use a more realistic model for non-Gaussianity supported by $f_{\\rm NL}$ type simulations, and because we also propose some new statistical tools, tests, and theoretical insights within this framework. In particular, while in precedent studies it has always been shown that the Gaussian correlation function of the excursion sets and peaks (a subset of the excursion sets) is easily distinguishable from a non-Gaussian one, even if the underlying bispectra are not statistically different (i.e. Kogut et al. 1995; Barreiro et al. 1998; Heavens \\& Gupta 2001), we suggest here that it may not be the case if the model of non-Gaussianity is of the local type, and the resolution adopted is not optimal. Our work is also motivated by another reason. In a previous analysis (Rossi et al. 2009), we compared the pixel clustering statistics -- properly extended to handle inhomogeneous noise -- against WMAP five-year data, and we detected deviations from the Gaussian theoretical expectations. In particular, we found a remarkable difference in the clustering of hot and cold pixels at relatively small angular scales. A similar trend has also been reported in the literature by Tojeiro et al. (2006), and by Hou, Banday \\& Gorski (2010), although at much larger scales. Whether or not this discrepancy may arise from primordial non-Gaussianity of the local type is another key question of this analysis. The layout of the paper is as follows. Section \\ref{NG theory} contains the theoretical tools developed and used in this study. In Section \\ref{fnl_model} we briefly describe the local $f_{\\rm NL}$ model. In Section \\ref{map_simulations} we explain how the simulated non-Gaussian maps are constructed. In Section \\ref{excursion_set_formalism} we provide the basic formalism for the excursion sets statistics, in the context of $f_{\\rm NL}$ scenarios. Expressions for the one- and two-dimensional probability distribution functions (PDFs) are given, under the assumption of weak non-Gaussianity; this is done via a perturbative approach by the multidimensional Edgeworth expansion around a Gaussian distribution function. Those PDFs are then used to characterize the number density and the clustering statistics above/below threshold as a function of $f_{\\rm NL}$ (some details are provided in Appendix \\ref{edgeworth_proxy}). In Section \\ref{history} we relate the excursion sets formalism to other commonly used topological estimators. In Section \\ref{NG maps analysis}, computations of the number density and the clustering statistics above/below threshold from non-Gaussian maps are presented and interpreted according to our theory predictions. Specifically, Section \\ref{nd_subsection} shows the abundance of the excursion set regions in a variety of ways, while in Section \\ref{nd_stat_subsection} we highlight some statistical tests developed using the number density. We also argue that there are optimal thresholds which can maximize the non-Gaussian contribution, as well as levels which do not allow to distinguish a Gaussian signal from a non-Gaussian one. In Section \\ref{clustering_subsection} we present the clustering of hot and cold pixels for one of the optimal temperature thresholds as a function of the smoothing scale, and in Section \\ref{clustering_stat_subsection} we propose a new statistical test derived from the clustering statistics. This procedure aims at minimizing the cosmic variance effect, and involves the computation of the power spectrum for any given CMB map. A final part (Section \\ref{NG_conclusions}) summarizes our findings, and highlights ongoing and future work. We leave in Appendix \\ref{noise_analytic}, \\ref{errors_analytic} and \\ref{spurious_ng} some technical details regarding experimental artifacts such as inhomogeneous noise, incomplete sky coverage, errorbar estimates and confusion effects caused by spurious non-Gaussianities; all these experimental complications will be examined in more detail in the forthcoming publications. ", "conclusions": "\\label{NG_conclusions} We have extended and applied the statistics of the excursion sets to models with primordial non-Gaussianity of the $f_{\\rm NL}$ local type. While in presence of Gaussian initial conditions many statistics based on geometrical and topological properties of the CMB temperature have been developed and well-studied, to date fewer analyses have been focused on geometrical properties of the CMB radiation in the presence of primordial non-Gaussianity. In particular, our work is the first extension of the excursion set formalism to local $f_{\\rm NL}$ type non-Gaussianity. From a large set of simulated full-sky non-Gaussian maps, we computed the number density and the spatial clustering of CMB patches above/below a temperature threshold (Section \\ref{NG maps analysis}). We found that a positive value of $f_{\\rm NL}$ enhances the number density of the cold CMB excursion sets (Figures \\ref{nd_smart_scale_fig}, \\ref{nd_difference_ratio_fig}) along with their clustering strength (Figures \\ref{clustering_fig}, \\ref{clustering_difference_ratio_fig}) and reduces that of the hot ones. We performed a thorough statistical analysis to evaluate the sensitivity of the two observables to the level of non-Gaussianity and to the smoothing resolution. We also provided the analytical formalism to interpret our results (Section \\ref{NG theory}). Expressions for the one- and two-dimensional PDFs (Equations \\ref{one_d_pdf_fnl} and \\ref{2d_edge_eq}) were obtained from a perturbative approach by the multidimensional Edgeworth expansion around a Gaussian distribution function, and used to characterize the abundance and clustering statistics as a function of $f_{\\rm NL}$ (Equations \\ref{nd_fnl_eq}, \\ref{corr_smart}, \\ref{corr_smart_gauss}). We showed that there are optimal thresholds which maximize the local $f_{\\rm NL}$ non-Gaussianity ($\\nu = 0.25, 0.50$ and $\\nu=2.00, 2.25$), as well as others ($\\nu=1.00$) which do not allow for a distinction between the Gaussian and the non-Gaussian signals (Figures \\ref{nd_sigma_distances_fig} and \\ref{clustering_sigma_units_fig}). We devised a new statistical test based of the number density (Section \\ref{nd_stat_subsection}), which combines two thresholds where departures from Gaussianity are most significant (Figure \\ref{nd_combined_fig} and Equation \\ref{nd_composite_eq}). We also proposed a new procedure aimed at minimizing the effect of cosmic variance (Section \\ref{clustering_stat_subsection}), which involves the clustering information alone (Figure \\ref{clustering_smart_smarter_test_fig}, Equations \\ref{cf_test_smart_hp_eq} and \\ref{cf_test_smart_cp_eq}). Although we focused here on $f_{\\rm NL}$ models of the local type, the statistical tools developed are more general and can be applied to describe any other type of non-Gaussianity. A typical example is represented by the curvaton model, for which the cubic term indicated as $g_{\\rm NL}$ can be large, while $f_{\\rm NL}$ can be negligible. Our technique can be applied to this case as well, and it is the subject of a forthcoming publication. This work was primarily motivated by our previous finding (Rossi et al. 2009), namely a remarkable difference in the clustering of hot and cold pixels at relatively small angular scales from the WMAP 5-yr data. We analyzed the possibility that this discrepancy may arise from primordial non-Gaussianity of the local $f_{\\rm NL}$ type (Section \\ref{clustering_subsection}), and concluded that only a large value of $f_{\\rm NL}$ would provide such a difference (Figure \\ref{clustering_fig}). Cosmic variance plays a crucial role within this statistical framework, so that the Gaussian correlation function of the excursion sets is not easily distinguishable from the non-Gaussian one, contrary to what was previously thought. In fact, while a distinct signature in the clustering of hot and cold pixels clearly emerges for a large $f_{\\rm NL}$ non-Gaussianity, particularly at angular scales of about 75 arcmin (around the Doppler peak), as expected this feature is reduced when $f_{\\rm NL}=100$. The clustering behavior is also strongly affected by the smoothing angle. These findings suggest that Gaussianity itself cannot be accurately constrained from the excursion set clustering statistics. In fact, if in principle the use of pixel-pixel correlation functions as a test of Gaussianity is very powerful, because there are no free parameters once the underlying power spectrum has been measured, this may not be the case if the non-Gaussian model is of the $f_{\\rm NL}$ local type, and $f_{\\rm NL}$ is small. Our study was focused on a few selected values of thresholds and two different statistics, so that the predicted constraints on $f_{\\rm NL}$ are wider than what one would get by combining several threshold levels and different smoothing angles. In this respect, our predicted constraints from the excursion sets are compatible with those of Smidt et al. (2010), obtained from the trispectrum. Since cosmic variance is the main obstacle in the analysis, we are considering derived statistics which could potentially beat its effect and maximize the non-Gaussian contribution. It is also important to adopt different and complementary statistical approaches, and not just a single view, because there is no such statistics which describes fully and uniquely the non-Gaussian nature of a sample. To this end, a lot of effort has recently gone into developing optimal estimators, and in this sense our statistical technique belongs to a class of topological estimators which may be considered ``sub-optimal'' for measuring non-Gaussianity. However, in reality all the geometrical methods complement and ``diagnose'' results obtained with bispectrum or trispectrum estimators. Moreover, geometrical techniques are often model-independent, easy to implement, with low computational cost, and they can retain information on the spatial distribution of the non-Gaussian signal. Also, they provide useful analytic insights and physical intuition. For example, the derivation and implementation of the analytical formula for the CMB Minkowski functionals in the limit of weak non-Gaussianity (Hikage, Komatsu \\& Matsubara 2006; Matsubara 2010) has allowed to obtain limits on various models, for which the optimal estimators are difficult to implement; at the moment, a limit on the primordial non-Gaussianity in the isocurvature perturbation is available \\textit{only} from the Minkowski functionals (Hikage et al. 2009). Note also that the concept of ``optimal'' is often misleading, as it requires a posteriori knowledge of the type of non-Gaussianity which is, at least in principle, unknown. The main question, instead, is whether or not it is possible to improve limits on $f_{\\rm NL}$ using the CMB data only. Including realistic effects in our simulations, such as inhomogeneous noise, point source contamination or foregrounds, so that we can compare our predictions with current observations, is subject of ongoing work (we provide some discussion in Appendices \\ref{noise_analytic} and \\ref{spurious_ng}). We present results of these investigations in a companion paper, where we are also consider more terms in the expansion (\\ref{fnl_expansion_eq}). Application of the formalism presented in Section \\ref{excursion_set_formalism} to peak rather than pixel statistics is a straightforward exercise, and is also the subject of another forthcoming publication. The Planck satellite with its increased sensitivity and resolution is expected to improve the measurements of most cosmological parameters by several factors compared to WMAP, and in synergies with future galaxy surveys (Colombo, Pierpaoli \\& Pritchard 2009). In fact, Planck gains a factor of 2.5 in angular resolution and up to 10 in instantaneous sensitivity with respect to WMAP, and it is nearly photon noise limited in the CMB channels (100-200 GHz). Repeating this analysis at the Planck resolution may then provide more stringent limits on $f_{\\rm NL}$ from the excursion set statistics, and is also the subject of work in progress." }, "1003/1003.3101_arXiv.txt": { "abstract": "{} {We derive the value of the dark matter density at the Sun's location ($\\rho_\\odot$) without globally mass-modeling the Galaxy.} {The proposed method relies on the local equation of centrifugal equilibrium and is independent of i) the shape of the dark matter density profile, ii) knowledge of the rotation curve from the galaxy center out to the virial radius, and iii) the uncertainties and the non-uniqueness of the bulge/disk/dark halo mass decomposition.} {The result can be obtained in analytic form and it explicitly includes the dependence on the relevant observational quantities and takes their uncertainties into account. By adopting the reference, state-of-the-art values for these, we find $\\rho_\\odot=0.43(11)(10)\\,$GeV/cm$^{3}$, where the quoted uncertainties are respectively due to the uncertainty in the slope of the circular-velocity at the Sun location and the ratio between this radius and the length scale of the stellar exponential thin disk.} {We obtained a reliable estimate of $\\rho_\\odot$, that, in addition has the merit of being ready to take into account any future change/improvement in the measures of the observational quantities it depends on.} ", "introduction": "Galaxy rotation curves (e.g.~\\citealt{rubin80,bosma81}) have unveiled the presence of a dark ``mass component'' in spirals. They are pillars of the paradigm of massive dark halos, composed of a still undetected kind of matter surrounding the luminous part of galaxies. The kinematics of spirals shows universal systematics~\\citep{PSS,URC2}, which seems to be at variance with the predictions emerging from simulations performed in the $\\Lambda$ cold dark matter ($\\Lambda$CDM) scenario, (e.g.~\\citealt{NFW96}), the currently preferred cosmological paradigm of galaxy formation (e.g.~\\citealt{gentile04}). Individual and coadded rotation curves (RCs) of spiral galaxies are also crucial to investigate frameworks alternative to the standard paradigm of cold collisionless DM in Newtonian gravity (e.g.~\\citealt{mond,bigr}). At the same time, dedicated searches of DM particle candidates have seen an important boost in recent years with relevant and costly experiments being planned and executed. The so-called direct-detection experiments look for the scattering of DM particles off the nuclei inside the detectors (e.g.\\ CDMS, XENON10, DAMA/LIBRA) by typically measuring the deposited energy or its annual modulation. Clearly in all these experiments the signal is proportional to the DM density in the Sun's region, $\\rho_\\odot$. On the other hand indirect-detection experiments (in particular Super-Kamionkande, AMANDA, IceCube and ANTARES) search for the secondary particles (neutrinos in these cases) produced by DM annihilations at the center of the Sun or Earth, where it is expected that DM accumulates after losing energy via scattering, possibly reaching a thermalized state. The expected signal in this case depends on the DM density inside these objects, which in turn is driven, via the capture mechanism, by the same halo DM density in the Sun region,~$\\rho_\\odot$. Therefore, in both these kinds of direct and indirect searches, an estimate of the the local density $\\rho_\\odot$ is very important for a precise estimate of the signal or at least reliable bounds on the DM cross-section vs mass to be compared with limits from other searches. What is then the value of $\\rho_{\\odot}$? A value of \\begin{equation} \\rho_\\odot =0.3\\, {\\rm GeV/cm}^{3} \\label{eq:one} \\end{equation} is routinely quoted in hundreds of papers, but how does this number come out? In which works do we find the details of its measure? It is worth observing that in most of the cases in the literature, the above value is given with no reference (e.g.~\\citealt{donato09,savage09}). Sometimes, the reference goes to a couple of seminal papers. Among them, the Particle Data Group Review \\citep{pdg08} indicates the above value ``within a factor of two or so'' and justifies it as coming from ``recent estimates based on a detailed model of our Galaxy''. However, the works cited are neither recent nor detailed and sometimes not even an independent estimation of $\\rho_{\\odot}$. The only exception is the work by \\cite{caldwell81} that devised what can be considered as the standard method (CO hereafter) to determine the value of $\\rho_{\\odot}$ from observations (see below). Their resulting value, $0.23^{+0.23}_{-0.12}$GeV/cm$^{3}$, arises however from very outdated kinematical observations and from a {\\it cored} (rather than a cuspy) halo distribution, so it is not a great support for equation~(\\ref{eq:one}). Similar conclusions can be drawn by looking at other influential reviews: the papers they cite to back up the value~(\\ref{eq:one}) either do not estimate this quantity or use very outdated observations. \\medskip In general, it is quite simple to infer the distribution of dark matter in spiral galaxies. Spiral's kinematics, in fact, reliably traces the underlying gravitational potential \\citep{PSS,URC2}. Then, from coadded and/or individual RCs, we can build suitable mass models that include stellar and gaseous disks along with a spherical bulge and a dark halo. More in detail, by carefully analyzing (high quality) circular velocity curves, with the help of relevant photometric and HI data, one can derive the halo density at any desired radius. The accuracy of the ``measurements'' is excellent and the results are at the core of the present debates on Galaxy formation (e.g.~\\citealp{gentile04,gentile05,deblok09}). \\medskip To measure $\\rho_\\odot$, instead, is far from simple, because the MW kinematics, unlike that of external galaxies, does not trace the gravitational potential straightforwardly. We do not directly measure the circular velocity of stars and gas but rather, at our best, the terminal velocity $V_T$ of the rotating HI disk, and this only inside the solar circle (e.g.~\\citealt{McClure07}). This velocity is related to the circular velocity $V(r)$, for $r3 R_D$ we have $F_{tot}(R_\\odot)\\gg F_D(R_\\odot)$, equation~(\\ref{eq:rhoDM1}) does not collapse and, as a bonus, the most uncertain term of the r.h.s.\\ of (\\ref{eq:rhoDM1}) is also the smaller one; 2) $\\omega_\\odot$ is very precisely measured; 3) $d \\alpha/dr|_{R_\\odot}\\simeq 0$; and 4) at the Sun's position the bulge density $\\rho_B(R_\\odot)$ is totally negligible, $< \\rho_H/50$ (e.g.~\\citealt{sofue2}). Thus, this method is very powerful for determining the value of the DM density at $R_{\\odot}$. The DM density at {\\it any} radii is obviously left to the standard mass modeling. The method is obviously simpler for a spherically symmetric DM halo, and can be further simplified by considering an infinitesimally thin disk for the distribution of stars in the Galaxy. However, below we also include the effects of a possible halo oblateness and disk thickness. Here, we anticipate that these effects, constrained by observations, are rather weak, of the order of a few percent, and are therefore irrelevant for this work. As a result, we obtain a reliable and model-independent determination of the local DM halo density, as well as of its intrinsic uncertainty. \\medskip In the next section we describe the method in detail and derive $\\rho_\\odot$ and its uncertainty as a function of the relevant observables. In the last section we discuss the results and draw the conclusions. In the Appendices, we explicitly describe the effect of the halo oblateness and disk thickness, and we comment on the inherent problems in the traditional determination of $\\rho_\\odot$. ", "conclusions": "In this work we have provided a model-independent kinematical determination of $\\rho_\\odot$. The method proposed here derives $\\rho_\\odot$ directly from the solution of the equation of centrifugal equilibrium, by estimating the difference between the `total' density and that of the stellar component. The method leads to an optimal kinematical determination of $\\rho_\\odot$, avoiding model-dependent and dubious tasks mandatory with the standard method, i.e., a) to assume a particular DM density profile and a specific dynamical status for the tracers of the gravitational potential, b) to deal with the non-negligible uncertainties of the global MW kinematics, c) to uniquely disentangle the flattish RC into the different bulge/disk/halo components. While the measure of $\\rho_\\odot$ can be performed in an ingenious way, it cannot escape the fact that it ultimately depends at least on three local quantities, the slope of the circular velocity at the Sun, the fraction of its amplitude due to the DM, and the ratio between the Sun galactocentric distance and the disk scale-length, whose uncertainty unavoidably propagates in the result. Two of these three quantities can be related by noting that the MW is a typical Spiral and using the relations available for these kind of galaxies~\\citep{URC2}, so that the final uncertainty can be slightly reduced. We found that some oblateness of the DM halo and the small finite thickness of the stellar disk play a limited role in the measure. However, we took them into account by the simple correction terms described. The resulting local DM density that we find, $\\rho_{\\odot}=(0.43 \\pm 0.11_{({\\alpha_\\odot})} \\pm0.10_{(r_{\\odot D})})\\, {\\rm GeV/cm^{3}}$, is still consistent with previous determinations, or slightly higher. However, the determination is free from theoretical assumptions and can be easily updated by means of equation~(\\ref{eq:10bis}) as the relevant quantities will become better known.\\footnote{Again, in the traditional method most of the uncertainty in the measure of $\\rho_\\odot$ discussed in the Introduction cannot be overcome by having more and better data.} A final comment is in order. The {\\it values} of $\\rho_\\odot$ found in previous studies by means of the traditional methods (e.g.\\ \\citealt{sofue2, weber09}) differ among themselves and also from the present value only by a small factor. This relatively good agreement in the values does not imply a concordance in the underlying mass models, in the various assumptions taken or in the data set employed, but is mainly due to the fact that $\\omega$ or equivalently $A-B$ (the well known combination of Oort constants) is measured with good precision. In fact, from $M_H(r)\\propto V^2_H(r)\\, r$ we have \\begin{equation} \\label{eq:stima} \\rho_\\odot = k^2 \\omega^2 (1+ 2\\alpha_H) \\end{equation} where $\\omega\\simeq 30\\,$km/s/kpc, $k^2$ is the fraction of the halo contribution to the circular velocity at $R_\\odot$ and $\\alpha_H\\equiv d{\\rm log} V_H/d{\\rm log} R$ is its unknown slope at the same radius. The quantities $k^2$ and $\\alpha_H$ are experimentally unknown. One can use different assumptions, mass modeling and data to get them, but he/she will always find that they range from 0.3 (max disk, \\citealt{PS90}) to 0.7, and from 0.2 to 0.8 (max disk). These are very large variations in terms of structural properties of spirals, but only mild ones in the determination of $\\rho_\\odot$, which is dominated by the term $\\omega^2$: for any reasonable value of $\\alpha_H$ and $k$, the density will be in the range $0.25\\,$GeV$<\\rho_\\odot<0.70\\,$GeV. An increase in precision beyond this scale estimate would require an accurate determination of $k$ and of the relative mass modeling, of difficult realization by means of traditional methods as discussed in the introduction. The present method will be able to determine very precisely $\\rho_\\odot$ through equation (\\ref{eq:11bis}) if improved/new measures of the relevant observational quantities also emerge. We thus believe that the value given in equation (\\ref{eq:result3}) reflects the present state-of-the-art knowledge of $\\rho_\\odot$ and of its uncertainty, and may result in being very useful in deriving reliable future bounds on the DM cross sections involved in direct and indirect DM searches. {\\small \\sect" }, "1003/1003.4384_arXiv.txt": { "abstract": "The final stage of terrestrial planet formation is known as the giant impact stage where protoplanets collide with one another to form planets. So far this stage has been mainly investigated by $N$-body simulations with an assumption of perfect accretion in which all collisions lead to accretion. However, this assumption breaks for collisions with high velocity and/or a large impact parameter. We derive an accretion condition for protoplanet collisions in terms of impact velocity and angle and masses of colliding bodies, from the results of numerical collision experiments. For the first time, we adopt this realistic accretion condition in $N$-body simulations of terrestrial planet formation from protoplanets. We compare the results with those with perfect accretion and show how the accretion condition affects terrestrial planet formation. We find that in the realistic accretion model, about half of collisions do not lead to accretion. However, the final number, mass, orbital elements, and even growth timescale of planets are barely affected by the accretion condition. For the standard protoplanetary disk model, typically two Earth-sized planets form in the terrestrial planet region over about $10^8$ years in both realistic and perfect accretion models. We also find that for the realistic accretion model, the spin angular velocity is about 30\\% smaller than that for the perfect accretion model that is as large as the critical spin angular velocity for rotational instability. The spin angular velocity and obliquity obey Gaussian and isotropic distributions, respectively, independently of the accretion condition. ", "introduction": "It is generally accepted that the final stage of terrestrial planet formation is the giant impact stage where protoplanets or planetary embryos formed by oligarchic growth collide with one another to form planets \\citep[e.g.,][]{w85,ki98}. This stage has been mainly studied by $N$-body simulations. So far all $N$-body simulations have assumed perfect accretion in which all collisions lead to accretion \\citep[e.g.,][]{acl99,c01}. However, this assumption would be inappropriate for grazing impacts that may result in escape of an impactor or hit-and-run. By performing Smoothed-Particle Hydrodynamic (SPH) collision simulations, \\cite{aa04} estimated that more than half of all collisions between like-sized protoplanets do not simply result in accumulation of a larger protoplanet, and this inefficiency lengthens the timescale of planet formation by a factor of 2 or more, relative to the perfect accretion case. The accretion inefficiency can also change planetary spin. \\cite{ki07} found that under the assumption of perfect accretion, the typical spin angular velocity of planets is as large as the critical spin angular velocity for rotational instability. However, in reality, the grazing collisions that have high angular momentum are likely to result in a hit-and-run, while nearly head-on collisions that have small angular momentum lead to accretion. In other words, small angular momentum collisions are selective in accretion. Thus, the accretion inefficiency may lead to slower planetary spin, compared with the perfect accretion case. The goal of this paper is to clarify the statistical properties of terrestrial planets formed by giant impacts among protoplanets under a realistic accretion condition. We derive an accretion condition for protoplanet collisions in terms of collision parameters, masses of colliding protoplanets and impact velocity and angle, by performing collision experiments with an SPH method. We implement the realistic accretion condition in $N$-body simulations and probe its effect to further generalize the model of terrestrial planet formation. We derive the statistical dynamical properties of terrestrial planets from results of a number of $N$-body simulations and compare the results with those in \\cite{kki06} and \\cite{ki07} where perfect accretion is adopted. In section 2, we outline the initial conditions of protoplanets and the realistic accretion condition. Section 3 presents our results, where we show the statistics of collision parameters and basic dynamical properties of planets. Section 4 is devoted to a summary and discussions. ", "conclusions": "\\label{sec:summary} We have investigated the basic dynamical properties of the terrestrial planets assembled by giant impacts of protoplanets by using $N$-body simulations. For the first time, we adopted the realistic accretion condition of protoplanets obtained by the SPH collision experiments. The basic dynamical properties have been studied statistically with numbers of $N$-body simulations. For the standard protoplanet system, the statistical properties of the planets obtained are the following: \\begin{itemize} \\item About half of collisions in the realistic accretion model do not lead to accretion. However, this accretion inefficiency barely lengthens the growth timescale of planets. \\item The numbers of planets are $\\langle n\\rangle \\simeq 3{\\rm -}4$ and $\\langle n_M\\rangle \\simeq \\langle n_a\\rangle \\simeq2$. The growth timescale is about $6{\\rm -}7\\times 10^7$ years. The masses of the largest and second-largest planets are $\\langle M_1\\rangle\\simeq 1.2M_\\oplus$ and $\\langle M_2\\rangle\\simeq 0.7M_\\oplus$. The largest planets tend to form around $\\langle a_1\\rangle\\simeq 0.8$AU, while $a_2$ is widely scattered in the initial protoplanet region. Their eccentricities and inclinations are $\\simeq 0.1$. These results are independent of the accretion model. \\item The RMS spin angular velocity for the realistic accretion model is about 30\\% smaller than that for the perfect accretion model that is as large as the critical spin angular velocity for rotational instability. The spin angular velocity and obliquity of planets obey Gaussian and isotropic distributions, respectively, independently of the accretion model. \\end{itemize} We confirm that except for the magnitude of the spin angular velocity, the realistic accretion model gives the same results as the perfect accretion model. This agreement justifies the use of the perfect accretion model to investigate the basic dynamical properties of planets except for the magnitude of the spin angular velocity. In the present realistic accretion condition, we do not consider the effect of the spin on the accretion condition that would potentially change collisional dynamics. It is difficult to derive an accretion condition in terms of spin parameters by SPH collision simulations since the collision parameter space becomes huge and it is almost impossible to cover all parameter space. The fragmentation of planets is not taken into account, either. Including the fragmentation may be able to further reduce the spin angular velocity of planets by producing unbound collisional fragments with high angular momentum. Furthermore, the collisional fragments can potentially alter orbital dynamics of planets through dynamical friction if their mass is large enough. The high eccentricities and inclinations of planets may be damped by dynamical friction from the collisional fragments. In order to take into account these effects and make the model of terrestrial planet formation more realistic, we plan to perform $N$-body simulations for orbital dynamics and SPH simulations for collisional dynamics simultaneously in a consistent way in future work. \\medskip We thank Shigeru Ida for his continuous encouragement. This research was partially supported by MEXT (Ministry of Education, Culture, Sports, Science and Technology), Japan, the Grant-in-Aid for Scientific Research on Priority Areas, ``Development of Extra-Solar Planetary Science,'' and the Special Coordination Fund for Promoting Science and Technology, ``GRAPE-DR Project.''" }, "1003/1003.2985_arXiv.txt": { "abstract": "{}{We present polarimetric observations in the UBVRI bands of 72 stars located in the direction of the medium age open cluster NGC 5617. Our intention is to use polarimetry as a tool membership identification, by building on previous investigations intended mainly to determine the cluster's general characteristics rather than provide membership suitable for studies such as stellar content and metallicity, as well as study the characteristics of the dust lying between the Sun and the cluster.} {The obsevations were carried out using the five-channel photopolarimeter of the Torino Astronomical Observatory attached to the 2.15m telescope at the Complejo Astron\\'omico El Leoncito (CASLEO; Argentina).} {We are able to add 32 stars to the list of members of NGC 5617, and review the situation for others listed in the literature. In particular, we find that five blue straggler stars in the region of the cluster are located behind the same dust as the member stars are and we confirm the membership of two red giants. The proposed polarimetric memberships are compared with those derived by photometric and kinematical methods, with excellent results. Among the observed stars, we identify 10 with intrinsic polarization in their light. NGC 5617 can be polarimetrically characterized with $ P_{max}= 4.40 \\%$ and $ \\theta_{v}= 73^\\circ.1$. The spread in polarization values for the stars observed in the direction of the cluster seems to be caused by the uneven distribution of dust in front of the cluster's face. Finally, we find that in the direction of the cluster, the interstellar medium is apparently free of dust, from the Sun's position up to the Carina-Sagittarius arm, where NGC 5617 seems to be located at its farthest border.} {} ", "introduction": "The polarimetric technique is a very useful tool for obtaining significant information (e.g. magnetic field direction,\\ $ \\lambda_{max}$, $P_{\\lambda_{max}}$) about the dust located in front of a luminous object. Open clusters are very good candidates for carrying out polarimetric observations, because previous photometric and spectroscopic studies of these clusters have provided detailed information about the color and luminosity of the main-sequence stars. In addition to the cluster physical parameters obtained by using those tools (e.g. age, distance, extinction, membership), the polarimetric data allow us study the location, size, and efficiency of the dust grains to polarize the starlight and the different directions of the Galactic magnetic field along the line of sight to the cluster. Since the open clusters are also spread within a fixed area, we can analyze the evolution in the physical parameters of the dust all over the region. In this framework, we conduct systematic polarimetric observations in a large number of Galactic open clusters. As part of this survey, we present a polarimetric investigation of the open cluster NGC 5617 (C1426-605). It is located at (l = $ 314\\fdg7$, b = $ -0\\fdg1$), covering a wide area of the sky of about 10 x 10 arc minutes. In the past it was investigated using photoelectric and photographic photometry (Lindoff 1968; Moffat $ \\&$ Vogt 1975; Haug 1978). Kjeldsen $ \\&$ Frandsen (\\cite{KF91}) and Carraro $ \\&$ Munari (\\cite{CM04}) presented CCD observations covering part of the area. In this last work deep CCD (BVI) photometry of the core region was performed to derive more accurate estimates of the cluster fundamental parameters by observing about 140 stars down to V = 17.5 mag. The analysis of two adjacent fields covering the central part of the cluster confirmed a previous mean value of the excess $ E_{B-V}$= 0.48 mag, found in the first of two CCD investigations, and a distance determination that located the open cluster at 2.0 $\\pm$ 0.3 kpc from the Sun. It is an intermediate-age open cluster ($8.2~ 10^7$~years) containing red giants and blue straggler stars (Ahumada \\& Lapasset \\cite{AL07}) in its surroundings, whose memberships of the cluster remain in doubt. ", "conclusions": "\\subsection{Fitting with Serkowski's law} To analyze the data, the polarimetric observations in the five filters were fitted for each star using Serkowski's law of interstellar polarization (Serkowski 1973) given by $$ P_{\\lambda}/P_{\\lambda max}=e^{-Kln^2(\\lambda_{max}/\\lambda)}. \\ \\ \\ \\ (1)$$ If polarization is produced by aligned interstellar dust particles, then we assume that, the observed data (in terms of wavelength within the bands UBVRI) can be reproduced accurately by Eq (1) and that each star has a separate $\\lambda_{max}$ and a $P_{\\lambda_{max}}$ value. To perform the fitting, we adopted $ K=~1.66 \\lambda_{max} + 0.01$, where $\\lambda_{max}$ is in $ \\mu$m (Whittet et al. \\cite{WMHRetal92}). For each star, we also computed the $\\sigma_{1}$ parameter (the unit weight error of the fit) in order to quantify the departure of our data from the ``theoretical curve'' of the Serkowski's law. In our scheme, when a star exhibits $ \\sigma_{1} > 1.80$ this is indicative of a non-interstellar origin (that is, an intrinsic polarization) in part of the measured polarization. The dominant source of intrinsic polarization is dust non-spherically distributed and, for classical Be stars, electron scattering. The $\\lambda_{max}$ values can also be used to test the origin of the polarization: objects with a $\\lambda_{max}$ much shorter than the average value for the interstellar medium (0.55 $\\mu m$, Serkowski et al. \\cite{SMF75}) are also likely to contain an intrinsic component of polarization (Orsatti et al. \\cite{OVM98}). The individual $P_{\\lambda_{max}}$, $\\sigma_{1}$, $ \\lambda_{max}$, and $\\bar{\\epsilon}$ values, together with the star identification from Haug (\\cite{H78}), are listed in Table 2. We excluded five stars with $ \\epsilon _{P_{max}}$ higher than 15$ \\%$: \\#58, 69, 137, 149, and 255. The mathematical expression used to obtain the individual $\\sigma_{1}$ values is found in this table as a footnote. According to Table 2, only 10 of the 67 stars exhibit signatures of intrinsic polarization: \\#154,~260,~270, and 274 (a group with very high $\\sigma_{1}$); \\#146,~330, and 333 (with lower values): and also two blue stragglers (\\#195,\\#261) and the red giant \\#227. Star 226 has $\\sigma_{1}$= 1.91 but this value was estimated using data for only 3 filters, so the detection of intrinsic polarization in the star is dubious. The use of the second criterion to detect intrinsic stellar polarization did not provide new candidates. Figure 3 shows, for some of these stars, both the polarization and position angle dependence on wavelength. For comparison purposes, the best fit Serkowski's law for an interstellar origin of the polarization has been plotted as a continuous line. In the individual plots, the ${P_\\lambda} $ values do not appear to fit this law and in some cases (e.g., \\#227,\\#261) there is evidence of a combination of two different polarization mechanisms. Most of the stars in the figure also show important rotations of the polarization position angle with $ \\lambda$. \\subsection{The $ Q_{v}\\ versus \\ U_{v}$ plot and membership review} The identification of members in a cluster is important, not only to distance and age determinations but also other studies such as those of stellar content and metallicity. In NGC 5617, as in many other clusters, probable cluster members were identified as those stars simultaneously having reconcilable positions in both color-color and color-magnitude diagrams. Several errors in membership assignment are possible using this kind of photometric approach, for example when dealing with photographic photometry, in particular the U-measures; and when studying intermediate-age clusters, where the evolved stars could not be located close to the ZAMS. In our cluster, the CCD plates of Kjeldsen $ \\&$ Frandsen (\\cite{KF91}) and Carraro $ \\&$ Munari (\\cite{CM04}) cover only the central region of the cluster, and for stars in the vicinities of the core only photographic measurements are possible. In addition, as found by these last authors, stars brighter than 12.5 mag are evolving away from the main sequence. Evolved members and background stars become mixed close to the ZAMS in the color-magnitude diagrams, and in this case the photometric identification of cluster members becomes difficult. The polarimetric technique can help us to solve membership problems. Different plots used in combination with photometric information can be useful for separating between members and non-members. One of those plots presents the Stokes parameters $ Q_{v}\\ versus\\ U_{v}$ for the V-bandpass, where $ Q_{v} = P_{v}~ cos(2\\theta_{v})$ and {\\bf $ U_{v} = P_{v}~ sin (2\\theta_{v})$} are the components in the equatorial system of the polarization vector $ P_{v }$, is shown in Figure 4. The plot illustrates the variations occurring in interstellar environments. Since the light from cluster members must have traversed a common sheet of dust, of particular polarimetric characteristics, the member data points should occupy similar regions of the figure. Non-member stars (frontside and background stars) should be located in the $ Q_{v}\\ versus \\ U_{v}$ figure somewhat apart from the region occupied by member stars, since their light must have traveled through different dust clouds from those affecting the light of member stars, of different polarimetric characteristics. The basic principle behind the use of polarimetry as a criterion for distinguishing members from non-members in a cluster is similar to that used in photometry to decide membership. The procedure is based on the assumption that member stars are located behind common dust clouds that polarize their light, while this is not valid for most non-member stars. Objects closer to the Sun, or located along the line of sight to the cluster, will have lower $ E_{B-V}$ and their light will be less polarized. If there are clouds between these stars and ourselves, and dust is orientated in a different direction, the final angle will not be the same as for the cluster's stars. In the Q versus U plot, these stars therefore are located in different regions. For stars located behind the cluster, the individual polarizations could be higher than those associated with the cluster if the dust has the same orientation, or it could be depolarized if the orientation is not the same. However, in both cases the location in the diagram Q-U will not be the same as that for the cluster's stars. And in the last case, the $ E_{B-V}$ of those stars will be higher and detectable in the efficiency diagram ( $ P_{v} $ versus $ E_{B-V}$). For example, the polarimetry became a very useful membership identification tool for stars belonging to the clusters NGC 6204 and Hogg 22, the last one located behind the first cluster. Stars of each object occupy different regions of the Q versus U diagram (see Fig. 6 of Mart\\'{\\i}nez et al. 2003), and since Hogg 22 is depolarized and has a higher $ E_{B-V}$'s, both clusters are separated in the polarization efficiency diagram (Fig. 5 of that work). In using polarimetry to decide memberships, we are able to use information in the literature derive using alternative methods (e.g., photometric, spectroscopic, proper motions) and we can determine in the Q - U and $ P_{v}$ versus $ E_{B-V}$ plots the regions where cluster and non-member stars are located. Because of previous considerations, we accept at first that any star in the non-member region is a very probable non-member and in the same way, any star located in the cluster region is also a very probable member of the cluster. We note that we do not apply this criterion to stars with detected intrinsic polarization for which we can only assign a dubious membership. Because this tool has been used in previous work, our opinion is that polarimetry can be as useful as photometry to identifications or, in most of the cases, a very useful complement. Members identified by applying this method are shown in both Figs. 4 and 5 using filled (for member) and open (for non-member) symbols. Circles are used for stars, squares for supergiants, and starred points for blue stragglers. Small symbols of any kind denote stars with intrinsic polarization. We could add an important number of new members to the previous list, and we have also been able to review the situation for other stars listed in the photometric investigations. The last column of Table 2 lists our conclusions. In particular, five of the observed blue straggler stars are behind the dust located in front of the member stars, but \\#195 may be a possible member because of its position in the $ Q_{v}\\ versus \\ U_{v}$ plot. As mentioned in the previous section, intrinsic polarization was detected for the star that could explain the position in Fig. 4. We also confirmed that the red giants \\# 116 and 227 are members of the open cluster, as asserted in the literature. To compare our proposed polarimetric memberships with those coming from other methods, we used the works of Mermilliod et al. (\\cite{MMU08}) and Frinchaboy \\& Majewski (\\cite{FM08}). In the first investigation, radial velocities of giant stars in the region of a cluster are compared with the cluster mean velocity to assign membership. We have four red giants in common with that study (\\#55, 116, 227, 347) and our membership results are in agreement. In the second investigation, the star proper motion (from Hipparcos and Tycho-2 catalogues), radial velocity, and spatial distribution are combined to detect cluster members. Among a group of seven stars in common (\\#55, 116, 180, 202, 227, 342, 347), there are membership discrepancies for only star \\# 202 ($ P_{max}$=4.83$\\%, \\theta_{v}$= 76$\\fdg1$), for which they find a probability of 51.2$\\%$ on the basis of its radial velocity but a 0$\\%$ probability using the proper motions. According to both Figs. 4 and 5, the star appears to be located behind the same sheet of dust as the remaining members, and for that we consider the star as a member, as the photometric plots suggest. Column 6 in Table 2 lists the memberships determined in these two works. In Fig. 4, stars \\#187, 270, and 333 are located far away from the member group. The first star is considered to be part of the cluster by some photometric studies but no membership information is provided by Mermilliod et al. (\\cite{MMU08}), Frinchaboy \\& Majewski (\\cite{FM08}), or Dias et al. (\\cite{DAFAL06}). The $ \\sigma_{1}$ is not indicative of abnormal polarization as to justify its position in the plot; and it can be seen in Fig. 5 that it has a low polarization relative to the remaining members. We propose that \\#187 is a frontside star, observed projected onto the central core of NGC 5617. The other two stars (\\#270 and 333) both display intrinsic polarization as mentioned in the following section. To derive mean values of polarization and polarization angle, we used 16 stars with similar Stokes parameters and free of intrinsic polarization: \\#79, 80, 93, 94, 100, 106, 116, 169, 180, 185, 186, 214, 225, 276, 294, and 302. We obtained $ P_{max}$ = 4.40$\\%$ and $ \\theta_{v}$ = 73$\\fdg$1$\\ \\pm$0.9 (both of them being the mean values for all 16). The mean $ \\lambda_{max}$ amounts to 0.53 $\\pm$ 0.03 $\\mu$m, the value associated with the ISM. \\subsection{Polarization efficiency} It is known that for the interstellar medium the polarization efficiency (ratio of the maximum amount of polarization to visual extinction) rarely exceeds the empirical upper limit, $$P_{\\rm max} \\leq~3~A_{v}~\\simeq~3~R_{v}~E_{B-V}~~~~~ (2)$$ obtained for interstellar dust particles (Hiltner 1956). The polarization efficiency indicates how much polarization is obtained for a certain amount of extinction and depends mainly on both the alignment efficiency and the magnetic field strength, and also on the amount of depolarization due to radiation traversing more than one cloud with different field directions. Figure 5 shows the plot of $ P_{max}\\ vs.\\ E_{B-V}$. The individual excesses $ E_{B-V}$ were obtained either from the literature or by dereddening the colors and using the relationship between either spectral type and color indexes (Schmidt-Kaler \\cite{SK82}). For stars with only photographic UBV measures, the calculated excesses may well be in error. It can be seen that, apart from five stars (from top to bottom: \\#333, 294, 261, 136, and 269), the remainder are located to the right of the interstellar maximum line, indicating that their polarizations are mostly due to the ISM. Star \\#333 has intrinsic polarization in Table 2 and, based on its position in the plot, we understand it to be a background star. Star \\#294, the second from top, has no evident indications of abnormal polarization in its light and we could not find any suitable explanation of the position in this figure. Most probably, the excess (calculated from photographic photometry) is affected by error. Stars \\# 261 (a blue struggler) and \\#136 are affected by intrinsic polarization as shown in Table 2; with respect to star \\#269, the calculated excess may well also be affected by error. As in Fig. 4, the two stars \\#187 and 333 are shown here as non-members. But regarding \\#270, we accept that this star is a member, beacuse even when its polarization includes an intrinsic part, its position in Fig. 5 matches those of the member stars, and its $ P_{max}$ and $ \\theta_{v}$ values (4.31 and 73$\\fdg$4, respectively) are coincident with the mean values for members. To calculate the polarization efficiency, we selected a group of 15 stars with $ P_{v}$ in the range from 4.20$ \\%$ to 4.95$ \\%$, and obtained a polarization efficiency of about 2.44, which is lower than the standard value for the interstellar dust (of about 5). This value indicates a very high efficiency of the dust that polarizes the light from the cluster stars. Neckel \\& Klare (1980) computed the interstellar extinction values and distances of more than 11000 O to F stars in the Milky Way. Their figure \\# 174 (314$^\\circ$, 0$ ^\\circ$) shows the variation in Av with increasing distance in the area of NGC 5617, starting at 1 kpc and showing that the absorption takes values of between 1.2 and 2.4 mag at the position of the cluster, in good agreement with the Av calculated from (2) for a $ P_{max}$ of 4.4$ \\%$. The nearest of our stars to the Sun is \\#187 (non-member) with a $ P_{v}$= 1.83$ \\%$ and a distance of about 1.2 kpc from us, which it implies that not in the Local but in the Carina-Sagittarius arm, in addition to the remaining non-members and the cluster itself. As can be seen in Fig. 5, there is a scatter in polarization values for the members of NGC 5617, which could be caused by either intracluster dust or the uneven distribution of dust in front of the cluster's face, as seen in any plate of the object. Since NGC 5617 is an intermediate-age open cluster ($8.2 10^7$ yr), we favor the second explanation of the scatter." }, "1003/1003.5277_arXiv.txt": { "abstract": "In the earlier work on the development of a model--independent data analysis method for determining the mass of Weakly Interacting Massive Particles (WIMPs) by using measured recoil energies from direct Dark Matter detection experiments directly, it was assumed that the analyzed data sets are background--free, i.e., all events are WIMP signals. In this article, as a more realistic study, we take into account a fraction of possible residue background events, which pass all discrimination criteria and then mix with other real WIMP--induced events in our data sets. Our simulations show that, for the determination of the WIMP mass, the maximal acceptable fraction of residue background events in the analyzed data sets of ${\\cal O}(50)$ total events is $\\sim$ 20\\%, for background windows of the entire experimental possible energy ranges, or in low energy ranges; while, for background windows in relatively higher energy ranges, this maximal acceptable fraction of residue background events can not be larger than $\\sim$ 10\\%. For a WIMP mass of 100 GeV with 20\\% background events in the windows of the entire experimental possible energy ranges, the reconstructed WIMP mass and the 1$\\sigma$ statistical uncertainty are $\\sim 97~{\\rm GeV}\\~^{+61\\%}_{-35\\%}$ ($\\sim 94~{\\rm GeV}\\~^{+55\\%}_{-33\\%}$ for background--free data sets). ", "introduction": "Currently, direct Dark Matter detection experiments searching for Weakly Interacting Massive Particles (WIMPs) are one of the promising methods for understanding the nature of Dark Matter and identifying them among new particles produced at colliders as well as reconstructing the (sub)structure of our Galactic halo \\cite{Smith90, Lewin96, SUSYDM96, Bertone05}. In order to determine the mass of halo WIMPs {\\em without} making any assumptions about their density near the Earth or their velocity distribution {\\em nor} knowing their scattering cross section on nucleus, a model--independent method by combining two experimental data sets with two different target nuclei has been developed \\cite{DMDDmchi-SUSY07, DMDDmchi}. This method builds on the earlier work on the reconstruction of the (moments of the) one--dimensional velocity distribution function of halo WIMPs, $f_1(v)$, by using data from direct detection experiments \\cite{DMDDf1v}. In the analysis of reconstructing $f_1(v)$, the moments of the WIMP velocity distribution function can be determined from experimental data directly with an {\\em unique} input information about the WIMP mass $\\mchi$. Hence, one can simply require that the values of a given moment of $f_1(v)$ determined by two experiments agree% \\footnote{ Note that, as demonstrated and discussed in Ref.~\\cite{DMDDmchi}, this condition requires an algorithmic procedure for matching the maximal cut--off energies of the analyzed data sets. }. This leads to a simple analytic expression for determining $\\mchi$ \\cite{DMDDmchi-SUSY07, DMDDmchi}, where each moment can in principle be used. Additionally, under the assumptions that the spin--independent (SI) WIMP--nucleus interaction dominates over the spin--dependent (SD) one and the SI WIMP coupling on protons is approximately the same as that on neutrons, a second analytic expression for determining $\\mchi$ has been derived \\cite{DMDDmchi}. Finally, by combining the first estimators for different moments with each other and with the second estimator, one can yield the best--fit WIMP mass as well as minimize its statistical uncertainty. In the work on the development of the model--independent data analysis procedure for the determination of the WIMP mass, it was assumed that the analyzed data sets are background--free, i.e., all events are WIMP signals. Active background discrimination techniques should make this condition possible. For example, the ratio of the ionization to recoil energy, the so--called ``ionization yield'', used in the CDMS-II experiment provides an event--by--event rejection of electron recoil events to be better than $10^{-4}$ misidentification \\cite{Ahmed09b}. By combining the ``phonon pulse timing parameter'', the rejection ability of the misidentified electron recoils (most of them are ``surface events'' with sufficiently reduced ionization energies) can be improved to be $< 10^{-6}$ for electron recoils \\cite{Ahmed09b}. Moreover, as demonstrated by the CRESST collaboration \\cite{CRESST-bg}, % by means of inserting a scintillating foil, which causes some additional scintillation light for events induced by $\\alpha$-decay of $\\rmXA{Po}{210}$ and thus shifts the pulse shapes of these events faster than pulses induced by WIMP interactions in the crystal, the pulse shape discrimination (PSD) technique can then easily distinguish WIMP--induced nuclear recoils from those induced by backgrounds% \\footnote{ For more details about background discrimination techniques and status in currently running and projected direct detection experiments, see e.g., Refs.~\\cite{Aprile09a, EDELWEISS-bg, % Lang09b} % }. However, as the most important issue in all underground experiments, the signal identification ability and possible residue background events which pass all discrimination criteria and then mix with other real WIMP--induced events in our data sets should also be considered. Therefore, in this article, as a more realistic study, we take into account different fractions of residue background events mixed in experimental data sets and want to study how well the model--independent method could reconstruct the input WIMP mass by using these ``impure'' data sets and how ``dirty'' these data sets could be to be still useful. The remainder of this article is organized as follows. In Sec.~2 we review the recoil spectrum of elastic WIMP--nucleus scattering and introduce two kinds of background spectrum used in our simulations. In Sec.~3 we first review briefly the model--independent method for the determination of the WIMP mass. Then we show numerical results of the reconstructed WIMP mass by using mixed data sets with different fractions of residue background events based on Monte Carlo simulations. We conclude in Sec.~4. Some technical details will be given in an appendix. ", "conclusions": "In this paper we reexamine the model--independent data analysis method introduced in Refs.~\\cite{DMDDmchi-SUSY07, DMDDmchi} for the determination of the mass of Weakly Interacting Massive Particles from data (measured recoil energies) of direct Dark Matter detection experiments directly by taking into account a fraction of residue background events, which pass all discrimination criteria and then mix with other real WIMP--induced events in the analyzed data sets. Differ from the maximum likelihood analysis described in Refs.~\\cite{Green-mchi07, Green-mchi08, Bernal08}, our method requires {\\em neither} prior knowledge about the WIMP scattering spectrum {\\em nor} about different possible background spectra; the unique needed information is the recoil energies recorded in {\\em two} direct detection experiments with {\\em two different} target nuclei. In Sec.~2 we considered first the measured energy spectrum for different WIMP masses with two forms of possible residue background spectrum: the {\\em target--dependent exponential} spectrum and the {\\em constant} spectrum. The exponential background spectrum contributes relatively more events to {\\em high} energy ranges once WIMPs are {\\em light} ($\\mchi~\\lsim~100$ GeV), and to {\\em low} energy ranges for {\\em heavy} WIMP masses ($\\mchi~\\gsim~100$ GeV); whereas the constant background spectrum contributes always relatively more events to {\\em high} energy ranges. As the consequence, the energy spectrum of all observed events looks more likely to be a scattering spectrum induced by {\\em heavier} WIMPs, once the spectrum of residue background events (induced perhaps by two or more different sources) is either exponential--like (and WIMPs are light) or approximately constant (for all WIMP masses); while if WIMPs are heavy and the residue background spectrum is approximately exponential, the measured energy spectrum would look more likely to be a scattering spectrum induced by {\\em lighter} WIMPs. In Sec.~3.2 the data sets generated in Sec.~2 have been analyzed for reconstructing the mass of incident WIMPs by using the model--independent method. With the exponential background spectrum, the input WIMP mass would be {\\em overestimated} once WIMPs are light ($\\mchi~\\lsim~100$ GeV), or, in contrast, would be {\\em underestimated} for heavy WIMPs ($\\mchi~\\gsim~100$ GeV). Our simulations show that, for background windows in the {\\em entire or low} experimental possible energy ranges, one could in principle reconstruct the WIMP mass with a maximal fraction of $\\sim$ 20\\% of residue background events in the analyzed data sets; whereas for background windows in {\\em high} energy ranges, the maximal acceptable fraction of residue backgrounds is only $\\sim$ 10\\%. Moreover, in order to check the need of a prior knowledge about an (exact) form of the residue background spectrum, we considered also the case with the constant background spectrum. In this rather extrem case, the WIMP mass would always be {\\em overestimated}, especially for heavy WIMPs ($\\mchi~\\gsim~100$ GeV). Our simulations give then a maximal acceptable fraction of \\mbox{$\\sim$ 5\\% -- 10\\%} of residue background events in the data sets for background windows in the {\\em entire or low} experimental possible energy ranges. Nevertheless, we found also that, by means of increased number of observed (WIMP--induced) events and improved background discrimination techniques \\cite{CRESST-bg, % EDELWEISS-bg}, % the WIMP mass could in principle be determined (pretty) precisely, no matter what kind of energy spectrum residue background events would have. On the other hand, in order to check the statistical fluctuation of the reconstructed WIMP mass with increased background ratio, we considered also the distribution of the deviation of the reconstructed WIMP mass from the true value. It was found in Ref.~\\cite{DMDDmchi} that, for a rather heavy WIMP mass of 200 GeV, the distribution of the deviation of the reconstructed WIMP mass is asymmetric and non--Gaussian, either with data sets of only a few ($\\cal O$(50)) events or with larger date sets (of $\\cal O$(500) events). However, our simulations with different background ratios show that, firstly, for both used (exponential and constant) background spectra, with increasing background ratio the distribution of the deviation of the reconstructed WIMP mass becomes more and more concentrated, although still asymmetric and non--Gaussian. Secondly, for the more realistic exponential background spectrum and using data sets with a larger number of total events, with increasing background ratio the distribution of the deviation becomes somehow more symmetric and Gaussian. This observation might be able to offer some new ideas for improving the algorithmic procedure for the reconstruction of the WIMP mass with a higher statistical certainty. In summary, our study of the effects of residue background events in direct Dark Matter detection experiments on the determination of the WIMP mass shows that, with currently running and projected experiments using detectors with $10^{-9}$ to $10^{-11}$ pb sensitivities \\cite{Baudis07a, Drees08, Aprile09a, Gascon09} and $< 10^{-6}$ background rejection ability \\cite{CRESST-bg, % EDELWEISS-bg, % Lang09b, % Ahmed09b}, % once two or more experiments with different target nuclei could accumulate a few tens events (in one experiment), we could in principle already estimate the mass of Dark Matter particle with a reasonable precision, even though there might be some background events mixed in our data sets for the analysis% \\footnote{ A possible first test could be a combination of the events observed by the CoGeNT experiment with their Ge detector with the events observed in the oxygen band of the CRESST-II experiment \\cite{Aalseth10, CRESST-Talk, Hooper10}. }. Moreover, two forms for background spectrum and three windows for residue background events considered in this work are rather naive. Nevertheless, one should be able to extend our observations/discussions to predict the effects of possible background events in their own experiment. Hopefully, this will encourage our experimental colleagues to present their (future) results not only in form of the ``exclusion limit(s)'', but also of the ``most possible area(s)'' on the cross section versus mass plan. \\subsubsection*" }, "1003/1003.5107_arXiv.txt": { "abstract": "{} {Massive stars are known to demonstrate significant spectroscopic and photometric variability over a wide range of timescales. However the physical mechanisms driving this behaviour remain poorly understood. Westerlund 1 presents an ideal laboratory for studying these processes in a rich, coeval population of post-main sequence stars and we present a pathfinding study aimed at characterising their variability.} {To do this we utilised the large body of spectroscopic and photometric data that has accumulated for Wd1 during the past decade of intensive studies, supplemented with the sparser historical observations extending back to the early 1960s.} {Despite the heterogeneous nature of this dataset, we were able to identify both spectroscopic and photometric variability amongst every class of evolved massive star present within Wd~1. Spectroscopic variability attributable to both wind asphericity and photospheric pulsations was present amongst both the hot and cool hypergiants and the former, also with the Wolf Rayets. Given the limitations imposed by the data, we were unable to determine the physical origin of the wind structure inferred for the OB supergiants, noting that it was present in both single pulsating and binary stars. In contrast we suspect that the inhomogineities in the winds of the Wolf Rayets are driven by binary interactions and, conversely, by pulsations in at least one of the cool hypergiants. Photospheric pulsations were found for stars ranging from spectral types as early as O9 I through to the mid F Ia$^+$ yellow hypergiants - with a possible dependence on the luminosity class amongst the hot supergiants. The spectroscopically variable red supergiants (M2-5 Ia) are also potential pulsators but require further observations to confirm this hypothesis. Given these findings it was therefore rather surprising that, with the exception of W243, no evidence of the characteristic excursions of both luminous blue variables and yellow hypergiants was found. Nevertheless, future determination of the amplitude and periodicity of these pulsations as a function of temperature, luminosity and evolutionary state holds out the tantalising possibility of constraining the nature of the physical mechanisms driving the instabilities that constrain and define stellar evolution in the upper reaches of the HR diagram. Relating to this, the lack of secular evolution amongst the cool hypergiants and the presence of both high-luminosity yellow hypergiants and red supergiants within Wd1 potentially place strong constraints on post-main sequence evolutionary pathways, with the latter result apparently contradicting current theoretical predictions for $>25$M$_{\\odot}$ stars at solar metallicites.} {} ", "introduction": "A large body of observational evidence exists that massive post-main sequence (MS) stars are subject to a variety of (pulsational) instabilities (e.g. Humphreys \\& Davidson \\cite{hd94}, Sterken \\cite{sterken77}, Burki \\cite{burki}, de Jager \\cite{dj98}, Wood et al. \\cite{wood}). Indeed, it has been supposed that such instabilities contribute to the absence of cool (T$<10$kK) hypergiants above log(L/L$_{\\odot}$)$\\sim$5.8, giving rise to the the Humphreys Davidson limit that spans the upper reaches of the HR diagram. Characterisation of such instabilities is important for two reasons. Firstly, one might hope to use the pulsational properties to investigate the internal structure of massive stars via asteroseismology. Secondly, there is increasing reason to suppose that MS mass loss is insufficient to permit the transition of high-mass stars from unevolved O-type objects to H-depleted Wolf Rayet stars, and instead that the H-rich mantle is shed in short-lived impulsive events associated with the `zoo' of transitional objects. Indeed observational evidence of this hypothesis is provided by the presence of fossil ejecta surrounding luminous blue variables (LBVs: e.g. Clark et al. \\cite{c03}), yellow hypergiants (YHGs; Castro-Carrizo et al. \\cite{castro}) and red supergiants (RSGs; Schuster et al. \\cite{schuster}) and the direct association of dramatic increases in mass loss with episodes of instability in both hot and cool luminous stars ($\\eta$ Car and \\object{$\\rho$ Cas}, respectively; Smith \\& Gehrz \\cite{smith98}, Lobel et al. \\cite{lobel}). Thus, determining the duration, duty cycle and mass loss rate of such events appears to be necessary for the construction of a theory of massive stellar evolution. These instabilities are manifest in both photometric and spectroscopic variability, but other physical processes, such as large and small scale wind inhomogineities, also lead to similar behaviour. Line driven winds are intrinsically unstable, leading to stochastic clumping (e.g. Owocki \\cite{ow2000}), while hydrodynamical simulatons by Cranmer \\& Owocki (\\cite{cranmer}) demonstrate that large-scale photospheric perturbations - originating in non-radial pulsations and/or magnetic fields - can lead to large scale rotating wind structures. Moreover, spectrophotometric variability is an important observational signature of wind collision zones in massive binaries (e.g. Lewis et al. \\cite{lewis}, Stevens \\& Howarth \\cite{stevens}). Of these phenomena, the degree of wind clumping present in massive stellar winds is critical to the observational determination of the mass loss rate, while the characterisation of the properties of the massive binary population is important for fields as diverse as star formation, the nature of SNe progenitors and the formation channels of X-ray binaries. Consequently, considerable observational effort has been expended in the identification and characterisation of variability in massive (binary) stars. Such studies have typically centred on field stars, with the attendant difficulties of the determination of stellar luminosity and hence mass and age. However the development of multiplexing spectrographs and the identification of a large number of massive young Galactic clusters - of which Westerlund 1 is an examplar (henceforth Wd1) - presents a unique opportunity to advance this field, offering a large coeval population of massive stars of identical initial metallicity and well constrained masses. First identified by Westerlund (\\cite{westerlund61}), Wd1 was subject to sporadic optical and near-IR photometric studies in the following four decades (Sect. 2), with the high reddening towards it preventing a spectroscopic study until that of Westerlund (\\cite{westerlund87}; West87). However, despite the discovery of a remarkable population of both early and late super/hypergiants, the cluster languished in relative obscurity until the turn of the century, when modern spectroscopic studies - motivated by the unexpected {\\em radio} properties of Wd1 - first identified a rich population of Wolf-Rayets (WRs; Clark \\& Negueruela \\cite{c02}) and subsequently revealed Wd1 to be the first example of a super star cluster identified in the Galaxy (Clark et al. \\cite{c05}; henceforth C05). In the following decade, Wd1 has been the subject of a large number of studies from X-ray to radio wavelengths aimed at determining both bulk cluster properties and those of the individual constituent stars. The combined dataset therefore offers the potential of investigating the variability of cluster members over a significant time period ($\\geq$40~yr for the brightest stars; Sect. 2) and we present the results of such an analysis in this manuscript. In Sect. 2 we present and describe the reduction of the new data included in this work and discuss the utilisation of data available in previous studies. In Sect. 3-7 we analyse the available data as a function of evolutionary state and provide a discussion and summary in Sect. 8. ", "conclusions": "Despite the disparate nature of the photospheric and spectroscopic dataset accumulated for Wd1 over the past 50 years, it presents a valuable resource to investigate the properties of the massive stellar population and in particular to characterise the prevalence and nature of stellar variability. This is of particular relevance given the unique census of evolved stars within Wd1, comprising examples of all post-MS stellar types thought to originate from stars with M$_{initial} \\sim$30-40M$_{\\odot}$. Given the limited sampling of the combined dataset it was somewhat surprising to find that stars in {\\em all} evolutionary states were variable, with this behaviour being attributable to both evolving wind structure(s) and stellar pulsations. Indeed these phenomena appear to be commonplace, with 26 of the OB super-/hypergiants and all but one of both the cool hypergiants and the Wolf Rayets included in this study being either photometrically and/or spectroscopically variable. Wind variability was most evident in the H$\\alpha$ line of the OB super-/hypergiants, with the line transiting from absorption to emission in W70, and pronounced P Cygni profiles developing in several other stars. Similar behaviour has been observed in field stars although, as with this population, the physical mechanism(s) driving the wind inhomogineities are poorly understood. In this respect the detection of both photospheric pulsations and wind variability in the mid B supergiants W70 \\& 71 and hypergiants W7 \\& 42a is of considerable interest and all four stars would benefit from further observations with a higher cadence to investigate potential causal links between the two phenomena (c.f. Markova et al. \\cite{markova}). Moreover, these results again emphasise the uncertainty inherent in determining mass loss rates for OB supergiants from a single epoch of H$\\alpha$ spectroscopic data. Significant wind structure is also inferred for the WRs investigated. Given the identification of binary markers such as the presence of hot dust, hard X-ray emission and, in the case of W13 and 239, reflex RV variations, we attribute this to the effect of the companion via wind/wind interaction, noting that no secular evolution is present in any of the stars. This is encouraging since it suggests that detailed modeling of these variation may yield the underlying binary properties of the systems (e.g. Stevens \\& Howarth \\cite{stevens}) enabling a comparison to the less evolved stellar populations within Wd1, hence illuminating the effects of binary interaction driven by stellar evolution. Wind diagnostics are more difficult to identify and study in the cool hypergiants. Nevertheless variable high velocity absorption in the blue wing of the metallic photospheric lines is present in W265, indicative of changes in the optical depth and hence physical properties of the wind. By analogy to $\\rho$ Cas, we attribute this to pulsationally driven mass loss. Variability in the H$\\alpha$ emission line profile is inferred for W12a and directly observed in W16a, although the morphology of this line is complex and hence the physical origin of this behaviour is consequently uncertain. Variations in both profile strength {\\em and} the RV of the line centroid were found for the photospheric lines of both the hot and cool evolved stellar populations, the combination indicative of pulsations rather than binary motion (e.g. Ritchie et al. \\cite{benbin}). Amongst the OB supergiants within Wd1, this behavour had previously been reported for B0.5-2 Ia stars; our results extend this to encompass the full spectral range present (O9-B4 Ia). Utilising the new spectroscopic callibrations of Negueruela et al. (\\cite{n09}) we find that the presence of photometric variability (Bonanos \\cite{bonanos}) appears to be a strong function of evolutionary phase, being more prevalent amongst the early-mid B Ia stars in comparison to the late O Iab-II stars. This is qualitatively similar to the conclusions drawn by Fullerton et al. (\\cite{fullerton}) regarding the occurence and magnitude of photospheric pulsations in OB stars and hence it is tempting to attribute the photometric variability to photospheric pulsations, such as those inferred for the $\\alpha$ Cygni variables. Unfortunately, we currently lack the high S/N and cadence observations that would enable pulsational mode identification. Amongst the cooler super-/hypergiants, pulsations were inferred for stars with spectral types ranging from mid B through to mid F. The best studied example - W265 - was observed to transit between F0-5 Ia$^+$ over the course of a single cycle ($\\sim$100~days); the combination of both timescale, magnitude and atmospheric depth velocity stratification of the pulsations being strongly reminiscent of those observed in $\\rho$ Cas, which have been attributed to non-radial modes (Lobel et al. \\cite{lobel94}, \\cite{lobel}). Although the data are sparse, both W26 and 237 display spectroscopic variability indicative of changes in spectral type, with the former being observed to evolve from M2-5 Ia. The finding that at least two of the four RSGs within Wd1 are unstable and by analogy with other extreme RSGs, possibly pulsating is of considerable interest given {\\em (i)} their extreme luminosity and hence high progenitor mass\\footnote{ A lower limit to their progenitor masses may be inferred from the cluster Main Sequence turnoff. This is an important consideration given that the mass/luminosity degeneracy for RSGs as they execute a red loop prevents such a determination for field stars (Meynet \\& Maeder \\cite{meynet}).}, {\\em (ii)} their rich circumstellar envelopes, which provide evidence for significant recent mass loss, possibly driven by these instabilities and {\\em (iii)} the rarity of this phenomenon (with only six other examples known in the Galaxy and Magellanic Clouds; Levesque \\cite{review}). Indeed, these preliminary findings emphasise the important role studies of young massive clusters with multiplexing spectrographs will play in understanding the physics of high mass stars. In particular the detection of photospheric pulsations extending from spectral types as early as O9 I, through F5 Ia$^+$ hypergiants and potentially as late as the M5 Ia supergiants confirms the finding of e.g. Burki (\\cite{burki}) that stars in this region of the HR diagram are unstable, but critically {\\em for a rich co-eval population of stars at a single metallicity rather than a heterogeneous field population.} Given the sample size - $\\sim$80 $\\sim$O8III-B2 Ia stars - future analysis of the full 2008-9 VLT/FLAMES dataset will provide valuable observational constraints on the location of the high temperature boundary of this region of instability in a statistically robust manner. Furthermore, observational determination of the amplitude and period of the pulsations as a function of temperature, luminosity and evolutionary phase will provide direct contraints on the underlying physical mechanism(s) driving this instability - for example strange- and/or g-modes (Kiriakidis et al. \\cite{kiri}, Saio et al. \\cite{saio}) - and, in conjunction with chemical abundance analysis, a detailed picture of the evolution of the internal structure of massive stars as they execute a red loop on their way to becoming Wolf Rayets. Similarly, high candence observations will enable an investigation into the role pulsations play in both the origin of wind structure in hot supergiants and in driving mass loss in the cool hypergiants. Indeed, the uniquely rich population of both hot and cool super-/hypergiants within Wd1 holds open the possibility of investigating the varying contribution of both line and pulsational driven mass loss as a function of stellar temperature and luminosity as stars evolve across the HR diagram. Relating to this, with the exception of the known LBV W243 (Ritchie et al. \\cite{benLBV}), we find no compelling evidence for other LBV like excursions or eruptions within Wd1, despite both the large number of transitional stars present and the evidence for past episodes of enhanced mass loss amongst a subset of both hot and cool stars (Dougherty et al. \\cite{d09}). H$\\alpha$ LPV similar to that observed in the LBV HD160529 was found for several of the B super-/hypergiants, but the poor temporal coverage meant that we could not determine whether these corresponded to the long term secular changes that characterise LBV excursions, or the short lived variations indicative of an asymmetric wind. However, while similar limitations also apply to the (non-) detection of events comparable to the outbursts of $\\rho$ Cas and the red loops of HR8752 in the YHG population, we would have expected to detect events comparable to the multi-decade excursions of e.g. M33 Var A, or the long term secular evolution of IRC +10 420. Given these results, it is of interest that a comparable analysis for the WNLh stars - which have been proposed to be quiescent LBVs - within the Galactic Centre and Quintuplet cluster similarly reveals little evidence for significant long term variability (Appendix B). Apparently transitional stars co-located with both `bona fide' LBVs and rapidly evolving cool hypergiants in the HR diagram (e.g. Clark et al. \\cite{c05b}) can enjoy significant periods of quiescence. These results are surprising given that the majority of the stars observed within Wd1, from the extreme-B hypergiants through to the RSGs, appear to be low-level variables and only borderline-stable. As such one might naively have expected them to be prey to such large scale instabilities, which presumably contribute to the lack of stars above the HD limit. In this respect the YHGs and RSGs are of particular interest since they are located at the low temperature boundary of the HD limit and hence appear to represent the most massive stars that may exist in this region of the HR diagram at solar metallicities. Drout et al. (\\cite{drout}) present detailed evolutionary predictions for the physical properties of such stars, which we may test against the population within Wd1. Our new spectral classifications confirm the expectation that YHGs evolving from $\\sim$40M$_{\\odot}$ stars should have relatively high temperatures ($>6000$~K), while the lack of secular variability for any of these stars over a period of $\\sim$50yrs is consistent with the expectation of a relatively long lifetime for this phase ($\\sim$10$^4$-10$^5$~yr). However at these luminosities YHGs are {\\em not} expected to evolve to cooler temperatures, but exactly this evolution is indicated by the presence of 4 RSGs within Wd1. An extention of the redwards loop to encompass a RSG phase would result in a corresponding reduction in the time spent as a YHG. While rapid evolution through the YHG phase might be consistent with the real time increase in stellar temperature observed for IRC +10 420 and HD 179821, the continuing absence of this behaviour from the cool hypergiant population of Wd1 will place increasingly stringent constraints on such a scenario. Therefore a synthesis of {\\em (i)} the current physical properties of the evolved stars within Wd1 - such as chemical composition, luminosity, temperature and, where observable, surface gravity, {\\em (ii)} the relative populations of the different evolutionary classes and {\\em (iii)} the characterisation of their (pulsational) instabilities and secular variability has great potential for constraining the physics of massive post-MS evolution. Finally, regarding the ensemble properties of Wd1, constraining the occurence and amplitude of the pulsationally (and binary) driven RV changes will be crucial in allowing the accurate determination of dynamical masses for unresolved extragalactic clusters (e.g. Gieles et al. \\cite{gieles}), where the velocity dispersion attributable to such processes may match or even dominate virial motions. Likewise determining the properties of photometric outbursts driven by pulsational instabilities in both hot and cool hypergiants will be important in determining the validity of photometric cluster mass estimates, particularly for the lower range of masses where the presence of such an event may dominate the integrated light." }, "1003/1003.1724_arXiv.txt": { "abstract": "We present an X-ray morphological and spectroscopic study of the pulsar B2224+65 and its apparent jet-like X-ray features based on two epoch \\textit{Chandra} observations. The main X-ray feature, which shows a large directional offset from the ram-pressure confined pulsar wind nebula (Guitar Nebula), is broader in apparent width and shows evidence for spectral hardening (at 95 percent confidence) in the second epoch compared to the first. Furthermore, the sharp leading edge of the feature is found to have a proper motion consistent with that of the pulsar ($\\sim$180 mas yr$^{-1}$). The combined data set also provides evidence for the presence of a counter feature, albeit substantially fainter and shorter than the main one. Additional spectral trends along the major and minor axes of the feature are only marginally detected in the two epoch data, including softening counter to the direction of proper motion. Possible explanations for the X-ray features include diffuse energetic particles being confined by an organized ambient magnetic field as well as a simple ballistic jet interpretation; however, the former may have difficulty in explaining observed spectral trends between epochs and along the feature's major axis whereas the latter may struggle to elucidate its linearity. Given the low counting statistics available in the two epoch observations, it remains difficult to determine a physical production scenario for these enigmatic X-ray emitting features with any certainty. ", "introduction": "As pulsars age, they lose their spin-down energy through both radiation and relativistic particle ejection. One may approximate the particle ejection in two forms: anisotropic winds from the pulsar surface and bi-polar jets. Previous X-ray studies have shown the wind component to be primarily equatorial and axisymmetric, producing observed tori in pulsar wind nebulae (PWNe), whereas the jet outflows are only observed in a few, though well studied, young pulsars \\citep[e.g. the Crab and Vela pulsars,][]{weiss00,pavlov03}. Only a few nearby and relatively old pulsars have been shown to have their PWNe, mostly enhanced by ram-pressure confinement \\citep[e.g.,][and references therein]{wang93,karga08}. For a PWN or similar pulsar ejection feature to be observed in X-rays, the synchrotron-emitting particles must have energies on order of 10$^2$ TeV, achievable in shocks of the pulsar ejecta as it encounters the local interstellar medium (ISM). X-ray investigations of pulsar outflows are therefore crucial in understanding their origin and particle acceleration, which can in turn provide insight into the origin of other, more exotic and distant, outflows (e.g. from Active Galactic Nuclei). A particularly striking case is that of the pulsar B2224+65 and its X-ray-emitting ``jet''. B2224+65 is one of the fastest pulsars known with a radio proper motion of $\\mu=182\\pm 3$ mas yr$^{-1}$ \\citep{harrison93}; corresponding to a transverse velocity of $864 \\rm~D_{kpc}~km~s^{-1}$, where $\\rm~D_{kpc}$ is the distance to B2224+65 in units of kpc and is uncertain in the range of $\\sim 1-2$ \\citep{chat04}. The rapid motion of the pulsar produces a bow shock nebula visible in H$\\alpha$ \\citep{cordes93,chat02,chat04} which has been dubbed the ``Guitar Nebula'' due to its peculiar morphology. Aside from its high proper motion, B2224+65 behaves like a normal radio pulsar, with a modest energy spin-down rate of $\\dot{E} = 10^{33.1}\\rm~ergs~s^{-1}$ and period $P=0.68$ s. X-ray studies of B2224+65 do not reveal significant emission concurrent with the H$\\alpha$ \\citep[however, see][]{romani97} but rather an unusual extended linear feature apparently stemming from B2224+65 and offset from the pulsar's direction of motion by $\\sim$118$\\degr$ \\citep{hui07,wong03,zavlin04}. The feature, extending $\\sim$2$\\arcmin$ away from the pulsar, was found to have an apparent non-thermal spectrum though the details as to its origin have thus far remained a mystery. Such a distinct, linear X-ray emitting feature has not been identified anywhere else in the Galaxy with the possible exception of the Galactic Center \\citep[e.g.][and references therein]{johnson09}. It was suggested by \\cite{karga08} that the extended feature may not be associated with B2224+65 but rather with a nearby point source, although no physical production scenario was developed for the feature. Through theoretical modeling of the linear X-ray emission by \\cite{band08}, it was proposed that the feature could be produced through highly energetic particles that are accelerated at, and subsequently escape, the pulsar wind terminal shock but are ultimately confined by a large-scale, organized interstellar magnetic field. However, there are a number of observational requirements to the \\cite{band08} model that have yet to be fulfilled, primarily a proper motion of the feature consistent with B2224+65. The uncertain association with the pulsar, combined with low photon statistics, has made it difficult to establish the physical process responsible for the feature. \\begin{deluxetable*}{lcccccc} \\tablecaption{\\textit{Chandra} ACIS-S Observations} \\tablehead{\\textit{Chandra} &R.A. (J2000) &Dec. (J2000) & Exposure &Roll Angle &OBS Date\\\\ OBS. \\# &(h~~m~~s) &(degree~~arcmin~~arcsec)&(s) &(degree) &(yyyy-mm-dd)} \\startdata 755 &22 25 48.439 & +65 35 08.09 & 47865 & 236.8 &2000-10-21 \\\\ 6691 &22 25 51.852 & +65 35 39.19 & 9968 & 182.7 &2006-08-29 \\\\ 7400 &22 25 51.428 & +65 35 35.13 & 35839 & 219.3 &2006-10-06 \\enddata \\tablecomments{The exposure represents the live time (dead time corrected) of cleaned data.} \\end{deluxetable*} \\begin{figure*} \\includegraphics{B2224+65_smooth.ps} \\caption{Merged epoch \\textit{Chandra} intensity map of B2224+65 in the 0.3-7.0 keV energy band. The image has been smoothed using an adaptive Gaussian filter (FWHM$\\la$1$\\arcsec$) and is presented with logarithmic scaling. Overlaid on the intensity map is the direction of B2224+65's proper motion along with detected X-ray sources. Also included are the regions used for spectral extraction of the pulsar and jet (solid line), the local background region (the area within the dotted line excluding sources and the dotted rectangle), and the regions used in constructing surface brightness profiles along the major (dot-dash line) and minor (dashed line) axis. These regions all have position angles of $\\sim$24.7$\\degr$ as determined from the bright edge of the jet within $\\sim$50$\\arcsec$ from B2224+65.} \\end{figure*} In this study, we examine two epoch data taken by the \\textit{Chandra} X-ray Observatory to determine any association between the extended, linear feature and B2224+65 as well as to place constraints on possible X-ray-emitting particle production scenarios. Throughout this manuscript, we will refer to the feature as the jet or jet-like feature based purely on its apparent morphology, implying no physical interpretation given its still uncertain origin. From the two epoch data, we first correct for any astrometric offset between observations and verify the motion of B2224+65. We then briefly examine the two epoch and combined spectroscopic information for the pulsar to see if there may have been any evolution in its spectrum. A morphological analysis of the jet will serve to detect any shift between epochs, providing evidence for (or against) an association between the jet and pulsar. Spectral analysis of the jet between epochs and across the major/minor axes may aid in identifying possible spectral trends, which can assist in determining the physical models behind the X-ray emission. Finally, we briefly discuss our findings in the context of the \\cite{band08} model and a simple ballistic jet model similar to that used to explain AGN jet emission. \\begin{figure*}[t] \\includegraphics[width=\\linewidth]{B2224+65_dif_image.eps} \\caption{Intensity maps of B2224+65 and the jet with the source region from Fig.~1 overlaid: 0.3-7.0 keV images in the first (a) and second (b) epochs as well as the merged images in the soft (0.3-1.5 keV; c) and hard (1.5-7.0 keV; d) bands. All images are adaptively smoothed similar to Fig.~1 and are shown here using the same relative logarithmic scale with the minimum set by the median image value and the maximum set a factor of 10$^3$ times higher.} \\end{figure*} ", "conclusions": "Based on our two epoch X-ray analysis of B2224+65 and the jet-like feature, we have verified the proper motion of the pulsar and have shown similar motion in the sharp leading edge of the jet, providing strong evidence for its association with the pulsar. Along the major axis of the jet, a region of extended X-ray emission (hereafter counter feature), detected with $\\sim$4$\\sigma$ confidence, is found near the pulsar and opposite the jet. The jet also appears wider, due to a significant ($\\sim$5$\\sigma$) region of excess emission, with a harder spectrum in the second epoch compared to the first whereas additional spectral evolution across the major and minor axes is more ambiguous. Given the low confidence for many of the spectral variations, including softening counter to the direction proper motion, the physical nature of the jet-like feature remains puzzling. Here, we briefly examine two candidate physical scenarios for the production of the jet based on current results. To begin with, let us first consider the model proposed by \\cite{band08}. This model assumes that the jet represents electrons that have been accelerated at the termination shock of the pulsar wind which then leak out of the bow shock and diffuse into an ambient medium containing a large-scale plane-parallel magnetic field. The orientation of the jet-like feature feature would then correspond to the direction of the magnetic field, which needs to be unusually strong ($\\sim 45\\mu$G). As stated by \\cite{band08}, this scenario would predict motion for the jet-like feature consistent with B2224+65, softening of the jet spectrum counter to the direction of proper motion due to synchrotron cooling of particles injected at different times and the presence of a counter feature. Our two epoch observations support many of these predictions though the spectral variation can not be positively confirmed. Additional spectral variations, such as hardening between epochs and away from the pulsar along the jet major axis, are difficult to explain under this scenario. If the jet spectrum was observed to soften away from the pulsar, then one may naturally expect synchrotron cooling to be responsible; however, the observed hardening implies acceleration within the jet feature, unlikely for particles confined by a magnetic field. Hardening between epochs may suggest variation in the acceleration of the energetic particles, possibly due to changes in the pulsar wind and/or ISM properties. However, the lack of spectral evolution for the pulsar advocates against any such deviations at the terminal shock if, in fact, a significant portion of the pulsar emission results from the ISM-pulsar wind interaction. Furthermore, the relative faintness of the counter feature posses a challenge for the \\cite{band08} model, as the energetic electrons should have no preference in which direction along the magnetic field they travel, though this issue may be rectified by internal anisotropies and asymmetry of the pulsar wind. Alternatively, the B2224+65 extended X-ray features may represent the externally processed particles from jets, akin to those associated with AGNs. In this scenario, the main jet-like feature, counter feature and the Guitar Nebula represent the pulsar jets and equatorial wind outflow where the latter is responsible for the formation of the Guitar Nebula trailing the pulsar. X-ray emission is not expected in the pulsar wind forming the Nebula as the particles are not energetic enough to generate detectable X-ray synchrotron radiation. The pulsar jets, intrinsically relativistic and containing highly energetic particles, may easily escape the bow shock but are ultimately confined by the large ram-pressure generated by the pulsar's motion. In this scenario, the spectrum of the jet-like feature would soften counter to the direction of proper motion, similar to the \\cite{band08} model, due to synchrotron cooling. If enhancement of the X-ray emission due to the high ram-pressure contributes significantly to the X-ray emission, then spectral variation between epochs may arise from density variations in the ISM resulting in post-processing variations of the energetic particles (i.e. through their acceleration and/or synchrotron cooling efficiency). ISM variation over time has been seen previously by \\cite{chat04} where they found a decrease in the local density by a factor of $\\sim$0.7 between 1994 and 2001; however, any such density contrast can not be directly confirmed given the current two epoch \\textit{Chandra} observations and the relatively static pulsar spectrum. Spatial variations of the ISM could also explain the spectral hardening away from the pulsar through similar post-processing of the energetic particles. The relative faintness of the counter feature compared to the main jet-like feature may be due to Doppler effects if they are indeed relativistic and aligned along the line of sight; specifically, if $\\beta\\cos\\theta\\sim0.4$ where $\\beta$ is the bulk flow velocity and $\\theta$ is the angle along the line of sight. Further constraints may be possible by invoking more detailed jet mechanics and synchrotron theory. The ballistic jet scenario was deemed unlikely by \\cite{band08}, as it would predict a ratio between the length and width of the jet-like feature ($L_{jet}$ and $W_{jet}$, respectively) that is too small to match observations ($L_{jet}\\sim0.6$ pc and $W_{jet}\\lesssim0.04$ pc for D$_{\\rm{kpc}}=1$) along with significant bending in the jet based on comparisons of the bow shock pressure, transverse momentum added to the jet through the ISM ram-pressure and its the initial momentum. Following \\cite{band08}, we may estimate the expected bending angle of the jet-like feature under the ballistic jet model by assuming that the pulsar imparts a constant fraction $\\eta$ of its spin-down energy into the jet, which has a transverse size of $W_{jet}$. Balancing the transverse and initial momenta, the angle $\\theta$ describing the bending of the jet may be given by \\begin{equation} \\tan\\theta=\\frac{W_{jet}L_{jet}\\sin(118\\degr)\\delta \\mu m_{H} n_{ISM}v^2_*}{\\eta\\dot{E}/c} \\end{equation} where $\\delta$ and $\\mu$ are the ISM ionization fraction and mean molecular weight, respectively, and $v_*$ is the pulsar proper motion. If we assume a fully ionized ISM composed mostly of hydrogen ($\\delta=1$ and $\\mu=0.5$) with number density $n_{ISM}\\approx 0.1\\rm~cm^{-3}$ and dimensions for the proposed jet as taken from current observations, then such a ballistic jet should trail behind the pulsar regardless of the value of $\\eta$; strongly suggesting against the jet interpretation. However, the actual jet may be very narrow ($W_{jet}\\ll 0.01$ pc), similar to AGN jets, and/or the ISM could be largely neutral in which case the expected bending angle would be significantly reduced. A largely neutral ISM would contribute little to the ram-pressure as the mean free path of neutral atoms is much larger than the width of the feature ($\\sim$0.1 pc for neutral hydrogen atoms with density $0.1\\rm~cm^{-3}$); however, this would impose a fine-tuning problem on the jet-ISM interaction where the number density for the small ionized fraction must still be large enough to generate the confining ram-pressure and accelerate the energetic particles to X-ray emitting energies. An additional complication of the jet interpretation is the misalignment of the spin axis (believed to be correlated with pulsar jets) and the direction of B2224+65's proper motion; nevertheless, the method by which pulsars or their progenitors acquire such high velocities remains uncertain \\citep[see, for example,][]{hills88,yu03}. Certainly, additional and more detailed theoretical modeling, which may require additional observations in order to confirm apparent spectral trends, will be necessary to further test the jet and \\cite{band08} models." }, "1003/1003.1662_arXiv.txt": { "abstract": "Extrasolar planets are a natural extension of the interacting binaries towards the companions with very small masses and similar tools might be used to study them. Unfortunately, the generally accepted treatment of the reflection effect in interacting binaries is not very suitable to study cold objects irradiated by hot objects or extrasolar planets. The aim of this paper is to develop a simple model of the reflection effect which could be easily incorporated into the present codes for modeling of interacting binaries so that they can be used to study above mentioned objects. Our simple model of the reflection effect takes into account the reflection (scattering), heating and heat redistribution over the surface of the irradiated object. The shape of the objects is described by the Roche potential and limb and gravity darkening can be taken into account. The orbital revolution and rotation of the planet with proper Doppler shifts for the scattered and thermal radiation are also accounted for. Subsequently, light-curves and/or spectra of exoplanets were modeled and the effects of the heat redistribution, limb darkening/brightening, (non-)grey albedo, and non-spherical shape were studied. Recent observations of HD189733b, WASP12b, and Wasp-19b were reproduced reasonably well. HD189733b has low Bond albedo and intense heat redistribution. Wasp-19b has low Bond albedo and low heat redistribution. We also calculate the exact Roche shapes and temperature distribution over the surface of all 78 transiting extrasolar planets known so far. It is found that the departures from the sphere vary considerably within the sample. Departures of about 1\\% are common. In some cases (WASP-12b, WASP-19b, WASP-33b) departures can reach about 14, 12, and 8\\%, respectively. The mean temperatures of these planets also vary considerably from 300 K to 2600 K. The extreme cases are WASP-33b, WASP-12b, and WASP-18b with mean temperatures of about 2600, 2430, and 2330 K, respectively. ", "introduction": "There are sophisticated computer codes for calculating and inverting light curves or spectra of binary stars with various shapes or geometry including the Roche model \\citep{lucy68,wd71,wood71,md72,rucinski73,hill79,pe81,zhang86, djurasevic92,drechsel94,lh96,hadrava97,oh00,bs02,pribulla04, pavlovski06,tamuz06}. The Wilson \\& Devinney (WD) code is most often used and is continuously being improved or modified \\citep{kallrath98,pz05}. The reflection effect is taken into account in most of these codes. The standard description of this effect is given in \\citet{wilson90}. This effect is understood in the following way: A surface element of star A is irradiated by many surface elements of star B. A fraction of impinging energy called bolometric albedo \\citep{rucinski69} is converted into the heat which rises the local temperature and re-radiates the energy on the day side of the object. Rest of the impinging energy is plunged into the star. Typically, bolometric albedo is set to 1 for radiative and 0.5 for convective envelopes. An increase of the temperature on the day side of one object triggers a secondary reflection effect on the second object and one, two or several iterations are allowed to converge to a final state. Roche model, limb darkening and gravity darkening are taken into account. This model does a good job for many interacting binaries. However, I argue that the above mentioned model of the reflection effect should be revisited. There has been a lot of progress in the field since the time this reflection effect was developed. New types of very cool objects such as brown dwarfs and extrasolar planets have been discovered. These new areas evolve rapidly and produce interesting results. In many respects, an extrasolar planet can be understood as an interacting binary with a small mass companion. It is thus an attractive idea to apply these new results from the extrasolar planets to interacting binaries and vice versa. For example, models of extrasolar planets can provide day-night heat redistribution and reflected light while models of interacting binaries can provide sophisticated Roche geometry. The standard model of the reflection effect faces several problems which prevent its application to cool objects irradiated by hot objects and to extrasolar planets. When it comes to very cold strongly irradiated objects a considerable amount of energy might be reflected off the surface. The above mentioned reflection effect in the interacting binaries neglects this reflected light which can be essential, especially at the shorter wavelength. Well known example of such effect are the planets and moons in our Solar system in the visible light. This reflected light bears the spectroscopic signatures of the hot irradiating star and is not converted into the heat and re-radiated. It is commonly taken into account in models of hot-Jupiters or planets in general \\citep{seager00,sudarsky05}. The definition of the albedo in the planetary sciences is almost the opposite of its meaning in the interacting binaries. In planets, the albedo is a fraction of the impinging energy which is reflected off the object and is not absorbed and converted into the heat. Reflected and re-radiated photons may also have completely different Doppler shifts depending on the mutual velocities of the two objects and observer. Moreover, some portion of the energy absorbed on the day side can be transferred to and irradiated from the night side. Various authors developed different parametrization of this effect in connection with the extrasolar planets \\citep{guillot96,burrows06,ca10}. Calculations of atmosphere models of extrasolar planets \\citep{hubeny03,barman05,burrows08,fortney08} demonstrate that there is a deep temperature plateau which turns off the convection in the atmospheres. Convection and vertical energy transport operates only at very deep layers. Observations of hot-Jupiters indicate that only a very small fraction of the impinging radiation ($10^{-4}$) could have been plunged into the object \\citep{burrows07}. Hydrodynamical simulations \\citep{dl08,showman09,mr09} of the atmospheres of extrasolar planets indicate that there are very strong horizontal currents and jets which can effectively redistribute and circulate the energy between the day and night side of the planet especially along the lines of constant latitudes. How much energy gets redistributed to the night side depends mainly on the complex structure and dynamics of the surface layers. The above mentioned effects and findings must be taken into account in the new or revisited model of the reflection effect when modeling very cold components of interacting binaries and extrasolar planets. At the same time, many transiting exoplanets are so close to their host stars that their shape may depart from the sphere and be best described by the Roche model. The main purpose of this paper is to develop a new simple model of the reflection effect which would consider the reflected (scattered) light, heating (absorption of the light), and heat redistribution over the surface and which could be used to model the cold components of interacting binaries and extrasolar planets. The model was included into our code {\\sc{shellspec}} \\citep{br04,budaj05}. This program was originally designed to calculate light-curves, spectra and images of interacting binaries immersed in a moving circumstellar environment which is optically thin. The code is freely available at {\\tt http://www.ta3.sk/$\\sim$budaj/shellspec.html} and might be used to study various effects observed or expected in the interacting binaries or extrasolar planets. A Fortran90 version of the code which also solves the inverse problem was created by \\citet{tkachenko10}. ", "conclusions": "I argue that the reflection effect in the interacting binaries must be revisited in order to describe properly the radiation from the cool object irradiated by the hot object. Subsequently, a new model for the reflection effect was proposed and applied to an extreme case of an interacting binary - extrasolar planet. The new model introduces several free parameters. Some of them are well known in the field of extrasolar planets. (1) Bond albedo $A_{B}=<0,1>$ which controls how much energy is reflected and how much is converted into the heat. $A_{B}=1$ means that all impinging energy is reflected off the surface and nothing gets converted into the heat. (2) Heat redistribution parameter $P_{r}=<0,1>$ which controls how much heat from a certain point is redistributed to other places and how much is re-radiated locally. $P_{r}=0$ means that all absorbed heat is re-radiated locally and nothing gets transported to other places. (3) Zonal temperature redistribution parameter $P_{a}=<0,1>$ is a measure of the effectiveness of the homogeneous heat redistribution over the surface versus the zonal distribution. $P_{a}=1$ means that the heat is homogeneously redistributed over the surface while $P_{a}=0$ means zonal redstribution and that the heat flows only along the parallels. Our model properly describes the shapes of the objects by means of the Roche potential and takes into account gravity and limb darkening. At the same time it takes into account the orbital revolution, rotation and proper Doppler shifts in the scattered and thermal radiation. It is demonstrated on HD189733b that the light-curves of transiting extrasolar planets are mainly sensitive to the heat redistribution parameter $P_{r}$ and not so sensitive to the zonal temperature redistribution parameter $P_{a}$ or to the limb darkening. However, light-curves of planets seen at very low inclinations are very sensitive also to the zonal temperature redistribution parameter $P_{a}$ as well as to the limb darkening. This planet has low Bond albedo and relatively intense heat redistribution. The effect of non-spherical shape on the light-curve can also be important. For highly distorted planets like WASP-12b this might cause a double humped curve with the amplitude of about 10\\% superposed on other types of variability. Observations of this planet at 0.9 micron cannot constrain the Bond albedo or heat redistribution well. The effect of the grey and non-grey albedo on the spectrum of the Wasp-19b was studied as well. It was demonstrated that non-blackbody model with fluxes given by non-irradiated atmosphere models \\citep{hb07,allard03} can fit the observations of \\citet{anderson10b} reasonably well. The planet has low Bond albedo and low heat redistribution. We calculated the exact Roche shapes of all currently known transiting exoplanets \\citep{schneider95} and found out that the departures from spherical symmetry may vary significantly. Departures of the order of 1\\% are common, and can exceed about 8\\% in most extreme cases like WASP-12b, WASP-19b, and WASP-33b. About 8\\% of transiting planets have departures more than 3\\% (all have semi-major axes smaller than 0.03 AU). Temperature redistribution over the surface of all these planets was also calculated. The mean temperatures of the planets also vary considerably from 300 K to 2600 K. The extreme cases are WASP-33b, WASP-12b, and WASP-18b with mean temperatures of about 2600, 2430, and 2330 K, respectively." }, "1003/1003.4271_arXiv.txt": { "abstract": "Previous observational studies of the infrared (IR)-radio relation out to high redshift employed any detectable star forming systems at a given redshift within the restricted area of cosmological survey fields. Consequently, the evolution inferred relies on a comparison between the average IR/radio properties of ({\\it i}\\,) very IR-luminous high-$z$ sources and ({\\it ii}\\,) more heterogeneous low(er)-$z$ samples that often lack the strongest IR emitters. In this report we consider populations of objects with comparable luminosities over the last 10\\,Gyr by taking advantage of deep IR (esp. {\\it Spitzer} 24\\,$\\mu$m) and VLA 1.4\\,GHz observations of the COSMOS field. Consistent with recent model predictions, both Ultra Luminous Infrared Galaxies (ULIRGs) and galaxies on the bright end of the evolving IR luminosity function do not display any change in their average IR/radio ratios out to $z$\\,$\\sim$\\,2 when corrected for bias. Uncorrected data suggested $\\sim$0.3 dex of positive evolution. ", "introduction": "The IR/radio properties of galaxies at successively higher redshift have been probed in the past decade using either statistical samples from cosmological survey fields \\citep[e.g.][hereafter: S10]{appleton04, frayer06, sargent10}, the stacking technique \\citep[e.g.][]{carilli08, ivison10} or dedicated samples of specific objects \\citep[e.g. sub-mm galaxies (SMGs);][]{kovacs06, hainline09, michalowski09}. Evolutionary studies, all based on samples poorly matched in terms of bolometric luminosity at low and high redshift, have provided conflicting results, concluding that the local IR-radio relation either does (e.g. \\citealp{garrett02, appleton04, ibar08, garn09}; S10) or does not \\citep[e.g.][]{seymour09, ivison10} hold out to high redshift. \\noindent Recently, predictions have been made for the redshift evolution of the IR/radio properties of star-forming galaxies having different luminosities and geometries \\citep[e.g. compact starbursts and normal star-forming disks;][]{murphy09b, lackithompson09}. The current generation of IR and radio observatories can directly detect the brightest of these systems over a significant fraction of Hubble time, provided that sufficiently large cosmological volumes are sampled. Here we make use of the VLA and {\\it Spitzer} coverage of the 2 deg$^2$ COSMOS field to construct (cf. \\S\\ref{sect:data}) a volume-limited sample of ULIRGs at $z<$\\,2 that allows a direct comparison of observations and theory. Our findings are presented in \\S\\ref{sect:results} and discussed in \\S\\,\\ref{sect:conclusions}. \\noindent We adopt the WMAP-5 cosmology \\citep[$\\Omega_m$\\,=\\,0.258, $\\Omega_{\\Lambda}$\\,+\\,$\\Omega_m$\\,=\\,1 and $H_0$\\,=\\,71.9 km\\,s$^{-1}$\\,Mpc$^{-1}$;][]{dunkley09}. ", "conclusions": "\\label{sect:conclusions} \\noindent We have presented the first investigation of the evolution of the IR-radio relation out to $z$\\,$\\sim$\\,2 for a statistically significant, volume-limited sample of IR-luminous galaxies. This advance became possible thanks to two factors: ({\\it i}\\,) the large area and deep mid-IR coverage of the COSMOS field, and ({\\it ii}\\,) the inclusion of flux limits in the analysis with appropriate statistical tools.\\\\ At redshifts $z$\\,$<$\\,2 the median TIR/radio ratio of ULIRGs remains unchanged if we compensate for biases. On the most basic level this implies that their magnetic fields, $B$, are sufficiently strong to ensure that cosmic ray electrons predominantly lose their energy through synchrotron radiation (rather than inverse Compton scattering off the CMB). Regarding the 0.3\\,dex increase of the uncorrected evolutionary signal as an upper limit implies that $B$\\,$\\gtrsim$\\,30\\,$\\mu$G \\citep[e.g.][Fig. 5]{murphy09b}, as expected for compact and strong starbursts. This conclusion applies to both SF systems and optically selected AGN hosts, consistent with the similar mean IR/radio ratios reported for these two classes of objects in S10. Our finding agrees with theoretical and numerical expectations that ULIRGs should follow the local IR-radio relation until at least $z\\sim$\\,2 \\citep{lackithompson09, murphy09b}. Moreover, it suggests that the lower IR/radio ratios frequently reported for high-$z$ SMGs (e.g. \\citealt{kovacs06, valiante07, murphy09a, michalowski09}, but see also \\citealt{hainline09}) are not typical of distant ULIRGs in general. \\noindent Our complete sample of `IR-bright' galaxies -- the population that resides on the evolving bright end of the TIR LF -- links high-$z$ ULIRGs to normal IR-galaxies (log($\\nicefrac{L_{\\rm TIR}}{L_{\\odot}})\\gtrsim$\\,10.5) in the local universe. The fact that the average IR/radio ratio of the latter is very similar to that of ULIRGs demonstrates that the similar IR/radio properties of existing SF samples at low and high redshift are not the fortuitous consequence of comparing objects in different luminosity ranges. While distant starbursts follow the same IR-radio relation as local sources, this has not yet been ascertained for galaxies with moderate SF rates ($\\lesssim$10\\,$M_{\\odot}$/yr) that cannot be directly detected with current radio and far-IR facilities. The recent stacking analyses of \\citet{seymour09} and \\citet{ivison10} measured a steady decline of average IR/radio ratios that begins at $z$\\,$<$\\,1 and continues out to $z$\\,$>$\\,2. Given that the average IR luminosities of their image stacks are comparable to those of our `IR-bright' sample these findings are at odds with our measurements and, as shown here, cannot be ascribed to a luminosity offset between low and high redshift sources. It is to be expected that the origin of the discrepancy -- possibly the different methodology, sample selection or SED evolution \\citep[e.g.][]{symeonidis09, seymour10} -- will soon be identified in upcoming EVLA and Herschel surveys by virtue of the increased sensitivity and/or wavelength coverage these observatories offer. The latter capability in particular will ensure more accurate measurements of radio spectral indices and a better sampling of dust emission well into the far-IR which is crucial to the determination of the dust temperatures in distant starbursts. These improvements should also lead to a decrease in the scatter of the IR-radio relation at high redshift, thereby reducing both the need for and the impact of bias corrections of the kind that were applied in this work." }, "1003/1003.1986_arXiv.txt": { "abstract": "{A 321.5 s modulation was discovered in 1999 in the X-ray light curve of HM~Cnc. In 2001 and 2002, optical photometric and spectroscopic observations revealed that HM~Cnc is a very blue object with no intrinsic absorptions but broad (FWHM$\\sim$1500 km s$^{-1}$) low equivalent width emission lines (EW$\\sim1\\div6$\\AA), which were first identified with the HeII Pickering series. The combination of X-ray and optical observations pictures HM~Cnc as a double degenerate binary hosting two white dwarfs, and possibly being the shortest orbital period binary discovered so far.} % {The present work is aimed at studying the orbital motion of the two components by following the variations of the shape, centroid and intensity of the emission lines through the orbit.} {In February 2007, we carried out the first phase resolved optical spectroscopic study with the VLT/FORS2 in the High Time Resolution (HIT) mode, yielding five phase bins in the 321 s modulation. } {Despite the low SNR, the data show that the intensity of the three most prominent emission lines, already detected in 2001, varies with the phase. These lines are detected at phases 0.2- 0.6 where the optical emission peaks, and marginally detected or not detected at all elsewhere. Moreover, the FWHM of the emission lines in the phase resolved spectra is smaller, by almost a factor 2, than that in the the phase-averaged 2001 spectrum.} {Our results are consistent with both the pulsed optical component and emission lines originating in the same region which we identify with the irradiated surface of the secondary. Moreover, we note that the EWs of the emission lines might be up to $-15 \\div -25$ \\AA, larger than thought before; these values are more similar to those detected in cataclysmic variables. All the findings further confirm that the 321s modulation observed in HM~Cnc is the orbital period of the system, the shortest known to date.} {} ", "introduction": "Double degenerate binaries (DDBs) are close binary star systems consisting of two white dwarfs (WD) or a WD and a He star. DDBs are expected to originate from main sequence stars which have experienced two common envelope phases, eventually exposing the CO or He cores of the original stars (for a review see Warner 1995). Among DDBs are AM CVns, systems in which either a low mass He~WD or a low mass non degenerate He burning star fills its Roche lobe and transfers matter to the companion (Nelemans et al. 2001a). AM CVn binaries are characterized by orbital periods in the 600-3000 s range and, in the case of non degenerate donors, they evolve from long to short orbital periods (till they reach the minimum orbital period of $\\sim$10 min), similarly to cataclysmic variables (CVs). They evolve toward longer orbital periods, otherwise (Nelemans et al. 2001a and reference therein). Gravitational radiation waves drive the mass transfer and double-peaked emission lines in the optical spectra of most AM CVns testifies to the presence of an accretion disk mediating the flow of matter from one star to the other. In 1999, Israel et al. (1999, hereafter I99) discovered an X-ray emitting source with a strong 321.5~s. Deep TNG and VLT observations were carried out to study the optical counterpart of HM~Cnc and its optical spectral features. HM~Cnc is a blue V=21.1 (B=20.7) star showing a $\\sim$15\\% pulsed fraction at the $\\sim$321.5 s X-ray period (Israel et al. 2002, hereafter I02; see also Ramsay et al. 2002a). The optical and X-ray pulsations displays a phase offset of $\\sim$0.20-0.25 (Barros et al. 2007) indicating that they originate from two different regions. The optical modulation also varied in shape with time, while its pulsed fraction was clearly wavelength dependent (being larger at longer wavelengths). The spectrum of HM~Cnc is characterized by a blue continuum with no intrinsic absorption lines (I02), but broad (FWHM$\\sim$1500 km s$^{-1}$), low equivalent width (EW $\\sim$1-6 \\AA) emission lines, which were first identified as lines from the HeII Pickering series. All these findings are fully compatible with HM~Cnc being an X-ray emitting, He-rich DDB, possibly a progenitor of the AM CVn systems, with the shortest orbital period ever recorded (I02). Due to the difference in EWs measured between the odd- and even-terms of the HeII Pickering series, it has been proposed that HM~Cnc may also be hydrogen-rich (Steiper et al. 2005 and references therein). Should this be confirmed, it would be diffcult to relate HM~Cnc with AM CVn systems which show only He emission/absorption lines in their spectra. Thanks to the 2001-2006 photometric monitoring, Israel et al. (2004) found a very accurate P and $\\dot{P}$ phase coherent solution, with $\\dot{P}= -$1.1$\\times$10$^{-3}$ s/yr implying that the orbit is decaying. The optical solution was sufficiently accurate that it could be extended backward in time so as to find a unique phase-connected solution encompassing ROSAT 1994 and 1995 observations. This yielded a 10-years baseline P-$\\dot{P}$ coherent solution giving P=321.53038(2)~s and $\\dot{P}=-$3.661(5)$\\times$10$^{-11}$~s s$^{-1}$ (for more details see Israel et al. 2004; see also Strohmayer 2005 and Barros et al. 2007). Yet, the nature of the X-ray emission detected from HM~Cnc and its twin source V407~Vul is still under debate. A number of models have been proposed (for a review see Cropper et al. 2004). Among these, is the DDB model with mass transfer proposed in two flavors: with a magnetic primary (polar-like, Cropper et al. 1998) and with a non-magnetic accretor (Algol-like, Marsh \\& Steeghs 2002, Ramsay et al. 2002b). In the latter model, the reason why a disk does not form is because the minimum distance of the gas stream from the center of mass of the system is smaller than the size of the accretor, resulting in the stream hitting the surface of the accretor. The unipolar inductor model, UIM, is a second accredited model and, according to it, the secondary star does not fill its Roche-lobe but, by moving across the primary star magnetic field, induces an electric current which heats up the B-field line footprints on the primary white dwarf. No mass transfer is expected within the UIM. A similar mechanism is likely responsible for the discovered Jupiter-Io interaction (Wu et al. 2002; Clarke et al. 1996; Dall'Osso et al. 2006, 2007). In the DDB scenario the X-ray emitting region is thought to be ``self-eclipsed'' by the primary object giving rise to the sharp and large amplitude X-ray modulation. In all models the source is expected to be one of the strongest sources of low frequency gravitational wave radiation, easily detectable in the future by the LISA mission (Nelemans et al. 2001b and references therein). Within the DDB scenario, the negative value of the period derivative suggests that matter transfer is not viable, therefore disfavoring the Polar-like and Algol-like scenarios, and the fact that the GW emission is driving the orbital period decay. Moreover, the $\\dot{P}$ value is slightly large for the gravitational wave emission mechanism as the only channel governing the period evolution and we cannot currently exclude that magnetic stresses are also present in the system (Dall'Osso et al. 2006, 2007). However, Deloye and Taam (2006) have revised AM~CVn formation and evolution, showing that fully degenerate donors and the accretion scenario can easily produce systems having the orbital periods of HM~Cnc and V407~Vul at early contact. In addition, they also revise the duration of the phase characterized by mass transfer and negative $\\dot{P}$, determining values of $\\tau_{\\dot{M}}$ which are comparable with those predicted by the UIM. An alternative way around the negative $\\dot{P}$ during accretion has been proposed by D'Antona et al. (2006), who suggest that HM~Cnc is an accreting DDB, which is characterized by p-p hydrogen shell burning in the donor, seen during a mass transfer phase. The nuclear reactions on the shell cause the secondary star to shrink in response to the mass loss. D'Antona et al. (2006) obtain a period derivative which is correct in both sign and magnitude. This scenario also accounts for the possible presence of H in the spectrum of HM~Cnc. The lack of (relatively large) optical polarization argues against the Polar-like scenario. We will not consider this model further in this paper (see Reinsch et al. 2004; Israel et al. 2004). Recently, Ramsay et al. (2007) reported the detection of transient radio emission from HM~Cnc, which can be accounted for only with maser emission in the framework of the UIM model. In 2006 we requested and obtained VLT-FORS2 observations in a new observing mode, aimed at carrying out the first ever phase-resolved spectroscopic study of the 5.4 min modulation of HM~Cnc. Our main goals included the disentangling of the different optical components (pulsed vs un-pulsed), the identification of the emission line forming region(s) and the detection of the emission lines Doppler modulation to provide the ultimate proof of HM~Cnc binary nature with orbital period OP=5.4~min. Here, we report on the results we obtained from this new spectroscopic campaign. In section 2 we describe the observation and the FORS2 HIT-S mode. In section 3 the obtained results, while in section~4 we report our conclusions. ", "conclusions": "The phase resolved spectra of HM~Cnc presented in this paper have shown that its time resolved study is possible and in particular it is feasible with FORS2 HIT-S mode offered at ESO-VLT, though care must be taken during the data reduction and the data analysis phase because of the low count level regime of the spectra and the deferred charges effect. Our phase resolved sequence has demonstrated the emission lines intensity varies across the period and that the emission lines in the phase-resolved spectra have smaller FWHM than those in the phase-averaged spectra. These are two strong observational constraints which however do not allow us -yet- to formulate any definitive conclusion about the location of the line forming region and on the nature of the period because of the low SNR and quality of our data. In particular, we cannot establish whether the emission lines arise from the X-ray emitting hot-spot or from the irradiated surface of the secondary star. A higher SNR series of spectra would allow us to 1) establish whether the ELR is visible at any phase or they gets really ``eclipsed''; 2) measure with smaller uncertainties the emission line centroids thus establishing whether they move across the period and the larger FWHM of the phase-averaged spectra is indeed an effect of Doppler broadening. Case~A and C, which locate the ELR on the primary hot-spot should reveal emission lines with no o small Doppler shift. While, on the contrary, our case~B configuration implies significant Doppler shift of the emission lines. All the discussed geometries are consistent both with unipolar inductor model and the direct impact accretor as those models do not prescribe any specific location of the ELR in their formulation. The face on-intermediate polar hypothesis, instead is not compatible with the variable emissions line intensity nor with their smaller FWHM in the phase-resolved spectra. Hence it does not fit any of the 3 cases discussed in section~4. In the double degenerate polar model, we expect that the ELR is on the primary and therefore it is consistent only with our configuration A or B. This model should be discarded in case large Doppler shift of the emission lines are measured. % Finally, our phase-resolved spectral analysis showed that the HeII odd-term series emission lines at 4199.8\\AA \\ and 4541.6\\AA \\ were not detected (at any phase). This is consistent with the EWs inferred from the 2001 phase-averaged spectrum and confirms that the odd-term series lines are fainter than the even-term series ones. The emission line centroids $\\lambda_c$ are not yet measured with sufficient accuracy to unambiguously address the issue of the presence of H in the spectrum of HM~Cnc. However, we note that recent theoretical (D'Antona et al. 2006) and observational (Reinsch et al. 2007) studies of HM~Cnc envisage the presence of H in DDB systems." }, "1003/1003.4047_arXiv.txt": { "abstract": "We present results from an on-going survey for the \\HI 21 cm line and the OH 18~cm lines in IR galaxies with the Arecibo 305~m Radio Telescope. The observations of 85 galaxies extracted from the 2 Jy IRAS-NVSS sample in the R.A.~(B1950) range 20$^{\\rm h}$--00$^{\\rm h}$ are reported in this paper. We detected the \\HI 21~cm line in 82 of these galaxies, with 18 being new detections, and the OH 18~cm lines in 7 galaxies, with 4 being new detections. In some cases, the \\HI spectra show the classic double-horned or single-peaked emission profiles. However, the majority exhibit distorted \\HI spectral features indicating that the galaxies are in interacting and/or merging systems. From these \\HI and OH observations, various properties of the sample are derived and reported. ", "introduction": "At bolometric luminosities above $10^{11} L_ {\\odot}$, infrared (IR) galaxies become the dominant population of extragalactic objects in the local universe ($z \\leq 0.3$). These galaxies are subdivided into three categories: luminous (LIRGs, $L_{\\rm IR} > 10^{11} L_ {\\odot}$), ultraluminous (ULIRGs, $L_{\\rm IR} >10^{12} L_ {\\odot}$), and hyperluminous (HyLIRGs, $L_{\\rm IR} > 10^{13} L_ {\\odot}$; \\citealp{SM96}). Even though these IR galaxies are relatively rare, comprising less than 6\\% of the total IR energy density in the local Universe \\citep{SN91}, some studies suggest that the majority of galaxies with $L_{\\rm B} > 10^{11} L_ {\\odot}$ go through a stage of intense IR emission \\citep{SOI87}. Most IR galaxies with $L_{\\rm IR} < 10^{11} L_{\\odot}$ are single, gas-rich spirals, and their IR emission can be accounted for by star formation. In the luminosity range $10^{11} < L_{\\rm IR} < 10^{12} L_ {\\odot}$, most of the galaxies are interacting/merging systems with enormous quantities of molecular gas ($\\sim 10^{10}~M_ {\\odot}$). At the lower end of this range, the bulk of the IR luminosity is due to warm dust grains heated by a nuclear starburst, while active galactic nuclei (AGN) become increasingly important at higher luminosities. Galaxies with $L_{\\rm IR} >10^{12} L_ {\\odot}$ are believed to be advanced mergers powered by a combination of starburst and AGN \\citep{SM96}. Previous \\HI observations of IR galaxies have revealed very broad absorption lines in ULIRGs, indicating rotation plus large amounts of turbulent gas \\citep{MIR82}. High angular resolution Very Large Array (VLA) and Very Long Baseline Interferometry (VLBI) observations show that these galaxies have the absorbing \\HI situated in the inner few hundred parsecs along the line of sight to the nuclear continuum sources \\citep{BGSM87,MOM03}. Several OH 18~cm absorption and megamaser (hereafter OHM) emission surveys of ULIRGs have also been published \\citep{B89,DG00,DG01,DG02}. \\citet{B89} concluded that the OHM emission usually occurs in galaxies with higher far-IR (FIR) luminosities and flatter 100--25~$\\mu$m spectra rather than in those with OH 18~cm absorption features. Here, we report results from an on-going spectroscopic survey with the Arecibo Radio Telescope\\footnote {The Arecibo Observatory is part of the National Astronomy and Ionosphere Center, which is operated by Cornell University under a cooperative agreement with the National Science Foundation.} targeting the \\HI 21~cm and the main and satellite OH 18~cm lines of 85 IR galaxies from the 2-Jy IRAS-NVSS sample \\citep{YRC01}. In this paper, we adopt $H_0$ = 71 km $\\rm s^{-1}$ $\\rm Mpc^{-1}$, $\\Omega_{M}$ = 0.27, and $\\Omega_{\\Lambda}$ = 0.73. ", "conclusions": "Here we report preliminary statistical analysis for the observed sample. More detailed analysis and conclusions will be presented after completion of the full survey with all 582 galaxies. Figure~5~({\\it top}) shows the well established radio--FIR correlation through a logarithmic plot of the 1.4~GHz continuum luminosities versus the FIR luminosities for the galaxies in our sample. The derived correlation coefficient is 88\\%. Figure~5~({\\it bottom}) shows a logarithmic plot of the 1.4~GHz continuum luminosities versus the total IR luminosities. The correlation coefficient here is 89\\%. This remarkably tight linear correlation between the total radio continuum emission and the IR (or FIR) luminosities is well known for ``normal'' galaxies where the main energy source is not due to a supermassive black hole \\citep{CON92}. The most obvious interpretation of this correlation is the presence of massive stars that provide both relativistic particles via subsequent supernova events, and heat the interstellar dust which radiates at IR (or FIR) wavelengths \\citep{HSR85,WK88,CON92}. Figure~6 shows logarithmic plots of the \\HI mass versus the FIR ({\\it top}) and IR ({\\it bottom}) luminosities for the observed sample. Both plots show extremely weak correlations with coefficients of 42\\% in each. Figure~7 shows a logarithmic plot of the \\HI mass versus the 1.4~GHz radio luminosity. The correlation coefficient here is 53\\%. These plots suggest that the total neutral gas content and star formation activity traced through the radio luminosities or the IR luminosities are only weakly correlated for this sample. This is consistent with the scenario that atomic gas has first to be converted into molecular gas to form stars, and that the molecular gas content itself correlates well with star formation \\citep{WB02}. In Table 7, we present the mean and median \\HI mass values of galaxies with \\HI 21 cm emission as a function of total-IR luminosity bins. The numbers reflect a general trend of higher \\HI mass values at higher IR luminosities, consistent with the weak correlation seen in Figure~6-{\\it bottom}. We utilize the values presented in this table in our notes on individual objects in \\S 6. In our observed sample, several galaxies show either \\HI absorption or both \\HI emission and absorption. Binning the sample in $L_{\\rm IR}$ (Table~8a) reveals that sources with higher IR luminosities have the greater likelihood of showing \\HI absorption. For instance, 38.5\\% of the sources with $L_{\\rm IR} \\geq 10^{11.50}~L_{\\odot}$ show \\HI absorption, while only 10.3\\% of sources with $10^{11.00}~L_{\\odot} \\leq L_{\\rm IR} \\leq 10^{11.49}~L_{\\odot}$, and 6.3\\% of sources with $10^{10.50}~L_{\\odot} \\leq L_{\\rm IR} \\leq 10^{10.99}~L_{\\odot}$ show \\HI absorption. No \\HI absorption is seen in sources with $L_{\\rm IR} \\leq 10^{10.49}~L_{\\odot}$ (see Table 8a). Thus, while we cannot say anything about traditionally defined ULIRGs ($L_{\\rm IR} \\geq 10^{12}~L_{\\odot}$) because of their rarity in our observed sample, it appears that galaxies with $L_{\\rm IR} \\geq 10^{11.50}~L_\\odot$ have a greater likelihood of showing \\HI absorption. We now explore whether this trend could arise due to a selection effect. While the flux density of \\HI emission is proportional to $L_{\\rm HI}/D^2$, where $D$ is the distance to the galaxy, the flux density dip of an absorption corresponding to a given optical depth is proportional to the background continuum flux density that is being absorbed. In Table~8b we present the calculated mean and median NVSS flux densities for the galaxies in our observed sample as a function of $L_{\\rm IR}$ bins, to investigate whether galaxies having $L_{\\rm IR}\\geq 10^{11.50}~L_{\\odot}$ preferentially have higher flux densities. Table~8b shows that the radio flux densities in our sample do not correlate with IR luminosity. However, we do find for those galaxies showing \\HI absorption in the two highest luminosity bins, that the mean flux densities are about double the mean values for all galaxies in those bins. As a second possible effect, the $L_{\\rm IR} \\geq 10^{11.50}~L_{\\odot}$ absorbers might have a higher covering factor than those of lower luminosity e.g., their continuum emission may mostly be in a compact nucleus vs.~more extended continuum emission for the others, with the absorption arising principally in the nuclear regions. However, whether this is the case or not, we do find a difference in the incidence of \\HI absorption between IR galaxies with $L_{\\rm IR} \\geq 10^{11.50}~L_{\\odot}$ and those with lower IR luminosities, i.e., the galaxies with higher IR luminosities have higher \\HI column densities on the lines of sight to the continuum sources, and/or their continuum emission is confined to more compact regions. In Table 9, we present the statistics of OH detections (absorption or emission) as a function of $L_{\\rm IR}$ bins. The total number of sources per $L_{\\rm IR}$ bin in this table excludes galaxies with OH main line spectra severely affected by RFI. Here we note that all 3 of the OH emitters have $L_{\\rm IR} \\geq 10^{11.50}~L_{\\odot}$, while the two detections for $10^{10.50}~L_{\\odot} \\leq L_{\\rm IR} \\leq 10^{10.99}~L_{\\odot}$ are both absorbers. However, despite the small numbers in this subgroup, it is strongly suggestive that OH-detected sources (emitters and absorbers) are found primarily in galaxies with $L_{\\rm IR} \\geq 10^{11.50}~L_{\\odot}$." }, "1003/1003.3451_arXiv.txt": { "abstract": "The NRAO VLA Sky Survey (NVSS) is the only dataset that allows an accurate determination of the auto-correlation function (ACF) on angular scales of several degrees for active galactic nuclei at $z \\simeq 1$. Surprisingly, the ACF is found to be positive on large scales while, in the framework of the standard hierarchical clustering scenario with Gaussian primordial perturbations it should be negative for a redshift-independent effective halo mass of order of that found for optically selected quasars. We show that a small primordial non-Gaussianity can add sufficient power on very large scales to account for the observed NVSS ACF. The best-fit value of the parameter $f_{\\rm NL}$, quantifying the amplitude of primordial non-Gaussianity of local type, is $f_{\\rm NL}=62 \\pm 27$ ($1\\,\\sigma$ error bar) and $254^\\circ$, a negative ACF, but, on the contrary, it is observed to be positive. Careful analyses of the NVSS sample \\citep[e.g.][]{BlakeWall} indicate that systematic offsets that may induce a spurious positive signal should be negligible for the sources with $S_{1.4\\rm GHz}>10\\,$mJy, selected for the present analysis. If so, the NVSS ACF may point at the presence of a small primordial non-Gaussianity that adds power to the largest scales. After marginalizing over all the other parameters, we find $25 < f_{\\rm NL} < 117$ at the $95\\%$ confidence level, compatible with bounds derived by other studies. The minimum halo mass turns out to be $M_{\\rm min}=10^{12.47\\pm 0.26}\\,h^{-1}\\,M_\\odot$ ($1\\,\\sigma$), remarkably close to the value found by \\citet{Croom} for optically selected QSOs. We have addressed the significance and the robustness of our findings by considering different bias models and by investigating the impact of gauge effects on large scales. Error bars were estimated by a jackknife re-sampling procedure, widely used in the literature \\citep{Scrantonetal03,XiaISW}. It is known to be robust and accurate for the diagonal elements of the covariance matrix but not as widely tested and calibrated for the off-diagonal ones, which cannot be neglected because neighboring ACF data points are highly correlated. From \\cite{Scrantonetal03}, we infer that jackknife can underestimate parameter errors by up to 30\\%. If the error bars were to be increased by this (maximal) amount, $f_{\\rm NL}$ would become compatible with zero at the $\\sim2\\,\\sigma$ confidence level. We conclude that our work should be seen as a ``proof of principle\", indicating that future surveys probing scales $\\sim 100$ Mpc at substantial redshifts can put stringent constraints on primordial non-Gaussianity \\citep[e.g.][]{carbone08,viel09}. {\\it Acknowledgments: } We thank the referee for constructive comments. We acknowledge the use of LAMBDA and the HEALPix package. Numerical analysis has been performed at the University of Cambridge High Performance Computing Service. Research is supported by the ASI Contract No. I/016/07/0 COFIS, the ASI/INAF Agreement I/072/09/0 for the Planck LFI, PRIN MIUR, MICCIN grant AYA2008-0353, FP7-IDEAS-Phys.LSS 240117, FP7-PEOPLE-2007-4-3-IRGn202182." }, "1003/1003.3445_arXiv.txt": { "abstract": "We report precise Doppler measurements of \\npllet\\ subgiants from Keck Observatory. All \\npllet\\ stars show variability in their radial velocities consistent with planet--mass companions in Keplerian orbits. The host stars have masses ranging from $1.1 \\leq M_\\star/M_\\odot \\leq 1.9$, radii $\\rstarF \\leq R_\\star/R_\\odot \\leq \\rstarE$, and metallicities $-0.21 \\leq$~\\feh~$ \\leq +0.26$. The planets are all more massive than Jupiter (\\msini~$ > 1$~\\mjup) and have semimajor axes $a > 1$~AU. We present millimagnitude photometry from the T3 0.4~m APT at Fairborn observatory for five of the targets. Our monitoring shows these stars to be photometrically stable, further strengthening the interpretation of the observed radial velocity variability. The orbital characteristics of the planets thus far discovered around former A-type stars are very different from the properties of planets around dwarf stars of spectral type F, G and K, and suggests that the formation and migration of planets is a sensitive function of stellar mass. Three of the planetary systems show evidence of long-term, linear trends indicative of additional distant companions. These trends, together with the high planet masses and increased occurrence rate, indicate that A-type stars are very promising targets for direct imaging surveys. ", "introduction": "The field of exoplanetary science recently reached a major milestone with the first direct-imaging detections of planetary systems around main sequence stars\\footnote{The planet candidate around the A-type star $\\beta$~Pic announced by \\citet{lagrange09} has not yet been confirmed by proper motion \\citep{fitz09,lagrange09b}. Similarly, there exists no proper motion follow-up of the faint object imaged around 1RXS J160929.1-210524 \\citep{lafreniere08}.}. \\citet{kalas08} detected a single planet-sized object with a semimajor axis $a \\approx 120$~AU, orbiting just inside of the dust belt around the nearby, young A4V star Fomalhaut. The young A5V dwarf star HR\\,8799 is orbited by a system of three substellar objects with semimajor axes $a = \\{24,38,68\\}$~AU \\citep{marois08}. These remarkable systems share a number of characteristics in common. Both host-stars are A-type dwarfs with stellar masses $> 1.5$~\\msun\\ surrounded by debris disks, the planets are super-Jupiters with masses $\\lesssim 3$~\\mjup, and the companions orbit far from their central stars with unexpectedly large semimajor axes (ranging from $20-120$~AU). That both systems were discovered orbiting A stars might at first seem unlikely, given that A stars make up less than 3\\% of the stellar population in the Solar neighborhood and because the star-planet contrast ratios are unfavorable compared to systems with fainter, less massive central stars. However, in light of recent discoveries from Doppler-based planet searches of massive stars it is becoming apparent that A dwarfs may in fact be ideal target stars for direct imaging surveys \\citep{hatzes03, setiawan05, reffert06, sato07, nied07, liu08, dollinger09}. Measurements of the frequency of giant planets around the ``retired'' counterparts of A-type dwarfs (subgiants and giants) have found that the occurrence of Jovian planets scales with stellar mass: A-type stars ($M_\\star \\gtrsim 1.5$~\\msun) are at least 5 times more likely than M dwarfs to harbor a giant planet \\citep{johnson07b,bowler10,johnson10a}. And just like the current sample of imaged planets, Doppler-detected planets around retired A stars are more massive \\citep[][]{lovis07} and orbit farther from their stars than do planets found around Sun-like, F, G and K (FGK) dwarfs \\citep{johnson07, sato08b}. Indeed, there is strong evidence that the orbital characteristics of planets around A stars are drawn from a statistical parent population that is distinct from those of planets around FGK dwarfs. \\citet{bowler10} performed a statistical analysis of planets detected in the Lick Subgiants Survey, which comprises 31 massive stars ($M_\\star \\gtrsim 1.5$~\\msun) monitored for the past 5 years . The mass-period distribution of exoplanets around FGK dwarfs is typically described by a double-power-law relationship, with the frequency of planets rising toward lower masses and remaining flat in logarithmic semimajor-axis bins from $\\sim 0.05$~AU to $\\sim5$~AU \\citep{tab02,lineweaver03, cumming08, johnson09rev}. Based on the 7 planet detections from the Lick survey, Bowler et al. concluded that the power-law indices of the distribution of planets around A stars and Sun-like stars differ at the 4-$\\sigma$ level; the planets in their sample all have \\msini~$ > 1.5$~\\mjup\\ and none orbit within 1~AU. However, their small sample size precluded a determination of the exact shape of the mass-period distribution. Fortunately, given the $26^{+9}_{-8}$\\% occurrence rate measured from the Lick survey, it will not take long to build a statistical ensemble comparable to the collection of planets around less massive stars. To increase the collection of planets detected around massive stars, and to study the relationships among the characteristics of stars and the properties of their planets, we are conducting a survey of massive subgiants at Keck and Lick Observatories. The decreased rotation rates and cooler surface temperatures of these evolved stars make them much more ideal Doppler-survey targets compared to their massive main-sequence progenitors \\citep{galland05}. The observed effects of stellar mass on the properties of planets have important implications for planet formation modeling \\citep{ida05a, kennedy08, kretke09, currie09, dr09}; the interpretation of observed structural features in the disks around massive stars \\citep{wyatt99,quillen06,brittain09}; and the planning of current and future planet search efforts, such as the Gemini Planet Imager \\citep[GPI;][]{gpi}, Spectro-Polarimetric High-contrast Exoplanets REsearch \\citep[SPHERE;][]{sphere}, the Near Infrared Coronographic Imager \\citep[NICI;][]{nici}, and {\\it Project 1640} \\citep{p1640}. Our Lick survey has resulted in the discovery of 7 new Jovian planets orbiting evolved stars more massive than the Sun \\citep{johnson06b,johnson07,johnson08a,peek09,bowler10}. In this contribution, we present the first \\npllet\\ planets discovered in the expanded Keck survey. ", "conclusions": "We report the detection of \\npllet\\ new Jovian planets orbiting evolved stars. These detections come from the sample of subgiants that we are monitoring at Lick and Keck Observatories. The host-stars have masses in the range $\\mstarG$~\\msun\\ to 1.9~\\msun, radii $\\rstarF \\leq R_\\star/R_\\odot \\leq \\rstarE$, and metallicities $-0.21 \\leq$~\\feh~$ \\leq +0.26$. Five of the host-stars have masses $M_\\star > 1.5$~\\msun, and are therefore the evolved counterparts of the A-type stars. We also derived a jitter estimate for our sample of evolved stars and find that subgiants are typically stable to within 5~\\ms. The observed jitter of subgiants makes them uniquely stable Doppler targets among massive, evolved stars \\citep{fischer03,hekker06}. \\citet{bowler10} found that the minimum masses and semimajor axes of planets around A stars are very different from those of planets around FGK stars. Their findings suggest that the formation and migration mechanisms of planets changes dramatically with increasing stellar mass. The planets reported in this work strengthen that conclusion. The five new planets we have discovered around stars with $M_\\star > 1.5$~\\msun\\ all orbit beyond 1~AU and have minimum masses \\msini~$> 1$~\\mjup. These properties contrast with those of planets orbiting less massive stars, which have a nearly flat distribution in $\\log{a}$ from $a = 0.05$~AU to $a = 1$~AU \\citep{cumming08}, and a steeply rising mass function with $d\\ln{N}/d\\ln{M_P} = -1.4$. Successful theories of the origin and orbital evolution of giant planets will need to account for the discontinuity between the distributions of orbital parameters for planets around Sun-like and A-type stars \\citep{kennedy08,currie09,kretke09}. The abundance of super-Jupiters (\\msini~$ > 1$~\\mjup) detected around massive stars bodes well for future direct-imaging surveys. In addition to harboring the massive planets that are predicted to be the most easily detectable in high-contrast images, A-type dwarfs have the added benefit of being naturally young. A 2~\\msun\\ star has a main-sequence lifetime of only $\\sim1$~Gyr, which means that Jovian planets in wide orbits will be young and thermally bright. Three of the planets in Table~\\ref{tab:planets} show linear velocity trends indicative of additional long-period companions. These linear trends provide clear markers of massive objects in wide orbits around nearby stars, and therefore warrant additional scrutiny from RV monitoring and high-contrast imaging." }, "1003/1003.4265_arXiv.txt": { "abstract": " ", "introduction": "Inflation \\cite{classics} can solve both the horizon and flatness problems of cosmology in an elegant and minimal way (for a recent pedagogical review, see \\cite{Baumann09}). Inflationary theories can also naturally explain the primordial density fluctuations that eventually collapse to give rise to the large-scale structure we see today. However, all known classes of inflationary models are potentially sensitive to Planck-suppressed corrections to the inflaton Lagrangian, which can yield slow-roll parameters of ${\\cal O}(1)$, stopping inflation. While this sensitivity to high-scale physics is true of all inflationary theories, a particularly stark example of UV sensitivity arises in so-called ``large-field models,\" where the inflaton enjoys a super-Planckian excursion in field space during inflation. Such models are of special interest because it is only in such models that one may obtain gravitational wave signatures that are observable in the forseeable future \\cite{Lyth} (for a thorough discussion, see \\cite{CMBPol}). But gaining control of the inflaton Lagrangian over this large range of field space clearly demands detailed knowledge of the structure of an infinite series of potential Planck-suppressed terms. One possible UV completion of particle physics and gravity is string theory. In recent years, as our understanding of the details of string compactification has grown, it has become a realistic possibility to enumerate precise corrections to candidate inflaton Lagrangians in various scenarios. Recent results in this direction include those of \\cite{Liam}, where possible quantum gravity corrections to D3-brane inflation models in warped throat geometries are determined, and those of \\cite{monodromy}, where a shift symmetry protects large-field inflation in theories with high-scale supersymmetry breaking.\\footnote{See \\cite{axions} for other papers attempting to use shift symmetries to justify large-field inflation in string theory, and \\cite{reviews} for more general reviews.} The UV sensitivity of all inflation models, and the especially stark sensitivity of models which predict observable gravitational waves, provides the principal motivation for trying to embed inflation in string theory. A second major paradigm in theoretical physics is supersymmetry (see e.g. \\cite{susyreview}). Many theorists believe that supersymmetry is the leading candidate to stabilize the Higgs mass and explain the physics of electroweak symmetry breaking. String compactifications very naturally give rise to models with low-energy supersymmetry (though this is by no means known to be a prediction of the framework); the low-energy theory is then a 4D ${\\cal N}=1$ supergravity. Therefore, in evaluating statements about the space of inflationary models in string theory, and in particular statements that correlate inflationary observables with particle physics observables directly tied to supersymmetry breaking, it is useful to first consider what is and is not possible in the context of string-inspired low-energy supergravities. Recently, a striking claim about the relation between the two most basic observables in inflation and in theories of supersymmetry and its breaking has been put forward \\cite{Kallosh04,CMB}\\footnote{See also~\\cite{quibble} which noted tensions cropping up between low-scale SUSY breaking and high-scale inflation.}. The most fundamental observable in inflation is the scale of inflation \\begin{equation} V = 3 M_{\\rm P}^2 H^2 \\end{equation} (or equivalently, the Hubble constant during inflation, $H$). It directly controls the tensor amplitude~\\cite{Lyth}, and is a major factor in setting the scale of density perturbations. The primary observable in any realistic supergravity model is the scale of supersymmetry breaking, which is captured by the gravitino mass $m_{3/2}$. Quantitatively, one has: \\begin{equation}\\label{gravitinomass} m_{3/2}^2 \\approx {e^{K}|W|^2 \\over M_{\\rm P}^4} \\end{equation} where $W$ is the expectation value of the superpotential and $K$ is that of the K\\\"ahler potential (and the two appear above in a combination which is invariant under K\\\"ahler transformations, as expected). The authors of \\cite{Kallosh04} study possibilities for inflation in one of the simplest toy models of moduli stabilization and supersymmetry breaking known in string theory \\cite{KKLT}. They claim that within this class of models, very simple arguments (which we shall review in section~\\ref{KL}) lead to the conclusion that one must have \\begin{equation} \\label{quack} H \\leq m_{3/2}~. \\end{equation} The basic extra microscopic requirement that leads to this constraint is that of volume modulus stabilization (as we shall explain in detail in section~\\ref{KL}). This is a ${\\it new}$ microscopic requirement that must be considered in inflationary models that arise in an extra-dimensional setting, like that of string theory; it is ${\\it a ~priori}$ only indirectly related to traditional questions of 4D inflaton dynamics, like the flatness of a candidate inflaton potential. We note that typical supersymmetric models of particle physics have $m_{3/2} \\leq {\\rm TeV}$ (sometimes far lower, coming all the way down to $10^{-2}\\ {\\rm eV}$ in models of low-scale gauge mediation). In contrast, typical models of inflation have a characteristic energy scale $V$ during inflation that often approaches the GUT scale. All models with observable gravitational waves predict $H \\geq 10^{14}\\ {\\rm GeV}$, and very few models of any sort have been proposed with $H$ smaller than the values of $m_{3/2}$ typical in low-scale gauge mediation (recall that one must do baryogenesis etc. sometime ${\\it after}$ inflation). Therefore, the constraint eq.~(\\ref{quack}) is rather unwelcome. It has been further argued that while one can (clearly) find more general low-energy Lagrangians generalizing that of \\cite{KKLT}, that allow one to circumvent eq.~(\\ref{quack}), rather significant fine-tuning in the moduli-stabilizing sector is required to obtain models that robustly allow $H \\gg m_{3/2}$. In this paper, we examine the conclusion of~\\cite{Kallosh04} by studying possibilities for large-field inflation in low-energy theories that incorporate the same model of moduli stabilization, but vary the nature of the inflationary sector. We do not work in the full framework of string theory, but we do incorporate all of the features of low-energy string models that led to the tension in \\cite{Kallosh04}. We find that a wide class of large-field models can nevertheless arise in this framework, with $m_{3/2} \\ll H$, and without significant fine-tuning of the moduli-stabilizing sector.\\footnote{For other work focused on related issues, see e.g. \\cite{Postma,BaOlkillKL,Covi,KLkillKL,Burgess,Shiu}, and for new ideas about using the universal supergravity Goldstino multiplet for inflation, see \\cite{Luis}.} We emphasize that the problem described in \\cite{Kallosh04}, and solved there only by significant fine-tuning in the moduli-stabilizing sector, is ${\\it different}$ from the problem of obtaining a stringy inflation sector with a flat inflaton potential; it comes instead largely from a constraint to avoid decompactification of the extra dimensions of string theory. It is this new problem that is the focus of our investigation. ", "conclusions": "In this paper, we wrote down a class of toy models of inflation in string-inspired supergravity that successfully achieve $ m_{3/2} \\ll H$. This shows that there is no general reason, even in simple models of moduli stabilization in string theory, that the gravitino mass should be tied to the scale of inflation. Hence, the problem discussed in \\cite{Kallosh04} is not generic within inflation models in string-inspired supergravity constructions. Rather, it is an artifact of studying very specific models. However, our work leaves several important open questions. Firstly, the models we presented are basically supergravity implementations of $\\varphi^{2n}$ chaotic inflation with various values of $n$. The values of $n$ which we require to accomodate a ${\\rm TeV}$ gravitino mass are ruled out by experiment (being clearly disfavored by their predictions for both $n_s$ and $r$). Values of $n$ which are consistent with cosmological data still yield a large hierarchy between $H$ and $m_{3/2}$, but it is not large enough to allow ${\\rm TeV}$ scale (or lower) gravitino mass. Therefore, finding models which have $m_{3/2} \\ll H$ and which are consistent with both precision cosmological data and low-energy supersymmetry remains an open problem. Secondly, our models are not derived in a top-down framework like string theory; the embedding of large-field inflation in string models with low-energy supersymmetry remains an unmet challenge. It would be very interesting to try and address these problems, either in the class of models similar to \\cite{KKLT}, or in alternatives based on e.g. \\cite{Conlon,eva,garyold,newone,newtwo,newthree,New,recentone,recenttwo}. Our basic idea is to write down natural theories which allow $\\langle W \\rangle$ to vary by many orders of magnitude during inflation, to free the Hubble scale of inflation from the final gravitino mass. This idea should allow many different implementations, and it seems quite plausible that some of them will yield models consistent with both cosmological observations and low-energy supersymmetry. \\bigskip \\centerline{\\bf{Acknowledgements}} \\bigskip We would like to thank R. Kallosh and A. Linde for bringing this problem to our attention, and G. Shiu, S. Trivedi and P. Yi for enjoyable informal discussions in the autumn quarter of 2008. We are indebted to L. Covi and L. McAllister for many insightful comments on a draft of this paper. This research was initiated at the Aspen Center for Physics during the workshop ``Fingerprints of the Early Universe\" in the summer of 2009, and we are grateful to the Center and the workshop participants for providing a suitably stimulating atmosphere. S.K. is also grateful to the Mitchell Institute for Fundamental Physics and Astronomy at Texas A \\& M for hospitality (and for hosting the stimulating Strings 2010 conference!) during the completion of this work. T.H. was supported in part by a Summer Research Fellowship (for undergraduate research) from the Stanford physics department. The research of S.K. is supported in part by NSF grant PHY-05-51164. The research of A.W. is supported by NSF grant PHY-0244728." }, "1003/1003.1730_arXiv.txt": { "abstract": "A stalled spherical accretion shock, such as that arising in core-collapse supernovae, is unstable to non-spherical perturbations. In three dimensions, this Standing Accretion Shock Instability (SASI) can develop spiral modes that spin-up the protoneutron star. Here we study these non-axisymmetric modes by combining linear stability analysis and three-dimensional, time-dependent hydrodynamic simulations with Zeus-MP, focusing on characterizing their spatial structure and angular momentum content. We do not impose any rotation on the background accretion flow, and use simplified microphysics with no neutrino heating or nuclear dissociation. Spiral modes are examined in isolation by choosing flow parameters such that only the fundamental mode is unstable for a given polar index $\\ell$, leading to good agreement with linear theory. We find that any superposition of sloshing modes with non-zero relative phases survives in the nonlinear regime and leads to angular momentum redistribution. It follows that the range of perturbations required to obtain spin-up is broader than that needed to obtain the limiting case of a phase shift of $\\pi/2$. The bulk of the angular momentum redistribution occurs during a phase of exponential growth, and arises from internal torques that are second order in the perturbation amplitude. This redistribution gives rise to at least two counter rotating regions, with the maximum angular momentum of a given sign approaching a significant fraction of the mass accretion rate times the shock radius squared $(\\dot{M}\\,r_{\\rm shock}^2\\sim 10^{47}$~g~cm$^2$~s$^{-1}$, spin period $\\sim 60$~ms). Nonlinear mode coupling at saturation causes the angular momentum to fluctuate in all directions with much smaller amplitudes. ", "introduction": "The explosion mechanism of massive stars is not well understood at present. Observational and theoretical evidence gathered thus far suggests that this mechanism is intrinsically asymmetric for all stars that form iron cores ($\\gtrsim 12M_\\sun$; see, e.g., \\citealt{wang08}, \\citealt{burrows07d} and \\citealt{janka07} for recent reviews). Stars that end up with O-Ne-Mg cores are currently found to explode in spherical symmetry via the \\emph{neutrino mechanism} (\\citealt{kitaura06}, \\citealt{burrows07c}, \\citealt{janka08}). One piece of phenomenology that a successful theory of core-collapse supernovae has to explain is the distribution of pulsar spins at birth. Population synthesis studies of radio pulsars generally assume normal or log-normal distributions with mean values ranging from a few ms \\citep{arzoumanian02} to a few $\\sim 100$~ms \\citep{FG-K06} in order to reproduce observations. This large range is due to the different assumptions made about input physics, such as the shape of the radio beam (e.g., \\citealt{FG-K06}). Independent constraints that combine \\emph{Chandra}, \\emph{XMM}, and \\emph{Swift} observations of historic supernovae with an empirical correlation between X-ray luminosity and spin down power tend to rule out a large pulsar population with initial periods $\\leq 40$~ms \\citep{perna08}. On the other hand, current stellar evolution calculations that account for magnetic torques predict, using conservation of angular momentum and an assumption for the mass cut, neutron star birth periods $\\lesssim 10$~ms \\citep{heger05}. However, these models suffer from large uncertainties in the input physics, which could also lead in principle to very slowly rotating pulsars \\citep{spruit98}. Axisymmetric core-collapse calculations indicate that there is an almost linear mapping between the rotation rate of the iron core and that of the resulting protoneutron star, with no robust braking mechanism in sight to bring the implied fast neutron star spins to values more in agreement with observations \\citep{ott06}. If most presupernova cores turn out to rotate slowly, then there is an alternative mechanism to generate periods $\\gtrsim 50$~ms, which arises from instabilities in the supernova shock itself \\citep{blondin07a}. Axisymmetric core-collapse supernova simulations that include neutrino transport and microphysics to several degrees of approximation find that the stalled postbounce shock undergoes large scale oscillatory motions that break the spherical symmetry of the system \\citep{burrows95,janka96,mezzacappa98,scheck06,ohnishi06, buras06a,buras06b,burrows06a,burrows07, scheck08,murphy08,ott08,marek09,suwa09}. When neutrino driven convection is suppressed, the instability of the shock persists in the form of an overstable sloshing cycle, with the most unstable modes having Legendre indices $\\ell=1$ and $2$ \\citep{BM03,BM06}. In the limit of short wavelength, this so-called Standing, Spherical, or Stationary Accretion Shock Instability (SASI) is driven by an interplay between advected and acoustic perturbations \\citep{F07}. There is no conclusive proof yet on the mechanism behind long wavelength modes, although considerations of the timescales \\citep{FT09a} and the case of planar geometry \\citep{foglizzo09,sato09} point to the advective-acoustic cycle as a relevant component. The SASI is weakened when rotation becomes dynamically important (e.g., \\citealt{ott08}). Three-dimensional simulations of the SASI have found that a spiral type of mode can develop, causing the flow to divide itself into at least two counter-rotating streams, leading to angular momentum redistribution \\citep{blondin07a}. \\citet{blondin07a} find that this process can spin-up a canonical neutron star to a period of $\\sim 50$~ms after $\\sim 0.1M_\\sun$ of material is accreted. Aside from the fact that this spin period is comparable to that obtained by population synthesis calculations, the most interesting results of \\citet{blondin07a} are that (i) no progenitor rotation is required for this mechanism to operate and, (ii) for weak progenitor rotation, this instability dominates the spin-up of the protoneutron star, imparting it along a different axis. One of the key questions that remains to be answered is how easy it is to excite these spiral modes. In contrast to \\citet{blondin07a} and \\citet{blondin07b}, the findings of \\citet{iwakami08}, \\citet{iwakami09a}, and \\citet{iwakami09b} suggest that with no rotation in the upstream flow, these modes are difficult to excite. \\citet{yamasaki08} find that, indeed, for a cylindrical accretion flow with rotation, prograde spiral modes have a higher growth rate than retrograde ones. They attribute this to a Doppler shift of the mode frequency induced by rotation. Part of the difficulty in comparing numerical results on this instability is that they have been obtained using different microphysics, numerical methods, and coordinate systems, and that the three-dimensional structure of these linear eigenmodes remains as of yet unexplored. The work of \\citet{blondin07b} examined spiral SASI modes with time-dependent hydrodynamical simulations restricted to a polar wedge around the equatorial plane. They found that sloshing and spiral modes are closely related, because the pressure perturbation rotates with the same frequency as a sloshing mode of the same Legendre index, and because by combining two counter-rotating spirals, the flow field of a sloshing mode is recovered. The work of \\citet{iwakami09b} also identified sloshing modes as the superposition of two counter-rotating spirals. The aim of this paper is to better understand spiral modes in the linear and nonlinear phase, particularly their three-dimensional spatial structure and angular momentum content. The structure of axisymmetric sloshing eigenmodes has been obtained previously through linear stability analysis \\citep{F07}, and verified to high precision with time-dependent axisymmetric hydrodynamic simulations \\citep{FT09a}. Starting from the findings of \\citet{blondin07b}, we show that spiral modes are most easily understood as sloshing modes out of phase. Their spatial structure is built using solutions to the differential system of \\citet{F07}, and their evolution is compared with results of time-dependent hydrodynamic simulations using Zeus-MP \\citep{hayes06}. In our calculations we do not add rotation to the flow anywhere, and employ an ideal gas equation of state, parametric cooling, and point-mass gravity, with no heating or nuclear dissociation below the shock. One of our main findings is that it is relatively simple to create spiral modes out of sloshing modes by changing their numbers, relative phases, and amplitudes, covering a broader parameter space than the limiting cases where the amplitudes are equal and the relative phase is $\\pi/2$. These modes survive to large amplitudes, resulting in a significant angular momentum redistribution below the shock. We set aside for now the question of whether this coherent superposition of sloshing modes is able to develop in a more realistic core-collapse context, where three-dimensional, nonlinear turbulent convection may likely act as a forcing agent \\citep{FT09b}. We revisit this issue in the discussion section. We also find that the bulk of the angular momentum redistribution generated by a spiral mode occurs during the phase of exponential growth. This spin-up of the flow arises from internal torques that are second order in the perturbation amplitude, and results in angular momenta of a characteristic magnitude at saturation. Non-linear mode coupling causes the angular momentum flux to become stochastic, fluctuating in all directions with an amplitude much smaller than that achieved during exponential growth. The structure of this paper is the following. In Section 2 we describe the physical model used, and summarize the most important elements that enter the linear stability calculation and time-dependent simulations, with details deferred to the Appendices. Section 3 describes how spiral modes can be constructed by superposing sloshing modes out of phase, and explores some of their features. Section 4 contains the results of time-dependent simulations of these modes, including both linear and nonlinear development. We conclude by summarizing our findings and discussing their implications for more realistic core-collapse models and neutron star spins. ", "conclusions": "In this paper we have studied the spiral modes of the SASI in the linear and nonlinear regime, when no rotation is imposed on the accretion flow, and when neutrino driven convection is suppressed. We have combined results from linear stability analysis and time-dependent simulations to understand the structure and evolution of these non-axisymmetric modes. Our main findings follow: \\newline \\noindent 1. -- Spiral modes are most easily understood as two or more sloshing modes out of phase. In the linear regime and in the absence of rotation, their three-dimensional structure can be obtained by linear superposition of known axisymmetric eigenfunctions. The parameters that determine a spiral mode are the number of sloshing modes involved, their relative phases, and amplitudes. Two modes of equal amplitude and relative phase equal to $\\pm \\pi/2$ is a limiting case of a more general class of non-axisymmetric mode. \\newline \\noindent 2. -- As long as the initial relative phase is not zero, spiral modes survive in the nonlinear phase. For the $\\ell=1$, $m=\\pm 1$ case, they seem to reach some type of equilibrium relative phase once all modes are excited in the fully saturated stage. This equilibrium phase does not seem to depend on the initial phase shift. Hence, the range of perturbations needed to excite spiral-like behavior leading to protoneutron star spin-up is broader than that needed to achieve spiral modes with a relative phase of $\\pi/2$. The absolute sign of the angular momentum component in the inner and outer region depends on the sign of the initial relative phase. \\newline \\noindent 3. -- The angular momentum redistribution in the linear and weakly nonlinear phase of an isolated spiral mode is caused by the spatial dependence of the angular momentum density, which consists of at least two counterotating regions. This division occurs because the eigenmodes have at least one radial node in the transverse velocity profile (see, e.g., Figure~\\ref{f:angmomdens_eigenfunctions}b-c). The angular momentum density is non-oscillatory, and increases in magnitude at nearly twice the growth rate of the spiral mode. For $\\ell=1$, we have found that the dominant term involves a second order perturbation to the azimuthal velocity, which couples to the background density. We see no apparent relation between the formation of the triple point at the shock and the angular momentum redistribution, other than a possible role in the saturation of the primary mode and hence a cutoff in the growth of the spin-up. \\newline \\noindent 4. -- The bulk of the angular momentum redistribution is achieved at the end of the phase of exponential growth. After $200$ dynamical times, our models achieve a maximum spin-up of at least $0.6\\dot{M}\\rs0^2$ along the axis of the primary spiral mode, independent of the resolution employed, but restricted to $\\ell=1$. Modes with $\\ell=2$ also lead to spin-up along the primary axis, but with a smaller magnitude. \\newline \\noindent 5. -- In the fully saturated stage, where all modes are excited due to nonlinear coupling, all components of the angular momentum fluctuate with a characteristic magnitude $\\lesssim \\dot{M}\\rs0^2/10$. Resolution seems to be relevant for capturing the long-term evolution of the spin-up, which can display secular increases in magnitude at low resolution. \\newline \\citet{blondin07a} found that the development of spiral modes at late times is a robust feature of the flow, as many different types of perturbation resulted in a similar outcome when the accretion flow is non-rotating. Our results tend to confirm this picture. This was not the case in the simulations of \\citet{iwakami09a}, which however include neutrino heating and thus convection, fundamentally altering the flow dynamics relative to the simpler case with no heating. The fact that most progenitor models are likely to show some degree of rotation will always make excitation of a prograde spiral mode more likely, as its growth rate is larger \\citep{yamasaki08}. Indeed, the development of SASI-like spiral modes has been observed in time dependent studies of accretion disks around a Kerr black hole \\citep{nagakura09} Are overdense shells with large angular extent realistic? Two- and three-dimensional compressible simulations of carbon and oxygen shell burning in a $23M_\\sun$ star find that density fluctuations of order $10\\%$ are obtained due to internal waves excited in the stably stratified layers that lie in between convective shells \\citep{meakin06,meakin07a,meakin07b}. These large amplitude fluctuations are due to the strong stratification, and are correlated over large angular scales \\citep{meakin07a}. The oxygen-carbon interface lies too far out in radius to reach the stalled shock during the crucial few hundred milliseconds after bounce. But if a similar phenomenon were to occur above the Si shell, it would have a definite impact on the evolution of the stalled shock. Given that typical SASI periods are $\\sim 30$~ms, all that is needed is a region of radial extent $\\lesssim v_{\\rm ff}/\\omega_{\\rm osc}\\sim 200M_{1.3}^{1/2}r_{\\rm 150}^{-1/2}$~km by the time it reaches the shock. A different question is whether this coherent superposition of linear modes can take place in an environment where turbulent neutrino-driven convection already operates. The convective growth time in the gain region is a few ms (e.g., \\citealt{fryer07}) compared to the several tens of ms required for the SASI to grow. A critical parameter determining the interplay between these two instabilities is the integral of the buoyancy frequency multiplied by the advection time over the gain region \\citep{foglizzo06}. If this dimensionless parameter $\\chi$ is less than 3, then infinitesimal perturbations do not have time to grow before they are advected out of the heating region, and a large amplitude perturbation is required to trigger convection. Using axisymmetric simulations with approximate neutrino transport and a realistic equation of state, \\citet{scheck08} found that for weak neutrino heating, the SASI overturns can actually trigger convection. On the other hand, using parametric simulations, \\citet{FT09b} found that when $\\chi > 3$, convection grows rapidly and the vorticity distribution reaches its asymptotic value before the SASI achieves significant amplitudes. Two-dimensional convection is volume filling, as the vorticity accumulates on the largest spatial scales, hence large scale convective modes could be exciting dipolar SASI modes. The interplay between these two instabilities is not well understood at present, hence discussion of spiral modes in this context will have to wait for further work. A few three-dimensional core-collapse simulations have been performed so far, some of these with highly dissipative numerical methods \\citep{fryer07,iwakami08}. These groups find that convection starts on smaller angular scales, with convective cells having roughly the size of the gain region \\citep{fryer07}, or being mediated by high-entropy bubbles with a range of sizes \\citep{iwakami08}. In exploding simulations, as the shock expands, convective cells increase size and a global $\\ell=1$ mode emerges \\citep{fryer07,iwakami08}. Persistent spiral modes do not seem to appear spontaneously, they need to be explicitly triggered \\citep{iwakami08}. Very recently, \\citet{nordhaus10a} have reported three-dimensional, parametric core-collapse simulations using Riemann solvers and covering the whole sphere. Their results are in line with previous studies in that they do not witness the development of large scale shock oscillations in the stalled phase, or spiral modes with noticeable amplitudes. Similarly, \\citet{wongwathanarat2010} have presented full-sphere parametric explosions with a Riemann hydrodynamic solver, including the contraction of the protoneutron star and grey neutrino transport. Their conclusions are similar to that of \\citet{nordhaus10a} in that no coherent spiral modes are observed, with angular momenta saturating at a few times $10^{46}$~g~cm$^2$~s$^{-1}$. \\citet{fryer07} observe specific angular momenta $\\lesssim 10^{13}$~cm$^2$~s$^{-1}$ imparted stochastically to the protoneutron star by anisotropic accretion. These values, and those of \\citet{wongwathanarat2010}, are in agreement with the angular momentum fluctuations we observe in our fully saturated SASI, which on average have a magnitude $\\sim 5\\times 10^{12} (f_{\\rm amp}/0.1)\\dot{M}_{0.3}(\\rs0/150~{\\rm km})^2 M^{-1}_{1.3}$~cm$^2$~s$^{-1}$." }, "1003/1003.5113_arXiv.txt": { "abstract": "{From the recent work of the Evolution and Seismic Tools Activity (ESTA, Monteiro et al. 2006; Lebreton et al. 2008), whose Task 2 is devoted to compare pulsational frequencies computed using most of the pulsational codes available in the asteroseismic community, the dependence of the theoretical frequencies with non-physical choices is now quite well fixed. To ensure that the accuracy of the computed frequencies is of the same order of magnitude or better than the observational errors, some requirements in the equilibrium models and the numerical resolutions of the pulsational equations must be followed. In particular, we have verified the numerical accuracy obtained with the Saclay seismic model, which is used to study the solar g-mode region (60 to 140$\\mu$Hz). We have compared the results coming from the Aarhus adiabatic pulsation code (ADIPLS), with the frequencies computed with the Granada Code (GraCo) taking into account several possible choices. We have concluded that the present equilibrium models and the use of the Richardson extrapolation ensure an accuracy of the order of $0.01 \\mu Hz$ in the determination of the frequencies, which is quite enough for our purposes.} ", "introduction": "The interior of the Sun has been very well studied thanks to the information provided by pressure-driven modes ($p$ modes). In the case of the dynamics of the solar interior, due to the very small number of non-radial $p$ modes penetrating inside the core, neither the rotation profile (e.g., Chaplin et al. 1999; Thomson et al. 2003; Garc\\'ia et al. 2001, 2004, 2008c) nor the dynamical processes (e.g., Mathis \\& Zahn 2004, 2005) are well constrained inside this region. On the other hand, gravity ($g$) modes would give us complete access to the solar core, in particular, to its dynamics (e.g., Mathur et al. 2008; Mathur et al. 2010). Gravity modes have been searched for a long time, almost since the beginning of helioseismology (e.g. Hill et al. 1991; Pall\\'e 1991). But there is currently no undisputed detection of individual $g$ modes for the Sun (Appourchaux et al 2010). However, some peaks (e.g. Gabriel et al. 2002; Jim\\'enez \\& Garc\\'ia 2009) and groups of peaks (Turck-Chi\\`eze et al. 2004; Garc\\'ia et al. 2008a) have been considered as reliable $g$-mode candidates as they are above than 90$\\%$ confidence level and they present several of their expected properties. Moreover, to increase the probability of detection, Garc\\'ia et al. (2007, 2008b) searched for the global signature of such modes instead of looking for individual $g$ modes. They have found the signature of the asymptotic-dipole $g$ modes with more than 99.99$\\%$ confidence level. The detailed study of this asymptotic periodicity revealed a higher rotation rate in the core than in the rest of the radiative region and a better agreement with solar models (Garc\\'ia et al. 2008b) computed with old-surface abundances (Grevesse, Noels, \\& Sauval 1993) compared to the new ones (Asplund, Grevesse \\& Sauval 2005). However, it was not possible to identify the sequence of individual peaks generating the detected signal because of the very small signal-to-noise ratio. Thus, to go further it is necessary to use theoretical $g$-mode predictions to guide our search (Broomhall et al. 2007; Garc\\'\\i a 2010). For this purpose, we need to know the limits of the modeled physical processes and quantities as well as the internal numerical errors of the codes used to compute the predicted frequencies. The accuracy of the present solar models has already been studied by Mathur et al. (2007) and Zaatri et al. (2007). They showed that models with different physical inputs and fixed surface abundances present differences in the frequencies of the $g$ modes that are below 1 $\\mu$Hz in the range [60, 140] $\\mu$Hz. In the present work we study the numerical errors introduced by the approaches followed by the oscillation codes used to compute the $g$-mode frequencies of the Sun. It is a direct application of the study done in the ESTA group (whose Task 2 is devoted to the pulsational code's comparison, see Moya et al. 2008) to the solar case and the calculation of the $g$-mode frequencies. This comparison makes it possible to fix the global uncertainties of the numerical schemes used (Moya et al. 2010a). ", "conclusions": "The search for $g$ modes has been a long quest as they would be the best probes of the solar core, thus representing a huge potential to better constrain its structure and dynamics. Up to now, a few candidates have been detected and recently the global properties of dipole $g$ modes have been detected with more than 99\\% confidence level. The next step in the search for individual $g$ modes would consist in being guided by the theoretical predictions of their frequencies obtained with an oscillation code for a given solar model. This is the reason why the accuracy of the frequencies calculated with numerical algorithms is important. In this paper we have taken the advantage of the previous studies of the ESTA group and we have tested the accuracy of the equilibrium model used for the search of $g$ modes in the Sun under changes of methodology in the numerical integration of the pulsational equations. A model based on the Saclay-Seismic model has been used as an input to the pulsational codes ADIPLS and GraCo. Two comparisons have been studied: $i$) an overview of the differences obtained along the complete frequency spectrum and, $ii$) a especial analysis of the $g$-mode region $[60,140]~\\mu$Hz. This first comparison has shown that among the different methodologies in several zones of the spectrum, the present equilibrium model provides differences of the order of 0.1~$\\mu$Hz. This also happens when the same choices are used for both codes. This means that we need a larger number of mesh points if we want to accurately fit the observed frequencies in these regions. On the other hand, the $g$-mode region presents an accuracy of the order of $\\pm~0.02~\\mu$Hz for any methodology choice when the Richardson extrapolation is used which is much better than the uncertainties given by the physical prescriptions used in the models. This study has shown that, if we want to pursue a search related to observed $g$ modes in this region and based on theoretical models, the numerical accuracy of the Saclay-Seismic model and the use of either ADIPLS or GraCo codes, in terms of number of mesh points and numerical accuracy of sensitive quantities, are enough. Thus, such a guided search will be mainly sensitive to uncertainties coming from the physical inputs of the models." }, "1003/1003.0670_arXiv.txt": { "abstract": "In this paper, we examine in detail the key structural properties of high redshift dark matter haloes as a function of their spin parameter. We perform and analyze high resolution cosmological simulations of the formation of structure in a LCDM Universe. We study the mass function, shapes, density profiles, and rotation curves for a large sample of dark matter haloes from $z = 15 - 6$. We also present detailed convergence tests for individual haloes. We find that high spin haloes have stronger clustering strengths (up to $25\\%$) at all mass and redshift ranges at these early epochs. High redshift spherical haloes are also up to $50\\%$ more clustered than extremely aspherical haloes. High spin haloes at these redshifts are also preferentially found in high density environments, and have more neighbors than their low spin counterparts. We report a systematic offset in the peak of the circular velocity curves for high and low spin haloes of the same mass. Therefore, estimating halo masses without knowledge of the spin, using only the circular velocity can yield errors of up to $40\\%$. The significant dependence of key structural properties on spin that we report here likely has important implications for studies of star formation and feedback from these galaxies. ", "introduction": "\\label{sec:intro} The currently favored model that describes the formation of structure in the Universe is the $\\Lambda$ cold dark matter (LCDM) paradigm. In this model, the initial density distribution of the Universe was nearly homogenous, with small Gaussian density perturbations imprinted during an inflationary epoch. These fluctuations expand linearly, until the over-dense regions undergo non-linear gravitational collapse to form bound dark matter haloes. These haloes form in a hierarchical fashion: small haloes form first, and then larger ones assemble later via merging. In the LCDM paradigm, baryons follow the dark matter. Since they can dissipate and cool, baryons condense, and eventually form observable galaxies in the centres of dark matter haloes. The properties of dark matter haloes in the context of the LCDM paradigm have been studied in detail using numerical simulations over the past couple of decades with increasing resolution \\citep[e.g.][]{Davis85, Frenk88, Efstathiou88, Katz99, Kauffmann99, Bullock01B, Frenk02, Springel05}. This approach has been very fruitful in providing us with a detailed picture of the assembly and growth of structure in the Universe. These theoretical studies provide the framework within which the role of baryons and details of galaxy formation can be probed. While collisionless dark matter in the LCDM paradigm interacts only gravitationally, baryons dissipate, have pressure, cool, form stars, and interact with radiation. These, and other effects, introduce complications when trying to understand the properties of dark matter haloes such as their mass, angular momentum, shape, and density profiles from observations of the baryonic component. There are, however two techniques that have allowed a more direct probe of the dark matter: gravitational lensing observations \\citep[e.g.][]{Fischer00, McKay02, Hoekstra03, Mandelbaum06, Limousin07,Parker07,Evans09}, and measurements of galaxy rotation curves \\citep[e.g.][]{Rubin85, Trimble87, Persic96, Salucci07}. Due to the difficulties and assumptions required to translate the observed baryonic properties to dark matter halo properties, cosmological N-body simulations offer a powerful tool to understand the properties and statistics of the dark matter haloes. Even with dark matter only numerical simulations, much has been learned about the assembly of dark matter haloes, including the halo mass function, halo clustering, halo shape and spin at low redshift \\citep[see, e.g.,][]{vitvit02, reed03, reed07, Bett07, Gao07, reed08, Maccio08, Faltenbacher09}. However, there have been few detailed studies of dark matter halo properties at high redshifts. This is partly due to the number of particles required to resolve high redshift, low mass haloes, and still match observations of larger haloes at lower redshifts. These restrictions until recently prevented the detailed study of a statistically significant sample of collapsed haloes at high redshifts. As the observational frontier is pushed to higher and higher redshifts with reports of the detection of galaxies out to $z \\approx 7-8$ \\citep{Oesch10}, a deeper understanding of the properties of the dark matter haloes that host these most distant galaxies is critical as well as extremely timely. A few recent studies have examined specific dark matter halo properties at higher redshifts. \\citet{Heitmann06}, \\citet{Warren06}, and \\citet{reed07} focus on the mass function of high redshift haloes. \\citet{Moore06} trace the spatial distribution of dark matter halos from $z=12$ to the present day to understand their effect on galaxy mass haloes today. \\citet{jch01} use low resolution simulations to determine the spin and shape parameters of dark matter haloes at $z=10$. In a recent study \\citep{Davis09} we reported the results of the first high redshift and high resolution study to follow the growth of angular momentum in dark matter haloes in the mass range $10^6 \\Msun$ to $10^8\\Msun$ from $z=15$ to $z=6$, a period spanning 700 Myrs of cosmic time. We found that the spin distribution at these early epochs can be fit by a log-normal distribution as at lower redshifts. In addition, we examined the two-point correlation function of haloes and found that higher spin haloes are more clustered by factors up to $25\\%$ compared to their low spin counterparts at a given mass. This finding extended across all mass bins and redshifts in our previous study, i.e. from $10^6 - 10^8 \\Msun$ and from $z=15 - 6$. This paper builds on our earlier work by investigating the role angular momentum and the environment play in the determination of structural properties of dark matter haloes at these epochs. In the LCDM paradigm, haloes acquire angular momentum by tidal torques from their neighbors \\citep{hoyle49, peebles69, doro70, white84}. This picture for the acquisition and growth of angular momentum has been shown to be generally accurate in N-body simulations wherein angular momentum initially grows linearly with time \\citep{barnes87} and then slows down at later times \\citep{sugerman00}. Linear theory, however, overpredicts the angular momentum when compared to the fully non-linear N-body simulations \\citep{barnes87,sugerman00,porci02}. In addition, as \\citet{vitvit02} point out, linear theory predicts the angular momentum of a halo at a given redshift, but not the angular momentum of any particular progenitor at an earlier redshift. Thus, it becomes impossible with linear theory to trace the evolution of a halo's angular momentum in a hierarchical Universe evolving via mergers. \\citet{vitvit02, maller02, Hetz06} all note that mergers do affect the spin of the halo in addition to the tidal torque model. \\citet{donghia07} study mergers and spin evolution explicitly and argue that mergers only affect the spin of unrelaxed haloes, and find that relaxed, isolated haloes show no correlation between spin and merger history. One way to study the acquisition of angular momentum is to correlate information about the environment with halo properties. Previous studies have shown that halo clustering strength depends on the angular momentum of the halo at low redshifts \\citep{Bett07, Gao07, Faltenbacher09}. \\citet{Avila05} find that galaxy mass haloes inside clusters have smaller spins than haloes in the field or voids, and \\citet{Reed05} find low specific angular momentum in subhaloes near the central host halo. Observations using the Sloan Digital Sky Survey (SDSS) show that the spin parameter, \\begin{equation} \\lambda = \\frac{J|E_{\\rm{tot}}|^{1/2}}{GM^{5/2}}, \\label{eq:lambda} \\end{equation} where $J$ is the total angular momentum, $E_{\\rm{tot}}$ is the total energy, and $M$ the halo mass, has little dependence on the local environmental density \\citep{cs08}. It remains unclear, however, whether the spin parameter derived from the baryonic disk model used in interpreting SDSS data correlates well with the host dark matter halo's spin which is what is assumed. These results are found using galaxy neighbors to trace the large scale tidal field \\citep{cs09}. As an example of the difficulties in relating baryonic properties to the host dark matter halo, \\citet{Quadri03} report how slight spatial offsets between the dark matter and the baryonic disk create disturbed lensing configurations which can be easily misinterpreted if the misalignment is not included in the mass reconstruction. Also, \\citet{Mandelbaum06} report the misalignment of light ellipticity with halo ellipticity in the SDSS catalog. These findings illustrate the complexities when inferring dark matter halo properties from baryonic observations. Assembly bias refers to the observation that the clustering strength of dark matter haloes depends on an additional parameter beyond just halo mass. Assembly bias has been studied in simulations by examining halo formation time, concentration, shape, triaxiality, velocity structure, and substructure content \\citep{Harker06, Wechsler06, Bett07, Gao07, Jing07, Wetzel07, Angulo08}. These previous works show convincingly that halo clustering depends on more than just halo mass. However, all of these studies have been at low redshifts ($z < 5$) and for massive haloes ($M > 1 \\times 10^{10} \\Msun$). In this paper, we extend previous work by examining the formation and growth of dark matter haloes at high redshift, with an emphasis on studying the role angular momentum and environment play in regulating the structural properties of dark matter haloes. We limit ourselves to haloes in the mass range of $10^{6} \\Msun$ to a few times $10^9 \\Msun$, and the redshift range $z=15$ to $z=6$. This allows us to focus on the dark matter haloes that will likely host the first generation of stars and galaxies. We outline our paper as follows. In Section \\ref{sec:lowz} we summarize studies of angular momentum and assembly bias at low redshift to provide the framework for our findings at high redshift. We describe our simulations in Section \\ref{sec:sims}, and present the results of convergence tests in Section \\ref{sec:CT}. Our results from the correlation of the spin parameter to the halo environment are presented in Section \\ref{sec:environ}, and that of the effect of angular momentum on halo structure in Section \\ref{sec:structure}. We conclude with a discussion of the implications of our results for high redshift galaxy formation. ", "conclusions": "\\label{sec:disc} Our key findings can be summarized as follows: \\begin{itemize} \\item We have measured the spin, concentration, circular velocity, sphericity, and triaxiality parameter for a statistically large sample of dark matter haloes at high redshift ($z > 6$). \\item High spin haloes at high redshift are $25\\%$ more clustered than their low spin counterparts at a given mass, and are more likely to be found in high density environments. \\item High spin haloes (with masses $\\le 10^7\\Msun$) have smaller maximum circular velocities than low spin haloes, leading to errors up to $40\\%$ in the derived enclosed mass. \\item High spin haloes at high redshift are more likely to be aspherical and prolate, similar to findings at low redshift. \\item Nearly spherical haloes are up to $50\\%$ more clustered than extremely aspherical haloes, while there appears to be no difference in the clustering strength based on the triaxiality of the haloes. \\end{itemize} Our findings have an impact in two general areas: the role of angular momentum in halo structure and formation, and the role of assembly bias at high redshift. Our findings show that angular momentum has a measurable correlation with structural properties, including the concentration, sphericity, and triaxiality. Also, haloes with higher spin are preferentially found in higher density environments. The finding that halo spin correlates with local environment at high redshift is important to the understanding of the evolving properties of the baryonic component of dark matter haloes. A correlation between spin and baryonic properties, such as formation time, disk rotational speed, or disk size, would be the specific consequences of the correlation with the environment. These correlations are likely to be significantly stronger at high redshift, before too many mergers have happened which could destroy any correspondence between the dark matter spin and the baryonic structure. The role of assembly bias in halo evolution has been discussed at low redshifts. Our work extends the study of assembly bias to high redshifts, when the first galaxies form. We see similar results to those at low redshift when studying the dependence of clustering on spin and sphericity of dark matter haloes. However, our results differ when looking at the triaxiality of haloes. We find that there is no difference in clustering strength between prolate and oblate haloes. Our findings are of particular importance now that galaxies are being found at these high redshifts \\citep{Oesch10}. The clustering of galaxies has been used to infer the masses of their host dark matter haloes \\citep[e.g., ][]{Quadri08}. However, we find that properties other than mass - in particular halo spin - affect the measured correlation function. This additional parameter will induce errors in mass estimates for dark matter haloes inferred purely from clustering measurements. This fits in with the results of \\citet{Quadri08}, who suggest that mass is likely not the only parameter that drives the interaction of haloes with their large-scale environment. In addition to mass measurements, assembly bias will play an important role in feedback at these high redshifts. At the highest redshifts, simulations show that Population III stars have a large impact on their environment due to radiative and supernova feedback \\citep{Johnson07, grief07, Whalen08A, Whalen08B}. One consequence of our findings is that if angular momentum affects the formation and evolution of these Pop III stars, their feedback effects will show an environmental bias. Thus, the distribution of metals and reionization will be more clustered than otherwise expected. We also expect that because these stars are the first baryonic objects to collapse, their properties can be expected to have a stronger relationship to their host dark matter halo than galaxies today, which have undergone multiple mergers. We intend to pursue the consequences of our findings on baryonic results in future work. Our results suggest that the angular momentum properties of dark matter haloes likely have consequences for the properties of the first stars and galaxies hosted by them." }, "1003/1003.0994_arXiv.txt": { "abstract": "A faint new radio source has been detected in the nuclear region of the starburst galaxy M82 using MERLIN radio observations designed to monitor the flux density evolution of the recent bright supernova SN\\,2008iz. This new source was initially identified in observations made between 1-5th May 2009 but had not been present in observations made one week earlier, or in any previous observations of M82. In this paper we report the discovery of this new source and monitoring of its evolution over its first 9 months of existence. The true nature of this new source remains unclear, and we discuss whether this source may be an unusual and faint supernova, a supermassive blackhole associated with the nucleus of M82, or intriguingly the first detection of radio emission from an extragalactic microquasar. ", "introduction": "The nearby \\citep[d=3.6\\,Mpc;][]{freedman94} star-forming galaxy M82 has been subject to frequent radio monitoring at centimetric wavelengths with the VLA from the early 1980s \\citep{b1,b2}, and with MERLIN from the early 1990s \\citep{b3,b4}. Of order 60 compact radio sources have been identified within the central kpc of M82, the majority of which are thought to be recent supernova remnants which have exploded within the last 2000 years. The origin of 46 of these objects has been determined by the study of their radio spectral indices; 30 are considered to be supernova remnants, and 16 are thought to be compact H{\\sc ii} regions \\citep{b5}. Radio monitoring at intervals of around a year has shown that there is an additional population of radio transient sources whose origin is unknown. To date two transient sources have been detected, and each for only a single monitoring epoch implying that their lifetimes are typically less than a year. \\citet{b6} detected the compact radio source 41.5+597 in M82 with the VLA in February 1981. At that epoch, the object had a flux density of 7.1\\,mJy and 2.6\\,mJy at wavelengths of 6 and 2\\,cm respectively, implying that the source possessed a steep radio spectral index of $\\alpha$=$-$0.9 (S=$\\nu$$^{\\alpha}$). By October 1983 the source had faded to below the detection threshold of their VLA monitoring observations with an upper limit of 1.5\\,mJy at 6\\,cm. In a series of 6\\,cm MERLIN monitoring observations starting in the early 1990s, no emission was found at the position of 41.5+597 to limits of $\\sim$60$\\mu$Jy, and in the deepest 6\\,cm MERLIN observations of M82 to date, made in 2002 \\citep{b7} no emission was found to a limit of 20\\,$\\mu$Jy. In July 1992 \\citet{b3} detected a second transient, 40.59+55.8 with MERLIN at 6\\,cm with a flux density of 1.2\\,mJy. Subsequent MERLIN monitoring at 21\\,cm in April/May 1993 failed to detect emission at the position of the transient to a limit of 300\\,$\\mu$Jy. Furthermore it was not detected by deep MERLIN imaging at 6\\,cm in February 1999 and April 2002 with limits of 35 and 21\\,$\\mu$Jy respectively \\citep{b5,b7}. Both 41.5+597 and 40.59+558 lie outside the dynamical centre of M82 and since neither has given rise to a radio supernova remnant, they may be examples of stellar binary microquasar systems. If so, they would be the first to be discovered in the radio outside the Milky Way. Recently, \\citet{brunthaler09a,b8} reported the detection of radio emission from a new bright supernova in M82 (SN\\,2008iz) which is thought to have flared during the last week of March 2008 \\citep{b10}. The appearance of SN\\,2008iz around 45 years after the previous supernova \\citep[43.31+592,][]{b11} is consistent with the radio supernova rate for M82 of a new supernova approximately every 15 to 30 years \\citep{b3,b7}. Subsequent enhanced MERLIN monitoring of M82 has resulted in the detection of a new radio source in the central region of the galaxy \\citep{b9}. This faint new radio source was discovered in observations taken 1-5th May 2009 and was not present in images taken $\\sim$1 week earlier, on the 25th April 2009 (see Fig.\\,\\ref{Maps1}). Using closely spaced MERLIN (and VLBI) observations between April 2009 and January 2010, it has been possible, for the first time, to study the detailed evolution of one of the M82 transient source population. \\vspace*{-0.5cm} \\section[]{Observations and Data Reduction} The detection of both this new source and the campaign of continued flux density monitoring of the evolving SN\\,2008iz motivated a series of radio monitoring observations of M82. MERLIN observations of M82 were made between late April 2009 and January 2010 at 4994 and 6668.4\\,MHz, and 1658\\,MHz. All observations were made in wide-field mode, with parallel hands of circular polarisation, measured over 16\\,MHz of bandwidth correlated into 32 frequency channels. The primary flux density calibrator 3C\\,286 was used to set the flux density scale and the unresolved bright calibrator OQ208 was used to calibrate the amplitudes and bandpass responses. Throughout each epoch observations were interspersed with scans on the nearby phase reference source J095910+693217, with an assumed position of RA 09$\\hour\\,59\\min\\,10\\fsec6391$, Dec 69$\\degr\\,32\\arcmin\\,17\\farcs723$ (J2000). Data from each epoch were independently reduced using standard methods applying phase corrections determined form the phase reference source, J095910+693217, and the data were weighted appropriately to account for the relative sensitivities of the individual antennas. Following calibration, a large field encompassing the entire radio extent of M82 at this resolution was imaged, using multiple imaging facets and fully accounting for wide-field imaging effects. At each different reference frequency all epochs were imaged in an identical manner and the images were restored with a circular Gaussian beam appropriate for the {\\it uv} spacing of the baseline lengths and the weighting applied to the gridded data during imaging. \\begin{center} \\begin{figure*} \\includegraphics[width=15.5cm,angle=0]{muxlow_fig2.eps} \\vskip -0.15cm \\caption{The {\\it left-hand panel} shows the radio light curve of SN\\,2008iz and the new transient source reported over the first 150 days of its existence. MERLIN 4994 and 1658\\,MHz data are plotted as black and blue crosses respectively. MERLIN observations at 6666.8\\,MHz are not plotted. 1.4\\,GHz eVLBI data are shown as blue stars, 5 and 1.6\\,GHz VLBA are shown as black and blue open diamonds respectively \\citep{brunthaler09d}. The 5\\,GHz light curve for the SN\\,2008iz \\citep{b10} derived from single dish Urumqi observations is shown in green. The flux densities measured from these MERLIN data at 5\\,GHz for two nearby compact remnants 44.01+59.6 and 43.31+59.2 which are known to have a constant flux density \\citep{b6,ulvestad94} are also shown as pale blue triangles. An enlargement of the shaded region showing the light curve for the new radio transient is shown in the {\\it right-hand panel}.} \\label{lightcurve1} \\end{figure*} \\end{center} \\vspace*{-1.0cm} ", "conclusions": "This new source could be any of the above possibilities although each of the proposed scenarios has difficulty in explaining all of the observed properties. At present this source has been detected for $>$9 months and shows no immediate signs of fading. Depending upon its longevity it may represent another example of a relatively short-lived faint radio source population in M82. If so it would be the third example seen in $\\sim$30\\,years of observations; thus a lower limit on their occurrence rate is $\\sim$1 every 10\\,years, depending on their lifetimes. If this population is associated with faint supernovae it will have significant implications on the radio derived supernova rate of M82. Alternatively if this source is associated with microquasar flare event and the occurrence of which is related to the host galaxies star-formation rate, we would expect to see a comparable event in our own Galaxy around 1 every 100\\,years. Regular monitoring observations with new sensitive, high resolution imaging arrays, such as e-MERLIN and the EVLA, will be required to determine the size and nature of any such a population, both M82 and other star-forming galaxies. Global VLBI observations at 1.6 and 5\\,GHz with milliarcsecond resolution were taken in late 2009 and are awaiting correlation. Images of the transient from these new data will constrain the nature of this exciting source. \\vspace*{-0.7cm}" }, "1003/1003.2169_arXiv.txt": { "abstract": "\\SDSS\\ is a white dwarf in a close orbit with a companion that has been suggested to be a neutron star. If so, it hosts the closest known neutron star, and its existence implies a great abundance of similar systems and a rate of white-dwarf neutron-star mergers similar to that of the type Ia supernova rate. Here, we present high signal-to-noise spectra of \\SDSS, which confirm an independent finding that the system is in fact composed of two white dwarfs, one relatively cool and with low mass, and the other hotter and more massive. With this, the demographics and merger rate are no longer puzzling (various factors combine to lower the latter by more than two orders of magnitude). We show that the spectra are fit well with a combination of two hydrogen model atmospheres, as long as the lines of the higher-gravity component are broadened significantly relative to what is expected from just pressure broadening. Interpreting this additional broadening as due to rotation, the inferred spin period is short, about 1 minute. Similarly rapid rotation is only seen in accreting white dwarfs that are magnetic; empirically, it appears that in non-magnetized white dwarfs, accreted angular momentum is lost by nova explosions before it can be transferred to the white dwarf. This suggests that the massive white dwarf in \\SDSS\\ is magnetic as well, with $B\\simeq10^5\\,$G. Alternatively, the broadening seen in the spectral lines could be due to a stronger magnetic field, of $\\sim\\!10^6\\,$G. The two models can be distinguished by further observations. ", "introduction": "An unexpected return from the Sloan Digital Sky Survey (SDSS; \\citealt{yaa+00}) has been the discovery of white dwarf binaries with short periods, by \\citet{bmt+09} and \\citet{mbt+09}. The massively multiplexed SDSS spectroscopic observations consist of several 15-minute integrations, which are combined to yield the final spectra. \\citet{mbt+09} and \\citet{bmt+09} took advantage of this approach and looked for rapid changes in radial velocity across the individual spectra. The first finds are already quite interesting. SDSS\\,J143633.29+501026.8 (orbital period, $P_{\\rm b}=$1.15\\,hr) and SDSS\\,J105353.89+520031.0 ($P_{\\rm b}=0.96$\\,hr) are both low mass ($\\sim$0.3\\,M$_\\odot$) DA white dwarfs orbiting an unseen companion. For both systems, the measured mass function can be reasonably explained by invoking a white dwarf for the secondary.\\footnote{We refer to the detected DA white dwarf as the photometric primary (or ``primary'') and the unseen/fainter companion as the photometric secondary (``secondary'').} The orbital periods are so short that the systems are expected to merge within a Hubble time. These two binaries increased the toll of such interesting double degenerates to seven \\citep{mbt+09}. The number is likely to increase further; e.g., from a targeted search of low-mass white dwarfs found in SDSS, \\citet{kba+09} find four short-period binaries (including the two from \\citealt{mbt+09}). These authors also discuss in detail the fate of these systems: since the combined masses are likely below the Chandrasekhar mass, most will not explode as type Ia supernovae, but rather become AM CVn systems or R CrB stars. SDSS\\,J125733.63+542850.5 (\\SDSS\\ hereafter) consists of a DA white dwarf primary in an orbit with period $P_{\\rm b}=4.6$\\,hr and no measurable eccentricity. The inferred companion mass (as constrained by the spectroscopically inferred white-dwarf mass and the mass function) is above the Chandrasekhar mass (in a probabilistic sense). If so, the companion is either a neutron star or a black hole, and, at the estimated distance, $d\\approx 50$\\,pc, it would be the nearest ultra-compact object known; see \\citet{bmt+09}. The importance of the proximity of \\SDSS\\ becomes apparent when we consider the distances to the nearest members of different sub-classes of neutron stars (see Table~\\ref{tab:NearestNS}): the nearest millisecond binary pulsar (a binary consisting of a low-mass white dwarf and a pulsar and in a circular orbit), PSR\\,J0437$-$4715; the thermally X-ray emitting middle aged neutron star, RX~J1856.5$-$3754; the nearest long period pulsar, PSR\\,J2144$-$3933; the nearest ordinary pulsars, PSR\\,J0108$-$1431 and PSR\\,0950+08; the nearest young pulsar, Vela and the nearest $\\gamma$-ray pulsar, Geminga. (For the ordinary pulsars, accounting for the usual beaming factor of 0.1, results in a distance of about 130\\,pc.) The distance $d_{\\mathrm{NS}}\\sim 150\\,$pc for the nearest neutron star(s) is in accord with the global demographics of neutron stars as well as demographics of specific classes of neutron stars (based on birth rates). Specifically for \\SDSS, if its unseen companion is truly a neutron star then, given the short period, the system would arguably be a binary millisecond pulsar. The local surface density of pulsars and millisecond pulsars is about 30\\,kpc$^{-2}$ \\citep{lml+98}. For millisecond pulsars the local density is between 30--45\\,kpc$^{-3}$ \\citep{cc97}. The expected distance to the nearest millisecond pulsar is thus 175--200\\,pc, consistent with the observations but inconsistent with the proximity of \\SDSS. \\begin{deluxetable}{lll} \\tablecaption{The Nearest Neutron Stars.\\label{tab:NearestNS}} \\tablewidth{\\hsize} \\tablehead{ &\\colhead{Distance}&\\\\ \\colhead{Name}&\\colhead{(pc)}&\\colhead{Ref.}} \\startdata PSRJ0437$-$4715 & $156.3\\pm1.3$ & \\citealt{dvt+08}\\\\ RXJ1856.5$-$3754 & $161^{+18}_{-14}$ & \\citealt{vkk07}\\\\ PSRJ2144$-$3933 & $165^{+17}_{-14}$ & \\citealt{dtb+09}\\\\ PSRJ0108$-$1431 & $240^{+124}_{-61}$ & \\citealt{dtb+09}\\\\ PSR0950+08 & $260^{+58}_{-5}$ & \\citealt{bbg+02}\\\\ Vela Pulsar & $286^{+19}_{-17}$ & \\citealt{dlr+03}\\\\ Geminga & $250^{+120}_{-62}$ & \\citealt{fwa07} \\enddata \\end{deluxetable} Thus, a neutron-star companion to \\SDSS\\ would be very puzzling, unless one assumes that its proximity to the solar system is a statistical fluke. Motivated by this, as well as by the alarming implications of a paper by \\cite{tks09}, we first reconsider the case for a neutron-star companion. A reading of the \\cite{bmt+09} paper shows that it rests primarily on the high mass for the primary, with kinematics offered as a supporting argument. We revisit these two issues in the next two sections, and then present observations showing that the system in fact hosts two white dwarfs. Before proceeding, we note that while writing up our results, a preprint by \\citet{mgs+10} appeared, which presented lower-quality but much more numerous spectra, from which the presence of a second component is also evident. These authors discuss in detail the preceding and future evolution of this binary. In this paper, we focus on the demographics, and on obtaining more accurate constraints on the white dwarfs in the system from our higher-quality spectra. ", "conclusions": "" }, "1003/1003.0446_arXiv.txt": { "abstract": "We present analyses of the physical conditions in the $z(\\OVI) = 0.22496$ and $z(\\OVI) = 0.22638$ multi-phase absorption systems detected in the ultraviolet HST/{\\it STIS} and {\\it FUSE} spectra of the quasar H~$1821+643$ ($m_V = 14.2, z_{em} = 0.297$). Both absorbers are likely associated with the extended halo of a $\\sim 2L_B^*$ Sbc-Sc galaxy situated at a projected distance of $\\sim 116~h_{71}^{-1}$~kpc from the sight line. The $z = 0.22496$ absorber is detected in {\\CII}, {\\CIII}, {\\CIV}, {\\OIII}, {\\OVI}, {\\SiII}, {\\SiIII} and {\\HI} ({\\Lya} - Ly$\\theta$) at $> 3\\sigma$ significance. The components of {\\SiIII} and {\\SiII} are narrow with implied temperatures of $T \\lesssim 3 \\times 10^4$~K. The low and intermediate ions in this absorber are consistent with an origin in a $T \\sim 10^4$~K photoionized gas with [Si/H] and [C/H] of $\\sim -0.6$~dex. In contrast, the broader {\\OVI} absorption is likely produced in collisionally ionized plasma under nonequilibrium conditions. The $z(\\OVI) = 0.22638$ system has broad {\\Lya} (BLA) and {\\CIII} absorption offset by $v = -53$~{\\kms} from {\\OVI}. The {\\HI} and {\\CIII} line widths for the BLA imply $T = 1.1 \\times 10^5$~K. For non-equilibrium cooling we obtain [C/H] $\\sim -1.5$~dex and $N(\\H) = 3.2 \\times 10^{18}$~{\\cmsq} in the BLA. The {\\OVI}, offset from the BLA with no detected {\\HI} or {\\CIII}, is likely collisionally ionized at $T \\sim 3 \\times 10^5$~K. From the observed multiphase properties and the proximity to a luminous galaxy, we propose that the $z=0.22496$ absorber is an extragalactic analog of a highly ionized Galactic HVC, in which the {\\OVI} is produced in transition temperature plasma ($T \\sim 10^5$~K) at the interface layers between the {\\it warm} ($T < 5 \\times 10^4$~K) HVC gas phase and the {\\it hot} ($T \\gtrsim 10^6$~K) coronal halo of the galaxy. The $z=0.22638$ {\\OVI} - BLA absorber could be tracing a cooling condensing fragment in the nearby galaxy's {\\it hot} gaseous halo. ", "introduction": "Quasar absorption line studies as well as galaxy formation models have shown that the baryons in the extended halos of galaxies exist in the form of structures with different masses, spatial scales, densities and temperatures. In the Milky Way, the multiphase composition of the halo is evident from the distribution of numerous neutral and {\\it warm} \\footnote {Throughout this paper we use the terms {\\it warm} and {\\it hot} in a manner that is consistent with the traditional definition of these terms in ISM astronomy. Thus, {\\it warm} refers to gas with $T \\sim (0.3 - 3) \\times 10^4$~K, and {\\it hot} refers to gas with $T > 10^6$~K. For intermediate temperatures of $T \\sim (0.5 - 10) \\times 10^5$~K, we use the term {\\it transition temperature} to covey that gas in this unstable temperature regime is likely not in a state of equilibrium.} high velocity gas clouds (HVCs) pressure supported by the {\\it hot} and diffuse corona \\citep{sembach03, collins04, fox04}. The sources of high velocity gas surrounding the Milky Way include tidally stripped gas during interactions with satellite galaxies, accreting gas from the IGM, galactic scale outflows, and fragmentations from the cooling of a {\\it hot} halo \\citep[see reviews by ][]{wakker97,richter06a}. Detecting multiphase gas in the gaseous halos of external galaxies provides an opportunity to trace these varied processes at higher redshifts. The multiphase nature of the absorber can only be fully understood through a combination of line diagnostics from low, intermediate and high ions. Close resemblance of the physical properties with the high velocity gas in the Galactic halo can be a crucial pointer towards the specific nature of these higher redshift systems. Among the high ions, {\\OVIdblt}~{\\AA}\\footnote{We use the oscillator strengths and the rest-wavelengths of electronic transitions from the atomic data given in Verner {\\etal} (1994, 1996) for $\\lambda < 912$~{\\AA} and \\citet{morton00, morton03}. The wavelengths are given as vacuum wavelengths rounded to the nearest natural number.} lines have been of particular importance for tracing interstellar gas with $T \\sim (1 -3) \\times 10^5$~K, due to the high relative abundance of oxygen and the large oscillator strength of the doublet lines ($f_{1032} = 0.133, f_{1038} = 0.067$). Ultraviolet spectroscopic observations of sight lines towards extragalactic objects have detected {\\OVIdblt}~{\\AA} absorption associated with many of the HVCs in the halo of the Milky Way. The {\\it FUSE} survey of \\citet{sembach03} found $\\sim 60 - 85$\\% of the sky covered by these high velocity {\\OVI} absorbing clouds. Some of this {\\OVI} is drawn from the same population as the neutral HVCs with $N(\\HI) \\gtrsim 2 \\times 10^{18}$~{\\cmsq} and thus is also detected in 21-cm radio emission above this column density limit. This neutral population has a sky covering fraction of $\\sim 30$\\%. The more highly ionized HVCs are detected through their absorption in the optical and UV spectra of background sources. Their non-detection in 21-cm emission constrains the {\\HI} column density to $N(\\HI) < 2 \\times 10^{18}$~{\\cmsq} \\citep{sembach99, wakker03, collins04, ganguly05}. The {\\OVI} is an excellent tracer of collisionally ionized gas since $E > 114$~eV energies are required for its production from lower ionization stages. The observational constraints set by the ionic column density ratio of {\\OVI} with low/intermediate ions such as {\\SiII}, {\\CII}, {\\SiIII}, {\\CIII} and other high ions such as {\\CIV}, {\\SiIV} and {\\NV} indicate that the ionization process in the high velocity gas phase traced by {\\OVI} is dominated by collisions of electrons with ions in a plasma with $T \\sim 2 \\times 10^5$~K. The gas at this temperature is susceptible to strong radiative cooling, with the {\\OVI} ion also acting as a major cooling agent \\citep{sutherland93}. In Galactic HVCs, the transition temperature phase with $T \\sim (1 - 3)~\\times~10^5$~K forms from the interaction of {\\it warm} photoionized gas with the {\\it hot} coronal halo, and is described by a temperature intermediate to the {\\it warm} and {\\it hot} phases \\citep{sembach03, fox04, fox05}. Our current knowledge on the physical state of the highly ionized HVCs in the Galactic halo provides useful indicators for understanding the nature of {\\OVI} absorption line systems detected at higher redshift. The redshift number density $dN/dz \\sim 15$ for absorbers with $W_r(\\OVI~\\lambda 1032) > 30$~m{\\AA} at $z < 0.5$ determined from observations \\citep{tripp08, danforth08} and supported by cosmological hydrodynamic simulations \\citep{tumlinson05, cen06, oppenheimer09} indicates a high frequency of incidence for {\\OVI} absorbers at low-$z$. Similarities in physical conditions with Galactic HVCs would suggest that at least some fraction of the population of the {\\OVI} absorbers are tracing high velocity gas in external galaxies. Evidence is emerging from absorber-galaxy pair studies in quasar fields as well as more general correlation studies between absorbers and databases of galaxies from wide field surveys in favor of {\\OVI} absorption preferentially tracing gas in the immediate vicinity (impact parameter, $\\rho \\lesssim 500~h_{71}^{-1}$~kpc) of one or more galaxies \\citep{sembach04, tumlinson05, stocke06, tripp06, wakker09, chen09}. Whether the absorption in all those cases is arising from a distinct gaseous structure embedded within the galaxy's halo, from a more diffuse gaseous envelope surrounding the galaxy or from an intergalactic filament networking the individual galaxies is not always evident. In a subset of nearby {\\OVI} absorbers detected at high velocities with respect to the LSR, the ionization properties, galactocentric distances and location in the region of the sky surrounding the Milky Way are all consistent with an origin in the general Local Group environment rather than the Galactic halo \\citep{sembach03}. It is therefore possible that some of the {\\OVI} absorbers detected at higher redshifts could also be tracing clouds within a group environment rather than the gaseous halo of one of the galaxies in the group. A detailed understanding of the physical conditions and metallicity in each absorber might help to distinguish its origin and location. In this paper we describe the astrophysical nature of two multiphase absorption systems detected along the sight line to the bright quasar H~$1821+643$, that were previously reported in \\citet{savage98}, \\citet{tripp00} and \\citet{tripp08}. The two absorption systems are within $\\Delta v \\sim 350$~{\\kms} of each other. Ground based imaging observations of this quasar field have identified galaxies in the vicinity of these absorbers \\citep{schneider92, savage98, tripp98a}. The organization of this paper is as follows. In Sec. 2 we provide details on the observations. Secs. 3 and 4 describe the observed properties of the two absorption systems. Ground based imaging observations of the H~$1821+643$ field and information on the galaxy identified in the vicinity of the absorbers are given in Sec. 5. Detailed investigations of the various ionization scenarios in each absorption system is considered in Sec. 6, followed by explanations on the astrophysical nature of these absorbers (Sec. 7). The main results are summarized in the last section. ", "conclusions": "Our interpretation of the astrophysical nature of the two absorbers, and the environments they trace, are meaningful in the light of some recent results connecting {\\OVI} absorption systems to galaxies. From a survey of galaxies in quasar fields with known {\\OVI} absorbers, \\citet{chen09} found that approximately $90$\\% (10/11) of the absorbing galaxies in their sample are emission line dominated. Those galaxies that are absorption dominated do not have {\\OVI} detected down to an equivalent width limit of $W_r(\\OVI~1032) \\lesssim 0.03$~{\\AA} in the spectrum of the background quasar. Thus {\\OVI} seems to be preferentially selecting gas-rich galaxies. The absorbing galaxies were also found to be part of group environments with their morphologies pointing to past events of interactions or satellite accretion. Such interactions were expected to contribute to the $\\sim 64$\\% covering fraction of {\\OVI} estimated from their sample. The star formation rate which we derive from the strength of the H$\\alpha$ emission in the spectrum of galaxy G suggests on going star formation, consistent with a gas-rich system (see Sec 5). It extends the possibility of the $z = 0.22496$ and $z = 0.22638$ {\\OVI} absorbers tracing gaseous structures which are associated with the extended environments of galaxy G. \\citet{tumlinson05b} have described the properties of two multiphase {\\OVI} absorbers along the line of sight to the low-$z$ QSO PG~$1211+143$. A galaxy survey of the surrounding field showed that the two absorbers are situated within $\\sim 150h^{-1}$~kpc of $\\sim L^*$ galaxies. One of those galaxies is part of a spiral-dominated group. The {\\OVI} in both those absorption systems are inconsistent with a photoionization origin. The $b(\\OVI)$ suggests gas temperatures of $T \\gtrsim 10^5$~K. Ion ratios in the two absorbers resemble the highly ionized Galactic HVCs. \\citet{tumlinson05} propose that these absorbers are related to the nearby galaxies by outflows or tidal streams created from interactions with unseen satellite galaxies. It is possible for the $z = 0.22496$ absorber along the H~$1821+643$ sight line to have an analogous origin. The multicomponent, multiphase absorption profile and the abundance ratio [Si/H] $\\sim -0.6$~dex in the photoionized gas agrees with a description of the absorber tracing tidal debris embedded in the halo of its nearby galaxy G. An extended geometry like that of the Magellanic stream would result in several conductive layers at regions where the high velocity gas interfaces with the galaxy's {\\it hot} corona. The line of sight intercepting such multiple interface layers would explain the significant velocity spread of $\\Delta v \\sim 200$~{\\kms} in {\\OVI} and {\\HI} absorption as well as the $N_a(\\OVI) \\sim 10^{14.3}$~{\\cmsq}. The physical nature of the $z = 0.22638$ system is relevant for other BLAs with metals detected in the same gas phase. BLAs are characterized by large line widths $b \\geq 40$~{\\kms} and shallow absorption $N(\\HI) < 10^{14}$~{\\cmsq} \\citep{richter04, richter06}. It is unclear whether the large line width is dominated by thermal or non-thermal effects. Thermal broadening would imply that BLAs are large reservoirs of baryons in the low-$z$ IGM. Detecting metals in the same phase as the broad {\\Lya} gas will enable an accurate temperature estimation. For the $z = 0.22638$ BLA, the $b(\\HI)$ and $b(\\CIII)$ imply $T \\sim 1.1 \\times 10^5$~K. The temperature indicates the gas to be heavily ionized ($f_{\\HI} = N(\\HI)/N(H) \\sim 10^{-5}$) with a substantial baryon column density of $N(\\H) \\sim 3.2 \\times 10^{18}$~{\\cmsq}. The $z = 0.16339$ BLA absorber with $b(\\HI) = 46.3~{\\pm}~1.9$~{\\kms} detected in the HE~$0226-4110$ spectrum (Lehner {\\etal}2006) resembles the $z = 0.22638$ absorber. In the $z = 0.16339$ system, there is a marginal ($2.9\\sigma$) {\\CIII} detection that appears shallow and broad, with no associated {\\OVI}. The lack of {\\OVI} detection points to low metallicity in the gas. If the gas is in CIE at $T = 1.3 \\times 10^5$~K suggested by the line width of {\\HI}, then the $N_a(\\CIII)/N_a(\\HI) \\sim 0.01$ ratio yields a carbon abundance of [C/H] $\\sim -1.8$~dex, an ionization correction of $f({\\HI}) \\sim 9 \\times 10^{-6}$ and a $N(\\H) \\sim 2.5 \\times 10^{19}$~{\\cmsq}. These values are comparable to the low abundance and the high baryon content for the BLA phase in the $z = 0.22638$ system. Interestingly, the $z = 0.16339$ BLA is also coincident in redshift ($z = 0.1630$) with a relatively bright galaxy ($m_B = 23.74$) at a projected separation of $\\sim 226~h_{71}^{-1}$~kpc \\citep{chen09}. The $|\\Delta v| \\sim 117$~{\\kms} velocity separation between the absorber and the galaxy favors an origin for the absorption in some high velocity gas associated with the galaxy. The similarities in derived properties between the $z = 0.22638$ and $z = 0.16339$ systems and their proximity to galaxies further emphasizes the likelihood of BLAs with metals in the same gas phase selecting low metallicity high velocity gas systems in external galaxies. The prospect for detecting such extragalactic absorbers is going to be significantly augmented in the near future with the advent of higher sensitivity observations using the Cosmic Origins Spectrograph. The $z = 0.22496$ and $z = 0.224638$ absorbers being extragalactic high velocity clouds has implications for our understanding of the nature of other {\\OVI} absorption systems. In a sample of 51 {\\OVI} systems at $z < 0.5$ identified along 16 sight lines, Tripp {\\etal}(2008) found 53\\% of the absorbers to be complex multiphase systems with significant velocity offset between the {\\OVI} and {\\HI} absorbing components. Even among the fraction of absorbers with closely aligned {\\OVI} and {\\HI} components, the temperature implied by the combined line widths of {\\HI} and {\\OVI} in 26\\% of the cases was $T \\gtrsim 4 \\times 10^4$~K. This temperature lower limit is inconsistent with photoionized gas. In such absorbers, nonequilibrium collisional ionization process is a distinct possibility for the production of {\\OVI}. The {\\OVI} could be tracing transition temperature plasma in high velocity gas clouds embedded within a galaxy halo. The fact that many {\\OVI} absorbers are clustered around galaxies further supports this possibility \\citep{wakker09, chen09}." }, "1003/1003.4445_arXiv.txt": { "abstract": "A photometric calibration of Kurucz static model atmospheres is used to obtain parameters of RR Lyrae stars: variation of stellar angular radius $\\vartheta$, effective temperature $T_{\\rm e}$, and gravity $g_{\\rm e}$ as a function of phase, interstellar reddening $E(B-V)$ toward the star, and atmospheric metallicity $M$. Photometric and hydrodynamic conditions are given to find the phases of pulsation when the quasi-static atmosphere approximation (QSAA) can be applied. The QSAA is generalized to a non-uniformly moving spherical atmosphere, and the distance $d$, mass ${\\cal M}$, and atmospheric motion are derived from the laws of mass and momentum conservation. To demonstrate the efficiency of the method, the $UBV(RI)_C$ photometry of SU Dra was used to derive the following parameters: $[M]=-1.60\\pm .10$~dex, $E(B-V)=0.015\\pm .010$, $d=663\\pm 67$~pc, ${\\cal M}=(0.68\\pm .03){\\cal M}_\\odot$, equilibrium luminosity, $L_{\\rm eq}=45.9\\pm 9.3L_\\odot$, $T_{\\rm eq}=6813\\pm 20$~K. ", "introduction": "Although the determination of the fundamental parameters of RR Lyrae (RR) stars is interesting in itself, it is also important from a practical point of view, because RR stars play a considerable role in establishing galactic and extragalactic distance scales. The Preston index and spectroscopic observations are used to determine their atmospheric metallicity $[M]$ and interstellar reddening $E(B-V)$. Since confirmed RR type components are not known in binary systems, the mass determination is based on both stellar evolution and pulsation theories. Due to the uncertainty of parallax data, mostly the Baade-Wesselink (BW) method is used to infer their distance $d$ \\citep{smit1}. In the BW analysis, and to determine $[M]$ and $E(B-V)$, the quasi-static atmosphere approximation (QSAA) is employed to interpret photometry and spectroscopy. The QSAA was introduced by \\citet{ledo1}: {\\it ``The simplest approach is to assume that at each phase, the atmosphere adjusts itself practically instantaneously to the radiative flux coming from the interior and to the effective gravity $g_{\\rm e}$ \\begin{equation}\\label{1.102} g_{\\rm e}=G{\\cal M}R^{-2}+{\\ddot R} \\end{equation} where $R,\\: \\mbox{and}\\: {\\ddot R}$ are the instantaneous values of the radius and acceleration, which is supposed uniform throughout the atmosphere''}, $G$ is the Newtonian gravitational constant, ${\\cal M}$ is the stellar mass, and dot is a differentiation with respect to time $t$. {\\it ``One may then build a series of static model atmospheres,''} and select one of them at each phase by spectroscopic or photometric observations. Its flux, colours, effective temperature $T_{\\rm e}$, and surface gravity $g_{\\rm e}$ are accepted as the atmospheric parameters of that phase providing basis for determination of other parameters like angular radius, mass, distance, etc. The subject of the present paper is the QSAA and its generalization to a non-uniform atmosphere. We will investigate the QSAA from the point of view of atmospheric emergent flux and hydrodynamics. By comparing the observed colour indices with those of static model atmospheres (\\citealt{cast1}, \\citealt{kuru1}) we select the phases when they coincide. Considering hydrodynamics, we do not construct a consistent dynamic model of an RR atmosphere. However, we find a better description of the pulsating atmosphere if we characterize it by pressure and density stratifications in addition to the two parameters $R$ and ${\\ddot R}$. We determine the fundamental parameters ${\\cal M}$ and $d$ from the hydrodynamic considerations without using the BW method or theories of stellar evolution and pulsation at all. Our method uses photometry as observational input; spectroscopy and radial velocity observations are not needed. In Section 2, conditions of QSAA are formulated for a spherically pulsating compressible stellar atmosphere with velocity gradient. Practical methods are described to determine $[M], E(B-V)$, $T_{\\rm e}(\\varphi)$, and $\\log g_{\\rm e}(\\varphi)$ from $UBV(RI)_C$ colours of Kurucz atmospheric models, $\\varphi$ being the pulsation phase. The laws of mass and momentum conservation are used to determine the mass and distance of the pulsating star. Section 3 presents the results obtained from the $UBV(RI)_C$ photometry of SU Dra. Discussion and conclusions are given in Sections 4 and 5. ", "conclusions": "Observed colours and magnitudes of a spherically pulsating star have been compared with those of static Kurucz model atmospheres to determine fundamental parameters of the star in the frame of the quasi-static atmosphere approximation. Photometric and hydrodynamic conditions have been formulated for the validity of the quasi-static atmosphere approximation in spherically pulsating stars. \\begin{itemize} \\item[(1)] The quasi-static atmosphere approximation has been generalized for a non-uniform, compressible atmosphere with radial velocity gradient. This is a step forward because the hitherto available uniform atmosphere approximation described the motion of the atmosphere by effective gravity and differentiating the sole parameter radius $R$ of optical depth zero with respect to time. \\item[(2)] Our combined photometric and hydrodynamic method uses the variation of effective gravity, angular radius, velocity, acceleration in the Euler equation. The complete input of the Euler equation has been derived from photometry and theoretical model atmospheres. Spectroscopic and radial velocity observations were used neither explicitly nor implicitly. This is a definite advantage in comparison with the Baade-Wesselink method, because less observational efforts are needed to get the fundamental parameters and the problematic conversion of observed radial velocities to pulsation velocities is not necessary. \\item[(3)] Concerning the interpretation of multicolour photometry, the inputs are identical with those of the Baade-Wesselink method. The present refinements are the quantitative conditions whether photometry can or cannot be interpreted by static model atmospheres. \\item[(4)] Firstly, the phases have been selected by the photometric conditions when the quasi-static atmosphere approximation is valid. Secondly, in a number of these phases, the laws of mass and momentum conservation have been applied in Euler formalism of hydrodynamics to determine mass and distance of the star from the motion of the atmospheric layers in the neighbourhood of zero optical depth. Afterwards, it was checked whether the hydrodynamic condition of the quasi-static atmosphere approximation was satisfied in the used phases, i.e. phases that satisfied the photometric condition but violated the hydrodynamic condition had to be excluded. {\\it Atmospheric dynamical mass} seems to be an appropriate term to indicate that the mass has been derived by a method which is completely different from a mass derived by pulsation or evolution theories \\citep{smit1}. \\item[(5)] As a by-product, a variation procedure has been given for estimating atmospheric metallicity and interstellar reddening toward a star from $UBV(RI)_C$ photometry. This method was successfully applied for the non-variable comparison star BD +67 708, giving $[M]=-0.77\\pm .03$, $E(B-V)=.00$, $\\vartheta=(1.913\\pm .002)\\times 10^{-10}$~rad, $\\log g=3.59\\pm .01$, $T_{\\rm e}=7505\\pm 5$~K. \\item[(6)] Using the $UBV(RI)_C$ photometry of the high amplitude RRab star SU Dra, the following fundamental parameters have been found: \\[[M]=-1.60\\pm .10,\\] \\[E(B-V)=0.015\\pm .010,\\] \\[d=(663\\pm 67)\\mbox{pc},\\] \\[R_{\\rm min}=4.46R_\\odot,\\:R_{\\rm max}=5.29R_\\odot,\\] \\[{\\cal M}=(0.68\\pm .03){\\cal M}_{\\odot},\\] \\[L_{\\rm eq}=(45.9\\pm 9.3)L_{\\odot},\\] \\[T_{\\rm eq}=(6813\\pm 20)\\mbox{K}.\\] $L_{\\rm eq}$ and $T_{\\rm eq}$ are approximate values, since they originate from averaging over the whole pulsation cycle containing phases in which the quasi-static atmosphere approximation is merely a first approximation. \\item[(7)] The internal motions of the atmosphere with respect to zero optical depth have been found to be significant in comparison with velocities and accelerations derived from the uniform atmosphere approximation. From the internal motions of the atmosphere, some constraints have been sketched for converting observed radial velocities to pulsation and centre of mass velocities. Estimations have been given for the error propagation in a Baade-Wesselink analysis. \\end{itemize}" }, "1003/1003.4818_arXiv.txt": { "abstract": "We study the effects of gravity waves, or $g$-modes, on hot extrasolar planets. These planets are expected to possess stably-stratified atmospheres, which support gravity waves. In this paper, we review the derivation of the equation that governs the linear dynamics of gravity waves and describe its application to a hot extrasolar planet, using \\HD\\ as a generic example. We find that gravity waves can exhibit a wide range of behaviors, even for a single atmospheric profile. The waves can significantly accelerate or decelerate the background mean flow, depending on the difference between the wave phase and mean flow speeds. In addition, the waves can provide significant heating ($\\sim\\! 10^2$ to $\\sim\\! 10^3$~K per planetary rotation), especially to the region of the atmosphere above about 10 scale heights from the excitation region. Furthermore, by propagating horizontally, gravity waves provide a mechanism for transporting momentum and heat from the dayside of a tidally locked planet to its nightside. We discuss work that needs to be undertaken to incorporate these effects in current atmosphere models of extrasolar planets. ", "introduction": "A stably-stratified atmosphere, characterized by a positive vertical entropy gradient, can support gravity waves, or $g$-modes. Gravity waves are oscillations which arise from the buoyancy of parcels in the fluid. Such waves are readily excited by flow over thermal and surface topography, convective and shear instabilities, and flow adjustment processes. They propagate through the atmosphere both horizontally and vertically. Gravity waves in the terrestrial atmosphere and ocean are much studied \\citep[e.g.,][]{Gossard1975,Gill1982}. They have also been observed on other Solar System bodies, such as Jupiter \\citep{Young1997} and Venus \\citep{Apt1980}. In the terrestrial atmosphere, a typical gravity wave has an energy flux of approximately 10$^{-3}$ to 10$^{-1}$~W~m$^{-2}$. Despite being small, compared to the total amount of absorbed solar flux ($\\sim$237~W~m$^{-2}$), gravity waves are responsible for significantly modifying---even dictating---large-scale flow and temperature structures. Several well known examples of this are the Quasi-Biennial Oscillation, reversal of mean meridional temperature gradient in the upper middle atmosphere, and generation of turbulence \\citep[e.g.,][]{Andrews1987}. We expect similar effects to be present in the atmospheres (and oceans) of extrasolar planets. Moreover, due to the greater irradiation and scale heights on them, the acceleration and heating effects of gravity waves can be much stronger on hot extrasolar planets. There has been much interest in modeling the atmospheric circulation of extrasolar planets \\citep[e.g.,][]{Joshi1997,Showman2002,Cho2003,Cho2008a,Burkert2005,Cooper2005,Dobbs-Dixon2008,Koskinen2007,Langton2007,Langton2008,Menou2009,Showman2008}. Accurate simulation of atmospheric circulation is crucial for interpreting observations of extrasolar planets, as well as for improving theoretical understanding in general. For this, the role of eddies and waves in transferring momentum and heat needs to be addressed \\citep{Cho2008b}. This has long been recognized in Solar System planet studies \\citep[e.g.,][]{Lindzen1990,Fritts2003}. The plan of the paper is as follows. In \\S\\ref{sec:theory} we derive the governing equation appropriate for linear monochromatic gravity waves on hot extrasolar planets. We also discuss a simple parameterization of the key non-linear process, saturation. In addition, we present solutions to the equation for simple isothermal atmospheres, with and without shear in the background mean flow. In \\S\\ref{sec:application} we extend the calculation to a physically more realistic situation, by using background flow and temperature profiles derived from a three-dimensional (3-D) hot--Jupiter atmospheric circulation simulation. This is the first such calculation to have been performed for extrasolar planets. Through this, the significant effects of gravity waves on hot extrasolar planet atmospheric mean flows are demonstrated. In this section, we also discuss a way in which gravity waves can transport momentum and heat horizontally---e.g., from the dayside to nightside on tidally locked planets. In \\S\\ref{sec:implications} we discuss the implications of our work for current extrasolar planet atmospheric modeling work. We conclude in \\S\\ref{sec:conclusion}. ", "conclusions": "\\label{sec:conclusion} Gravity wave propagation and momentum and energy deposition are complicated by the environment in which the wave propagates. For example, spatial variability of the background winds causes the wave to be refracted, reflected, focused, and ducted. Additionally, temporal variability of the background winds cause the wave to alter its phase speed. Still further complications arise due to the wave's ability to generate turbulence, which can modify the source or serve as a secondary source, and the wave's interaction with the vortical (rotational) mode of the atmosphere. Many of these issues are as yet not well-understood and are currently areas of active research. In this work, we have emphasized only some of these issues. We have shown that gravity waves propagate and transport momentum and heat in the atmospheres of hot extrasolar planets and that the waves play an important role in the atmosphere. They modify the circulation through exerting accelerations on the mean flow whenever the wave encounters a critical level or saturates. They also transport heat to the upper stratosphere and thermosphere, causing significant heating in these regions. Moreover, through ducting, they also provide a mechanism for transporting heat from the dayside of tidally synchronized planets. Before relying on GCMs for quantitative descriptions of hot extrasolar planet atmospheric circulations, further work needs to be performed to ensure that the effects of important sub-scale phenomena, such as the gravity waves discussed here, are accurately parameterized and included in the GCMs." }, "1003/1003.2263_arXiv.txt": { "abstract": "{ The study of low-degree high-frequency waves in the Sun can provide new insight into the dynamics of the deeper layers of the Sun. Here, we present the analysis of the velocity observations of the Sun obtained from the Michelson and Doppler Imager (MDI) and Global Oscillations at Low Frequency (GOLF) instruments on board Solar and Heliospheric Observatory ({\\em SOHO}) spacecraft for the major flare event of 2003 October 28 during the solar cycle 23. We have applied wavelet transform to the time series of disk-integrated velocity signals from the solar surface using the full-disk Dopplergrams obtained from MDI. The wavelet power spectrum computed from MDI velocity series clearly shows that there is enhancement of high-frequency global waves in the Sun during the flare. We do observe this signature of flare in the Fourier Power Spectrum of these velocity oscillations. However, the analysis of disk-integrated velocity observations obtained from GOLF shows only feeble effect of flare on high-frequency oscillations.} ", "introduction": "\\sloppy The normal modes of oscillations of the Sun peak in the frequency regime 2-4 mHz and are known as $p$ modes. Apart from the normal $p$ modes, researchers have found the presence of high-frequency oscillations (frequencies higher than the solar-photospheric acoustic cut off at $\\sim$ 5.3 mHz) in the solar-acoustic spectrum \\citep{libb88a, libb88b, garcia98, chap03, jim05, karoff08}. Unlike the $p$ modes, the driving force for these high-frequency solar oscillations is still not clearly understood. On one hand, \\cite{balm90} suggest that the high-frequency waves are partly reflected by the sudden change in temperature at the transition region between the chromosphere and the corona, while \\cite{kumar91} explain these high-frequency waves as an interference phenomenon between ingoing and outgoing waves from a localized source just beneath the photosphere. \\sloppy Just after the advent of helioseismology in 1970s, \\cite{wolff72} suggested that large solar flares can stimulate free modes of oscillation of the entire Sun, by causing a thermal expansion that would drive a compression front to move into the solar interior. First of all, \\cite{haber88} reported an average increase in the power of intermediate-degree modes after a major flare (of class X13/3B) using a few hours of solar-oscillations data. However, \\cite{braun90} could not detect acoustic-wave excitation from an X-class flare. \\cite{koso98} reported the first detection of ``solar quakes'' inside the Sun, caused by the X2.6 flare of 1996 July 9. Following this result, \\cite{donea99} found an acoustic source associated with a flare using seismic images produced with helioseismic-holography technique. Application of ring-diagram analysis showed that the power of the global $p$ modes appears to be larger in several flare-producing active regions as compared with the power in non-flaring regions of similar magnetic field strength \\citep{ambastha03}. Further, \\cite{donea05} have reported emission of seismic waves from large solar flares using helioseismic holography. Some of the large solar flares have been observed to produce enhanced high-frequency acoustic velocity oscillations in localized parts of active regions \\citep{kumar06}. \\cite{venkat08} observed co-spatial evolution of seismic sources and H-alpha flare kernels and large downflows in the seismic sources during the flare event of 28 October 2003. A search for a correlation between the energy of the low-degree $p$ modes and flares using velocity observations of the Sun remained inconclusive \\citep{foglizzo98, gavry99, chap04, ambastha06}. The study of low-degree high-frequency (LDHF) waves in the Sun is important as this can bring more constraints on the rotation profile between 0.1 and 0.2~$R_\\odot$ \\citep{garcia_new08, mathur08}. The study of the effect of flares on such LDHF waves can provide a clue for the origin of these waves. Recently, \\cite{karoff08} have reported that the correlation between X-ray flare intensity and the energy in the acoustic spectrum of disk-integrated intensity oscillations (as observed with VIRGO (Variability of Solar IRradiance and Gravity) (\\cite{frohlich95}) instrument on board {\\em SOHO}) is stronger for high-frequency waves than for the well known 5-minute oscillations. In this study, we have searched for the effects of flares in the time series of disk-integrated velocity signals from the solar surface using the full-disk Dopplergrams obtained from the MDI instrument. We have also looked for these effects in disk-integrated velocity observations obtained from the GOLF (Global Oscillation at Low Frequency) \\citep{gabriel95, gabriel97} instrument on board {\\em SOHO}. These studies have been applied to the major solar flare of 2003 October 28 (of class X17.6/4B) that occurred in the solar cycle 23. Wavelet and Fourier analyses of MDI velocity observations clearly indicate the enhancement in high-frequency global waves in the Sun during the flare. However, this signature of flare is weaker in the case of GOLF as compared to MDI data. ", "conclusions": "Earlier attempts to find a correlation between the energy of these high-frequency oscillations and flares using disk-integrated velocity observations of the Sun had remained inconclusive \\citep{gavry99, chap04}. In our analysis, the enhancement of high-frequency power is clearly seen in the MDI velocity data (and a feeble enhancement is also seen in the GOLF velocity data) during the major flare event of 2003 October 28. It is in good agreement with the flare related enhancements reported by \\cite{karoff08} in disk-integrated intensity oscillations as observed with VIRGO. Although, in our results the flare induced enhancement signals are seen with a low probability (around 50\\%), this signature is larger than that seen in non-flaring condition for both the MDI and GOLF data. It is evident from the various models proposed for the generation of high-frequency waves that the amount of energy that is stored in the high-frequency waves is extremely low compared to the amount of energy stored in the normal $p$ modes which are powered by the strong turbulence in the convection zone of the Sun. Therefore, it is believed that the flare energy will have a larger relative effect at high frequency as the other sources of its excitation are much smaller. A more detailed discussion on this can be found in \\cite{kumar2010}. These observations provide us the opportunity to further investigate the possible excitation of global high-frequency waves by local tremors due to major solar flares. As a future study, we plan to correlate the epochs of enhancement of high-frequency waves with episodes of flare energy release seen in hard X-ray observations of the Sun. {\\bf Acknowledgments} The use of data from the MDI and the GOLF instruments on board {\\em SOHO} spacecraft is gratefully acknowledged. The {\\em SOHO} is a joint mission under cooperative agreement between ESA and NASA. This work has been partially supported by the CNES/GOLF grant at the Service d'Astrophysique (CEA/Saclay). We are thankful to Douglas Gough, John Leibacher, Frank Hill, P. Scherrer, Robertus Erdelyi, H. M. Antia, A. Kosovichev and Christoffer Karoff for useful discussions related to this work." }, "1003/1003.0099_arXiv.txt": { "abstract": "{Multiple rebrightenings have been observed in the multiband afterglow of GRB 030329. Especially, a marked and quick rebrightening occurred at about t $\\sim$ 1.2 $\\times$ 10$^{5}$ s. Energy injection from late and slow shells seems to be the best interpretation for these rebrightenings. Usually it is assumed that the energy is injected into the whole external shock. However, in the case of GRB 030329, the rebrightenings are so quick that the usual consideration fails to give a satisfactory fit to the observed light curves. Actually, since these late/slow shells coast freely in the wake of the external shock, they should be cold and may not expand laterally. The energy injection then should only occur at the central region of the external shock. Considering this effect, we numerically re-fit the quick rebrightenings observed in GRB 030329. By doing this, we were able to derive the beaming angle of the energy injection process. Our result, with a relative residual of only 5\\% --- 10\\% during the major rebrightening, is better than any previous modeling. The derived energy injection angle is about 0.035. We assume that these late shells are ejected by the central engine via the same mechanism as those early shells that produce the prompt gamma-ray burst. The main difference is that their velocities are much slower, so that they catch up with the external shock very lately and manifest as the observed quick rebrightenings. If this were true, then the derived energy injection angle can give a good measure of the beaming angle of the prompt $\\gamma$-ray emission. Our study may hopefully provide a novel method to measure the beaming angle of gamma-ray bursts. ", "introduction": "\\label{sect:intro} GRB 030329, which is one of the brightest gamma-ray bursts (GRBs), has a fluence of about 1.18 $\\times$ 10$^{-4}$ ergs$\\cdot$ cm$^{-2}$ (\\citealt{ric03}; \\citealt{van04}). It is also very close to us. The redshift is z = 0.168 (\\citealt{gre03}). At the same time, GRB 030329 is unambiguously confirmed to be associated with supernova (\\citealt{hjo03}; \\citealt{sta03}; \\citealt{mat03}). Because of these important characteristics, this GRB has attracted a lot of attentions. Numerous and detailed multiband afterglow observations have been accumulated. Multiple rebrightenings were observed in the afterglow of GRB 030329. Especially, a quick and marked rebrightening occurred at about t $\\sim$ 1.2 $\\times$ 10$^{5}$ s. Several models have been proposed to explain the quick and marked rebrightenings. (i) Density-jump model: if the external shock encounters a sudden density variance of the circum-burst medium, the emission of the afterglow may be enhanced temporarily (\\citealt{wan00}; \\citealt{laz02}; \\citealt{nak03}; \\citealt{dai03}; \\citealt{tam05}). But for GRB 030329, it has been argued that the density-jump model could not produce the observed rebrightenings (\\citealt{nak03}; \\citealt{wil04}; \\citealt{pir04}; \\citealt{hua06}). (ii) The two-component jet model is another kind of choice (\\citealt{ber03}). But \\cite{hua06} made a detailed numerical calculation and found that although the two-component jet model could basically reproduce the overall R-band light curve, it was unable to explain the steep rebrightenings. (iii) \\cite{wil04} discussed a possibility that the rebrightening might be due to a supernova component, but their model is unlikely to produce multiple rebrightenings. Actually, the energy-injection model seems to be the best explanation for the steep rebrightenings (\\citealt{hua06}). \\cite{gra03} have made a general analysis on this interpretation and \\cite{hua06} presented a detailed numerical study. In Huang et al.'s study, it is interesting to note that although the fitting result was much better than previous models, the theoretical rebrightenings still could not be as steep as observations. This problem was probably caused by their simple assumption that the energy carried by the late (and initially slower) shells was injected into the whole external shock homogeneously. In reality, it was probable that the late/slow shells, which coast freely in the wake of the previous external shock, should be cold and do not expand laterally, as illustrated by \\cite{gra03}. So when the late/slow shells catch up with the external shock, the energy should only be injected into the central region of the external shock. We call this scenario as a localized energy-injection scenario. We conjecture that the energy-injection angle can be derived from the rapidness of the observed rebrightenings. Since these late/slow shells basically may have the same origin as those early shells that produce the internal shocks and the main GRB, we argued that this method could be used to hint the degree of collimation in GRBs. This will be a novel way to measure the beaming angle of GRBs, in addition to the traditional jet-break timing method. In this paper, we will try to reproduce the multiband afterglow light curves of GRB 030329 numerically, based on the localized energy-injection scenario. The structure of our paper is organized as follows. Section 2 is a detailed description of the model. In Section 3, we study the effects of various parameters on the rapidness of the rebrightening. In Section 4, we simulate the observed multiband afterglow light curves of GRB 030329 in the framework of our model. Section 5 presents a brief discussion and the conclusion. ", "conclusions": "\\label{sect:Conclusion} The rapid rebrightening is an interesting feature of the afterglow of GRB 030329. Conventional homogeneous energy-injection model cannot account for the rapidness of the observed rebrightenings. In this study, it is clearly shown that the localized energy-injection model can satisfactorily resolve the rapidness problem. By using a two-component jet model, together with 5 localized energy-injections, we can satisfactorily reproduce the overall and multiband afterglow light curves of GRB 030329. In our scenario, the narrow component gives birth to the prompt GRB emission and the early afterglow, while the wide component contributes significantly to the late afterglow. The energy-injections are induced by 5 late shells, which should be ejected by the central engine via the same or similar mechanism as the narrow jet component, but at a slightly later period and with much lower velocities. These shells coast freely in the empty environment after the external shock. They keep to be cold and do not experience any lateral expansion until they catch up with the external shock and inject their energies into the GRB remnant. The energy-injection is a prompt process, and most importantly, is a localized collision. It can lead to a very rapid increase in the afterglow brightness, barely subjected to the smearing of the equal arrival time surface effect. The initial beaming angle is an important parameter of GRBs. It can provide important clues for the central engine (\\citealt{cl01}; \\citealt{zh07}; \\citealt{gao10}). According to our best fit, the energy-injection angle is about 0.035 in the case of GRB 030329. In our modeling, this angle also equals to the half opening angle of the narrow jet component. We also showed that the rapidness of the rebrightening is mainly controlled by the energy-injection angle. Other parameters, such as $\\epsilon_{\\rm e}$, $\\epsilon_{\\rm B}^{2}$, $p$, and $n$ only have minor or negligible effect. So, we suggest that the rapidness of the observed rebrightenings can basically be used to derive a measure of the beaming angle of GRBs. This is a novel and hopeful method, independent to the conventional jet-break timing method. Comparing with the jet-break timing method, the advantage of our method is that it does not rely on the assumption of the speed of lateral expansion, which itself is quite unclear currently. The impact of another factor, i.e. the density of the circum-burst environment, is also very week in our method. However, the main restriction of our method is that it is applicable only when rapid rebrightenings due to energy-injections were observed in the afterglow. In our modeling of the five observed rebrightenings of GRB 030329, the beaming angles of the 5 energy-injections are set to be constant. The fitting results are good. It means that the ejecting angle of the central engine might not vary significantly over the active period. However, in the future, it is possible that we could find new examples which need to vary the energy-injection angle. Such information will also provide interesting constraints on the central engine. In our model, we have assumed that the energy-injection angle equals to the beaming angle of the narrow jet that produces the prompt GRB. Although this assumption is a natural hypothesis, it is still possible that in reality they might not equal, since the real central engine may be much more complicated as compared with our imagination. Even in this case, the energy-injection angle derived from observations still can give useful hints on the characteristics of the central engine. In our current study, in order to get an acceptable fit to the overall and multiband afterglow of GRB 030329, we have adopted a two-component jet geometry. However, note that such a two-component configuration is not always necessary for localized energy-injection. For example, in other GRBs, it is quite possible that the prompt GRB jet may only have one component. This component can expand laterally during its deceleration, still leaving a relatively clean environment behind it. So, late shells will still coast freely after the external shock. They keep to be cold and will not experience sideways expansion before the final energy-injection. For GRB 030329, in order to get an acceptable fit to the overall and multiband afterglow, we have adopted a $p$ value that is less than 2. However, as stated in section 3, $p$ does not influence the rapidness of the rebrightening. So $p<2$ is not necessary for the localized energy-injection model. In other GRBs, different $p$ can be adopted according to observations. In many GRBs, multiple optical and/or X-ray flares are observed to be superposed on the early afterglow light curves. The rising and falling of these flares are even more rapid than that in the case of GRB 030329. Although the nature of these flares is still largely unknown, we suggest that useful information on the beaming angle of these flares can also be derived by fitting the rising and falling profiles. \\normalem" }, "1003/1003.0861_arXiv.txt": { "abstract": "{} {We study the dwarf spheroidal galaxies in the nearby M81 group in order to construct their photometric metallicity distributions and to investigate the potential presence of population gradients. We select all the dwarf spheroidals with available Hubble Space Telescope\\,/\\,Advanced Camera for Surveys archival observations, nine in total.} {We interpolate isochrones so as to assign a photometric metallicity to each star within a selection box in the color-magnitude diagram of each dwarf galaxy. We assume that the dwarf spheroidals contain mainly an old stellar population. In order to search for metallicity gradients, we examine the spatial distribution of two stellar populations that we separate according to their metallicities.} {As a result, we present the photometric metallicity distribution functions, the cumulative histograms and smoothed density maps of the metal-poor and metal-rich stars as well as of the intermediate-age stars.} {From our photometric data we find that all the dwarf spheroidals show a wide range in metallicities, with mean values that are typical for old and metal-poor systems, with the exception of one dwarf spheroidal, namely IKN. Some of our dwarf spheroidals exhibit characteristics of transition-type dwarfs. Compared to the Local Group transition type dwarfs, the M81 group ones appear to have mean metallicity values slightly more metal-rich at a given luminosity. All the dwarf spheroidals considered here appear to exhibit either population gradients or spatial variations in the centroids of their metal-poor and metal-rich population. In addition, there are luminous AGB stars detected in all of them with spatial distributions suggesting that they are well mixed with the old stars.} ", "introduction": "\\label{sec:introduction} The dwarf galaxies within our Local Group (LG) have been the subject of intensive spectroscopic and photometric observations in different wavelength regimes and thus are well studied objects. Their study has been facilitated by the proximity of these dwarfs so that individual stars may even be resolved down to the main sequence turn-off, depending on their distance. Thus, extending the studies to dwarf galaxies in nearby groups with different environment and comparing their properties are of great importance in order to understand the main drivers of their evolution. In addition, the derived properties can provide a way to constrain models of galaxy formation as well as chemical evolutionary models In this respect, the \\object{M\\,81 group} is an interesting target: despite several differences, it bears close resemblance to our LG. The similarity of the M\\,81 group to our LG lies in its binary structure (Karachentsev et al.~\\cite{sl_m81distances}), while its difference is mainly due to the recent interactions between its dominant consituents as revealed by the formation of tidal tails and bridges detected in HI observations (Appleton, Davies \\& Stephenson \\cite{sl_appleton}; Yun, Ho \\& Lo \\cite{sl_yun}). With a mean distance of $\\sim$3.7~Mpc (Karachentsev et al.~\\cite{sl_m81distances}), the M\\,81 group is one of the nearest groups to our own LG. It consists of about 40 dwarfs of both early-type and late-type, with the addition of 12 recently discovered dwarf candidates (Chiboucas, Karachentsev \\& Tully \\cite{sl_chiboucas}). Here we focus on the dwarf spheroidal galaxies (dSphs) in the M\\,81 group with available Hubble Space Telescope (HST)\\,/\\,Advanced Camera for Surveys (ACS) archival data. The dSphs are objects with low surface brightness and poor in gas content. For a summary of their properties we refer to Grebel, Gallagher \\& Harbeck (\\cite{sl_grebel}; and references therein). We use their color-magnitude diagrams (CMDs) to derive the photometric metallicity distribution functions (MDFs) and search for the potential presence of population gradients in the M\\,81 group dSphs. The use of the CMD to infer the star formation histories and MDFs is a very powerful tool. For nearby groups at distances, where individual red giants are beyond the reach of spectroscopy even with 8-10~m class telescopes, CMDs are the best means to derive evolutionary histories. With the use of HST observations of adequate depth, the upper part of the red giant branch (RGB) can be resolved into single stars. Many studies have derived the photometric MDFs of distant LG dSphs (for example Cetus by Sarajedini et al.~\\cite{sl_sarajedini}; And\\,VI and And\\,VII by Grebel \\& Guhathakurta \\cite{sl_grebel99}) from their CMDs. A similar work to derive the photometric MDFs for dwarf galaxies in nearby groups has not been done so far. The search for radial population gradients in LG dwarf galaxies has been favoured by the fact that the resolved stellar populations reach the horizontal branch or extend even below the main-sequence turn-off depending on the distance of the dwarf, permitting one to use a variety of different stellar tracers. There are several studies for population gradients in the LG dwarfs and as an example of such studies we refer to the work done by Hurley-Keller, Mateo \\& Grebel (\\cite{sl_hurley-keller99}), Harbeck et al.~(\\cite{sl_harbeck}), Battaglia et al.~(\\cite{sl_battaglia06}) (photometric) and Tolstoy et al.~(\\cite{sl_tolstoy}), Koch et al.~(\\cite{sl_koch06}) (spectroscopic). There is not any study so far searching for population gradients in nearby group dwarf galaxies. This paper is structured as follows. In \\S2 we present the observations, in \\S3 we show our results, in \\S4 we discuss our main findings and in \\S5 we present our conclusions. ", "conclusions": "\\label{sec:conclusions} We use the CMDs of nine dSphs in the M\\,81 group to construct their photometric MDFs. These MDFs show populations covering a wide range in metallicity with low mean metallicities indicating that these are metal-poor systems. All MDFs show a steeper fall-off at their high-metallicity end than toward their low-metallicity end indicating that galactic winds may play a role in shaping their distribution. We compute the mean metallicity, $\\langle$[Fe/H]$\\rangle$, and the mean metallicity weighted by the metallicity error, $\\langle$[Fe/H]$\\rangle_{w}$, along with their corresponding standard deviations. The most metal-rich dSph in our sample is IKN even though it is the least luminous galaxy in our sample. IKN's comparatively high metallicity may indicate that it is a tidal dwarf galaxy or that it suffered substantial mass loss in the past. We do not see any correlation between the $\\langle$[Fe/H]$\\rangle$ and the deprojected distance from the M\\,81 galaxy, R. We use the mean metallicity weighted by the metallicity errors, $\\langle$[Fe/H]$\\rangle_{w}$, to select two stellar populations having metallicities above and below that value. For these two stellar populations we construct cumulative histograms, as a way to search for population gradients in metallicity. We find that some dSphs show strong metallicity gradients, while others do not. In dSphs with radial metallicity gradients the more metal-rich populations are more centrally concentrated. Furthermore, we study the spatial (i.e., two-dimensional) distribution of our defined metal-rich and metal-poor stellar populations. This refined look no longer assumes radial symmetry, and we now find that in some dwarfs the metal-rich population is more centrally concentrated, while others show offsets in the centroid of the two populations. By examining the distribution of the luminous AGB stars, we conclude that, for the majority of the dSphs, these stars have mostly extended distributions, indicating that they have been well-mixed with the metal-poor stellar population. We do not find any correlation between the fraction of luminous AGB stars and the deprojected distance from the M\\,81 galaxy. While present-day distances may not be indicative of the dwarfs' position in the past and while their orbits are unknown, the apparent lack of a correlation between distance and evolutionary history may suggest that the evolution of the dwarfs was determined to a large extend by their internal properties and not so much by their environment. Finally, there are some M\\,81 dSphs that straddle the transition region between LG dSphs and dIrrs in the metallicity-luminosity relation. We may be observing low-luminosity transition-type dwarfs moving toward the dSph locus. Interestingly, these dwarfs are slightly more luminous than the bulk of the LG transition dwarfs. Perhaps some of the M\\,81 dwarfs experienced gas stripping during the recent interactions between the dominant galaxies in the M\\,81 group." }, "1003/1003.3652_arXiv.txt": { "abstract": "SNS is a MATLAB-based software library written to aid in the design and analysis of receiver architectures. It uses electrical scattering matrices and noise wave vectors to describe receiver architectures of arbitrary topology and complexity. It differs from existing freely-available software mainly in that the scattering matrices used to describe the receiver and its components are analytic rather than numeric. This allows different types of modeling and analysis of receivers to be performed. Non-ideal behavior of receiver components can be parameterized in their scattering matrices. \\textbf{SNS} enables the instrument designer to then derive analytic expressions for the signal and noise at the receiver outputs in terms of parameterized component imperfections, and predict their contribution to receiver systematic errors precisely. This can drive the receiver design process by, for instance, allowing the instrument designer to identify which component imperfections contribute most to receiver systematic errors, and hence place firm specifications on individual components. Using \\textbf{SNS} to perform this analysis is preferable to traditional Jones matrix-based analysis as it includes internal reflections and is able to model noise: two effects which Jones matrix analysis is unable to describe. \\textbf{SNS} can be used to model any receiver in which the components can be described by scattering matrices. Of particular interest to the sub-mm and terahertz frequency regime is the choice between coherent and direct detection technologies. Steady improvements in mm and sub-mm Low Noise Amplifiers (LNAs) mean that coherent receivers with LNAs as their first active element are becoming increasingly competitive, in terms of sensitivity, with bolometer-based receivers at frequencies above $\\sim 100$~GHz. As an example of the utility of \\textbf{SNS}, we use it to compare two polarimeter architectures commonly used to perform measurements of the polarized Cosmic Microwave Background: differencing polarimeters, an architecture commonly used in polarization sensitive bolometer-based polarimeters; and pseudo-correlation polarimeters, an architecture commonly used in coherent, LNA-based, polarimeters. We parameterize common sources of receiver systematic errors in both architectures and compare them through their Mueller matrices, which encode how well the instruments measure the Stokes parameters of the incident radiation. These analytic Mueller matrices are used to demonstrate the different sources of systematic errors in differencing and correlation polarimeters. ", "introduction": "Many fields of astrophysics aim to measure increasingly faint signals. For instance, there is great interest at present in detecting and characterizing the B-mode \\cite{Zaldarriaga:1997p247} of the polarized Cosmic Microwave Background (CMB). The strength of this signal is not yet determined by theory, but a strong upper limit is 170~nK \\cite{Page:2007p205}, a tiny fraction of the CMB total intensity signal ($\\sim2.7$~K). An instrument built to detect very faint signals will almost certainly be heavily affected by systematic errors. It is increasingly important to be able to model the effects of receiver systematic errors on the measured signal, and on the receiver sensitivity. We want to be able predict the level of receiver systematic errors, show their impact on the gathered data, and make quantitative comparisons between different receiver architectures. Previous analytic and semi-analytic approaches to characterizing systematic effects in receivers have usually employed Jones matrices to describe receiver components and Mueller matrices to characterize the effects of receiver systematics on the observed Stokes parameters \\cite{Hu:2003p1726, ODea:2007p1678}. The use of Jones matrices to describe individual receiver components has several shortcomings. Only the forward path of the signal through the instrument is modeled -- internal scattering caused by reflections from poorly matched components is not included; and Jones matrix modeling is unable to describe component noise, and hence receiver sensitivity. Modeling a receiver with a full analytic description of the outputs and sensitivity in terms of individual component parameters allows us to identify which parameters of each component in a receiver are most important, and concentrate our efforts on improving them. This paper introduces a technique and software for developing full analytic descriptions of receiver outputs and sensitivities in terms of lab-measureable errors in individual components. In this technique components are modeled by electrical scattering matrices. When describing a network of components with Jones matrices the forward-path cascaded response can be obtained through simple matrix manipulation and multiplication. The scattering matrix formulation does not share this simplicity of calculation: only the case of cascaded 2-port devices is amenable to a relatively simple analytic solution. This paper describes an algorithm for calculating the response of arbitrarily connected networks of components. We present software which implements this algorithm, and apply it to two common polarimeter architectures: differencing polarimeters, and pseudo-correlation polarimeters. This software allows us to make robust analytic comparisons of receiver architectures. Errors in individual receiver components can be parameterized and propagated into the description of the receiver performance, e.g. the instrument Mueller matrix. We hence have a powerful tool for guiding the instrument design process and diagnosing the causes of non-ideal instrument behavior. ", "conclusions": "We have presented {\\verb SNS }, a MATLAB-based software library written to aid in the design and analysis of receiver architectures. It uses electrical scattering matrices and noise wave vectors to describe receiver architectures of arbitrary topology and complexity. We use {\\verb SNS } to compare two polarimeter architectures commonly used to perform measurements of the polarized CMB: differencing polarimeters, an architecture commonly used in PSB-based polarimeters; and pseudo-correlation polarimeters, an architecture commonly used in coherent polarimeters. This analysis highlights the differing sources of systematic error in these architectures: $I$ to $P$ leakage in pseudo-correlation polarimeters is almost immune to gain imbalance, but sensitive to OMT cross polarization; while $I$ to $P$ leakage in differencing polarimeters is immune to OMT cross polarization, but very sensitive to power detection imbalance. We show how {\\verb SNS } can be used to calculate analytical expressions for the receiver noise temperature of arbitrary receivers. Analytic expressions for the receiver temperature of a pseudo-correlation polarimeter are derived, and are found to be consistent with those obtained from conventional receiver temperature calculations." }, "1003/1003.6077_arXiv.txt": { "abstract": "We have developed a one-dimensional thermochemical kinetics and diffusion model for Jupiter's atmosphere that accurately describes the transition from the thermochemical regime in the deep troposphere (where chemical equilibrium is established) to the quenched regime in the upper troposphere (where chemical equilibrium is disrupted). The model is used to calculate chemical abundances of tropospheric constituents and to identify important chemical pathways for CO-CH$_4$ interconversion in hydrogen-dominated atmospheres. In particular, the observed mole fraction and chemical behavior of CO is used to indirectly constrain the Jovian water inventory. Our model can reproduce the observed tropospheric CO abundance provided that the water mole fraction lies in the range $(0.25 - 6.0)\\times10^{-3}$ in Jupiter's deep troposphere, corresponding to an enrichment of 0.3 to 7.3 times the protosolar abundance (assumed to be H$_{2}$O/H$_{2}=9.61\\times10^{-4}$). Our results suggest that Jupiter's oxygen enrichment is roughly similar to that for carbon, nitrogen, and other heavy elements, and we conclude that formation scenarios that require very large ($>\\, $8 times solar) enrichments in water can be ruled out. We also evaluate and refine the simple time-constant arguments currently used to predict the quenched CO abundance on Jupiter, other giant planets, and brown dwarfs. ", "introduction": "The water abundance in Jupiter's deep atmosphere provides important clues for solar system formation and evolution and reveals conditions in the solar nebula at the time of giant-planet formation \\citep[e.g.,][]{lunine2004}. In addition, the planetary water inventory has important implications for the cloud structure, energy balance, thermal structure, and chemistry of the Jovian troposphere. Unfortunately, the deep water abundance is difficult to obtain by remote sensing methods because H$_{2}$O is expected to condense near the $\\sim$5 bar level in Jupiter's cold upper troposphere \\citep[e.g., see][]{taylor2004}, and clouds and other opacity sources limit the depth to which infrared and other wavelength radiation can penetrate. The only \\textit{in-situ} measurement of the Jovian water abundance, by the \\textit{Galileo} Probe Mass Spectrometer (GPMS), pertains to a meteorologically anomalous ``hot-spot'' region characterized by low cloud opacity, low mixing ratios of condensable species, high thermal emission, and a water abundance which increased with depth \\citep[e.g.,][]{orton1996,orton1998,niemann1998,ragent1998,sromovsky1998,wong2004}. Thus, it is unclear whether the probe descended deep enough to sample the well-mixed water abundance below the cloud base, and the \\textit{Galileo} probe value of H$_{2}$O/H$_{2}=4.9\\pm1.6\\times10^{-4}$ at the 19-bar level \\citep{wong2004} is generally considered to be a lower limit for Jupiter's O/H inventory \\citep[e.g.,][]{roosserote2004,taylor2004,wong2004,depater2005}. For this reason, chemical models must be used to determine the deep water abundance in Jupiter's atmosphere until further measurements become available, such as microwave observations from the \\textit{Juno} mission or future deep-atmosphere probes \\citep[e.g.,][]{bolton2006,atreya2004}. Several investigators have used clever methods to estimate the deep H$_{2}$O abundance of the giant planets by considering the observed tropospheric abundance of CO and other trace constituents and by investigating how H$_{2}$O chemistry and atmospheric transport can influence the abundance of these trace species \\citep[e.g.,][]{prinn1977,fegley1985apj,fegley1988,lodders1994,lodders2002,bezard2002,visscher2005}. However, \\citet{prinn1977}, and all subsequent modelers who used their kinetic schemes, were limited by a lack of key chemical kinetics data, and some of the initial kinetic assumptions have been shown to be incorrect \\citep{dean1987,yung1988,griffith1999,bezard2002,cooper2006}. In addition, some of the transport time-scale arguments in the earlier works have been shown to have inappropriate assumptions \\citep{smith1998,bezard2002}. These problems offset each other to an extent, such that models using earlier assumptions yield reasonable results \\citep[cf.][]{visscher2005,bezard2002}. \\citet{fegley1988} found that H$_{2}$O/H$_{2}= (0.46-5.8)\\times10^{-3}$ is consistent with CO kinetics and atmospheric mixing on Jupiter, whereas \\citet{bezard2002} derived H$_{2}$O/H$_{2}=(0.34-15.3)\\times10^{-3}$ using revised kinetic and transport time-scale parameters \\citep{page1989,yung1988,smith1998} and improved CO observations. Further refinement of this estimate was not possible in the B\\'ezard et al.~analysis due to uncertainties in reaction kinetics, convective mixing rates, and the back-of-the-envelope time-scale arguments used to derive the H$_{2}$O abundance. Here we attempt to improve the determination of the deep water abundance in Jupiter's atmosphere by taking advantage of recent updates in thermodynamic parameters and reaction rate coefficients and by using a numerical model to provide a more rigorous quantitative test of the simple kinetic vs. transport time-scale approach. With a numerical model, we implicitly solve the continuity equations for all tropospheric constituents, considering reaction kinetics and atmospheric transport. As a result, our model tracks the transition from the thermochemical regime in the deep troposphere (where chemical equilibrium is established), to a quenched regime in the upper troposphere (where chemical equilibrium is disrupted). In contrast to previous studies, we make no \\textit{a priori} selection of the reaction mechanism, nor of the rate-determining step for the chemical conversion of CO into CH$_{4}$. Instead, we input a full suite of chemical kinetic reactions connecting the different relevant tropospheric species, and allow the dominant chemical pathways for the conversion of CO$\\rightarrow$CH$_{4}$ to be identified from our model results. Furthermore, we make no assumptions about the mixing length scale but instead model atmospheric transport via diffusion for an assumed eddy diffusion coefficient profile. We explore the effects of the tropospheric water abundance on the chemical behavior of CO and other oxidized carbon gases and use the observed CO abundance to indirectly constrain the water inventory in Jupiter's deep atmosphere. The paper is organized as follows. We begin with a description of our chemical model and the CO chemical constraint in \\S\\ref{s Computational Approach}. We present our model results in \\S\\ref{s Model Results}, derive an estimate of Jupiter's deep water abundance using CO as an observational constraint, identify the dominant chemical pathways involved for CO$\\rightarrow$CH$_{4}$ conversion in Jupiter's troposphere, and discuss the chemical behavior of other oxygen-bearing carbon species. In \\S\\ref{s: Discussion} we discuss implications of our results for constraining Jupiter's total oxygen inventory and planetary formation scenarios, and conclude with a summary in \\S\\ref{s: Summary}. ", "conclusions": "\\label{s: Discussion} \\subsection{Implications for Jupiter's Total Oxygen Inventory} Water vapor is the dominant oxygen-bearing gas throughout Jupiter's atmosphere and is much more abundant than other oxygen-bearing gases (CO, OH, etc.). For this reason, the H$_{2}$O abundance in the troposphere is expected to be representative of the majority of Jupiter's total oxygen inventory. However, some oxygen ($\\sim$20\\%) is removed from the gas phase by oxide formation (rock) in the deep atmosphere \\citep[e.g.,][]{fegley1988,lodders2004apj,visscher2005}. This removal must be considered before evaluating the bulk planetary oxygen abundance and, in turn, the heavy-element composition of planetesimals during Jupiter's formation. The fraction of oxygen removed by rock depends upon the abundance of all rock-forming elements (Mg, Si, Ca, Al, Na, K, Ti) relative to the total oxygen abundance. Following the method of \\citet{visscher2005}, and using updated solar abundances from \\citet{lodders2009}, the relative abundances of water vapor, total oxygen, and the rock-forming elements can be written as \\begin{equation}\\label{eq: enrichment mass balance} E_{\\textrm{H}_{2}\\textrm{O}}=1.261E_{\\Sigma{\\textrm{O}}}-0.261E_{\\textrm{rock}} \\end{equation} where $E_{i}$ represents the enrichment (over solar ratios) for each component $i$. This expression serves as a general mass-balance constraint for the relative abundances of water, oxygen, and rock over a range of heavy element enrichments in Jupiter's interior \\citep{visscher2005}. To derive this expression, it was assumed that all of the rock-forming elements are equally enriched, and $E_{\\textrm{rock}}=2.74\\pm0.65$ was adopted based upon the ``deep'' tropospheric abundance of sulfur \\citep[as H$_{2}$S;][]{wong2004}, which behaves as a rock-forming element in meteorites \\citep[][]{lodders2004apj}. Using $E_{\\textrm{rock}}=2.74\\pm0.65$ in equation (\\ref{eq: enrichment mass balance}) along with the \\textit{Galileo} entry probe H$_{2}$O abundance of $E_{\\textrm{H}_{2}\\textrm{O}}=0.51\\pm0.17$ (see Table \\ref{tab: gas abundances}) yields a total oxygen abundance (characterized by $E_{\\Sigma\\textrm{O}}$) of $0.97\\pm0.40$ times the solar abundance of $\\Sigma\\textrm{O}/\\textrm{H}_{2}=1.212\\times10^{-3}$. Using our nominal model result of $E_{\\textrm{H}_{2}\\textrm{O}}=2.5\\pm0.5$, equation (\\ref{eq: enrichment mass balance}) yields a total oxygen enrichment of $E_{\\Sigma\\textrm{O}}=2.5\\pm0.8$. Further considering a 3x uncertainty in reaction kinetics and a range of plausible $K_{zz}$ values from $4\\times10^{7}$ to $1\\times10^{9}$ cm$^{2}$ s$^{-1}$, our water constraint (0.3 -- 7.3x solar) gives a total oxygen inventory of 0.7 to 6.5 times the solar $\\Sigma\\textrm{O}/\\textrm{H}_{2}$ ratio in Jupiter's interior. This value represents the bulk oxygen inventory of Jupiter's interior consistent with CO chemistry, and includes oxygen as water plus oxygen bound in rock. \\subsection{Implications for Planetary Formation} All viable giant-planet formation models must consider how heavy elements become entrained in the planet during its formation and evolution. An important constraint for such models is therefore the observed atmospheric abundances of gases such as CH$_{4}$, NH$_{3}$, H$_{2}$S, PH$_{3}$, and H$_{2}$O, which are taken to represent the planetary elemental inventories of C, N, S, P, and the majority of planetary oxygen, respectively. As illustrated in Fig.~\\ref{fig: enrichments}, observations of Jupiter's atmosphere show that C, N, S, P, Ar, Kr, and Xe are enhanced relative to solar element-to-hydrogen ratios \\citep{mahaffy2000,wong2004,lodders2004apj} by factors of 2-4. The enrichment in heavy elements is generally believed to be consistent with the core-accretion model for giant planet formation \\citep{mizuno1980}, in which a rock or rock-ice core initially forms and continues to grow through the accretion of solid planetesimals until it is massive enough to capture nebular gas \\citep{bodenheimer1986,lissauer1987,pollack1996}. In this scenario, the observed heavy element enrichments on Jupiter arise from degassing of the initial core material and the continued accretion of solid planetesimals, which will most likely vaporize before reaching the core \\citep[e.g.,][]{pollack1986}. What remains unclear is the source and composition of the planetesimals which provided the enrichment, and several scenarios have been proposed to explain the observed heavy-element abundances. One distinguishing characteristic of Jovian formation scenarios is the predicted water inventory in Jupiter's deep atmosphere. For example, trapping of heavy elements by hypothetical solar-composition icy planetesimals \\citep{owen1999,atreya2003,owen2006} would be expected to give an enrichment in oxygen (as H$_{2}$O) around 3$\\pm$1 times solar, similar that for the other heavy elements. Trapping of heavy elements in the form of clathrate hydrates near the snow line \\citep[e.g.][]{lunine1985,gautier2001,hersant2004} would yield higher water abundances, with predicted enrichments ranging from $\\sim$6 \\citep{mousis2009} to $\\sim$8 \\citep{alibert2005} to as high as $\\gtrsim$17 \\citep{gautier2001,gautier2001b} or $\\sim$19 \\citep{hersant2004} times the solar H$_{2}$O/H$_{2}$ ratio of $9.61\\times10^{-4}$. Accretion of carbon-rich planetesimals behind a nebular ``tar line'' \\citep{lodders2004apj} would give subsolar water abundances similar to that observed by the \\textit{Galileo} probe (0.51x solar). Our indirect constraint of a deep water abundance 0.3--7.3 times the solar H$_{2}$O/H$_{2}$ ratio from our kinetic-transport model is plausibly consistent with each of these formation mechanisms but precludes clathrate-hydrate scenarios that would require large ($>\\, $8x) water enrichments in Jupiter's deep atmosphere \\citep[e.g.,][]{gautier2001,gautier2001b,hersant2004,alibert2005}. We have developed a comprehensive thermochemical kinetics and diffusion model for Jupiter which correctly transitions between equilibrium chemistry in the deep troposphere and quenched/disequilibrium chemistry in the upper troposphere. We use this numerical model to compute the vertical abundance profiles for all carbon- and oxygen-bearing atmospheric constituents and to explore the chemical behavior of CO and other oxidized carbon species in Jupiter's deep atmosphere. We find that carbon monoxide is reduced to CH$_{4}$ via a mechanism similar to that proposed by \\citet{yung1988}; however, our model indicates that the rate-limiting reaction for CO reduction in Jupiter's atmosphere is H$_2$ + CH$_{3}$O $\\rightarrow$ CH$_{3}$OH + H rather than Yung et al.'s proposed reaction H + CH$_3$O + M $\\rightarrow$ CH$_3$OH + M. We also confirm the original analytic prediction of \\citet{prinn1977} that the mole fraction of CO will ``quench'' and remain constant with altitude when kinetic reaction rates can no longer compete with atmospheric mixing. This quenching occurs at the $\\sim$400 bar (1000 K) level in our nominal model. Carbon monoxide is not the only species to quench; virtually all atmospheric constituents will quench at some point where temperatures become low enough to inhibit the kinetics. Our kinetic-transport model quantitatively confirms the convenient, back-of-the-envelope time-scale approach currently used to explore quenched disequilibrium chemistry on giant planets and brown dwarfs \\citep[e.g.,][]{prinn1977,lewis1984,fegley1985apj,fegley1988,lodders1994,lodders2002,griffith1999,bezard2002,visscher2005}. We find that the the time-scale approach is valid for estimating Jupiter's water inventory, provided that the correct rate-limiting reaction is considered (which we find to be reaction R863, H$_2$ + CH$_{3}$O $\\rightarrow$ CH$_{3}$OH + H) and provided that the mixing length $L$ is determined via the procedure advocated by \\citet{smith1998}. Using the CO abundance reported by \\citet{bezard2002} as our observational constraint, our model-data comparisons indirectly constrain the Jovian deep water abundance to lie in the range 0.3--7.3 times the solar H$_{2}$O/H$_{2}$ ratio of $9.61\\times10^{-4}$. Our results suggest that the enrichment for oxygen (as H$_{2}$O) is similar, to within uncertainties, as that for carbon, nitrogen, and other heavy elements --- giant-planet formation scenarios that require very large ($>\\, $8x) enrichments in the water abundance (such as some clathrate-hydrate formation scenarios) are precluded. The subsolar water abundance (0.51x solar) measured by the \\textit{Galileo} entry probe \\citep{wong2004} remains plausibly consistent with the observed tropospheric abundance of carbon monoxide if relatively rapid vertical mixing (e.g., $K_{zz}\\ \\gtrsim 1\\times10^{9}$ cm$^2$ s$^{-1}$) prevails in Jupiter's deep troposphere. We will not be able to narrow our estimated range of Jovian deep water enrichments without experimental confirmation of the rate coefficient for the reaction of methoxy with H$_2$ (CH$_3$O + H$_2$ $\\rightarrow$ CH$_3$OH + H) at high pressures, for temperatures near 1000 K, and our results are subject to revision as updated kinetics data become available. Perhaps more importantly, we need a better understanding of appropriate diffusion coefficients that can be used to represent convective processes under tropospheric conditions on Jupiter. The \\textit{Juno} mission may provide microwave data \\citep{janssen2005,bolton2006} that can be used to test our model prediction regarding the Jovian deep water abundance. We point out that oxygen species are not the only constituents to quench in our model; the quenching of nitrogen species like N$_{2}$ and HCN is interesting in its own right and will be the subject of a future investigation. Our kinetics-transport model can easily be applied to other giant planets and brown dwarfs. Of particular interest is (1) constraining the deep water abundance on the other giant planets in our own solar system from current or future tropospheric CO observations (which must be able to constrain the vertical profile to separate the contributions arising from possible internal and external sources), as no current plans to send multiprobe missions to these planets are on schedule for the near future, and (2) predicting the vertical variation of observable species in brown dwarfs and extrasolar giant planets, as these atmospheres are unlikely to be in complete thermochemical equilibrium." }, "1003/1003.2148_arXiv.txt": { "abstract": "We construct general relativistic models of stationary, strongly magnetized neutron stars. The magnetic field configuration, obtained by solving the relativistic Grad-Shafranov equation, is a generalization of the {\\em twisted torus} model recently proposed in the literature; the stellar deformations induced by the magnetic field are computed by solving the perturbed Einstein's equations; stellar matter is modeled using realistic equations of state. We find that in these configurations the poloidal field dominates over the toroidal field and that, if the magnetic field is sufficiently strong during the first phases of the stellar life, it can produce large deformations. ", "introduction": "After the discovery of the Soft Gamma Repeaters and Anomalous X-ray Pulsars \\cite{M79,MS95}, a model of these sources was proposed according to which they are neutron stars with a very strong magnetic field; these {\\it magnetars} would have a surface field as large as $\\sim 10^{14}\\!-\\!10^{15}$ G, and internal fields about ten times larger \\cite{DT}. However, a clear picture of the structure, dynamics and evolution of magnetars is still missing. For instance we do not know whether the toroidal components of the field prevail on poloidal ones and how intense they are. Consequently, we do not know how large the deformation induced by the magnetic field on the star is, an information which is essential if one is interested in the gravitational wave emission of these sources. A deeper knowledge of the structure of strongly magnetized neutron stars would also help understanding various astrophysical processes involving magnetars (intense activity in the X- and gamma- spectra, quasi-periodic oscillations, eventually gamma-ray bursts). In a recent paper \\cite{C09}, to be referred to as Paper I hereafter, we constructed stationary models of non rotating neutron stars endowed with a strong magnetic field, in the framework of General Relativity (GR). In these models the poloidal field extends throughout the star and in the exterior, whereas the toroidal field is confined into a torus-shaped region inside the star, where the field lines are closed. It is worth reminding that these {\\it twisted torus} configurations have been found to be a quite general outcome of dynamical simulations of the early evolution of magnetized stars, in the framework of Newtonian gravity. Furthermore, due to magnetic helicity conservation, they appear to be stable on dynamical time-scales \\cite{BS0,BN,BS}, are not significantly affected by rotation \\cite{Y06}, and do not depend on the initial angle between the rotation and the magnetic axes \\cite{GR06}. In Lander \\& Jones \\shortcite{LJ} twisted torus configurations were studied in Newtonian gravity; the maximal relative strenght of the toroidal and poloidal components and the induced stellar deformation were evaluated using a polytropic equation of state (EOS) to model neutron star matter. In Paper I, we studied the twisted torus configurations in GR using a more realistic EOS. We considered a relation between the poloidal and the toroidal components of the field which is linear in the flux function, and estimated their ratio by determining the configuration of minimal energy at fixed magnetic helicity, under the assumption that the contribution of the $l>1$ multipoles is minimum outside the star. In the present paper, we reconsider the above assumptions: the higher multipoles contribution is not assumed {\\it a priori} to be minimum outside the star, and we allow for a more general parametrization of the relation between toroidal and poloidal fields. We determine the configuration of minimal energy at fixed magnetic helicity, and evaluate the stellar deformations induced by the twisted torus field by solving the perturbed Einstein equations including all relevant higher order multipoles. As in Paper I, the magnetized fluid is described in the framework of ideal MHD, which is accurate only in the first few hours of the star life, when the crust is still liquid and the matter in the core has not yet undergone a phase transition to the superfluid state. Since the characteristic Alfv\\`en time is of the order of $\\sim 0.01\\!-\\!10$ s, the magnetized fluid could reach a stationary state while the matter is still liquid and not yet superfluid\\footnote{We remark that even in presence of a stable stratification of the chemical composition, a magnetic field as strong as $B\\gtrsim 10^{15}$ is still allowed to evolve throughout the star on a dynamical time-scale \\cite{TM01}.}. The magnetic field induces quadrupolar deformation on the stellar shape and, as we shall later show, for magnetic fields as high as those observed in magnetars this deformation would be large; when the crust forms, it would maintain this deformed shape. The magnetic field would subsequently evolve on time-scales of the order of $\\sim10^3\\!-\\!10^5$ years due to dissipative effects like ohmic decay, ambipolar diffusion and Hall drift \\cite{GR92,WT,PG07}. The structure of the paper is as follows. In Section \\ref{TTMFC} we generalize the twisted torus magnetic field configurations introduced in Paper I by dropping the assumption that the contribution from the $l>1$ multipoles outside the star is minimum, and using a more general parametrization of the functional relation between toroidal and poloidal fields. In Section \\ref{ell} we determine the stellar deformations induced by the magnetic field. In Section \\ref{conclusions} we draw our conclusions. ", "conclusions": "In this paper we construct relativistic models of non-rotating, stationary stars, with a twisted torus magnetic field configuration. We extend the work done in Paper I, by removing the assumption of minimal contribution from multipoles higher than $l=1$, by considering more general forms of the function $\\beta(\\psi)$ which describes the ratio between toroidal and poloidal components, and by evaluating the deformation that the magnetic field induces on the star. We find that the non-minimal contribution of the $l>1$ multipoles, and the more general parametrization of the function $\\beta$, yield some interesting differences with respect to the magnetic field configurations found in Paper I: the new configurations have a much smaller poloidal field near the symmetry axis, and a larger toroidal field near the stellar surface. In any event, the toroidal field never contributes to more than $\\sim13\\%$ of the total magnetic energy stored inside the star. Since the poloidal field always prevails, in the twisted torus configurations the quadrupole ellipticity of the star $\\varepsilon_Q$ is always positive, and its maximum value is obtained in the purely poloidal limit. As shown by eq. (\\ref{finalformula}), which summarizes our results on the stellar deformation, $\\varepsilon_Q$ depends on the equation of state of matter: less compact stars can have larger deformations. We remark that the ellipticities given by eq. (\\ref{finalformula}) are larger than the bounds derived by evaluating the maximal strain that the crust can sustain \\cite{UCB00,Haskell06}. These bounds do not apply to the case we study, because in our case the magnetic field is assumed to reach a stationary configuration during the first few seconds of the neutron star life, when the star is still fluid and no crust has formed yet; therefore, if the field is sufficiently strong, the deformation it induces can be large, and may persist as the star cools down and the crust freezes in a non-spherical shape (see also the discussion in Haskell {\\it et al.} \\shortcite{Haskell}, Colaiuda {\\it et al.} \\shortcite{C08}). Recent results from the LIGO-Virgo collaboration \\cite{LIGO08,LIGO09} put an upper limit on the ellipticity of the Crab pulsar, which should be $\\varepsilon_Q\\lesssim 10^{-4}$. Our study indicates that strongly magnetized neutron stars with ellipticities of this order of magnitude may exist, provided this strong deformation was built up before the crust was formed, and the magnetic field was sufficiently strong. In order to further substantiate this scenario, an important issue which remains to be clarified is whether the twisted torus configurations we find are stable. This issue will be the subject of a future investigation." }, "1003/1003.2462_arXiv.txt": { "abstract": "Just as big bang nucleosynthesis allows us to probe the expansion rate when the temperature of the universe was around 1 MeV, the measurement of gravity waves from electroweak scale first order phase transitions may allow us to probe the expansion rate when the temperature of the universe was at the electroweak scale. We compute the simple transformation rule for the gravity wave spectrum under the scaling transformation of the Hubble expansion rate. We then apply this directly to the scenario of quintessence kination domination and show how gravity wave spectra would shift relative to LISA and BBO projected sensitivities. ", "introduction": "The detection of gravity waves (GWs) generated during a scalar sector's first order phase transition (PT) represents an interesting future possibility \\cite{Steinhardt:1981ct,Witten:1984rs,Kosowsky:1991ua,Kosowsky:1992rz,Kamionkowski:1993fg,Apreda:2001us,Kosowsky:2001xp,Dolgov:2002ra,Nicolis:2003tg,Caprini:2006jb,Grojean:2006bp,Randall:2006py,Huber:2007vva,Caprini:2007xq,Megevand:2008mg,Espinosa:2008kw,Huber:2008hg,Caprini:2009fx,Caprini:2009yp,Kusenko:2009cv,Megevand:2009ut,Ashoorioon:2009nf,Das:2009ue,Kahniashvili:2009mf,Kehayias:2009tn,Durrer:2010xc}. During a first order PT, bubbles of true vacuum nucleate, stir up the cosmological fluid, and collide, producing GWs. First order PTs at the electroweak scale in particular have received much attention because of their possible connection to a well motivated electroweak baryogenesis scenario in which the baryon asymmetry is generated during an electroweak phase transition (EWPT) \\cite{Shaposhnikov:1987tw,Kuzmin:1985mm}. Studies of this scenario are particularly timely given that Tevatron and the LHC are actively probing the Higgs sector responsible for electroweak symmetry breaking. If the Higgs boson or any degree of freedom that can be responsible for a first order PT at the electroweak scale is found in the ongoing experiments, future experiments may eventually be able to measure all of the parameters necessary to give an accurate prediction for the GWs. It is important to emphasize that even if the electroweak symmetry breaking is not a first order PT, a typical beyond the standard model scalar sector has multiple degrees of freedom, and some of these can undergo a first order PT at a temperature near the electroweak scale. Just as the relative isotope abundance measurements have led to a constraint on the expansion rate during big bang nucleosynthesis, the measurement of GWs may allow us to constrain any non-standard expansion rate during the time of EWPT. Several detailed computations of the gravity wave spectrum exist, and each give varying degrees of dependence on the Hubble expansion rate. However, to our knowledge, previous work does not sufficiently discuss the general dependence of the Hubble expansion rate to directly answer the following question: how would the observed gravity wave spectrum change if the Hubble expansion during the EWPT was changed from that of pure relativistic degrees of freedom? We compute a simple transformation rule for the gravitational wave spectrum in terms of $\\xi\\equiv H_{*}/H_{U}$, where $H_{U}$ is the expansion rate which assumes radiation domination and $H_{*}$ is the actual expansion rate:\\begin{equation} \\frac{d\\rho_{GW}(k)}{d\\ln k}\\rightarrow\\frac{1}{\\xi^{2}}\\frac{d\\rho_{GW}(k/\\xi)}{d\\ln k}.\\label{eq:main-result}\\end{equation} where $\\frac{d\\rho_{GW}(k)}{d\\ln k}$ is the spectrum computed assuming radiation domination. This immediately implies the following: 1) The peak frequency of the spectrum will shift from the standard scenario frequency $f_{p}$ as $f_{p}\\rightarrow\\xi f_{p}$, and 2) the peak amplitude of the spectrum will be suppressed from the standard scenario amplitude $\\mathcal{A}_{p}$ as $\\mathcal{A}_{p}\\rightarrow\\mathcal{A}_{p}/\\xi^{2}$. The intuition for the amplitude is that less source contributes to the gravitational wave at a typical spacetime point today because of the smaller intersection of the past null boundary with the approximately compact time support of the source. The intuition for the frequency shift is that all conformal symmetry breaking scales relevant for the observable frequency range is controlled by the Hubble expansion rate. We then apply this scaling relationship to the results of \\cite{Huber:2008hg,Caprini:2007xq} to compute how the gravity wave spectrum will transform due to the assumption of the existence of a quintessence kination dominated phase \\cite{Caldwell:1997ii,Peebles:1987ek,Kamionkowski:1990ni,Salati:2002md,Rosati:2003yw,Profumo:2003hq,Pallis:2005hm}. Such assumptions are interesting because as pointed out by \\cite{Salati:2002md} (related scenarios were also suggested before by \\cite{Kamionkowski:1990ni,Barrow:1982ei}), the freeze-out abundance of thermal relics can be strongly enhanced in scenarios in which the energy density is dominated by the kinetic energy of the quintessence field (kination domination) during the time of freeze-out, but dilutes away by the time of big bang nucleosynthesis (BBN). Such kination dominated freeze-out scenarios are then consistent with standard cosmology and predict that the standard relic abundance computed from the parameters extracted from collider measurements will be mismatched from the relic abundance deduced by observational cosmology. Thus, this scenario has interesting implications for physics models with thermal dark matter candidates (e.g. models with low energy supersymmetry such as the minimal supersymmetric extension of the SM (MSSM), technicolor models, models with large/warped extra dimensions, or certain classes of little Higgs models), which will be probed at the LHC and other experiments in the foreseeable future (for collider implications of this class of models, see for example \\cite{Chung:2007cn}). Furthermore, as pointed out by \\cite{Chung:2007vz,Profumo:2004ty}, large annihilation cross sections compatible with the dark matter explanation of the excess positrons \\cite{Adriani:2008zr,Beatty:2004cy} can be compatible with the right thermal relic abundance since the effective boost factor coming from the kination scenario can easily be as large as $10^{3}$. Finally, since the measurement of CMB B-mode polarization can almost model independently falsify this scenario, this scenario can be nontrivially checked with a variety of probes. In particular, the EWPT gravity wave probe in conjunction with dark matter cosmology can represent a smoking gun probe of the scenario if the gravity wave is measurable \\cite{Chung:2007vz} and colliders can eventually measure the requisite short-distance parameters with sufficient accuracy. The order of presentation is as follows. In the next section, we present the main analytic result of this paper. Sec.~\\ref{sec:survey} focuses on checking explicit consistency of our result with the existing literature on explicit gravity wave spectrum computations. In Sec.~\\ref{sec:kinationdominationexample}, we apply the transformation relation to the quintessential kination scenario and give plots showing how the transformed spectra look relative to the projected sensitivities of LISA and BBO. Sec~\\ref{sec:caveat} discusses all of the caveats associated with the analytic result. We then conclude with a summary. The appendices contain some of the details used throughout the paper. As far as conventions are concerned, we use the reduced Planck's constant $M_{p}=2.4\\times10^{18}$ GeV and also assume a flat FRW background metric $ds^{2}=a^{2}(t)(dt^{2}-|d\\vec{x}|^{2}).$ ", "conclusions": "Summary} In this paper, we have presented an analytic transformation rule Eq.~(\\ref{eq:spectrum shift}) that is useful for understanding how the gravitational wave spectrum generated through an electroweak scale first order phase transition at a fixed temperature $T_*$ would change if the expansion rate of the universe during the phase transition were different from that inferred from the assumption of pure radiation domination. We have explored the remarkable robustness of the scaling relationship with respect to many computational uncertainties in the gravity wave spectrum. We apply this transformation rule to the example of a universe having a quintessential kination dominated phase and find as expected a strong sensitivity to the single phenomenological parameter controlling this scenario. In principle, this scaling relationship together with dark matter properties measured by colliders can be used to overconstrain this single phenomenological parameter. Unfortunately, if the current technology of gravity wave computations is correct, then we find that any measurement of the gravity wave spectrum at the level of the projected BBO sensitivity can rule out any appreciable boost factors relevant for reconciliations between various sets of data such as that between colliders and cosmology and/or indirect detection (such as that relevant \\cite{Chung:2007vz} for PAMELA data \\cite{Adriani:2008zr}). Nonetheless, using the results of this work, any future gravity wave detection experiments measuring phase transition induced gravity waves can understand their measurement's sensitivity to the expansion rate of the universe. It would indeed be exciting to have an observational anchor on the expansion rate of the universe when the universe is as hot as $100$ GeV, just as isotope abundance measurements allow us to have an observational anchor on the expansion rate at a temperature of $1$ MeV in the context of big bang nucleosynthesis." }, "1003/1003.0961.txt": { "abstract": "Spitzer data at 24, 70, and 160~$\\mu$m and ground--based H$\\alpha$ images are analyzed for a sample of 189 nearby star--forming and starburst galaxies to investigate whether reliable star formation rate (SFR) indicators can be defined using the monochromatic infrared dust emission centered at 70 and 160~$\\mu$m. We compare recently published recipes for SFR measures using combinations of the 24~$\\mu$m and observed H$\\alpha$ luminosities with those using 24~$\\mu$m luminosity alone. From these comparisons, we derive a reference SFR indicator for use in our analysis. Linear correlations between SFR and the 70~$\\mu$m and 160~$\\mu$m luminosity are found for L(70)$\\gtrsim$1.4$\\times$10$^{42}$~erg~s$^{-1}$ and L(160)$\\gtrsim$2$\\times$10$^{42}$~erg~s$^{-1}$, corresponding to SFR$\\gtrsim$0.1--0.3~M$_{\\odot}$~yr$^{-1}$, and calibrations of SFRs based on L(70) and L(160) are proposed. Below those two luminosity limits, the relation between SFR and 70~$\\mu$m (160~$\\mu$m) luminosity is non--linear and SFR calibrations become problematic. A more important limitation is the dispersion of the data around the mean trend, which increases for increasing wavelength. The scatter of the 70~$\\mu$m (160~$\\mu$m) data around the mean is about 25\\% (factor $\\sim$2) larger than the scatter of the 24~$\\mu$m data. We interpret this increasing dispersion as an effect of the increasing contribution to the infrared emission of dust heated by stellar populations not associated with the current star formation. Thus, the 70 (160) $\\mu$m luminosity can be reliably used to trace SFRs in large galaxy samples, but will be of limited utility for individual objects, with the exception of infrared--dominated galaxies. The non--linear relation between SFR and the 70 and 160~$\\mu$m emission at faint galaxy luminosities suggests a variety of mechanisms affecting the infrared emission for decreasing luminosity, such as increasing transparency of the interstellar medium, decreasing effective dust temperature, and decreasing filling factor of star forming regions across the galaxy. In all cases, the calibrations hold for galaxies with oxygen abundance higher than roughly 12$+Log$(O/H)$\\sim$8.1. At lower metallicity the infrared luminosity no longer reliably traces the SFR because galaxies are less dusty and more transparent. ", "introduction": "The star formation rate (SFR) is one of the principal parameters that needs to be measured in star--forming regions and galaxies, in order to characterize their evolution. Extensive efforts have been made over the past couple of decades to derive SFR indicators from luminosities at a variety of wavelengths, spanning from the UV, where the recently formed massive stars emit the bulk of their energy, to the infrared, where the dust-reprocessed light from those stars emerges, to the radio, which is mostly a tracer of supernova activity \\citep[e.g.,][for earlier papers, see \\citet{Kennicutt1998}]{Kennicutt1998,Yun2001, Kewley2002, Ranalli2003, Hirashita2003, Bell2003, Kewley2004, Calzetti2005, Schmitt2006, Moustakasal2006, AlonsoHerrero2006, Calzetti2007, Salim2007, Persic2007, Rosa2007, Kennicutt2007, Bigiel2008, Rieke2009, Calzetti2009}. In recent years, investigations of monochromatic (i.e., based on a single band measurement) SFR indicators based on the infrared emission from galaxies have experienced a new resurgence, thanks to the high--sensitivity and high--angular resolution data provided by the Spitzer Space Telescope, which have yielded both deeper distant galaxy surveys and more accurate information on the relation between dust and stellar emission in nearby galaxies. Deep surveys are often characterized by limited information across the infrared wavelength range and monochromatic SFRs are an important tool for these projects. The rest--frame mid--infrared emission from dust in galaxies, in particular the emission detected in the 8~$\\mu$m and 24~$\\mu$m Spitzer bands, has been analyzed by a number of authors \\citep[][]{Roussel2001,Forster2004,Boselli2004,Calzetti2005, Wu2005, AlonsoHerrero2006, PerezGonzalez2006, Rellano2007, Calzetti2007, Zhu2008, Rieke2009, Salim2009}, and a general correlation (but also a number of caveats) between mid--IR infrared emission and SFR has been found. Over the next few years, new facilities, both from space (e.g., the Herschel Space Telescope) and from the ground (ALMA and the Large Millimeter Telescope, to mention just two) will open new windows of sensitivity at even longer wavelengths than those explored by Spitzer, and will, in turn, provide even more powerful tools for probing the evolution of the rate at which galaxies have assembled their %luminous gas and dust components. Deep surveys will be able to probe the dust emission from galaxies at rest--frame infrared wavelengths that are close to the peak emission, and to the bulk of the infrared energy budget. Herschel will, for instance, enable us to probe the peak dust emission from galaxies ($\\sim$60--150~$\\mu$m) up to redshift z$\\sim$2, while the Rayleigh--Jeans tail of the dust emission will be the dominion of sub--millimeter and millimeter facilities. Exploring the viability, and limitations, of using the monochromatic emission close to the infrared peak and at longer wavelengths as SFR indicators is thus timely for providing a reference for those future surveys. In this paper we investigate the use of the 70~$\\mu$m and 160~$\\mu$m emission, the two longest wavelength Spitzer bands, from nearby galaxies as SFR diagnostics. The present paper is organized as follows: Section~2 introduces the sample of local star--forming galaxies and the data used for the present analysis; Section~3 describes the quantities used in this work; Section~4 compares a variety of existing SFR indicators in the optical/infrared, to derive `reference' SFRs for our galaxies, which are then compared with the Spitzer 70~$\\mu$m and 160~$\\mu$m emission from the same galaxies and with expectations from models in Sections~5 and 6, respectively. A discussion of the results and the conclusions are given in Section~7. Throughout the paper, we adopt a value of the Hubble constant H$_0$=70~km~s$^{-1}$~Mpc$^{-1}$. ", "conclusions": "Monochromatic SFR indicators have the undeniable advantage of immediate application, especially when observational data cover only a limited range of wavelengths, as is often the case for distant galaxies. Within this framework, Spitzer observations of nearby galaxies are offering a unique opportunity to test the applicability and limitations of a number of SFR indicators based on the emission from the dust heated by stars. The main goal of this paper has been to investigate whether long wavelength infrared fluxes at 70 and 160~$\\mu$m can be reasonably calibrated as SFR indicators. Previous works \\citep{Roussel2001,Boselli2004,Forster2004,Calzetti2005,Dale2005,Wu2005,AlonsoHerrero2006,PerezGonzalez2006,Calzetti2007,Dale2007,Rellano2007,Zhu2008,Rieke2009} have concentrated on the short--wavelength infrared, in the range $\\sim$5--40~$\\mu$m, commonly referred to as the mid--IR. Although this wavelength range is below the peak infrared emission and only accounts for a few to a few tens of percent of the total IR emission, the dust heated by the hot, massive stars in a recent star formation episode can have high effective temperatures and may preferentially emit at the shorter infrared wavelengths. This general picture translates into a correlation between SFRs and luminosities in the mid--IR, such as L(8) and L(24), centered on the Spitzer 8~$\\mu$m and 24~$\\mu$m bands. However, the correlation is also non--linear, and careful analysis shows that other contributors, in addition to the SFR, determine the short infrared wavelength luminosity, including transiently heated dust by single UV or optical photons in the general radiation field of a galaxy \\citep{Haas2002,Boselli2004,Bendo2008}, variations in the dust effective temperature \\citep[for L(24), see][]{Calzetti2007}, and, especially for L(8), gas metallicity \\citep{Engelbracht2005,Rosenberg2006,Wu2006,Madden2006,Jackson2006,Draine2007}. \\subsection{Establishing a reference SFR indicator} In the present paper, we first establish a ``reference'' SFR calibration against which we can compare the 70 and 160~$\\mu$m luminosities and LSDs. The adopted reference SFR is given in equation~17. We choose this particular three--part formulation after comparing a variety of published 24~$\\mu$m and composite H$\\alpha$ plus 24~$\\mu$m calibrations. Above a luminosity L(24)$\\sim$5$\\times$10$^{43}~$erg~s$^{-1}$, the contribution of direct (unabsorbed by dust) massive stars emission to SFR measurements is less than $\\sim$15\\% (Figure~\\ref{fig8}). % in this regime, many calibrations using L(24) to %measure SFRs converge to common or similar mean values (Figure~\\ref{fig6}). Thus, For L(24)$\\gtrsim$5$\\times$10$^{43}~$erg~s$^{-1}$, the 24~$\\mu$m emission is a reliable SFR indicator {\\em for whole galaxies}. However, in this luminosity regime, the 24~$\\mu$m dust emission is also self--absorbed \\citep{Rieke2009}, and a non--linear correction to the calibration needs to be applied \\citep[third term of equation~17, re--calibrated from the original calibration of][]{Rieke2009}. Below a luminosity L(24)$\\sim$5$\\times$10$^{43}~$erg~s$^{-1}$, the galaxies become transparent in the mid--infrared, and, as the luminosity further decreases, they become transparent at UV and optical wavelengths \\citep[e.g.][]{Bell2003,Buat2007}. At L(24)$\\sim$4$\\times$10$^{42}~$erg~s$^{-1}$, the contribution to the SFR from emission at optical (H$\\alpha$) and IR (24~$\\mu$m) wavelengths is roughly equal (Figure~\\ref{fig7}). Thus, below L(24)$\\sim$5$\\times$10$^{43}~$erg~s$^{-1}$, the 24~$\\mu$m luminosity becomes increasingly insufficient, by itself, to fully characterize the SFR of a whole galaxy, as the portion of the star formation unabsorbed by dust becomes an increasingly significant contribution to the total SFR. Additionally, a correlation between the effective temperature and the luminosity/area of the thermal dust emission is present in HII~regions for decreasing total SFRs \\citep{Calzetti2007}, and in whole galaxies as well \\citep[][and Figure~\\ref{fig4}]{Chanial2007}. This correlation is further seen for the galaxies in our sample by plotting the ratio L(70)/L(24) as a function of $\\Sigma_{24}$ (Figure~\\ref{fig17}, left); the star--forming and starburst galaxies show a decreasing IR ratio for increasing LSD, as is expected for increasing effective dust temperature. This trend is less tight, albeit still present to some extent, when the L(70)/L(24) ratio is plotted as a function of luminosity (Figure~\\ref{fig17}, right), rather than luminosity/area. The use of total luminosity introduces a scatter because galaxies are weighted not only by their intrinsic level of activity but also by their size (a low--activity, low IR/area, but large, galaxy can have the same total luminosity as a high--activity, but small, galaxy). The relation between the effective temperature and the luminosity/area is present in galaxies up to an IRSD $\\Sigma_{TIR}\\sim$0.3--1$\\times$10$^{43}$~erg~s$^{-1}$~kpc$^{-2}$ (L(TIR)$\\approx$1--3$\\times$10$^{44}$~erg~s$^{-1}$$\\sim$3--8$\\times$10$^{10}$~L$_{\\odot}$ or L(24)$\\approx$2--6$\\times$10$^{43}$~erg~s$^{-1}$$\\sim$0.6--2$\\times$10$^{10}$~L$_{\\odot}$), independent of the metal content of the galaxy (Figures~\\ref{fig3}, left, and \\ref{fig17}). This can be further seen in Figure~\\ref{fig18}, where we plot L(24)/L(TIR) as a function of IRSD only for those galaxies in our sample that have oxygen abundance 12$+$Log(O/H)$>$8.5. There are 66 normal star--forming and starburst galaxies above this oxygen abundance value (or about 1/2 of the normal star--forming and starburst galaxies in our high--metallicity sample), and the 29 LIRGS. A similar trend to that of Figure~\\ref{fig4} of L(24)/L(TIR) as a function of IRSD is observed for this restricted, higher--metallicity, sample, implying that the effect is not primarily driven by the abundance of dust in the galaxy. The combination of the two effects discussed above, increasing transparency of the galaxy and decreasing effective dust temperature for decreasing luminosity/area, may account for why different published calibrations for SFR(24) diverge when extrapolated towards the low luminosity end of our sample (Figure~\\ref{fig6}). A more robust approach appears to involve measuring SFRs by combining the 24~$\\mu$m emission and the observed H$\\alpha$ emission, the first probing the dust--absorbed star formation and the second the unabsorbed one \\citep{Kennicutt2007,Calzetti2007,Kennicutt2009}. Even in this case, though, a transition in the proportionality factor between L(24) and L(H$\\alpha$)$_{obs}$ needs to be implemented for estimating accurate SFRs (equation~17, first and second terms). The transition point is marked by the condition L(H$\\alpha$)$_{obs}$/(0.020 L(24))$\\approx$1, which corresponds, for decreasing luminosity, to the transition from galaxies dominated by dust--obscured star formation to those dominated by unobscured star formation \\citep[e.g.,][]{Buat2007}. The transition corresponds to a change of roughly 50\\% in the fraction of L(24) to be added to L(H$\\alpha$)$_{obs}$, from 0.031 to 0.020 for L(24)$\\lesssim$4$\\times$10$^{42}~$erg~s$^{-1}$ (first and second relation in equation~17), or SFR$\\sim$1~M$_{\\odot}$~yr$^{-1}$. Our results thus indicate that around this SFR, the stellar population that dominates the heating of the dust transitions from being a relatively young one (constant star formation over the past $\\sim$100~Myr), typical of starburst events, to a much more evolved one (constant star formation over the past $\\sim$10~Gyr, or exponentially decreasing star formation), typical of widespread, but isolated, star formation amid an otherwise quiescent galaxy. In the latter case, the stellar populations not associated with the current star formation events contribute, in a proportionally larger fraction than the starburst case, to the heating of the 24~$\\mu$m--emitting dust, whose emission contribution to the SFR calibration needs then to be accordingly removed, on average, by reducing L(24) via a smaller multiplicative factor \\citep{Kennicutt2009} than the starburst or HII~region case \\citep{Calzetti2007}. \\subsection{The Emission at 70 and 160~$\\mu$m as SFR Indicators} Our reference SFR calibration (equation~17) enables the testing of other IR monochromatic luminosities as potential SFR indicators, over a wide range of luminosity surface densities. In this work, we have investigated the applicability of L(70) and L(160), i.e., the long wavelength Spitzer bands, as such indicators. These two wavelength regions are attractive because each contains a larger fraction of the total infrared emission than the 24~$\\mu$m emission. L(70) is 50\\% of L(TIR) for infrared LSD $\\Sigma_{TIR}\\gtrsim$10$^{41}$~erg~s$^{-1}$~kpc$^{-2}$, and L(160) is about 20\\% or larger fraction of L(TIR) for the entire luminosity range investigated in this work (Figure~\\ref{fig5}), while L(24) only contains $\\sim$10--20\\% of the total infrared emission (Figure~\\ref{fig4}). At the high luminosity LIRGs regime, both L(70) and L(160) represent a larger fraction of L(TIR) than expected from models \\citep[e.g.,][]{Draine2007}. This indicates a `cooler' (larger effective temperature) IR SED for those systems than predicted from models, an effect already previously observed \\citep{Dunne2001}, and in line with the re--emission of the energy absorbed in the mid--IR wavelength range. Galaxies in the LIRG luminosity regime (L(TIR)$\\gtrsim$1--3$\\times$10$^{44}$~erg~s$^{-1}$ or $\\Sigma_{TIR}\\gtrsim$0.3--1$\\times$10$^{43}$~erg~s$^{-1}$~kpc$^{-2}$) have their energy output dominated by a star formation event that is evolving in a dusty region of sufficient optical depth that even light emerging in the mid--IR gets absorbed, and that light is re--emitted at longer infrared wavelengths, thus producing an overall cool IR SED, specifically cooler than expected from models where the emission is emerging from a transparent (at IR wavelengths) medium. This requires that the dust attenuation at 24~$\\mu$m be substantial, $\\approx$1~mag, which translates to A$_V\\sim$15--50~mag, depending on the extinction curve adopted \\citep{DraineLi2007, Flaherty2007}. Such large extinctions are not uncommon in bright infrared galaxies. \\citet{Genzel1998} determines that Ultraluminous InfraRed Galaxies (ULIRGs) have extinction values in the range A$_{V,screen}$=5--50~mag or A$_{V,mixed}$=50--1000~mag. Although LIRGs are typically an order of magnitude less luminous than ULIRGs, dust extinctions are still expected to be large in these systems. The 70~$\\mu$m emission from galaxies correlates linearly with the SFR for luminosities L(70))$\\gtrsim$1.4$\\times$10$^{42}$~erg~s$^{-1}$, and a calibration is given in equation~22. Similarly, L(160) is linearly correlated with the SFR for luminosities L(160)$>$10$^{42}$~erg~s$^{-1}$ (equation~25). \\subsection{Scatter and the Impact of Metallicity} Both calibrations, however, have significant scatter, $\\sigma_{70}\\sim$0.2~dex and $\\sigma_{160}\\sim$0.4~dex, around the mean trend, and larger than the scatter, $\\sim$0.12--0.16~dex, of calibrations based on L(24) or a mix of L(24) and L(H$\\alpha$)$_{obs}$ \\citep{Rieke2009,Kennicutt2009}. Furthermore, the scatter increases for increasing wavelength, and appears to decrease for increasing luminosity (Figures~\\ref{fig9} and \\ref{fig15}). To further investigate these trends, we subdivide the high--metallicity sample into two sub--samples of oxygen abundance higher and lower than 12$+$Log(O/H)$=$8.5. This separating value is chosen after Figure~5 of \\citet{Engelbracht2008}, which shows that galaxies above a metallicity of 12$+$Log(O/H)$=$8.5 tend to have a roughly constant f$_{\\nu}$(70)/f$_{\\nu}$(160) flux ratio (albeit with a large scatter), while galaxies below that metallicity value tend to have an increasing ratio for decreasing oxygen abundance. Figure~\\ref{fig19} shows the high--metallicity data from our sample divided into the two sub--samples. The scatter of the data for both $\\Sigma_{70}$ and $\\Sigma_{160}$ at constant $\\Sigma_{SFR}$ decreases when the sample is restricted to the highest metallicity data (12$+$Log(O/H)$>$8.5); specifically, faint $\\Sigma_{70}$ and $\\Sigma_{160}$ galaxies at fixed $\\Sigma_{SFR}$ drop from the sub--sample. This is in the expected direction, as seen earlier in this paper when analyzing the impact of low--metallicity galaxies (Figure~\\ref{fig12}). The results of Figure~\\ref{fig19} should be considered preliminary at this stage, since the oxygen abundances for a number of high metallicity galaxies are uncertain (sometimes by about 0.2~dex). Nevertheless, they confirm and support the results of section~5 for SFR(70), in the sense that as the galaxy metallicity increases and its dispersion in the sample decreases, $\\Sigma_{70}$ tends to approach a 1--to--1 relation with $\\Sigma_{SFR}$ (Figure~\\ref{fig19}, left). Conversely, in the same metallicity range, $\\Sigma_{160}$ tends to show a flattening in its relation with $\\Sigma_{SFR}$ and the dispersion of the datapoints around the mean trend remains large (Figure~\\ref{fig19}, right). As already discussed in section~6, this is in agreement with the expectation that the emission at long infrared wavelengths receives an increasing contribution (and an increasing scatter) from dust heated by evolved stellar populations unassociated with the current star formation event \\citep{Lonsdale1987}. Figure~\\ref{fig19} (right) further stresses that the calibration of SFR(160), equation~25, should only be used for L(160)$\\gtrsim$2$\\times$10$^{42}$~erg~s$^{-1}$, equivalent to $\\Sigma$(160)$\\gtrsim$10$^{41}$~erg~s$^{-1}$~kpc$^{-2}$. The calibrations just quoted apply to galaxies with mean gas--phase metallicity 12$+Log$(O/H)$\\gtrsim$8.1. The scatter becomes even larger, and strongly asymmetric, when galaxies of lower oxygen abundance value are included (Figure~\\ref{fig12}). Indeed, the general expectation is that lower metallicity galaxies will in general contain less dust and be more transparent than high metallicity ones. In this case, the infrared will be a poor tracer of the current SFR, as most of the light produced by recently formed stars will emerge unabsorbed by dust. Thus, calibrations of SFR(70) and SFR(160) need to be applied with an awareness of these limitations. A hybrid approach, which combines, e.g., L(70) with L(H$\\alpha$)$_{obs}$, could likely provide a `remedy' for the dependence of far--infrared SFR indicators on metallicity, similarly to what has been already verified for L(24) and L(8), where the H$\\alpha$ observed luminosity accounts for the unabsorbed portion of the SFR \\citep{Calzetti2007, Kennicutt2007, Zhu2008, Kennicutt2009}. However, our current data do not enable us to test such hybrid indicator. We already use, in our analysis, the combination of L(24) and L(H$\\alpha$)$_{obs}$ to trace SFR in galaxies in an unbiased fashion; this indicator is dominated by the H$\\alpha$ luminosity at low SFSDs, as would be any hybrid combination of L(70) and L(H$\\alpha$)$_{obs}$, thus resulting in a degeneracy for any test we may attempt. \\subsection{Limits of Validity for the Calibrations} For the high-metallicity galaxies, the presence of a contribution to the long wavelength infrared emission from dust heated by populations unassociated with the current star formation becomes more evident for decreasing galaxy luminosity (Figures~\\ref{fig11} and \\ref{fig15}), but it also competes with at least three more effects: (1) the effective dust temperature as measured from the IR SED tends to decrease (Figures~\\ref{fig4}, \\ref{fig5}, and \\ref{fig17}); (2) the overall interstellar medium becomes progressively more transparent, thus a lower fraction of the SFR is traced in the infrared; (3) the area filling factor of star forming regions within galaxies decreases down to $\\sim$10\\%--50\\% (depending on the dust extinction and stellar population models adopted), contributing to a change in the proportionality between $\\Sigma_{70}$ or $\\Sigma_{160}$ with $\\Sigma_{SFR}$ relative to that of more luminous galaxies. These four effects all contribute, at various levels, to change the relation between $\\Sigma_{70}$ (or $\\Sigma_{160}$) with $\\Sigma_{SFR}$ from linear to non--linear (e.g., equation~21), thus complicating any SFR calibration for luminosities below L(70)$\\sim$1.4$\\times$10$^{42}$~erg~s$^{-1}$ (L(160)$\\sim$2$\\times$10$^{42}$~erg~s$^{-1}$). These luminosity values correspond to SFR$\\sim$0.1--0.3~M$_{\\odot}$~yr$^{-1}$, thus they are not extremely restrictive especially in the context of current studies of distant galaxy populations. \\subsection{Additional Uncertainties} So far, we have not considered another potential source of heating for the dust associated with the diffuse interstellar medium of galaxies: photons leaking out of HII regions. About 20\\%--50\\% of the integrated H$\\alpha$ luminosity in galaxies is associated with the diffuse interstellar medium \\citep{Reynolds1990,Ferguson1996,Wang1997,Martin1997}, and those photons will produce dust heating. Since photons leaking out of HII regions are associated with the current star formation, in our analysis we have assumed that the diffuse ionized gas follows the same scaling relations as the star--forming regions, albeit possibly with different absolute values (e.g., the diffuse photons may heat the dust to typically lower temperatures than those spatially associated with HII regions). This assumption is justified by the findings of \\citet{Wang1998}, according to which the fraction of diffuse-to-total ionized gas remains relatively constant in galaxies of increasing SFR, from normal star--forming to starbursts, once the diffuse gas surface brightness is normalized by the mean star formation rate per unit area of the galaxy. The infrared emission from the galaxies in our sample is dominated by dust heated by current star formation, with, in some cases, a non--negligible contribution from evolved stellar populations. Central AGNs do not constitute a dominant heating source in the present sample, and their impact on our calibrations could not be investigated. Thus, the application of equations~22 (or 25) to galaxy populations first requires the evaluation of any contamination from dust heating by AGNs. Finally, all our derivations are predicated on the assumption that a universal stellar IMF can be applied to all galaxies. While this is generally considered a reasonable assumption, recent investigations using the Sloan Digital Sky Survey galaxy sample \\citep{Hoversten2008}, a complete sample of nearby galaxies \\citep{Lee2009b}, and HI--selected galaxies \\citep{Meurer2009} suggest that the galaxy--integrated IMF may become increasingly bottom--heavy in low luminosity and/or low surface brightness galaxies. Changes in the IMF will clearly impact the calibration of SFR indicators across the full electromagnetic wavelength range, and establishing whether such variations are present and what physical parameters may be causing them is an avenue of future investigation." }, "1003/1003.3134_arXiv.txt": { "abstract": "{} {We have studied the molecular content of the circumstellar environs of symbiotic stellar systems, in particular of the well know objects R Aqr and CH Cyg. The study of molecules in these stars will help to understand the properties of the very inner shells around the cool stellar component, from which molecular emission is expected to come. } {We have performed mm-wave observations with the IRAM 30m telescope of the \\doce\\ \\juc\\ and \\jdu, \\trece\\ \\juc\\ and \\jdu, and SiO \\jcc\\ transitions in the symbiotic stars R Aqr, CH Cyg, and HM Sge. The data were analyzed by means of a simple analytical description of the general properties of molecular emission from the inner shells around the cool star. Numerical calculations of the expected line profiles, taking into account the level population and radiative transfer under such conditions, were also performed. } {Weak emission of \\doce\\ \\juc\\ and \\jdu\\ was detected in R Aqr and CH Cyg; a good line profile of \\doce\\ \\jdu\\ in R Aqr was obtained. The intensities and profile shapes of the detected lines are compatible with emission coming from a very small shell around the Mira-type star, with a radius comparable to or slightly smaller than the distance to the hot dwarf companion, 10$^{14}$ -- 2\\,$\\times$\\,10$^{14}$\\,cm. We argue that other possible explanations are improbable. This region probably shows properties similar to those characteristic of the inner shells around standard AGB stars: outwards expansion at about 5 -- 25 \\kms, with a significant acceleration of the gas, temperatures decreasing with radius between about 1000 and 500 K, and densities $\\sim$ 10$^9$ -- 3\\,$\\times$\\,10$^8$ cm$^{-3}$. Our model calculations are able to explain the asymmetric line shape observed in \\doce\\ \\jdu\\ from R Aqr, in which the relatively weaker red part of the profile would result from selfabsorption by the outer layers (in the presence of a velocity increase and a temperature decrease with radius). The mass-loss rates are somewhat larger than in standard AGB stars, as often happens for symbiotic systems. In R Aqr, we find that the total mass of the CO emitting region is $\\sim$ 2 -- 3 $\\times$ 10$^{-5}$ \\ms, corresponding to \\mloss\\ $\\sim$ 5 $\\times$ 10$^{-6}$ -- 10$^{-5}$ \\my, and compatible with results obtained from dust emission. Taking into account other existing data on molecular emission, we suggest that the small extent of the molecule-rich gas in symbiotic systems is mainly due to molecule photodissociation by the radiation of the hot dwarf star. } {} ", "introduction": "Symbiotic stellar systems (SSs) are very close binary systems composed by a cold red giant and a very hot dwarf companion. The strong interaction between both stars yields a number of interesting and sometimes spectacular phenomena, like the ejection of fast collimated flows and the formation of high-excitation, bipolar nebulae \\citep[see e.g.][and references therein]{corradi03}. They are thus a very attractive laboratory for studying various aspects of stellar evolution in binary systems and of the circumstellar chemistry and structure under these (extreme) conditions. Some symbiotic systems show strong IR excesses, due to the emission of dust grains formed in the circumstellar envelopes (CSEs) ejected by the cold primary. In these objects, the ejection of circumstellar gas from the primary is thought to be particularly important. However, molecules are abundant only in the innermost regions of the CSEs around SSs. Bands of CO, TiO, and other molecules are often detected in SSs, coming from the outer atmosphere of the cool primary, as in standard red giants. However, to our knowledge, molecular low- and intermediate-excitation emission (SiO masers, CO thermal lines, etc), known to come from the circumstellar shells, had been only detected in two SSs: R Aqr and H1-36 \\citep{ivison94,ivison98,schwarz95,seaquist95}. SiO maser emission (v$>$0 lines at mm wavelengths) in R Aqr is relatively 'normal', compared to that observed in other AGB stars \\citep[e.g.][]{pardo04,cotton04,kamohara10}. We recall that SiO maser emission comes from the innermost circumstellar regions, in our case at about 5 $\\times$ 10$^{13}$ cm. H$_2$O masers have been also detected in these same two objects, with characteristics that suggest that H$_2$O emission forms only at a few stellar radii \\citep{seaquist95,ivison94,ivison98}. In R Aqr, an H$_2$O-rich shell extending also $\\sim$ 5 $\\times$ 10$^{13}$ cm has been detected from observations of H$_2$O vibrational bands \\citep{ragland08}. Molecules characteristic of the outer shells are yet weaker and very rarely detected. OH masers have been well detected only in H1-36 \\citep{ivison94,seaquist95}. The CO thermal lines, which in general come from still outer shells (often within about 10$^{17}$ cm), are extremely weak in SSs. Previous to our work, only a very tentative detection of R Aqr had been reported by \\cite{groenewegen99}. In fact, our observations (see next sections) show that the feature detected in R Aqr by these authors is mostly due to baseline ripples in the spectrum, since the actual intensity of the line is about three to four times lower. It seems well established that the detectability of molecular emission in SSs increases for lines coming from very compact regions around the AGB star and for systems showing relatively large distances between the stars \\citep{schwarz95,ivison98}. Therefore, the lack of detections of molecular lines is probably due to photodissociation by the UV radiation from the hot companion or to dynamical disruption of the emitting regions (see further discussion in Sect.\\ 5). The properties of the innermost shells around red giants, within about 10 stellar radii, from which molecular lines in SSs seem to come, are not very well known even for isolated stars. Both theoretical and observational studies \\citep[e.g.][]{hinkle82,hofner98,andersen03,sandin08} suggest that relatively high temperatures, between 500 and 1000 K, and densities, 10$^8$ -- 10$^9$ cm$^{-3}$, are present. These inner shells are thought not to show the fast expansion characteristic of the outer regions. They are probably pulsating, due to shocks originated from the photospheric pulses, or show incipient expansion, since in these regions dust grains are being formed and radiation pressure efficiently acts onto them. In symbiotic systems, pulsation and outwards acceleration may also dominate the dynamics of the inner circumstellar layers, but the gravitational effects of the secondary cannot be neglected, since SSs are interacting systems. In particular, it is remarkable that SSs present mass-loss rates systematically larger than those of isolated AGB stars \\citep[e.g.][and references therein]{miko99}, which may be due to such gravitational effects. In this paper, we present observations of molecular mm-wave lines in three SSs: R Aqr, CH Cyg, and HM Sge. \\doce\\ emission is well detected in R Aqr and CH Cyg. R Aqr and CH Cyg are bright and nearby symbiotic stars, extensively studied over the whole spectral range. The distance to R Aqr has been accurately determined from recent VLBI measurements of its parallax by Kamohara et al.\\ (2010); we will adopt their distance value, $D$ = 214$^{+45}_{-32}$. A distance of $244^{+49}_{-35}$ pc was measured for CH Cyg from Hipparcos data \\citep{vanleu07}. In both systems, the hot component is an accreting white dwarf showing spectacular activity: irregular accretion-powered outbursts accompanied by massive outflows and jets (e.g.\\ Kellogg et al.\\ 2007, Karovska et al.\\ 2007, and references therein). The cool component in R Aqr is a Mira variable with a pulsation period of $387^{\\rm d}$, {whereas in CH Cyg it is an M7\\,III semiregular variable with a complex variability} \\citep[e.g.][and references therein]{gromadzki09,mikolajewski92}. Both have relatively well known orbital parameters \\citep{gromadzki09,hinkle09}. Aqr and CH Cyg have the longest orbital periods measured in well studied symbiotic systems, respectively 43.6 yr and 15.6 yr. The orbital solution for R Aqr \\citep{gromadzki09} implies that the average component separation is $\\sim$ 2.25 $\\times$ 10$^{14}$ cm (15 AU). During our observations the component separation was $\\sim$ 17.6 and 16.5 AU, in May 2008 and May 2009, respectively. In the case of CH Cyg, the orbital elements from Hinkle et al.\\ (2009) yield an average component separation of $\\sim$ 9 AU (1.35 $\\times$ 10$^{14}$ cm), and during the May 2009 observation the separation was $\\sim 9.8$ AU. The red giant radius in both systems is known from interferometric measurements, $\\sim 1.9$ AU in R Aqr (Gromadzki \\& Miko{\\l}ajewska 2009, and references therein), and $\\sim 1.2$ AU in CH Cyg (e.g. Dyck et al.\\ 1998), with uncertainty set mostly by the uncertainty in their distances, $\\sim 15\\, \\%$ in both cases. The symbiotic nova HM Sge is composed of a Mira variable with a pulsation period of 527$^{\\rm d}$, embedded in an optically thick dust shell, and a white dwarf companion slowly declining from a thermonuclear nova outburst started in 1975 (Belczynski et al.\\ 2000, and references therein; Muerset \\& Nussbaumer 1994). The orbital period is unknown, likely higher than $\\sim 100$ yr. Eyres et al.\\ (2001) measured the binary component positions using HST images, and estimated a projected binary separation of $40 \\pm 9$ mas, and a position angle of the binary axis of $130 \\pm 10$ degrees, in agreement with that suggested by Schmid et al.\\ (2000) based on spectropolarimetry. Unfortunately, the distance to HM Sge is rather uncertain -- published values range from 0.3 to 3.2 kpc (e.g.\\ Richards et al.\\ 1999). Eyres et al.\\ adopted $D$ = 1.25 kpc, and a component separation of 50 AU. {The recently revised period-luminosity relation for single Miras (Whitelock et al.\\ 2008) would place HM Sge at $D$ = 2.5 kpc; this estimate is however uncertain due to the peculiar nature of the object, particularly because of the low amplitude and poor periodicity of its light curve (see the AAVSO database). } As we will see, the CO lines in these systems are extremely weak, typically about 100 times weaker than for standard AGB stars. We will argue that this low intensity, as well as the other main properties of the detected lines, show that the CO emission only comes from very inner circumstellar regions around the cool stellar component, closer than about 1--2 10$^{14}$ cm. ", "conclusions": "We present observations of the \\doce\\ and \\trece\\ \\jdu\\ and \\juc\\ transitions and of the SiO \\jcc\\ one in three symbiotic stellar systems (SSs), R Aqr, CH Cyg, and HM Sge. \\doce\\ \\jdu\\ and \\juc\\ emissions were detected in R Aqr and CH Cyg. An accurate line profile of \\doce\\ \\jdu\\ in R Aqr was obtained. The observed lines are very weak. If we compare them with those usually observed in standard evolved stars with similar properties \\citep[see e.g.][]{bujetal92}, the CO lines in SSs are $\\sim$ 100 times weaker, even if we also consider objects, like young planetary nebulae, in which molecule photodissociation seems to be already efficient. The weak intensities of the observed lines suggests that CO emission comes from a very compact region, much smaller than the region responsible for low-$J$ CO emission in circumstellar envelopes (CSEs) around standard AGB stars (typical radius larger than 10$^{16}$ cm). This result is compatible with observations (mostly nondetections) in SSs of other molecular lines, like the SiO, H$_2$O and OH masers, see Sect.\\ 1. We have shown, from very general considerations, Sects.\\ 3, 4, that the detected CO intensities are coincident with those expected for emission coming from a circumstellar region comparable in radius to (or slightly smaller than) the distance between the stellar components of the system. We assumed that these molecule-rich layers in SSs have physical conditions similar to those of the inner layers around standard AGB stars. Particularly satisfactory is the fitting of the lines observed in the most intense source, R Aqr, in which the stars are separated by typically 2.3 $\\times$ 10$^{14}$ cm, see Sect.\\ 1. We have also performed model calculations of the CO emission of these shells, assuming emitting region radii of that order and the physical conditions and dynamics expected in gas placed at similar distances around standard AGB stars (Sects.\\ 3, 4). Our calculations reproduce the observed intensities, within the observational uncertainties. We can also reproduce the \\doce\\ \\jdu\\ line profile observed in R Aqr, see Sect.\\ 4.1 and Fig.\\ 4. Our model explains in particular the asymmetry observed in the profile, as a result of selfabsorption in an optically thick line and in the presence of significant gas acceleration and systematic temperature decrease with radius. All these properties are expected to be present in the very inner layers from which we are assuming that the CO emission comes. {The data obtained in CH Cyg are worse and their analysis is more uncertain.} To explain the observed intensities, we deduce a typical radius of about 10$^{14}$ cm, again slightly smaller than the distance between the stars (Sects.\\ 1, 4.2). In this region, the physical conditions seem also comparable to those typical of AGB stars, although the expansion velocity, $\\sim$ 25 \\kms, is slightly larger than those usual in CSEs around standard red giants. We so conclude that the nebulae around the SSs R Aqr and CH Cyg include a very compact molecule-rich region, with radii $\\sim$ 10$^{14}$ -- 2\\,$\\times$\\,10$^{14}$ cm, from which the molecular emission detected in these objects comes. This region is in extent comparable to or slightly smaller than the region within both components of the binary system. We deduce in these regions outwards expansion velocities of about 5 -- 25 \\kms, with a significant acceleration of the gas, temperatures decreasing with radius between about 1000 and 500 K, and densities $\\sim$ 10$^9$ -- 3\\,$\\times$\\,10$^8$ cm$^{-3}$. The general physical conditions and dynamics in this inner shell would be then similar to those typical of standard AGB stars. This result is supported by the 'normal' observational properties of SiO maser emission from R Aqr (Sects.\\ 3.4, 4), which show that the CSE around this star is relatively normal (i.e.\\ that its main properties are comparable to those of standard AGB stars) up to distances $\\sim$ 5 10$^{13}$ cm, at least. {We find relatively high values of the gas mass and mass-loss rate. In R Aqr, we estimate that the total mass of the CO emitting region is $\\sim$ 2 -- 3 $\\times$ 10$^{-5}$ \\ms, and that the value of \\mloss\\ ranges between 5 10$^{-6}$ and 10$^{-5}$ \\my, in agreement with results obtained from dust emission. These relatively high values therefore confirm the general trend of SSs to present significantly higher mass-loss rates than standard, isolated evolved stars \\citep{miko99, miko02}. We must take into account, however, that the comparison between mass-loss rates in the very extended and isotropic winds from standard AGB stars and the values of \\mloss\\ deduced for our sources is difficult. Here, we are really detecting gas which is confined within the orbit of the companion, and its future dynamics is difficult to know. It is probable that most of this material will never leave the system (in this case we would not be dealing with a true mass ejection) or will leave it by means of outbursts or bipolar jets. It is also possible that the process of mass ejection in these stars is significantly affected by the companion and cannot be easily compared with that in isolated AGB stars. Hydrodynamical simulations \\citep{gav02,pods07} show indeed that, in systems with orbital parameters similar to those of CH Cyg and R Aqr, the radius of the dust shell is comparable with the Roche lobe radius and a wind Roche-lobe overflow will occur. In any case, we think we can safely conclude that the mass of the gas in the close surroundings of R Aqr and CH~Cyg is relatively high, and that its presence is probably related to the gravitational effects of the nearby companion, even if these processes are not well understood. } We note that the few symbiotic systems showing molecular emission (as R Aqr and CH Cyg) are objects showing particularly large distances between the stars \\citep{schwarz95}. It is then probable that in other SSs the molecule-rich region is an extremely small shell tightly surrounding the Mira-type component, which would explain its very weak molecular emission. From the properties we derive for the CO-rich region and the nondetection of SiO thermal lines (Sect.\\ 4.3), whose emitting region must be significantly smaller than for CO, we suggest that the small extent of the molecule-rich gas in these sources is mainly due to molecule photodissociation by the hot dwarf star, which must be significantly more efficient for low-abundance molecules. Ionized gas has been found to be very extended in SSs, particularly in the well studied R Aqr \\citep[see e.g.][]{corradi03, hollis99a, hollis99b}, where bipolar and disk-like ionized-gas structures are known to extend over more than one arcminute (several 10$^{17}$ cm), apparently without any molecule-rich counterpart. It is obvious that the relatively standard inner shells around the red giant show very different properties than the spectacular outer parts of the nebula. Therefore, disruption of the CSE due to the strong stellar interaction in SSs must also play an important role in the confinement of gas (in particular, molecule-rich gas) in a very compact and dense shell. \\cite{seaquist95} proposed that the powerful wind from the companion may sweep much of the material ejected by the Mira star. This effect would explain the strong discontinuity found between the general properties of the regions inside and outside the orbit of the companion and, following \\cite{seaquist95}, could increase the efficiency of photodissociation in the rest of the nebula, helping to explain the lack of molecular emission in SSs." }, "1003/1003.4522_arXiv.txt": { "abstract": "We report the detection of Cd~\\textsc{i} ($Z =$~48), Lu~\\textsc{ii} ($Z =$~71), and Os~\\textsc{ii} ($Z =$~76) in the metal-poor star BD~$+$17~3248. These abundances are derived from an ultraviolet spectrum obtained with the Space Telescope Imaging Spectrograph on the \\textit{Hubble Space Telescope}. This is the first detection of these neutron-capture species in a metal-poor star enriched by the $r$-process. We supplement these measurements with new abundances of Mo~\\textsc{i}, Ru~\\textsc{i}, and Rh~\\textsc{i} derived from an optical spectrum obtained with the High Resolution Echelle Spectrograph on Keck. Combined with previous abundance derivations, 32 neutron-capture elements have been detected in BD~$+$17~3248, the most complete neutron-capture abundance pattern in any metal-poor star to date. The light neutron-capture elements (38~$\\leq Z \\leq$~48) show a more pronounced even-odd effect than expected from current Solar system $r$-process abundance predictions. The age for \\bd\\ derived from the Th~\\textsc{ii}/Os~\\textsc{ii} chronometer is in better agreement with the age derived from other chronometers than the age derived from Th~\\textsc{ii}/Os~\\textsc{i}. New Hf~\\textsc{ii} abundance derivations from transitions in the ultraviolet are lower than those derived from transitions in the optical, and the lower Hf abundance is in better agreement with the scaled Solar system $r$-process distribution. ", "introduction": "\\label{intro} Steady progress has been made over the last half-century toward understanding how the heaviest elements in the Universe are produced. For the elements heavier than the iron (Fe) group, the vast majority of isotopes are produced by the successive addition of neutrons to existing nuclei on timescales that are slow or rapid relative to the average $\\beta^{-}$ decay rates. These are referred to as the slow ($s$) and rapid ($r$) neutron ($n$) capture processes, respectively (see, e.g., \\citealt{truran02} and \\citealt{sneden08} for discussion of these processes). The basic physical principles of these reactions are well known. The \\spro\\ involves isotopes near the valley of $\\beta$ stability, so the properties relevant to understanding the nature of the \\spro\\ (e.g., \\ncap\\ cross sections, half-lives, etc.) can be studied in laboratories on Earth (see \\citealt{cowan91} and references therein). Phenomelogical or nuclear reaction models can then be constructed to predict the general abundance pattern produced by the \\spro\\ (see \\citealt{busso99}). When applied to the Solar system (S.~S.) heavy element abundance distribution, the \\spro\\ abundances can be subtracted from the total abundances to reveal the \\rpro\\ component (e.g., \\citealt{seeger65,kappeler89,arlandini99}). Due to the more energetic nature of the \\rpro\\ and the exotic, short-lived nuclei involved, reaction networks for the \\rpro\\ were not tractable until only recently (see \\citealt{kratz07}). To evaluate and verify detailed nucleosynthesis models, abundance patterns must be accurately characterized for as many elements as possible in locations beyond the S.~S. In this Letter, we report abundance estimates for neutral cadmium (Cd~\\textsc{i}, $Z =$~48), singly ionized lutetium (Lu~\\textsc{ii}, $Z =$~71), and singly ionized osmium (Os~\\textsc{ii}, $Z =$~76) in the near-ultraviolet (NUV) spectrum of the \\rpro\\ enriched metal-poor star \\bd. This is the first clear detection of Cd and Lu in a metal-poor star enriched by the \\rpro. Combined with previous abundance derivations \\citep{cowan02,cowan05,sneden09} and several other new abundances derived from the optical spectrum of this star, 32 \\ncap\\ elements have been detected in \\bd, the most complete \\ncap\\ pattern in any metal-poor star. In the metal-poor star \\hd, which is relatively deficient in the heavy \\ncap\\ elements, we also report tentative detections for Cd~\\textsc{i} and Lu~\\textsc{ii}, as well as an upper limit for Os~\\textsc{ii}. Finally, we use these new abundances to differentiate among the various techniques used to predict the \\rpro\\ abundance pattern. ", "conclusions": "\\label{results} In Figure~\\ref{solarr} we display the abundance distribution for the \\ncap\\ elements in \\bd. The S.~S.\\ \\rpro\\ abundance distribution, calculated as residuals from a classical model of the \\spro\\ \\citep{sneden08}, is shown for comparison. When normalized to the Eu abundance in \\bd, this distribution is a superb match to the stellar abundances for $Z \\geq$~56, as has been demonstrated previously (e.g., \\citealt{cowan02}). \\begin{figure} \\begin{center} \\includegraphics[angle=0,width=3.4in]{f4.eps} \\end{center} \\caption{ \\label{solarr} Neutron-capture abundance distributions in \\bd\\ and \\hd. Detections are indicated by filled symbols, and upper limits are indicated by downward-facing open triangles. The bold blue line in the top panel represents the core collapse supernova high entropy neutrino wind calculations of \\citet{farouqi09} (estimated from their Figure~3), normalized to Sr ($Z =$~38). The solid line in each panel represents the scaled S.~S.\\ \\rpro\\ abundance distribution \\citep{sneden08}, normalized to Eu ($Z =$~63), and the dotted line in the lower panel represents this same distribution normalized instead to Sr. Abundances are taken from \\citet{cowan02,cowan05}, \\citet{honda06}, \\citet{roederer09}, \\citet{sneden09}, and the present study. Several elements have been renormalized to a common set of laboratory log($gf$) values, and the abundances of \\hd\\ have been renormalized to our abundance scale as described in the caption of Figure~\\ref{fezr}. } \\end{figure} The detection of Cd extends the suite of lighter \\ncap\\ elements observed in metal-poor stars more than halfway between the 1st and 2nd \\rpro\\ peaks (roughly $A \\sim$~80 and 130, respectively). For Sr to Cd ($Z =$~38--48, missing only the short-lived isotopes of Tc, $Z =$~43), there is a very pronounced even-odd abundance pattern in \\bd, much more so than is predicted by the scaled S.~S.\\ \\rpro\\ distribution or the \\rpro\\ residuals derived from the (Solar metallicity) stellar \\spro\\ model of \\citealt{arlandini99} (not shown). This effect was first noticed in the Pd and Ag abundance pattern of three \\rpro\\ enriched stars observed by \\citet{johnson02}. The large-scale dynamical network calculations from the core collapse supernova high-entropy neutrino wind model of \\citet{farouqi09} reproduce this pattern better for Sr--Pd but still underestimate the even-odd effect for Ag and Cd. This general agreement is encouraging, but additional calculations and comparisons are warranted. Lu is the final member of the rare earth elements (REE) to be unambiguously detected in \\rpro\\ enriched metal-poor stars. The Lu/Eu (or, more generally, Lu/REE) ratio predicted by the scaled S.~S.\\ \\rpro\\ distribution is in reasonable, though not perfect, agreement with our derived Lu abundance. Additional Lu abundance derivations for other metal-poor, \\rpro\\ enhanced stars are required to assess whether the predicted Lu abundance or the stellar measurement (or both) is in error. Each of the neutral and singly-ionized states of Os have now been detected in \\bd, and our abundance of Os~\\textsc{ii}, log~$\\epsilon = +$0.03, is in fair agreement with an updated Os~\\textsc{i} abundance derived from three optical and NUV lines, log~$\\epsilon = +$0.25. Os is the heaviest stable element that can be detected in its singly-ionized state in \\bd. If the Os~\\textsc{ii} abundance should prove reliable and its uncertainty can be reduced, this has the potential to offer two significant improvements for nuclear cosmochronometry. The only radioactive isotopes practical for age dating the material in old stars are \\iso{232}{Th} and \\iso{238}{U}, both of which can only be detected as first ions. Abundance uncertainties are minimized when considering ratios of two elements in the same ionization state. When predicting the initial production ratios, the uncertainty is generally smallest when the two elements are as close in mass number as possible. Previously, the heaviest singly-ionized reference element has been Hf~\\textsc{ii}, whose stable isotopes are separated by 52--55 mass units from \\iso{232}{Th}; the stable \\rpro\\ isotopes of Os are only separated by 40--44 mass units from \\iso{232}{Th}. Adopting the range of production ratios from \\citet{kratz07}, the Th~\\textsc{ii}/Os~\\textsc{i} ratio predicts an age range of 15.7--21.5~Gyr, whereas Th~\\textsc{ii}/Os~\\textsc{ii} predicts an age range of 5.4--11.2~Gyr. The latter is in better agreement with the age predicted from other chronometer pairs (e.g., Th~\\textsc{ii}/Eu~\\textsc{ii}, which predicts an age range of 7.9--12.3~Gyr). The present uncertainty in our Os~\\textsc{ii} abundance translates to an age precision of 14~Gyr, but if the Os~\\textsc{ii} uncertainty could be reduced to 0.10~dex the age precision would improve to 4.7~Gyr. The Hf~\\textsc{ii} abundance derived in \\bd\\ from 4 lines in the NUV is marginally lower (log~$\\epsilon = -$0.76~$\\pm$~0.08, $\\sigma =$~0.14) than that derived from 6 lines in the optical (log~$\\epsilon = -$0.57~$\\pm$~0.03, $\\sigma =$~0.08). Previous analyses of the Hf~\\textsc{ii} abundance in \\rpro\\ enriched metal-poor stars have revealed that the stellar Hf~\\textsc{ii} \\rpro\\ abundance is higher by 0.15--0.25~dex than that predicted by the scaled S.~S.\\ \\rpro\\ distribution \\citep{lawler07,roederer09,sneden09}. Our Hf~\\textsc{ii} NUV abundance is in excellent agreement with the scaled S.~S.\\ \\rpro\\ Hf/REE ratio. Many of the transitions used to derive the REE stellar \\rpro\\ abundance distribution from the optical spectral range arise from 0.0~eV lower levels, but only one of the 12 Hf~\\textsc{ii} lines employed by \\citet{lawler07} has a 0.0~eV lower level. Two of the four lines used to derive our NUV Hf~\\textsc{ii} abundance arise from 0.0~eV levels, with log~$\\epsilon = -$0.68 from just these two transitions. Perhaps by using these transitions we have mitigated a subtle systematic effect present in the computation of the Hf~\\textsc{ii} abundance relative to other REE. This might imply that other stellar Hf~\\textsc{ii} \\rpro\\ abundances---rather than the predicted S.~S.\\ \\rpro\\ abundances---warrant minor revisions. Again adopting the range of production ratios from \\citet{kratz07}, the NUV Th~\\textsc{ii}/Hf~\\textsc{ii} ratio predicts an age range of 5.8--18.5~Gyr, whereas the optical Th~\\textsc{ii}/Hf~\\textsc{ii} ratio predicts an age of 14.7--27.4~Gyr. The precision is 5--6~Gyr in each measurement, but clearly the lower Hf~\\textsc{ii} abundance derived from the NUV lines provides a more realistic age estimate for \\bd. Figure~\\ref{solarr} also displays the abundance distribution for \\hd, which is known to be deficient in the heavy \\ncap\\ elements. The scaled S.~S.\\ \\rpro\\ distribution is a poor fit to the abundance pattern whether normalized to the 1st \\rpro\\ peak or the REE \\citep{honda06}. (No reasonable \\spro\\ distribution, or combination of $r$- and \\spro\\ distributions, matches either.) \\hd\\ may be an example of enrichment by the so-called ``weak'' \\rpro, which produces small amounts of light \\ncap\\ material and steadily-decreasing amounts of heavier material \\citep{honda06,wanajo06}. Our Cd abundance in \\hd\\ suggests that the downward abundance trend continues in the region between the 1st and 2nd \\rpro\\ peaks. Our Os upper limit in this star is not strong enough to exclude a scaled S.~S.\\ \\rpro\\ pattern between the REE and the 3rd \\rpro\\ peak. The detection of these three new species in \\bd\\ is only a first step in understanding how and in what amount these elements were produced. By examining their abundances in other metal-poor stars enriched to different levels by the \\rpro, we may gain a better sense of any systematic offsets affecting the present analysis. These uncertainties must be minimized to take full advantage of these species as constraints on \\ncap\\ nucleosynthesis models and meaningful age probes for the \\ncap\\ material in metal-poor stars." }, "1003/1003.2903_arXiv.txt": { "abstract": "{} {The aim is to investigate the star-formation and LINER (Low Ionization Nuclear Emission Line Region) activity within the central kiloparsec of the galaxy NGC\\,1614. In this paper the radio continuum morphology, which provides a tracer of both nuclear and star-formation activity, and the distribution and dynamics of the cold molecular and atomic gas feeding this activity, are studied. In particular, the nature of an R$\\approx$300\\,pc nuclear ring of star-formation and its relationship to the LINER activity in NGC\\,1614 is addressed.} {A high angular resolution, multi-wavelength study of the LINER galaxy NGC\\,1614 has been performed. Deep observations of the CO 1-0 spectral line were performed using the Owens Valley Radio Observatory (OVRO). These data have been complemented by extensive multi-frequency radio continuum and H{\\sc i} absorption observations using the Very Large Array (VLA) and Multi-Element Radio Linked Interferometer Network (MERLIN).} {Toward the center of NGC\\,1614, we have detected a ring of radio continuum emission with a radius of 300\\,pc. This ring is coincident with previous radio and Pa$\\alpha$ observations. The dynamical mass of the ring based on H{\\sc i} absorption is $ 3.1 \\times 10^9 \\, \\rm M_{\\odot}$. The peak of the integrated CO 1-0 emission is shifted by 1$''$ to the north-west of the ring center. An upper limit to the molecular gas mass in the ring region is $\\sim1.7 \\times 10^9 \\, \\rm M_{\\odot}$. Inside the ring, there is a north to south elongated 1.4 GHz radio continuum feature, with a nuclear peak. This peak is also seen in the 5 GHz radio continuum and in the CO.} {We suggest that the $R=300$ pc star forming ring represents the radius of a dynamical resonance - as an alternative to the scenario that the starburst is propagating outwards from the center into a molecular ring. The ring-like appearance is probably part of a spiral structure. Substantial amounts of molecular gas have passed the radius of the ring and reached the nuclear region. The nuclear peak seen in 5 GHz radio continuum and CO is likely related to previous star formation, where all molecular gas was not consumed. The LINER-like optical spectrum observed in NGC~1614 may be due to nuclear starburst activity, and not to an Active Galactic Nucleus (AGN). Although the presence of an AGN cannot be excluded.} ", "introduction": "NGC~1614 is a luminous infrared galaxy (LIRG) at a distance of 64 Mpc (for $H_0$=75 \\kms \\permpc). The galaxy is barred and interacting (morphologically classified as type {SB(s)c~pec}) and its nuclear optical spectrum shows both starburst and LINER (Low Ionization Nuclear Emission Line Region) activity. Although there is still no single consensus to what is powering the emission in LINERs, nuclear starburst or AGN (Active Galactic Nucleus) activity are two major candidates (e.g. Ho 1999; Terashima \\etal\\, 2000; Alonso-Herrero \\etal\\, 2000). Bar-driven inflow of gas in galaxies has been suggested to trigger nuclear starbursts as well as feeding AGNs (e.g. Simkin, Su \\& Schwarz 1980; Scoville \\etal~1985). However, the radial inflow of gas along the bar may be slowed down at certain radii, often associated with inner Lindblad resonances (ILR) (Combes 1988; Shlosman, Frank \\& Begelman 1989). \\par Previous studies of NGC~1614 include e.g. an optical, near-infrared, radio continuum and {\\sc{hi}} study by Neff \\etal~(1990). They suggested that the spectacular structure with tidal tails or plumes is the result of an earlier interaction with another galaxy of comparatively modest mass and impact velocity. Alonso-Herrero \\etal~(2001), hereafter AH2001, have investigated NGC~1614, in the optical and near-infrared. They detected a starburst nucleus of about 45 pc in diameter based on deep CO stellar absorption. The nucleus is surrounded by a $\\sim$600 pc diameter ring of current star formation, which is revealed in Pa$\\alpha$ line emission. Just outside the star forming ring, a dust ring is indicated by its extinction shadow in $H-K$. Neither Neff \\etal~(1990) nor AH2001 find any indication of an AGN. Scoville \\etal~(1989) used the Owens Valley Radio Observatory (OVRO), which contained three elements at that time, to map NGC~1614 in CO 1--0 at a resolution of $4'' \\times 6''$. They found an unresolved CO-concentration with a mass of $6 \\times 10^9 $ M$_{\\odot} $. \\begin{figure*} \\centering \\includegraphics[width=17cm]{11538fg1.eps} \\caption{Overlay of OVRO CO 1--0 integrated intensity contours over a F606W WFPC2 {\\it{HST}} image. The contours are in percent of the peak value of 34 Jy\\,km\\,s$^{-1}$ per beam of $2.''75\\times2.''40$. The levels are 8, 11, 16, 23, 32, 45, 64 and 90\\%. The greyscale of the optical image is logarithmic and arbitrary. The astrometrical alignment of the {\\it{HST}} image with the CO map is within $1.''5$ (see Sect.~\\ref{coresults}).} \\label{hst} \\end{figure*} We have used OVRO to study the distribution and dynamics of the molecular gas at $2''$ resolution. We have also studied the morphology of the 1.4 GHz radio continuum as well as the distribution and dynamics of the neutral gas via H{\\sc i} in absorption at arcsecond resolution, using the Very Large Array (VLA). To study the morphology of the nuclear radio continuum, we used the Multi Element Radio Linked Interferometer Network (MERLIN) and obtained maps at higher (0.5 arcsecond) resolution, at 1.4 GHz and 5 GHz. The purpose of the studies was to provide information about the feeding and nature of the central activity giving rise to the LINER like spectrum in NGC~1614. In particular, we would like to address the question whether the LINER activity is due to an AGN or to a nuclear starburst. ", "conclusions": "\\begin{enumerate} \\item We detect $ 3 \\times 10^9 \\, \\rm M_{\\odot}$ of molecular gas within a north to south slightly elongated bar like structure with a size of $\\sim 2.2 \\times 1.5$ kpc. \\item The north to south rotation of the central kpc molecular gas is consistent with an inclined disc / bar that follows the expected rotational direction based on the assumption of trailing spiral arms. \\item There is a molecular extension to the northeast, associated with a crossing dust lane. \\item The molecular gas in the central kpc is double peaked, with peaks at $R=300$ pc. There is a third, central peak at lower integrated intensities, probably associated with a nuclear starburst/AGN. \\item We have detected a MERLIN and VLA 5 GHz radio continuum ring with a radius of 300 pc. This ring is coincident with previous radio continuum and Pa$\\alpha$ observations. We conclude that the radio continuum ring originates in recent star formation. The CO peak intensity is shifted 1$''$ to the north-west of the center of the starburst ring. \\item The MERLIN 1.4 GHz radio continuum is triple peaked, with a peak separation of $\\sim 1''$ (300 pc). The brightest peaks occur in the northern and southern part of the star forming ring with a radius of 300 pc. The other peak, which is almost as bright, occur inside the star forming ring, close to its center. The brightness temperature of maximum 53 000 K of the peaks are consistent with a cluster of supernova remnants, although an AGN can not be ruled out for the peak close to the center. \\item The position velocity diagram of the H{\\sc i}-absorption was used to calculate a dynamical mass of $ 3.5 \\times 10^9 \\, \\rm M_{\\odot}$ in the central $R=1''$ (300 pc). An upper limit to the molecular gas mass in the same region is $\\sim 1.3 \\times 10^9 \\, \\rm M_{\\odot}$. This gives a gas mass fraction of $\\sim 0.4$ or less in the central $R=1''$ (300 pc). \\item A solid body rotation curve was fitted to the H{\\sc i}-absorption velocity field, out to a radius of $2''$ (600 pc). The residuals are blueshifted over all of the fitted area. We interpret this as an outflow or a superwind, with a deprojected velocity of $\\sim$160 \\kms. \\item We suggest that gas may flow to the nucleus along a dynamically decoupled bar inside the star forming ring (or spiral), which we see tentative evidence for in the 1.4 GHz MERLIN map. The existence of such a bar should be investigated with CO observations and higher angular resolution. \\item The LINER like spectrum is likely related to the shock fronts associated with either the large scale super wind, or with shock fronts on smaller scales related to supernova remnants from previous bursts of star formation. We conclude that the LINER activity observed in NGC~1614 is probably due to starburst activity, and not to AGN-activity. \\end{enumerate}" }, "1003/1003.2418_arXiv.txt": { "abstract": "We perform numerical calculations of the expected transit timing variations (TTVs) induced on a Hot-Jupiter by an Earth-mass perturber. Motivated by the recent discoveries of retrograde transiting planets, we concentrate on an investigation of the effect of varying relative planetary inclinations, up to and including completely retrograde systems. We find that planets in low order (E.g. 2:1) mean-motion resonances (MMRs) retain approximately constant TTV amplitudes for $0<\\,^{\\circ}i<170\\,^{\\circ}$, only reducing in amplitude for $i > 170\\,^{\\circ}$. Systems in higher order MMRs (E.g. 5:1) increase in TTV amplitude as inclinations increase towards $45\\,^{\\circ}$, becoming approximately constant for $45 < i < 135$, and then declining for $i > 135\\,^{\\circ}$. Planets away from resonance slowly decrease in TTV amplitude as inclinations increase from 0 to 180, where-as planets adjacent to resonances can exhibit a huge range of variability in TTV amplitude as a function of both eccentricity and inclination. For highly retrograde systems ($135\\,^{\\circ} < i \\leq 180\\,^{\\circ}$), TTV signals will be undetectable across almost the entirety of parameter space, with the exceptions occurring when the perturber has high eccentricity or is very close to a MMR. This high inclination decrease in TTV amplitude (on and away from resonance) is important for the analysis of the known retrograde and multi-planet transiting systems, as inclination effects need to be considered if TTVs are to be used to exclude the presence of any putative planetary companions: absence of evidence is not evidence of absence. ", "introduction": "Of the known extra-Solar planets, more than 60 transit the host star. Of these systems, at least four show evidence for an external companion (GJ 436, \\citealt{Maness07}; HAT-P-13, \\citealt[][]{Bakos09}; HAT-P-7, \\citealt[][]{Winn09} and CoRoT-7 \\citealt[][]{Queloz09}). If the transiting planet were the only planet in the system, then the period between each successive transit would be constant (neglecting complicating effects such as general relativity, stellar oblateness and tides). The presence of additional planets in the system (which themselves may or may not transit the star) can cause perturbations to the orbit of the transiting planet, leading to detectable transit timing variations (TTVs) of the known transiting planet \\citep{Miralda-Escude02,Holman05,Agol05}. Other studies have extended the TTV method to demonstrate the feasibility of detecting planetary moons \\citep{Kipping09a,Kipping09b} and Trojan companions \\citep{Ford07b} in extra-solar planetary systems. However, practical use of TTVs as a detection tool requires the solution of the difficult ``inverse problem'', i.e. given a particular TTV profile, can one reconstruct (or at least restrict) the mass and orbit of the unseen perturber? This problem is non-trivial, as numerous different perturber mass-orbit configurations can lead to degenerate TTV solutions \\citep[e.g.][]{Nesvorny08}. Several investigations have considered (and ruled out) planets more massive than the Earth in certain orbits / regions of parameter space close to many of the known transiting planets \\citep[e.g. ][]{Steffen05,Agol07,Alonso08,Bean08}, while \\citet{Nesvorny08} and \\citet{Nesvorny09} have developed an approximate analytic method to try and tackle the inverse problem in a more general manner. However, most previous investigations have focused on analyzing the effect of prograde, coplanar companions. Even where explicit investigations of inclination effects have been conducted \\citep[e.g.][]{Nesvorny09}, the investigations have been restricted to prograde cases, examining only relative inclinations significantly lower than $90\\,^{\\circ}$. There are several reasons to consider TTVs of inclined systems. Observationally, it is possible to measure the sky-projected angle between the spin vector of a star, and the orbital angular momentum vector of a planet transiting that star via the Rossiter-McLaughlin effect. Measurements over the past year have revealed that a number of the known transiting planets are strongly inclined, or even retrograde (e.g. HAT-P-7b, \\citealt{Narita09} and \\citealt{Winn09}, as well as WASP-17b, \\citealt{Anderson09}). Although the mechanism(s) driving the creation of such retrograde planets remains unclear, models used to explain high \\emph{eccentricity} exo-planets via planet-planet scattering \\citep{Rasio96b,Weidenschilling96,Lin97b,Levison98,Papaloizou01,Moorhead05,Chatterjee08,Juric08,Ford08a} and/or Kozai oscillations \\citep[E.g.][]{Kozai62,Takeda05,Nagasawa08} can naturally excite large orbital \\emph{inclinations} (although the quantitative details of the distributions can differ greatly). Other mechanisms such as inclination-pumping during migration \\citep[e.g.][]{Lee09} may also contribute. In addition, previous dynamical studies have suggested that systems of prograde and retrograde planets are \\emph{more} stable than standard prograde-prograde cases \\citep{Gayon08,Gayon09,Smith09}, essentially because the planets spend less time ``close together'' and thus perturbations are smaller. Given the observational evidence for highly inclined and retrograde systems, and the suggestion of enhanced stability in such systems, we investigate and quantify the hypothesis that retrograde systems will have a significantly reduced TTV profile compared to a standard prograde case. ", "conclusions": "We have investigated the TTV signals for systems of highly inclined and retrograde planets. We find that: \\begin{itemize} \\item In the vicinity of exterior MMRs the inclination dependence is complex: low order resonances maintain a high TTV amplitude for all regions $i<170\\,^{\\circ}$, declining in amplitude only for low eccentricity cases close to $i=180\\,^{\\circ}$, where-as higher order resonances display an increase in TTV amplitude as inclinations rise from 0 to 45 degrees. Moreover, the regions immediately adjacent to MMRs show extreme sensitivity to changes in perturber $a$, $e$ and $i$. \\item Exterior perturbers away from resonances tend to show a slow decrease in TTV amplitude with increasing inclination, although regions adjacent to resonances can show remarkably complex behaviour. \\item Perturbing planets on interior orbits display a slightly different behaviour: Away from resonance the amplitude remains approximately constant with inclination, but around MMRs the perturbations become \\emph{stronger} as the inclination increases towards $90\\,^{\\circ}$ before decreasing again beyond $90\\,^{\\circ}$. \\end{itemize} We note that: \\begin{itemize} \\item Absence of evidence is not evidence of absence: Planets in retrograde orbits should be expected to produce markedly reduced TTV signals as compared to the standard prograde case. For an Earth-mass perturber in an anti-aligned orbit, almost the entirety of the sample parameter space would result in a very small or undetectable TTV signal (unless the planet happened to be fortuitously located precisely on a strong MMR). \\item Retrograde orbits may be a natural way to explain transiting systems in which little or no TTV signal is observed, but in which the radial velocity observations point towards the existence of an additional planetary companions (E.g. GJ 436, HAT-P-13, etc). \\item In addition to the TTV considerations in this work, it is important to acknowledge that inclined orbits in multi-planet systems will precess, leading to variations in transit duration. More work is required to try and understand whether a combination of TTV signals with transit \\emph{duration} variation signals could remove some of the degeneracies inherent in this problem. \\item Many additional dependencies (E.g. sampling period, perturber mass, relative mean anomaly, etc, etc) can significantly alter the expected TTV signal for a particular system. We defer the provision of a detailed investigation of such matters to a companion paper, \\emph{Veras et al. 2010, in prep.} \\item Future work could also investigate in more detail the regions \\emph{above} the Hill stability curve, seeking to identify the (likely large) TTVs for any stable systems which exist in that region. \\end{itemize}" }, "1003/1003.0008.txt": { "abstract": "The distribution of cold gas in dark matter haloes is driven by key processes in galaxy formation: gas cooling, galaxy mergers, star formation and reheating of gas by supernovae. We compare the predictions of four different galaxy formation models for the spatial distribution of cold gas. We find that satellite galaxies make little contribution to the abundance or clustering strength of cold gas selected samples, and are far less important than they are in optically selected samples. The halo occupation distribution function of present-day central galaxies with cold gas mass $> 10^{9} h^{-1} M_{\\odot}$ is peaked around a halo mass of $\\approx 10^{11} h^{-1} M_{\\odot}$, a scale that is set by the AGN suppression of gas cooling. The model predictions for the projected correlation function are in good agreement with measurements from the HI Parkes All-Sky Survey. We compare the effective volume of possible surveys with the Square Kilometre Array\\footnote{http://www.skatelescope.org} with those expected for a redshift survey in the near-infrared. Future redshift surveys using neutral hydrogen emission will be competitive with the most ambitious spectroscopic surveys planned in the near-infrared. ", "introduction": "Cold gas is central to galaxy formation yet little is known about how much there is in the Universe at different epochs and how this gas is distributed in dark matter haloes of different mass. The primary probe of atomic hydrogen, 21cm line emission, is incredibly weak. It is only in recent years that a robust and comprehensive census of atomic hydrogen (HI) in the local universe has been made possible through the HI Parkes All Sky Survey (Barnes et~al. 2001; Zwaan et~al. 2003, 2005). This work is being extended to lower mass systems by the ALFALFA survey (Giovanelli et~al. 2007). Despite this progress, the highest redshift direct detection of HI in emission is very firmly confined to the local Universe at $z=0.34$ (Lah et~al. 2009, see also Verheijen et~al. 2007). Information about cold gas in the high redshift Universe is restricted to absorption lines in quasar spectra (e.g. Peroux et~al. 2003). However, over the coming decade, this situation is expected to change dramatically with the construction of new, more sensitive radio telescopes such as the pathfinders for the Square Kilometre Array, MeerKAT (Booth et~al. 2009) and ASKAP (Johnston et~al. 2008), and the Square Kilometre Array itself (Schilizzi, Dewdney \\& Lazio 2008). The SKA will revolutionise our understanding of galaxy formation and cosmology, uncovering the HI Universe out to high redshifts. One of the major science goals is to better characterise the evolution of dark energy with redshift. The SKA is expected to provide competitive constraints on the nature of dark energy through high accuracy measurement of large-scale structure in the galaxy distribution over a lookback time representing a significant fraction of the age of the Universe (Albrecht et~al. 2006). This conclusion currently rests on very uncertain calculations which we seek to place on a firmer, more physical footing in this paper. Modelling the abundance and clustering of HI sources is challenging. A number of possible approaches have been tried; empirical modelling, which relies upon the observations of HI in the Universe, the fully numerical approach, which uses cosmological gas dynamics simulations to model the HI content of galaxies from first principles and semi-analytical modelling, which we use in this paper. Empirical estimates have been attempted despite the paucity of observational results for guidance (Abdalla \\& Rawlings 2005; Abdalla, Blake \\& Rawlings 2010). Such calculations require an assumption about the evolution of the HI mass function over a broad redshift interval. The only constraint on this assumption is the integrated density of HI, which can be compared with the results inferred from quasar absorption features, which themselves require corrections for unseen low column density systems and dust extinction (Storrie-Lombardi et~al. 1996). The empirical approach does not predict the clustering of HI sources. Further assumptions and approximations are necessary to extend this class of modelling so that predictions can be made for galaxy clustering. Another layer of approximation in this class of modelling has been motivated by observations which suggest that HI sources tend to avoid the centres of clusters and that clusters do not boast an important population of satellites (e.g. Waugh et~al. 2002; Verheijen et~al. 2007). This led Marin et~al. (2009) to make a one-to-one connection between halo mass and HI mass. However, the nature of the relation is uncertain and several possibilities are explored by Marin et~al. based on different assumptions about the evolution of the HI mass function. Ideally, a physically motivated model which follows the sources and sinks of cold gas is needed. Gas dynamic simulations are computationally expensive and are typically restricted to small computational volumes, which makes it impossible to accurately follow the growth of structure to the present day. An example is provided by Popping et~al. (2009), who carry out a smoothed particle hydrodynamics simulation in a $32 h^{-1}$Mpc box. The HI mass function in the simulation is in very poor agreement with the observational estimate of Zwaan et~al. (2005), underpredicting the abundance of galaxies of HI mass $10^{10} M_{\\odot}$ by a factor of 30, which the authors put down to the small computational volume, and overpredicting low mass systems by a factor of two. Clustering predictions are limited to scales smaller than a few Mpc due to the small box size. Furthermore, it is important to be aware that gas dynamic simulations do not have the resolution to follow all of the processes in galaxy formation directly and in all cases resort to what are essentially semi-analytical rules to treat sub-resolution physics. Currently the most promising route to making physical and robust predictions for the HI in the Universe is semi-analytical modelling of galaxy formation (see Baugh 2006). This type of model includes a simplified but physically motivated treatment of the processes which control the amount of cold gas in a galaxy: gas cooling, galaxy mergers, star formation and reheating of gas by supernovae. These calculations are quick and can rapidly cover the haloes in a cosmological volume. Baugh et~al. (2004) presented predictions for the mass function of cold gas galaxies in the {\\tt GALFORM} semi-analytical model of Cole et~al. (2000). One issue which must be dealt with is that the models predict only the total mass of cold gas, which includes helium, and both atomic and molecular hydrogen. Baugh et~al. assumed a fixed ratio of molecular to atomic hydrogen. Obreschkow \\& Rawlings (2009) developed an empirical model based on observations and theoretical arguments by Blitz \\& Rosolowsky (2006) in which this ratio could vary from galaxy to galaxy. Obreschkow \\& Rawlings applied this ansatz to the semi-analytical model of de Lucia \\& Blaizot (2007). In the first paper in this series, we compared the predictions of a range of semi-analytical models for the mass function of HI (Power et~al. 2009). Despite the different implementations of the physical ingredients used in the models and the different emphasis placed on various observations when setting the model parameters, the predictions show generic features. Power et~al. found that there is surprisingly little variation in the predicted HI mass function with redshift, and that the models make similar predictions for the rotation speed and size of HI systems. The models predict the mass of cold gas and so a conversion is required to turn this into a HI mass. Currently the most uncertain step is the assumption about what fraction of hydrogen is in atomic form and what fraction is molecular. Power et~al. presented predictions for two cases, one in which all model galaxies are assumed to have a fixed molecular to atomic hydrogen ratio ($H_{2}/$HI) and the other in which this ratio varies from galaxy to galaxy, depending upon the local conditions in the galactic disk (Blitz \\& Rosolowsky 2006). The assumption of a variable $H_{2}/$HI ratio results in a dramatic reduction in the number of HI sources in the tail of the redshift distribution. In this paper we look at the distribution of cold gas in galaxies as a function of halo mass. In particular we look at the halo occupation distribution (HOD) for HI galaxies, which gives the mean number of galaxies of a given HI mass as a function of dark matter halo mass, and the clustering of HI galaxies. Using this information, we assess the potential of the SKA to measure the baryonic acoustic oscillation (BAO) signal. We briefly review the {\\tt GALFORM} model in Section 2, explaining the differences between the four models that we consider. We then look at the halo occupation distribution of cold gas galaxies in Section 3, in which we also present predictions for the clustering of cold gas galaxies at different redshifts and compare to measured clustering at the present day. In Section 4 we compare the performance of future redshift surveys in the optical and using HI emission for measuring the properties of the dark energy. We present a summary along with our conclusions in Section 5. ", "conclusions": "\\label{Summary} The cold gas content of galaxies and its variation with halo mass lie at the core of the galaxy formation process. The amount of cold gas in a galaxy is set by the balance between a number of competing processes. The cold gas supply comes from the cooling of gas from the hot halo and the accretion of cold gas following mergers with other galaxies. Star formation and supernova feedback act as sinks of cold gas. Semi-analytical simulations model all of these processes in the context of structure formation in the dark matter and so are ideally suited to make predictions for the distribution of cold gas in haloes of different mass. Since the models can make a wide range of predictions, their parameters are set by the requirement that a variety of observed galaxy properties be reproduced, not just the local HI data. The model predictions can be tested by measurements of the clustering of HI-selected galaxy samples, and are necessary to plan surveys to measure the large-scale structure of the Universe with the next generation of radio telescopes. In this paper we have compared the predictions for the distribution of cold gas in dark matter haloes of four versions of the Durham semi-analytical galaxy formation model, {\\tt GALFORM}. The Bower et~al. (2006) and Font et~al. (2008) models are publicly available from the Millennium Archive. These models overpredict the local abundance of galaxies as a function of their cold gas mass. This excess was straightforward to fix, with the primary adjustment made to the model star formation timescale. This modified model, based on Bower et~al. (2006) was still able to reproduce the quality of match to the optical luminosity functions enjoyed by Bower et~al. We also considered a galaxy formation model set in a different cosmology, to take advantage of a N-body simulation with a large enough box size to accurately model baryonic acoustic oscillations. This model also adopted a modified star formation timescale to better match the local HI mass function. The model predictions have several features in common. In agreement with observations, satellite galaxies are relatively unimportant in samples selected by their cold gas mass. This is true even in the Font et~al. (2008) model in which satellites retain some of their hot haloes, depending on their orbit within the main halo, and can continue to accrete cooling gas. Samples constructed according to a cold gas mass threshold are dominated by central galaxies in haloes around $10^{11} h^{-1} M_{\\odot}$. The halo occuption distribution of central galaxies is peaked in halo mass, rather than being a step function, as is the case for optical samples. As the cold gas mass cut is increased, the width of the central galaxy HOD increases and the amplitude drops. The peaked nature of the HOD of central galaxies is due to suppression of gas cooling in masses haloes following heating by AGN. We found the same general form for the HOD in a model by de Lucia \\& Blaizot 2007, in which the implementation of AGN/radio mode feedback is different from that in {\\tt GALFORM}. The relative importance of central and satellite galaxies has an impact on the form of the predicted correlation function. The correlation function of a galaxy sample selected by cold gas mass is remarkably similar on small scales in real and redshift space. For pair separations in excess of a few Mpc, the redshift space correlation function has a higher amplitude than in real space, as expected given the effective bias of the sample (Kaiser 1987). In contrast, for an optically selected sample with the same number density of galaxies, the correlation steepens in real space for $r < 1 h^{-1}$Mpc and is damped in redshift space on these scales, due to the greater influence of satellite galaxies in massive haloes. On larger scales there is a more modest boost in the clustering amplitude in redshift space, due to the larger effective bias of the optical sample. The clustering predictions are in reasonable agreement with the measurements by Meyer et~al. (2007). The clustering in the modified version of the Bower et~al. model (MHIBow06) best agrees with the HIPASS results. One of the primary science goals of the Square Kilometre Array (SKA) is to make a high precision measurement of large-scale structure in the galaxy distribution. By measuring the apparent size of baryonic acoustic oscillations (BAO) at a particular redshift, the cosmological distance to that redshift can be derived, thereby constraining the equation of state of the dark energy. By combining the galaxy formation model with a very large volume N-body simulation ($ 1 h^{-3} {\\rm Gpc}^{3}$), we have been able to demonstrate that galaxy samples constructed on the basis of cold gas mass can trace the BAO with the same fidelity as an near-infrared selected sample with the same number density of galaxies. The key remaining question is how effectively do HI and optical redshift surveys sample the available geometrical volume and how does this translate into an error on the dark energy equation of state parameter? The effective survey volume varies substantially between HI surveys of different duration and for different assumptions about the split between atomic and molecular hydrogen. However, at least for the case of a cosmological constant, these differences occur in a redshift range which has little impact on the derived error on the equation of state. We find that HI surveys are comparable to the most ambitious near-infrared spectroscopic surveys currently under discussion, and will give a factor of two smaller error on $w$ than a slitless H-$\\alpha$ redshift survey; all are bone fide Stage V experiments in the Dark Energy Task Force nomenclature (Albrecht et~al. 2006). The uncertainty in the ratio of molecular to atomic hydrogen is one of the major uncertainties at present, and leads to larger differences in the predicted counts of HI emitters than the choice of galaxy formation model. The fraction of molecular hydrogen is thought to depend upon the local conditions in the interstellar medium. This question requires further modelling (e.g. Krumholz, McKee \\& Tumlinson 2009), augmented by observations of the HI and CO distribution in nearby galaxies, for example by HI surveys on the SKA pathfinder MeerKAT and CO measurements using the Atacama Large Millimeter/submillimeter Array (Wootten 2008)." }, "1003/1003.2304_arXiv.txt": { "abstract": "We investigate cosmological constraints on an energy density contribution of elastic dark matter self-interactions characterized by the mass of the exchange particle $m_\\mrm{SI}$ and coupling constant $\\alpha_\\mrm{SI}$. Because of the expansion behaviour in a Robertson-Walker metric we investigate self-interacting dark matter that is warm in the case of thermal relics. The scaling behaviour of dark matter self-interaction energy density ($\\varrho_\\mrm{SI}\\propto{a^{-6}}$) shows that it can be the dominant contribution (only) in the very early universe. Thus its impact on primordial nucleosynthesis is used to restrict the interaction strength $m_\\mrm{SI}/\\!\\sqrt{\\alpha_\\mrm{SI}}$, which we find to be at least as strong as the strong interaction. Furthermore we explore dark matter decoupling in a self-interaction dominated universe, which is done for the self-interacting warm dark matter as well as for collisionless cold dark matter in a two component scenario. We find that strong dark matter self-interactions do not contradict super-weak inelastic interactions between self-interacting dark matter and baryonic matter ($\\sigma_\\mrm{A}^\\mrm{SIDM}\\ll\\sigma_\\mrm{weak}$) and that the natural scale of collisionless cold dark matter decoupling exceeds the weak scale ($\\sigma_\\mrm{A}^\\mrm{CDM}>\\sigma_\\mrm{weak}$) and depends linearly on the particle mass. Finally structure formation analysis reveals a linear growing solution during self-interaction domination ($\\delta\\propto{a}$); however, only non-cosmological scales are enhanced. ", "introduction": "introduction\\protect} In the past decades high-precision observations allowed the development of a standard model of cosmology: $\\Lambda$CDM. Its main statements are that we are living in a flat universe ($\\Omega_\\mrm{tot}^0={1.0052}\\pm{0.0064}$), dominated by the {\\textquoteleft}dark' components: dark energy ($\\Omega_\\mrm{DE}^0={0.721}\\pm{0.015}$) and non-baryonic dark matter ($\\Omega_\\mrm{DM}^0={0.233}\\pm{0.013}$) \\cite{Hinshaw_2009}.\\\\ The necessity of a dark energy component comes from the acceleration of the universe expansion, inferred from a high redshift Hubble diagram of type Ia supernovae as standard candles \\cite{Kowalski_2008} and radio galaxies as standard yardsticks \\cite{Daly_2009}. A recent and impressive proof for the existence of dark matter (DM) can be deduced from observations of colliding galaxy clusters. Optical and near infrared observations of the galaxies, X-ray emission of the upheated intergalactic plasma and gravitational lensing of the mass distribution show the necessity of a non-visible DM component, which dominates the mass budget \\cite{Clowe_2006, Bradac_2008}. An overview of DM physics and particle candidates can e.g.\\ be found in recent reviews \\cite{Baltz_2004, Bergstroem_2009_A, Bertone_2005, Taoso_2008}. Numerical structure formation simulations in the $\\Lambda$CDM framework (from one of the first and most popular \\cite{Navarro_1996} to the most recent \\cite{Springel_2008, Navarro_2010}) show an impressive agreement with observations and have therefore become a cornerstone of modern cosmology. Nevertheless they also reveal two shortcomings that are worth being taken seriously. First, simulations predict scale-independently a large number of substructures in collisionless cold dark matter (CDM) halos, which exceed on galactic scales clearly the number of yet observed Milky Way satellites \\cite{Klypin_1999, Moore_1999}. One explanation is that reionization could prevent formation of visible baryonic structures in the smallest DM halos (e.g.\\ \\cite{Shapiro_2004}). Second, simulations show cusps in the center of collisionless CDM halo density profiles. But observations of dwarf spheroidal galaxies -- which have a huge mass-to-light ratio and hence are objects suited to study DM properties without perturbing baryonic effects, rotation curves of high spatial resolution and large extension of low luminosity spiral galaxies, and the universal rotation curve for spiral galaxies indicate a constant DM halo core density (e.g.\\ \\cite{Gilmore_2007, Gentile_2004, Salucci_2007}). An overview about processes that might lead from intrinsic cuspy CDM distributions to the observed cored ones gives e.g.\\ Ref.\\ \\cite{deBlok_2010}. An idea to avoid both mismatches of the CDM scenario is to introduce strong elastic DM self-interactions \\cite{Spergel_2000}. A recent overview concerning collisional DM is given in Ref.\\ \\cite{Taoso_2008}. Here, we want to concentrate on the most important facts that are also relevant for this work. The original self-interaction strength proposed by Ref.\\ \\cite{Spergel_2000} is $\\sigma_\\mrm{SI}/m_\\mrm{DM}=0.45-450\\,\\mrm{cm^2/g}$ (self-interaction cross-section over DM particle mass). But in Ref.\\ \\cite{Yoshida_2000} it was demonstrated that cross-sections generating reasonable dwarf galaxy cores predict too large galaxy cluster cores. Ref.\\ \\cite{Donghia_2003} showed that independent of the dependence on the halo velocity dispersion self-interacting cross-sections cannot solve the satellite problem accurately. The most reliable constraint on the self-interaction strength comes again from observations of colliding galaxy clusters \\cite{Markevitch_2004, Randall_2008}. The nonexistence of an offset between the galaxy distribution and the gravitational lens mass peak, and the subcluster mass-to-light ratio allow to constrain $\\sigma_\\mrm{SI}/m_\\mrm{DM}<0.7\\,\\mrm{cm^2/g}$. This result nearly completely rules out the formerly proposed self-interaction strength. The strongest limit on the collisional character of DM ($\\sigma_\\mrm{SI}/m_\\mrm{DM}<0.02\\,\\mrm{cm^2/g}$) can be inferred from the observed ellipticity of DM halos and the property of DM collisions to make halos spherical \\cite{Miralda_2002}, but one has to take into account its model dependence \\cite{Markevitch_2004, Randall_2008}. Another approach that avoids the satellite and cuspy halo problems is to provide the DM particles with a finite thermal streaming velocity, to achieve a cut-off in the power spectrum and smearing of the innermost, highest density halo regions \\cite{Bode_2001, Sommer_2001} (see also \\cite{Gilmore_2007}). This means to introduce warm dark matter (WDM) particles. A lower bound on the DM particle mass can be determined from the Lyman-$\\alpha$ forest ($m_\\mrm{WDM}\\gtrsim4\\,\\mrm{keV}$, \\cite{Viel_2008}) and gravitational lensing ($m_\\mrm{WDM}\\gtrsim2.2\\,\\mrm{keV}$, \\cite{Miranda_2007}) of high redshift quasars (given limits are for thermal relics). Ref.\\ \\cite{Boyarsky_2009} showed that these boundaries can be lowered considerably in a set-up of mixed cold and warm DM, which we also consider in this work. So we follow Ref.\\ \\cite{Boyanovsky_2008} in using $1-10\\,\\mrm{keV}$ as a typical mass range for WDM particles in the following. Other solutions proposed are stronger CDM annihilations (\\cite{Kaplinghat_2000}, but see also \\cite{Beacom_2007}) or a coupling between quintessence dark energy and CDM (e.g.\\ \\cite{Baldi_2010}). Present attempts to enlarge the phenomenology of DM physics are e.g.\\ strong DM baryonic matter interactions \\cite{Cyburt_2002} or a dark radiation (electromagnetism) between DM particles \\cite{Ackerman_2009}. We introduce in this work an energy density contribution of elastic dark matter self-interactions. Despite the fact that self-interacting dark matter (SIDM) may not solve the shortcomings of the collisionless approach, the motivation for this work is to explore new, interesting cosmological consequences of an additional energy density contribution of DM self-interactions within the above mentioned constraints. Interestingly enough, an interaction strength of $\\sigma_\\mrm{SI}/m\\sim1\\,\\mrm{cm^2/g}$ still corresponds to strong interactions ($\\sigma_\\mrm{strong}\\sim10\\,\\mrm{fm^2}$) between nucleon-like particles ($m\\sim1\\,\\mrm{GeV}$).\\\\ In Sec.\\ \\ref{sec:SIDM model} we introduce our idea of a self-interaction energy density contribution $\\varrho_\\mrm{SI}$. Energy density scaling according to the Friedmann equations and its equation of state ($p_\\mrm{SI}=\\varrho_\\mrm{SI}$) imply that the self-interaction contribution shows the steepest decrease with the scale factor ($\\varrho_\\mrm{SI}\\propto{a^{-6}}$) and thus can (only) have a direct impact on the very early universe. Its proportionality to the SIDM particle density ($\\varrho_\\mrm{SI}\\propto{n_\\mrm{SIDM}^2}$) leads us to consider warm self-interacting dark matter (WSIDM) in the case of thermal relics to have the correct scaling behaviour ($n_\\mrm{SIDM}\\propto{a^{-3}}$). But this does not rule out a second collisionless CDM component.\\\\ After finally defining our parameter set, we use in Sec.\\ \\ref{sec:constraints} today's DM energy density $\\Omega_\\mrm{DM}^0$ to constrain the parameters characterising the SIDM particle properties and primordial nucleosynthesis limits on an additional energy density contribution to constrain the self-interaction strength $m_\\mrm{SI}/\\!\\sqrt{\\alpha_\\mrm{SI}}$. We find that it depends inversely on the SIDM particle mass ($m_\\mrm{SI}/\\!\\sqrt{\\alpha_\\mrm{SI}}\\propto{1/m_\\mrm{WDM}}$) but can be at least as strong as for strong interactions ($m_\\mrm{SI}/\\!\\sqrt{\\alpha_\\mrm{SI}}\\sim100\\,\\mrm{MeV}$).\\\\ In Sec.\\ \\ref{sec:decoupling} we analyse the consequences on DM decoupling in a universe dominated by the self-interaction energy density contribution. The annihilation cross-section of WSIDM $\\sigma_\\mrm{A}^\\mrm{WDM}$ is inverse proportional to the elastic self-interaction strength ($\\sigma_\\mrm{A}^\\mrm{WDM}\\propto\\sqrt{\\alpha_\\mrm{SI}}/m_\\mrm{SI}$) and rather low ($\\sigma_\\mrm{A}^\\mrm{WDM}\\ll\\sigma_\\mrm{weak}$) while the natural scale for the annihilation cross-section of a collisionless CDM component $\\sigma_\\mrm{A}^\\mrm{CDM}$ exceeds the weak scale and depends beside the self-interaction strength also on the particle mass $m_\\mrm{CDM}$. This casts new light on the {\\textquoteleft}WIMP miracle' and coincides with Fermi-LAT and PAMELA data (e.g.\\ \\cite{Grasso_2009, Bergstroem_2009_B}). We use the unitary bound and neutrino induced constraints on the DM annihilation cross-section to again limit the self-interaction strength.\\\\ Another consequence of an early self-interaction dominated epoch may concern structure formation. We show in Sec.\\ \\ref{sec:struct formation} that a relativistic analysis of linear perturbation theory reveals a linear growing solution $\\delta\\propto{a}$ of self-interaction dominated SIWDM and also of collisionless CDM in a mixed model during self-interaction domination. However, only non-cosmological scales ($M\\lesssim10^{-3}M_\\odot$) can be enhanced and a small observable effect could only be present with fine-tuned parameters.\\\\ Finally we summarize our results in Sec.\\ \\ref{sec:conclusions}. ", "conclusions": "Conclusions} In this paper we have analysed constraints on an energy density contribution of elastic dark matter self-interactions $\\varrho_\\mrm{SI}$, characterized by the mass of the exchanged particle $m_\\mrm{SI}$ and the coupling constant $\\alpha_\\mrm{SI}$.\\\\ The scaling of energy densities implied that the self-interaction contribution decreases as $\\varrho_\\mrm{SI}\\propto{a^{-6}}$ and thus can only have a direct impact on the very early universe. As the energy density scales with the number density squared due to interactions, self-interacting dark matter has to be warm in the case of thermal relics to give the correct scaling behaviour $n_\\mrm{SIDM}\\propto{a^{-3}}$. Note that this does not rule out a second collisionless cold dark matter component.\\\\ We used today's dark matter energy density and the allowed radiation energy density during primordial nucleosynthesis to constrain the parameters characterising the warm self-interacting dark matter particle properties. The dependence of the primordial $^4\\mrm{He}$ abundance on the dark matter self-interaction energy density contribution at neutron to proton number ratio freeze-out allowed to constrain the self-interaction strength $m_\\mrm{SI}/\\!\\sqrt{\\alpha_\\mrm{SI}}$, which depends inversely on the self-interacting dark matter particle mass ($m_\\mrm{SI}/\\!\\sqrt{\\alpha_\\mrm{SI}}\\propto{1/m_\\mrm{WDM}}$) but can be at least as strong as the strong interaction scale ($m_\\mrm{SI}/\\!\\sqrt{\\alpha_\\mrm{SI}}\\sim100\\,\\mrm{MeV}$). Furthermore, our constraint on the dark matter self-interaction strength has a trivial dependence on the relative amount of self-interacting warm dark matter ($m_\\mrm{SI}/\\!\\sqrt{\\alpha_\\mrm{SI}}\\propto{F_\\mrm{WDM}^0}$). We also analyzed dark matter decoupling in a universe dominated by the self-interaction energy density contribution. The annihilation cross-section of warm self-interacting dark matter $\\sigma_\\mrm{A}^\\mrm{WDM}$ is inverse proportional to the elastic self-interaction strength ($\\sigma_\\mrm{A}^\\mrm{WDM}\\propto\\sqrt{\\alpha_\\mrm{SI}}/m_\\mrm{SI}$) and much smaller than $\\sigma_\\mrm{weak}$. The natural scale for the annihilation cross-section of a collisionless cold dark matter component $\\sigma_\\mrm{A}^\\mrm{CDM}$ exceeds the weak scale ($\\sigma_\\mrm{A}^\\mrm{CDM}>\\sigma_\\mrm{weak}$) and depends linearly on the particle mass $m_\\mrm{CDM}$ ($\\sigma_\\mrm{A}^\\mrm{CDM}\\propto{m_\\mrm{CDM}\\times\\sqrt{\\alpha_\\mrm{SI}}/m_\\mrm{SI}}$). This casts new light on the {\\textquoteleft}WIMP miracle' and coincides with the Fermi-LAT and PAMELA data. The unitary bound and neutrino induced constraints on the dark matter annihilation cross-section allowed to disfavour the combination of superstrong elastic warm dark matter self-interactions ($m_\\mrm{SI}/\\!\\sqrt{\\alpha_\\mrm{SI}}\\lesssim1\\,\\mrm{MeV}$) together with very heavy thermal relic cold dark matter particle masses ($m_\\mrm{CDM}\\sim10\\,\\mrm{TeV}$).\\\\ A relativistic analysis of linear perturbation theory reveals a linear growing solution $\\delta\\propto{a}$ of self-interaction dominated warm dark matter and also of collisionless cold dark matter in a mixed model during self-interaction domination. However, only non-cosmological scales ($M\\lesssim10^{-3}M_\\odot$) can be enhanced and a small observable effect could only be present with fine-tuned parameters." }, "1003/1003.5148_arXiv.txt": { "abstract": "We study the evaporation of stars from globular clusters using the simplified Chandrasekhar model [S. Chandrasekhar, Astrophys. J. {\\bf 97}, 263 (1943)]. This is an analytically tractable model giving reasonable agreement with more sophisticated models that require complicated numerical integrations. In the Chandrasekhar model: (i) the stellar system is assumed to be infinite and homogeneous (ii) the evolution of the velocity distribution of stars $f(v,t)$ is governed by a Fokker-Planck equation, the so-called Kramers-Chandrasekhar equation (iii) the velocities $|v|$ that are above a threshold value $R>0$ (escape velocity) are not counted in the statistical distribution of the system. In fact, high velocity stars leave the system, due to free evaporation or to the attraction of a neighboring galaxy (tidal effects). Accordingly, the total mass and energy of the system decrease in time. If the star dynamics is described by the Kramers-Chandrasekhar equation, the mass decreases to zero exponentially rapidly. Our goal is to obtain {\\it non-perturbative} analytical results that complement the seminal studies of Chandrasekhar, Michie and King valid for large times $t\\rightarrow +\\infty$ and large escape velocities $R\\rightarrow +\\infty$. In particular, we obtain an exact semi-explicit solution of the Kramers-Chandrasekhar equation with the absorbing boundary condition $f(R,t)=0$. We use it to obtain an explicit expression of the mass loss at any time $t$ when $R\\rightarrow +\\infty$. We also derive an exact integral equation giving the exponential evaporation rate $\\lambda(R)$, and the corresponding eigenfunction $f_{\\lambda}(v)$, when $t\\rightarrow +\\infty$ for any sufficiently large value of the escape velocity $R$. For $R\\rightarrow +\\infty$, we obtain an explicit expression of the evaporation rate that refines the Chandrasekhar results. More generally, our results can have applications in other contexts where the Kramers equation applies, like the classical diffusion of particles over a barrier of potential (Kramers problem). ", "introduction": "In a seminal paper, Chandrasekhar \\cite{chandra} developed a Brownian theory of stellar dynamics in order to determine the rate of escape of stars from globular clusters. Small groups of stars tend to approach a statistical equilibrium state (described by the Boltzmann distribution) as a result of stellar encounters. However, high energy stars are not bound to the system and escape to infinity. For an isolated system, the average escape velocity for all stars in the cluster is fixed by the virial theorem according to the equation $R=2v_{rms}$ where $v_{rms}=\\langle v^2\\rangle^{1/2}$ is the root-mean-square velocity (RMS) \\cite{bt}. If the system, e.g. a globular cluster, is submitted to tidal forces from a neighboring galaxy, the escape velocity can be smaller. Therefore, stars clusters tend to slowly evaporate. This evaporation was first studied by Ambartsumian \\cite{ambartsumian} and Spitzer \\cite{spitzer} using phenomenological arguments. They estimated the evaporation time by removing a fraction $\\gamma=7.38\\ 10^{-3}$ of stars every relaxation time, where $\\gamma$ is the fraction of particles in a Maxwellian distribution that have speeds exceeding twice the RMS velocity. In a more precise treatment, Chandrasekhar \\cite{chandra} described the ``collisional'' evolution of a stellar system by a Fokker-Planck equation (the nowadays called Kramers equation) involving a diffusion term in velocity space modeling the erratic motion of the stars and a friction term that appears to be necessary to drive the system towards the Boltzmann distribution predicted by statistical mechanics (fluctuation-dissipation theorem). The diffusion coefficient and the ``dynamical friction\", satisfying the Einstein relation, were independently justified from kinetic theory by explicitly calculating the first and second moments of the velocity increment suffered by a star during a succession of binary encounters. In order to account for evaporation, Chandrasekhar imposed as a boundary condition that the distribution function $f(v,t)$ vanishes when the star velocity reaches a maximum value $|v|=R$. He then reduced the problem to the study of an eigenvalue equation in a bounded domain of velocities $|v|\\le R$. The fundamental eigenvalue gives the exponential evaporation rate of the stars from the cluster for $t\\rightarrow +\\infty$ and the associated eigenfunction gives the quasi-stationary distribution function of the system. This distribution function is close to the Boltzmann distribution for low velocities and tends to zero at the escape velocity. Chandrasekhar solved the eigenvalue problem by transforming the Kramers equation into a Schr\\\"odinger equation (with imaginary time) for a quantum oscillator in a box and by expanding the solutions of that equation in the form of a series. He obtained (semi-explicit) analytical results in the $R\\rightarrow +\\infty$ limit or, equivalently, for a small evaporation rate $\\lambda(R)\\rightarrow 0$. In his first treatment \\cite{chandra}, he assumed that the diffusion coefficient is constant and in a more exact theory \\cite{chandramore}, he took into account the dependence of the diffusion coefficient with the velocity. His work was followed by Spitzer \\& H\\\"arm \\cite{sh} who determined the escape rate (eigenvalue) and the quasi-stationary distribution function (eigenfunction) numerically for any value of the escape velocity $R$. Then, Michie \\cite{michie} and King \\cite{king} obtained for $R\\rightarrow +\\infty$ a simple analytical expression of the quasi-stationary distribution function in the form of a lowered isothermal distribution which vanishes at the escape velocity. This leads to the so-called Michie-King model \\cite{bt} that is asymptotically valid in the limit $R\\rightarrow +\\infty$. The Chandrasekhar model described previously is based on simplifying assumptions. It is first assumed that the system is spatially homogeneous and infinite while globular clusters are highly inhomogeneous and limited in space. On the other hand, the collisional evolution of the system is modeled by the Kramers equation while a more relevant equation is the gravitational Landau equation that is the standard kinetic equation of stellar dynamics \\cite{bt,spitzerbook,paddy,epjb}. The Kramers equation corresponds to a canonical description in which the system is assumed to be in contact with a thermal bath from which it can extract energy so that the temperature $T$ is fixed. Alternatively, the Landau equation corresponds to a microcanonical description in which the system is assumed to be isolated so that the energy $E$ is conserved. Since globular clusters are isolated Hamiltonian systems (up to the slow evaporation process), the microcanonical description appears to be more relevant. Therefore, when we take into account spatial inhomogeneity and model the encounters in a self-consistent way, the proper model to consider is formed by the gravitational Landau equation coupled to the Poisson equation. In order to go beyond the limitations of the Chandrasekhar model and obtain more accurate rates of escape, the astrophysicists have performed numerical simulations of stellar systems. Different types of simulations have been performed. They solved (i) the $N$-body Hamiltonian problem associated to the Newton equations \\cite{al} (ii) the hydrodynamic moments of the Landau equation \\cite{larson,lbe} (iii) the $N$-body problem where the effect of encounters is modeled by Monte Carlo methods \\cite{henonmc,spitzermc,shapiro}, and (iv) the orbit-averaged-Fokker-Planck equation \\cite{cohn}. These methods are reviewed in the books of Spitzer \\cite{spitzerbook} and Binney \\& Tremaine \\cite{bt} and in the reviews \\cite{mh,ls}. In these numerical works, the spatial inhomogeneity of the cluster is properly taken into account. These simulations have led to the following scenario for the evolution of globular clusters \\cite{bt,spitzerbook}. In a first regime, a self-gravitating system initially out-of-mechanical equilibrium undergoes a process of {\\it violent collisionless relaxation} towards a virialized state\\footnote{This form of relaxation is appropriate to account for the actual structure of elliptical galaxies whose dynamics is encounterless for the timescales of interest \\cite{bt}.}. In this regime, the dynamical evolution of the cluster is described by the Vlasov-Poisson system and the phenomenology of violent relaxation has been described by H\\'enon \\cite{henonvr}, King \\cite{kingvr} and Lynden-Bell \\cite{lb}. Numerical simulations that start from cold and clumpy initial conditions generate a quasi stationary state (QSS) that fits the de Vaucouleurs $R^{1/4}$ law quite well \\cite{vanalbada}. The inner core is almost isothermal (as predicted by Lynden-Bell \\cite{lb}) while the velocity distribution in the envelope is radially anisotropic and the density profile decreases like $r^{-4}$ \\cite{sb,hm}. One success of Lynden-Bell's statistical theory of violent relaxation is to explain the isothermal core without recourse to ``collisions''. By contrast, the structure of the halo cannot be explained by Lynden-Bell's theory as it is a result of an {\\it incomplete relaxation}. On longer timescales, encounters between stars must be taken into account and the dynamical evolution of the cluster is governed by the Vlasov-Landau-Poisson system which is the standard model of stellar dynamics. This collisional regime is appropriate to understand the actual structure of globular clusters. In this regime, the system passes through a succession of quasi stationary states (QSS) that are steady states of the Vlasov equation slowly evolving in time due to the cumulative effect of encounters. The first stage of the collisional evolution is driven by {\\it evaporation}. Due to a series of weak encounters, the energy of a star can gradually increase until it reaches the local escape energy; in that case, the star leaves the system\\footnote{There can also be a process of {\\it ejection} \\cite{henonejection} in which a single close encounter produces a velocity change that is sufficient to eject the star out of the cluster. However, it can be shown that this process is less efficient than evaporation.}. Numerical simulations \\cite{spitzerbook} show that during this regime the system reaches a quasi-stationary state that slowly evolves in amplitude due to evaporation as the system loses mass and energy. This quasi stationary distribution function (DF) is close to the Michie-King model. The system has a core-halo structure. The core is isothermal while the stars in the outer halo move in predominantly radial orbits. Therefore, the distribution function in the halo is anisotropic. The density follows the isothermal law $\\rho\\sim r^{-2}$ in the central region (with a core of almost uniform density) and decreases like $\\rho\\sim r^{-7/2}$ in the halo \\cite{bt}. Due to evaporation, the halo expands while the core shrinks as required by energy conservation. At some point of the evolution, when the energy passes below a critical value (or when the density contrast becomes sufficiently high), the system undergoes an instability related to the Antonov \\cite{antonov} instability and the gravothermal catastrophe takes place \\cite{lbw}. This instability is due to the negative specific heat of the inner system that evolves by losing energy and thereby growing hotter (see reviews in \\cite{paddy,revue}). This leads to {\\it core collapse} \\cite{bt}. Mathematically speaking, core collapse would generate a finite time singularity: if the evolution is modeled by the orbit-averaged-Fokker-Planck equation, Cohn \\cite{cohn} finds that the collapse is self-similar, that the central density becomes infinite in a finite time and that the density behaves like $\\rho\\sim r^{-2.23}$ (if the evolution is modeled by the Landau-Poisson system, it is argued in \\cite{lk} that the density behaves like $\\rho\\propto r^{-3}$ in the final stage of the collapse). In reality, if we come back to the $N$-body system, core collapse is arrested by the formation of binary stars. These binaries can release sufficient energy to stop the collapse \\cite{henonbinary} and even drive a re-expansion of the cluster in a post-collapse regime \\cite{inagakilb}. Then, in principle, a series of gravothermal oscillations should follow \\cite{bettwieser}. In practice, the processes of evaporation and core collapse take place simultaneously so that it is difficult to isolate the effect of any single process in the evolution of a globular cluster. Concerning the evaporation process, the Princeton code was the first code to yield reliable evaporation rates \\cite{princeton} giving $t_{evap}=-N(dN/dt)^{-1}\\simeq 300 t_{rh}$ for isolated clusters\\footnote{Tidal forces from the Galaxy can increase the evaporation rate \\cite{chevalier}.}. These results are not very different from those obtained with the spatially homogeneous Kramers-Chandrasekhar equation. Our goal in this paper is not to make a realistic modeling of stellar systems but rather to consider simple models of evaporation that are {\\it analytically tractable}. Therefore, we shall use the Chandrasekhar model which yields a reasonable description of the evaporation process in globular clusters and which can be studied analytically. Chandrasekhar solved the problem perturbatively: he first considered the long time limit $t\\rightarrow +\\infty$ so that the distribution function $f_R(v,t)$ is dominated by the contribution of the fundamental eigenmode $f_R(v)e^{-|\\lambda(R)| t}$ and then took the limit $R\\rightarrow +\\infty$ to obtain an approximate expression of the quasi-stationary distribution $f_R(v)$ (fundamental eigenfunction) and escape rate $\\lambda(R)$ (fundamental eigenvalue). In this paper, we shall reconsider the Chandrasekhar problem on a new angle which allows to obtain {\\it non-perturbative} results. In particular, we find an exact semi-explicit solution of the Kramers equation with boundary condition $f(v,t)=0$ when $|v|=R$. This solution $f(v,t)$ depends on the remaining mass in the cluster $M_0(t)$ which satisfies an autonomous equation. We use this general formula to obtain: (i) the mass $M_0(t)$ for any fixed time $t$ in the limit $R\\rightarrow +\\infty$, (ii) an exact integral equation for the eigenvalue $\\lambda(R)$ of the fundamental mode (evaporation rate) valid for any sufficiently large $R$, (iii) an exact explicit expression of the fundamental eigenfunction valid for any sufficiently large $R$, and (iv) an explicit asymptotic expression of the evaporation rate when $R\\rightarrow +\\infty$. Therefore, our approach complements Chandrasekhar's original work and offers new perspectives. Our main results are, however, restricted to the Kramers equation, i.e. a Fokker-Planck equation with constant diffusion coefficient and quadratic potential (linear friction). A different approach that goes beyond these limitations (but which is restricted to the asymptotic limits $t\\rightarrow +\\infty$ and $R\\rightarrow + \\infty$) is developed in Appendix \\ref{sec_comp}. Let us finally note that our approach is not limited to the astrophysical problem mentioned above but that it can have applications in different area. First of all, Chandrasekhar's study of the rate of escape of stars from globular clusters is closely connected to the classical Kramers \\cite{kramers} problem for the escape rate of a Brownian particle across a potential barrier that has many applications in physics and chemistry (surprisingly, Chandrasekhar \\cite{chandra} did not mention this connection). In that case, the problem is usually formulated in $d=1$ dimension. On the other hand, Chandrasekhar's procedure has been used in the context of planet formation \\cite{vortex} in order to determine the rate of escape of dust from large-scale vortices (assumed to be present in the solar nebula) due to turbulence. In that case, the problem is two-dimensional. In view of the fundamental nature of the mathematical problem, it seems relevant to develop our formalism in arbitrary dimension of space $d$ in order to cover a wide range of possible applications. ", "conclusions": "\\label{sec_conclusion} In this paper, we have obtained new analytical results for the escape of stars from globular clusters in the framework of the Chandrasekhar model. These results can also have applications to other physical systems described by the Kramers equation with parabolic potential and absorbing boundary conditions (i.e. the classical Kramers problem). We have first obtained an autonomous equation for the mass loss [see \\fref{equation-masse}] and a semi-explicit expression of the distribution function $f(v,t)$ [see \\fref{solution-exacte-d}] that are valid for an arbitrary escape velocity $R$ and time $t$. We have simplified the expression of the mass loss in the limit $R\\rightarrow +\\infty$ for any fixed time $t$ [see \\fref{M0Rgrand}]. We have also used the semi-explicit expression of the distribution function $f(v,t)$ to obtain an exact integral equation for the fundamental eigenvalue $\\lambda(R)$ [see \\fref{relation-lambda}] and for the fundamental eigenfunction $f_{\\lambda}(v)$ [see \\fref{eigenfunction}] that are valid for any sufficiently large $R$. This is an interesting complement to the perturbative results derived by Chandrasekhar \\cite{chandra} that are valid in the asymptotic limit $R\\rightarrow +\\infty$. We have obtained the explicit behavior of the fundamental eigenvalue in the limit $R\\rightarrow +\\infty$ [see \\fref{lambda-asymptotique}] which improves upon the result of Chandrasekhar expressed in the form of a series. Additional asymptotic results for the fundamental eigenvalue and fundamental eigenfunction are given in Appendix \\ref{sec_comp}. Finally, we have illustrated our results with numerical simulations [see Sec. \\ref{sec_figures}]. Of course, our approach is based on several approximations. As previously discussed, it assumes that the medium is infinite and spatially homogeneous and that the encounters between stars can be described by the Kramers equation (Chandrasekhar's model). Furthermore, our analytical results (except those of Appendix \\ref{sec_comp}) are valid only when the diffusion coefficient in the Fokker-Planck equation is constant while a more exact description of encounters between stars would involve a velocity dependent diffusion coefficient \\cite{chandramore,michie,king}. It remains therefore a challenging issue to extend the mathematical theory of the escape of stars from globular clusters in the case of more realistic models. \\vskip0.2cm \\begin{acknowledgement} M. Lemou was supported by the French Agence Nationale de la Recherche, ANR JC MNEC. \\end{acknowledgement} \\appendix" }, "1003/1003.0779_arXiv.txt": { "abstract": "Rapidly rotating neutron stars can be unstable to the gravitational-wave-driven CFS mechanism if they have a neutral point in the spectrum of nonaxisymmetric $f$-modes. We investigate the frequencies of these modes in two sequences of uniformly rotating polytropes using nonlinear simulations in full general relativity, determine the approximate locations of the neutral points, and derive limits on the observable frequency band available to the instability in these sequences. We find that general relativity enhances the detectability of a CFS-unstable neutron star substantially, both by widening the instability window and enlarging the band into the optimal range for interferometric detectors like LIGO, VIRGO, and GEO-600. ", "introduction": " ", "conclusions": "" }, "1003/1003.2917_arXiv.txt": { "abstract": "Explosions of sub-Chandrasekhar-mass white dwarfs are one alternative to the standard Chandrasekhar-mass model of Type Ia supernovae. They are interesting since binary systems with sub-Chandrasekhar-mass primary white dwarfs should be common and this scenario would suggest a simple physical parameter which determines the explosion brightness, namely the mass of the exploding white dwarf. Here we perform one-dimensional hydrodynamical simulations, associated post-processing nucleosynthesis and multi-wavelength radiation transport calculations for pure detonations of carbon-oxygen white dwarfs. The light curves and spectra we obtain from these simulations are in good agreement with observed properties of Type Ia supernovae. In particular, for white dwarf masses from 0.97--1.15~M$_{\\odot}$ we obtain $^{56}$Ni masses between 0.3 and 0.8~M$_{\\odot}$, sufficient to capture almost the complete range of Type Ia supernova brightnesses. {Our optical light curve rise times, peak colours and decline timescales display trends which are generally consistent with observed characteristics although the range of $B$-band decline timescales displayed by our current set of models is somewhat too narrow.} In agreement with observations, the maximum light spectra of the models show clear features associated with intermediate mass elements and reproduce the sense of the observed correlation between explosion luminosity and the ratio of the Si~{\\sc{ii}} lines at $\\lambda6355$ and $\\lambda5972$. We therefore suggest that sub-Chandrasekhar mass explosions are a viable model for Type Ia supernovae for any binary evolution scenario {leading to explosions} in which the optical display is dominated by the material produced in {a detonation} of the primary white dwarf. ", "introduction": "In recent years, considerable work has been devoted to the study of the Chandrasekhar-mass ($M_\\mathrm{Ch}$) model of Type~Ia supernovae (SNe~Ia). As shown by \\cite{arnett1971a}, prompt detonations of $M_\\mathrm{Ch}$ carbon/oxygen (C+O) white dwarfs (WDs) in hydrostatic equilibrium mainly produce iron group elements (IGEs). Thus, they cannot account for the significant amounts of intermediate-mass elements (IMEs; e.g.\\ silicon and sulphur) responsible for the features which dominate the maximum light spectra. To obtain these, pre-expansion of the WD is necessary such that burning partially takes place under low-density conditions where IMEs can be synthesized. One way of achieving this is provided by models in which the flame ignites as a deflagration which releases sufficient energy to expand the star before a deflagration-to-detonation transition occurs {(\\citealt{khokhlov1991a})}. An alternative to this pre-expansion is the detonation of a sub-Chandrasekhar mass (sub-$M_\\mathrm{Ch}$) WD starting from a hydrostatic configuration. Here, a variety of density profiles can be realized, determined by the WD mass. Close to $M_\\mathrm{Ch}$, the detonation produces primarily IGEs and few IMEs, while for less massive WDs more IMEs and less IGEs will be synthesized. Detonation of a sub-$M_\\mathrm{Ch}$ WD cannot occur spontaneously but must be triggered by external compression. The most widely discussed mechanism for sub-$M_\\mathrm{Ch}$ explosions has been the \\emph{double detonation} model. Here, a C+O~WD accretes from a companion star and develops a helium-rich outer shell. This may occur for binaries with helium-rich donors or hydrogen-rich donors where the accreted hydrogen is burned to helium. If the helium-shell becomes sufficiently massive, it can become unstable and detonate. Subsequent compression of the core by inward propagating shocks may produce a secondary carbon detonation which explodes the WD (e.g. \\citealt{woosley1986a,fink2007a}). Detonations in helium-rich surface layers have also been discussed for the case of rapid dynamical mass transfer in binary systems containing a C+O~WD with a helium-rich WD companion \\citep{guillochon2009}. In that case instabilities in the accretion seed dense knots which, by impacting on the underlying WD surface, might trigger a detonation in the accreted helium leading to a potential secondary core detonation. It has also been speculated that sub-$M_\\mathrm{Ch}$ explosions may arise during violent accretion in mergers of C+O~WD binaries. Here, the C+O accretion may lead to an edge-lit detonation or carbon flashes that trigger a core detonation (see e.g.\\ \\citealt{shigeyama1992} but for a different result see \\citealt{loren2009b}). Most previous work on testing sub-$M_\\mathrm{Ch}$ models has focused on cases in which the core detonation is triggered by detonation in an overlying massive shell ($\\sim0.2{\\;}\\mbox{M}_{\\odot}$) of helium (e.g. \\citealt{woosley1994b, livne1995a, hoeflich1996a,hoeflich1996b,nugent1997}). In those models burning in the helium shell synthesizes significant masses of $^{56}$Ni in the outer ejecta, leading to spectra and light curves in conflict with observations. As noted in those studies, however, these conclusions are strongly dependent on the influence of the shell material. In particular, they may not be applicable if such a layer is absent (or much less massive) or if its post-burning composition lacks $^{56}$Ni. Recently, \\cite{bildsten2007a} suggested that detonation of the helium-shell may be possible for a shell with mass as low as $\\sim0.055{\\;}\\mbox{M}_{\\odot}$ around a $1.025{\\;}\\mbox{M}_{\\odot}$ C+O core and that the burning produces only $0.012{\\;}\\mbox{M}_{\\odot}$ of $^{56}$Ni along with some lighter IGEs (\\citealt{guillochon2009} find that even lower atomic-number burning products can dominate in their helium detonations). Even for the low shell masses of \\cite{bildsten2007a}, \\cite{fink2010} find that a secondary core detonation is possible. To date, sub-$M_\\mathrm{Ch}$ explosions in the absence of a nickel-rich outer layer have not been studied in detail. \\cite{shigeyama1992} investigated the explosion dynamics of sub-$M_\\mathrm{Ch}$ detonations and concluded that their characteristic properties were consistent with SNe~Ia but they did not perform realistic radiative transfer simulations. A full treatment of any class of sub-$M_\\mathrm{Ch}$ explosion model requires realistic hydrodynamical and nucleosynthesis simulations of the accretion phase, triggering mechanism and subsequent explosion. Here, however, we present a simple numerical experiment that is relevant to any class of sub-$M_\\mathrm{Ch}$ explosion model. We consider pure detonations of sub-$M_\\mathrm{Ch}$ C+O~WDs with different masses, neglecting the question of how this detonation is initiated. This allows us to investigate the idealized case of sub-$M_\\mathrm{Ch}$ detonation scenarios in which the observable display is dominated by material produced in the core explosion. Our goal is to determine the extent to which the least ambiguous component of the system, namely the detonation of a sub-$M_\\mathrm{Ch}$ C+O~WD, could lead to explosions which are consistent with observations of SNe~Ia. ", "conclusions": "The sub-$M_\\mathrm{Ch}$ model for SNe~Ia has much to commend it. First, it has already been suggested by empirical modelling of bolometric SNe~Ia light curves \\citep{stritzinger2006a} that differing ejecta masses may be required for different SNe~Ia, a property which the sub-$M_\\mathrm{Ch}$ model may explain. Secondly, population synthesis studies predict large numbers of binary systems with accreting C+O~WDs: \\cite{ruiter2009a} estimate a Galactic rate of $\\sim10^{-3}{\\;}\\mbox{yr}^{-1}$ for possible explosions of sub-$M_\\mathrm{Ch}$ C+O~WDs accreting from helium-rich companions. This is much higher than their estimate of the Galactic rate for single-degenerate $M_\\mathrm{Ch}$ explosions ($0.6-1.4\\times10^{-4}{\\;}\\mbox{yr}^{-1}$) {and comparable to their estimate of the WD-WD merger rate ($1-2\\times10^{-3}{\\;}\\mbox{yr}^{-1}$) in systems that exceed $M_\\mathrm{Ch}$. For comparison, the observed Galactic rate of SNe~Ia is $(4\\pm2)\\times10^{-3}{\\;}\\mbox{yr}^{-1}$ (\\citealt{cappellaro99}).} Moreover, sub-$M_\\mathrm{Ch}$ models provide a simple physical parameter which could account for the range of observed brightnesses: the mass of the exploding C+O~WD\\@. This parameter allows for a possible link between the typical brightness of a SN Ia and the stellar population in which it resides. For example, if it can be shown that explosions in binary systems with larger $M_\\mathrm{WD}$ are more often found among young stellar populations relative to their less massive $M_\\mathrm{WD}$ counterparts, the observed correlation of SN Ia brightness with host galaxy type (e.g., \\citealt{howell2001b}) might be explained. Here we have shown that detonations of sub-$M_\\mathrm{Ch}$~WDs lead to explosions which give a reasonable match to several properties of SNe~Ia. Specifically, WDs with masses between $\\sim1$ and $\\sim1.2{\\;}\\mbox{M}_{\\odot}$ can reproduce a wide range of brightness with light curves that have rise times and peak colours in roughly the correct range. In addition, the models reproduce the characteristic spectral features present around maximum light and the observed trend for a higher velocity at the inner boundary of the IME-rich layer in brighter SNe~Ia \\citep{mazzali2007a}. Although our {pure-C+O models} yield light curves that fade too fast after maximum, {the models predict a width-luminosity relation which behaves in the observed sense and we would argue that the combination of uncertainties in radiative transfer simulations and details of the nucleosynthesis (which is sensitive to the progenitor composition) can systematically affect the decline timescale.} Thus there is potential for even better agreement with improved modelling. {There are several additional observational constraints that our current models do not address but which should be considered in future studies. For example, off-centre detonation might lead to observable effects associated with departures from spherical symmetry (e.g. \\citealt{fink2010}). Chemical inhomogeneity of the pre-explosion WD could affect the explosive nucleosynthesis: in particular, significant gravitational settling of $^{22}$Ne \\citep{bildsten2001,garcia2008} might yield a layered ejecta structure with a central concentration of neutron-rich isotopes as favoured by observations (e.g. \\citealt{hoeflich2004a,gerardy2007a}).} In conclusion, detonations of naked sub-$M_\\mathrm{Ch}$ C+O~WDs yield light curves and spectra which are in qualitatively good agreement with the observed properties of SNe~Ia. The critical question remains whether or not realistic progenitor scenarios in which the optical display is dominated by such an explosion can be established: {it must involve detonation of a WD with a density profile similar to those of our toy models without producing large masses of high-velocity IGEs.} Any sub-$M_\\mathrm{Ch}$ scenario which meets these criteria will likely be promising in accounting for the observed characteristics of SNe~Ia." }, "1003/1003.0829.txt": { "abstract": "We use two independent methods to reduce the data of the surveys made with RATAN-600 radio telescope at 7.6\\,cm in 1988--1999 at the declination of the SS433 source. We also reprocess the data of the ``Cold'' survey (1980--1981). The resulting RCR (RATAN COLD REFINED) catalog contains the right ascensions and fluxes of objects identified with those of the NVSS catalog in the right-ascension interval $7^h \\le$ R.A. $ < 17^h$. We obtain the spectra of the radio sources and determine their spectral indices at 3.94 and 0.5\\,GHz. The spectra are based on the data from all known catalogs available from the CATS, Vizier, and NED databases, and the flux estimates inferred from the maps of the VLSS and GB6 surveys. For 245 of the 550 objects of the RCR catalog the fluxes are known at two frequencies only: 3.94\\,GHz (RCR) and 1.4\\,GHz (NVSS). These are mostly sources with fluxes smaller than 30\\,mJy. About 65\\% of these sources have flat or inverse spectra ($\\alpha > -0.5$). We analyze the reliability of the results obtained for the entire list of objects and construct the histograms of the spectral indices and fluxes of the sources. Our main conclusion is that all 10--15\\,mJy objects found in the considered right-ascension interval were already included in the decimeter-wave catalogs. ", "introduction": "In 1991 by Parijskij et al.~\\cite{pa1:Soboleva_n} a catalog of 7.6-cm radio sources in the \\mbox{$4^h \\le$ R.A. $< 22^h$} interval of right ascensions (RC-catalog) and 31-cm radio sources in the \\mbox{$4^h \\le$ R.A. $< 13^h$} interval at the declination of \\mbox{$Dec_{2000}=5\\degr \\pm 20'$} was published. This catalog was based on the results of observations made in \\mbox{1980--1981} with the RATAN-600 radio telescope in the meridian and azimuth of $30\\degr$ \\cite{h:Soboleva_n,ber:Soboleva_n,pa2:Soboleva_n}. The coordinate calibration was based on the then most accurate UTRAO (365 MHz) catalog\\footnote{The data for the sky strip of interest was kindly provided to us by Prof.~J.~N.~Douglas prior to publication.}. After the publication of the NVSS and FIRST catalogs (1.4 GHz, VLA) \\cite{co1:Soboleva_n,fr:Soboleva_n} the objects of the RC catalog were compared with the objects of the former two catalogs and 20--25\\% of RC objects proved to be impossible to cross identify with NVSS objects \\cite{zh:Soboleva_n}. Several additional observation sets were carried out from 1987 through 1999 at the Northern sector of RATAN-600 in order to refine the RC catalog and, in particular, the fluxes and coordinates of its sources. The observations, like earlier, were made at the declination of SS433. The declination varied from cycle to cycle because of precession. However, these variations proved to be too small to allow the declinations of RCR objects to be found with sufficient accuracy. The results of the reduction of these observations in the right-ascension interval \\mbox{$2^h \\le {\\rm R.A.} < 7^h$} and \\mbox{$17^h \\le {\\rm R.A.} < 22^h$} can be found in our earlier paper~\\cite{so1:Soboleva_n}. In this paper we report the results of the reduction of the 7.6-cm observations made in 1987--1999 in the right-ascension band $7^h \\le {\\rm R.A.}< 17^h$. In addition, we also report the results of our rereduction of the records obtained in the ``Cold'' experiment in 1980. We uses NVSS objects to calibrate the right ascensions. We also used the declinations of sources from NVSS catalog. Almost the entire observed region has been studied with a VLA with a resolution of 5.4$''$ (the FIRST catalog, $ 8^{h}11^{m} \\le {\\rm R.A.} < 16^{h}26^{m}$), and the results of these observations allowed us to refine the structure of the radio sources. We report the list of objects (the RCR catalog) found within the right-ascension band mentioned above and identified with NVSS objects~\\cite{co1:Soboleva_n}. We separately discuss the reliability of the identification of our objects with the objects of the NVSS catalog and the statistical conclusions based on the rereduced RC catalog. We pay special attention to objects with peculiar spectra and to objects with fluxes known only at two frequencies: \\mbox{1.4\\,GHz} (NVSS) and 3.94 GHz (RCR). ", "conclusions": "We used RATAN-600 observations of the sky band at the declination $Dec \\sim 5\\degr$ in the right-ascen\\-sion interval $7^h \\le$ R.A. < $17^h$ in 1987--1999 combined with reprocessing of the ``Cold'' survey data (1980--1981) to compile a list (the RCR catalog) of 550 objects identified with objects of the NVSS catalog. This list includes 18 blends and 15 double sources. Data reduction was performed using two independent methods. We determined the fluxes, right-ascensions, and spectral indices of every object of our list and constructed the histograms of spectral indices and fluxes for different samples of sources \\footnote{We did not determine the declinations of the sources but adopted them from the NVSS catalog.}. We reconstructed the spectra using all the catalogs available from the CATS, Vizier, and NED databases and the flux estimates based on the maps of the VLSS and GB6 surveys. These estimates are useful primarily for the reconstruction of the sources spectra for that fluxes are known only at two frequencies: 3.94~GHz (RCR) and 1.4~GHz (NVSS). Such objects make up for about 50\\% of the RCR catalog (245 objects). These are mostly sources with fluxes not exceeding $30$\\,Jy, about 65\\% of them have flat or inverse spectra ($\\alpha > -0.5$). The histograms of the fluxes and spectral indices and the average spectral index of this sample of sources are affected by selection, since objects with steep spectra cannot be detected due to the limited sensitivity of the survey. We analyzed the reliability of the results obtained. We demonstrated that use of two different methods of data reduction yields more accurate results both for the right ascensions and fluxes of the sources. We therefore believe that the technique of independent reduction of observational data has proved to be \\mbox{successful.} Note in conclusion that the study of spectral indices at centimeter-wave frequencies is closely linked to the problem of eliminating selection effects. For such studies, i.e., for the complete analysis of the spectra of objects in deep decimeter-wave sky surveys (NVSS \\cite{co1:Soboleva_n}, FIRST \\cite{backer:Soboleva_n}), the sensitivity at centimeter-wave frequencies should be one to two orders of magnitude higher than the sensitivity of decimeter-wave surveys, i.e., at the level of or better than several tens of $\\mu$Jy. So far, such a high sensitivity at centimeter-wave frequencies could have been achieved only within small sky areas. The aim of sky surveys made with RATAN-600 radio telescope \\cite{h:Soboleva_n,bu:Soboleva_n} is to obtain more comprehensive information about the spectral indices of decimeter-wave sources. These surveys serve as an intermediate link between deep VLA surveys and low-sensitivity all-sky surveys. The main and very important conclusion of such surveys is that we found no objects within the right ascension interval considered---at least at the 10--15~mJy level---that had not been previously included into decimeter-wave catalogs. Ninety per cent of the RCR sources identified with NVSS objects and having fluxes ${\\rm F_{1.4}} \\ge $ 100 mJy lie in the $\\Delta Dec = \\pm6'$ band of the survey. In the same band 72\\% and 16\\% of the sources were identified with NVSS objects with fluxes \\mbox{20 mJy < ${\\rm F_{1.4}}$<100 mJy} and ${\\rm F_{1.4}} \\le $ 20 mJy, respectively. In the broader band \\mbox{$\\Delta Dec = \\pm20'$} the number of sources identified with NVSS objects decreases down to 70\\%, 37\\%, and 6\\%, respectively. Thus so far for the small population of NVSS objects all the centimeter-wave surveys including ``Cold'' and RZF provide the data mostly for objects with synchrotron self absorption (BL Lac, QSR, AGN). The new epoch in this field is expected to start the operation of ALMA and SKA instruments at millimiter and about 1-cm waves, respectively. The catalog is available at \\linebreak {\\tt http://cdsavc.u-strasbg.fr./viz-bin/\\linebreak cat?J/other/AstBu/65.42}." }, "1003/1003.0409_arXiv.txt": { "abstract": "Differences in masses inferred from dynamics, such as velocity dispersions or X-rays, and those inferred from lensing are a generic prediction of modified gravity theories. Viable models however must include some non-linear mechanism to restore General Relativity (GR) in dense environments, which is necessary to pass Solar System constraints on precisely these deviations. In this paper, we study the dynamics within virialized structures in the context of two modified gravity models, $f(R)$ gravity and DGP. The non-linear mechanisms to restore GR, which $f(R)$ and DGP implement in very different ways, have a strong impact on the dynamics in bound objects; they leave distinctive signatures in the dynamical mass-lensing mass relation as a function of mass and radius. We present measurements from N-body simulations of $f(R)$ and DGP, as well as semi-analytical models which match the simulation results to surprising accuracy in both cases. The semi-analytical models are useful for making the connection to observations. Our results confirm that the environment- and scale-dependence of the modified gravity effects have to be taken into account when confronting gravity theories with observations of dynamics in galaxies and clusters. ", "introduction": "\\label{sec:intro} Gravity, as described by General Relativity (GR), is remarkably weakly constrained in the present day on scales larger than a few AU. Though measurements from binary pulsar timing to the cosmic microwave background (CMB) and big bang nucleosynthesis are all consistent with GR, there is still room for order unity deviations in the cosmos today, on scales of kpc and larger. Thus, testing gravity on cosmological scales is an interesting frontier, and the focus of much current research \\cite{ZhangEtal,JainZhang,SongKoyama,KnoxSongTyson,Schmidt08,Song2006,JainZhang,Tsujikawa2008,ZhangfR,SongEtalDGP,Schmidt07,Uzan09}. Any gravity theory that attempts to be complete has to satisfy stringent Solar System constraints, and has to locally match the predictions of GR to within one part in $10^5$ there. Only a few consistent models which modify gravity appreciably on large scales, but restore GR locally are known. Two of them will be the subject of this study: $f(R)$ gravity \\cite{Caretal03,NojOdi03,Capozziello:2003tk,Sotiriou:2008rp}, and the DGP model \\cite{DGP1}. Within certain bounds placed by the CMB and expansion history measurements in addition to Solar System tests, both theories can be made to satisfy all current constraints on gravity (including the observation of an accelerating expansion). In both models there exists a non-linear mechanism to restore GR in high-density environments: the \\textit{chameleon effect} for $f(R)$, and the \\textit{Vainshtein mechanism} for DGP. Furthermore, all currently known consistent modifications of gravity on large scales include some variant of either of these mechanisms. In order to be able to constrain these models with cosmological data, it is crucial to correctly include the non-linear mechanisms. Recently, N-body simulations of $f(R)$ \\cite{HPMpaper} and DGP \\cite{DGPMpaper,ScII,DGPMpaperII} have been done which self-consistently solve the non-linear field equations together with the growth of structure (see also \\cite{KW} for the first study of the DGP case, using a different approach). In principle, it has become possible with these simulations to unlock the wealth of observations available on non-linear scales to probe gravity, albeit in a necessarily model-dependent way. It is well known that the additional degrees of freedom present in modified gravity theories generically affect the {\\it dynamical} potential, which governs the propagation of non-relativistic bodies, differently than the {\\it lensing} potential, which governs the propagation of massless particles such as light (e.g., \\cite{BekSan94}). Thus, comparing dynamical with lensing mass estimates is an interesting and quite generic probe of modifications to gravity. In this paper, we study the signatures of $f(R)$ and DGP in dynamical observables such as velocity dispersions, compared to lensing which measures essentially the ``true'' mass (i.e. the integral over the rest-frame density) in both models. Constraints on the difference between dynamical and lensing potential are often phrased in terms of the post-Newtonian parameter $\\gamma_{\\rm PPN}$ [\\refeq{gamma} below], in analogy to Solar System tests. In general, however, the departures from GR cannot be encapsulated by a single parameter but are functions of scale, time and the local environment. In particular, this is the case for both $f(R)$ and DGP. Hence, we introduce a more generally applicable quantity $\\g$ [\\refeq{gdef}] which is defined directly via the modified forces, and is well suited for predictions in the context of $f(R)$ and DGP as well as for constraints from observations. Velocities of extragalactic objects are measured through their redshifts $z$, which receive a contribution $|\\D z| = v_{\\parallel}/c$ from the line-of-sight velocity $v_{\\parallel}$. In the cosmological context, there are two regimes where the dynamics of matter can be understood fairly easily: on very large scales, linear perturbation theory in the matter density is valid, simplifying the theoretial predictions. Large-scale velocity fields can be measured through the redshift distortion of the power spectrum, which thus offers a probe of the dynamical potential \\cite{Sc04,ZhangLBD07}. On small scales, most of the observable matter lies in gravitationally bound dark matter halos. In this regime, for relaxed systems, the velocity distribution of collisionless objects such as dark matter, galaxies or stars is related to the dynamical potential by the virial theorem. For collisional particles such as diffuse gas, this relation is given by hydrostatic equilibrium. The virial or thermal velocities can be observed as velocity dispersion of stars in galaxies, galaxies in clusters, or as X-ray or Sunyaev-Zeldovich signal from diffuse gas in clusters. Also, the redshift-space matter power spectrum on small scales is a probe of virial velocities \\cite{PD96,Sc04}. This paper is concerned with the latter regime, and our goal is to study the dynamics of matter in halos. Since these are highly non-linear systems, rigorous results can only be obtained via N-body simulations. We therefore present measurements from the modified gravity simulations of $f(R)$ and DGP \\cite{HPMpaper,DGPMpaper,DGPMpaperII}. However, for many practical purposes including comparison with observations, it is necessary to go beyond the simulation results which have limited resolution and cover only a few points in the parameter space of the models. Thus, a sufficiently accurate semi-analytic model of the dynamics in modified gravity is desirable to bridge the gap with observations. Fortunately, we can make some justified assumptions which simplify the problem greatly: first, since we are concerned with sub-horizon scales, we employ the quasi-static approximation, neglecting time derivatives and assuming the halos are in steady state. Further, we assume spherically symmetric halos. While certainly not realistic, deviations from spherical symmetry are not expected to affect the results qualitatively. Throughout, we will assume a Navarro-Frenk-White (NFW) \\cite{NFW} profile, although all derivations can easily be generalized to different profiles. The problem is then reduced to finding the solution of the field equations for a spherically symmetric mass, and calculating the modified gravitational force. The accuracy of this simplified model can then be benchmarked with the simulation results. The paper is structured as follows. In \\refsec{th}, we introduce our main observable, the modified gravitational force strength, and present the theoretical expectations and semi-analytic models for $f(R)$ and DGP. \\refsec{sim} contains the simulation results and comparisons with the theoretical models. We then discuss the application to observations in \\refsec{obs}. We conclude in \\refsec{concl}. ", "conclusions": "\\label{sec:concl} In this paper, we have studied the dynamics of matter within bound cosmic structures, i.e. dark matter halos, in $f(R)$ and DGP. The potential governing matter dynamics can differ from the lensing potential by $20-30$\\% in these models. These unique signatures of modified gravity can be observed by comparing dynamical and lensing mass estimates of clusters or galaxies. Furthermore, they strongly influence the observed abundance of massive clusters when measured via dynamical mass proxies such as X-rays or the SZ effect. For example, the enhancement of the cluster abundance in $f(R)$ (with respect to $\\L$CDM) at a fixed {\\it dynamical} mass can be roughly twice that measured if the mass is based on lensing measurements. These signatures in the dynamics are also relevant for large-scale structure observations, such as the redshift-space power spectrum or correlation function on small scales. However, since halos are highly non-linear objects, the peculiar chameleon and Vainshtein mechanisms play a crucial role, as they are necessary in order to restore General Relativity in high-density environments. Thus, the dynamics in these models can only be rigorously studied through N-body simulations which include the non-linear mechanisms of $f(R)$ and DGP consistently. In the case of $f(R)$, the chameleon mechanism is triggered once the depth of the potential well is comparable to the background value of the scalar field. The suppression of the force modifications within a halo thus depends not only on the halo mass but also its environment. Consequently, we found significant scatter from halo to halo in the force modification $\\g$ measured in the $f(R)$ simulations. Furthermore, we identified a subset of halos which are in the close vicinity of massive neighbors, and which show a much stronger suppression of the force modifications than expected for isolated halos. In the majority of cases however, the simulation results confirm the basic expectation that halos are ``unscreened'' below a certain threshold mass determined by the potential well and the field value, whereas GR is restored at higher masses. Furthermore, a simple model based on the spherically symmetric solution of the field equations provides a good match to the scale- as well as mass-dependence of the force modifications in $f(R)$. In DGP, the non-linear suppression of the force modifications through the Vainshtein mechanism is much less dependent on halo mass and details of the large scale environment. Instead, the crucial quantity is the average mass density within a given radius. Thus, uncertainties in the semi-analytic predictions for DGP are mainly due to the density profile, and are already quite small. When taking into account the force resolution of the simulations, our predictions provide an excellent fit to the simulation measurements. Since the basic assumptions of the model, in particular spherical symmetry seem to hold well, we expect that force modifications can be predicted very accurately in DGP, provided the density profile is known sufficiently well. Given that our semi-analytic models appear to capture the mass- and scale-dependence of the modified forces correctly for both $f(R)$ and DGP, they can be useful in extending predictions beyond the limits of resolution and parameter space of the simulations. This will be necessary in particular for the comparison with observations. While this study is specific to $f(R)$ and DGP, it shows the qualitative features expected in observations of dynamics from viable modified gravity models, which employ a non-linear mechanism to restore GR locally. In the outer regions of massive clusters, as well as in lower mass objects, these models generally predict order unity deviations from GR. Observations in this regime thus offer the perspective of closing the last remaining loopholes for significant modifications to gravity on large scales. \\vspace*{1cm}" }, "1003/1003.4536_arXiv.txt": { "abstract": "We study the kinematics of GALEX-selected \\Ha\\ knots in the outer disk (beyond $R_{25}$) of NGC 628 (M74), a galaxy representative of large, undisturbed, extended UV (Type 1 XUV) disks. Our spectroscopic target sample of 235 of the bluest UV knots surrounding NGC 628 yielded 15 \\Ha\\ detections ($6\\%$), roughly the number expected given the different mean ages of the two populations. The measured vertical velocity dispersion of the \\Ha\\ knots between $1 - 1.8 R_{25}$ ($13.5 - 23.2$ kpc) is $< 11$ \\kms. We assume that the \\Ha\\ knots trace an `intermediate' vertical mass density distribution (between the isothermal sech$^2(z)$ and exponential distributions) with a constant scaleheight across the outer disk ($h_z$ = 700 pc) and estimate a total surface mass density of 7.5 $M_{\\sun}$ pc$^{-2}$. This surface mass density can be accounted for by the observed gas and stars in the outer disk (little or no dark matter in the disk is required). The vertical velocity dispersion of the outer disk \\Ha\\ knots nearly matches that measured from older planetary nebulae near the outskirts of the optical disk by Herrmann et al., suggesting a low level of scattering in the outer disk. A dynamically cold stellar component extending nearly twice as far as the traditional optical disk poses interesting constraints on the accretion history of the galaxy. ", "introduction": "The outskirts of a galaxy are expected to host signatures of disk formation and hierarchical accretion because of the long dynamical times at these radii. \\cite{TO92} argued that the thinness and coldness of inner galactic disks (inside the optical radius, $R_{25}$) posed significant problems for hierarchical models of galaxy formation. Qualitatively, these concerns become more pronounced if one can establish that cold galactic disks extend to even larger radii. Various studies have revisited the cold disk problem, generally finding that accretion events are less destructive than originally envisioned. Most recently, \\cite{Kazan09} studied the dynamical response of thin galactic disks to bombardment by cold dark matter substructure out to large radii in fully self-consistent, dissipationless $N$-body simulations. They found that disks survive these bombardments, but that they produce considerable thickening and heating at all radii, substantial flaring, and an increase in the stellar surface density in the disk outskirts (the latter due to outward radial migration of old stars during the growth and redistribution of disk angular momentum). Observations of the dynamical state of outer stellar disks \\citep{Christlein08} demonstrate that outer disks generally continue to obey the flat rotation curves of inner disks, with no increase in the in-plane velocity dispersion. Here we present the face-on kinematics of one nearby galaxy, NGC 628 (M74), and compare to both neutral hydrogen and existing stellar measurements at smaller radii. NGC 628 is a prototypical Grand Design spiral galaxy \\citep[type SA(s)c, dynamical mass within the studied region $\\sim3.3\\sci{11} M_{\\sun}$, assuming $v_{rot} = 200$ \\kms\\ out to $2.3 R_{25}$;][]{Thilker07, KB92}, and is by far the largest and most massive member of its small group (the brightest member after NGC 628, UGC1176, is $\\sim4.5$ mag fainter and over 125 kpc away). NGC 628 shows a standard exponential optical light profile to $R_{25}$, with only a very slight possible downbending in the profile to $\\sim1.3 R_{25}$, the extent of the deep optical observations \\citep{Natali92}. Using ultraviolet (UV) imaging, \\cite{Thilker07} classified NGC 628 as a Type 1 extended UV (XUV) disk, due to the structured, UV-bright emission complexes seen in the outer disk. Because the stellar disk appears largely undisturbed (both in optical and UV imaging, and, as we will show, from the kinematics) and because NGC 628 dominates the dynamics of its local environment, we treat it as representative of large, isolated spiral galaxies in the nearby universe. The outer disks of galaxies have received much recent attention, both observational \\citep{Thilker07, GildePaz07, ZC07, Christlein08, HF09, Trujillo09, Herrmann3} and theoretical \\citep{Bush08, Roskar08a, Roskar08b, Kazan09}. The surge in interest in outer disks has been fueled by recent ultraviolet (UV) observations of nearby disks with the GALEX satellite \\citep{Martin05, Thilker07}. Previously, \\cite{Ferg98} had used deep \\Ha\\ imaging to discover star formation in an extended component around three nearby galaxies, yet the ubiquity of this component in other disks remained largely unrecognized until the UV observations. This neglect stemmed in part from the fact that \\Ha\\ traces a limited subpopulation of the outer disk, namely those regions with O and B stars. Even relatively young regions may lack OB stars because they are older than 10 Myr or because they simply did not form such massive stars \\citep{Ferg98, HF09, PA08}. As such, the ubiquity of outer disk star formation is somewhat concealed. These barriers are removed with UV observations and we now know that many nearby galaxies host young outer disk stellar populations \\citep[$>30\\%$;][]{Thilker07, ZC07}. While galactic interactions can dramatically increase the level of outer disk star formation \\citep[see the well-known cases of M83, NGC 4625 and M94;][]{Thilker05, GildePaz05, Trujillo09}, even isolated galaxies show low-levels of ongoing star formation in their outer parts \\citep{Ferg98, Christlein08, HF09}. However, the \\Ha\\ knots provide the bright emission lines that make it possible to measure the kinematics we present here. While the gaseous outer disk of NGC 628 has been well-studied, for example by \\cite{KB92}, the stellar component near the outskirts of the disk remains poorly characterized. Observers have acquired deep broadband imaging of the outer disk of NGC 628 \\citep[e.g.][]{Natali92}, however the diffuse light becomes difficult to reliably trace fainter than $\\sim28$ mag arcsec$^{-2}$ in $V$, or beyond $\\sim1.3 R_{25}$. To sidestep the difficulty of obtaining detailed kinematics from low surface brightness emission, \\cite{Herrmann3} use planetary nebulae (PNe) across the face of NGC 628 to trace the underlying stellar distribution and to estimate the total mass density of the disk using the kinematic approach of \\cite{vdKruit88} (specifically, by measuring the vertical velocity dispersion of a tracer of the disk mass distribution to estimate the underlying surface mass density). In principle, the PNe provide an excellent approach to the problem, with large numbers of PNe available for analysis. Unfortunately for our purposes, the \\cite{Herrmann3} study of PNe in NGC 628 is mostly constrained to the inner disk (only two of their PNe are beyond $R_{25}$; their study of M83 and especially M94 yield more PNe in those outer disks, however those are interacting systems for which the kinematics are complicated by recent events). The current work represents the first to analyze the disk kinematics out to $1.8 R_{25}$ in NGC 628. Nevertheless, the results from \\cite{Herrmann3} at the largest radii ($\\sim R_{25}$) provide an interesting comparison to the results found here, especially because the PNe trace a population of objects that is on average at least $100\\times$ older than the \\Ha\\ knots considered here. We use multiobject spectroscopy to search for \\Ha\\ emission associated with blue GALEX sources in the outer disk of NGC 628. We measure the dispersion of the \\Ha\\ knot vertical velocity distribution, estimate the surface density of the outer disk, and determine if the baryonic material can account for the inferred mass. We then compare the results from the young \\Ha\\ knots to those from the PNe and consider the implications on disk evolution from the two measurements. Section 2 presents our sample selection, observations and data reductions. Section 3 presents our analysis, results and discussion, including our measurement of the vertical velocity dispersion of the \\Ha\\ knots and the surface mass density of the outer disk. Section 4 presents a summary and our conclusions. ", "conclusions": "We have measured the kinematics of \\Ha\\ knots in the outer disk of NGC 628. We find the stellar disk (traced by the \\Ha\\ knots) to have a low velocity dispersion, suggesting an undisturbed extended stellar disk. NGC 628 differs from the better-known outer disks of M83 and M94, which were likely accentuated by recent disturbances from neighbors, but resembles (both in optical and kinematic profiles) the outer disks detected in other relatively undisturbed nearby edge-on disks by \\cite{Christlein08}, supporting the idea that outer disk star formation can be a low-level and ongoing phenomenon in isolated galaxies. We find $\\sigma_z$ of the \\Ha\\ knots to be $< 11$ \\kms\\ between $1 - 1.8 R_{25}$ ($13.5 - 23.2$ kpc). We adopt a scaleheight similar to the known flaring gas profiles of outer disks ($h_z = 700$ pc) and estimate a mass density $\\Sigma = 7.5$ M$_{\\sun}$ pc$^{-2}$ that can be entirely explained by the observed gas and stars in the outer disk. If the \\Ha\\--hosting disk is actually much thinner than the flaring gas disk, more dark matter in the outer regions would be allowed. Assuming that outer disk star formation has been going for a Hubble time does not violate either the surface brightness nor surface mass constraints. The high incidence of outer disks \\citep[c.f.][]{Christlein08} suggests that the star formation is not a rare phenomenon -- here we show that current limits cannot exclude long-lived outer disk star formation. Finally, the velocity dispersion of PNe towards the outer disk of NGC 628 \\citep{Herrmann3} is nearly the same as that of the \\Ha\\ knots (the discrepancy grows slightly when considering the dispersion of the \\Ha\\ knots as a strict upper limit, although any plausible difference remains small). This can result either if the PN population is rather young (so that scattering has not had a chance to enlarge the dispersion significantly), if there is very little scattering, or if scattering occurs primarily in a single (or few) discrete events that occurred prior to the creation of the bulk of the PNe. We do not expect the outer disk to be exclusive young (see above) and we do expect some level of scattering in outer disks (if not from the classical spiral arms and molecular clouds of inner disks, then perhaps from satellites and dark halo substructure). The solution may be found in the \\cite{Kazan09} simulations, which show that outer disk heating is dominated by the most massive infall event of halo substructure onto the disk, so that stellar populations of different ages do not necessarily have different velocity dispersions. A larger sample of outer disk kinematic measurements could be used to constrain the rate and impact of such infall events." }, "1003/1003.3911_arXiv.txt": { "abstract": "GRB afterglows are among the best examples of astrophysical sources requiring a true multiwavelength observational approach. Radiation processes and main physical ingredients can only be disentangled with knowledge of their spectral and temporal properties through the largest possible band, i.e. from the X-ray down to radio. We now briefly review some of the most relevant observational findings recently obtained through optical observations. ", "introduction": "Optical (and NearInfraRed, NIR) observations are crucial for GRB afterglow studies. They contribute to multiwavelength study of afterglows making use of mature technologies guaranteeing high sensitivity. Moreover, already now and more and more in the near future, there will be specialized instruments able to compete with the capabilities of space-borne soft X-ray telescopes. On the other hand, optical observations are also among the best tools to observe the GRB host galaxies and the associated supernov\\ae, thus allowing us to study the possible progenitors and in general the physical conditions driving the occurrence of GRBs. ", "conclusions": "Optical follow-up of GRB afterglows is still a key ingredient for afterglow studies. In the recent years there has been an increasing attitude by the observers to share their data producing papers with essentially all the available information. There are therefore several events with optical coverage comparable in quality to the \\textit{Swift}-XRT soft X-ray coverage. In addition advanced facilities like GROND and RINGO, although equipping relatively small size telescopes, are providing exciting new data. In the near future it is likely that a new generation of intermediate size telescope will be designed, built and managed to allow a continuous and high quality coverage of GRB evolution, from a few seconds from the high-energy event down to the detection limits. The late-time evolution is also revealing to be a precious diagnostic and it will be followed by the largest facilities with better emphasis than in the past making a profitable use of the increasing number of 8m-class telescopes now operational." }, "1003/1003.1313_arXiv.txt": { "abstract": "{The generation of magnetic flux in the solar interior and its transport from the convection zone into the photosphere, the chromosphere, and the corona will be in the focus of solar physics research for the next decades. With 4\\,m class telescopes, one plans to measure essential processes of radiative magneto-hydrodynamics that are needed to understand the nature of solar magnetic fields. One key-ingredient to understand the behavior of solar magnetic field is the process of flux emergence into the solar photosphere, and how the magnetic flux reorganizes to form the magnetic phenomena of active regions like sunspots and pores. Here, we present a spectropolarimetric and imaging data set from a region of emerging magnetic flux, in which a proto-spot without penumbra forms a penumbra. During the formation of the penumbra the area and the magnetic flux of the spot increases. First results of our data analysis demonstrate that the additional magnetic flux, which contributes to the increasing area of the penumbra, is supplied by the region of emerging magnetic flux. We observe emerging bipoles that are aligned such that the spot polarity is closer to the spot. As an emerging bipole separates, the pole of the spot polarity migrates towards the spot, and finally merges with it. We speculate that this is a fundamental process, which makes the sunspot accumulate magnetic flux. As more and more flux is accumulated a penumbra forms and transforms the proto-spot into a fully-fledged sunspot.} ", "introduction": "The generation of the solar magnetic flux is thought to take place in the solar interior, either in the convection zone pro\\-per or in the tachocline beneath it. The generated magnetic flux becomes buoyant and rises upward through the convection zone \\citep{parker1979book}. Regions of flux emergence appear in the photosphere, and processes of structure formation take place. As a product pores and sunspots are observed in the photosphere. One particularly interesting process is the formation of a sun\\-spot penumbra. What is the crucial ingredient that initiates the transformation from a proto-spot into a sun\\-spot with a full-size penumbra? Where does the magnetic flux come from that is needed to form the penumbra? Our knowledge on how the penumbra forms is rather poor. \\citet{zwaan1992} describes the formation of a sunspot with penumbra as coalescence of the existing pores \\citep[refering to][]{mcintosh1981, bumba+suda1984}. The spatial resolution and polarimetric sensitivity of those observations did not allow to trace the smaller magnetic patches, e.g., magnetic knots or bright points. Zwaan compiles the following results: the penumbra forms section by section; each section completes within an hour; and the penumbra nearly closes leaving a gap toward the inside of the active region. The formation of sunspots is attended by magnetic flux emergence. In this respect, elongated granules are essential since they are signatures of flux emergence \\citep{bray+loughhead1964}. This link has recently also been demonstrated in MHD simulations \\citep[e.g.,][]{tortosa+moreno2009, cheung+etal2008}. Such aligned and elongated granules are compatible with the top of the magnetic loop passing through the photosphere \\citep{strous+etal1996}. We have recently presented a unique data set about the formation of a penumbra \\citep{schlichenmaier+al2010a}. There we have established that the total area of the spot increases, and that this increase is exclusively due to an area increase of the penumbra (see also Fig.~\\ref{fig:1} discussed below). This indicates that the magnetic flux of the spot is increasing, and the immediate question is: Where does the magnetic flux come from? In this paper we investigate how the site of emerging magnetic flux is linked to the growth of the sunspot. \\begin{figure*}% \\includegraphics[width=13cm]{schliche_fig1.png} \\hfill \\parbox[b]{3.7cm}{ \\caption{Snapshot images in the G-band document the evolution of a sunspot penumbra at 08:32, 09:33, and 13:06 UT. Tickmarks are in arcsec.}} \\label{fig:1} \\end{figure*} ", "conclusions": "We report on flux emergence in the close vicinity of a spot. The proto-spot evolves into a sunspot with a well developed penumbra. We argue that the increase of the magnetic flux of the proto-spot is due to the magnetic flux that emerges through elongated granules in the immediate vicinity of the spot. We present one example in which an elongated granule is associated with the signature of a rising magnetic flux loop. While one foot-point of the loop migrates towards the spot, the other moves in the opposite direction. From this we infer that the small-scale flux emergence contributes to the growth of the sunspot. Zwaan (1992) reviews the observations of the sunspot formation. He shows that a sunspot forms by coalescence of pores. Here we present an alternative way of flux accumulation to form a sunspot. The small-scale bipolar loops that appear in the emergence site can significantly contribute to the total magnetic flux of the sunspot. Within 4:40 h of observation the spot area grows from 230 arcsec$^2$ to 360 arcsec$^2$, but we do not see a pore that merges with the sunspot. Instead, we witness elongated granules that carry magnetic flux in form of small-scale loops. The distance between the footpoints of some of the small-scale bipoles increases gradually. Finally, the footpoint that has the same polarity as the spot merges with it. From this we envisage that the increase in the area and magnetic flux of the sunspot during our time sequence is due to merging of footpoints from small-scale bipoles: the footpoints with the proper polarity migrate towards the spot, the footpoints with the other polarity migrate towards the other polarity of the active region. In our case this process contributes one third of the total flux of the sunspot. We have no information on how the pre-existing proto-spot was formed. It may well be that it was formed by merging pores. Hence, both processes are probably relevant for the formation of a fully-fledged sunspot. We find many examples of elongated granules and rising bipoles, for which the magnetic foot-point closer to spot has the polarity of the spot. This seems to indicate that the small-scale emergence at the surface is rooted in a larger coherent structure beneath the surface, as seen in the simulations by \\citet{cheung+etal2008}. Mysteriously, all these small-scale bipoles are reassembled at the surface to form a sunspot. Further investigation will help to understand the associated processes of reassembling the magnetic flux into a large coherent structure. It is also interesting to note that while the magnetic flux merges with the proto-spot on one side, the penumbra forms on the opposite side of the spot. How does the opposite side `know' that flux is accumulated? How does the spot propagate the signal of the incoming magnetic flux to its other end? These questions will be addressed in a further investigation of our spectropolarimetric data set." }, "1003/1003.1255_arXiv.txt": { "abstract": "Sunspots are the most spectacular manifestation of solar magnetism, yet, 99\\% of the solar surface remains 'quiet' at any time of the solar cycle. The quiet sun is not void of magnetic fields, though; they are organized at smaller spatial scales and evolve relatively fast, which makes them difficult to detect. Thus, although extensive quiet Sun magnetism would be a natural driver to a uniform, steady heating of the outer solar atmosphere, it is not clear what the physical processes involved would be due to lack of observational evidence. We report the topology and dynamics of the magnetic field in very quiet regions of the Sun from spectropolarimetric observations of the Hinode satellite, showing a continuous injection of magnetic flux with a well organized topology of $\\Omega$-loop from below the solar surface into the upper layers. At first stages, when the loop travels across the photosphere, it has a flattened (staple-like) geometry and a mean velocity ascent of $\\sim3$ km/s. When the loop crosses the minimum temperature region, the magnetic fields at the footpoints become almost vertical and the loop topology ressembles a potential field. The mean ascent velocity at chromospheric height is $\\sim12$ km/s. The energy input rate of these small-scale loops in the lower boundary of the chromosphere is (at least) of $1.4\\times 10^6-2.2\\times 10^7$ erg cm$^{-2}$ s$^{-1}$. Our findings provide empirical evidence for solar magnetism as a multi-scale system, in which small-scale low-flux magnetism plays a crucial role, at least as important as active regions, coupling different layers of the solar atmosphere and being an important ingredient for chromospheric and coronal heating models. ", "introduction": " ", "conclusions": "" }, "1003/1003.3250_arXiv.txt": { "abstract": "{We present new ATCA 21-cm line observations of the neutral hydrogen in the nearby radio galaxy Centaurus~A. We image in detail (with a resolution down to 7\\arcsec , $\\sim 100$ pc) the distribution of \\HI along the dust lane. Our data have better velocity resolution and better sensitivity than previous observations. The \\HI extends for a total of $\\sim 15$ kpc. The data, combined with a titled-ring model of the disk, allow to conclude that the kinematics of the \\HI is that of a regularly rotating, highly warped structure down to the nuclear scale. The parameters (in particular the inclination) of our model are somewhat different from some of the previously proposed models but consistent with what was recently derived from stellar light in a central ring. The model nicely describes also the morphology of the dust lane as observed with Spitzer. There are no indications that {\\sl large-scale} anomalies in the kinematics exist that could be related to supplying material for the AGN. Large-scale radial motions do exist, but these are only present at larger radii ($r > 6$ kpc). This unsettled gas is mainly part of a tail/arm like structure. The relatively regular kinematics of the gas in this structure suggests that it is in the process of settling down into the main disk. The presence of this structure further supports the merger/interaction origin of the \\HI in Cen~A. From the structure and kinematics we estimate a timescale of $1.6 - 3.2 \\times 10^8$~yr since the merging event. No bar structure is needed to describe the kinematics of the \\Hi . The comparison of the timescale derived from the large-scale \\HI structure and those of the radio structure together with the relative regularity of the \\HI down to the sub-kpc regions does not suggest a one-to-one correspondence between the merger and the phase of radio activity. Interestingly, the radial motions of the outer regions are such that the projected velocities are {\\sl redshifted} compared to the regular orbits. This means that the blueshifted absorption discovered earlier and discussed in our previous paper cannot be caused by out-moving gas at large radius projected onto the centre. Therefore, the interpretation of the blueshifted absorption, together with at least a fraction of the redshifted nuclear absorption, as evidence for a regular inner disk, still holds. Finally, we also report the discovery of two unresolved clouds detected at 5.2 and 11~kpc away (in projection) from the \\HI disk. They are likely an other example of left-over of the merger that brought the \\HI gas.} ", "introduction": "The assembly of early-type galaxies and its connection to the presence of an active nucleus (and in particular a radio-loud nucleus) is still a matter of debate. Gas-rich galaxy mergers and interactions could play an important role in providing the fuel to trigger nuclear activity (including the radio phase). However, recent studies of radio galaxies have shown that the activity in some of these galaxies may be associated instead with the {\\sl slow} accretion of (hot) gas. This is thought to be the case in particular, for edge-darkened radio galaxies where advective low efficiency/rate flows could be the dominant mode of accretion \\citep[see e.g.][]{allen06,balmaverde08}. One way to address the issue of the relation between interaction/merger and AGN, is by studying the gas in radio galaxies. Nearby sources are prime targets as the morphology and kinematics of the gas --- both at large scales and in the nuclear region --- can be studied in detail and related to each other. Atomic hydrogen can be a useful tracer for these purposes. In particular, \\HI gas found at kiloparsec scales, or larger, can be used to trace the evolution of the host galaxies (e.g.\\ major merger vs.\\ small accretions). On the other hand, \\HI in absorption is one of the diagnostics to probe the central regions of active galactic nuclei and the interaction of the AGN with its immediate environment. \\HI absorption studies offer the unique possibility of tracing the nature of the accretion or the energetics of the outflow \\citep[e.g.][]{morganti05}. Unfortunately, the number of radio-loud galaxies where the atomic neutral hydrogen can be studied in such details both in emission {\\sl and} in absorption is limited. Due to its proximity\\footnote{At the assumed distance of Cen~A of 3.8~Mpc \\citep{harris10}, 1~arcmin corresponds to 1.1~kpc.}, the nearby radio galaxy Centaurus~A (Cen~A) is a unique object where to study the different phases of the gas, including \\Hi , in a radio-loud AGN. This can be done on spatial scales ranging from the inner kpc, where the influence of the AGN can still be relevant, up to many kpc where the signature of past merger(s) might be still visible through unsettled gas. Cen A harbours a well known radio source classified as a Fanaroff-Riley type-I source \\citep[FR-I,][]{fanaroff74} with a total power $P_{2.7\\rm{GHz}}=10^{24.26}$~W~Hz$^{-1}$. The total extent of the radio source reaches more than 8\\degr , making Cen A the largest extragalactic radio source in the sky \\citep{cooper65}. The radio source shows a variety of features on different scales. A radio jet extends from sub-parsec scales \\citep[e.g.][]{jones96,horiuchi06} to a projected distance of $\\sim 6$ kpc from the nucleus, ending in an inner radio lobe \\citep{burns83,clarke92}. Vast regions of fainter emission extend beyond the inner lobes: the middle lobe extends out to 40 kpc, while the outer lobe is even more diffuse extending out to more than 500 kpc \\citep[see Fig.~1 in][]{morganti99}. The different orientation and morphologies of the various structures could be the result of precession \\citep[e.g.][]{haynes83} of the central engine combined with strong interaction with the surrounding medium. Alternatively, the nuclear activity in Cen~A has been recurrent, although recent spectral index studies suggest that the large-scale structure is currently undergoing (or has very recently undergone) some particle injection event \\citep[][and refs. therein]{hardcastle09}. Most of the cold and warm gas in Cen~A is concentrated along the strongly warped dust lane \\citep[e.g.][]{nicholson92,quillen06A}. The neutral hydrogen present in this region (both in emission and absorption) has been studied by e.g. \\citet{vdh83,vangorkom90,sarma02}. The morphology and kinematics of the outer regions of the \\HI disk suggest that some of the gas has not yet settled into regular rotating orbits \\citep{vangorkom90}. \\HI has also been found well outside this disk. At a distance between 10 and 15 kpc from the nucleus, \\citet{schiminovich94} have found a partial ring structure with a smooth N-S velocity gradient and rotating perpendicular to the gas located along the dust lane (but co-rotating with the stellar body). In the north-eastern region of this ring, young stars have been detected that have possibly formed by the interaction between the radio jet and the neutral hydrogen in the ring \\citep{oosterloo05}. Cen~A was shaped through hierarchical merging over many Gyr as revealed by the spread in ages of globular clusters \\citep[see e.g.][]{woodley10}. The last major merger event happened about 5~Gyr ago \\citep{peng04}. The prominent warped dust and gas disk in the central region, and the outer partial ring structure, suggest that Cen~A experienced a recent merging event where its large elliptical body merged with a smaller gas-rich galaxy \\citep{baade54, tubbs80}. \\citet{quillen93} calculated a timescale of about 200 million years since the core of the infalling galaxy merged with the centre of the elliptical galaxy. This is consistent with the timescales derived by the presence of tidal debris and by the shell-like features containing atomic hydrogen \\citep{schiminovich94,peng02} and molecular gas \\citep{charmandaris00}. \\begin{table*} \\caption{Summary of observations.} \\label{reducprop} \\centering \\begin{tabular}{c l l} \\hline\\hline Total on-source integration time & & $3 \\times 12$~h\\\\ Observation dates & & 12/14/22 April 2005\\\\ Bandwidth & & 16~MHz\\\\ Number of channels in the cubes & & 512\\\\ Channel spacing ($dV$) & & 6.6~\\kms\\\\ Velocity resolution ($\\Delta v$, after Hanning) & & 13.2~\\kms\\\\ Shortest/longest baseline & & 77~m/6~km\\\\ Bandpass/flux calibrator & & PKS~1934--638\\\\ Gain/phase calibrator & & PKS~1315--46\\\\ \\hline Data cube & high resolution & low resolution\\\\ Weighting scheme & uniform & robust 0\\\\ Beam (HPBW) & $8.1\\arcsec\\times 6.8$\\arcsec & $19.6\\arcsec\\times 18.2$\\arcsec \\\\ Beam position angle & --12.3$^{\\circ}$ & --48.3$^{\\circ}$\\\\ rms noise (mJy~beam$^{-1}$) in the cubes & 1.3 & 1.2\\\\ rms noise ($10^{19}$~atoms~cm$^{-2}$) & 30.92 & 4.45\\\\ rms noise ($M_{\\odot}$~pc~$^{-2}$) & 2.48 & 0.36\\\\ Peak flux continuum (Jy~beam$^{-1}$) & 4.41 & 4.95\\\\ rms noise (mJy~beam$^{-1}$) continuum & 32.2 & 26.7\\\\ \\hline \\end{tabular} \\end{table*} \\begin{figure*} \\centering \\includegraphics[width=0.7\\textwidth, angle=270]{14355f01.ps} \\caption{\\HI emission in greyscale overlaid with the low resolution radio continuum (thin, grey contour levels: 0.5~Jy~beam$^{-1}$ increasing with a factor of 1.5) and absorption contour levels on top (thick contours, red in the online version). Absorption contour levels: $0.33\\times 10^{19}$~cm$^{-2}$ increasing with a factor 3. Emission contour level: $1.5\\times10^{20}$~cm$^{-2}$.} \\label{m0.optical} \\end{figure*} \\HI absorption is detected against the nucleus, the northern jet and the southern lobe \\citep[e.g.][]{roberts70, vdh83,vangorkom90, sarma02}. The absorption against the jet and southern lobe has been interpreted as gas that is part of the large-scale disk which is roughly perpendicular to the jet axis \\citep[$PA=50$\\degr ,][]{tingay98}. Since no absorption is found against the northern radio lobe, this lobe is likely pointing towards us whereas the southern radio lobe points away. Van Gorkom et al. (1990) and \\citet{burns83} concluded that the northern part of the gas disk is (at least partly) behind the radio source and the southern part of the \\HI disk is in front. For the component of \\HI absorption against the nucleus, earlier observations have shown the presence of redshifted absorption against the (unresolved) core \\citep[e.g.][]{vdh83,vangorkom90,sarma02}, but our recent results, discussed in \\citet[][Paper~I]{morganti08} and in this paper, have shown that blueshifted absorption against the core is also present. Despite the many studies devoted to Cen~A, there are a number of open questions that are essential for the understanding of the formation and evolution of Cen~A and its AGN. For example, the relation between the age of the radio structures and the timescale of the merger is not completely clear. In particular, it is unclear what triggers the current AGN activity (i.e. the activity producing the inner radio lobe) and whether it is related to the merging event. Thus, it is important to obtain the morphology and kinematics of the gas disk to the largest possible radii and investigate the structure of the disk. This can provide an estimate of the timescale of the merger and allow to investigate the presence and kinematics of unsettled gas (e.g. whether significant radial motions are present and, if so, on which scale). At the same time, it is important to investigate the distribution of \\HI down to the most inner parts of Cen~A to understand whether kinematical signatures of gas connected with the fuelling of the AGN are present. The large amount of \\HI present in Cen~A and its proximity allow to obtain deeper observations which can be used to investigate the questions above. To achieve this, we have performed new spectral-line observations of Cen~A that have higher spatial and velocity resolution and better sensitivity than those of previous studies. These new observations have already allowed us to detect blueshifted absorption against the nucleus while also the redshifted absorption extends to higher velocities compared to previously known. These results have been presented in Paper~I. This gas has been interpreted as part of a circumnuclear disk, likely the counterpart of what was already detected in CO and other lines, leaving still open the question of what triggers the current AGN activity. To further investigate this issue, we present in this paper the study of the kinematical state of the \\HI disk and the results from 3D modelling. The paper is organised in the following way. In Sect.~2 we present the new observations while in Sect.~3 we describe the results. In Sect.~4 we derive and discuss tilted-ring models of the \\HI disk that are based on the 3D data cube that describe the strongly warped \\HI disk as well as previously published data. We show that the \\HI is dominated by rotation, with radial motions giving a significant contribution in the outer parts. Furthermore, we argue that the blue- and redshifted absorption against the nucleus must be mainly caused by gas close to the nucleus. In Sect.~5 we discuss the consequences of the observational and modelling results for the formation and evolution process of Cen~A, the merger timescales involved and the implications for the AGN activity. A short summary is given in Sect.~6. \\begin{figure*} \\centering \\includegraphics[width=0.90\\textwidth]{14355f02.eps} \\caption{Top panel: Total intensity contours over-plotted on DSS image. Contour levels: 0.2, 0.4, 0.8, 1.6 (absorption only), 3.2, 6.4, 12.8, 25.6 and $51.2\\times 10^{20}$~atoms~cm$^{-2}$ for the emission (black) and absorption (white). $T_{\\rm{spin}}=100$~K. The horizontal line indicates the position of the pv-slices in Fig.~\\ref{pv.plot2} and \\ref{pvplots.models}. Bottom left panel: spectrum of \\HI absorption cloud. Bottom middle panel: spectrum of the nuclear absorption. Bottom right panel: zoom-in of bottom middle panel.} \\label{mom0.spec} \\end{figure*} ", "conclusions": "The new data presented here have allowed to image in detail the \\HI emission and absorption of the warped disk of Cen~A and to identify new features in the gas. In this section we discuss the implications of the \\HI results for the formation and evolution (e.g. merger history) of Cen~A and the different phases of AGN activity. \\begin{figure} \\centering \\includegraphics[height=8cm, angle=270]{14355f10.ps} \\includegraphics[height=8cm, angle=270]{14355f11.ps} \\caption{Position angle (top panel), inclination (bottom panel) of our \\HI model (solid line) superimposed to the values from the \\citet{quillen06A} model (long dashed line, similar to those of \\citet{nicholson92}. In addition, also the values from the stellar modelling of \\citet{kainulainen09} are plotted (dotted line).} \\label{incl} \\end{figure} \\subsection{Modelling implications: The large-scale \\HI distribution} Our data confirm that the structure and kinematics of the \\HI in Cen~A is dominated by rotation, down to the central regions, and most of the \\HI (but clearly not all) is confined to a fairly regular, warped disk (for $r<6$~kpc). In addition to this, we detect outer unsettled gas that appears to be located in an arm-like structure. Therefore, some of the gas in emission and absorption is observed below the dust-lane and it is not reproduced by our modelling. This arm shows a smooth velocity gradient - similar for absorption and emission but with an offset in velocity between the two. Because of the regular velocity characteristics of this structure, it is likely that it will become part of the outer part of the rotating disk in the near future. This type of structure is expected as a result of a merger or accretion as shown in many numerical simulations \\citep[e.g.][]{hibbard96,barnes02}. Thus, the \\HI further support the hypothesis of Cen~A being the result of a relatively recent merger of an elliptical galaxy with a smaller gas-rich (SMC-like?) disk galaxy \\citep[e.g.][see also Sect.~1]{baade54,quillen93}. Using the results of our modelling, we can derive a rough estimate of how long ago the merger must have taken place to allow a large fraction of the \\HI to settle down in the warped disk we observe. We can assume that the gas located at larger radii than we can model (i.e. $r>350$\\arcsec ) has performed at least one or two revolutions since the merger took place giving a rough estimate resulting in $1.6-3.2\\times 10^8$~yr. This time estimate is in agreement, inside the uncertainties, with \\citet[][$2\\times 10^8$~yr]{quillen93} and \\citet[][$3\\times10^8$~yr]{peng02}, while it is shorter than what derived by \\citet{sparke96} - $7.5\\times 10^8$~yr - that would correspond to almost five revolutions. Such a long timescale would be sufficient to relax most of the gas in the disk, in contradiction with our \\HI results. The merger hypothesis does find further support from the partial ring structure discovered by \\citet[][see also Sect.~1]{schiminovich94}. In that respect the existence of additional, low mass ($<10^6$~\\Msun ) \\HI clouds is not surprising. The velocities of the two newly discovered \\HI clouds (see Sect.~3.3) are in rough agreement with the spatial velocity distribution of the ``outer ring'' structure. Clouds and filaments of ionised gas have been found in this region along the jet direction \\citep[e.g.][]{dufour78,morganti91}. For the origin of these structures, two hypothesis have been put forward. In addition to the merger origin, an alternative scenario was the cooling condensations in a tenuous, thermally unstable hot component of the interstellar medium gas. This scenario was somewhat favoured to explain the filaments of ionised gas, as it would explain the highly turbulent velocity structure observed there \\citep[see also][]{oosterloo05}. However, such high velocity gradients are not seen in the \\HI profile of the newly detected clouds which could mean that the clouds were not hit by the jet. Whether the newly discovered \\HI clouds are a cooling condensation or are more likely the left over of the interaction/merger discussed above, remains unclear. As a final remark, we would like to note that also the distribution of the \\HI confirms what was found by other studies of the molecular gas and dust \\citep{nicholson92,quillen06A}, i.e. the lack of gas in region between 10\\arcsec ~and 45\\arcsec ~(0.2~kpc to 0.8~kpc) while on the smallest scales circumnuclear disks are found (Paper~I and ref. therein). The peak in the surface brightness around 150\\arcsec ~from the centre suggests a ring-like structure in the distribution of \\HI in Cen~A (see Fig.~\\ref{model.parameters}). \\HI ring-like structures are often detected around early-type galaxies and there are a number of possible scenarios for the formation of such a structure (see e.g. Donovan et al. 2009 for a summary and references). A ring morphology can be the result of the presence of a bar structure but we found no evidence for such a structure in Cen~A and indeed many other galaxies with gas rings do not show bars. In some cases, \"burning out\" of the gas in the central regions of the galaxies has been proposed, but the famous example of IC~2006 \\citep{schweizer89} shows that also gas accretion can form these structures. \\subsection{Merger and different phases of radio activity} The kinematics of the \\HI in Cen~A is dominated by rotation down to the nuclear regions. The discovery that the nuclear absorption is also partly blueshifted, (see Paper~I) has questioned the interpretation that \\HI against the nucleus represent evidence for the {\\sl direct} fuelling of the AGN. Furthermore, our analysis presented here shows that the nuclear blueshifted absorption is not due to gas at large radius seen in projection against the core. The nuclear absorption is likely originating from a nuclear disk that represents the \\HI counterpart of the structure seen from the molecular gas (see Paper~I for a detailed discussion). This nuclear disk is part of the overall warped disk structure. Thus, our results confirm what suggested in Paper~I that the nuclear \\HI gas does not constitute direct evidence of gas infall into the AGN. However, it is worth mentioning that radial motions (in addition to rotation) on the scale of tens of pc have been detected in the high resolution, near-IR data obtained by \\citet{neumayer07}. Surprisingly they found that for higher excitation lines [SiVI] and [FeII] the velocity pattern is increasingly dominated by a non-rotating component, elongated along the radio jet. Interestingly, these non-rotational motions were detected {\\sl along} the jet with redshifted velocities (compared to the systemic) seen on the main-jet side and blueshifted on the counter-jet side. These motions (stronger in [SiVI]) can be explained as backflow of gas accelerated by the plasma jet. This gas can perhaps be involved in the fuelling of the AGN. Whether a similar situation is also occurring in the \\HI gas can only be investigated with deep VLBI observations. In order to investigate in more detail whether a link can be found between the merger/accretion event and the nuclear activity, it is important to compare the timescales of these events. Cen~A shows a complex radio structure with lobes on different scales, from a few kpc of the inner lobe (see Fig.~\\ref{m0.optical}) to tens of kpc for the middle lobe and up to several hundred kpc for the outer lobe (Sect.~1). These three radio structures could be either the result of precession \\citep[][see also Sect.~1]{haynes83} or the nuclear activity in Cen~A has been recurrent. However, the large-scale structure is currently experiencing a particle injection event, as recent spectral-index studies suggest \\citep[][and refs. therein]{hardcastle09}. The ages of these lobes are not easy to estimate. According to \\citet[]{saxton01} the age of the northern {\\sl middle} lobe can be estimated to be about $1.4\\times 10^8$~yr assuming this lobe has been created by an old episode of jet activity that has been {\\sl interrupted} by the disruption of the jet by the merger. Since then, this structure has been rising buoyantly away from the nucleus. If this is the case, the merger that has brought the \\HI would actually be the cause of the disruption of the old radio activity. The timescale of the merger derived from the \\HI could support this idea. The timescale of the inner radio lobe, $\\sim 5-10\\times 10^6$~yr \\citep[][and refs. therein]{croston09}, would indicate that it takes quite some time for the system to recover from this disrupting event and for the gas to settle again in a steady accretion flow. A (possible) time-delay between the merger and the onset of the AGN has also been detected in a few other nearby radio galaxies \\citep[e.g.][]{tadhunter96,emonts06,davies07,wild10}. It is worth noting that if the scenario proposed by \\citet[]{saxton01} is correct, the delay between the disruptive merger and the onset of a new phase of activity could be related to the time required by the X-ray halo to build up again \\citep[see][]{sansom00}. On the other hand, recent spectral-index studies have suggested different timescales for the large-scale radio structure. High frequency studies show a spectral break for the southern lobe that would suggest a life time of $3 \\times 10^7$ yrs \\citep[see][]{hardcastle09}. Furthermore, no spectral break is observed in the northern lobe (both the giant and the middle lobe), suggesting that they have very recently undergone some particle injection event \\citep[][and refs. therein]{alvarez00,hardcastle09}. Thus, if the timescales derived from the spectral index are correct, it would indicate that that the {\\sl overall} radio emission starts relatively late compared to the merger. In both cases, the comparison of the timescale derived from the \\HI and the radio structure, together with the relatively regularity of the \\HI down to the sub-kpc regions, do not suggest a one-to-one correspondence between the merger and the phase of radio activity. In particular, if the timescales derived from spectral index are correct, {\\sl more than one radio burst can be produced while the merger is settling in}. This would indicate that, while the merger could be responsible for bringing gas in the vicinity of the nucleus, the accretion mechanism that fuels the AGN is likely not related to the merger. This is in agreement with the idea that in radio galaxies of relatively low radio power (i.e. FR-I type) the fuelling of the black hole proceeds directly from the hot phase of the interstellar medium in a manner analogous to the Bondi accretion. As suggested by a number of authors \\citep[see e.g.][]{allen06,hardcastle07} the hot, X-ray emitting phase of the intergalactic medium (IGM) is sufficient to power the jets of several nearby, low-power radio galaxies. In the case of Cen~A, a mass accretion rate $\\dot M_{Bondi} = 6.4 \\times 10^{-4}$\\Msun ~yr$^{-1}$ and a Bondi efficiency of $\\sim 0.2$ \\% has been derived from Chandra observations \\citep{evans04}. However, it would decreases if the new value of the BH mass derived by \\citet{neumayer07} and \\citet{cappellari09} is used. These values are consistent with what is found for other FR-I radio galaxies \\citep{balmaverde08} while the efficiency appears to be lower than the \"canonical\" 0.1 for the high efficiency optically thick accretion disks. On the other hand, and to complicate the picture, it is known \\citep[from the Fe K$\\alpha$ line and the large column density measured in the X-ray towards the nucleus][]{evans04} that Cen A has large quantities of cold gas in a molecular torus at about 1pc from the black hole, more typical of FR-II galaxies." }, "1003/1003.2640_arXiv.txt": { "abstract": "% The diffusion of astrophysical magnetic fields in conducting fluids in the presence of turbulence depends on whether magnetic fields can change their topology or reconnect in highly conducting media. Recent progress in understanding fast magnetic reconnection in the presence of turbulence is reassuring that the magnetic field behavior in computer simulations and turbulent astrophysical environments is similar, as far as the magnetic reconnection is concerned. This makes it meaningful to perform MHD simulations of turbulent flows in order to understand the diffusion of magnetic field in astrophysical environments. These simulations support the concept of reconnection diffusion, which describes the ability of magnetic fields to get removed from magnetized clouds and cores in the process of star formation. ", "introduction": "% Magnetic reconnection describes the ability of magnetic field to change its topology. The famous Sweet-Parker model of reconnection (Sweet 1958, Parker 1957) (see Figure~\\ref{recon1}, upper panel) produces reconnection rates which are smaller than the Alfv\\'en velocity by a square root of the Lundquist number, i.e. by $S^{-1/2}\\equiv (LV_A/\\eta)^{-1/2}$, where $L$ in this case is the length of the current sheet. Thus this scheme produces reconnection at a rate which is negligible for most of astrophysical circumstances. If the Sweet-Parker were proven to be the only possible model of reconnection, it would have been possible to show that MHD numerical simulations do not have anything to do with real astrophysical fluids. Fortunately, faster schemes of reconnection are available. The first model of fast reconnection proposed by Petschek (1964) assumed that magnetic fluxes get into contact not along the astrophysically large scales of $L$, but instead over a scale comparable to the resistive thickness $\\delta$, forming a distinct X-point, where magnetic field lines of the interacting fluxes converge at a sharp point to the reconnection spot. The stability of such a reconnection geometry in astrophysical situations is an open issue. At least for uniform resistivity, this configuration was proven to be unstable and to revert to a Sweet-Parker configuration (see Biskamp 1986, Uzdensky \\& Kulsrud 2000). Recent years have been marked by the progress in understanding some of the key processes of reconnection in astrophysical plasmas. In particular, a substantial progress has been obtained by considering reconnection in the presence of the Hall-effect (see Shay et al. 1998). The condition for which the Hall-MHD term becomes important for the reconnection is that the ion skin depth $\\delta_{ion}$ becomes comparable with the Sweet-Parker diffusion scale $\\delta_{SP}$. The ion skin depth is a microscopic characteristic and it can be viewed as the gyroradius of an ion moving at the Alfv\\'en speed, i.e. $\\delta_{ion}=V_A/\\omega_{ci}$, where $\\omega_{ci}$ is the cyclotron frequency of an ion. For the parameters of the interstellar medium (see Table~1 in Draine \\& Lazarian 1998), the reconnection is collisional (see further discussion in Yamada 2006). A radically different model of reconnection was proposed in Lazarian \\& Vishniac (1999, henceforth LV99). The middle and bottom panels of Figure~\\ref{recon1} illustrate the key components of LV99 model\\footnote{The cartoon in Figure~\\ref{recon1} is an idealization of the reconnection process as the actual reconnection region also includes reconnected open loops of magnetic field moving oppositely to each other. Nevertheless, the cartoon properly reflects the role of the 3-dimensionality of the reconnection process, the importance of small-scale reconnection events, and the increase of the outflow region compared to the Sweet-Parker scheme.}. The reconnection events happen on small scales $\\lambda_{\\|}$ where magnetic field lines get into contact. As the number of independent reconnection events that take place simultaneously is $L/\\lambda_{\\|}\\gg 1$ the resulting reconnection speed is not limited by the speed of individual events on the scale $\\lambda_{\\|}$. Instead, the constraint on the reconnection speed comes from the thickness of the outflow reconnection region $\\Delta$, which is determined by the magnetic field wandering in a turbulent fluid. The model is intrinsically three dimensional as both field wandering and simultaneous entry of many independent field patches, as shown in Figure~\\ref{recon1}, are 3D effects. The magnetic reconnection speed becomes comparable with $V_A$ when the scale of magnetic field wandering $\\Delta$ becomes comparable with $L$. \\begin{figure}[!t] \\begin{center} \\includegraphics[width=.6 \\columnwidth]{fig1.eps} \\caption{{\\it Upper plot}: Sweet-Parker model of reconnection. The outflow is limited by a thin slot $\\Delta$, which is determined by Ohmic diffusivity. The other scale is an astrophysical scale $L\\gg \\Delta$. {\\it Middle plot}: Reconnection of weakly stochastic magnetic field according to LV99. The model that accounts for the stochasticity of magnetic field lines. The outflow is limited by the diffusion of magnetic field lines, which depends on field line stochasticity. {\\it Low plot}: An individual small scale reconnection region. The reconnection over small patches of magnetic field determines the local reconnection rate. The global reconnection rate is substantially larger as many independent patches come together. From Lazarian et al. 2004.} \\label{recon1} \\end{center} \\end{figure} The LV99 model was successfully tested in Kowal et al. (2009) and we shall use this model to justify the concept of ``reconnection diffusion'' below. The advantage of LV99 model is that it is applicable to both collisional and collisionless plasmas. ", "conclusions": "\"Reconnection diffusion\" is a process which is expected to happen in the presence of turbulent reconnection predicted by LV99 model. As magnetic turbulence is ubiquitous in astrophysical fluids, we expect \"reconnection diffusion\" to be ubiquitous as well. The process may remove magnetic flux from star forming clouds and clumps on times much faster than it is allowed by the traditional ambipolar diffusion. We note that \"reconnection diffusion\" should be distinguished from the concept of \"turbulent magnetic diffusivity\" which is frequently discussed in the framework of kinematic, i.e. without backreaction, dynamo (see Parker 1979). The former concept is based on the tested idea of fast reconnection of strong magnetic fields in the presence of weak turbulence, while the latter concept assumes the efficient diffusion turbulent of magnetic field without taking into account its backreaction, as if the magnetic field were a passive scalar. A usual ``justification'' of the ``turbulent magnetic diffusivity'' is that turbulence mixes magnetic field opposite polarity on the very small scales, which is the process prohibited on the energetic grounds for any dynamically significant field. Therefore we claim that the \"turbulent magnetic diffusivity\" is an erroneous idea, while the \"reconnection diffusion\" is a well founded concept. \"Reconnection diffusion\" has direct relation to the problem of dissipation in accretion discs discussed in Shu et al. (2006, henceforth SX06). They have found that the dissipation there should be about four orders of magnitude larger than the Ohmic dissipation in order to solve the magnetic flux problem in these systems. They then appealed to a hyper-resistivity concept in order to explain the higher dissipation of magnetic field in a turbulent environment. We feel, however, that the hyper-resistivity idea is poorly justified (see criticism of it in Lazarian et al. 2004). At the same time, fast 3D reconnection can provide the magnetic diffusivity that is required for removal of the magnetic flux. This is what, in fact, was demonstrated in the present set of numerical simulations. It is worth mentioning that, unlike the actual Ohmic diffusivity, magnetic diffusivity mediated by fast reconnection does not transfer the magnetic energy {\\it directly} into heat. The lion share of the energy is being released in the form of kinetic energy, driving turbulence. The annihilation of the magnetic field happens in LV99 model, as in any model of fast reconnection, over a small fraction of the volume. This fraction goes to zero as the resistivity goes to zero. Magnetic turbulence induced by reconnection eventually dissipates energy, resulting in the medium heating. If the system is initially laminar, this potentially can result in flares of reconnection and the corresponding diffusivity. Similar to SX06, we expect to observe the heating of the media. Indeed, although we do not expect to have Ohmic heating, the kinetic energy released due to magnetic reconnection is dissipated locally and therefore we expect to observe heating in the medium. Our setup for gravity can be seen as a toy model representing the situation in SX06. In the broad sense, our work confirms that a process of magnetic field diffusion that does not rely on ambipolar diffusion is efficient. We showed that the higher the strength of the gravitational force, the lower is the flux-to-mass ratio in the central region (compared with the mean value in the computational domain). This could be understood in terms of the potential energy of the system. When the potential is higher, more it is energetically favorable is to pile up of matter near the center of gravity, decreasing the total potential energy of the system. When the turbulence is increased, there is an initial trend to remove more magnetic flux from the center (and consequently more inflow of matter into the center), but for the highest value of the turbulent velocity in our experiments, there is a trend to remove material (together with magnetic flux) from the center, reducing the role of the gravity, due the fact that the gravitational energy became small compared to the kinetic energy of the system. Our results also showed that when the gas is less magnetized (higher $\\beta$, or higher values of the Alfv\\'enic Mach number $M_{A}$), the turbulent diffusion of magnetic flux is more efficient, but the central flux-to-mass ratio relative to external regions is smaller for more magnetized models (low $\\beta$), compared to less magnetized models. That is, the contrast $B / \\rho$ between the inner and outer radius is higher for lower $\\beta$ (or $M_{A}$). If the turbulent diffusivity of magnetic field may explain the results in SX06, one may wonder whether one can remove magnetic field this way not only from the class of systems studied by SX06, but also from less dense systems. For instance, it is frequently assumed that only ambipolar diffusion is important for the evolution of subcritical magnetized clouds Tassis \\& Mouschovias (2005). Our study indicates that this conclusion may require modification in the presence of turbulence. While the concept of \"reconnection diffusion\" describes the diffusion of magnetic field in turbulent media, it is also closely connected to other important astrophysical concepts, i.e. turbulent advection of heat in the presence of turbulence (see Cho et al. 2003). If reconnection were slow, the mixing motions required by the turbulent advection would be difficult to explain. Thus, the fast diffusion of magnetic field induced by turbulence and the turbulent advection of heat in magnetized plasmas are interconnected." }, "1003/1003.0229_arXiv.txt": { "abstract": "We have investigated the empirical lag-luminosity relation in the Gamma-ray Burst (GRB) source-frame. We selected two energy bands ($100-200$ keV and $300-400$ keV) in the GRB source-frame, which after redshift correction, lie in the observer-frame energy range of the $Swift$ Burst Alert Telescope (BAT). The spectral lags between these energy channels are then presented as a function of the isotropic peak luminosity of the GRBs in the sample. ", "introduction": "Spectral lag is a common feature in Gamma-ray Bursts (GRBs). The lag is defined as the difference in time of arrival of high and low energy photons and is considered positive when the high-energy photons arrive earlier than the low energy ones. Norris et al. reported a correlation between spectral lag and the isotropic peak luminosity of GRBs based on a limited sample~\\cite{norris2000}. Subsequently, various authors have studied the lag-luminosity relation using arbitrary observer-frame energy bands of various instruments~\\cite{Ukwatta2009lag,Hakkila2008}. \\begin{figure}[htp] \\centering \\includegraphics[width=12cm, angle=0]{ukwattatn_fig1.eps}% \\caption{Fixed energy bands at the GRB source-frame are projected to various energy bands at the observer-frame, depending on the redshift.} \\label{fig01} % \\end{figure} Typically the spectral lag ($\\tau$) is extracted in two arbitrary energy bands in the observer-frame. However, because of the redshift (z) dependance of GRBs, the two energy bands can correspond to multiple energy bands in the source-frame thus introducing a variable energy dependant factor, which is difficult to take into account. We avoid this difficulty by defining two energy bands ($100-200$ keV and $300-400$ keV) in the GRB source-frame and projecting these two bands into the observer-frame using the relation $E_{\\rm observer}=E_{\\rm source}/(1+z)$. For our sample of GRBs, after projecting to the observer-frame, the selected energy bands lie in the $Swift$ Burst Alert Telescope (BAT)~\\cite{Gehrels2004} energy range ($15-350$ keV; see Fig.~\\ref{fig01}). Then we extracted the spectral lags between these energy channels, correct them for cosmological time dilations using the relation, $\\tau_{\\rm source}=\\tau_{\\rm observer}/(1+z)$ and plot them as a function of the isotropic peak luminosity. ", "conclusions": "We have investigated the spectral lag between $400-300$ keV and $200-100$ keV energy bands at the GRB source-frame by projecting these bands to the observer-frame. This is a step forward in the investigation of lag-luminosity relation since all previous investigations used arbitrary observer-frame energy bands. The correlation coefficient of $\\sim \\, -0.76$ shows a significant improvement over the average correlation coefficient of $\\sim \\, -0.68$ reported in the reference \\cite{Ukwatta2009lag}." }, "1003/1003.3470_arXiv.txt": { "abstract": "Spectroscopic and photometric observations show that many globular clusters host multiple stellar populations, challenging the common paradigm that globular clusters are ``simple stellar populations'' composed of stars of uniform age and chemical composition. The chemical abundances of second-generation (SG) stars constrain the sources of gas out of which these stars must have formed, indicating that the gas must contain matter processed through the high-temperature CNO cycle. First-generation massive Asymptotic Giant Branch (AGB) stars have been proposed as the source of this gas. In a previous study, by means of hydrodynamical and N-body simulations, we have shown that the AGB ejecta collect in a cooling flow in the cluster core, where the gas reaches high densities, ultimately forming a centrally concentrated subsystem of SG stars. In this Letter we show that the high gas density can also lead to significant accretion onto a pre-existing seed black hole. We show that gas accretion can increase the black hole mass by up to a factor of 100. The details of the gas dynamics are important in determining the actual black hole growth. Assuming a near-universal seed black hole mass and small cluster-to-cluster variations in the duration of the SG formation phase, the outcome of our scenario is one in which the present intermediate-mass black hole (IMBH) mass may have only a weak dependence on the current cluster properties. The scenario presented provides a natural mechanism for the formation of an IMBH at the cluster center during the SG star-formation phase. ", "introduction": "\\label{sec:intro} If globular clusters follow the same ``$M_{BH}-\\sigma$'' relation between central black hole mass and effective velocity dispersion as is observed in galaxies (Tremaine et al. 2002), many clusters should host intermediate-mass black holes (IMBHs) with masses in the range $10^2 \\ltorder M_{BH}/M_{\\odot} \\ltorder 10^4$. Unfortunately, hard observational evidence for IMBHs in globular clusters has proved elusive. Currently the strongest case is the massive cluster G1 in M31, where dynamical, X-ray, and radio studies are consistent with a mass $M_{BH}\\sim2\\times10^4~$ (Gebhardt et al. 2005, Ulvestad et al. 2007). In our own Galaxy, a measurement of an IMBH mass $M_{BH} \\sim 4\\times10^4 M_{\\odot}$ has been reported for $\\omega$ Centauri (Noyola et al. 2008), although this result has been questioned in a recent study (Anderson \\& van der Marel 2010, van der Marel \\& Anderson 2010) suggesting that the IMBH signature becomes much weaker when a more accurate estimate of the cluster center is used. Recently Ibata et al. (2009) have reported evidence of a density and kinematic cusp in the core of M54, a cluster located at the center of the Sagittarius dwarf galaxy; their models suggest that the cusp could be due to a $\\sim 10^4 M_{\\odot}$ IMBH. For most Galactic globular clusters, kinematic and structural data currently provide only upper limits on possible IMBH masses (McLaughlin et al. 2006, Pasquato et al. 2009, van der Marel \\& Anderson 2009). However, these upper limits are larger than the values of $M_{BH}$ expected from the $M_{BH}-\\sigma$ relation, and do not exclude the possibility that many globular clusters host IMBHs with masses in the above range (van der Marel \\& Anderson 2009). The presence of an IMBH in a globular cluster has important consequences for the cluster's structure, kinematics, internal energetics and long-term dynamical evolution (see e.g. Baumgardt et al. 2004, Heggie et al. 2007, Trenti et al. 2007, Miocchi 2007; see also McMillan 2008 for a review). In addition, Eddington accretion onto IMBHs has been proposed as the mechanism powering some ultraluminous X-ray sources (e.g. Farrell et al. 2009), and IMBHs have long been recognized as potential rich sources of gravitational waves detectable by LISA (Miller 2009). For these reasons, despite the lack of strong observational constraints, the formation and consequent properties of IMBHs have become a topic of considerable theoretical and observational interest. One possible IMBH formation mechanism is direct collapse of a very massive ($\\geq 250 M_{\\odot}$) Population III star (see e.g. Bond et al. 1984, Madau \\& Rees 2001 and references therein; but see Abel et al. 2002 for hydrodynamical simulations suggesting that these stars would form in isolation). For non-primordial or more massive IMBHs, the leading formation mechanisms center on dynamical processes in dense star clusters. These include runaway mergers of massive stars in clusters with high central densities (Portegies Zwart \\& McMillan 2002, G\\\"urkan et al. 2004, Portegies Zwart et al. 2004), and merging of black holes in binaries (see O'Leary et al. 2006 and references therein). However, numerous potential problems affecting these models have been pointed out in the recent literature, suggesting that the net growth rate and hence the final mass of the resulting IMBH may be much lower than suggested by earlier dynamical simulations (Yungelson et al. 2008, Glebbeek et al. 2009). Thus, while there appear to be numerous plausible stellar, binary, and dynamical pathways to the formation of relatively ``low-mass'' ($\\sim100-200 M_{\\odot}$) IMBHs in clusters, there is currently no clear consensus on a mechanism for the production of high-mass ($10^3-10^4 M_{\\odot}$) IMBHs. In this Letter we propose a new scenario for the growth of IMBHs in early globular clusters. It is a natural consequence of the physical conditions that may have obtained in cluster cores during the formation of the second generation (hereafter SG) stars now observed in many globular clusters. Specifically, we show that accretion by a pre-existing seed black hole of just a small fraction of the gas from which the SG stars formed can lead to black holes in the high-mass ($10^3-10^4 M_{\\odot}$) IMBH range. The structure of this Letter is as follows. In Section \\ref{sec:multip} we briefly review the observational evidence for multiple stellar populations in globular clusters, and summarize the main elements of our model for the formation and evolution of multiple populations. In Section \\ref{sec:bh}, we present our scenario for IMBH growth. We discuss some consequences of this scenario in Section \\ref{sec:concl}. ", "conclusions": "\\label{sec:concl} Several aspects of the scenario just presented merit further discussion and exploration. We consider first the effect on the cooling flow of the energy radiated by the accreting black hole. To address this issue we follow the analysis of black hole outflow presented by King (2003, 2005), in which it is shown that the wind produced by the black hole sweeps up the surrounding gas into an expanding shell. Because of effective energy losses due to inverse Compton scattering, the shell expands in the momentum-conserving regime. Under these conditions (see King 2003, 2005), in order for the momentum flux from a black hole accreting at the Eddington rate to overcome gravity and expel the shell from the cluster, the black hole mass must be larger than \\begin{equation} M_{BH,crit} = \\frac{f_g\\kappa\\sigma^4}{\\pi G^2}. \\end{equation} Here, $\\kappa$ is the electron scattering opacity, $f_g$ is the fraction of the total cluster mass in the form of gas, and it has been assumed that the gas is embedded in an isothermal system with velocity dispersion $\\sigma$. Assuming $\\sigma \\gtorder 20~\\hbox{km/s}$ (the smallest value for which the AGB ejecta can be retained) and $f_g\\approx0.05$ (D'Ercole et al. 2008), we find $M_{BH,crit} \\gtorder 10^4 M_{\\odot}$, much larger than the assumed mass of our initial seed black hole. Thus feedback from the accreting black hole is {\\em not} sufficient to alter significantly the global dynamics of the cooling flow or the SG star formation process. On the other hand, the effects of feedback may significantly alter the local gas dynamics in the vicinity of the black hole. For example, recent detailed 2-D hydrodynamical models of accretion onto a 100 $M_{\\odot}$ seed black hole in a dense protogalactic cloud (Milosavljevic et al. 2009) show intermittent accretion, reducing the net accretion rate to about $1/3$ of Eddington. Such a reduction in the mean accretion rate in our case would reduce the growth of the seed black hole to a factor $\\sim2-5$ for the range of $\\eta$ values adopted in \\S\\ref{sec:bh}. We note here that during the accretion phase the black hole will be very luminous and might be observed as a (possibly intermittent) ultraluminous X-ray source (ULX). We caution however that, before linking this scenario to the ULXs observed in nearby young massive clusters (see e.g. Portegies Zwart et al. 2010), one must first verify whether these clusters (1) meet the conditions to form a seed black hole, (2) are sufficiently massive and (3) have the structural properties required for SG star formation. A second important issue is the possible relation between the IMBH mass and the current structural properties of the parent cluster. Any correlation between the mass of the seed black hole and the early properties of the cluster should be preserved during the Eddington growth phase, provided that the duration $\\Delta T$ of the SG star formation episode is approximately independent of the cluster environment. The mass of the seed black hole depends on its formation mechanism. If it is the result of stellar evolution in a massive Pop. III progenitor, a roughly universal mass seems the most likely result (see e.g. Madau \\& Rees 2001). Early dynamical simulations of the runaway merger scenario suggested that the mass of the runaway should scale with the mass of the cluster (or the cluster core) in which it forms (G\\\"urkan et al. 2004, Portegies Zwart et al. 2004). However, as mentioned in \\S\\ref{sec:intro}, strong stellar winds may well limit the actual mass attained, possibly to as little as a few hundred solar masses, again largely independent of the initial cluster mass. For the black hole growth phase to occur, the cluster must {\\em initially} have been massive and/or compact enough to retain the AGB ejecta that subsequently flowed into the cluster core and formed the SG population. As discussed in D'Ercole et al. (2008), multiple population clusters must have undergone an early phase of strong FG mass loss and structural evolution; it is therefore not straightforward to connect the initial properties of a cluster during the growth of the seed black hole with the current properties of the cluster hosting an IMBH. However, given the loss of most of the FG population and the possible near-universality of the seed black hole mass, the most likely outcome of the scenario described here seems to be one in which the present IMBH mass has only a weak dependence on current cluster properties. $~~~~$\\\\ {\\bf Acknowledgments.} EV and SM were supported in part by NASA grants NNX07AG95G, NNX08AH15G, and NNX10AD86G and by NSF grant AST-0708299. AD acknowledges financial support from italian MIUR through grant PRIN 2007 (prot. 2007JJC53X). FD was supported by PRIN MIUR 2007 'Multiple Stellar Populations in Globular Clusters: Census, Characterization and Origin' (prot. n. 20075TP5K9)." }, "1003/1003.3193_arXiv.txt": { "abstract": "Identified as extinction features against the bright Galactic mid-infrared background, infrared dark clouds (IRDCs) are thought to harbor the very earliest stages of star and cluster formation. In order to better characterize the properties of their embedded cores, we have obtained new 24\\,\\um, 60--100\\,\\um, and sub-millimeter continuum data toward a sample of 38 IRDCs. The 24\\,\\um\\, \\Spitzer\\, images reveal that while the IRDCs remain dark, many of the cores are associated with bright 24\\,\\um\\, emission sources, which suggests that they contain one or more embedded protostars. Combining the 24\\,\\um, 60--100\\,\\um, and sub-millimeter continuum data, we have constructed broadband spectral energy distributions (SEDs) for 157 of the cores within these IRDCs and, using simple gray-body fits to the SEDs, have estimated their dust temperatures, emissivities, opacities, bolometric luminosities, masses and densities. Based on their \\Spitzer/IRAC 3--8\\,\\um\\, colors and the presence of 24\\,\\um\\, point source emission, we have separated cores that harbor active, high-mass star formation from cores that are quiescent. The active `protostellar' cores typically have warmer dust temperatures and higher bolometric luminosities than the more quiescent, perhaps `pre-protostellar', cores. Because the mass distributions of the populations are similar, however, we speculate that the active and quiescent cores may represent different evolutionary stages of the same underlying population of cores. Although we cannot rule out low-mass star-formation in the quiescent cores, the most massive of them are excellent candidates for the `high-mass starless core' phase, the very earliest in the formation of a high-mass star. ", "introduction": "The earliest phase of isolated low-mass star-formation occurs within Bok Globules. Viewed against background stars, Bok globules are identified as isolated, well-defined patches of optical obscuration and have typical visual extinctions, \\av, of 1--25 mag \\citep{Bok47}. The individual precursor to a low-mass star, referred to as the `pre-protostellar core', are found within Bok Globules. These pre-protostellar cores are typically small ($\\sim$0.05\\,pc) and dense (10$^{5}$--10$^{6}$\\,\\cmc), with low temperatures ($\\sim$10\\,K) and low masses (0.5--5\\,\\Msun; e.g., \\citealp{Myers83,Ward-Thompson94}). Because many low-mass star-forming regions are nearby, their pre-protostellar cores and protostars have been studied extensively. Moreover, these studies have the added benefit of achieving sufficiently high spatial resolutions to distinguish the individual protostars, allowing one to study and characterize their various evolutionary stages, i.e., Class-0, I, II, and III \\citep{Lada84,Adams87,Andre94}. These stages are characterized by IR spectral energy distributions (SEDs) corresponding to increasing black-body temperatures. Moreover, differences in the shapes of their SEDs trace the gradual removal of circumstellar material surrounding the central protostar as it evolves and emerges from its natal core. While these early evolutionary stages are well-characterized for low-mass star-formation, it has been difficult to classify high-mass protostars in a similar manner. The combination of large distances to high-mass star-forming regions, their rarity and rapid evolution, and the fact that most high-mass stars form deeply embedded in dense molecular clumps and within clusters with many lower-mass stars nearby, makes their identification and separation difficult. However, in a recent study combining \\MSX, \\IRAS, and sub-millimeter data toward 42 regions of high-mass star formation, \\citet{Molinari08} have investigated the evolution of SEDs for young, high-mass protostars and has attempted to classify them via their SEDs. They find that objects in apparently different evolutionary stages occupy different areas in a bolometric luminosity versus envelope mass diagram, in a similar manner to the low-mass regime and, thus, conclude that high-mass star-formation may be a scaled up analog to low-mass star-formation. If this is the case, then perhaps the early stages of high-mass star-formation can also be characterized via differences in the SEDs of their dense molecular cores. In order to better characterize the properties of high-mass pre-protostellar and protostellar cores, a large sample of cores in the very earliest stages of high-mass star-formation is required. High-mass stars and clusters form from cold, dense molecular clumps within giant molecular clouds \\citep{Blitz91,Blitz99,Lada03}. Recent studies suggest that the cold precursors of warm cluster-forming molecular clumps can be identified as `infrared dark clouds' (IRDCs; \\citealp{Simon-msxgrs,Rathborne06}). Surveys of the Galactic Plane at mid-IR wavelengths made with the \\ISO\\, and \\MSX\\, satellites first identified IRDCs as dark extinction features seen in absorption against the bright mid-IR emission arising from the Galactic background \\citep{Perault96,Carey98,Hennebelle01}. IRDCs are ubiquitous across the Galaxy \\citep{Simon-catalog} and are typically long and very filamentary. They correspond to the densest parts of much larger giant molecular clouds \\citep{Simon-msxgrs} and are characterized by high densities ($>10^{5}$\\,cm$^{-3}$), high column densities ($\\sim$10$^{23}$--10$^{25}$\\,cm$^{-2}$), and low temperatures ($<25$\\,K; \\citealp{Egan98,Carey98,Carey00}). Recently, \\cite{Rathborne06} conducted a survey of the 1.2\\,mm continuum emission toward 38 IRDCs using MAMBO-II on the IRAM 30\\,m telescope. These IRDCs were selected from the catalog of \\MSX\\, IRDC candidates \\citep{Simon-catalog} and all have known kinematic distances, determined via the morphological match of \\tcol\\, emission to the mid-IR extinction \\citep{Simon-msxgrs,Jackson06}. In these IRDCs \\cite{Rathborne06} found 190 cores, 140 of which are cold, compact cores. These cold, compact cores have typical sizes of $<$~0.5~pc and masses of $\\sim$~120\\,\\Msun\\, \\citep{Rathborne06}. Indeed, millimeter and sub-millimeter observations toward other IRDCs suggest such compact cores are ubiquitous within IRDCs (e.g., \\citealp{Lis94,Carey00,Garay04,Ormel05,Beuther05}). Because IRDCs are cold, their thermal dust emission will peak in the far-IR/sub-millimeter regime. At these wavelengths the dust emission is optically thin, making this regime the best for probing their internal structure and revealing their star-forming cores. If IRDCs are the high-mass analogue to Bok globules and the cold precursors to cluster-forming molecular clumps, then these dense cores may be the precursors to the stars \\citep{Rathborne06}. IRAM and JCMT molecular line spectra, \\Spitzer\\, 3--8\\,\\um\\, and 24\\,\\um\\, continuum images, and GBT water and methanol maser spectra toward a sample of cores within IRDCs reveal that most contain little evidence for active star-formation, such cores are called `quiescent' \\citep{Chambers09}. However, some of the IRDC cores do appear to be actively forming stars as they show broad molecular line emission, shocked gas, bright 24\\,\\um\\, emission, and strong water and methanol maser emission (e.g., \\citealp{Rathborne05,Wang06,Chambers09}). Each of these tracers provides independent evidence for star-formation, either indirectly (from the interaction between the protostar and the surrounding core traced through the broad line emission and the shocked gas) or more directly (from the heating of the dust surrounding the central protostar traced via the bright 24\\,\\um\\, emission). It is likely, therefore, that these particular cores contain protostars. Indeed, a number of protostars have already been identified within IRDCs and span a range in mass, from low- and intermediate-mass \\citep{Carey00,Redman03} to high-mass \\citep{Beuther05,Rathborne05,Pillai06,Wang06}. In order to characterize the cores within IRDCs, we have conducted a large, multi-wavelength observational survey, combining IR, sub-millimeter, and millimeter continuum data. These data, when combined to make broadband SEDs, provide estimates of the core dust temperatures, dust emissivities, opacities, bolometric luminosities, and masses. Because the cores have SEDs that peak in the far-IR, to date, it has been very difficult to estimate bolometric luminosities of the individual protostars due to the uncertain extrapolation from much longer and/or shorter wavelengths. Indeed, many previous studies using \\MSX\\, and \\IRAS\\, fluxes have been forced to make assumptions about the relative contributions to the far-IR flux from individual objects within a star-forming region (e.g. \\citealp{Molinari08}). Thus, the inclusion of the sub-millimeter and far-IR data is critical to determine reliable SEDs for the cores and to better constrain these parameters. In this paper we provide the first mid-IR, far-IR, sub-millimeter, and millimeter SEDs for a large sample of cores within IRDCs. ", "conclusions": "To characterize the physical properties of cores within IRDCs we have obtained new 24\\,\\um, 60--100\\,\\um, and sub-millimeter continuum data toward a sample of 38 IRDCs. These IRDCs contain 190 compact cores, 140 of which are dark at 8\\,\\um\\, and are cold, compact and dense. The \\Spitzer/MIPS 24\\,\\um\\, images reveal that while the IRDCs remain dark, many of their cores are associated with bright 24\\,\\um\\, emission sources. The sub-millimeter continuum data elucidate both the large- and small-scale structure of the IRDCs. Because emission at millimeter/sub-millimeter wavelengths can trace either temperature and/or density enhancements one needs an additional method to distinguish between these two parameters. Using the presence or absence of 24\\,\\um\\, point source emission, in combination with their \\Spitzer/IRAC 3--8\\,\\um\\, colors, we have classified the cores into five groups: red, active, intermediate, quiescent, and blue. We find that, of the 190 cores, 35 can be classified as red, 38 as active, 32 as intermediate, 79 as quiescent, and 6 as blue. From gray-body fits to their spectral energy distributions (SEDs) we have determined the dust temperatures, emissivities, opacities, bolometric luminosities, and masses for a large sample of the IRDC cores. The derived distributions of the dust temperatures, luminosities, and masses for the different groups of cores reveals that the dust temperatures and luminosities are higher for those cores that show active, high-mass star formation compared to those cores that are more quiescent. Lower dust temperatures and luminosities are expected for the quiescent cores because they presumably have no high-mass internal source to significantly heat the dust. Comparing the derived masses for the core samples, however, we find that their mass distributions are similar. We interpret this similarity to be a result of evolutionary differences: the cooler quiescent cores may be the pre-protostellar precursors to the warmer, more active protostellar cores. Using their derived bolometric luminosities, we estimate that $\\sim$10\\% of the cores that show evidence for star formation may contain high-mass protostars. If the quiescent cores are indeed devoid of star formation, then the most massive of these are excellent candidates for the `high-mass starless core' phase, a very early phase in the formation of a high-mass star. Because of their distances, we cannot yet rule out the possibility that many of these cores may contain low-mass stars. Observations with ALMA will be crucial to address this issue. Nevertheless, our study supports the idea that IRDCs harbor the very earliest evolutionary stages in the formation of high-mass stars and, thus, clusters." }, "1003/1003.4250_arXiv.txt": { "abstract": "The optical-infrared afterglow of the LAT-detected long duration burst, GRB 090902B, has been observed by several instruments. The earliest detection by ROTSE-IIIa occurred 80 minutes after detection by the GBM instrument onboard the {\\it Fermi} Gamma-Ray Space Telescope, revealing a bright afterglow and a decay slope suggestive of a reverse shock origin. Subsequent optical-IR observations followed the light curve for 6.5 days. The temporal and spectral behavior at optical-infrared frequencies is consistent with synchrotron fireball model predictions; the cooling break lies between optical and XRT frequencies $\\sim$ 1.9 days after the burst. The inferred electron energy index is $p = 1.8 \\pm 0.2$, which would however imply an X-ray decay slope flatter than observed. The XRT and LAT data have similar spectral indices and the observed steeper value of the LAT temporal index is marginally consistent with the predicted temporal decay in the radiative regime of the forward shock model. Absence of a jet break during the first 6 days implies a collimation-corrected $\\gamma$-ray energy $E_{\\gamma} > 2.2\\times10^{52}\\rm$ ergs, one of the highest ever seen in a long-duration GRBs. More events combining GeV photon emission with multi-wavelength observations will be required to constrain the nature of the central engine powering these energetic explosions and to explore the correlations between energetic quanta and afterglow emission. ", "introduction": "The recently launched {\\it Fermi} Gamma-Ray Space Telescope with the on-board Gamma-ray Burst Monitor (GBM) and Large Area Telescope (LAT) instruments \\citep{atwood09, meegan09} in conjunction with the {\\it Swift} narrow field instruments \\citep{gehrels04} have opened a new window to understand the physical mechanisms that generate GeV photons in very energetic GRBs and the relation to lower energy components of the afterglow \\citep{band09}. Since the {\\it Fermi} launch more than a year ago, only 14 GRBs have been detected by the LAT while more than $\\sim$ 350 bursts have been seen by the GBM during the same period. Optical afterglows have been detected for 7 of the 14 LAT events starting from $\\sim$ 300 s to a few hours after the burst. The origin of these high-energy photons and their possible correlation to afterglow emission is still debated \\citep[and references therein]{zou09}. GRB 080916C \\citep{abdo09b}, GRB 090510 \\citep{abdo09c, max10} and GRB 090902B \\citep{abdo09a} are among the brightest LAT bursts indicating some signatures consistent with the synchrotron forward shock models \\citep{kumar09a, kumar09b, ghirlanda10}. However, in the case of GRB 090902B, the deviation of the burst spectrum from the Band function and the observed large amplitude variability at very short time-scales \\citep{abdo09a} does not support the afterglow origin of the LAT data. High energy photons from GRBs have previously been observed by the EGRET detector and has shown evidence for deviations from synchrotron models \\citep{hurley94, gonzalez03}. The bright GRB 090902B (trigger 273582310) was detected by the GBM on 2009 Sep 2nd at 11:05:08.31 UT with an initial error box radius of 2-3 degrees centered at R.A. = 17${^h}$ 38${^m}$ 26${^s}$, Dec. = +26$\\degr$ 30$\\arcmin$ and a burst duration of 21.9 s in the energy band 50 - 300 keV \\citep{bissaldi09}. The burst was one of the brightest at LAT energies with a power-law spectral distribution {at both} low and high energies \\citep{depalma09}. The detailed analysis of the LAT and GBM data has been presented in \\citet{abdo09a}. The observed features in the prompt burst spectrum have also shown evidence for an underlying photospheric thermal emission \\citep{ryde10}. ToO observations with {\\it Swift} started $\\sim$ 12.5 hours after the GBM trigger. The X-ray afterglow was detected within the LAT error-circle by the XRT \\citep{kennea09}, the UVOT \\citep{swenson09} and later by several other ground-based multiwavelength facilities. The burst redshift, $z$ = 1.822, was determined by the Gemini-North telescope \\citep{cucchiara09}. The afterglow was also seen at radio frequencies by the WSRT \\citep{vander09} and by the VLA \\citep{chandra09}. The details of the optical-infrared (IR) observations along with the temporal and spectral properties of the afterglow are described in the next section. In \\S 3, we discuss the observed properties of the afterglow and comparisons to various models. These results are summarized in \\S 4. Throughout the paper, we use the usual power-law representation of flux density, $f{_\\nu}(t) \\propto \\nu^{-\\beta} t^{-\\alpha}$, for the regions without spectral breaks where $\\alpha$ and $\\beta$ are the power-law temporal decay and spectral indices, respectively. For the cosmological calculations we have used the cosmological parameters $H_0=71~\\rm km~s^{-1}Mpc^{-1}$, $\\Omega_{M}=0.27$, $\\Omega_{\\Lambda}=0.73$. Errors are quoted at the 1-sigma level unless otherwise stated. ", "conclusions": "\\begin{figure}[t] \\vspace{-1.3cm} \\includegraphics[scale=0.45]{Fig2.eps} \\caption{Multi-wavelength spectral energy distribution of the GRB 090902B afterglow at optical-IR and XRT frequencies derived at 1.9 days post-burst. The epoch has been chosen to allow the best possible spectral coverage. The observations in the $u$ band might be effected by Ly-$\\alpha$ and hence were excluded to determine the spectral index $\\beta_O$. } \\end{figure} \\subsection{The Optical Brightness and LAT-Detected Bursts} The power-law decays seen in other early optical afterglows of GRBs \\citep{panaitescu08, oates09b} suggest that the single observed data point at $\\sim$ 1.4 hours is unlikely to be a flaring feature at such late times. For an observation $\\sim$ 1.4 hours after the burst, the ROTSE detection at $m_R \\sim$ 16.4 mag is remarkably bright. This is best qualified by the statistical study of a large ensemble of bursts afterglows by \\citet{akerlof07}. They found that the temporal evolution of the brightness distribution is well described by a power-law exponent, $\\alpha \\sim$ 0.7. With that behavior, a magnitude of 16.4 at 1.4 hours would have evolved from $m_R \\sim$ 15.2 at t = 1000s and thus lie among the top 5 \\% of all bursts. Even one-magnitude errors in the ROTSE measurement would not substantially modify this conclusion. At later times, the optical afterglow must drop with a much steeper slope of $>$ 1.6. If this behavior was manifested earlier, the brightness of GRB 090902B was even more pronounced. The apparent temporal decay index, $\\alpha >$ 1.6, between 1.4 to 12.5 hours is steeper than the value of $\\alpha = 0.90 \\pm 0.08$ for epochs $>$ 12.5 hours by more than 7-$\\sigma$ and is consistent with the dominance of reverse shock origin \\citep{sari99a, kobayashi00, zhang03} as seen recently for the energetic `naked eye' GRB 080319B, on similar time scales \\citep[e.g.] [] {bloom09, pandey09b}. Comparison of the observed apparent optical brightness of GRB 090902B at $\\sim$ 1.4 hours to a much larger sample of pre-{\\it Swift} and {\\it Swift} optical afterglows \\citep[See Figure 1] [] {kann07} also indicates that the GRB 090902B was one of the brightest at such early epochs. A handful of other examples of long-duration GRBs, detected by LAT and with measured redshift values (GRB 080916C \\citep{greiner09}; GRB 090323 \\citep{updike09}; GRB 090328 \\citep{oates09a}; GRB 090926 \\citep{haislip09} and GRB 091003 \\citep{gronwall09} have been observed at optical frequencies starting $\\sim$ 16 to 26 hours post-burst. Comparison of the observed apparent optical brightness of the LAT-detected bursts 16-26 hours post-burst to the sample published in \\citet{kann07} indicates that the optical brightness of LAT-detected GRBs are typical except GRB 090926 \\citep{haislip09} which was one of the brightest even at $\\sim$ 20.0 hours post-burst. Thus, the late time behavior of LAT-detected GRBs at optical frequencies is not unusual. \\subsection{Afterglow Models and GRB 090902B} The derived values of temporal and spectral indices from multi-wavelength data can be compared with the closure relations \\citep{price02} to discriminate between interstellar medium (ISM) and wind ambient profiles \\citep{sari98, chevalier00} and to infer the location of the cooling break, $\\nu_c$. For the observed values of $\\alpha$ ($>$ 12.5 hours) and $\\beta_O$ at optical frequencies, the closure relation $\\alpha = 3\\beta/2$ is satisfied within errors in case of the ISM model \\citep{price02} for the observed frequencies $\\nu < \\nu_{c}$. Also, the value of the temporal decay index at XRT frequencies is steeper than for the optical, which clearly rules out the wind model and requires the ISM model with $\\nu_{c}$ between optical and XRT frequencies. The location of $\\nu_{c}$ below XRT frequencies implies that the electron energy index is $p = 1.8\\pm0.2$, deduced solely using the value of $\\beta_X$. The value of $p$ and the determined temporal slopes at optical frequencies are also consistent with the closure relation $\\alpha = 3(p-1)/4$ within errors, valid for the spectral regime $\\nu_{m} < \\nu < \\nu_{c}$ in case of the ISM model. However, the observed temporal decay index at XRT frequencies is inconsistent with the ISM model closure relation $\\alpha = (3p-2)/4$ (for $\\nu > \\nu_{c}$) by 2.8$\\sigma$ and requires a steeper value of $p$ than estimated. The predicted value of $p$ for $\\nu > \\nu_{c}$ is 2.4$\\pm$0.05 using the XRT temporal decay index $\\alpha$ = 1.30$\\pm$0.04. The afterglow properties favor an evolution of $\\nu_{c}$ between the XRT and optical frequencies during the observations with an expected $\\beta_O$ of $0.4\\pm0.1$ for $p = 1.8\\pm0.2$. The relatively shallower value of $\\beta_O$ than observed can be attributed to a moderate amount of extinction $A_V$ = 0.20$\\pm$0.06 mag for SMC-like dust and assuming $\\leq$ 20\\% of the $u$ flux is affected by Ly-$\\alpha$ at the SED epoch using the method described in \\citet{perley08}. The present optical-IR and XRT data have determined the value of $\\nu_c$ that is contrary to the assumption of \\citep{kumar09b} that $\\nu_{c}$ lies above the XRT frequencies and thus implies a steeper value of $p$. In the light of above discussions, the published radio data at 4.8 GHz \\citep{vander09} and 8.46 GHz \\citep{chandra09} of GRB 090902B near the SED epoch was used to constrain location of the self absorption frequency $\\nu_a$. The expected value of the spectral index between 4.8 and 8.46 GHz will be $\\sim$ -0.4, closer to the expected $\\nu^{-1/3}$ spectral regime for $\\nu_a < \\nu_m$ in the case of slow-cooling forward shock model \\citep{sari98}. The effect of scintillation has not been taken into account which might modify the flux values for the observed frequencies at early epochs. Using the observed flux values at the radio frequencies and assuming $\\nu_m$ to be $< 5.0\\times10^{14}$ Hz at $\\sim 10^{4}$ s after the burst, the estimated value of the peak synchrotron flux at the SED epoch is $\\sim$ 0.5 milliJy. The values of the peak synchrotron flux and $\\nu_m$ at the epoch of SED are used to constrain the value of $\\nu_a$ using equation 4.9 of \\citet{sari01}. The calculated value of $\\nu_a < 10^{8}$ Hz is below the observed radio frequencies and in agreement with the slow-cooling model for $\\nu_a < \\nu_m$ at the epoch of the SED. The analysis also indicates no signature of a possible jet-break before or during the period of our afterglow observations. For the measured fluence between 10 keV and 10 GeV \\citep{abdo09a}, the inferred value of the isotropic equivalent energy is $E_{\\gamma}^{iso\\rm}=3.6\\times10^{54}\\rm$ ergs assuming a gamma-ray efficiency $\\eta_\\gamma =$ 0.2 and the circumburst density $n =$ 1 cm$^{-3}$ \\citep{frail01}. Based on the observed properties of the burst, if we limit the jet-break time to be greater than 6 days after the burst, the value of the jet opening angle is $\\theta_j > 0.11$ rad which gives the collimation corrected energy $E_{\\gamma} > 2.2\\times10^{52}\\rm$ ergs, one of the highest ever inferred \\citep{cenko10}. With known values of $p$, the measured XRT flux at 1 day after the burst and using the description given in \\citet{freedman01}, the isotropic fireball energy carried by electrons is $\\epsilon_eE = 3.1\\times10^{54}\\rm$ ergs, where $\\epsilon_e$ is the fraction of shock energy carried by relativistic electrons, comparable to $E_{\\gamma}^{iso\\rm}$ \\citep[see also] []{starling09, tanvir09}. The constraint on the energetics of the burst is comparable to the energy budget in the case of magnetars \\citep{usov92, starling09, cenko10} and could also be accommodated within the ``collapsar'' origin of GRBs \\citep{mac99}. The energy estimates of more LAT-detected GRBs in the future will help towards a better understanding the nature of the central engine powering these energetic events. \\subsection{Onset of the GeV Afterglow} The detection of many delayed photons at energies $>$ 1 GeV, the observed high value of the isotropic $\\gamma$-ray energy and the very early peak time seen in the LAT light curve constrain the value of the bulk Lorentz factor $\\Gamma$ to be $\\sim 1000$ \\citep{abdo09a}. Such high values of $\\Gamma$ have also been estimated in the case of other LAT bursts, GRB 080916C \\citep{abdo09b, greiner09} and GRB 090510 \\citep{abdo09c, ghirlanda10}. Along with the afterglow properties discussed in the previous sections, the very high value of $\\Gamma$ and the very early peak in the LAT light curve provide a good opportunity to test the LAT temporal decay and spectral properties \\citep{abdo09a} for an early onset of the afterglow in terms of synchrotron shock models \\citep{sari98, sari99b}. Under the synchrotron fireball model, a power-law distribution can be assumed in both time and spectral domains. Based on the discussions in the previous section and assuming $\\nu_c \\sim 2\\times10^{16}$ Hz at the epoch of SED, the extrapolated values of $\\nu_m$ and $\\nu_c$ at 100.0 s are $< 5.0\\times10^{17}$ Hz and $\\sim 8\\times10^{17}$ Hz respectively (assuming a temporal scaling of $\\nu_c \\propto t^{-1/2}$ and $\\nu_m \\propto t^{-3/2}$), both below 10 keV. At the SED epoch, the observed flux density of 0.03$\\pm$0.01~$\\mu$Jy (at 2.88 keV) will give rise to an extrapolated flux density of 450$\\pm$150~$\\mu$Jy at 100.0 s assuming the temporal decay of $\\sim$ 1.3. The 1 GeV flux density calculated at 100.0 s \\citep{abdo09a} is $\\sim$ 0.004 $\\mu$Jy. These flux densities at 1 GeV and 2.88 keV imply a spectral index of $\\sim$ 0.9 at 100.0 s, in agreement with the XRT spectral index at the epoch of SED and the LAT spectral index within 2-$\\sigma$ \\citep{abdo09a, ghisellini10}. This indicates that for GRB 090902B, both XRT and LAT frequencies share the same spectral regime with $\\nu_c$ below XRT frequencies under the synchrotron model, although the temporal index of $\\sim$ 1.5, observed at LAT frequencies is steeper than the XRT temporal decay index 1.3$\\pm$0.04. However, our results show that around 100.0 s, $\\nu_m < \\nu_c$ and the observed temporal decay index at LAT frequencies is marginally consistent with the expected temporal decay index of (2 - 6p)/7 in the radiative case of the synchrotron model \\citep{sari98}. Recently, based on the bolometric afterglow luminosity estimates for radiative fireballs, the expected temporal decay index $t^{10/7}$ for the LAT frequencies \\citep{ghisellini10} is also close to the observed LAT temporal index of $\\sim$ 1.5. In the case of another LAT-detected GRB 080916C, the value of spectral index at LAT frequencies is $1.1\\pm0.1$ \\citep{ghisellini10} and the value of GeV flux density at 100.0 s is $\\sim$ 0.006 $\\mu$Jy \\citep{abdo09b}. Using the XRT data analysis published in \\citep{greiner09}, the extrapolated value of 2.88 keV flux density at 100.0 s is 250$\\pm$50~$\\mu$Jy. For GRB 080916C, the spectral index between 2.88 keV and 1 GeV at 100.0 s comes out to be $\\sim$ 0.9, close to the spectral index seen at the LAT frequencies \\citep{ghisellini10}. This also indicates that $\\nu_c$ is between XRT and LAT frequencies at 100.0 s within the assumptions of the afterglow model proposed by \\citep{greiner09}. In the case of GRB 090510, the multi-wavelength SED at 100.0 s after the burst supports the afterglow origin of LAT data in terms of synchrotron forward shock model with a possible energy injection at optical and XRT frequencies \\citep{max10}. The observed values of XRT flux for GRB 080916C, GRB 090510 and GRB 090902B are typical for other well observed Swift GRBs\\footnote[1]{http://www.swift.ac.uk/xrt\\_curves} \\citep{zheng09} at similar time scales. Based on above discussion, the evidence in favor of synchrotron forward shock model for the observed GeV emission for GRB 090902B, GRB 090510 and GRB 080916C is consistent with the predictions made by \\citep{zou09}. However, the hard photon index at LAT frequencies and the evidence for reverse shock emission in early optical data cannot rule out the possibilities of synchrotron self-Compton emission at LAT frequencies \\citep{wang01} and the Klein-Nishina suppression of high energy electrons at early times \\citep{wang10}. Such processes would require theoretical modeling which is beyond the scope of this paper." }, "1003/1003.3646_arXiv.txt": { "abstract": "{Degenerate ignition of helium in low-mass stars at the end of the red giant branch phase leads to dynamic convection in their helium cores. One-dimensional (1D) stellar modeling of this intrinsically multi-dimensional dynamic event is likely to be inadequate. Previous hydrodynamic simulations imply that the single convection zone in the helium core of metal-rich Pop I stars grows during the flash on a dynamic timescale. This may lead to hydrogen injection into the core, and a double convection zone structure as known from one-dimensional core helium flash simulations of low-mass Pop III stars. } {We perform hydrodynamic simulations of the core helium flash in two and three dimensions to better constrain the nature of these events. To this end we study the hydrodynamics of convection within the helium cores of a 1.25 \\Msun metal-rich Pop I star (Z=0.02), and a 0.85 \\Msun metal-free Pop III star (Z=0) near the peak of the flash. These models possess single and double convection zones, respectively.} {We use 1D stellar models of the core helium flash computed with state-of-the-art stellar evolution codes as initial models for our multidimensional hydrodynamic study, and simulate the evolution of these models with the Riemann solver based hydrodynamics code Herakles which integrates the Euler equations coupled with source terms corresponding to gravity and nuclear burning.} {The hydrodynamic simulation of the Pop I model involving a single convection zone covers 27 hours of stellar evolution, while the first hydrodynamic simulations of a double convection zone, in the Pop III model, span 1.8 hours of stellar life. We find differences between the predictions of mixing length theory and our hydrodynamic simulations. The simulation of the single convection zone in the Pop I model shows a strong growth of the size of the convection zone due to turbulent entrainment. Hence we predict that for the Pop I model a hydrogen injection phase (\\ie hydrogen injection into the helium core) will commence after about 23 days, which should eventually lead to a double convection zone structure known from 1D stellar modeling of low-mass Pop III stars. Our two and three-dimensional hydrodynamic simulations of the double (Pop III) convection zone model show that the velocity field in the convection zones is different from that predicted by stellar evolutionary calculations. The simulations suggest that the double convection zone decays quickly, the flow eventually being dominated by internal gravity waves.} {} ", "introduction": "\\label{sect:intro} Runaway nuclear burning of helium in the core of low mass red giant stars leads to convective mixing and burning on dynamic time scales. One dimensional evolutionary simulations (which assume time scales much longer than the dynamical ones) may miss key features of this rapid phase that could have significant effects on the further evolution of the stars. Furthermore, 1D hydrodynamical simulations of this intrinsically multi-dimensional event is likely to be inadequate. Our previous hydrodynamic simulations \\citep{Mocak2008, Mocak2009} imply that a $1.25\\,$\\Msun solar-like star can experience injection of hydrogen into its helium core during the canonical core helium flash near its peak. Hydrogen injection results from the growth of the convection zone (which is sustained by helium burning) due to turbulent entrainment on a dynamic timescale \\citep{MeakinArnett2007}, and probably occurs for all low-mass Pop I stars, as the properties of their cores are similar at the peak of the core helium flash \\citep{SweigertGross1978}. An obvious consequence of this scenario is that the convection zones are enlarged in these stars. Whether they fail to dredge up nuclear ash to the atmosphere shortly after the flash is still unclear. However, such a dredge up could explain the Al-Mg anticorrelation found in red giants at the tip of the red giant branch \\citep{Shetrone1996a, Shetrone1996b, Yong2006}. In 1D simulations one has to manipulate the properties of the core helium flash to achieve such a dredge up, \\eg by changing the ignition position of the helium in the core \\citep{PaczynskiTremaine1977} or by forcing inward mixing of hydrogen \\citep{Fujimoto1999}. Canonical (1D) stellar evolution calculations predict hydrogen injection during the core helium flash and subsequent dredge-up of nuclear ashes to the atmosphere only for Pop III \\footnote{They are supposed to be the first stars in the Universe and seem to be extremely rare, as the most metal-poor star discovered up to now has a metallicity of [Fe/H] $\\sim -5.5$ \\citep{Frebel2005}.} and extremely metal-poor (EMP; with intrinsic metallicities [Fe/H] $\\lesssim -4$) stars. This is a promising scenario for explaining the peculiar abundances of carbon and nitrogen observed in Galactic EMP Halo stars \\citep{Campbell2008}. If these stars are assumed to be polluted by accretion of CNO-rich interstellar matter they will possibly experience hydrogen injection but no subsequent dredge-up, because a large CNO metallicitiy (as compared to the intrinsic [Fe/H] metallicity) in the stellar envelope influences the ignition site of the first major core helium flash, and thereby the occurrence of the dredge-up \\citep{Hollowell1990}. The helium abundance adopted in the stellar models also seems to influence the process of hydrogen injection itself as shown by \\citet{SchlattlCassisiSalaris2001}, while the same authors find that the injection process seems to be independent of the assumed convection model. Stellar models with a higher intrinsic metallicity, \\ie [Fe/H] $> -4$, do not inject hydrogen into the helium core, and consequently there is also no dredge-up of CNO-rich nuclear ashes to the atmosphere \\citep{FujimotoIben1990, Hollowell1990, Campbell2008}. Whether this is the final answer remains unclear, however, as \\citet{Fujimoto1999} with his semi-analytic study and a postulated hydrogen injection followed by a dredge-up could show that such a scenario can explain some peculiarities observed in the spectra of red-giant stars (related to CNO elements and $^{24}$Mg) with metallicities as large as [Fe/H] $\\lesssim -1$. There exist two main reasons why hydrogen injection episodes occur only in Pop III and EMP stars: ($i$) these stars possess a flatter entropy gradient in the hydrogen burning shell, and ($ii$) they ignite helium far off center, relatively close to the hydrogen-rich envelope \\citep{FujimotoIben1990}. However, Pop II and Pop I stars could also mix hydrogen into the helium core during the core helium flash \\footnote{According to \\citet{SchlattlCassisiSalaris2001} the occurrence of a hydrogen flash is favored by a higher electron degeneracy in the helium core, which leads to helium ignition closer to the hydrogen shell. } \\begin{itemize} \\item if the flash was more violent, and thus the helium convection zone wider \\citep{DespainSalo1976, Despain1981}. This scenario is disfavored by the facts that the flash is less violent in stars with higher metallicity as less energy is needed to lift the degeneracy of the less massive cores \\citep{SweigertGross1978}, and that helium ignition occurs at smaller densities \\citep{FujimotoIben1990}, \\item or if the entropy gradient between the hydrogen and helium burning shell was sufficiently shallow \\citep{Iben1976, Fujimoto1977}. A small entropy gradient would allow the convective shell in the helium core to reach the hydrogen layers even though the flash itself would not be very violent. This scenario is also disfavored as solutions to the stellar structure equations are robust with many different groups getting very similar results \\ie no hydrogen injection \\citep{FujimotoIben1990, Hollowell1990, Campbell2008} \\item or if a growth of the helium convection zone through turbulent entrainment at the convective boundaries \\citep{Mocak2008, Mocak2009} could be sustained for a sufficiently long period of time. \\end{itemize} A hydrogen injection phase also occurs in low-mass metal deficient stars on the asymptotic giant branch (AGB) at the beginning of the thermally pulsing stage (TPAGB) \\footnote{If hydrogen injection occurs at the tip of the red giant branch branch (RGB), it does not occur on the AGB -- the star evolves like a normal thermally pulsating Pop I or II star \\citep{SchlattlCassisiSalaris2001}}. Hydrogen injection is found to occur in more massive stars ($M \\gtrsim 1.3\\,$\\Msun) with low metallicity during the TPAGB \\citep{Chieffi2001, Siess2002, Iwamoto2004}, in ``Late Hot Flasher'' stars experiencing strong mass loss on the RGB \\citep{Brown2001, CassisiSchlattl2003}, and in H-deficient post AGB (PAGB) stars. These events are referred to with various names in the literature. Here we use the nomenclature ``dual flashes'' \\citep{Campbell2008}, since they all have in common simultaneous hydrogen and helium flashes. Dual flash events often lead to a splitting of the single helium convection zone (HeCZ) into two parts (double convection zone): one sustained by helium burning, and a second one by hydrogen burning via CNO cycles (Fig.\\,\\ref{fig.kipd}). Double convection zones are structures which are commonly encountered in stellar models, but their hydrodynamic properties have so far only been studied for the oxygen and carbon burning shell of a $23\\,$\\Msun star by \\citet{Meakin2006}. \\begin{figure} \\includegraphics[width=0.99\\hsize]{fig1-Ver5-labels.pdf} \\caption{ {\\it{Upper panel:}} Kippenhahn diagram of a stellar evolutionary calculation during the core helium flash of a $0.85\\,$\\Msun Pop III star with convection zones sustained by helium (He-rich) and hydrogen (H-rich) or CNO burning (grey shaded regions, except for the convective envelope). The border between the helium and hydrogen rich layers is given by the solid blue curve. The location of the initial model SC (model number 9120) - black vertical line. {\\it{Lower panel:}} the temporal evolution of the helium (dotted-blue) and hydrogen (solid-red) luminosity as a function of time. } \\label{fig.kipd} \\end{figure} In the following we describe two-dimensional (2D) and three-dimensional (3D) hydrodynamic simulations of a helium core during the core helium flash with a single convection zone (Pop I; in 3D only) and a double convection zone (Pop III, in 2D and 3D), respectively. Previous studies have indicated that there is a strong interaction between the adjacent shells of a double convection zone by internal gravity waves \\citep{Meakin2006}. We introduce the stellar models used as input for our hydrodynamic simulations in Sect.\\,\\ref{sect:regime}, briefly discuss the physics included in our simulations in Sect.\\,\\ref{sect:input}, and give a short description of our hydrodynamics code and the computational setup in Sect.\\,\\ref{sect:hcode}. Subsequently, we present and compare the results of our 2D and 3D hydrodynamic simulations in Sect.\\,\\ref{sect:tmpevol}. In particular, we discuss turbulent entrainment at the convective boundaries for our single convection zone model, the temporal evolution of its kinetic energy density, and how our results compare with the predictions of mixing-length theory (MLT). We proceed similarly for our hydrodynamic double convection zone model, except for turbulent entrainment for reasons which become clear later. Finally, a summary of our findings is given in Sect.\\,\\ref{sect:sum}. ", "conclusions": "\\label{sect:sum} We have performed and analyzed a 3D hydrodynamic simulation of a core helium flash near its peak in a Pop I star possessing a single convection zone (single flash) sustained by helium burning. The simulation covers 27\\,hrs of stellar life, or roughly 100 convective turnover timescales. In addition, we performed and analyzed 2D and 3D simulations of the core helium flash near its peak in a Pop III star which has a double convection zone (dual flash) sustained by helium and CNO burning, respectively. These simulations cover only 1.8\\,hrs and 0.39\\,hrs of stellar life, respectively, as convection dies out shortly after it appears. The convective velocities in our hydrodynamic simulation of the single flash model and those predicted by stellar evolutionary calculations agree approximately. Contrary to our previous findings, the temperature gradient in the convection zone remains superadiabatic, probably because of the increased spatial resolution of these simulations as compared to our old models. As expected, the simulation shows that the convection zone grows on a dynamic timescale due to turbulent entrainment. This growth can lead to hydrogen injection into the helium core as predicted by stellar evolutionary calculation of extremely metal-poor or metal-free Pop III stars. Hydrogen injection leads to a split of the single convection zone into two parts separated by a supposedly impenetrable radiative zone. Our hydrodynamic simulations of the double convection zone show that the two zones vanish as their convective motion decays very fast. However, this result may be caused by an insufficient spatial grid resolution or probably because the conditions represented by the stabilized initial model are a bit different from those of the original stellar model. While the convective velocities in our 2D hydrodynamic models do not match those predicted by stellar evolutionary calculations for the double convection zone at all, a rough agreement is found in our 3D model for the velocities in the inner convection zone sustained by helium burning." }, "1003/1003.0416_arXiv.txt": { "abstract": "{The study of planetary nebulae in the inner-disk and bulge gives important information on the chemical abundances of elements such as He, N, O, Ar, Ne, and on the evolution of these abundances, which is associated with the evolution of intermediate-mass stars and the chemical evolution of the Galaxy. We present accurate abundances of the elements He, N, S, O, Ar, and Ne for a sample of 54 planetary nebulae located towards the bulge of the Galaxy, for which 33 have the abundances derived for the first time. The abundances are derived based on observations in the optical domain made at the National Laboratory for Astrophysics (LNA, Brazil). The data show a good agreement with other results in the literature, in the sense that the distribution of the abundances is similar to those works. } \\resumen{} \\addkeyword{Galaxy: abundances} \\addkeyword{Galaxy: evolution} \\addkeyword{planetary nebulae: general} \\addkeyword{techniques: spectroscopic} \\begin{document} ", "introduction": "\\label{sec:intro} The bulge of the Galaxy shows a large metallicity dispersion. The study of the metallicity distribution from K giants, as done by Rich~(\\citeyear{rich88}), shows values from 0.1 to 10 Z$_{\\odot}$. More recently, Rich \\& Origlia~(\\citeyear{rich05}) find an $\\alpha$-enhancement at the level of $+0.3$ dex relative to the solar composition stars for 14 M giants and within a narrow metallicity range around [Fe/H$]=-0.2$. Zoccali et al.~(2006) and Lecureur et al.~(\\citeyear{lecu07}) find that bulge stars have larger values of [O/Fe] and [Mg/Fe] when compared to thin and thick disk stars. This is the signature of a chemical enrichment by massive stars, progenitors of type II supernovae, with little or no contribution from type Ia supernovae, showing a shorter formation timescale for the bulge than both thin and thick disks. In this context, planetary nebulae (PNe) are an important tool for the study of the chemical evolution of galaxies. The understanding of this stage of stellar evolution allows us to grasp how the Galaxy originated and developed. As an intermediate mass star evolution product, PNe offer the possibility of studying both elements produced in low and intermediate mass stars, such as helium and nitrogen, and also elements which result from the nucleosynthesis of large mass stars, such as oxygen, sulfur and neon, which are present in the interstellar medium at formation epoch of the PNe stellar progenitor. Regardless of the fact that the chemical abundances obtained from PNe are relatively accurate, their distances are subject of discussion even nowadays. Excluding a few PNe whose distances are determined from direct methods such as trigonometric parallax or in cases where there is a binary companion in the main sequence, most PNe have their distances derived from nebular properties (see e.g. Maciel \\& Pottasch~\\citeyear{maciel80}, Cahn et al.~\\citeyear{cahn92}, Stanghellini et al.~\\citeyear{stangh08}). These uncertainties in the distances of PNe make the study of the chemical properties with respect to the galactocentric distance a difficult task. In spite of the uncertainties, statistical distance scales are still the best tool to study the chemical abundance patterns in the Galaxy from the point of view of PNe, as e.g. done by Maciel \\& Quireza~(\\citeyear{maciel99}), Maciel et al.~(\\citeyear{maciel06}), Perinotto \\& Morbidelli~(\\citeyear{peri06}), and Gutenkunst et al.~(\\citeyear{guten08}). Since the bulge and the disk may have different evolution histories, described for example by the disk inside out formation model (Chiappini et al.~\\citeyear{chiap01}) or by the multiple infalls scenario (Costa et al.~\\citeyear{costa05},\\citeyear{costa08}), we should expect these differences reflected on the chemical properties of each component. Indeed, bulge and disk display different chemical abundance patterns like the radial abundance gradients found in the disk (Carigi et al. \\citeyear{carigi05}; Daflon \\& Cunha~\\citeyear{daflon04}; Andrievsky et al.~\\citeyear{andri04}; Maciel et al.~\\citeyear{maciel05},\\citeyear{maciel06}), or the large abundance distribution found in the bulge (Rich~\\citeyear{rich88}, Zoccali et al.~\\citeyear{zoc03,zoc06}). On the other hand, Chiappini et al.~(\\citeyear{chiap09}) made a comparison between abundances from PNe located at the bulge, inner-disk and Large Magellanic Cloud. Their results do not show any clear difference between bulge and inner-disk objects. Some other previous studies of the Galactic bulge based on abundances of PNe such as Ratag et al.~(\\citeyear{ratag92}), Cuisinier et al.~(\\citeyear{cuisin00}), Escudero \\& Costa (\\citeyear{escu01}), Escudero et al. (\\citeyear{escu04}), and Exter et al. (\\citeyear{exter04}), find that bulge PNe have an abundance distribution similar to disk PNe, showing that He, O, Si, Ar, and Ca have a normal abundance pattern, favouring therefore a slower Galactic evolution than that indicated by stars. In conclusion, the study of chemical abundances in the inner region of the Galaxy is still an open question, especially regarding the bulge-disk connection. The goal of this paper is to report new spectrophotometric observations for a sample of PNe located in the inner-disk and bulge of the Milky Way Galaxy, aiming to derive their nebular physical parameters and chemical abundances, as has been done by our group (see e.g. Costa et al. \\citeyear{costa96,costa00}, Escudero \\& Costa \\citeyear{escu01}, Escudero et al.~\\citeyear{escu04}, and references therein), as part of a long-term program to derive a large sample of chemical abundances of southern PNe. As a result, our database has become one of the largest in the literature with a very homogeneous observational setup, reduction and analysis procedures, which is necessary to perform large scale statistical studies. In this work, 33 objects have their abundances derived for the first time. Additionally, objects in common with other samples are used to compare our data with previously data already published. The comparison of the final abundances with those obtained in other multi-object studies allowed us to assess the accuracy of the new abundances. This paper is organized as follows: in \\S~\\ref{sec:obs_red} the details of the observations and data reduction procedures are presented. In \\S~3 we describe the process of determination of chemical abundances and the new abundances are listed. In \\S~4, a comparison is made between the abundances obtained in this work and those taken from the literature. Finally, in \\S~5 the main conclusions are presented. ", "conclusions": "In summary, this work reports an important result concerning PNe and the chemical evolution of the Galaxy. We present the analysis of chemical abundances of a PNe sample located towards the galactic bulge. New chemical abundances were derived through spectrophotometric observations made at the 1.60 m telescope of the LNA-Brazil, comprising the elements He, N, S, O, Ar, and Ne. 54 PNe were considered, among which 33 objects have their abundances derived for the first time. A comparison between the chemical abundances from this work and abundances obtained from the literature was performed. The analysis shows that the distributions of abundances are similar but not identical. Some objects of this work are listed in other investigations and a direct comparison between these abundances shows that the differences are of the order of 0.2 dex, indicating that the distinct methods used to derive the abundances are the main source of this difference. With the present results, we intend to enlarge the number of planetary nebulae with accurate chemical abundances, providing a large and homogeneous set of chemical abundances, contributing to the understanding of this stage of star evolution as well as the study of the chemical evolution of the inner Galaxy. \\bigskip {\\it Acknowledgements. Part of this work was supported by the Brazilian agencies \\emph{FAPESP} and \\emph{CNPq}. O.C. would like to acknowledge FAPESP for his graduate fellowship (processes 05/03194-4 and 07/07704-2). }" }, "1003/1003.2763_arXiv.txt": { "abstract": "We observed two secondary eclipses of the exoplanet WASP-12b using the Infrared Array Camera on the \\textit{Spitzer Space Telescope}. The close proximity of WASP-12b to its G-type star results in extreme tidal forces capable of inducing apsidal precession with a period as short as a few decades. This precession would be measurable if the orbit had a significant eccentricity, leading to an estimate of the tidal Love number and an assessment of the degree of central concentration in the planetary interior. An initial ground-based secondary eclipse phase reported by \\citeauthor{lopez:2009} (0.510 {\\pm} 0.002) implied eccentricity at the 4.5\\math{\\sigma} level. The spectroscopic orbit of \\citeauthor{hebb:2009} has eccentricity 0.049 {\\pm} 0.015, a 3\\math{\\sigma} result, implying an eclipse phase of 0.509 {\\pm} 0.007. However, there is a well documented tendency of spectroscopic data to overestimate small eccentricities. Our eclipse phases are 0.5010 {\\pm} 0.0006 (3.6 and 5.8 {\\microns}) and 0.5006 {\\pm} 0.0007 (4.5 and 8.0 {\\microns}). An unlikely orbital precession scenario invoking an alignment of the orbit during the {\\em Spitzer} observations could have explained this apparent discrepancy, but the final eclipse phase of \\citeauthor{lopez:2010} (0.510 {\\pm} \\sbp{-0.006}{+0.007}) is consistent with a circular orbit at better than 2\\math{\\sigma}. An orbit fit to all the available transit, eclipse, and radial-velocity data indicates precession at \\math{<1\\sigma}; a non-precessing solution fits better. We also comment on analysis and reporting for {\\em Spitzer} exoplanet data in light of recent re-analyses. \\if\\submitms y \\else \\comment{\\hfill\\herenote{DRAFT of {\\today} \\now}.} \\fi ", "introduction": "\\label{intro} When exoplanets transit (pass in front of) their parent stars as viewed from Earth, one can constrain their sizes, masses, and orbits \\citep{charbonneau:2007, winn:2009}. Most transiting planets also pass behind their stars (secondary eclipse). This allows atmospheric characterization by measurement of planetary flux and constrains orbital eccentricity, \\math{e}, through timing and duration of the eclipse \\citep{kallrath:1999}. WASP-12b is one of the hottest transiting exoplanets discovered to date, with an equilibrium temperature of 2516 K for zero albedo and uniform redistribution of incident flux \\citep{hebb:2009}. It also has a 1.09-day period, making it one of the shortest-period transiting planets. The close proximity to its host star \\citep[0.0229 {\\pm} 0.0008 AU,][]{hebb:2009} should induce large tidal bulges on the planet's surface. Tidal evolution should quickly circularize such close-in orbits \\citep{mardling:2007}. \\citet{hebb:2009} calculate a circularization time for WASP-12b as short as 3 Myr, much shorter than the estimated 2 Gyr age of WASP-12 or even the circularization times estimated for other hot Jupiters, given similar planetary tidal dissipation, though this calculation was based on a formalism \\citep{goldreich:1966} that ignores the influence of stellar tides and the coupling of eccentricity and semi-major axis in the evolution of the system. The influence of stellar tides could prolong the dissipation timescale to well over the age of the system \\citep{jackson:2008}. The non-Keplerian gravitational potential may cause apsidal precession, measurable as secondary eclipse and transit timing variations over short time scales. WASP-12b also has an abnormally large radius (\\math{R\\sb{\\rm{p}}} = 1.79 {\\pm} 0.09 Jupiter radii, \\math{R\\sb{\\rm{J}}}, \\citealp{hebb:2009}) compared to those predicted by theoretical models \\citep{bodenheimer:2003, fortney:2007} and to other short-period planets. Tidal heating models assume non-zero \\math{e}, and the heating rate can differ substantially for different values of \\math{e}. WASP-12b's inflated radius may result from tidal heating, but this is difficult to justify if the orbit is circular \\citep{li:2010}. Ground-based observations by \\citet{lopez:2009} detected a secondary-eclipse phase for WASP-12b of 0.510 {\\pm} 0.002, implying an eccentric orbit at the 4.5\\math{\\sigma} level (\\citealp{lopez:2010} revised the uncertainty to \\sbp{-0.006}{+0.007}). Radial velocity data \\citep{hebb:2009} find \\math{e} = 0.049 {\\pm} 0.015, a 3\\math{\\sigma} eccentricity, and predict an eclipse phase of 0.509 {\\pm} 0.007. Given an eccentric orbit and the fast predicted precession time scale, WASP-12b makes an excellent candidate for the first direct detection of exoplanetary apsidal precession. Such precession has been detected many times for eclipsing binary stars \\citep{kreiner:2001}. Against an orbit established by transit timings, precession would be apparent in just two eclipses, if sufficiently separated in time. For eccentric orbits, the eclipse-transit interval can differ from the transit-eclipse interval, and for precessing orbits this difference varies sinusoidally over one precession period. If the difference is insignificant, it places an upper limit on \\math{e \\cos \\omega}, where \\math{\\omega} is the argument of periapsis. In the case of WASP-12b, which is expected to precess at a rate of 0.05{\\degrees} d\\sp{-1} \\citep{ragozzine:2009}, if the orbit is observed when \\math{\\omega \\sim \\pm 90 \\degrees} and the effect on the eclipse timing is maxmized, and assuming a timing precision of 0.0007 days, then secondary eclipse observations situated five months apart could detect precession at the 3\\math{\\sigma} level (see Equation \\ref{eq:gimenmod}). We note that the method of \\citealp{batygin:2009}, based on the work of \\citealp{mardling:2007} and extended to the three-dimensional case by \\citealp{mardling:2010}, is an indirect assessment of apsidal precession, since no orbital motion is actually observed. The technique, which only applies to multi-planet systems with a tidally affected inner planet and a nearby, eccentric, outer planet, cannot currently be applied to WASP-12b. Paired with the \\citeauthor{lopez:2009} data, our \\textit{Spitzer Space Telescope}\\/ \\citep{werner:2004} eclipse observations provide a one-year baseline. \\textit{Spitzer's}\\/ high photometric precision also allows an accurate assessment of \\math{e \\cos \\omega}. One can solve for \\math{e} and \\math{\\omega} separately given \\math{e \\sin \\omega} from precise radial velocity data. The following sections present our observations; photometric analysis; a dynamical model that considers parameters from this work, the original and revised parameters of \\citeauthor{lopez:2009}, \\citeauthor{hebb:2009}, new transit times from the Wide-Angle Search for Planets (WASP), and transit times from a network of amateur astronomers; and our conclusions. \\if\\submitms y \\clearpage \\fi \\begin{figure}[htb] \\if\\submitms y \\setcounter{fignum}{\\value{figure}} \\addtocounter{fignum}{1} \\newcommand\\fignam{f\\arabic{fignum}.eps} \\else \\newcommand\\fignam{figs/wa012bs41_preflash.ps} \\fi \\includegraphics[width=\\columnwidth, clip]{f1.eps} \\figcaption{\\label{fig:preflash} Preflash light curve. These are channel-4 (8 {\\microns} data, analyzed with aperture photometry at the pixel location of the eclipse observations. The preflash source is bright compared to WASP-12, which allows the array sensitivity to ``ramp'' up before the science observations. Without a preflash, similar observations generally show a steeper and longer ramp in the eclipse observations.} \\end{figure} \\if\\submitms y \\clearpage \\fi \\if\\submitms y \\clearpage \\fi \\begin{figure*}[thb] \\if\\submitms y \\setcounter{fignum}{\\value{figure}} \\addtocounter{fignum}{1} \\newcommand\\fignama{f\\arabic{fignum}a.eps} \\newcommand\\fignamb{f\\arabic{fignum}b.eps} \\newcommand\\fignamc{f\\arabic{fignum}c.eps} \\else \\newcommand\\fignama{figs/WASP-12b-raw.ps} \\newcommand\\fignamb{figs/WASP-12b-bin.ps} \\newcommand\\fignamc{figs/WASP-12b-norm.ps} \\fi \\strut\\hfill \\includegraphics[width=0.32\\textwidth, clip]{f2a.eps} \\includegraphics[width=0.32\\textwidth, clip]{f2b.eps} \\includegraphics[width=0.32\\textwidth, clip]{f2c.eps} \\hfill\\strut \\figcaption{\\label{fig:lightcurves} Raw (left), binned (center), and systematics-corrected (right) secondary eclipse light curves of WASP-12b in the four IRAC channels, normalized to the mean system flux within the fitted data. Colored lines are the best-fit models; black curves omit their eclipse model elements. A few initial points in all channels are not fit, as indicated, to allow the telescope pointing and instrument to stabilize. } \\end{figure*} \\if\\submitms y \\clearpage \\fi \\if\\submitms y \\clearpage \\fi \\begin{figure}[thb] \\if\\submitms y \\setcounter{fignum}{\\value{figure}} \\addtocounter{fignum}{1} \\newcommand\\fignama{f\\arabic{fignum}a.eps} \\newcommand\\fignamb{f\\arabic{fignum}b.eps} \\newcommand\\fignamc{f\\arabic{fignum}c.eps} \\newcommand\\fignamd{f\\arabic{fignum}d.eps} \\else \\newcommand\\fignama{figs/wa012bs11-noisecorr.ps} \\newcommand\\fignamb{figs/wa012bs21-noisecorr.ps} \\newcommand\\fignamc{figs/wa012bs31-noisecorr.ps} \\newcommand\\fignamd{figs/wa012bs41-noisecorr.ps} \\fi \\includegraphics[width=0.22\\textwidth, clip]{f3a.eps} \\hfill \\includegraphics[width=0.22\\textwidth, clip]{f3b.eps} \\newline \\strut\\newline \\includegraphics[width=0.22\\textwidth, clip]{f3c.eps} \\hfill \\includegraphics[width=0.22\\textwidth, clip]{f3d.eps} \\figcaption{\\label{fig:noisecorr} Root-mean-squared (RMS) residual flux \\textit{vs.}\\/ bin size in each channel. This plot tests for correlated noise. The straight line is the prediction for Gaussian white noise. Since the data do not deviate far from the line, the effect of correlated noise is minimal.} \\end{figure} \\if\\submitms y \\clearpage \\fi ", "conclusions": "\\label{sec:concl} The timing of the {\\em Spitzer}\\/ eclipses is consistent with a circular orbit, and our best fit, including RV data and transit and eclipse times, does not detect precession. Although the \\citet{lopez:2010} eclipse phase is now marginally consistent with zero eccentricity, we note that this 0.9-{\\micron} observation could be affected by a wavelength-dependent asymmetry in the planet's surface-brightness distribution that manifests itself as a timing offset \\citep{knutson:2007}. This offset has a maximum possible value of \\math{R\\sb{\\rm{p}}/v\\sb{\\rm{p}} \\approx 9} minutes, where \\math{v\\sb{\\rm p}} is the planet's orbital velocity. This is somewhat smaller than the observed variation in eclipse timing between \\citet{lopez:2010} and \\textit{Spitzer}. While we have not yet measured precession, the possible prolateness should be measurable in high-accuracy, infrared transits and eclipses \\citep{ragozzine:2009}, such as we expect will be available from the James Webb Space Telescope. This would provide another constraint on interior structure, one that does not depend on an elliptical orbit. As this paper was in late stages of revision, \\citet{croll:2010} published three ground-based secondary eclipses and \\citet{husnoo:2010} produced additional radial velocity data. These datasets are consistent with a circular orbit for WASP-12b. As the quality of these data attests (the signal-to-noise ratio of the eclipse depth in channel 1 is over 29, second only to that for HD 189733b), WASP-12b has emerged as a highly observable exoplanet. \\citet{madhu:2010} report our analysis of the planet's atmospheric composition. Its phase curves, already in {\\em Spitzer's}\\/ queue, will enable the first observational discussion of atmospheric dynamics on a prolate planet. \\if\\submitms y \\clearpage \\fi \\atabon\\begin{deluxetable}{rr@{\\,{\\pm}\\,}rr@{\\,{\\pm}\\,}rr@{\\,\\,}} \\tablecaption{\\label{tab:orbit2} Revised Orbital Fits} \\tablewidth{0pt} \\tablehead{ \\colhead{Parameter} & \\mctc{No Precession} & \\mctc{With Precession} } \\startdata \\comment{ NO PRECESSION ERR PRECESSION ERR } \\math{e \\sin \\omega\\sb{0}\\tablenotemark{a} } & -0.063 & 0.014 & -0.065 & 0.015 \\\\ \\math{e \\cos \\omega\\sb{0}\\tablenotemark{a} } & 0.0011 & 0.00072 & -0.0036 & 0.0045 \\\\ \\math{e} & 0.063 & 0.014 & 0.065 & 0.015 \\\\ \\math{\\omega\\sb{0}} (\\degree) & -89.0 & 0.8 & -93 & 5 \\\\ \\math{\\dot{\\omega}} (\\degree d\\sp{-1}\\tablenotemark{a}) & 0 & 0 & 0.017 & 0.019 \\\\ \\math{P\\sb{s}} (days) \\tablenotemark{a} & 1.0914240 & 3\\math{\\times 10\\sp{-7}} & 1.0914315 & 7\\math{\\times 10\\sp{-6}} \\\\ \\math{P\\sb{a}} (days) & 1.0914240 & 3\\math{\\times 10\\sp{-7}} & 1.0914872 & 7\\math{\\times 10\\sp{-5}} \\\\ \\math{T\\sb{0}} (MJD)\\tablenotemark{a,b} & 508.97683 & 0.00012 & 508.97685 & 0.00012 \\\\ \\math{K} (ms\\sp{-1}\\tablenotemark{a}) & 225 & 4 & 224 & 4 \\\\ \\math{\\gamma} (ms\\sp{-1}\\tablenotemark{a}) & 19087 & 3 & 19088 & 3 \\\\ BIC & \\mctc{90.1} & \\mctc{92.8} \\enddata \\tablenotetext{a}{MCMC Jump Parameter.} \\tablenotetext{b}{MJD = JD - 2,454,000.} \\end{deluxetable}\\ataboff \\if\\submitms y \\clearpage \\fi \\placetable{tab:orbit2}" }, "1003/1003.0766_arXiv.txt": { "abstract": "We present the first mid-IR study of galaxy groups in the nearby Universe based on {\\em Spitzer} MIPS observations of a sample of nine redshift-selected groups from the XMM-IMACS (XI) project, at $z=0.06$. We find that on average the star-forming (SF) galaxy fraction in the groups is about 30\\% lower than the value in the field and 30\\% higher than in clusters. The SF fractions do not show any systematic dependence on group velocity dispersion, total stellar mass, or the presence of an X-ray emitting intragroup medium, but a weak anti-correlation is seen between SF fraction and projected galaxy density. However, even in the densest regions, the SF fraction in groups is still higher than that in cluster outskirts, suggesting that preprocessing of galaxies in group environments is not sufficient to explain the much lower SF fraction in clusters. The typical specific star formation rates (SFR/$M_\\ast$) of SF galaxies in groups are similar to those in the field across a wide range of stellar mass ($M_\\ast>10^{9.6}\\msun$), favoring a quickly acting mechanism that suppresses star formation to explain the overall smaller fraction of SF galaxies in groups. If galaxy-galaxy interactions are responsible, then the extremely low starburst galaxy fraction ($<1\\%$) implies a short timescale ($\\sim0.1$ Gyr) for any merger-induced starburst stage. Comparison to two rich clusters shows that clusters contain a population of massive SF galaxies with very low SFR (14\\% of all the galaxies with $M_\\ast>10^{10}\\,\\msun$), possibly as a consequence of ram pressure stripping being less efficient in removing gas from more massive galaxies. ", "introduction": "The local galaxy population presents a clear bimodality in many different properties: blue galaxies with active star formation and late-type morphologies vs red, quiescent galaxies with early-type morphologies \\citep[e.g.,][]{Strateva01,Baldry04,Balogh04}. This bimodality is ubiquitous, extending from galaxy clusters to groups and to the general field \\citep{Lewis02,Gomez03}. The physical origin of this bimodality remains one of the most puzzling questions of galaxy formation and evolution. Is the difference of the two distinct populations seeded in the early stages of galaxy formation (the so-called \\lq nature\\rq\\ scenario), or is it the end result of a transformation driven by environment (the \\lq nurture\\rq\\ scenario)? Strong evidence favoring the nurture scenario comes from the drastic change of the fraction of galaxies in these two populations in different environments: the fraction of passive galaxies increases with increasing galaxy density. However, such a correlation alone does not directly imply a nurture scenario. The fraction of galaxies in different populations also depends on galaxy stellar mass \\citep{Kauffmann03}, which could have a different distribution in low- and high-density regions seeded at the time of galaxy formation. Hence, to fully understand galaxy evolution, we need to disentangle the stellar-mass and environment dependence \\citep[e.g.,][]{Baldry06,Iovino09,Kovac09} and identify the mechanisms responsible for establishing the bimodality in galaxy properties. Galaxy groups, as the most common galaxy associations, contain about 50\\% of the galaxy population at the present day \\citep{Geller83, Tully88, Eke04, Eke05}. The characteristic depth of the potential wells of groups is similar to those of individual galaxies, and the velocities of galaxies within groups are only a few hundred \\kms, comparable to the internal velocity of galaxies. Under these circumstances, galaxies interact strongly with one another, and with the group as a whole. Such interactions could transform the morphology of galaxies, induce starbursts, and thereby turn active, late-type galaxies into quiescent, early-type galaxies. In addition, the group environment could also transform galaxies via ram pressure stripping of their hot gas halos, eventually suffocating star formation \\citep{Rasmussen06a,Kawata08,McCarthy08}. Hence, not only are groups the most common environmental phase experienced by galaxies during their evolution, they also have the potential to strongly affect large populations of galaxies and thereby help to explain the ubiquitous bimodality in galaxy properties. Another important implication of galaxy transformations in groups is the pre-processing of galaxies before they fall into clusters \\citep{Zabludoff98}. In a hierarchical structure formation scenario, clusters are built up by the accretion of smaller structures, e.g., isolated galaxies and groups. However, it is still under debate if the majority of cluster galaxies were ever located in groups before being acquired by clusters \\citep{Berrier09, McGee09}. If the fraction of cluster galaxies accreted in groups is substantial, preprocessing of galaxy properties by the group environment could play a major role in forming the predominantly passive population in clusters. Additional physical mechanisms which only work efficiently in the cluster environment, such as galaxy harassment and ram pressure stripping of cold galactic gas, could further affect galaxy properties but would possibly only be of secondary importance. Despite the importance of the group environment for the global galaxy population, our understanding of groups is still very limited. There have been many studies of group galaxies trying to address these questions. Due to the low galaxy density contrast against the field, many previous group studies focused on X-ray luminous groups, which are mostly virialized groups \\citep{Mulchaey98,Zabludoff98}. However, when referring to the majority of galaxies as being located in groups, this applies specifically to {\\it optically} selected groups, typically identified through a \\lq\\lq friends-of-friends\\rq\\rq \\ analysis of galaxy redshift survey data \\citep{Eke04,Balogh04}. Such groups span a wide range of evolutionary states, including systems in the process of collapsing (like the Local Group), systems in the throes of strong galaxy interactions or subgroup mergers, and fully virialized systems. To fully assess the importance of the group environment on galaxy evolution, complete samples of galaxy groups extending to poor systems are essential. There have also been studies of optically selected groups based on large-area galaxy surveys \\citep{Balogh04,Weinmann06}. However, due to the limited survey depth, many of the poor groups in those studies only contain a handful of known group members. To fully understand the dynamics and galaxy content of each group, we need to probe to significantly fainter optical magnitudes than is usually possible with such large surveys. For this purpose, we have started the XMM-IMACS (XI) Groups Project (\\citealt{Rasmussen06b}; hereafter Paper~I) to carry out a multi-wavelength study of a statistically representative nearby group sample. We selected 25 groups with velocity dispersion $\\sigma <500$ \\kms\\ from the group catalog of \\citet{Merchan02}, which was carefully derived from a friends-of-friends analysis of the 2dF redshift survey \\citep{Colless01}. The groups are selected in a narrow redshift range ($0.0610^{9.6}\\msun$. Similar to our result, \\citet{Tyler10} also found the specifc SFRs of group and field galaxies of intermediate redshift ($0.3$~1~Gyr). If it is the dominant process responsible for the deficit of SF galaxies, we would expect to see many SF galaxies with suppressed SFRs. This is inconsistent with the similar typical specific SFR we found in our group and field samples across a large stellar mass range, suggesting that strangulation is probably not the dominant mechanism for galaxies with $M_\\ast>10^{9.6}\\msun$ in the poor groups studied here. In X-ray luminous groups, however, this mechanism may become more important. \\subsection{Comparing Groups with Clusters} In hierarchical structure formation, galaxy clusters are assembled from lower mass halos and the galaxies in clusters might have been residing in group environments before they finally fell into clusters. This makes ``preprocessing'' in groups potentially important for cluster galaxies \\citep{Zabludoff98}. However, the importance of preprocessing depends on the accretion history of clusters. Using CDM $N$-body simulations, \\citet{Berrier09} claimed that the majority of cluster galaxies (70\\%) have never resided in a group environment before they fell into the cluster, and therefore that preprocessing in group environments could not play a significant role in explaining the difference between cluster and field galaxies. This result seems to be at odds with the observation that about half of the galaxy population resides in group environments in the nearby Universe \\citep{Eke04,Eke05,McGee09}, because it would mean that clusters preferentially accrete isolated galaxies rather than galaxies in groups. The apparent discrepancy, again, is likely related to differences in the adopted group definition. \\citet{Berrier09} define group members as the halos within the virial overdensity boundary of the hosting dark matter halos. This definition results in a much smaller fraction ($\\sim15\\%$) of galaxies residing in group environments compared to those found by the friends-of-friends algorithms typically used to identify groups in observational samples. Given the small fraction of galaxies residing in groups according to the former definition, it is no surprise that the majority of the galaxies falling into clusters have never been preprocessed within a group environment by this definition. On the other hand, if we follow the much more relaxed group definition as adopted in friends-of-friends algorithms, we would expect a much higher fraction of cluster galaxies to have been part of the group environment prior to falling into clusters. However, even if all the cluster galaxies have been in groups, the lower fraction of SF galaxies in clusters in a large stellar mass range compared to the fraction found in the XI groups suggests that preprocessing in groups is not sufficient to explain the deficiency of SF galaxies in clusters. This is reinforced by the fact that even in the highest density regions of our groups, the SF fraction is still higher than the fraction found in the outskirts of clusters. To explain the low SF fractions in clusters, either an environmental mechanism that works in group environments must continue to work in clusters, or some other cluster-specific environmental mechanism must be invoked. In either case, further processing of SF galaxies within the cluster environment is required. Not only are the SF fractions of cluster galaxies on average smaller than in groups, but the SF galaxies in clusters have smaller specific SFRs at $M_\\ast>10^{10}\\msun$. Such differences strongly support further SF suppression in clusters. \\citet{Bai06,Bai09} found that the IR luminosity function of nearby rich clusters has the same shape at the bright end ($L_{\\rm IR}>10^{43}$~\\ergss) as that of field galaxies. The galaxies that contribute to the bright end of the IR luminosity function are galaxies with SFR\\,$>0.3 \\msunyr$, which constitute the upper envelope of the specific SFR vs stellar mass distribution of cluster galaxies (cf.\\ Figure~\\ref{f_ssfr}). The similarity of the bright-end shape of the IR luminosity function in different environments is corroborated by the similar upper envelope of the specific SFR vs stellar mass distribution found here. \\citet{Bai09} suggest this similarity points to a fast acting SF suppression mechanism in clusters, for example, ram pressure stripping, that produces few galaxies in transition. However, the comparison of the specific SFRs of cluster SF galaxies to those in the groups and field does reveal a population of massive SF galaxies ($M_\\ast>10^{10}\\msun$), with suppressed but not totally extinguished SF, predominantly seen in clusters. This difference does not necessarily exclude ram pressuring stripping as an important mechanism in clusters, however, because massive galaxies are less vulnerable to such processes and could still retain some of their gas following a stripping event. The residual gas in those massive galaxies could sustain low-level SF for a much longer time. We have presented the first mid-IR study of nearby groups with complete optical spectroscopy and X-ray data. The nine groups in our sample span a wide range in velocity dispersion (100--500~\\kms), X-ray properties, richness, and galaxy distribution. These groups are typical of the galaxy groups that make up more than half of the galaxy population in the nearby Universe, and they are likely covering a wide range of evolutionary states. We analyzed the SF properties of the group galaxies from their MIPS 24~\\micron\\ emission and tested for correlations with global group properties. We also compared the SF properties of the group galaxies with those of cluster and field galaxies at the same redshifts to investigate the environmental effect on galaxy evolution. Our major results are summarized as follows: 1) The projected galaxy distributions of the nine groups show large variations, from more concentrated and circular distributions in the most massive groups ($\\sigma\\sim500$~\\kms) to filamentary structures in the least massive ones ($\\sigma\\sim100$~\\kms). This variation, along with the lack of a dominant BGG in the group center and the generally low level of extended X-ray emission, suggests that some of these systems are not yet virialized but still in the process of collapsing. 2) On average, the SF galaxy fraction (SFR\\,$>0.1\\msunyr$, $M_{R}<-20$) in our group sample is about 30\\% lower than in a comparison field sample extracted from the same data, and 30\\% higher than in our comparison cluster sample. The SF fraction of our groups does not show a strong systematic dependence on group global properties such as velocity dispersion, total stellar mass, or the presence of detectable diffuse X-ray emission. However, the two groups with the smallest velocity dispersion ($\\sigma\\approx100$~\\kms) do show the highest SF fractions, at a level comparable to that of the field. These conclusions remain unchanged if only considering the ``healthy'' SF galaxy fractions (SFR/$M_\\ast>$0.2[SFR/$M_\\ast$]$_{\\rm field}$, $M_\\ast>10^{9.6}\\msun$). 3) There is no strong dependence of the SF fraction on the radial distance from the group center. The 24~\\micron\\ SF fraction in the groups is, at all radii, larger than the corresponding fraction in the outskirts ($\\sim R_{vir}$) of rich clusters at similar redshifts. There is a weak trend of SF fractions decreasing with increasing projected galaxy density, with the lowest density regions having an SF fraction comparable to the field population. Even in the highest density regions of groups, the SF fraction is still larger than the SF fraction in the outer regions of clusters. In addition, the healthy SF fractions of cluster galaxies across a large stellar mass range are all at least 20\\% lower than those in groups. These pieces of evidence strongly suggest that preprocessing of galaxies in group environments prior to infall into clusters is not a sufficient explanation for the lower fraction of SF galaxies in clusters and that further processing by the cluster environment is required. 4) The typical specific SFRs of SF galaxies in groups are very similar to that in the field across a wide mass range ($M_\\ast>10^{9.6}\\msun$), favoring a quickly acting mechanism that suppresses star formation to explain the overall smaller fraction of SF galaxies in groups. The healthy SF fractions in groups show a possible dip at $M_\\ast \\approx 10^{10.1}\\msun$, corresponding to the bump seen in the optical luminosity function of poor groups \\citep{Miles04}. This is consistent with the speculation that galaxy merging by dynamical friction preferentially depletes intermediate-luminosity galaxies, which become subject to rapid gas consumption during the interaction, resulting in a population of galaxies of high stellar mass and a relatively low fraction of healthy SF galaxies. If galaxy-galaxy interactions are responsible for the deficit of SF galaxies in groups, then our non-detection of a significant starburst population among current group members does indeed imply a short time-scale for any merger-induced starburst stage ($\\sim0.1$ Gyr). This agrees well with a supernova feedback model that assumes an isothermal state for the star-forming gas. 5) At $M_\\ast>10^{10}\\msun$, the SF galaxies in clusters show much lower typical specific SFRs than galaxies in groups and the field, due to a more significant population of massive galaxies with very low SFR (14\\% of all the cluster galaxies with $M_\\ast>10^{10}\\msun$). This could result from ram pressure stripping being less efficient in removing gas from more massive cluster galaxies, allowing such galaxies to sustain low-level star formation fueled by a residual gas reservoir." }, "1003/1003.4373_arXiv.txt": { "abstract": " ", "introduction": "Blazars include BL Lacertae (BL Lac) objects and flat spectrum radio quasars (FSRQs). They are the most extreme class of active galactic nuclei (AGNs) and exhibit strong variability at all wavelengths of the whole electron-magnetic (EM) spectrum, strong polarization from radio to optical wavelengths, and are usually core dominated radio structures. These extreme properties are generally interpreted by many authors as a consequence of non-thermal emission from a relativistic jet oriented close to the line of sight \\citep{bland79,urry95}. Blazars vary on the diverse time scales (Gupta et al., 2008 and references therein). Variability has been one of the most powerful tools in revealing the nature of blazars. Understanding variation behaviour is one of the major issues of AGNs studies. The variability lags between different energy bands provide very important constrains for the interpretation of the emission components. In recent years, correlations of the variability in different energy regions have been widely studied \\citep{raite01,raite08,dai06,marsh08,areva08,chatt08,villa09,bonni09}. The radio source S5 0716+714 is one of the brightest and best-studied BL Lacertae objects. It was discovered in 1979 \\citep{kuhr81}, and was classified as a BL Lac object because of its featureless spectrum and its strong optical polarization \\citep{bierm81}. By optical imaging of the underlying galaxy, its redshift of z = 0.31$\\pm$0.08 was derived recently \\citep{nilss08}. It is a highly variable BL Lac object in the whole EM spectrum on diverse time scales \\citep{heidt96,ghise97,villa00,raite03,pian05,bach05,bach06, nesci05,ostor06,monta06,fosch06,wu07,wagne96,stali06}. The correlations and time lags between different energy bands in this source have been studied by some authors. A correlation was claimed for a selected range of data at 6 and 3 cm wavelengths, and optical wavelengths by \\cite{quirr91}. But 2 cm data from the same epoch published later by \\cite{quirr00} does not show evidence for correlated variability, casting doubt on the reliability of the claimed correlation between radio and optical bands \\citep{bigna03}. A good correlation was noticed between the light curves in the different passbands \\citep{sagar99}. An upper limit to the possible delay between $B$ and $I$ variations (10 minutes) was determined by \\cite{villa00}. With the data taken on January 8, 1995, \\cite{qian00} found an upper limit of the time lag of 0.0041 days between variations in the $V$ and $I$ bands. \\cite{raite03} found that the variations between different radio bands are very correlated, and the flux variations at the lower radio frequencies are delayed with respect to those at the higher frequencies. They also found weak correlations between the optical and radio emissions. \\cite{stali06} analyzed two night data, found that $V$ and $R$ are correlated with a small time lag (about 6 and 13 minutes, respectively, for April 7 and 14, 1996) and the variation at $V$ leading that at $R$ on both nights. \\cite{chen08} studied the gamma-optical correlation. Their result suggested a possible delay in the gamma-ray flux variations with respect to optical ($R$ band) variations of the order of 1 day. \\cite{fuhrm08} confirmed the existence of a significant correlation across all their observed radio-bands. The time delay between the two most separated bands ($\\lambda _{3mm}$ and $\\lambda _{60mm}$) is about 2.5 days. It must be noted that the time delay of 2.5 days is determined from an almost monotonic increase in flux density observed over a time range of 11 days. The determination of time lag can be used to study the geometry, kinematics and physical conditions in the inner regions of AGNs. This is attributed to light travel-time effect \\citep{villa00}. It is very important to search for detectable delays between optical bands themselves, even if these can be expected to be very small, if any. Our goals in this paper are to investigate the correlations between different optical light curves. The paper is arranged as follows: in Sect. 2, the light curves are presented; then in Sect. 3, the correlations and the time lags between different optical passbands are presented; after this, the discussion and conclusions are given in Sect. 4. ", "conclusions": "It is well known that the light curves between variant wave bands may have a short time lag. According to the inhomogeneous jet models, time delays are expected between the emission in different energy bands, as plasma disturbances propagate downstream \\citep{georg98}. Multi-wavelength monitoring of blazars shows that flares usually begin at high frequencies and then propagate to lower frequencies, implying that high-frequency synchrotron emission arises closer to the core than low-frequency synchrotron emission does \\citep{ulric97,marsc01}. High energy electrons emit synchrotron radiation at high frequencies and then cool, emitting at progressively lower frequencies and resulting in time lags between high and low frequencies \\citep{bai05}. The small time lags in optical regimes may be result of very small frequency intervals, and may indicate that the photons in these wavelengths should be produced by the same physical process. The above results have been obtained with a light-curve time resolution of the same magnitude as the possible lag. To check the method and the result mentioned above, we calculate the $B$-$B$ band autocorrelation DCF (Fig. \\ref{fig:dcf3}). The result of Monte Carlo simulations suggests that the time delay is $-0.0^{-16.3}_{+15.5}$ minutes. For the $V$-$V$, $R$-$R$ and $I$-$I$ autocorrelations, $\\tau$s are $0.0^{-11.4}_{+11.6}$, $0.1^{-11.5}_{+12.0}$ and $-0.0^{-11.5}_{+11.8}$ minutes(see Fig. \\ref{fig:dcf3}). They are all consistent with the expectation of zero lags. Furthermore, we shift $BVRI$ band light curves by a time of $t_{shift}$ and generate the corresponding artificial light curves of $B_{shift}$, $V_{shift}$, $R_{shift}$ and $I_{shift}$. Using the same process, we search for the time lag between the light curve and the shifted light curve. The results are listed in Table \\ref{tab:test}. The first column is for the shifted time , the second to the last are for the lags of $B$-$B_{shift}$, $V$-$V_{shift}$, $R$-$R_{shift}$ and $I$-$I_{shift}$ determined by the method introduced above. From the results listed in Table \\ref{tab:test}, one can see that the lags between two correlated light curves can be well determined. In addition to shifting the light curves to calculate the displaced Auto-DCF peaks, we shift the light curves with respect to those at each other wavelength by the time lags determined. The resulting DCFs are more consistent with zero lag. So, the time lags between different optical $BVRI$ band light curves are more reliable. The time lag of 10.2 minutes between variations of $B$ and $I$ band is consistent with the result found by \\cite{villa00}. They derived an upper limit of 10 minutes between $B$ and $I$ variations. The 5.8-minute lag between the band $V$ and $I$ is in good agreement with the 0.0041-day lag determined by \\cite{qian00}. For the $V$ and $I$ band variations, \\cite{stali06} reported time lags of about 6 and 13 minutes on two individual nights. Between the $V^{'}$ and $R^{'}$ band variations, \\cite{wu07} obtained the lags of 2.34 $\\pm$ 5.25 minutes on JD 2,453,737 and -0.01 $\\pm$ 1.88 minutes on JD 2,453,742. Because the time lags are very short, shorter than their typical sampling interval, they concluded that they did not detect an apparent time lag. In this analysis, it also must be pointed out that the time lags between the different optical bands are of the same orders of the light-curve time resolution, and the uncertainties are larger than the lags. So, the lags should be treated with caution, and much more significant results can be achieved only with much denser monitoring data. Two ways were suggested by \\cite{villa00} to improve the time resolution of the light curves, one way is to enlarge the telescope, another is to synchronized use two or more telescopes. In summary, the variations between different optical bands have been analyzed and the time lags and their uncertainties are convincingly determined. The results suggest that the variations are correlated very strongly. Considering the errors, every one of the values agrees with zero-lag well. While looking at the whole set of lags, it is possible to say that there is a hint of the variations of high frequency bands leading those of low frequency bands in the order of a few minutes." }, "1003/1003.1283_arXiv.txt": { "abstract": "We study the primordial perturbations generated during a stage of single-field inflation in Einstein-aether theories. Quantum fluctuations of the inflaton and aether fields seed long wavelength adiabatic and isocurvature scalar perturbations, as well as transverse vector perturbations. Geometrically, the isocurvature mode is the potential for the velocity field of the aether with respect to matter. For a certain range of parameters, this mode may lead to a sizable random velocity of the aether within the observable universe. The adiabatic mode corresponds to curvature perturbations of co-moving slices (where matter is at rest). In contrast with the standard case, it has a non-vanishing anisotropic stress on large scales. Scalar and vector perturbations may leave significant imprints on the cosmic microwave background. We calculate their primordial spectra, analyze their contributions to the temperature anisotropies, and formulate some of the phenomenological constraints that follow from observations. These may be used to further tighten the existing limits on the parameters for this class of theories. The results for the scalar sector also apply to the extension of Ho\\v{r}ava gravity recently proposed by Blas, Pujol\\`as and Sibiryakov. ", "introduction": "The enigmas of dark matter and cosmic acceleration have motivated the exploration of theories where gravity is ``modified\" at large distances. On the other hand, the range of possibilities for constructing such theories is severely limited by the requirement of general covariance, and for that reason, most of the proposed alternatives to General Relativity (GR) can in fact be cast as GR coupled to new fields.\\footnote{ A counterexample is the DGP brane-world scenario, where gravity is modified in the infrared by a continuum of Kaluza-Klein gravitons \\cite{DGP}. Because of the continuum, DGP cannot be formulated as a standard four dimensional GR with additional fields. See also \\cite{galileon} and \\cite{galileoncov}, for recent related proposals in the four dimensional context.} Cosmic acceleration may be due to a scalar field slowly rolling down a potential \\cite{Peebles:1987ek,mod}, or simply sitting in one of its local minima \\cite{bopo}. Alternatively, it can be driven by the non-minimal kinetic term of a k-essence scalar field with a Lagrangian of the form $p(X,\\phi)$, where $X=\\partial_\\mu\\phi\\partial^\\mu\\phi$. This form is quite versatile, and can be used to mimic cosmic fluids with a wide range of possibilities for the effective equation of state and speed of sound, including those which are characteristic of dark energy and cold dark matter \\cite{kessence}. The gradient of the k-essence field, $\\partial_{\\mu} \\phi$, is a time-like vector which spontaneously breaks Lorentz invariance, in a way that is parametrically independent of its effects on the time evolution of the background geometry. In particular, Lorentz invariance can be spontaneously broken by $\\partial_\\mu \\phi$ while the background spacetime remains maximally symmetric, a situation which is known as ghost condensation \\cite{ghostcondensation}. Still, the ``fluid\" responds to the gravitational pull of ordinary matter, leading to modifications of the long range potentials. More generally, theories with a massive graviton can be written in a covariant form as GR coupled to a set of ``St\\\"uckelberg\" scalar fields $\\phi^A$ with non-minimal kinetic terms, whose gradients have non-vanishing expectation values \\cite{massive,phases}. Depending on the interactions and the expectation values of the condensates, this can describe different phases of massive gravity. Aside from the Lorentz preserving Fierz-Pauli phase \\cite{massive} (see also \\cite{slavaali}), Lorentz breaking phases have been investigated in \\cite{phases}. Some of these have interesting phenomenology, such as the absence of ghosts in the linearized spectrum, a massive graviton with just two transverse polarizations, and weak gravitational potentials which differ from those in standard GR by terms proportional to the square of the graviton mass \\cite{lbmass,phases,pheno}. Additional fields of spin 2 have been considered in bi-gravity (or multi-gravity) theories \\cite{Damour:2002ws}, where space-time is endowed with several metrics interacting with each other non-derivatively. Due to general covariance, only one of the gravitons in the linearized spectrum stays massless, while the remaining ones acquire masses proportional to the non-derivative interaction terms. Lorentz invariance can be broken spontaneously even in cases where all metrics are flat, provided that their light-cones have different limiting speeds. This leads to phenomenology \\cite{lbbig} similar to that of certain phases of Lorentz breaking massive gravity referred to above \\cite{lbmass,phases,pheno}, of which multigravity can be seen as a particular realization. Finally, additional vector fields have received considerable attention in cosmology. Effective field theories for vectors are strongly constrained by stability requirements. Typically, those with non-trivial cosmological dynamics contain a massive ghost \\cite{ArmendarizPicon:2009ai}, which can be removed from the spectrum by sending its mass to infinity. This amounts to imposing a fixed-norm constraint on the vector, which in turn forces a Lorentz-breaking vacuum expectation value. This led Jacobson and Mattingly to dub this type of models Einstein-aether theories \\cite{Jacobson:2000xp,mond}. Their low-energy excitations are the Goldstone bosons of the broken Lorentz symmetry,\\footnote{In theories with spontaneously broken spacetime symmetries, the number of Goldstone bosons does not generally agree with the number of broken generators. However, if the order parameter that breaks the spacetime symmetry is spacetime-independent (as the constant aether field), then both numbers do agree \\cite{Low:2001bw}.} which will participate in the dynamics of the long range gravitational interactions. An interesting recent development is the proposal by Ho\\v{r}ava \\cite{hg} that a Lorentz-breaking theory of gravity may be renormalizable and UV complete. The breaking of Lorentz invariance in this case is implemented by introducing a preferred foliation of space-time, but no additional structure. As pointed out in \\cite{bps1}, any theory with a preferred foliation can be written in a generally covariant form by treating the time parameter which labels the different hypersurfaces as a St\\\"uckelberg scalar field ${\\cal T}$. The foliation is considered to be physical, but not the parametrization, and therefore the covariant theory should be invariant under field redefinitions ${\\cal T}\\to f({\\cal T})$. In other words, the Lagrangian can have a dependence on the unit normal to the hypersurfaces, but not on the magnitude of the gradient ${\\cal T},_\\mu$ (in contrast with the examples of k-essence and ghost condensation mentioned above). From this observation, Blas, Pujol\\`as and Sibiryakov showed \\cite{bps2} that Ho\\v{r}ava gravity could be extended by including in the action all terms compatible with reparametrization symmetry, and consistent with power counting renormalizability. Interestingly enough, this extension also cured certain problems in the scalar sector of the original proposal (such as instabilities and strong coupling at low energies \\cite{bps1}). Jacobson \\cite{jbpsh}, has recently clarified the relation between the Einstein-aether theory and this extended version of Ho\\v{r}ava gravity, which he dubbed BPSH gravity. In particular, he pointed out that any solution of Einstein-aether where the vector field is hypersurface orthogonal is also a solution of the low energy limit of BPSH gravity. Since the aether only interacts gravitationally, any signal of it must be proportional to a power of $(E/M_P)^2$, where $M_P$ is the reduced Planck mass, and $E$ is an energy scale. Thus, even though the aether contains massless fields, its presence is hard to detect. In that respect, inflation provides an interesting window to probe the aether and its implications. During inflation, short-scale vacuum fluctuations of light fields are transferred to cosmological distances, where they may leave an observable imprint. It is thus natural to look for signatures of Einstein-aether on the spectrum of primordial perturbations, which is the subject to which we devote this article. Previous work on this subject has been done in Refs. \\cite{Lim:2004js,Li:2007vz}, although in a somewhat narrower region of parameter space and with somewhat different conclusions. In the scalar sector, we find that there is a primordial isocurvature mode, which can be interpreted as the velocity potential for the aether with respect to matter. Depending on the aether parameters, this mode can grow on superhorizon scales, leading to a large random velocity field for the aether. Similar results apply to the transverse vector sector. These perturbations may thus be of phenomenological interest. We also find that the isocurvature mode is strongly correlated with the usual adiabatic mode, which corresponds to curvature perturbations in the co-moving slicing. For previous work on the impact of adiabatic scalar perturbations on the cosmic microwave background radiation (CMB) and large scale structure in (generalized) aether theories, see \\cite{Zuntz:2008zz, Zuntz:2010jp}. While this paper was being prepared, an interesting related paper by Kobayashi, Urakawa and Yamaguchi appeared \\cite{yuko}, which analyzes the post-inflationary evolution of the adiabatic scalar mode in BPSH theory. Where we overlap, our conclusions agree. The plan of the paper is the following. In Section \\ref{sec:Theory} we review the basics of Einstein-aether theory and the homogeneous cosmological solutions. Sections \\ref{sec:tensor perturbations}, \\ref{sec:scalar perturbations} and \\ref{sec:vector perturbations} are devoted to the analysis of tensor, scalar and transverse vector perturbations respectively. Section \\ref{ncmbs} analyzes the contribution of vector modes to the CMB spectrum. Readers familiar with the details Einstein-aether or BPSH theory are encouraged to jump directly to the concluding Section \\ref{sumcon}, for a self contained summary of the main results. Appendix A summarizes the existing bounds on the parameters of Einstein-aether theories. Appendix B discusses the equations of motion for the scalar sector of the theory in the longitudinal gauge. Appendix C deals with the canonical reduction of the scalar sector to the two physical degrees of freedom (a necessary step for the proper normalization of the vacuum fluctuations). Appendix D contains a derivation of the long wavelength adiabatic and isocurvature scalar modes, for generic matter content and expansion history. Appendix E derives the CMB temperature anisotropies due to vector modes. ", "conclusions": "\\label{sumcon} In this article we have studied cosmological perturbations in Einstein-aether theory, where the scalar and transverse vector sectors of general relativity are enlarged by an additional dynamical field each. We find that inflation can induce sizable perturbations in both of these new massless fields on observable scales. Our analysis also applies to the low energy limit of BPSH gravity, where the transverse vector is missing by construction \\cite{jbpsh}. For the purposes of summarizing our results, we shall assume that the aether parameters $c_i$\\, ($i=1,\\ldots,4$) are small. This is natural, since they can be thought of as proportional to the square of the ratio of the symmetry breaking scale $M$ to the Planck scale $c_i \\sim (M/M_P)^2\\ll 1$. To motivate the choice of the range of parameters we shall use below, let us recall that in the Einstein-aether theory, the effective gravitational constant on small scales $G_N$ can be different from the effective gravitational constant which appears in the Friedmann equation $G_\\mathrm{cos}$. We shall call $\\alpha$ and $c_{14}$ the parameters which relate these two constants to the bare Newton's constant $G$. They are given in Eqs. (\\ref{abbreviations}) as linear combinations of the standard $c_i$. In terms of $\\alpha$ and $c_{14}$ the effective gravitational constants are given by \\begin{equation} G = \\left(1-{\\alpha\\over 2}\\right) G_\\mathrm{cos} = \\left(1+{c_{14}\\over 2}\\right) G_N. \\end{equation} Note that for $\\alpha+c_{14}=0$ we have $G_\\mathrm{cos} = G_{N}$. The difference $G_\\mathrm{cos}-G_{N}$ is constrained by nucleosynthesis to be less than 10 \\%, so it seems natural to consider the range \\begin{equation}\\label{deltarange} |\\tilde{\\kappa}| \\ll 1, \\quad {\\rm where}\\quad \\tilde{\\kappa} \\equiv -\\left(1+{\\alpha \\over c_{14}}\\right). \\end{equation} This range guarantees the similarity of $G_\\mathrm{cos}$ and $G_{N}$, but it is typically more restrictive than required by the nucleosynthesis bound, since $c_{14}\\sim (M/M_P)^2$ is naturally small. If the parameters are such that we are outside of the range (\\ref{deltarange}), the effects we are investigating would be either too small to be of phenomenological interest, or too large to be compatible with observations. The main results of the paper are the following. First, we find that in the scalar sector, aside from the standard adiabatic mode $\\zeta$ (which corresponds to the curvature of surfaces of constant matter density), there is an additional isocurvature mode which can be important for phenomenology. Geometrically, the isocurvature mode can be described as the differential e-folding number $\\delta N$ which separates the surfaces of constant matter density from the surfaces orthogonal to the aether. This plays the role of a velocity potential $v$ for the aether with respect to matter. At the time of horizon exit during inflation, the amplitudes of $\\delta N$ and $v$ are comparable to that of the standard adiabatic mode $\\zeta$: \\begin{equation} v \\sim \\delta N \\sim \\zeta \\sim {H \\over M_P} \\epsilon^{-1/2} \\quad ({\\rm horizon\\ exit}). \\end{equation} Here $H$ is the Hubble rate and $\\epsilon \\ll 1$ is the slow roll parameter during inflation, which is independent of the aether parameters. After horizon crossing, the curvature perturbation $\\zeta$ stays constant, while the behaviour of $\\delta N$ depends on the parameter $\\tilde{\\kappa}$ defined above. For $\\tilde{\\kappa} < 0$, the isocurvature perturbation slowly decays on large scales, while for $\\tilde{\\kappa} > 0$ it grows. On the other hand, the velocity perturbation is given by $v \\sim (k/\\dot a) \\delta N$, where $k$ is the co-moving wave number and $\\dot a$ is the derivative of the scale factor with respect to proper time. Hence, during inflation, when $\\dot a$ grows, the long wavelength velocity field decays, roughly in proportion to the inverse of the scale factor. After inflaton, the universe decelerates and the velocity field grows again. At the time of horizon reentry, on cosmologically relevant scales, we have \\begin{equation} v \\sim \\delta N \\sim e^{N\\tilde{\\kappa}} \\zeta \\sim e^{N\\tilde{\\kappa}}\\ 10^{-5} \\lesssim 1,\\quad ({\\rm horizon\\ reentry}) \\end{equation} where $N\\sim 60$ is the number of e-foldings of inflation since the time when the cosmological scale first crossed the horizon. The last inequality indicates the limit of validity of the linear approximation. Note that for $\\tilde{\\kappa}=0$, the isocurvature perturbation and the velocity field of the aether are comparable to $\\zeta \\sim 10^{-5}$ at horizon reentry. However, with $\\tilde{\\kappa} \\lesssim 10/N$, we can have $\\delta N \\lesssim 1$. If $\\tilde{\\kappa}$ is large enough to saturate the inequality, this still allows for mildly relativistic speeds for the aether field $v\\sim 1$ within the observable universe. Similar results hold for the vector sector. Denoting by $V$ the transverse component of the aether velocity field with respect to matter, we find that on superhorizon scales $$ V \\sim \\left({\\epsilon \\over c_{14}}\\right)^{1/2} v. $$ Hence, if $c_{14} < \\epsilon$ (which seems quite natural if the scale of Lorentz symmetry breaking is low), the vector contribution to the velocity field will be dominant with respect to that of the longitudinal component. On the other hand, in a theory such as BPSH, the transverse component $V$ is missing, and the scalar part $v$ is the dominant one. We also find that the longitudinal gauge gravitational potentials $\\phi$ and $\\psi$ can be different even for the adiabatic mode. On superhorizon scales, we find that this effect (which can be attributed to anisotropic stress of the aether energy momentum tensor) is of order: \\begin{equation} (\\phi-\\psi)_\\mathrm{adiab} \\sim \\phi_\\mathrm{adiab}\\ c_{13} \\sim \\zeta\\ c_{13} \\sim 10^{-5}c_{13}. \\end{equation} where $c_{13} \\sim (M/M_P)^2$ is another combination of the aether parameters $c_i$, given in Eqs. (\\ref{abbreviations}). Physically, this parameter can be expressed in terms of the propagation speed of tensor modes $c_{13}=c_t^{-2}-1$. The isocurvature mode contributes maximally to the anisotropic stress, but the potential due to the isocurvature mode is suppressed by $c_{13}$: \\begin{equation} (\\phi-\\psi)_\\mathrm{isoc} \\sim \\phi_\\mathrm{isoc} \\sim c_{13}\\ \\delta N. \\end{equation} Since $\\delta N$ can be larger than $\\zeta$, the anisotropic stress can be dominated by the isocurvature mode. The anisotropic stress on observable scales is suppressed from its value at horizon crossing, due to the dynamics of the aether on subhorizon scales. For $\\tilde{\\kappa}=0$, the effect scales like $k^{-2}$ for modes that crossed the horizon during the matter era. For modes that crossed the horizon during the radiation era, the behaviour changes to $k^{-1}$. Current constraints on $\\phi-\\psi$ on cosmological scales are not very restrictive, and $|c_{13}|\\lesssim 1$ seems to be allowed by observations. The aether manifests itself in PPN parameters through frame dependent effects, which cause anisotropies in the gravitational field of bodies which move with respect to the aether. In this way, the velocity field generated during inflation might be detectable. It should be noted, however, that it seems difficult with present technology to observe the statistical properties of the random field from this particular type of observations. Even if the velocity field were relativistic on cosmological scales $v\\sim 1$, it falls with scale as $k^{-2}$. In particular, the component which varies on scales of the order of 100 Mpc would then be below the virial velocity $v_\\mathrm{vir} \\sim 10^{-3}$ of objects bound in galaxies, and it seems unlikely that we can directly sample frame dependent effects in objects which are located at distances larger than that. On the other hand, at the relatively small distances where the observation of frame dependent effects is accessible, we may still detect a large but fairly homogeneous velocity field, even one that is much larger than the virial velocity of bound objects. Finally, we have computed the contribution of transverse vector fields $V$ to the angular power spectrum of CMB anisotropies. We find that for $\\tilde{\\kappa}=0$, the spectrum of multipole coefficients $C^{V}_\\ell$ has the same shape as that of tensor modes. The amplitude, on the other hand, is related to the spectrum $C^{h}_\\ell$ for tensor modes and $C^{\\zeta}_\\ell$ for the adiabatic scalar mode as \\begin{equation}\\label{comparativa} C^{V}_\\ell \\sim\\ {c_{13}^2 \\over c_{14}}e^{2 N \\tilde{\\kappa}} C^{h}_\\ell\\ \\sim\\ {\\epsilon\\ c_{13}^2\\over c_{14}}e^{2 N \\tilde{\\kappa}} C^{\\zeta}_\\ell. \\end{equation} This means that the vector modes in Einstein-aether theory can easily dominate the signal from tensor modes. The analysis of polarization induced by the vector modes is therefore of phenomenological interest, and is left for further research. Moreover, we know that the CMB is well-fit with a primordial spectrum of \\emph{scalar} adiabatic perturbations. This imposes additional phenomenological restriction amongst the parameters $c_{13}$ and $\\tilde{\\kappa}$ of Einstein-aether theories, of the form \\begin{equation} \\tilde{\\kappa} \\lesssim {1\\over 2 N}\\ln\\left|{c_{14}\\over\\epsilon\\ c_{13}^2}\\right|. \\end{equation} So far, we have not included the constraints which follow from the frame-dependent effects on the PPN parameters. These are summarized in Appendix \\ref{sec:other constraints}, and take the form \\begin{equation}\\label{ppnconst} \\omega \\, \\alpha_1\\lesssim 10^{-7},\\quad \\omega^2\\alpha_2\\lesssim 10^{-13}. \\end{equation} Here, $\\omega= \\max \\{V,v,v_\\mathrm{vir}\\}$, is the velocity of the aether with respect to the object whose gravitational field is being tested at post-Newtonian order, and $v_\\mathrm{vir} \\sim 10^{-3}$ is the typical virial velocity for bound objects with respect to the CMB frame. The post-Newtonian parameters $\\alpha_1$ and $\\alpha_2$ are combinations of the four aether parameters $(\\alpha,c_{14},c_+, c_-)$. Here, following \\cite{Jacobson:2008aj}, we have introduced $c_+\\equiv c_{13}=c_1+c_3$ and $c_- \\equiv c_1-c_3$. Phenomenologically, it is possible to set $\\alpha_1=\\alpha_2=0$, which determines $\\alpha$ and $c_{14}$ as functions of the other two parameters in the model, \\begin{equation} \\alpha=-c_{14}= - 2 {c_+ c_- \\over (c_+ +c_-)}. \\label{restriction} \\end{equation} The parameters $c_+$ and $c_-$ remain rather unconstrained by observations. Stability requirements and superluminality (or Cherenkov) constraints are satisfied provided that $-1\\leq c_{+} \\leq 0$, $c_{+} /3(1+c_{+}) \\leq c_{-} \\leq 0$. Constraints from radiation damping in binary systems determine further constraints on the $(c_+,c_-)$ plane, but a sizable coefficient \\begin{equation} |c_{13}|\\lesssim 1 \\end{equation} still seems to be allowed by all observations \\cite{Jacobson:2008aj}. This is important, since the gravitational effects of the aether are suppressed by this coefficient. For instance the contribution of vectors to the CMB anisotropies is of order \\begin{equation} C_\\ell^V \\sim c_{13}^2 V^2, \\end{equation} where $V\\lesssim 1$ is the aether velocity field. Hence, the observability of the effect depends crucially on $c_{13}$ being sufficiently large. This brings us to the question of fine tuning amongst the parameters of the model. In a low energy theory, one might have expected all dimensionless parameters to be of the same order, $$ c_i \\sim (M/M_P)^2. $$ Observability of $C_\\ell^V$ requires an inequality of the form $c_{13} V \\gtrsim 10^{-6}$, which would be natural provided that \\begin{equation}\\label{osbs} (M/M_P)_\\mathrm{obs}^2 \\gtrsim 10^{-6} V^{-1}. \\end{equation} On the other hand, in Eq. (\\ref{restriction}) we have adjusted the parameters so that $\\alpha_1=\\alpha_2=0$, but the actual restriction (\\ref{ppnconst}) is of the form $\\alpha_2 \\lesssim 10^{-13}\\omega^{-2}$. Hence, $\\alpha_2$ must be well below the natural scale (\\ref{osbs}) by a considerable suppression factor \\begin{equation} \\alpha_2 \\lesssim 10^{-7} \\omega^{-1} (V/\\omega) (M/M_P)_\\mathrm{obs}^2, \\end{equation} with $10^{-3}<\\omega<1$. In the classical theory, the parameter $\\alpha_2$ can always be chosen by hand to have any particular value. However, in an effective field theory (EFT) a parameter is considered to be finely tuned or technically unnatural if quantum corrections to it are larger than the desired renormalized value of the parameter. The question, therefore is whether the very small values of $\\alpha_2 \\ll (M/M_P)^2$ are stable or not under quantum corrections. Withers \\cite{Withers} has recently analyzed the Einstein-Aether theory as an EFT, with the conclusion that the parameters $c_i$ receive only negligible logarithmic corrections. A similar result may hold in BPSH theory \\cite{bps2}. This subject is left for further study. To conclude, the results presented here show that the preferred frame singled out by the aether field $A^{\\mu}$, or by the preferred foliation of the BPSH theory, may have picked up a large random velocity field seeded by quantum fluctuations during inflation. Depending on the parameters, this may even be mildly relativistic on cosmological scales. The effects of this velocity field may be detectable in observations of frame dependent PPN effects, or in specific features in the CMB spectrum such as a sizable contribution from vector modes. These issues deserve further investigation." }, "1003/1003.3415_arXiv.txt": { "abstract": "According to the no-hair theorem, an astrophysical black hole is uniquely described by only two quantities, the mass and the spin. In this series of papers, we investigate a framework for testing the no-hair theorem with observations of black holes in the electromagnetic spectrum. We formulate our approach in terms of a parametric spacetime which contains a quadrupole moment that is independent of both mass and spin. If the no-hair theorem is correct, then any deviation of the black-hole quadrupole moment from its Kerr value has to be zero. We analyze in detail the properties of this quasi-Kerr spacetime that are critical to interpreting observations of black holes and demonstrate their dependence on the spin and quadrupole moment. In particular, we show that the location of the innermost stable circular orbit and the gravitational lensing experienced by photons are affected significantly at even modest deviations of the quadrupole moment from the value predicted by the no-hair theorem. We argue that observations of black-hole images, of relativistically broadened iron lines, as well as of thermal X-ray spectra from accreting black holes will lead in the near future to an experimental test of the no-hair theorem. ", "introduction": "The no-hair theorem establishes the remarkable property of general-relativistic black holes that their spacetimes and hence all of their characteristics are uniquely determined by precisely three parameters: their mass, spin, and charge (Israel 1967, 1968; Carter 1971, 1973; Hawking 1972; Robinson 1975; Mazur 1982). The spacetimes of such black holes are described by the Kerr-Newman metric (Newman et al. 1965), which reduces to the Kerr metric (Kerr 1963) in the case of uncharged black holes. This theorem relies on the cosmic censorship conjecture (Penrose 1969) as well as on the physically reasonable assumption that the exterior metric is free of closed timelike loops. If these requirements are met, then all astrophysical black holes should be described by the Kerr metric. Black holes are commonly believed to be the final states of the evolution of sufficiently massive stars at the end of their lifecycle. The gravitational collapse of such stars leads to the formation of a black hole (Oppenheimer \\& Snyder 1939; Penrose 1965), and any residual signature of the progenitor other than its mass and spin are radiated away by gravitational radiation (Price 1972a, 1972b). This scenario provides an astrophysical mechanism with which Kerr black holes can be generated. Almost all nearby galaxies harbor dark objects of high mass and compactness at their center (Kormendy \\& Richstone 1995) including our own galaxy (Ghez et al. 2008; Gillessen et al. 2009) providing strong evidence that black holes are realized in nature. In addition, the measurement of the orbital parameters of many galactic binaries supports the claim that they contain stellar-mass black holes (e.g., McClintock \\& Remillard 2006). Despite the large amount of circumstantial evidence, there has been no direct proof, so far, for the existence of an actual event horizon. An event horizon, one of the most striking predictions of general relativity, is a virtual boundary that causally disconnects the interior of a black hole from the exterior universe. The presence of an event horizon in black-hole candidates has only been inferred indirectly, either from the properties of advection dominated accretion flows or from the lack of observations of Type I X-ray bursts (Narayan, Garcia, \\& McClintock 1997, 2001; Narayan \\& Heyl 2002; McClinock, Narayan, \\& Rybicki 2004; Broderick et al. 2009; see discussion in Psaltis 2006). The alternative hypothesis that these dark compact objects are not described by the Kerr metric but perhaps by a solution of the Einstein equations with a naked singularity (e.g., Manko \\& Novikov 1992) and, therefore, violate the no-hair theorem, is still possible within general relativity. Alternatively, these dark objects might be stable stellar configurations consisting of exotic fields (e.g., boson stars, Friedberg, Lee, \\& Pang 1987; gravastars, Mazur \\& Mottola 2001; black stars, Barcel\\'{o} et al. 2008). Finally, the fundamental theory of gravity may be different from general relativity in the strong-field regime, and the vacuum black-hole solution might not be described by the Kerr metric at all (e.g., Yunes \\& Pretorius 2009; Sopuerta \\& Yunes 2009; c.f. Psaltis et al. 2008). As a result, testing the no-hair theorem allows us not only to verify the identification of dark compact objects in the universe with Kerr black holes but to test the strong-field predictions of general relativity, as well. The exterior spacetime of a black-hole candidate is usually best described in terms of its multipole moments. The mass and the spin of a black hole can be identified as the first two such moments. As a consequence of the no-hair theorem, mass and spin already specify all higher multipole moments of a black-hole spacetime completely. Therefore, measuring three independent multipole moments of the spacetime around a black hole suffices to test the no-hair theorem (Ryan 1995). Recent advances in instrumentation have opened new horizon towards probing in detail the immediate vicinity of black holes. Very-long baseline interferometry (VLBI) has allowed us to observe the first images of the inner accretion flows of black holes and revealed evidence for sub-horizon scale structures (Doeleman et al. 2008). The International X-Ray Observatory (IXO) will measure the detailed profiles of relativistically broadened iron lines in AGN (see, e.g., Brenneman et al. 2009). Moreover, the Laser Interferometer Space Antenna (LISA) will probe supermassive black holes near their event horizons via extreme mass-ratio inspirals (EMRIs) (see, e.g., Hughes 2006). On the theoretical front, several approaches have been developed to date in order to extract information about black-hole spacetimes with gravitational-wave observations. This can be done by analyzing several observables that are parametrized as functions of the multipole moments (Ryan 1995). It requires a suitable choice of the metric that either allows for a determination of multipole moments of order $l\\geq2$ (Glampedakis \\& Babak 2006; Gair, Li, \\& Mandel 2008) or of metric perturbations (Collins \\& Hughes 2004; Vigeland \\& Hughes 2010). The principles of a general algorithm for stationary axisymmetric vacuum spacetimes that admit full integrability of the geodesic equations have been outlined by Brink (2008; 2009 and references therein). In this series of papers, we investigate a framework for testing the no-hair theorem with observations in the electromagnetic spectrum. We formulate our tests in terms of a quasi-Kerr metric (Glampedakis \\& Babak 2006) that incorporates an independent quadrupole moment but reduces smoothly to the familiar Kerr metric when the deviation at the quadrupole order is set to zero. The central object constitutes a quasi-Kerr black hole (similar to the term bumpy black hole coined by Collins \\& Hughes 2004): it deviates from the Kerr metric in (at least) one multipole moment and allows for a test that distinguishes a Kerr black hole from a different type of object. We compute observables that can be measured directly with either current or near-future instruments and allow for the extraction of at least three independent multipole moments. This approach allows for a two-fold test: If the central object is indeed a black hole, a deviation from the Kerr metric must be zero. If, however, the deviation is measured to be nonzero, then it either has to be a different type of object or general relativity itself breaks down in the strong-field regime very close to the black hole (c.f., Collins \\& Hughes 2004; Hughes 2006). In particular, our framework will enable us to perform a test of the no-hair theorem with the measurement of black-hole images, relativistically broadened iron lines emitted from accretion disks, and the measurement of the innermost stable circular orbit (ISCO) from the continuum disk spectra. All of these observables depend explicitly on the multipole moments of the spacetime. Under the assumption that the compact objects are Kerr black holes, these observables have been used to measure the spins of the black holes (e.g., Zhang, Cui, \\& Chen 1997; Brenneman \\& Reynolds 2006; Broderick et al. 2009). In the more general case, their properties and respective shapes likewise can be used to constrain the quadrupole moments. Similar approaches have been suggested based on timing observations of pulsar black hole binaries (Wex \\& Kopeikin 1999) and on observations of stellar orbits in the vicinity of Sgr A* (Will 2008; Merritt et al. 2009). This paper is structured as follows: In Section~2 we briefly summarize techniques to measure deviations from the Kerr metric and to extract multipole moments of a given black-hole spacetime. We develop our framework for a test of the no-hair theorem in Section~3 and analyze various relevant properties of the underlying quasi-Kerr metric in Section~4. ", "conclusions": "In this series of papers, we investigate a framework for testing the no-hair theorem with observations of black holes in the electromagnetic spectrum. Our framework can be viewed as either a null-hypothesis test of the no-hair theorem within general relativity or as a self-consistent test of general relativity itself in the strong-field regime. If the multipole moments of a black-hole candidate are measured to be different from the moments of the Kerr metric, there are two possibilities. If general relativity is assumed, the astrophysical object is not a black hole. But if it is known to possess a horizon, then both the no-hair theorem and general relativity are incorrect. If the measured multipole moments coincide with the Kerr multipole moments, general relativity and the no-hair theorem may be correct or not. A definite answer from the extraction of the multipole moments alone is not possible, because other theories of gravity likewise predict the Kerr metric as a black-hole spacetime (Psaltis et al. 2008). In this first paper, we formulated our tests based on the quasi-Kerr metric (Glampedakis \\& Babak 2006) which contains an independent quadrupole moment and deviates smoothly from the Kerr metric at the quadrupole order. Since the no-hair theorem admits exactly two independent multipole moments, a measurement of three moments allows us to test the no-hair theorem (Ryan 1995). General-relativistic black holes must have the multipole structure of the Kerr metric. If a different set of multipoles is detected, then the compact object cannot be a black hole within general relativity. We estimated the range of validity of the quasi-Kerr metric and demonstrated the dependence of various properties of this spacetime on both the spin and the quadrupole moment of the black hole. We analyzed in detail the effects of light bending, photon redshift, and the respective locations of the ISCO and the circular photon orbit, all of which are of critical importance for observables of accretion flows around astrophysical black holes. In particular, we showed that the radius of the ISCO and the amount of gravitational lensing experienced by photons are altered significantly for already moderate changes of the quadrupole moment. We identified several observational approaches within our framework that will test the no-hair theorem within the next few years. Among these are imaging observations of accretion flows around black holes and, especially, Sgr A* using VLBI techniques, as well as the precise measurements of the spectra of iron lines and accretion disks from AGN with IXO. In the following papers we will explore in detail the prospects of testing the no-hair theorem with such observations. We thank Avery Broderick, Abraham Loeb, Scott Hughes, Kostas Glampedakis, Emanuele Berti, Nico Yunes, Sarah Vigeland, and Daniel Marrone for helpful discussions. We also thank Sarah Vigeland and Scott Hughes for sharing their paper with us in advance of publication. DP thanks the ITC at the Harvard-Smithsonian Center for Astrophysics for their hospitality. This work was supported by the NSF CAREER award NSF 0746549." }, "1003/1003.2719_arXiv.txt": { "abstract": "We consider the problem of self-regulated heating and cooling in galaxy clusters and the implications for cluster magnetic fields and turbulence. Viscous heating of a weakly collisional magnetised plasma is regulated by the pressure anisotropy with respect to the local direction of the magnetic field. The intracluster medium is a high-beta plasma, where pressure anisotropies caused by the turbulent stresses and the consequent local changes in the magnetic field will trigger very fast microscale instabilities. We argue that the net effect of these instabilities will be to pin the pressure anisotropies at a marginal level, controlled by the plasma beta parameter. This gives rise to local heating rates that turn out to be comparable to the radiative cooling rates. Furthermore, we show that a balance between this heating and Bremsstrahlung cooling is thermally stable, unlike the often conjectured balance between cooling and thermal conduction. Given a sufficient (and probably self-regulating) supply of turbulent power, this provides a physical mechanism for mitigating cooling flows and preventing cluster core collapse. For observed density and temperature profiles, the assumed balance of viscous heating and radiative cooling allows us to predict magnetic-field strengths, turbulent velocities and turbulence scales as functions of distance from the centre. Specific predictions and comparisons with observations are given for several different clusters. Our predictions can be further tested by future observations of cluster magnetic fields and turbulent velocities. ", "introduction": "Early X-ray observations of galaxy clusters and the intracluster medium (ICM) indicated radiative losses large enough to lead to cooling flows \\citep[see reviews by][]{sarazin86,sarazin88}. Mass deposition rates $\\dot{M}_{\\rm CF}$ due to presumed cooling flows were estimated to be as much as $\\sim 10^3~{\\rm M}_\\odot~{\\rm yr}^{-1}$ in some clusters \\citep{cb77,fn77,mb78}. The cooling-flow model also predicted copious iron line emission from temperatures between $10^6$ and $10^7~{\\rm K}$. However, at the time there was little direct evidence for mass dropout in any spectral band other than X-rays \\citep[for a review, see][]{fabian94}. More recent high spectral resolution X-ray observations failed to detect the expected iron-line emission and constrained the central temperature to be $\\sim 1/3$ of the bulk cluster temperature \\citep{pf06}. The spectroscopically determined mass deposition rates are $\\lesssim 0.1~\\dot{M}_{\\rm CF}$ \\citep[e.g.][]{vf04}. This is despite the fact that the cooling time at $r\\lesssim 100~{\\rm kpc}$ is less than a Hubble time in $\\gtrsim 70\\%$ of clusters \\citep[e.g.][]{esf92,pfeajw98,spo06,vkfjmmv06}. This discrepancy is the so-called `cooling flow problem.' Some heating mechanism must therefore be balancing radiative cooling in the `cool-core' clusters. A wide variety of heating and heat transport schemes have been considered, including thermal energy from the outer regions of the cluster being transported to the central cooling gas by conduction \\citep[e.g.][]{bc81,tr83,bm86,vsfaj02,fvm02,zn03}; energy in jets, bubbles, cosmic rays, outflows and/or radiation from a central active galactic nucleus (AGN) either resulting in turbulent diffusion of heat \\citep[e.g.][]{dj96,kn03a} and/or being thermalised via dissipation of turbulent motions, sound waves and/or gravitational modes \\citep[e.g.][]{lf90,bt95,co01,rbb04a,rbb04b,vd05,frtd05,brh05,cr07,mbbb08}; dynamical friction from galaxy wakes \\citep[e.g.][]{schipper74,ly76,miller86,jdkbm90}; or some combination of these \\citep[e.g.][]{brueggen03,dc05,co08}. Much of recent work has focused on either thermal conduction, AGN heating, or both. Thermal conduction alone cannot be the solution to the cooling flow problem across the full range of masses due to its steep temperature dependence \\citep{vf04,kaastra04,ppkf06}. Even in hot systems where conduction is potent, fine tuning of the suppression factor $f$ (the fraction by which Spitzer conductivity is reduced due to, e.g., tangled magnetic field lines) is required \\citep{bd88}. Moreover, if the thermal conduction has the same temperature dependence as the Spitzer conductivity (i.e. if $f$ is a constant) for a given ICM atmosphere, the resulting equilibria are thermally unstable (e.g. \\citealt{bd88,soker03,kn03b}; see further discussion in Sections \\ref{sec:balance} and \\ref{sec:conduction}). More recently, however, there has been renewed interest in the possibility that thermal conduction may provide sufficient heating to stably counteract the effects of radiative cooling. This has gone hand-in-hand with a dramatic increase in our understanding of `dilute' (i.e. only weakly collisional) plasmas, due to the appreciation that even very weak magnetic fields introduce an anisotropy into heat fluxes. One important consequence is that the criterion for convective instability changes to one of temperature, rather than entropy, increasing downwards \\citep{balbus00,balbus01}. \\citet{quataert08} generalised Balbus's (2000) analysis and found that a heat-flux buoyancy-driven instability (HBI) occurs for upwardly-increasing temperature profiles as well, so long as the magnetic field is not entirely horizontal (orthogonal to gravity). \\citet{br08} conjectured that the nonlinear HBI is self-regulating and drives a reverse convective thermal flux, both of which may mediate the stabilisation of cooling cores. Numerical simulations of the HBI have been performed by \\citet{pq08}, \\citet{brbp09} and \\citet{pqs09} with applications to the ICM. It was discovered that the HBI acts to rapidly reorient field lines to insulate the core, undermining the role of thermal conduction and causing a cooling catastrophe to occur. Subsequent work has shown that a moderate amount of turbulent driving may help regulate the HBI and allow thermal channels to remain open, potentially stabilising the core against collapse \\citep{scqp09,pqs10,ro10}. There are several reasons to believe that AGNs play an important role in regulating cooling. In many clusters, the AGN energy output inferred from radio-emitting plasma outflows and cavities is similar to the X-ray cooling rate of the central gas \\citep{fabian00,ksh06,mn07,forman07}. Moreover, $\\gtrsim 70\\%$ of cool-core clusters harbour radio sources at their centres, while $\\lesssim 25\\%$ of clusters without cool cores are radio loud \\citep{burns90}, providing strong circumstantial evidence for a connection between the processes that fuel the radio emission (such as AGNs) and the X-ray emission from the cooling gas. Models of self-regulated heating from AGN have been constructed \\citep[e.g.][]{co01,rb02,bm03,kb03,hb04,go08,bs09} in which AGN activity is triggered by cooling-induced gas accretion toward cluster centres, increasing AGN heating and halting the collapse. Episodic outflows from AGN are thought not only to quench cooling and condensation in clusters, but also to limit the maximum luminosity of galaxies and regulate the growth of black holes at their centres \\citep{binney05}. While AGN activity is fundamentally linked with the observed presence of radio bubbles and/or X-ray cavities \\citep[e.g.][]{brmwn04,df06}, it is currently unclear how the AGN energy is actually thermalised (e.g. see the introduction of \\citealt{vd05} for a review of possibilities). This question can only be answered once knowledge of the effective viscosity of the ICM is acquired. The ICM hosts subsonic turbulence and magnetic fields with energy density comparable to that of the motions. Both of these should affect the viscosity of the ICM \\citep[see review by][]{sc06}. In particular, the presence of a magnetic field alters the form of the viscosity when the ratio of the ion cyclotron and collision frequencies is much greater than unity \\citep{braginskii65}, a condition amply satisfied in galaxy clusters. As a result, the transport properties of the ICM become strongly dependent on both the geometry and strength of the magnetic field, as well as on microscale plasma instabilities that are likely to occur ubiquitously in the ICM \\citep[e.g. firehose and mirror;][]{sckhs05,lyutikov07,sckrh08,scrr10,rsrc10}. In this paper, we investigate the effect these plasma effects might have on the large-scale transport properties of the ICM. Specifically, we argue that parallel viscous heating, due to the anisotropic damping of turbulent motions, is regulated by the saturation of microscale plasma instabilities (e.g., firehose and mirror) and can balance radiative cooling in the cool cores of galaxy clusters in such a way as to ensure thermal stability (Section \\ref{sec:coolheat}). Given observed densities and temperatures, this balance implies specific values for central magnetic-field strengths and radial profiles of the rms magnetic field that are in good agreement with current observational estimates and that lend themselves to testing by future observations (Section \\ref{sec:profiles}). We also show that, under the reasonable assumption that turbulent kinetic and magnetic energies are comparable to one another, cluster profiles for the turbulent velocity and the characteristic turbulence scale may be derived. The specific case of A1835 is considered as a typical example in Section \\ref{sec:a1835}. Since the fundamental plasma-physical processes we appeal to are in principle universal in both cool-core and non-cool-core (i.e. unrelaxed) clusters, we also calculate in Section \\ref{sec:coma} the magnetic-field strengths and turbulence characteristics for some non-cool-core clusters. They turn out to be in good agreement with current observational estimates. Thermal conduction and the robustness of our results with respect to it are briefly discussed in Section \\ref{sec:conduction}. In Section \\ref{sec:discussion}, we close with a discussion of our results and their limitations. ", "conclusions": "\\label{sec:discussion} In this paper, we have introduced a model for regulating cooling in cluster cores in which turbulence, magnetic fields and plasma physics all play crucial roles. Our findings are fundamentally based on an appreciation that, whatever the source of effective viscosity in the ICM, it is certainly not hydrodynamic. Instead, it is set by the microscale plasma-physical processes, which are inevitable in a weakly collisional, magnetised environment such as the ICM. Assuming that microscale plasma instabilities pin the pressure anisotropy at its marginally stable value, we derive an expression for the heating rate due to parallel viscous dissipation of turbulent motions. For typical conditions in a variety of cluster cores, this rate is of the right magnitude to balance radiative cooling. Moreover, this source of heating turns out to be thermally stable. Put simply, the viscous stress in the ICM is a dynamic quantity that responds to local changes in temperature, density and magnetic-field strength in such a way as to prevent runaway heating or cooling. A basic qualitative outline of how this occurs is given in Fig. \\ref{fig:outline}. If true, what we have conjectured constitutes a physical mechanism that allows clusters to develop stable non-isothermal temperature profiles and avoid a cooling catastrophe -- an outcome that has remained elusive in models involving balancing the cooling by thermal conduction. Of the five quantities -- density, temperature, rms magnetic-field strength, rms turbulent velocity, characteristic turbulence scale -- given radial profiles of any two, we can predict the rest. The most reliable, as well as readily observable, predictions that follow from our theory concern the (rms) magnetic-field strength. For typical electron densities and temperatures, we predict magnetic-field strengths in the range $\\sim 1$ -- $10~\\mu{\\rm G}$, well within current observational constraints. For the specific clusters discussed here (e.g. A1835, Hydra A, A2199, A2382, A2255), a balance between parallel viscous heating and radiative cooling results in field strengths and profiles that are quite reasonable and in good agreement with current observational estimates where available. It would be interesting to test our model further via analysis of Faraday rotation maps. Another prediction is the manner in which the magnetic-field strength $B$ depends on electron density $n_{\\rm e}$ and temperature $T$ (eq. \\ref{eqn:Bprofile}): \\begin{equation} B\\propto n^{1/2}_{\\rm e} T^{3/4} \\quad \\textrm{(for cool-core clusters).} \\end{equation} In cool-core clusters, this suggests that the oft-employed assumption of an exclusive relationship between $B$ and $n_{\\rm e}$ should be relaxed. It is encouraging that, in isothermal clusters, the implied scaling \\begin{equation}\\label{eqn:bscalingiso} B\\propto n^{1/2}_{\\rm e} \\quad \\textrm{(for isothermal clusters)} \\end{equation} has already been observationally inferred to hold in A2382 \\citep{gmgpgrcf08}, Coma \\citep{bfmggddt10} and A665 \\citep{vmgfgob10}. A caveat here is that a local heating-cooling balance in such clusters may take very long to establish and so should be treated with due scepticism. A test of whether this conjecture is reasonable would be to determine whether equation (\\ref{eqn:bscalingiso}) is satisfied systematically in a large sample of isothermal clusters. \\begin{figure} \\centering \\includegraphics[width=3.2in]{KSCBS10_fig4.eps} \\caption{Qualitative outline of the processes responsible for thermal balance between parallel viscous heating and radiative cooling. Energy sources (AGNs, mergers, etc.) inject energy into the ICM, some or all of which (depending on the effective impedance of the ICM) is absorbed by the plasma and converted into turbulence. This changes the magnetic-field strength, giving rise to both a fluctuation dynamo and pressure anisotropies. The fluctuation dynamo presumably saturates in equipartition between the turbulent magnetic and kinetic energies. The pressure anisotropies excite microscale instabilities, whose effect is to pin the pressure anisotropy at marginal stability. The pressure anisotropy determines the viscous stress and therefore the heating rate. This heating balances radiative cooling, giving rise to a stable thermal equilibrium that is maintained by energy injection by self-regulated sources (such as AGNs) and/or self-regulated energy absorption by the turbulence.} \\label{fig:outline} \\end{figure} There are several peripheral consequences of our predictions that warrant further discussion. Firstly, the relatively large magnetic-field strengths inferred from observed rotation measures and predicted by a balance between parallel viscous heating and radiative cooling are stable or at least marginally stable to the HBI. The stability criterion for the HBI \\citep{quataert08} may be written as a lower-bound on the magnetic-field strength: \\begin{eqnarray}\\label{eqn:hbi} \\lefteqn{B \\gtrsim 3\\left(\\frac{g}{10^{-8}~{\\rm cm~s}^{-2}}\\right)^{1/2} \\left(\\frac{n_{\\rm e}}{0.1~{\\rm cm}^{-3}}\\right)^{1/2}\\left(\\frac{r}{100~{\\rm kpc}}\\right)^{1/2} }\\nonumber\\\\*\\mbox{} &&\\times \\left(\\frac{{\\rm d}\\ln T/{\\rm d}\\ln r}{0.1}\\right)^{1/2}~\\mu{\\rm G} . \\end{eqnarray} Even under the favourable conditions we have chosen for the electron density $n_{\\rm e}$, the radial distance from the cluster centre $r$ and the temperature slope $d\\ln T/d\\ln r$, the magnetic fields predicted in this paper exceed the stability limit (\\ref{eqn:hbi}). Indeed, \\citet[\\S~3.3]{brbp09} and \\citet[\\S~5.7]{pqs09} found that simulated clusters with magnetic-field strengths $B\\sim 3~\\mu{\\rm G}$ demonstrated a delayed cooling catastrophe due to the stabilisation of HBI modes by magnetic tension. By contrast, the rms magnetic-field strength used in the \\citet{pqs10} simulation, where a combination of turbulent stirring and HBI-governed thermal conduction gave a cool core in long-term thermal balance, was $10^{-9}~{\\rm G}$. More observational estimates of magnetic-field strengths in a variety of clusters would clearly be very useful to help realistic modelling and testing of theories. Secondly, the same process that is responsible for stably heating the ICM in our model may also influence the turbulent diffusion of metals throughout the cores of galaxy clusters. That our values for the turbulent diffusion coefficient $\\kappa_{\\rm turb}$ are comparable to those inferred in a variety of clusters is encouraging. However, \\citet{rcbf05} found that diffusion coefficients that increase with radius (as ours does) imply abundance profiles that are too centrally peaked compared to those observed. It is therefore worth investigating the diffusion of metals using our predicted scaling $\\kappa_{\\rm turb} \\propto n^{-1}_{\\rm e} T^{5/2}$. Thirdly, turbulent heat diffusion may become dominant relative to collisional electron heat conduction if the Spitzer conductivity suppression factor $f\\lesssim 0.03$ (see eqns \\ref{eqn:kprofile} and \\ref{eqn:ksp1}). However, for typical density and temperature profiles of cool-core clusters, neither form of heat diffusion seems to be important relative to parallel viscous heating (see Section \\ref{sec:conduction}). It is prudent to reiterate here our assumptions and their limitations: \\begin{enumerate} \\item We have implicitly assumed that there is enough turbulent energy so that, if it is thermalised via parallel viscous heating, it can offset cooling. This requires either the source of the turbulent energy or the amount of energy locally accepted by the plasma and converted into turbulent motions to have some knowledge of the cooling rate and to be self-regulating. AGNs are a natural candidate for providing a self-regulating energy source, whether the stirring is due to AGN-driven jets, bubbles, weak shocks/sound waves, gravitational modes and/or cosmic-ray--buoyancy instabilities. Observationally, a large majority of cool-core clusters harbour radio sources at their centres, and the AGN energy output inferred from radio-emitting plasma outflows and cavities is often similar to the X-ray cooling rate of the central gas \\citep[e.g.][]{bcsrm10}. However, it may not be necessary for the energy source to be fine-tuned or tightly self-regulating. Our scheme for heating the ICM can deal with excess turbulent energy via the following two considerations. First, not all of the power provided by the external stirring has to be locally thermalised via turbulence, as the turbulence may have an effective `impedance' and only accept the amount of power that can be locally viscously dissipated without triggering the microinstabilities; the remaining power could be transported elsewhere. Second, the parallel viscous heating rate (eq. \\ref{eqn:heating2}), coupled with the assumption of equipartition kinetic and magnetic energies (eq. \\ref{eqn:equipartition}), is naturally self-regulating: any increase in turbulent energy implies an increase in magnetic energy, which implies an increased viscous dissipation rate. Both of these regulation mechanisms require there to be a sufficient amount of turbulence to pin the pressure anisotropy at its marginal stability threshold. As it happens, there seems to be no dearth of turbulent power \\citep[e.g.][]{cfjsb04}. The fact that some of this power can be thermalised via parallel viscous dissipation in a thermally stable way provides an attractive alternative to other heating models that suffer from stability issues. Further work on the excitation and replenishment of turbulence in a weakly collisional ICM is clearly needed. \\item The heating is assumed to be all due to parallel collisional viscosity. In principle, energy could cascade through the parallel viscous scale and onwards to collisionless scales via Alfv\\'{e}nic turbulence and then heat the plasma via microphysical dissipation at the ion and electron Larmor scales \\citep[see][and references therein]{scdhhqt09}. This possibility has been ignored here. \\item The predicted turbulent velocity, obtained by assuming approximate equipartition between kinetic and magnetic energies, is probably a lower limit on the actual turbulent velocity since $U_{\\rm rms}$ is unlikely to be smaller than the Alfv\\'{e}n speed. Unfortunately, a more exact estimate of $U_{\\rm rms}$ would require a detailed understanding of turbulent dynamo saturation in the ICM, far beyond the scope of the present work. \\item Our estimates of the characteristic turbulence scale and turbulent diffusion coefficient are perhaps the more uncertain of our predictions, as they depend on the rate of transfer of energy from the external driving sources into the turbulence. This encodes what we have referred to above as the effective `impedance' of the turbulence, the detailed physics of which is poorly understood. \\end{enumerate} All these concerns clearly require further theoretical work. However, the fact that our predictions are not too far from current observational estimates or plausible expectations lends us hope that these concerns might be of secondary importance. While the exact numbers predicted in this paper should be taken with a grain of salt, one cannot help being encouraged by the fact that they seem to be quite reasonable without any fine-tuning of adjustable prefactors. The basic conjecture that the saturation of microscale plasma instabilities endows the ICM with a thermally stable source of viscous heating seems robust. If our predictions are confirmed, this would constitute strong evidence that microphysical plasma processes play a decisive role in setting the large-scale structure and evolution of galaxy clusters." }, "1003/1003.0236_arXiv.txt": { "abstract": "It seems generic to have vacua with lower dimensionality than ours. We consider the possibility that the observable universe originated in a transition from one of these vacua. Such a universe has anisotropic spatial curvature. This may be directly observable through its late-time effects on the CMB if the last period of slow-roll inflation was not too long. These affect the entire sky, leading to correlations which persist up to the highest CMB multipoles, thus allowing a conclusive detection above cosmic variance. Further, this anisotropic curvature causes different dimensions to expand at different rates. This leads to other potentially observable signals including a quadrupolar anisotropy in the CMB which limits the size of the curvature. Conversely, if isotropic curvature is observed it may be evidence that our parent vacuum was at least 3+1 dimensional. Such signals could reveal our history of decompactification, providing evidence for the existence of vastly different vacua. ", "introduction": "\\label{Sec:Intro} Our current understanding of cosmology and high energy physics leaves many questions unanswered. One of the most fundamental of these questions is why our universe has three large dimensions. This may be tied to the more general question of the overall shape and structure of the universe. In fact, it is possible that our universe was not always three dimensional or that other places outside of our observable universe have a different dimensionality. There are surely long-lived vacua where one or more of our three dimensions are compactified, since this does not even rely on the presence of extra-dimensions and indeed happens in the Standard Model \\cite{ArkaniHamed:2007gg}. Eternal inflation can provide a means to populate these vacua, and naturally leads to a highly inhomogeneous universe on very long length scales. Further, it seems likely that these lower-dimensional vacua are at least as numerous as three dimensional ones since there are generally more ways to compactify a greater number of spatial dimensions. If we do indeed have a huge landscape of vacua (e.g.~\\cite{Bousso:2000xa}) then it seems all the more reasonable that there should be vacua of all different dimensionalities and transitions between them (see e.g.~\\cite{Krishnan:2005su, Carroll:2009dn, BlancoPillado:2009di, BlancoPillado:2009mi, Wu:2003ht}). We will ignore the subtle issues of the likelihood of populating those vacua (the ``measure problem\"). Instead we will focus on the possibility of observing such regions of lower dimensionality since surely such a discovery would have a tremendous effect on our understanding of cosmology and fundamental physics. Our compact dimensions are generically unstable to decompactification \\cite{Giddings:2004vr}. Thus it seems possible that the universe began with all the dimensions compact (the starting point in \\cite{Brandenberger:1988aj, Greene:2009gp} for example). In this picture our current universe is one step in the chain towards decompactifying all dimensions. Of course, eternal inflation may lead to a very complicated history of populating different vacua, but in any case, it seems reasonable to consider the possibility that we came from a lower dimensional ``ancestor\" vacuum. We will assume that prior to our last period of slow-roll inflation our patch of the universe was born in a transition from a lower dimensional vacuum. Our universe then underwent the normal period of slow-roll inflation. For our signals to be observable we will assume that there were not too many more than the minimal number of efolds of inflation necessary to explain the CMB sky. This may be reasonable because this is very near a catastrophic boundary: large scale structures such as galaxies would not form if inflation did not last long enough to dilute curvature sufficiently \\cite{Vilenkin:1996ar, Garriga:1998px, Freivogel:2005vv}. Since achieving slow-roll inflation is difficult and the longer it lasts the more tuned the potential often is, there may be a pressure to be close to this lower bound on the length of inflation. We will actually use the energy density in curvature, $\\Omega_k$, in place of the number of efolds of inflation. The observational bound requires that $\\Omega_k \\lesssim 10^{-2}$ today (this corresponds to $\\sim 62$ efolds for high scale inflation). The existence of galaxies requires $\\Omega_k \\lesssim 1$ today (corresponding to $\\sim 59.5$ efolds if we use the bound from \\cite{Freivogel:2005vv}). Thus $\\Omega_k$ may be close to the observational bound today. Other, similar arguments have also been made for a relatively large curvature today \\cite{Bozek:2009gh}. Most signals of the presence of other vacua, e.g.~bubble collisions \\cite{Aguirre:2007an, Chang:2007eq, Chang:2008gj}, also rely on this assumption. These signals have also mostly been explored assuming that the other vacua are all 3+1 dimensional. While an important first step, this seems like a serious oversimplification. We find interesting differences in the case that our parent vacuum was lower dimensional. In particular, our universe can be anisotropic, with different spatial curvatures in the different directions. This anisotropic curvature dilutes exponentially during inflation, making the universe appear very isotropic at early times. However, this curvature ($\\Omega_k$) grows at late times, leading to several observable effects. This anisotropic curvature sources an anisotropy in the Hubble expansion rate, since the different dimensions expand at different rates. The most interesting signal is an anisotropy in the normal CMB curvature measurement. The angular size of a ``standard ruler\" now appears to depend on the orientation of that ruler. In the CMB this shows up as unexpected correlations between modes of all angular sizes. Unlike the normal curvature measurement, this anisotropic curvature measurement is not degenerate with the scale factor expansion history and is thus easier to measure. This anistropic curvature also leads to a significant quadrupolar anisotropy in the CMB which constrains the size of $\\Omega_k$. There are possibly other observables from 21 cm measurements, direct measurements of the Hubble expansion (e.g.~from supernovae), or from searches looking for nontrivial topology of the universe. ", "conclusions": "\\label{Sec: Conclusions} A universe produced as a result of bubble nucleation from an ancestor vacuum which has two large dimensions and one small, compact dimension is endowed with anisotropic curvature $\\Omega_k$. Such an anisotropic universe is also produced in the case when our 3+1 dimensional universe emerges from a transition from a 1+1 dimensional vacuum. In this case, depending upon the curvature of the compact dimensions, the resulting universe can have either positive or negative curvature along two dimensions, with the other remaining flat. The geometry of the equal time slices of the daughter universe are such that two of the directions are curved while the other dimension is flat. Immediately after the tunneling event, the energy density of the universe is dominated by this anisotropic curvature $\\Omega_k$. This curvature drives the curved directions to expand differently from the flat direction, resulting in differential Hubble expansion $\\Delta H$ between them. The expansion of the universe dilutes $\\Omega_k$ until it becomes small enough to allow slow roll inflation. At this time, the universe undergoes a period of inflation during which the curvature $\\Omega_k$ and the differential Hubble expansion $\\Delta H$ are exponentially diluted. However, during the epochs of radiation and matter domination, the curvature red shifts less strongly than either the radiation or the matter density. Consequently, the fractional energy density $\\Omega_k$ in curvature grows with time during these epochs. This late time emergence of an anisotropic curvature $\\Omega_k$ also drives a late time differential Hubble expansion $\\Delta H$ in the universe. These late time, anisotropic warps of the space-time geometry are all proportional to the current fractional energy density in curvature, $\\Omega_{k_0}$. They can be observed in the present epoch if inflation does not last much longer than the minimum number of efolds required to achieve a sufficiently flat universe ($\\sim 65$ efolds for high scale inflation). Anisotropic curvature leads to the warping of the angular size of standard rulers. This warping is a function of both the angle and orientation of the ruler in the sky. Consequently, this effect is immune to degeneracies from the expansion history of the universe since it affects rulers that are along the same line of sight but oriented differently. The CMB is also warped by the anisotropic curvature. In addition to the geometric warping, the differential Hubble expansion $\\Delta H$ also preferentially red shifts the energies of the CMB photons. This energy shift differentially changes the monopole temperature of the CMB giving rise to a quadrupole in the CMB. Furthermore, since the anisotropic curvature is a late time effect, it affects all the modes that can be seen in the CMB. Consequently, this effect leads to statistical anisotropy on all angular scales. This effect is different from other signatures of anisotropy considered in the literature \\cite{Gumrukcuoglu:2007bx, Gumrukcuoglu:2008gi}. Previous work has concentrated on the correlations that are produced due to the initial anisotropy in the universe at the beginning of inflation. Since these modes are roughly stretched to the Hubble size today, these initial anisotropies only affect the largest modes in the sky and are hence low l effects in the CMB. The late time anisotropy however warps the entire sky and leads to statistically robust correlations on all angular scales. The anisotropies in the pre-inflationary vacuum can however lead to other interesting signatures, for example in the gravitational wave spectrum \\cite{Gumrukcuoglu:2008gi}. These signatures are an independent check of this scenario. Anisotropies that affect all angular scales have also been previously considered \\cite{Ackerman:2007nb, Boehmer:2007ut}. These required violations of rotational invariance during inflation and the anisotropy emerges directly in the primordial density perturbations. In our case, the density perturbations are isotropic and the anisotropy observed today is a result of a late time warp of the space-time. Anisotropic curvature is already more stringently constrained than isotropic curvature. While isotropic curvature is bounded to be $\\lessapprox 10^{-2}$, it is difficult for anisotropic curvature to be much larger than $\\sim 10^{-4}$ without running afoul of current data, in particular, the size of the CMB quadrupole. Since the measurement of curvature is ultimately limited by cosmic variance $\\sim 10^{-5}$, there is a window between $ 10^{-5} \\lessapprox \\Omega_{k_0} \\lessapprox 10^{-4}$ that can be probed by upcoming experiments, including Planck. Future cosmological measurements like the 21 cm experiments will significantly improve bounds on the curvature of the universe. A discovery of isotropic curvature would be evidence suggesting that our ancestor vacuum had at least three large space dimensions. On the other hand, a discovery of anisotropic curvature would be strong evidence for the lower dimensionality of our parent vacuum. The anisotropy produced from such a transition has a very specific form due to the symmetries of the transition. It leads to correlations only amongst certain modes in the CMB (for example, only $A^{20}_{ll}$ and $A^{20}_{l, l-2}$). This distinguishes it from a generic anisotropic $3+1$ dimensional pre-inflationary vacuum which will generically have power in all modes. In these scenarios, it is also natural for the universe to have non-trivial topology. The existence of a non-trivial topological scale within our observable universe can be searched for using the classic ``circles in the sky\" signal. If both the non-trivial topology and anisotropic curvature can be discovered, implying a period of inflation very close to the catastrophic boundary, it would be powerful evidence for a lower dimensional ancestor vacuum. A discovery of these effects would establish the existence of vacua vastly different from our own Standard Model vacuum, lending observational credence to the landscape." }, "1003/1003.0693_arXiv.txt": { "abstract": "Star formation-driven outflows are a critically important phenomenon in theoretical treatments of galaxy evolution, despite the limited ability of observational studies to trace galactic winds across cosmological timescales. It has been suggested that the strongest \\mgii\\ absorption-line systems detected in the spectra of background quasars might arise in outflows from foreground galaxies. If confirmed, such ``ultra-strong'' \\mgii\\ (\\usmgii) absorbers would represent a method to identify significant numbers of galactic winds over a huge baseline in cosmic time, in a manner independent of the luminous properties of the galaxy. To this end, we present the first detailed imaging and spectroscopic study of the fields of two \\usmgii\\ absorber systems culled from a statistical absorber catalog, with the goal of understanding the physical processes leading to the large velocity spreads that define such systems. Each field contains two bright emission-line galaxies at similar redshift ($\\Delta v \\la 300$ \\kms) to that of the absorption. Lower-limits on their instantaneous star formation rates (SFR) from the observed \\oii\\ and \\hb\\ line fluxes, and stellar masses from spectral template fitting indicate specific SFRs among the highest for their masses at these redshifts. Additionally, their 4000\\AA\\ break and Balmer absorption strengths imply they have undergone recent ($\\sim 0.01$ - 1 Gyr) starbursts. The concomitant presence of two rare phenomena -- starbursts and \\usmgii\\ absorbers -- strongly implies a causal connection. We consider these data and \\usmgii\\ absorbers in general in the context of various popular models, and conclude that galactic outflows are generally necessary to account for the velocity extent of the absorption. We favour starburst driven outflows over tidally-stripped gas from a major interaction which triggered the starburst as the energy source for the majority of systems. Finally, we discuss the implications of these results and speculate on the overall contribution of such systems to the global SFR density at $z \\simeq 0.7$. ", "introduction": "In the local Universe, large scale gas outflows are observed to arise in galaxies exhibiting high surface densities of star formation. While the precise roles of such outflows, including galactic ``superwinds'', in galaxy evolution are still being determined, simulations suggest that the balance between outflows and the accretion of cool gas is one of the primary mechanisms by which star formation is regulated in individual halos (e.g., Oppenheimer et al., 2009; Brooks et al., 2009). At the current epoch, the highest star formation rate (SFR) surface densities -- and therefore galactic winds -- are preferentially found in relatively low-mass halos, such as those hosting dwarf starburst galaxies. However, the mass of halos containing the highest specific star formation rates (sSFRs) are thought to increase with increasing look-back time, consistent with ``top-down'' galaxy formation scenarios (Cowie et al., 1996; Neistein et al., 2006). Tracking the occurrence of galactic winds across cosmic time, particularly in manners unbiased by luminosity or halo mass, would provide a powerful way to study such models. Furthermore, outflows are believed to be required to explain a wide variety of astrophysical observations, from the shape of the galaxy luminosity function (Benson et al., 2003; Khochfar et al., 2007), to the stellar mass-metallicity relation (Tremonti, 2004; Erb et al., 2006; Brooks et al., 2007; Finlator \\& Dav\\'{e}, 2008), the large extent of dust and metals in galactic halos and in the intergalactic medium (Scannapieco, Ferrara, \\& Madau, 2002; Oppenheimer \\& Dav\\'{e}, 2006; Kobayashi et al., 2007), and many other related phenomena. Despite their clear importance in the galaxy formation process, outflows have generally been overlooked theoretically and, until recently, have proved difficult to study observationally at redshifts $z \\simeq 1 - 4$, the epoch when the Universe formed most of its stars and superwinds were ubiquitous. Surveys have identified outflowing gas from galaxies at redshift $z>0$ through strong, blue-shifted resonance-line absorption in their spectra arising in low-ion gas entrained in the flows (e.g., Pettini et al., 2001; Shapley et al., 2003; Tremonti, Moustakas, \\& Diamond-Stanic, 2007; Wiener et al., 2009). However, there are two important limitations inherent in methods which rely only on spectra of the outflow hosts. First, they provide no information on the location of the outflowing gas; it is only presumed that the material reaches the IGM. Secondly, these surveys searched for evidence of outflows in spectra of either the brightest galaxies at the relevant redshift, known star-forming galaxies, or known post-starburst galaxies. What is needed is the reverse experiment: a survey of the galaxies from which known large-scale outflows originate. This begs the question: how does one identify a galactic wind without {\\it a priori} knowledge of the galaxy itself? Quasar absorption lines may offer such an opportunity, as they select galaxies based on the gas absorption cross section, with no direct dependence on emission from the galaxy. The physical processes that determine the properties of intervening low-ion quasar absorption line systems are not well understood. While it has long been known that such absorbers can in general be identified with individual galaxies (e.g., Bergeron \\& Boiss\\'{e}, 1991; Steidel, Dickinson, \\& Persson, 1994) correlations between emission (i.e., of the galaxy) and absorption (i.e., strength and velocity structure of the absorbing gas) properties have been elusive (Kacprzak et al., 2007) and/or inconclusive (e.g., Steidel et al., 2002; Kacprzak et al., 2010). The strongest absorbers, which have until recently been neglected due to their relative scarcity, may hold important clues. For example, Bond et al.\\ (2001) considered the velocity profiles of the strongest \\mgii\\ absorbers known at the time (rest equivalent widths \\W$\\sim 2$\\AA) measured with high-resolution spectroscopy, and proposed that such systems may arise in galactic superwinds. Detections of outflows through broad low-ion absorption in the spectra of starbursting galaxies suggest that galactic winds {\\it can} result in very strong \\mgii\\ absorption along a sightline past a galaxy. However, this does not necessarily imply that all or even any of the strongest intervening absorbers detected in quasar spectra actually {\\it do} arise in the winds of foreground galaxies. Indeed, models have been proposed that account for the observed distribution of \\mgii\\ absorption strength without relying on outflows (e.g., Tinker \\& Chen 2008). Alternatively, the huge kinematic spreads that define the strongest systems may be due simply to the chance intersection of the sightline with multiple ``normal'' \\mgii\\ absorbing galaxies in a rich group or cluster, as first suggested by Pettini et al.\\ (1983). Very large \\mgii\\ absorber surveys (e.g., Nestor et al., 2005; Prochter, Prochaska, \\& Burles, 2006; Quider et al., 2010), which are now becoming available, have uncovered large numbers (e.g., $>600$ in Quider et al.) of ``ultra-strong'' \\mgii\\ (\\usmgii) systems with \\W$\\ge 3$\\AA. Catalogs of such systems afford the opportunity to explore in depth the proposed \\usmgii\\ absorber-galactic wind connection. Recent work has already given support for a connection between the strongest \\mgii\\ systems and star forming galaxies, which is usually considered as support of an outflow scenario. For example, by stacking thousands of relatively-shallow Sloan Digital Sky Survey (SDSS) images of the fields of strong \\mgii\\ absorption systems Zibetti et al.\\ (2007) demonstrated the strongest systems are associated with bluer galaxies closer to the sightline to the background quasar compared to weaker systems. Similarly, Bouch\\'{e} et al.\\ (2007) have detected strong H$\\alpha$ emission at the absorption redshift towards strong \\mgii\\ absorbers and Rubin et al.\\ (2009) have identified an \\usmgii\\ absorber in the spectrum of a background galaxy that they identify with a wind from a foreground galaxy. Perhaps the most compelling evidence suggesting a connection between \\mgii\\ absorbers and star formation is the relation between \\W\\ and \\oii\\ emission discussed by M\\'{e}nard et al.\\ (2009), wherein they demonstrate that the strongest \\mgii\\ absorbers are on average associated with the highest \\oii\\ luminosity densities and are therefore likely to be hosted by vigorously star-forming galaxies. Nestor et al.\\ (2007; hereafter NTRQ) published the first imaging survey aimed specifically at the strongest \\mgii\\ absorption systems, including images of the fields of thirteen moderate-redshift ($0.42 < z < 0.84$) \\usmgii\\ systems. These revealed bright galaxies at relatively low impact parameter to the absorption sightline (compared to the fields of most \\mgii\\ absorbers). While consistent with the outflow model in general and, e.g., the results of Zibetti et al., in particular, detailed study of the galaxies associated with \\usmgii\\ absorbers is needed to test this putative connection. If the ``ultra-strong'' nature of these absorbers is indeed linked to galactic winds, we expect to find evidence of recent, high mass fraction starbursts in one or more of the low impact parameter (low-$b$) galaxies. In this paper, we present the results of an imaging and spectroscopic study of the galaxies detected in two \\usmgii\\ absorber fields from the NTRQ sample, conducted to test both the \\usmgii\\ absorber/outflow connection and alternative models. In \\S2 we describe the selection of targets, observations and reductions of our data, and our basic observational results. In \\S3 we discuss the properties of the galaxies determined to be at similar redshift to the \\usmgii\\ systems. We present further discussion in \\S4 and \\S5, and summarize the paper in \\S6. Throughout the paper we assume a cosmology with $\\Omega_M = 0.3$, $\\Omega_\\Lambda = 0.7$ and $H_0 = 70$~{\\mbox{km\\,s$^{-1}$\\,Mpc$^{-1}$} and state magnitudes in the AB system. ", "conclusions": "\\label{sect:Discusion} We undertook the observations discussed above with the goal of uncovering the physical mechanism driving the extreme absorption velocity spreads that define \\usmgii\\ systems. In this section, we test the star formation-driven galactic wind model in light of the results of these observations, and find abundant circumstantial evidence in its favour. We conclude the section by considering other popular models, and find tidally-stripped gas from interacting galaxies is also consistent with the observational results for some \\usmgii\\ systems. \\subsection{Evidence from Bayesian probability} First, we consider the possibility that the presence of \\mbox{(post-)}starburst galaxies in the vicinities of ultra-strong \\mgii\\ absorption are coincidental. Indeed, the existence of a strong \\mgii\\ absorber requires a galaxy at $z \\simeq z_{abs}$. However, galaxies with such properties as we find for \\mbox{G07-1}, \\mbox{G07-2}, \\mbox{G14-1} and \\mbox{G14-2} are uncommon (e.g., Bell \\& de Jong 2001; Feulner et al., 2005; Wild et al., 2009). Similarly, the presence of a galaxy produces a likelihood of detecting a \\mgii\\ absorber with $z \\simeq z_{gal}$, but \\mgii\\ systems with \\W $\\ge 3.6$\\AA\\ account for only 1 per cent of strong (\\W $\\ge 0.3$\\AA) systems, and those with \\W $\\ge 5.6$\\AA\\ only 0.05 per cent. Thus it is exceedingly unlikely that the (post-)starburst nature of the galaxies and the ultra-strong nature of the \\mgii\\ absorption are unrelated. This does not necessarily imply the starbursts are driving a wind that is responsible for the \\usmgii\\ absorption. For example, it is conceivable that they share a common cause, such as a major interaction stripping gas out to large galactic radii while simultaneously triggering a starburst, or that such an interaction channeled gas to the galaxy centers and fed galactic nuclear activity, which in turn drove the outflow. \\subsection{Evidence from kinematic absorption spread} \\label{Sec:kin} The minimum rest-frame velocity width of a completely saturated, opaque \\mgii\\ $\\lambda2796$ feature is given by $\\Delta v_{min}= (\\Delta\\lambda/\\lambda)\\times c = $(\\W$/$1\\AA)$\\times 107$\\,\\kms. For \\usmgii\\ systems (i.e., 3\\AA\\ $<$ \\W\\ $\\la$ 6\\AA), this corresponds to 320 \\kms\\ $\\la \\Delta v_{min} \\la$ 640 \\kms. Actual kinematic spreads are typically larger due to partially non-opaque profiles. As discussed in \\S\\ref{Sec:targets}, the absorber towards Q0747+305 (figure~\\ref{Fig:0747}) and Q1417+011 (figure~\\ref{Fig:1417}) have observed profiles consistent with $\\Delta v \\simeq 390$~\\kms, and $\\Delta v \\approx 1000$~\\kms, respectively. Thus, the physical mechanism behind \\usmgii\\ systems must be one that is able to produce cool gas that, along a single line-of-sight, continuously spans many hundreds of \\kms\\ with a total dynamic spread of up to $\\ga$~1000~\\kms. Such a dynamic range can naturally be obtained along a sightline passing though a galactic wind at some angle to the outflow orientation. \\subsection{Timescale, distance, and velocity consistency} The impact parameters of the four $z_{gal} \\simeq z_{abs}$ galaxies in the present study are comparable to those of the general population of \\mgii\\ absorbers (see, e.g., Kacprzak et al., 2007). They are, however, at the extreme end of the $b$-distribution for \\usmgii\\ systems (NTRQ). It is therefore worthwhile to consider the inferred velocities and timescales of the putative outflows in light of the distances needed to be traversed by outflowing material to cover the sightline to the QSO. The observed absorption velocity spread arises from projections of the outflow velocity (of sufficient columns of low-ion gas) onto the sightline, and thus depends on the unknown geometry. However, order of magnitude estimates of the outflow speeds can be made by considering the red- and blue-most velocity extents relative to the galaxy systemic velocity (e.g., figures \\ref{Fig:0747} and \\ref{Fig:1417}) together with a conic outflow geometry featuring opening angles between $\\simeq 45^{\\deg}$ and $100^{\\deg}$ (Veilleux, Cecil, Bland-Hawthorn, 2005). While such a calculation is overly simplistic, it should give an order of magnitude estimate for the outflow speeds which can be interpreted together with our estimations of the age of the most recent starburst. For the Q0747+305 field, we estimate an outflow velocity of $v \\sim 300$ - 600~\\kms. This range is comparable to our approximation of the escape velocities for these galaxies (\\S \\ref{Sec:props}). If outflows from both galaxies contribute to the observed profile, then the velocity could be as low as $v \\sim 200$~\\kms. At these velocities, the gas would have to have flowed for a minimum of $\\sim100$ - 500~Myrs to have reached the distance to the sightline. These timescale estimates are well below our estimate of a $\\sim 1$~Gyr age for the bursts in \\mbox{G07-1} and \\mbox{G07-2}. Thus, the relatively large impact parameters in this field are completely consistent with a scenario involving outflows driven by the starburst event. The starbursts in the Q1417+011 field are likely much younger. As can be seen in figure \\ref{Fig:1417}, however, much larger velocities are also necessary to explain the absorption. For an outflow to account for the absorption, we estimate $v \\sim 800$ - 1000~\\kms, which exceeds our approximations of the escape velocities. The velocity estimate changes little if both galaxies contribute, as they are at similar redshift and both $z\\beta_{pair}$), but softer spectral slope at low frequency ($\\alpha_{baryonic}<\\alpha_{pair}$). The electrostatic field, although screened, also plays a significant role for spectrum of emitted radiation. In the up-stream region ahead of the shock, where streaming instabilities are of early stage type, screened electrostatic effects are likely to produce heavily observer dependent effects. For global shock simulations with non-periodic boundary conditions such a non-isotropic acceleration mechanism will be sustained. Our results establish tractability of an alternative approach to radiation modeling -- the jitter radiation approach. It contrasts the standard synchrotron approach; assuming a pure synchrotron mechanism in observed relativistic streaming phenomena, we would infer --- incorrectly --- that large scale homogenous magnetic fields were present, when rather it might be small scale fields, in (a priori) non-magnetized relativistic shocked outflows. A tight correspondence exists between on one hand the spectral evolution of radiation emitted from electrons in optically thin collisionless plasmas and, on the other, the field spectral evolution in the plasma due to the instabilities that create and mediate such radiation. Care should be exhibited with reduced dimensionality modeling; as 3D models become increasingly affordable they are to be preferred over 2D models, as spectral diagnostics can be distorted as a consequence of reduced dimensionality---filament entanglement is for example only possible in 3D. Attention should also be given to choices of plasma constituents (electrons, protons, ions, etc.). To conclude, the radiation spectral modeling results presented here reveal limitations that need to be considered while reverse engineering spectra observed from, for example, relativistic astrophysical outflows, GRBs, relativistic reconnection sites, and from laser-plasma interaction in laboratory astrophysics, when examining likely physical scenarios for their match to observed spectra." }, "1003/1003.1376_arXiv.txt": { "abstract": "We infer the 3D magnetic structure of a \\emph{transient horizontal magnetic field} (THMF) during its evolution through the photosphere using SIRGAUS inversion code. The SIRGAUS code is a modified version of SIR (Stokes Inversion based on Response function), and allows for retrieval of information on the magnetic and thermodynamic parameters of the flux tube embedded in the atmosphere from the observed Stokes profiles. Spectro-polarimetric observations of the quiet Sun at the disk center were performed with the Solar Optical Telescope (SOT) on board \\emph{Hinode} with \\ion{Fe}{1} 630.2~nm lines. Using repetitive scans with a cadence of 130~s, we first detect the horizontal field that appears inside a granule, near its edge. On the second scan, vertical fields with positive and negative polarities appear at both ends of the horizontal field. Then, the horizontal field disappears leaving the bipolar vertical magnetic fields. The results from the inversion of the Stokes spectra clearly point to the existence of a flux tube with magnetic field strength of $\\sim400$~G rising through the line forming layer of the \\ion{Fe}{1} 630.2~nm lines. The flux tube is located at around $\\log\\tau_{500} \\sim0$ at $\\Delta t$=0~s and around $\\log\\tau_{500} \\sim-1.7$ at $\\Delta t$=130~s. At $\\Delta t$=260~s the horizontal part is already above the line forming region of the analyzed lines. The observed Doppler velocity is maximally 3~km~s$^{-1}$, consistent with the upward motion of the structure as retrieved from the SIRGAUS code. The vertical size of the tube is smaller than the thickness of the line forming layer. The THMF has a clear $\\Omega$-shaped-loop structure with the apex located near the edge of a granular cell. The magnetic flux carried by this THMF is estimated to be $3.1\\times10^{17}$ Mx. ", "introduction": "The SOT spectropolarimeter \\citep[SP,][]{Tsuneta2008SoPh,Suematsu2008SoPh,Ichimoto2008SoPh,Shimizu2008SoPh} on board \\emph{Hinode} \\citep{Kosugi2007} revealed that there are ubiquitous horizontally inclined magnetic fields with strengths of a few hundred G in the internetwork regions \\citep{Orozco2007,Lites2007}. \\citet{Centeno2007} and \\citet{Ishikawa2008} reported the temporal evolution of such horizontal magnetic fields both in the quiet Sun and in plage regions. The evolution commonly starts with the appearance of linear polarization signals (horizontal fields) inside granular cells, followed by the appearance of the opposite-polarity circular polarization signals (vertical fields) at both ends of the horizontal component. The time evolution indicates that the horizontally inclined magnetic fields are part of a small magnetic loop. The sizes of the horizontal magnetic fields are smaller than the size of the granule where they appear. Statistical analyses of a large number of the horizontal fields indicate that the life time of the horizontal magnetic fields ranges from 1 to 10 minutes and the mean lifetime is about 4 minutes \\citep{Ishikawa2009,Ishikawa_Hinode2}. Their occurrence rate is high, and more than approximately 10\\% of the granules have embedded horizontal magnetic fields \\citep{Ishikawa2008}. The movie of linear polarization signals \\citep{Ishikawa2009} shows that these horizontal magnetic fields are highly frequent and transient, and the term ''transient horizontal magnetic fields'' (THMFs) is used for them. Earlier reports of flux emergence in the quiet Sun \\citep{Lites1996,DePontieu2002,MartinezGonzalez2007} appear to be examples of THMFs. The properties of the THMFs described above are the same in the quiet Sun and the plage region. There are granular-sized horizontal fields in the polar regions, and these magnetic fields may be the same as THMFs \\citep{Tsuneta2008ApJ,Itoh2009}. However, there seem to be a few distinct patterns in their disappearance. For example, \\citet{Ishikawa2008} report an example of a THMF disappearing in the intergranular lane, while \\citet{Centeno2007} show a THMF which disappears inside a bright granule. This is confirmed by \\citet{Ishikawa_Hinode2} who report that normalized continuum intensities where THMFs appear are above $\\sim$1.0, corresponding to the bright granular regions. However, the normalized continuum intensities where THMFs disappear are both higher and lower than $\\sim$1.0, i.e. THMFs do not necessarily reach the intergranular lanes. There also appears to be two types of THMFs; with or without apparent bipolar vertical fields after their disappearance. This variety of THMF behavior suggests a few different scenarios with regard to their disappearance: (1) THMFs somehow submerge with the downward convective motion; (2) They become fragmentary due to the convective flow, and their Stokes signals are below the SP detection level; (3) THMFs go through the line formation layer of the \\ion{Fe}{1}~630.2~nm lines and reach higher atmospheric layers \\citep{MartinezGonzalez2009}. Given the ubiquity of THMFs in the solar photosphere, the question of how they appear and disappear in such a short time becomes one of the important issues in solar magnetohydrodynamics (MHD). A series of sophisticated MHD simulations \\citep[e.g.,][]{Vogler2007, Abbett2007, Schussler2008, Isobe2008, Steiner2008} have been carried out with varieties of initial and boundary conditions. A detailed comparison of the observed properties of THMFs with the numerical simulation results has just begun. It is thus desirable to clarify the structure and evolution of the enigmatic small-scale magnetic fields in the photosphere, and possibly in the chromosphere. Spectropolarimetric data obtained with the SOT have been analyzed thus far mainly with inversion codes based on a Milne-Eddington atmosphere \\citep[e.g.,][]{Lites2007, Orozco2007, Ishikawa2009}. Milne-Eddington inversion provides us with various physical parameters (magnetic field vector and Doppler velocity, etc.), which have to be constant throughout the atmosphere (except for the source function), and synthesize only symmetric Stokes profiles. However, the observed Stokes profiles of the horizontal magnetic fields usually have area-asymmetries. Such asymmetric profiles suggest the presence of velocity and magnetic field gradients along the line of sight (LOS) \\citep[e.g.,][]{Solanki1988,Sanchez1992}. Thus using the Milne-Eddington inversion, physically important information about the small-scale horizontal magnetic fields may be lost. We are using an inversion code based on the SIR code \\citep[Stokes Inversion based on Response function,][]{RuizCobo1992,RuizCobo1994,DelToroIniesta1996,DelToroIniest2003}. This code allows for the change of atmospheric parameters along the LOS and thus fits the observed asymmetric Stokes profiles. \\citet{MartinezGonzalez2007} used the SIR inversion code to analyze a snapshot of the magnetic loops using the spectropolarimetric data of 1.56 $\\mu$m taken with the Tenerife Infrared Polarimeter. However, the used settings of the inversion assumed the plasma parameters to be constant with height in the atmosphere, except in magnetic field inclination and temperature. Recently, \\citet{Gomory2009} studied the temporal evolution of a similar event applying SIR inversions. However, they kept all magnetic parameters constant with height. Thus, these analyses resemble a Milne-Eddington inversion with a realistic temperature distribution. We fully exploit the inversion code without such constraints to analyze the observed Stokes profiles asymmetries and with a better spatial resolution. This allows us to follow the temporal evolution of the stratified atmosphere. In this paper, we study a THMF event whose properties are similar to those reported by \\citet{Centeno2007}. We carry out the full Stokes inversion with the SIRGAUS code \\citep{BellotRubio2003}, which is a modified version of SIR. The SIRGAUS code treats explicitly a magnetic flux tube embedded in the atmosphere with a Gaussian function, and is expected to be more suitable to obtain vertical (LOS) structure of the THMF. Indeed, we successfully identify an isolated flux tube in the photosphere, and track the emergence of the $\\Omega$-shaped flux tube through the photosphere. We describe statistical significance of our result in detail. With the clear identification of the upward moving flux tube, we are able to obtain various physical parameters that characterize the flux tube. Such information on the flux tube becomes available for the first time with the use of the SIRGAUS code. We finally discuss the magnetohydrodynamic properties of the THMF. ", "conclusions": "\\subsection{3D magnetic structure and evolution} We have analyzed a magnetic phenomena in which horizontal magnetic field patch appears first at the edge of a granular cell at $\\Delta t=0$ s. Later at $\\Delta t=130$ s, the size of the patch harbouring horizontal fields becomes larger in size, and vertical magnetic fields with opposite polarity appear at both ends of the magnetic patch. Then (at $\\Delta t=260$ s), the horizontal magnetic field disappears, while the vertical magnetic fields remain and separate. The whole process takes about 4 min. This event corresponds to a typical example of the so called transient horizontal magnetic field (THMF). We have successfully applied the SIRGAUS inversion code to this event and analyzed its 3D magnetic structure. The results show that the flux tube occupies only a fraction of the line formation region of the \\ion{Fe}{1} lines at 630.2~nm and that it has an upward movement. The magnetic flux tube rises from $\\log\\tau \\sim 0$ (0~km) to $\\log\\tau \\sim -1.7$ (244~km) between the first and the second scan that are separated in time by 130~s. The Doppler velocity of the flux tube consistently indicates the upward motion. The magnetic field strength of the flux tube is about 400~G. The azimuth direction of the magnetic field is aligned with the axis of the flux tube, and the flux tube does not show any helical structure. So far, to determine whether the observed magnetic loops correspond to the rising $\\Omega$ loops or to submerging U loops, Doppler velocity measurement have been used \\citep[e.g.,][]{MartinezGonzalez2009}. Our analysis with SIRGAUS presented here shows that the flux tube with shallow but clear $\\Omega$ shape rises. We observe only a small part of the shallow $\\Omega$ loop above the continuum formation layer and this might be just a part of a larger-scale magnetic structure located in the bulk of the convection zone. \\subsection{Force balance of flux tube} \\label{forcebalance} The field strength of the flux tube is $\\sim$400 G in the first scan ($\\Delta t=0$~s). Taking the velocity of the convective flow to be $v=2\\times10^{5}$~cm~s$^{-1}$ and the mean density at the base of the photosphere ($\\log\\tau\\sim0$) to be $\\rho=$2.7$\\times$ 10$^{-7}$~g~cm$^{-3}$, we find that the equipartition field strength is $B_{eq}=(4\\pi \\rho)^{1/2}v=370$~G. Therefore, the inferred field strengths are comparable to $B_{eq}$. This is consistent with the fact that the flux tube seems to maintain its integrity during its rise from the lower photosphere upwards. In this section, we compare the buoyancy force, $F_{b}=(\\rho_{e}-\\rho_{i})g$, with the magnetic tension force, $F_{T}=B^{2}/4\\pi L$, acting on the flux tube at $\\Delta t=$130~s. Here $\\rho_{i}$ and $\\rho_{e}$ are the plasma density inside and outside the flux tube respectively, $g$ is the gravity acceleration of $2.7\\times10^{4}$~cm~s$^{-2}$, and $L$ is the curvature radius of the $\\Omega$ loop. Here we use all the plasma parameters (temperature, field strength, and the density) at the height where the flux tube is located and the 3D geometrical scale of the flux tube (the length and the cross section). From the separation of the footpoints in the second scan (which is about 1320~km) and from the averaged magnetic field inclination in the footpoints (about 47$^\\circ$), we estimate $L$ to be 1250~km. The comparison is made at the apex of the flux tube which is located above $\\log\\tau\\sim-1.0$. There the temperature $T$ is $\\sim$4950~K (corresponding to the averaged value from $\\log\\tau=-1.0$ to $-$2.0 at pixel 2-14). Assuming that the temperature is the same inside and outside of the flux tube, $F_{b}$ can be written as $F_{b}=\\frac{B^{2}}{8\\pi} \\cdot \\frac{m_{p}g}{2k_{B}T}$, where $m_{p}$ is the proton mass, and $k_{B}$ the Boltzmann constant. The ratio between $F_{b}$ and $F_{T}$ is then $\\frac{F_{b}}{F_{T}}=\\frac{m_{p}gL}{4k_{B}T}\\sim2$. This clearly indicates that the buoyancy force, $F_{b}$, participates in the emergence process of the magnetic flux tube. If we suppose that the stationary state is reached, then the buoyancy force, the magnetic tension force, and the resultant drag force associated with the rising motion are balanced: \\begin{align} (\\rho_{e}-\\rho_{i})g \\cdot al-\\frac{B^{2}}{4\\pi L} \\cdot al-0.5C_{d}\\rho_{e}u^{2}l & = \\frac{B^{2}}{8\\pi} \\cdot \\frac{m_{p}g}{2k_{B}T}\\cdot al-\\frac{B^{2}}{4\\pi L} \\cdot al-0.5C_{d}\\rho_{e}u^{2}l \\notag\\\\ & = 0. \\label{eq:force} \\end{align} With $C_{d}$ being the aerodynamic coefficient of the order of unity \\citep{parker1979}, and $u$ the relative velocity of the flux tube with respect to the velocity of the surrounding plasma. $a$ and $l$ are the vertical extent and the horizontal size of the flux tube, respectively. We take $\\rho_{e}$ to be $7.78\\times10^{-8}$~g~cm$^{-3}$ (corresponding to the mean density from $\\log\\tau=-1.0$ to $-2.0$ at pixel 2-14). The relative velocity $u$ obtained from Eqn.~(\\ref{eq:force}) is then 2.3~km~s$^{-1}$, which is consistent with the observed relative Doppler velocity associated to the flux tube. As shown in section~\\ref{CrossShape}, the vertical extent of the flux tube ($a=190$ km) is smaller than its lateral size ($l=360$~km). Therefore, the cross section of the flux tube has a flattened shape. \\citet{magara2001} studied the emergence of large flux tubes with larger field strengths which form active regions. He found that the top of the loop decelerates when it reaches photospheric layers while the bottom of the flux tube continues to rise. The reason is that the photosphere is a convectively stable layer. Therefore the cross section of the tube changes from the circular shape to horizontally extended shape. The same process may take place for granular-sized magnetic flux tubes. \\citet{Steiner2008} pointed out that an emerging flux tube can get pushed to the middle and upper photosphere by overshooting convection, forming atmospheric layers full of horizontal fields. This may also contribute to the flattening of the flux tube. \\subsection{Magnetic flux and energy} The horizontal magnetic flux $\\Phi_{H}$ is estimated to be $3.1\\times10^{17}$ Mx near the apex of the flux tube. Recently \\citet{Ishikawa2009} reported that the occurrence rate of THMFs is about $1.1\\times10^{-10}$ km$^{-2}$ s$^{-1}$ in a quiet Sun region. Also, \\citet{Tsuneta2008ApJ} and \\citet{Itoh2009} have shown the presence of horizontal magnetic fields ubiquitously in the polar region, and that there is no difference on the intrinsic magnetic field strengths between the quiet Sun and the polar region. \\citet{Lites2009} suggested that the uniform fluctuations of longitudinal magnetogram signals over the whole disk observed by \\citet{Harvey2007} are the same phenomenon as the horizontal fields observed by \\emph{Hinode}. All these results together indicate that THMFs would have the same occurrence rate all over the solar surface. Combining these observations, we conclude that THMFs could carry a total magnetic flux of $1.8\\times10^{25}$~Mx per day to the higher atmosphere what corresponds to a total of $7\\times10^{28}$~Mx in one solar cycle (11 year). This estimation is based on the assumption that all THMFs reach higher atmospheric layers and that all of them carry the same amount of magnetic flux $\\Phi_{H}$ (given in this paper). Although our calculations are based only upon a single event, similar values have been found also by \\citet{Jin2009} and \\citet{MartinezGonzalez2009}. The magnetic flux carried by THMFs is much larger than the total magnetic flux in sunspot regions during one solar cycle which corresponds to about $10^{25}$ Mx \\citep{Harvey1993}. \\citet{Ishikawa2009} have estimated the magnetic energy flux carried by THMFs, and found that the amount of energy flux is comparable to the total chromospheric and coronal energy loss, assuming that all THMFs reach above the photosphere. In order to perform a more precise estimation of the total magnetic flux and the total magnetic energy that THMFs could carry to higher atmospheric layers, we need to know what percentage of THMFs reach the chromosphere or higher layers in addition to the precise estimation of the magnetic flux of each THMF. Such statistical analysis was done by \\citet{MartinezGonzalez2009}, who pointed out that 23\\% of the horizontal fields that emerge with clear bipolar footpoints in the photosphere reach the low chromosphere. \\subsection{Disappearance of THMFs} We show that the THMF analyzed in this paper is rising through the line forming layer of the \\ion{Fe}{1} 630.2~nm lines and that it reaches higher atmospheric layers. The flux tube essentially maintains its integrity during the rising motion. At the beginning of this paper, we addressed three different possibilities of the THMFs disappearance. The observed event is clearly into the third of the suggested mechanisms, i.e., the THMF goes through the photosphere reaching the layers above the line forming region of \\ion{Fe}{1} 630.2~nm lines. Some THMFs do not show any conspicuous footpoints in their lifetime \\citep{Ishikawa2008, Jin2009}. Such THMFs possibly harbor weaker magnetic fields than the equipartition field strength. Thus, the convective motion may destroy the flux tube before it reaches the upper photosphere, or it may force it to submerge below the photosphere. %" }, "1003/1003.1230_arXiv.txt": { "abstract": "We present high angular resolution ($\\theta_{syn} \\lesssim 0\\rlap.{''}2$) observations of the 23.1-GHz methanol (CH$_3$OH) transition toward the massive star forming region NGC 7538 IRS 1. The two velocity components previously reported by Wilson et al. are resolved into distinct spatial features with brightness temperatures ($T_B$) greater than $10^4$ K, proving their maser nature. Thus, NGC 7538 IRS 1 is the third region confirmed to show methanol maser emission at this frequency. The brighter 23.1-GHz spot coincides in position with a rare formaldehyde (H$_2$CO) maser, and marginally with a 22.2-GHz water (H$_2$O) maser, for which we report archival observations. The weaker CH$_3$OH spot coincides with an H$_2$O maser. The ratio of $T_B$ for the 23.1-GHz masers to that of the well-known 12.2-GHz CH$_3$OH masers in this region roughly agrees with model predictions. However, the 23.1-GHz spots are offset in position from the CH$_3$OH masers at other frequencies. This is difficult to interpret in terms of models that assume that all the masers arise from the same clumps, but it may result from turbulent conditions within the gas or rapid variations in the background radiation field. ", "introduction": "\\label{sec:intro} NGC 7538 is a Galactic star forming region located at a distance of $2.65^{+0.12}_{-0.11}$ kpc \\citep{Mosca09}. Multiple infrared sources were discovered in the vicinity by \\cite{Wynn74}. The brightest of these, IRS~1, was recently resolved in the near-IR and mid-IR by \\cite{Kraus06} and \\cite{DBM05}, respectively. The centimeter free-free emission has a bipolar structure \\citep{Camp84,Sandell09,Zhu10} and shows variability \\citep{Ramiro04}. Also, the recombination lines from the ionized gas are unusually broad \\citep{Gaume95,Sewi04,KZK08}. NGC 7538 IRS~1 is an exceptionally rich maser source. Maser emission has been detected in hydroxyl (OH), water (H$_2$O), ammonia (NH$_3$), methanol (CH$_3$OH), and formaldehyde (H$_2$CO) (e.g., Gaume et al. 1991, hereafter G91; Hutawarakorn \\& Cohen 2003; Hoffman et al. 2003, hereafter H03; Kurtz et al. 2004). CH$_3$OH emission at 23.1~GHz, from the $9_2-10_1$ $A^+$ transition, was reported for NGC 7538 by \\cite{Wil84} (hereafter W84). They speculated that the emission was maser in nature, based on observations taken with the Effelsberg radio telescope ($\\theta_{beam} \\approx 43\\arcsec$). The well-known 6.7-GHz ($5_1-6_0$ $A\\tothe +$) and 12.2-GHz ($2_0-3_{-1}$ $E$) class II CH$_3$OH masers have also been detected toward this source (Minier et al. 2000, hereafter M00; Minier et al. 2002; Moscadelli et al. 2009, hereafter M09). Evidence for a circumstellar disk in IRS~1, based on a velocity gradient within a linear structure of 6.7-GHz and 12.2-GHz methanol maser spots, was presented by \\cite{Pesta04, Pesta09} \\citep[see however][]{DBM05}. The warm molecular gas surrounding the hypercompact \\HII region (IRS 1) also has dynamics indicative of rotation \\citep{Klaass09}. The 23.1-GHz class II CH$_3$OH maser is quite rare. Until now, only two regions harboring these masers have been confirmed: W3(OH) and NGC 6334 F (W84; Menten et al. 1985; Menten et al. 1988; Menten \\& Batrla 1989). \\cite{Cragg04} observed 50 southern star-forming regions and detected 23.1-GHz maser emission in only one --- the previously known NGC 6334 F. The 4.8-GHz ($1_{10}-1_{11}$) H$_2$CO maser in NGC 7538 IRS 1 is also quite rare; at present, only seven of these masers have been found in the Galaxy \\citep[][and references therein]{Araya08}. And NH$_3$, although ubiquitous in massive star forming regions, is not a common maser. Various transitions and isotopes are known to present maser emission \\citep[e.g.,][]{Schilke91,Hofn94,GM09}, but to date, the metastable $^{15}$NH$_3$ (3,3) maser has only been found in NGC 7538 IRS~1 \\citep[G91;][]{Mauer86,John89,WW96}. Because the original 23.1-GHz CH$_3$OH detection by W84 was never pursued at higher angular resolution, and because of the presence in IRS~1 of not one but possibly three rare maser species, we observed NGC 7538 IRS~1 with the goals of confirming the maser nature of the 23.1-GHz CH$_3$OH emission, and accurately locating this emission with respect to the other masers. In \\S ~2 we describe the observations and the data reduction procedure. We present the results in \\S ~3 and provide a discussion in \\S ~4. ", "conclusions": "\\label{sec:disc} \\subsection{Nature of the emission} \\label{sec:nature} W84 did not have sufficient angular resolution to place a useful limit on the source brightness temperature. Nevertheless, they argued that the emission was of maser origin, based on absorption in the CH$_3$OH~$10_1-9_2 \\, A^-$ line. This absorption implied that the excitation temperature of the CH$_3$OH~$9_2-10_1 \\, A^+$ emission line was greater than the background temperature of at least 1000~K. With our $250\\times$ higher angular resolution we calculate brightness temperatures of $T_N>1.3\\times10\\tothe5$ K and $T_S \\sim 4\\times10\\tothe4$ K for components $N$ and $S$, respectively. Such high brightness temperatures cannot reflect the kinetic temperature of the molecular gas, thus we prove the maser nature of the emission. The relatively broad linewidths may indicate that the masers are saturated; alternatively, it could indicate the presence of unresolved maser components. Three objects are now confirmed to harbor 23.1-GHz CH$_3$OH maser emission: NGC 7538 IRS 1 (this paper), W3(OH) \\citep{Menten85,Menten88}, and NGC 6334F \\citep{MB89}. \\subsection{Comparison with other masers} \\label{sec:corr} Figure 3 shows an overlay of the free-free continuum with the locations of the masers we discuss here. The 6.7-GHz and 12.2-GHz CH$_3$OH masers reported by M00 ({\\it right} panel) and the 12.2-GHz CH$_3$OH masers reported by M09 ({\\it left} panel) are plotted in separate panels for clarity. The positions of M00 were shifted by 18 mas to the west and south to account for the proper motions reported by M09 for CH$_3$OH masers. Although not shown in Figure 3, OH masers reported by Hutawarakorn \\& Cohen (2003) are located slightly south of the IRS~1 core, coincident with the more diffuse continuum emission seen at the bottom of Figure 3. The most striking feature is the general clustering of the many different maser species around IRS~1. The bright core of continuum emission from IRS~1 has linear dimensions of order 1000 AU. Scattered across this area are the various masers of class II CH$_3$OH, H$_2$O, NH$_3$, and H$_2$CO. There are three spatial coincidences at the present angular resolutions. One is the 23.1-GHz component $S$, which coincides with H$_2$O maser M8. Another is the 23.1-GHz component $N$, which coincides with component I of the 4.8-GHz formaldehyde maser (which is itself resolved into two components: Ia and Ib, separated by $\\sim 60$ AU; H03) and is also marginally associated with the H$_2$O maser M7. The third association is H$_2$O maser M6 with component II of the H$_2$CO masers. Table 3 lists the positions, offsets, and absolute positional uncertainty of the coincident masers. The 6.7 or 12.2-GHz CH$_3$OH spot closest to a 23.1-GHz maser is component 1 of \\cite{Mosca09}, with a separation of 65 mas (or 1.3$\\sigma$) from our component $N$. As we are confident with our estimates of the astrometric precision, we still rule out an association between these features. The velocities of the 23.1-GHz CH$_3$OH masers ($V_N\\approx-56.1$ $\\kms$, $V_S\\approx-59.1$ $\\kms$) are similar to those of the other masers in the vicinity of IRS~1. The 6.7-GHz CH$_3$OH masers have LSR velocities in the range $V_{6.7}=[-61.6,-55.0]$ $\\kms$, and the cross-power spectrum of those masers for the entire field shows two velocity components centered at $\\approx -56$ $\\kms$ and $\\approx -58$ $\\kms$ (M00). M00 also report several 12.2-GHz CH$_3$OH masers coincident with 6.7-GHz masers; the velocities of the former are $\\approx -58$ $\\kms$. M09 reports more CH$_3$OH masers at 12.2 GHz than does M00; all of the additional masers are in the velocity interval $[-61.9,-55.8]$ $\\kms$. The 4.8-GHz H$_2$CO masers also have two velocity components. The components Ia and Ib have almost the same LSR velocity ($V_{Ia} \\approx -57.9$ $\\kms$, $V_{Ib} \\approx -57.8$ $\\kms$), while component II is at $V_{II}\\approx -60.2$ $\\kms$ (H03). The $\\tothe{15}$NH$_3$ (3,3) masers have velocities in the range $V_{\\tothe{15}NH_3}=[-52.8,-61.7]$ $\\kms$ (G91). The H$_2$O masers reported here have a velocity range of $V_{H_2O}=[-53,-63]$ $\\kms$ (see Table 1). The typical errors in the determination of the velocity features are $\\sim 0.1$ $\\kms$ or better. For the positional associations the corresponding velocities are very similar. CH$_3$OH maser $S$ is centered at $-59.11$ $\\kms$, and H$_2$O maser M8 spans the range $[-60.8,-59.2]$ $\\kms$. CH$_3$OH maser $N$ is centered at $-56.09$ $\\kms$ while H$_2$CO masers Ia/Ib are at $-57.9/-57.8$ $\\kms$. The corresponding H$_2$O maser, M7, spans the range $[-58.2,-55.2]$ $\\kms$, which includes both the CH$_3$OH and H$_2$CO maser velocities. Hence, the maser counterparts coincide both spatially and kinematically. \\subsection{Comparison with models} \\label{sec:models} Although the 23.1-GHz methanol masers do coincide with water and formaldehyde masers, they {\\it do not} coincide with 6.7-GHz or 12.2-GHz methanol masers. Fig. 3 shows that this holds when the positions of the 23.1-GHz CH$_3$OH masers are compared to either the masers reported by M00 ($right$) or by M09 ($left$). The absolute position uncertainties are $\\sim 0\\rlap.{\"}01$ or better for M00, and $\\sim 0\\rlap.{\"}001$ for M09. It is not surprising that regions masing at 6.7 or 12.2 GHz might not show 23.1-GHz maser emission --- the models of \\citet{Cragg05} indicate that the former transitions mase over a wider range of physical parameters than the latter transition. But it {\\it is} suprising that 23.1-GHz masers would not be accompanied by 6.7 or 12.2-GHz masers: if conditions are adequate for 23.1-GHz masing, they should also be adequate for 6.7 and 12.2-GHz masing. Moreover, current models of methanol masers \\citep{Sobo97a,Sobo97b,Cragg04,Cragg05} predict brightness temperatures that are orders of magnitude higher for the 6.7 and 12.2-GHz masers compared to the 23.1-GHz maser. In light of current models, then, the puzzle is why the locations of the 23.1-GHz masers do not also show 6.7 or 12.2-GHz emission. A possible explanation is the non-simultaneity of the observations. However, the observations of M09 were made within days to months of ours, and the maser spots were consistent for all the epochs of M09 (Mark Reid, personal communication). Changing physical conditions as a function of time are to be expected in the dynamic and turbulent medium of star formation regions \\citep[] [and references therein]{BP07}. In fact, the turbulent medium may be intimately related to the production of the masers \\citep{Sobo98}. Most of these changes will occur on a dynamic timescale, however, which is typically orders of magnitude longer than astronomical monitoring. Much more rapid change can be produced in the radiation fields. Recent models of \\HII regions with accretion activity predict the existence of significant variations in the cm free-free flux on timescales as short as a few years \\citep{Peters10}. Fast variations in the radio flux have been reported for IRS1 by \\cite{Ramiro04}, which is one of the best candidates for an \\HII region still undergoing accretion \\citep{Klaass09,Sandell09,Zhu10}. Moreover, all of the three known 23.1-GHz CH$_3$OH maser regions have background free-free emission detected at radio wavelengths (Menten et al. 1988; Carral et al. 2002; this paper). The extended frequency coverage of the Expanded VLA (EVLA) will facilitate simultaneous multi-transition maser studies of this type of objects at high angular resolution. Finally, we underscore that if the $N$ 23.1-GHz maser and the 4.8-GHz H$_2$CO are close in space, then similar general conditions may give rise to both rare masers. \\cite{BdJ81} developed a pumping model for the IRS~1 H$_2$CO maser based on free-free emission from the \\HII region, and \\cite{Pratap92} concur that this model is adequate to explain the 4.8-GHz maser in IRS~1. Nevertheless, the Boland \\& de Jong model also predicts a bright H$_2$CO maser at 14.5~GHz, and this maser transition was not detected by \\cite{Hoff03}. Furthermore, \\cite{Araya07} showed that this model is not able to explain the H$_2$CO maser in another region of high-mass star formation: IRAS 18566+0408. We note that the consideration of non-simultaneous observations applies equally to these coincident masers as it did to the non-coincident methanol masers mentioned above. This caveat notwithstanding, the possible correlation of these two rare maser species may prove helpful in better defining the pumping mechanism of each." }, "1003/1003.3017_arXiv.txt": { "abstract": "A systematic investigation of the relationship between different redshift estimation schemes for more than 91\\,000 quasars in the Sloan Digital Sky Survey (SDSS) Data Release 6 (DR6) is presented. The publicly available SDSS quasar redshifts are shown to possess systematic biases of $\\Delta z/(1+z)$$\\ge$0.002 (600\\kms) over both small ($\\delta z$$\\simeq$0.1) and large ($\\delta z$$\\simeq$1) redshift intervals. Empirical relationships between redshifts based on i) \\caii H \\& K host galaxy absorption, ii) quasar \\oii $\\lambda$3728, iii) \\oiii $\\lambda\\lambda$4960,5008 emission, and iv) cross-correlation (with a master quasar template) that includes, at increasing quasar redshift, the prominent \\mgii $\\lambda$2799, \\ciii $\\lambda$1908 and \\civ $\\lambda$1549 emission lines, are established as a function of quasar redshift and luminosity. New redshifts in the resulting catalogue possess systematic biases a factor of $\\simeq$20 lower compared to the SDSS redshift values; systematic effects are reduced to the level of $\\Delta z/(1+z)$$\\le$10$^{-4}$ (30\\kms) per unit redshift, or $\\le$2.5$\\times$10$^{-5}$ per unit absolute magnitude. Redshift errors, including components due both to internal reproducibility and the intrinsic quasar-to-quasar variation among the population, are available for all quasars in the catalogue. The improved redshifts and their associated errors have wide applicability in areas such as quasar absorption outflows, quasar clustering, quasar-galaxy clustering and proximity-effect determinations. ", "introduction": "\\label{sec:intro} The Sloan Digital Sky Survey (SDSS) \\citep{2000AJ....120.1579Y} has produced a revolution in both the volume and quality of spectroscopic data available for quasars. The Data Release 5 (DR5) \\citep{2007ApJS..172..634A} and Legacy Data Release 7 (DR7) \\citep{2009ApJS..182..543A} with their associated quasar catalogues \\citep[respectively]{2007AJ....134..102S, 2010AJ....XXX..XXXX} provide intermediate resolution ($R$$\\sim$2000), moderate signal-to-noise ratio (SNR) (SNR$\\sim$15 per 69\\kms pixel), spectra of unprecedented homogeneity, covering essentially the entire ''optical'' wavelength region ($\\lambda$=3800--9180\\,\\AA). The quality of the Schneider et al. quasar catalogues is truly impressive, with errors in redshift identification reduced to the 0.01 per cent level and individual redshift estimates, resulting primarily from the SDSS spectroscopic pipeline \\citep[and the SDSS DR7 website\\footnote{http://www.sdss.org/dr7/algorithms/redshift\\_type.html}]{2002AJ....123..485S}, are accurate to of order $\\Delta z/(1+z)$$\\sim$0.002. The publication of even individual quasar redshifts, based on moderate resolution spectra, to such accuracy was a significant achievement prior to the mid-1990s, further highlighting the advance represented by the SDSS. Notwithstanding the quality of the SDSS quasar spectra and the associated redshift estimates, important scientific investigations, including the clustering of quasars themselves \\citep[e.g.][]{2002MNRAS.335..459C, 2007AJ....133.2222S}, the cross-correlation of quasars and other object populations \\citep[e.g.][]{2009MNRAS.397.1862P}, the proximity effect \\citep[e.g.][]{1988ApJ...327..570B, 2008MNRAS.391.1457K}, the origin and properties of associated absorbers \\citep[e.g.][]{2008MNRAS.386.2055N, 2008MNRAS.388..227W, 2009MNRAS.392.1539T} benefit significantly both from reduced systematics in redshift determinations and the reliable assignment of redshift uncertainties for individual quasars. In this paper we present the determination of new redshifts and associated error estimates for more than 89\\,500 quasars from the SDSS DR6 \\citep{2008ApJS..175..297A}. Our redshift determinations suffer from much smaller systematic effects compared to the default values from the SDSS spectroscopic pipeline. Specifically, systematics are reduced by more than an order of magnitude to 1.0$\\times$10$^{-4}$ in $\\Delta z$/(1+$z$) per unit redshift, or, equivalently, 30\\kms per unit redshift\\footnote{The quasar research community has normally quantified redshift errors in terms of $\\Delta z$/(1+$z$), whereas researchers studying galaxies conventionally specify uncertainties in kilometres per second. The improvements possible in redshift determination made possible by the SDSS spectra are such that the `kilometres per second' parameterisation is increasingly attractive and we specify the main results using both schemes.}. A detailed comparison of redshifts derived from \\caii H\\&K absorption, \\oii $\\lambda\\lambda$3727,3729 emission, \\oiii $\\lambda\\lambda$4960,5008 emission and cross-correlation with a new quasar template spectrum, provides greatly improved error estimates for individual quasar redshifts. The error estimates incorporate both the uncertainties resulting from the properties of the SDSS spectra, quantified using the very large number of multiple spectra present in the SDSS, and the intrinsic quasar-to-quasar dispersion. The resulting catalogue will allow significant advances in many studies that rely on the determination of systemic quasar redshifts with small systematics and well-determined uncertainties. The paper is structured as follows. Section 2 describes the quasar sample, before the features of the quasar redshifts available from the SDSS spectroscopic pipeline are illustrated in Section 3. Section 4 includes a description of the procedures involved in generating a master quasar template for the cross-correlation redshift estimates. Section 5 then describes the procedures employed to provide the new redshift estimates for the SDSS quasars. An assessment of the consistency of the different redshift estimates is given in Section 6 and redshift estimates, based on different rest-frame wavelength regions, are placed onto the same `systemic' reference system. A critical assessment of the internal and external reliability of the new quasar redshifts is presented at this point. Twenty-one centimetre radio observations of the majority of the SDSS quasars are available from the Faint Images of the Radio Sky at Twenty centimetres \\citep[FIRST,][]{1995ApJ...450..559B}. Section 7 contains a description showing how the new redshift estimation scheme allows spectral energy distribution (SED) dependent composite spectra (for FIRST-detected quasars in this case) to be constructed, producing significantly improved redshifts. The resulting redshift catalogue, including well-determined error estimates for each quasar, is described in Section 8. A short discussion, including consideration of the origin of the differences with published redshifts and an independent test of the new redshifts follows in Section 9. The paper concludes with a brief summary of the conclusions as Section 10. We adopt the same convention as employed in the SDSS and use vacuum wavelengths throughout the paper. Absolute magnitudes are calculated in a cosmology with $H_0$=70\\kms, $\\Omega_M$=0.3 and $\\Omega_\\Lambda$=0.7. ", "conclusions": "A systematic investigation of the relationship between different redshift estimation schemes for more than 91\\,000 quasars in the Sloan Digital Sky Survey (SDSS) Data Release 6 (DR6) is presented. Empirical relationships between redshifts based on i) \\caii H \\& K host galaxy absorption, ii) quasar \\oii $\\lambda$3728, iii) \\oiii $\\lambda\\lambda$4960,5008 emission, and iv) cross-correlation (with a master quasar template) that includes, at increasing quasar redshift, the prominent \\mgii $\\lambda$2799, \\ciii $\\lambda$1908 and \\civ $\\lambda$1549 emission lines, are established as a function of quasar redshift and luminosity. New redshifts in the resulting catalogue possess systematic biases a factor of $\\simeq$20 lower compared to the SDSS redshift values; systematic effects are reduced to the level of $\\Delta z/(1+z)$$\\le$10$^{-4}$ (30\\kms) per unit redshift, or $\\le$2.5$\\times$10$^{-5}$ per unit absolute magnitude. It is important to realise that there will be systematic redshift trends present as a function of the quasar SEDs and the specific example of FIRST-detected quasars (Section~\\ref{radio_qsos}) provides an example, related to the radio-properties of the quasar SEDs. One of the primary motivations of this work is to facilitate further studies of SED-dependent systematic emission line properties, working from redshift estimates whose properties as a function of redshift and absolute magnitude are well understood. Equally important as the new redshift determinations, well-determined empirical estimates of the quasar-to-quasar dispersion in redshifts are available for each method of redshift estimation and a combined internal+population uncertainty is provided for every quasar in the catalogue. The improved redshifts and their associated errors have wide applicability in areas such as quasar absorption outflows, quasar clustering, quasar-galaxy clustering and proximity-effect determinations." }, "1003/1003.5541_arXiv.txt": { "abstract": "\\noindent Source confusion defines a practical depth to which to take large-area extragalactic surveys. 3D imaging spectrometers with positional as well as spectral information, however, potentially provide a means by which to use line emission to break the traditional confusion limit. In this paper we present the results of our investigation into the effectiveness of mid/far infrared, wide-area spectroscopic surveys in breaking the confusion limit. We use SAFARI, a FIR imaging Fourier Transform Spectrometer concept for the proposed JAXA-led SPICA mission, as a test case. We generate artificial skies representative of 100 SAFARI footprints and use a fully-automated redshift determination method to retrieve redshifts for both spatially and spectrally confused sources for bright-end and burst mode galaxy evolution models. We find we are able to retrieve accurate redshifts for 38/54\\% of the brightest spectrally confused sources, with continuum fluxes as much as an order of magnitude below the 120 $\\mu$m photometric confusion limit. In addition we also recover accurate redshifts for 38/29\\% of the second brightest spectrally confused sources. Our results suggest that deep, spectral line surveys with SAFARI can break the traditional photometric confusion limit, and will also not only resolve, but provide redshifts for, a large number of previously inaccessible galaxies. To conclude we discuss some of the limitations of the technique, as well as further work. ", "introduction": "The cosmic infrared background (CIB) peaks at $\\sim 150$ $\\mu$m and comprises the total infrared (IR) emission from all sources in the sky (eg. \\cite{dole-01}, \\cite{elbaz-02}), integrated over all time. It has been found to contain as much energy as the combined optical/UV extragalactic background, suggesting that half of all light emitted by stars and active galactic nuclei (AGN) is absorbed by dust before we are able to observe it in the optical \\citep{hauser-dwek}. Locally, the IR output of typical galaxies is only one third of their optical output \\citep{soifer-neugebauer-91}, which implies strong evolution in the IR properties of galaxies as one moves to high redshift. The CIB has been well-studied using many instruments including ISOCAM (ISO), MIPS (Spitzer), ISOPHOT (ISO) and SCUBA(JCMT) at 15, 24, 160 and 850 $\\mu$m (\\cite{elbaz-02}, \\cite{papovich-04}, \\cite{juvela-00} and \\cite{smail-02} respectively). Although the peak of the CIB lies at $\\sim150$ $\\mu$m, it has yet to be resolved into individual sources at these wavelengths: our understanding of the make-up of the CIB therefore relies on extrapolation. Observations in the mid and far-infrared (MIR and FIR; typically defined as lying in the wavebands 5-30 and 30-1000 $\\mu$m respectively) from ground based telescopes are difficult if not impossible because of the high opacity of the Earth's atmosphere at these wavelengths. Ground based telescopes are only sensitive to narrow wavebands in the MIR/FIR where the atmosphere's transmission is higher, thus wide band observations in the MIR/FIR must typically be made from space. As a result, FIR telescopes are smaller in diameter than their optical counterparts (up to a few meters), with a resulting angular resolution that is low compared to both optical and radio telescopes/interferometers. A direct consequence of this low angular resolution is that completely resolving the CIB into discrete sources in the FIR is very difficult, if not impossible, because of source confusion. Source confusion may be defined as the degradation of the quality of photometry of sources clustered on a scale to the order of the telescope beam size (eg. \\cite{scheuer-57}). The confusion limit sets the useful depth to which large-area extra-galactic surveys should be taken. For example, the Herschel mission \\citep{pilbratt-04}, successfully launched in May 2009, has a 3.5 m diameter mirror which realizes an angular resolution of 8'' at 120 $\\mu$m. At these wavelengths, the confusion limit for such a mirror is estimated to be around $\\sim$5 mJy (eg. \\cite{dole-04}, \\cite{jeong-06}). Surveys at 24 $\\mu$m suggest that at these flux levels one will only be able to resolve at most $\\sim$~50\\% of the CIB \\citep{dole-04}. It is possible to reduce the confusion limit through making observations with a larger diameter mirror. However, due to the practical limitations of high angular resolution FIR imaging there is a limit on how much we can reduce confusion noise. One way to break the confusion limit makes use of the extra dimension of wavelength, to which one has access in spectroscopic surveys. Discrete sources can be identified by relatively bright, narrow-band emission lines: thereby allowing redshifts to be determined. A preliminary study to explore the efficacy of blind, wide area spectroscopic surveys in resolving FIR sources is described in \\citet{clements-07}. In their work an artificial `sky' was populated using template FIR spectral energy distributions (SEDs) of a selection of different types of galaxy, to which were added FIR emission lines of strengths derived from ISO-LWS observations (eg. \\cite{negishi-01}). The sources were redshifted and assigned luminosities according to the evolutionary models of \\citet{pearson-05}, \\citet{pearson-07} and \\citet{pearson-k-09}. Observations of the `sky' were made using the instrumental parameters (eg. sensitivity/noise levels, spectral resolution, field of view (FoV), beam size) of SAFARI, a FIR imaging Fourier Transform Spectrometer concept for the proposed JAXA-led SPICA (Space Infrared Telescope for Astronomy and Astrophysics) \\citep{swinyard-08} mission. It will offer the large FoV and high spectral resolution required to break the confusion limit using spectroscopy. According to its current specifications SPICA will have a 3.5 m diameter mirror, and therefore will be subject to the same confusion noise as Herschel. The primary mirror will, however, be cooled to $<$6 K and so will offer a great leap in sensitivity over Herschel. SAFARI will cover the waveband 35 to 210 $\\mu$m with varying resolution, including R$\\sim$1000 - (at 120 $\\mu$m, $\\Delta\\lambda = 0.176$ $\\mu$m, when run in SAFARI's higher resolution mode) - which matches the typical width of an extragalactic MIR/FIR emission line. Estimates of source redshift were made by hand by locating the position of the strongest emission line in each spectrum. The strongest lines typically observed in the FIR are the [OI] and [CII] lines at 63.18 and 157.74 $\\mu$m respectively. If one assumes a typical dust temperature of 35 K, then the strongest lines shortward and longward of the peak of the SED will be [OI] and [CII] respectively. Beyond $z=2.5$ these lines are shifted out of SAFARI's observable waveband, therefore the redshifts of more distant sources than this are irretrievable. If these two lines were the only ones present in the FIR then by comparing the evaluated source redshifts with the model input redshifts, one can assess the efficiency of this blind-line method. It was found that when looking at a patch of simulated `sky' equal to one SAFARI FoV, it was possible to retrieve accurate redshifts for sources with 120 $\\mu$m continuum fluxes as much as a factor of $\\sim$10 below the traditional continuum confusion limit. Sources with 120 $\\mu$m flux $S_{120\\mu m} > 1$ mJy and at redshifts $z < 2.5$ were retrieved with 100\\% accuracy. The use of blind spectral line surveys to resolve FIR sources is not without its own limitations and type of confusion. Line, or spectral confusion occurs when multiple sources are observed in a single telescope beam: the spectra from two or more objects are effectively scrambled, and it can become difficult to determine which lines are emitted by which objects. As a result, source redshifts become hard to extract. The work described in \\citet{clements-07} made use of model spectra with FIR emission lines only. To assess the true viability of using spectral line surveys to break the confusion limit requires the inclusion of MIR emission lines in the `sky' model, as sources will, in general, have both FIR and MIR emission lines. Inclusion of these shorter wavelength lines will enable the recovery of sources with redshifts of $z > 2.5$, beyond which the [OI] and [CII] emission lines at 63.18 and 157.74 $\\mu$m, respectively, are shifted out of the SAFARI waveband. By including MIR lines however, one increases the problem of line confusion, and so assigning lines to individual, but spatially unresolved, sources becomes more problematic. In this paper we examine a much larger model `sky' populated with more realistic template spectra with both FIR and MIR emission lines and employ a new automated method of evaluating source redshifts in a time efficient manner. We also implement a method to extract the redshifts of multiple sources clustered in a single spatial bin. Through the implementation of this method we investigate how effectively we can break the traditional photometric confusion limit. In sections~\\ref{sec:generate_sky} and~\\ref{sec:generate_cube} we describe the generation of the artificial sky used to test the source recovery technique, which in turn is outlined in section~\\ref{sec:method}. This is followed in section~\\ref{sec:results} by a quantitative assessment of the retrieval and error rates of the redshifts output, and a discussion of the results in section~\\ref{sec:discussion}. ", "conclusions": "We have found that our PCC redshift determination method is capable of resolving sources (i.e. determining a unique redshift) more than an order of magnitude fainter than the traditional continuum confusion limit, however the efficacy of our method is higher for brighter sources. In this work we have used the PCC method on models based around the SAFARI instrument for SPICA, however the same technique could be used on any sensitive imaging spectrometer. The bright-end and burst mode evolution models include sources up to redshifts of $z\\sim$4 and 5 respectively. At these redshifts we are still able to determine redshifts for sources using the PCC method. We have not yet tested to see at which redshift the PCC method begins to fail, this may be the subject of future work. The evolutionary models we have investigated in this work have the bulk of their populations at fluxes fainter than the confusion limit. By employing our PCC method we are therefore greatly increasing the number of galaxies that we are able to uniquely identify in any single observation. This presents us with a better statistical sample with which to compare observed source counts and redshift distributions with those presented in different evolutionary models. We therefore have the potential to reliably distinguish different evolutionary models and our observations with much smaller area surveys, and therefore within shorter observing times. Future work should include the following: 1) Use of the PAH features to identify sources and determine their redshifts. This will allow us to decrease our spectral resolution, thus increasing instrument sensitivity. 2) Investigation of the viability of taking into account sources with atypical line strengths. 3) More realistic angular resolution modeling where spatial resolution varies across the waveband. 4) Investigation of the efficiency of the method when implemented on a wider range of evolutionary models. 5) A more quantitative analysis of where our ability to retrieve redshifts from single and combined spectra begins to break down is being conducted. 6) It should also be noted that since the work described in this paper was conducted the technical specifications of SAFARI have changed somewhat (eg. waveband, sensitivity), therefore the results should be re-checked with more up to date modeling of the SAFARI instrument. These topics are currently being investigated. \\begin{flushleft} {\\bf Acknowledgements} \\end{flushleft} We would like to thank the anonymous referee for excellent comments and suggestions which markedly improved the clarity of the paper. GR would like to acknowledge an STFC postgraduate studentship, KGI funding from RCUK and DC was funded in part by STFC." }, "1003/1003.0687_arXiv.txt": { "abstract": "Many of the cosmological tests to be performed by planned dark energy experiments will require extremely well-characterized photometric redshift measurements. Current estimates for cosmic shear are that the true mean redshift of the objects in each photo-z bin must be known to better than $0.002(1+z)$, and the width of the bin must be known to $\\sim0.003(1+z)$ if errors in cosmological measurements are not to be degraded significantly. A conventional approach is to calibrate these photometric redshifts with large sets of spectroscopic redshifts. However, at the depths probed by Stage III surveys (such as DES), let alone Stage IV (LSST, JDEM, Euclid), existing large redshift samples have all been highly (25-60\\%) incomplete, with a strong dependence of success rate on both redshift and galaxy properties. A powerful alternative approach is to exploit the clustering of galaxies to perform photometric redshift calibrations. Measuring the two-point angular cross\\hyp{}correlation between objects in some photometric redshift bin and objects with known spectroscopic redshift, as a function of the spectroscopic z, allows the true redshift distribution of a photometric sample to be reconstructed in detail, even if it includes objects too faint for spectroscopy or if spectroscopic samples are highly incomplete. We test this technique using mock DEEP2 Galaxy Redshift survey light cones constructed from the Millennium Simulation semi-analytic galaxy catalogs. From this realistic test, which incorporates the effects of galaxy bias evolution and cosmic variance, we find that the true redshift distribution of a photometric sample can, in fact, be determined accurately with cross\\hyp{}correlation techniques. We also compare the empirical error in the reconstruction of redshift distributions to previous analytic predictions, finding that additional components must be included in error budgets to match the simulation results. This extra error contribution is small for surveys which sample large areas of sky ($> \\sim $10-100 degrees), but dominant for $\\sim 1$ square degree fields. We conclude by presenting a step-by-step, optimized recipe for reconstructing redshift distributions from cross\\hyp{}correlation information using standard correlation measurements. ", "introduction": "\\label{sec:intro} For many years it was thought that the expansion of the universe should be slowing due to the gravitational attraction of matter, but measurements of Type Ia supernovae and other observations have shown that the expansion rate is in fact accelerating \\citep{1998AJ....116.1009R,1999ApJ...517..565P}. This accelerating expansion is generally attributed to an unknown component of the energy density of the universe commonly referred to as ``dark energy.'' One of the goals of future cosmological probes (e.g. LSST, JDEM, and Euclid) \\citep{2001ASPC..232..347T, 2005ASPC..339...95T,2009arXiv0901.0721A, 2010arXiv1001.3349B}, is to determine constraints on dark energy equation of state parameters, e.g. $w \\equiv P/\\rho$ and $\\mathrm{w_a} \\equiv dw/da$ \\citep{2007IJMPD..16.1581J}, where $P$ is the pressure from dark energy, $\\rho$ is its mass density, and $a$ is the scale factor of the universe (normalized to be 1 today). In order for these experiments to be successful, we require information about the redshift of all objects used to make measurements. However, it is impractical to measure spectroscopic redshifts for hundreds of millions of galaxies, especially extremely faint ones. We can measure the redshift of many more objects from photometric information, e.g. by using a large set of spectroscopic redshifts to create templates of how color varies with redshift \\citep{1995AJ....110.2655C}. However current and future spectroscopic surveys will be highly incomplete due to selection biases dependent on redshift and galaxy properties \\citep{2006MNRAS.370..198C}. Because of this, along with the catastrophic photometric errors\\footnote{such as contamination from overlapping or unresolved objects; this is a frequent problem in deep surveys, particularly at high redshifts, cf. Newman et al. 2010} that can occur at a significant ($\\sim 1\\%$) rate \\citep{2009ApJ...699..958S, 2010MNRAS.401.1399B}, photometric redshifts are not as well understood as redshifts determined spectroscopically. If future dark energy experiments are to reach their goals, it is necessary to develop a method of calibrating photometric redshifts with high precision \\citep{2006astro.ph..9591A, 2006MNRAS.366..101H, 2006ApJ...636...21M}. Current projections for LSST cosmic shear measurements estimate that the true mean redshift of objects in each photo-z bin must be known to better than $\\sim0.002(1+z)$ \\citep{2006ApJ...644..663Z, 2006JCAP...08..008Z, 2006ApJ...652..857K, 2006AIPC..870...44T} with stringent requirements on the fraction of unconstrained catastrophic outliers \\citep{2010arXiv1002.3383H}, while the width of the bin must be known to $\\sim0.003(1+z)$ \\citep{2009arXiv0912.0201L}. In this paper we test a new technique for calibrating photometric redshifts measured by other algorithms, which exploits the fact that objects at similar redshifts tend to cluster with each other. If we have two galaxy samples, one with only photometric information and the other consisting of objects with known spectroscopic redshifts, we can measure the angular cross\\hyp{}correlation between objects in the photometric sample and the spectroscopic sample as a function of spectroscopic $z$. This clustering will depend on both the intrinsic clustering of the samples with each other and the degree to which the samples overlap in redshift. Autocorrelation measurements for each sample give information about their intrinsic clustering, which can be used to break the degeneracy between these two contributions. The principal advantage of this technique is that, while the two sets of objects should overlap in redshift and on the sky, it is not necessary for the spectroscopic sample to be complete at any given redshift. Therefore it is possible to use only the brightest objects at a given $z$, from which it is much easier to obtain secure redshift measurements, to calibrate photometric redshifts. Even systematic incompleteness (e.g. failing to obtain redshifts for galaxies of specific types) in the spectroscopic sample is not a problem, so long as the full redshift range is sampled. This method is effective even when the two samples do not have similar properties (e.g. differing luminosity and bias). We here describe a complete end-to-end implementation of cross\\hyp{}correlation methods for calibrating photometric redshifts and present the results of applying these algorithms to realistic mock catalogs. Throughout the paper we assume a flat $\\Lambda$CDM cosmology with $\\Omega_m$=0.3, $\\Omega_{\\Lambda}$=0.7, and Hubble parameter $H_0=100h$ km s$^{-1}$ Mpc$^{-1}$, where we have assumed $h$=0.72, matching the Millennium simulations, where it is not explicitly included in formulae. In \\S \\ref{sec:datasets} we describe the catalog and data sets used to test cross\\hyp{}correlation methods. In \\S \\ref{sec:method} we provide a description of the reconstruction techniques used in detail, and in \\S \\ref{sec:results} we provide the results of the calculation. In \\S \\ref{sec:conclusion} we conclude, as well as give a more concise description of the steps taken, providing a recipe for cross\\hyp{}correlation photometric redshift calibration. ", "conclusions": "\\label{sec:conclusion} Section \\ref{sec:method} has described in detail the steps we took to recover the redshift distribution, $\\phi_p(z)$, of a photometric sample by cross\\hyp{}correlating with a spectroscopic sample of known redshift distribution. We will now summarize the procedure used to make this calculation, to facilitate its application to actual data sets. \\begin{list}{\\labelitemi}{\\leftmargin=1em} \\item \\textbf{Obtain the necessary information for each sample; RA, dec and redshift for the spectroscopic sample, and RA and dec for the photometric sample.} \\item \\textbf{Create the random catalogs for each sample. (\\S \\ref{sec:autospec}-\\ref{sec:crossphi})} \\item \\textbf{Calculate the data-data, data-random, and random-random paircounts for each correlation function.} \\begin{list}{\\labelitemi}{\\leftmargin=1em} \\item For $w_p(r_p)$: bin the spectroscopic sample and its corresponding random catalog in redshift. In each spectroscopic z-bin, calculate $\\Delta r_p$ and $\\Delta \\pi$ for each pair and bin the pair separations into a grid of $\\log(r_p)$ and $\\pi$. Then sum the paircounts in the $\\pi$ direction. (\\S \\ref{sec:autospec}) \\item For $w_{pp}(\\theta)$: using the '$p$' sample and its random catalog, calculate $\\Delta \\theta$ for each pair and bin the pair separations into log($\\theta$) bins. (\\S \\ref{sec:autophot}) \\item For $w_{sp}(\\theta,z)$: bin the spectroscopic sample and its corresponding random catalog in redshift. For each spectroscopic z-bin, calculate the pair separations, $\\Delta \\theta$, for pairs between the '$s$' and '$p$' samples and their random catalogs and bin them into $\\log(\\theta)$ bins. (\\S \\ref{sec:crossphi}) \\end{list} \\item \\textbf{Use the paircounts to calculate the correlation functions using standard estimators (e.g. Landy \\& Szalay). (\\S \\ref{sec:autospec}-\\ref{sec:crossphi})} \\item \\textbf{Calculate the parameters of $w_p(r_p)\\ (r_{0,ss}(z), \\gamma_{ss}(z))$ and $w_{pp}(\\theta)\\ (A_{pp}, \\gamma_{pp})$ by fitting as described above. (\\S \\ref{sec:autospec}-\\ref{sec:autophot})} \\item \\textbf{Use the autocorrelation parameters along with an initial guess of $r_{0,pp}$ (e.g. $r_{0,pp}\\sim r_{0,ss}$) to calculate $r^{\\gamma_{sp}}_{0,sp}(z)=(r^{\\gamma_{ss}}_{0,ss}r^{\\gamma_{pp}}_{0,pp})^{1/2}$. (\\S \\ref{sec:autophot})} This gave a more accurate reconstruction of $\\phi_p(z)$ (reducing $\\chi^2$ by 33\\%) than the assumption $r_{0,pp}=$ constant; in fact, a calculation of $\\xi_{pp}(r)$ from the simulation sample directly showed $r_{0,pp}$ to have similar behavior to $r_{0,ss}$. Using a linear fit of $r_{0,ss}(z)$ and $\\gamma_{ss}(z)$ reduced $\\chi^2$ by $\\sim32\\%$ compared to utilizing the noisier reconstructed values in each z-bin. \\item \\textbf{Estimate $\\gamma_{sp}=(\\gamma_{ss}+\\gamma_{pp})/2$}. Using this $\\gamma_{sp}$, calculate the amplitude, $A_{sp}(z)$, of $w_{sp}(\\theta,z)$ by fitting as described above. (\\S \\ref{sec:crossphi}) We fit over the range $0.001^{\\circ} <\\theta< 0.1^{\\circ}$. We found that fitting over this smaller $\\theta$ range resulted in smaller errors in the amplitude, $A_{sp}(z)$, which reduced the error in $\\phi_p(z)$ for each z-bin by $\\sim25\\%$ on average. We fix $\\gamma_{sp}$ because of degeneracies between $\\gamma_{sp}$ and $A_{sp}$ when fitting them simultaneously. This degeneracy is especially strong in regions where $\\phi_p(z)$ is small. We also tried modeling $\\gamma_{sp}$ as constant with z using the arithmetic mean of $\\gamma_{ss}(z=0.77)$ and $\\gamma_{pp}$; however, that method increased the $\\chi^2$ of the final fit by $\\sim20\\%$. \\item \\textbf{Combining the results of the last two steps and the assumed cosmology, calculate $\\phi_p(z)$ using equation \\ref{eq:phi}. (\\S \\ref{sec:crossphi})} We also tried calculating $\\phi_p(z)$ using the integrated cross\\hyp{}correlation function, $\\tilde w(z)$, integrating to an angle equivalent to a comoving distance $r_{max}=10h^{-1}$ Mpc (Newman 2008); however, that method produced inferior results. \\item \\textbf{Using $\\phi_p(z)$, along with the calculated $A_{pp}$ and $\\gamma_{pp}$, in equation \\ref{eq:wpp} gives a new $r_{0,pp}$, which is then used to recalculate $r^{\\gamma_{sp}}_{0,sp}(z)$. Putting this back into equation \\ref{eq:phi} gives a new $\\phi_p(z)$. This is repeated until convergence is reached. (\\S \\ref{sec:crossphi})} \\item \\textbf{To recover the underlying/universal distribution of objects of the type selected for the photometric sample, rather than the distribution within the specific fields chosen for observation, correct for sample/cosmic variance using the fluctuations in the redshift distribution of the spectroscopic; i.e., construct a smooth function describing the overall redshift distribution of the spectroscopic sample, $\\langle dN_s/dz\\rangle$, and divide $\\phi_p(z)$ by the ratio $dN_s/dz/\\langle dN_s/dz\\rangle$. (\\S \\ref{sec:errorest})} \\end{list} We have shown in this paper that by exploiting the clustering of galaxies at similar redshifts we can accurately recover the redshift distribution of a photometric sample using its angular cross\\hyp{}correlation with a spectroscopic sample of known redshift distribution, using mock catalogs designed to match the DEEP2 Galaxy Redshift Survey. This test includes the impact of realistic bias evolution and cosmic variance. Our error estimates for the recovered mean and standard deviation of the distribution are larger than those predicted previously, but improvements could be obtained either by using more optimal correlation function estimators or by surveying the same number of galaxies distributed over a wider area of sky. Based on these tests we expect that this technique should be able to deliver the performance needed for dark energy experiments. In a recent paper \\citep{2009arXiv0910.3683S}, cross\\hyp{}correlation techniques were applied to mock data generated by populating a single time slice of an N-body dark matter simulation using various halo models. They develop a pipeline for calculating the redshift distribution of a photometric sample using cross\\hyp{}correlation measurements and the autocorrelation of a spectroscopic sample, $\\xi_{ss}(r,z)$. They do not attempt to model the bias although they do examine how varying the bias of the two samples affects the reconstruction (i.e. using radically different halo models). The catalogs constructed to test their method are significantly larger in volume than our individual mock catalogs, and while the number of objects in their photometric sample is comparable to ours, their spectroscopic sample is much smaller, which would be expected to lead to larger errors (Newman 2008), as observed. Another major difference is the use of a smoothness prior in reconstruction, which was not done here. While \\citet{2009arXiv0910.3683S} found that cross\\hyp{}correlation techniques were generally successful in reconstructing redshift distributions, these conclusions were primarily qualitative due to the limited sample sizes and source densities of the mock samples used, along with less-optimal correlation measurement techniques. In this paper, we have used simulations which include much less massive halos, allowing us to perform quantitative tests of cross\\hyp{}correlation techniques using sample sizes and source densities comparable to those which will be used in realistic applications. Several techniques for calibrating photometric redshifts using only photometric data have also been developed \\citep{2006ApJ...651...14S, 2009arXiv0910.4181Z, 2010arXiv1002.2266B, 2009arXiv0910.2704Q}; in general, such techniques require priors or assumptions on biasing which can be relaxed or tested in spectroscopic cross\\hyp{}correlation measurements. In \\cite{2009arXiv0910.2704Q}, spectroscopic/photometric cross\\hyp{}correlation techniques have now been applied to real data using the COSMOS dataset. Using data from a single field, they are able to determine typical photo-z uncertainties well, even when ignoring the effects of bias evolution. However, when constraining catastrophic photo-z errors, methods which ignore these effects should break down, as bias evolution should be a much greater problem over broad redshift intervals than in the core of the photo-z error distribution. In future work, we will explore alternate methods of measuring correlation functions that are invariant to the variance in the integral constraint (e.g. \\cite{2007MNRAS.376.1702P}). This should reduce errors in the measurement of the redshift distribution, which we found to be larger than expected due to extra variance terms in the correlation function measurements not considered previously. We also plan to test this technique with mock catalogs in which photometric redshifts have been 'measured' on simulated LSST photometry, rather than simply assuming a redshift distribution. We will also apply this method to real data using photometric and spectroscopic samples from the AEGIS survey \\citep{2007ApJ...660L...1D}. The authors wish particularly to thank Darren Croton for developing the mock catalogs used and making them publicly available. We would also like to thank Andrew Hearin, Arthur Kosowsky, David Wittman, Michael Wood-Vasey, Andrew Zentner, Tony Tyson, Gary Bernstein, Nikhil Padmanabhan, Dragan Huterer, Hu Zhan, and in particular Alexia Schulz for useful discussions during the course of this work, and also Ben Brown and Brian Cherinka for their technical expertise in performing these calculations. This work has been supported by the United States Department of Energy." }, "1003/1003.0825_arXiv.txt": { "abstract": "The particle distribution function that describes two interpenetrating plasma streams is re-investigated. It is shown how, based on the Maxwell-Boltzmann-J\\\"uttner distribution function that has been derived almost a century ago, a counterstreaming distribution function can be derived that uses velocity space. Such is necessary for various analytical calculations and numerical simulations that are reliant on velocity coordinates rather than momentum space. The application to the electrostatic two-stream instability illustrates the differences caused by the use of the relativistic distribution function. ", "introduction": "Plasma physics---both MHD (magnetohydrodynamics) and kinetic theory---are based on the knowledge of a distribution function that provides statistics about the average velocity direction (bulk flow) and the deviation from that mean value (known as temperature). Whereas MHD calculations are based on the use of a Maxwellian velocity distribution, such is not the case in kinetic theory, where the distribution function can be calculated using, e.\\,g., the Vlasov equation. Kinetic theory \\citep[e.\\,g.,][]{rs:rays} is mostly used when dilute plasmas are considered that do not satisfy the condition of frequent binary particle collisions so that no Maxwellian velocity distribution is established. In various and already historical work, the classic Maxwellian velocity distribution has been generalized to a relativistic gas (see, e.\\,g., \\citealt{jut11:dyn}; \\citealt{jut11:max}; \\citealt{syn57:gas}; \\citealt{fra79:col}). The result was the MBJ (Maxwell-Boltzmann-J\\\"uttner) distribution function (\\citealt{jut11:dyn}; \\citealt{jut11:max}), which is essentially given through \\be f\\propto\\exp\\!\\left(-\\alpha\\sqrt{1+\\f p^2}\\right) \\ee where $\\alpha$ is a temperature-related parameter. However, there are cases when the use of a distribution function in momentum space is (i) not suitable for the analytical calculations at hand; and/or (ii) not implemented in the numerical (simulation) code. A recent example is the investigation of LIDAR (LIght Detection And Ranging) Thomson scattering systems in ITER (International Thermonuclear Experimental Reactor) plasmas \\citep{bea08:tho}. The calculation was based on a formula \\citep{hut87:pla} for the scattered power per unit solid per unit angular frequency, and thus the particle distribution function had to be expressed in terms of velocity variables rather than momentum variables. Another example is the PIC (Particle-In-Cell) simulation code \\textsc{Tristan} (originally \\citealt{bun93:pic}, see also, e.\\,g., \\citealt{sak04:mag}), which is used for the investigation of plasma instabilities in the context of astrophysical scenarios such as the generation of magnetic fields at shock wave sites. A basic example of such instabilities is the generalized filamentation (or, originally, Weibel) instability (\\citealt{wei59:wei}; \\citealt{fri59:wei}). In this short Note, particle distribution functions are considered that describe two interpenetrating plasma streams. Such distributions are widely used in plasma astrophysics, because in the right reference frame, all outflow motion into ambient media and shock wave sites can be described by counterstreaming flows \\citep[e.\\,g.,][]{tau05:cov}. Examples are solar, stellar, and galactic winds, and the interaction of relativistic jets such as that from GRBs (gamma-ray bursts) and AGNs (active galactic nuclei) with the interstellar medium. Especially in the latter cases, one needs a distribution function that accounts for the relativistic effects, which can modify the instability rates significantly \\citep{usr08:wei}. In Sec.~\\ref{dist}, it will be shown how the relativistic distribution function for counterstreams in momentum space, which has been derived \\citep{tau05:cov} from the MBJ distribution above, can be rewritten in terms of velocity variables. In Sec.~\\ref{twostr}, the application to the two-stream instability will illustrate the significant differences between the relativistic and the non-relativistic counterstreaming distribution. Finally, the results are summarized in Sec.~\\ref{summ}. ", "conclusions": "\\label{summ} It has been known for a long time that, for relativistic temperatures (or, to be more accurate, for relativistic thermal velocities), the classic Maxwellian distribution function becomes increasingly inaccurate. Such is escpecially the case for distribution functions that describe plasmas streaming with high velocities, as is the case for two counterpropagating components. Because there exist wide application ranges of such distributions, it is both appropriate and necessary to take care of a precise basic construction. In this short Note, it has been shown how a particle distribution function that describes two interpenetrating plasma streams can be generalized to relativistic velocities. Based on the MJB distribution function, an expression has been derived that explicitely uses velocity coordinates. Such is advantageous both for analytical calculations and for numerical simulations that rely on velocity coordinates. As an example to illustrate the application of the relativistic counterstreaming distribution and, at the same time, to demonstrate the both qualitative and quantitative differences that result from the use of a correct relativistic distribution function, the electrostatic two-stream instability has been chosen. It has been shown that, due to the sharply emphasized anisotropy profile of the relativistic distribution function, the growth rate is significantly higher and the unstable wavenumber range is more extended as was the case for the non-relativistic distribution. In future work, the distribution function proposed here should be applied mainly to numerical simulations. The differences between non-relativistic, semi-relativistic, and fully relativistic initial distribution functions have to be worked out. Only then can one assess the necessity of a, for analytical calculations rather unwieldy, distribution function as that proposed here." }, "1003/1003.3401_arXiv.txt": { "abstract": "We consider the three-body decays of gravitino dark matter in supersymmetric scenarios with bilinear $R$-parity violation. In particular, gravitino decays into $\\ell W^*$ ($\\ell f\\bar f'$) and $\\nu Z^*$ ($\\nu f\\bar f$) are examined for gravitino masses below $M_W$. After computing the gravitino decay rates into these three-body final states and studying their dependence on supersymmetric parameters, we find that these new decay modes are often more important than the two-body decay, into a photon and a neutrino, considered in previous works. Consequently, the gravitino lifetime and its branching ratios are substantially modified, with important implications for the indirect detection of gravitino dark matter. ", "introduction": "Recent measurements indicate that about $25\\%$ of the energy density of the Universe is made up of a mysterious form of non-baryonic matter known as cold dark matter \\cite{Dunkley:2008ie}. To explain it, physics beyond the Standard Model is required. Supersymmetric extensions of the Standard Model are the most promising scenarios that may account for the dark matter. They are motivated, in addition, by the hierarchy problem and by the unification of the gauge couplings. Within supersymmetric models, R-parity conservation was long believed to be a prerequisite for supersymmetric dark matter, for it guarantees the stability of the lightest supersymmetric particle (LSP) --the would-be dark matter candidate. It was pointed out in \\cite{Takayama:2000uz}, however, that if the gravitino --the superpartner of the graviton that arises in local supersymmetric theories-- is the LSP, it can be a suitable dark matter candidate in R-parity violating models. Indeed, even though the gravitino is unstable in such scenarios, its lifetime may be much longer than the age of the Universe. Such a long lifetime is the result of the gravitino feeble interactions, which are suppressed by the Planck scale and by small R-parity violating couplings. A scenario with gravitino dark matter and R-parity violation is even favoured by thermal leptogenesis, as it alleviates the tension between the large reheating temperatures required by leptogenesis and the constraints from Big-Bang Nucleosynthesis \\cite{Buchmuller:2007ui}. In R-parity violating models, therefore, the primordial gravitino produced in the early Universe is a viable and well-motivated dark matter candidate. An important feature of these R-parity violating models is that, unlike their R-parity conserving counterparts, gravitino dark matter can be indirectly detected. In fact, the small fraction of gravitinos that have decayed until today constitutes a source of high energy cosmic rays \\cite{Bertone:2007aw,Ishiwata:2008cu}. In their decays, gravitinos may produce gamma rays \\cite{Bertone:2007aw,Ibarra:2007wg}, neutrinos \\cite{Covi:2008jy} and antimatter \\cite{Ibarra:2008qg} that could be observed in present and future experiments. Present data from FERMI, for instance, already constrains the parameter space of these models \\cite{Choi:2009ng}. The indirect detection signatures of gravitino dark matter strongly depend on the gravitino lifetime and on its branching ratios. In previous works, it has been assumed that they are determined by the two-body decays of the gravitino. We will show, in this paper, that that is not always the case. For gravitino masses below $M_W$, the three-body final states $\\ell W^*$ and $\\nu Z^*$, where * denotes a virtual particle, typically give a large contribution to the gravitino decay width, modifying in a considerable way the gravitino lifetime and its branching ratios. These results have important implications for scenarios of decaying gravitino dark matter, especially regarding their indirect detection prospects. ", "conclusions": "We have studied the decays into $\\tau W^*$ and $\\nu_\\tau Z^*$ of gravitino dark matter in the context of supersymmetric models with bilinear R-parity violation. We computed the gravitino decay rates and showed that these previously neglected processes typically give large contributions to the gravitino decay width. They actually dominate the decay rate over a wide range of gravitino and gaugino masses. In a future work, we will study in detail the implications of these results for gravitino dark matter." }, "1003/1003.4826_arXiv.txt": { "abstract": "We present new kinetic-gasdynamic model of the solar wind interaction with the local interstellar medium. The model incorporates several processes suggested by \\cite{mccomas09} for the origin of the heliospheric ENA ribbon -- the most prominent feature seen in the all sky maps of heliospheric ENAs discovered by the Interstellar Boundary Explorer (IBEX). The ribbon is a region of enhanced fluxes of ENAs crossing almost the entire sky. Soon after the ribbon's discovery it was realized (McComas et al., 2009) that the enhancement of the fluxes could be in the directions where the radial component of the interstellar magnetic field around the heliopause is close to zero (Schwadron et al., 2009). Our model includes secondary charge exchange of the interstellar H atoms with the interstellar pickup protons outside the heliopause and is a further advancement of the kinetic-gasdynamic model by \\cite{malama06} where pickup protons were treated as a separate kinetic component. \\cite{izmod09} have shown in the frame of Malama's model that the interstellar pickup protons outside the heliopause maybe a significant source of ENAs at energies above 1 keV. The difference between the current work and that of \\cite{izmod09} is in the assumption of no-scattering for newly created pickup protons outside the heliopause. In this limit the model produces a feature qualitatively similar to the ribbon observed by IBEX. ", "introduction": "The collision of the supersonic solar wind with the interstellar plasma flow results in formation of a complex interaction region or heliospheric interface. This region includes the termination and, possibly, bow shocks decelerating the solar wind (SW) and interstellar plasma, respectively, and the heliopause separating the two plasmas. The region of the heated SW behind the termination shock (TS) is known as the inner heliosheath, while the region behind the heliopause is called the outer heliosheath. The local interstellar medium (LISM) is a partly ionized medium consisting mainly of neutral atoms. It has become evident within recent years that the interstellar atoms have a pronounced effect on the global structure of the interface region and on the physical processes operating in the heliosphere. Apart from the fact that the position and shape of the TS and heliopause are significantly determined by the action of the atoms, they give rise to a specific hot population of pickup ions (PUIs). The first direct measurements of pickup helium \\citep{moebius85} and pickup hydrogen \\citep{gloeckler93} showed that the velocity distributions of the PUIs differ in significant ways from the velocity distributions of primary solar wind ions. The first measurements of the IBEX (Interstellar Boundary Explorer) spacecraft \\citep{mccomas09, fuselier09, funsten09, schwadr09} show results that were entirely unexpected. The objective of the IBEX mission is to image the complex interaction between the local interstellar medium (LISM) and the outflowing solar wind by measuring the fluxes of energetic neutral atoms (ENAs) originating in the outer parts of our heliosphere and beyond. The first scan of the whole sky showed that maxima of ENA fluxes form a long ($\\sim 250-300^{\\circ}$) and narrow ribbon-like feature that was not predicted by any model prior to the IBEX observations. The speed of the original interstellar atoms entering the heliosphere is $\\sim$26.4 km/s, which for hydrogen atoms corresponds to the kinetic energy of about 3 eV. Some portion of the atoms experiences charge exchange with shock heated solar wind protons and PUIs, and a new population of energetic atoms, created as a result of this process, has the broad energy distribution extending over several keV. These ENAs represent the energy distributions of the parent charged particles and, therefore, when measured at the Earth's orbit, can be used as a remote sensing of the ions in the interaction region. Current theoretical models of the SW/LISM interaction fall into two categories: standard models which assume instantaneous assimilation of pickup ions in the SW \\citep{baranov93}, and the compound or multi-component model \\citep{malama06} in the framework of which the pickup particles are considered as separate isotropic (in the solar wind rest frame) populations with their specific energy distributions. \\cite{izmod09} presented an extension of the \\cite{malama06} model by introducing a non-thermal population of pickup protons in the interstellar medium. These authors showed that the interstellar pickup protons form significant fluxes of ENAs dominating at energies above $\\simeq 1$ keV. Although the multi-component models are more comprehensive, all of the current numerical models predict that the ENA fluxes have maxima near the upwind direction of the heliosphere and minima at the flanks, though, of course, the position of maxima can slightly deviate from the upwind direction due to effects of the interstellar magnetic field and solar wind asymmetry (e.g. Izmodenov et al. 2009). \\cite{mccomas09} presented six possible concepts for the formation of the ribbon observed by IBEX. Among these concepts was the idea that neutralized solar wind propagates out beyond the heliopause, becomes ionized, gyrates about interstellar magnetic field lines, and then charge exchanges again to become ENAs. Some of these ENAs move back in toward the Sun where they can be imaged by IBEX. The advantage of this mechanism is that it produces sharply peaked ENA emissions in directions roughly perpendicular to the interstellar magnetic field beyond the heliopause - the same alignment inferred by comparing the IBEX ribbon to an MHD simulation of the heliosphere (Schwadron et al., 2009). Another concept suggested by \\cite{mccomas09} is that compression of the interstellar magnetic field beyond the heliopause may cause ions to align preferentially perpendicular to the interstellar magnetic field through conservation of the first adiabatic invariant and conservation of energy. This will also lead to a special orientation of peaked ENA emissions perpendicular to the interstellar magnetic field, and may help to explain both how the ribbon is formed and the even more surprising fine structure observed in it (McComas et al., 2009). The basic idea of secondary ENA generation of the IBEX ribbon was further examined by \\cite{heerick10}. These authors assumed that pickup protons in the outer heliosheath have and retain a partial shell distribution and that their re-neutralization is effectively instantaneous. This approach is significantly different from ours in this study since we solve consistently for the motion of pickup protons along magnetic field lines in the scatter-free limit, and thus include the motion of PUIs along the field line between their pickup and reneutralization. \\begin{figure}[t] \\noindent\\includegraphics[width=7cm]{fig1.eps} \\caption{The spatial distribution of the interstellar magnetic field around the heliospause in the $\\mathbf{BV}$ plane. The arrows show direction of the magnetic field, while the color indicates the magnetic field magnitude. The angle between $\\mathbf{B}$ and $\\mathbf{V}$ far from the heliopause equals $20^{\\circ}$ and the magnitude of $B$ is 4.4 $\\mu G$.} \\label{fig1} \\end{figure} ", "conclusions": "Here we have reported a new model without scattering, but including the effects of ion transport for the pickup protons generated in the region outside of the heliopause by charge exchange of the thermal interstellar protons and heliospheric ENAs. The results of the model yield a feature qualitatively similar to the IBEX ribbon. In future studies the results of simulations will be quantitatively compared to IBEX ENA observations. These further studies need to take into account ENAs for the inner heliosheath in proper kinetic way as it was done in \\cite{malama06}. Acknowledgements. The work of S.C., D.A., Y.M. was supported in part by Rosnauka under goskontrakt 02.740.11.5025 and by the RFBR in the frames of the projects 08-02-91968-DFG, 10-01-00258, 10-02-01316 and the Program for Basic Researches of OEMMPU RAS. V.I. was supported by President Grant MD-3890.2009.2 and Dynastia Foundation. Work in the USA was supported by the IBEX mission, which is a part of NASA's Explorer Program." }, "1003/1003.1859_arXiv.txt": { "abstract": "We introduce and discuss an effective model of a self-gravitating system whose equilibrium thermodynamics can be solved in both the microcanonical and the canonical ensemble, up to a maximization with respect to a single variable. Such a model can be derived from a model of self-gravitating particles confined on a ring, referred to as the self-gravitating ring (SGR) model, allowing a quantitative comparison between the thermodynamics of the two models. Despite the rather crude approximations involved in its derivation, the effective model compares quite well with the SGR model. Moreover, we discuss the relation between the effective model presented here and another model introduced by Thirring forty years ago. The two models are very similar and can be considered as examples of a class of minimal models of self-gravitating systems. ", "introduction": "Systems of classical particles mutually interacting via gravitational forces can model the behavior of many objects in the universe (globular clusters, elliptical galaxies, clusters of galaxies) as long as other interactions are negligible compared to the gravitational ones \\cite{BinneyTremaine}. When the number $N$ of particles is large the direct numerical simulation of such systems is a heavy task \\cite{Nbody} and it would be reasonable and useful to approach them via equilibrium statistical mechanics. However, self-gravitating systems do not have a ``true'' thermal equilibrium for two main reasons \\cite{IspolatovCohen}: $(i)$ the gravitational potential is singular for vanishing distance between two particles, making (at least part of) the system collapse in states with infinite density and $(ii)$ particles that do not collapse tend to escape the system (evaporation). From a physical point of view the first problem can be easily solved. No real system exists where the only non-negligible interaction is classical gravity at {\\em all} length scales: either the interacting ``particles'' are macroscopic bodies like stars or galaxies, or quantum effects must be taken into account below a certain length scale. In both cases, a length scale exists (the size of the bodies or the scale where quantum effects set in) below which the potential has no longer the classical gravitational form. If one is not interested in small-scale details, the potential can be regularized by replacing it with a softened one at short distances \\cite{softening,fermions} or by directly considering self-gravitating fermions \\cite{fermions}. To solve the second problem is less straightforward and one is forced to (somehow artificially) confine the particles in a finite volume. However, on physical grounds such an approximation is reasonable since in many cases the evaporation rate is slow compared to the other time scales in the system \\cite{FanelliMerafinaRuffo}. A regularized and confined self-gravitating system has a thermal equilibrium in both the canonical and the microcanonical ensemble \\cite{Kiessling}. Such a system can thus be studied within the framework of equilibrium statistical mechanics, although its behaviour in the two ensembles is very different: it can be considered a prototype of systems with ensemble inequivalence \\cite{CampaDauxoisRuffo}, showing e.g.\\ a core-halo phase with negative specific heat in the microcanonical ensemble which is replaced by a discontinuous phase transition from a clustered to a gas phase in the canonical ensemble \\cite{DeVega}. Canonical and microcanonical MonteCarlo simulations of a full self-gravitating system in three spatial dimensions are heavy \\cite{DeVega}, although they may be a little easier than the direct integration of the equations of motion. This suggests to look for simplified models which may be easier to study; one main simplification comes from considering models which are effectively one-dimensional. Many such models have been introduced in the last decades, ranging from the sheet model \\cite{HohlFeix} and the shell model \\cite{YoungkinsMiller} to the self-gravitating ring (SGR) model \\cite{SGR1}. The latter is particularly interesting because the interaction among the particles is given by the full three-dimensional gravitational potential (regularized at small distances), while the particles are confined on a ring. This yields a behaviour that is qualitatively very close to that found in three-dimensional systems, although allowing a much simpler study. In the limit of an infinite number of particles, the model can be studied in the mean-field approximation with a very efficient numerical technique \\cite{SGR2}, showing that in the microcanonical ensemble there is a phase transition separating a homogeneous high-energy phase from a clustered phase. An independent analytical argument supporting the existence of such a transition has been given in \\cite{gravring}. The aim of the present paper is to introduce and discuss an effective model which approximates the SGR model and is exactly solvable (up to a maximization with respect to a single variable which has to be performed numerically). Although the approximations involved are rather crude, the behaviour of this effective model is very close to that of the SGR model, not only at a qualitative level but also at a semi-quantitative one: this suggests that the approximations made do capture most of the important physics. The paper is organized as follows. After recalling the main features of the SGR model we introduce and discuss our effective model, and we present its solution in both the microcanonical and the canonical ensembles; the details of the calculations are reported in two appendices. Then, we compare our model with one introduced by Thirring forty years ago \\cite{Thirring}: the two models are indeed very similar, although the latter was not aimed at approximating any particular explicit model. We end with some remarks and a discussion of some open issues. ", "conclusions": "We have discussed how to derive an effective model of a self-gravitating system starting from the SGR model introduced and studied in Refs.\\ \\cite{SGR1,SGR2}. Such an effective model is solvable up to a numerical maximization in a single variable (which plays the r\\^ole of an order parameter) in both the microcanonical and the canonical ensembles. Following a suggestion coming from the study of the dynamics of the SGR model \\cite{SGR1,SGR2}, the effective model assumes that particles can be split into two families: cluster particles, all of which are interacting with each other by harmonic forces, and gas particles, which feel only a constant potential due to both the cluster particles and the mean field of the other gas particles. The fraction of gas particles, $n_g$, is the ``order parameter'' to be determined self-consistently. Despite the rather crude approximations used to derive the effective model, the results for the thermodynamic quantities are quite close to those found for the SGR model using the numerical method developed in \\cite{SGR2}, especially for small values of the softening parameter $\\alpha$, where the behaviour of the two systems is closer to that of an ``ideal'' self-gravitating system. Although one could expect an agreement at small energies, the agreement is definitely good also at energies up to and above the transition between the phase with negative specific heat and the homogeneous phase, which is well reproduced by our results. It is interesting that such a simple toy model as ours is able to reproduce the thermodynamics of the SGR model even at quantitative level. A qualitative disagreement appears at larger values of the softening parameter, where in the SGR model the entropy becomes concave and there is no longer ensemble inequivalence: in our effective model the entropy appears to be always nonconcave, for any value of $\\alpha$; however, the approximations made are no longer justified when $\\alpha$ is not small. Clearly, there is room for improvements: the caloric curve $T(\\varepsilon)$ is well reproduced when $\\alpha$ is small, but its shape in the low-energy part---see e.g.\\ Fig.\\ \\ref{fig_high}---shows that there are important anharmonic effects which should be taken into account. We have also compared our model with another model introduced by Thirring in 1970 \\cite{Thirring}. In that model one has a cluster of particles confined in a finite region of space, all interacting with each other via a constant attractive potential. This region of space is enclosed in larger volume where particles can escape becoming free particles. When the cluster extension is small, the thermodynamics of the two models is very similar, also in the low energy region dominated by the regularization of the potential, although the nature of the cluster is different in the two models. This shows, in our opinion, that ``cluster + gas'' toy models like these are good candidates as minimal models of self-gravitating systems. An interesting feature of our model (and of the Thirring model too, although this had not been noticed before, to the best of our knowledge) is the presence of a low-energy region where the fraction of gas particles $n_g$ is very small and stays very small up to a certain energy where it starts rising rapidly. There is a mathematical reason why $n_g$ can not be exactly zero in a finite region of energy, i.e., the degeneracy factor associated with counting the number of ways of constructing a state with a given fraction of gas particles. This counting assumes that all the gas particles are independent, which is clearly an oversimplification. One could wonder whether a more refined counting may imply the presence of a singularity at low energy, bounding a phase with $n_g\\equiv 0$, whose existence for the SGR model has been conjectured in \\cite{gravring}. A hint in this direction comes from assuming that the gas particles are all correlated like bosons (which is wrong, but is somehow the opposite situation to that considered here), so that states obtained from each other by interchanging two gas particles should not be considered as distinct and no degeneracy factor would be present in the density of states. Such a calculation has been performed in \\cite{cesare} and yields a sharp transition at a finite energy density and a phase with $n_g\\equiv 0$. The drawback is that the agreement with the thermodynamics of the SGR model is definitely worse. It is tempting to think that maybe a weaker transition occurs, and that it can be described by a ``proper'' counting of the degeneracy of the gas particle states." }, "1003/1003.4151_arXiv.txt": { "abstract": "We introduce a method for calculating the probability density function (PDF) of a turbulent density field in three dimensions using only information contained in the projected two-dimensional column density field. We test the method by applying it to numerical simulations of hydrodynamic and magnetohydrodynamic turbulence in molecular clouds. To a good approximation, the PDF of log(normalised column density) is a compressed, shifted version of the PDF of log(normalised density). The degree of compression can be determined observationally from the column density power spectrum, under the assumption of statistical isotropy of the turbulence. ", "introduction": "The probability density function (PDF) of the density field in molecular clouds is a key ingredient in most analytic models of star formation (Padoan \\& Nordlund 2002; Krumholz \\& McKee 2005; Elmegreen 2008; Hennebelle \\& Chabrier 2008; Padoan \\& Nordlund 2009; Hennebelle \\& Chabrier 2009). Knowledge of the density PDF is required to determine the overall star formation rate or efficiency. In some models, the density PDF is of central importance in determining the emergent stellar initial mass function. The majority of models assume a lognormal density PDF (V{\\'a}zquez-Semadeni 1994), with the width of the PDF increasing with the Mach number of the turbulence (Padoan, Nordlund, \\& Jones 1997; Passot \\& V\\'{a}zquez-Semadeni 1998; Federrath, Klessen, \\& Schmidt 2008). Observational knowledge of density fields in molecular clouds (let alone the PDF) is very limited. We do not have access to the density field in three dimensions (3D) but instead can only view the projected column density field in two dimensions (2D). Use of molecular tracers of different critical density can in principle yield some information, but even the most sophisticated excitation analyses provide only a single ``density'' per line-of-sight whereas the transverse variations in ``density'' so measured must imply comparable, or perhaps greater, fluctuations in density along the line-of-sight. A route to the PDF could involve (for example) the measurement of mass exceeding a range of critical densities from a suite of tracers. While the amount of data involved here would likely be prohibitive for nearby clouds of large angular extent, it can be usefully applied in an extragalactic context (e.g. Krumholz \\& Thompson 2007). Column density fields traced by extinction of background stars are perhaps the most robust way of acquiring constraining data on the density PDF in nearby clouds. Column density PDFs can be constructed (e.g. Cambresy 1999; Lombardi 2009; Kainulainen et al 2009) but the relation between these and the 3D density PDF is currently unknown. Compression of the PDF due to line-of-sight averaging is expected, and there are some indications that a lognormal density PDF will project into a (less broad) lognormal column density PDF (Ostriker, Stone, \\& Gammie 2001; V{\\'a}zquez-Semadeni \\& Garc{\\'i}a 2001, Federrath et al 2009). The degree of compression is presently unknown, but recent work by Brunt, Federrath, \\& Price (2010; hereafter BFP) demonstrated how to calculate the normalised density variance in 3D from information contained solely in the column density field. Comparison of the measured normalised column density variance with the inferred normalised density variance can provide some information on the degree of compression. In this paper, we introduce a method by which the 3D density PDF can be constructed from measurements made on the column density field alone. Measurements of the normalised column density variance, the column density power spectrum, and the column density PDF are required, and can be combined to construct an estimate of the 3D density PDF. ", "conclusions": "We have introduced and tested a simple method for reconstructing the probability density function (PDF) of a 3D turbulent density field using information present solely in the projected (observable) column density field in 2D. The method builds on a previously established method to calculate the 3D normalised density variance, recently presented by Brunt, Federrath, and Price (BFP, 2010). To a good approximation, the PDF of log(normalised column density) is a compressed, shifted version of the PDF of log(normalised density), but can deviate significantly in the extreme tails. The compression factor, $\\xi$, can be derived observationally from the column density power spectrum, assuming statistical isotropy, using the BFP method." }, "1003/1003.4198_arXiv.txt": { "abstract": "{The formation of massive stars is a highly complex process in which it is unclear whether the star-forming gas is in global gravitational collapse or an equilibrium state supported by turbulence and/or magnetic fields.} {By studying one of the most massive and dense star-forming regions in the Galaxy at a distance of less than 3 kpc, i.e. the filament containing the well-known sources DR21 and DR21(OH), we attempt to obtain observational evidence to help us to discriminate between these two views.} {We use molecular line data from our $^{13}$CO 1$\\to$0, CS 2$\\to$1, and N$_2$H$^+$ 1$\\to$0 survey of the Cygnus X region obtained with the FCRAO and CO, CS, HCO$^+$, N$_2$H$^+$, and H$_2$CO data obtained with the IRAM 30m telescope.} {We observe a complex velocity field and velocity dispersion in the DR21 filament in which regions of the highest column-density, i.e., dense cores, have a lower velocity dispersion than the surrounding gas and velocity gradients that are not (only) due to rotation. Infall signatures in optically thick line profiles of HCO$^+$ and $^{12}$CO are observed along and across the whole DR21 filament. By modelling the observed spectra, we obtain a typical infall speed of $\\sim$0.6 km s$^{-1}$ and mass accretion rates of the order of a few 10$^{-3}$ M$_\\odot$ yr$^{-1}$ for the two main clumps constituting the filament. These massive clumps (4900 and 3300 M$_\\odot$ at densities of around 10$^5$ cm$^{-3}$ within 1 pc diameter) are both gravitationally contracting. The more massive of the clumps, DR21(OH), is connected to a sub-filament, apparently 'falling' onto the clump. This filament runs parallel to the magnetic field.} {All observed kinematic features in the DR21 filament (velocity field, velocity dispersion, and infall), its filamentary morphology, and the existence of (a) sub-filament(s) can be explained if the DR21 filament was formed by the convergence of flows on large scales and is now in a state of global gravitational collapse. Whether this convergence of flows originated from self-gravity on larger scales or from other processes cannot be determined by the present study. The observed velocity field and velocity dispersion are consistent with results from (magneto)-hydrodynamic simulations where the cores lie at the stagnation points of convergent turbulent flows. } ", "introduction": "\\label{intro} Our goal in this series of papers is to investigate the generic link between the large-scale (several 10 pc) molecular-cloud spatial structure (Schneider et al. 2010), medium-scale ($<$10 pc) dynamic fragment properties (this paper), and the occurrence of high-mass star formation on small (0.1-0.5 pc) scales (Bontemps et al. \\cite{bontemps2010}, Csengeri et al. \\cite{csengeri2010}). To achieve this, we make use of our multiwavelength study in Cygnus X (see below). This region has already been shown to represent an excellent laboratory for studies of high-mass star formation. Molecular clouds (MCs) are the birthplaces of low- and high-mass stars. Understanding their formation and evolution is essential to understanding in general the star formation (SF) process. There are currently two different scenarios for the evolution of MCs and stars, a {\\sl dynamic} and a {\\sl quasi-static} view. In the dynamic context, MCs are transient, rapidly evolving entities that are not in equilibrium, and the spatial and velocity structure of the cloud is determined by compressible supersonic turbulence (see e.g., Mac Low \\& Klessen \\cite{maclow2004} for an overview). The driving sources of turbulence can be diverse and their relative importance remains a subject of debate. Large-scale mechanisms such as supernovae explosions, should be the most importantXS, based on theoretical arguments extensively discussed by MacLow \\& Klessen (\\cite{maclow2004}). Turbulence may also occur during the formation process of molecular clouds (Vazquez-Semandeni et al. \\cite{vaz2002}, Klessen \\& Hennebelle \\cite{klessen2010}) within large-scale colliding flows of mostly atomic gas in the galactic disk, generated by dynamic compression in the interstellar medium or other large-scale instabilities (see, e.g., Hennebelle \\& Audit \\cite{hennebelle2007}, Vazquez-Semadeni et al. \\cite{vaz2008}, Hennebelle et al. \\cite{hennebelle2008a}, Banerjee et al. \\cite{banerjee2009}). High velocity compressive flows form dense structures at stagnation points that may collapse to form stars/clusters (Ballesteros-Paredes et al. \\cite{ball2003}, Vazquez-Semadeni et al. \\cite{vaz2007}, \\cite{vaz2009}, Heitsch et al. \\cite{heitsch2008}). In this case, the velocity field and velocity dispersion of molecular clumps may contain the signature of the external convergent motion {\\sl and} the compressive, gravitational contraction (this was already proposed by Goldreich \\& Kwan \\cite{goldreich1974}). The star-forming core is then the dense post-shock region with a more quiescent velocity dispersion than the turbulent flow (Klessen et al. \\cite{klessen2005}). This scenario is valid for both, low- and high-mass star formation. However, there is no direct observational confirmation of the existence of {\\sl convergent flows}. Studying the dynamics of HI and molecular gas may be a possibility (Brunt \\cite{brunt2003}), though it is not clear what observational signatures are to be expected. An indirect argument for molecular cloud formation out of colliding HI flows was provided by Audit \\& Hennebelle (\\cite{audit2009}), who showed that their models closely reproduce the observed clump mass spectra and Larson-relations. It has also been noted that the lifetime of the cloud in the gravoturbulent framework is short (one dynamical crossing time, i.e. $\\sim$10$^7$ yr for giant molecular clouds) and the star formation process is rapid. In the quasi-static view (see McKee \\& Ostriker \\cite{mckeeostriker2007} and references therein), MCs are formed by large-scale self-gravitating instabilities such as spiral density waves. (High-mass) star formation is approached in the 'turbulent core' model (McKee \\& Tan~\\cite{mckee2003}) in which the star-forming cores are supersonically turbulent. Turbulent magnetic and thermal pressure supports the clump against self-gravity. Because of energy injection from newly formed stars, the clumps and most of the cores may maintain their equilibrium, unless they quasistatically evolve to a gravitationally unstable state to finally form stars (if the magnetic flux diffuses out of the clump by ambipolar diffusion). In the case of low-mass stars, the collapse of a rotating cloud of gas and dust leads to the formation of an accretion disk through which matter is channeled onto a central protostar. For stars with masses higher than about 8 M$_\\odot$ this mechanism of star formation is less clear due to the strong radiation field that pushes against infalling material and may halt accretion. However, theoretical work (Yorke \\& Sonnhalter \\cite{yorke2002}; Krumholz, McKee, Klein \\cite{krumholz2005}; Peters et al. \\cite{peters2010}) has shown that outflows lead to anisotropy in the stellar radiation field, which reduce the radiation effects experienced by gas in the infalling envelope. Thus, massive stars may therefore be able to form by a mechanism similar to that by which low mass stars form. Peters et al. (\\cite{peters2010}) even show that clustered SF is a natural outcome of massive SF even in the presence of radiative feedback. However, clustered SF is also a natural result of the {\\sl competitive accretion} scenario (Bonnell \\& Bate \\cite{bonnell2006}), in which the fragmentation of a turbulent cloud produces stars with masses of the order of the Jeans mass within a common gravitational potential. These stars, located close to the centre of the potential, accrete at much higher rates than isolated stars and become massive. Both scenarios have in common that massive ($>$10\\,M$_\\odot$) stars form only from the coldest, and densest, cores (size scale $<$0.1 pc) of molecular clouds and that high infall/accretion rates are required to overcome feedback processes (ionizing radiation, jets, outflows). In the dynamic framework, stars form as the dynamical evolution of the MC progresses with gas continuously being accreted. Individual clumps ($\\sim$0.1--0.5 pc) fragment into hundreds of protostars, possibly competing for mass in the central regions of the cluster (Bonnell \\& Bate \\cite{bonnell2006}). Support from magnetic fields (Hennebelle \\& Teyssier \\cite{hennebelle2008b}) and/or radiation (Krumholz \\cite{krumholz2006}, Bate \\cite{bate2009}) could drastically limit the level of fragmentation and channel the global infall to fewer, more massive protostars (Bontemps et al. \\cite{bontemps2010}). From the observational point of view, indications of {\\sl global collapse} have been detected using molecular lines. Spectral profiles of high density tracers, usually combining at least an optically thin and an optically thick line, are good probes of infalling gas (e.g. Myers et al. \\cite{myers1996}). Observations in low-mass star-forming regions (e.g. Lee et al. \\cite{lee2003}) contain a mixture of infall (blue-shifted emission and/or redshifted self-absorption in the optically thick line) and outflow asymmetry. This sort of line profile was also found in the high-mass star-forming region W43 (Motte et al.~\\cite{motte2003}). Peretto et al. (\\cite{peretto2006}) proposed that a peculiar velocity discontinuity could be the result of some large-scale motion, originating from self-gravity (Peretto et al. \\cite{peretto2007}). The theory of {\\sl turbulent core formation}, on the other hand, predicts the existence of massive pre-stellar cores that have not been detected so far, not even in sensitive and extensive dust continuum studies (e.g. Motte et al. \\cite{motte2007}). In this paper, we shortly introduce the Cygnus X region and the DR21 filament in Sect.~\\ref{cygnusx} and describe the molecular line oberservations in Sect.~\\ref{obs} and show maps and spectra in Sect.~\\ref{results}. Section~\\ref{analysis} presents an analysis of the kinematic structure and the physical properties of the DR21 filament and in Sect.~\\ref{discuss} we use our findings to test the conditions of the different high-mass star formation and turbulence models. Section~\\ref{summary} summarizes the paper. \\begin{figure*}[ht] \\includegraphics[angle=0,width=140mm]{figure1_paper.jpg} \\caption [] {Zeroth moment map of CS 2$\\to$1 emission (red contours, mapping area indicated by a continuous blue line) in Cygnus X overlaid on a moment map of $^{13}$CO 1$\\to$0 emission (mapping area indicated by a dashed line). Both maps were obtained with the FCRAO and have an angular resolution of $\\sim$50$''$. Emission was integrated in the velocity range --10 to 20 km s$^{-1}$ and the grey scaling ranges from 1 to 15 K km s$^{-1}$. Red contour levels for CS are at 0.5,1,2,4 K km s$^{-1}$. The DR21 filament is indicated, as well as some prominent objects in Cygnus X (S106, AFGL2591, W75N, the most massive stars of the OB2 cluster). } \\label{overview} \\end{figure*} \\begin{figure*}[ht] \\includegraphics[angle=0,width=80mm]{moment_cs_paper_dr21.jpg} \\hskip 1cm \\includegraphics[angle=0,width=80mm]{moment_n2h_paper_dr21.jpg} \\caption [] {Zeroth moment maps of CS 2$\\to$1 (left, contour levels are 0.5, 1, 2, 4 K km s$^{-1}$) and N$_2$H$^+$ 1$\\to$0 (right, contour levels are 0.5, 1.5, 3 K km s$^{-1}$) emission in Cygnus X North in the velocity range --10 to 20 km s$^{-1}$ observed with the FCRAO. Thermal HII regions (DR17--23) are indicated by red stars, mm-continuum sources from Motte et al. (\\cite{motte2007}) by green triangles. The latter correspond very well with peaks of N$_2$H$^+$ emission.} \\label{fcrao-total} \\end{figure*} \\begin{table}[ht] \\label{obs-table} \\caption{Observing parameters of the molecular line data obtained with the FCRAO and IRAM 30 m telescope: Column one and two indicate the molecular transition and frequency, followed by the half power beam width (HPBW) in arcsec, $\\eta_{\\rm mb}$ is the main beam efficiency, $T_{\\rm sys}$ the system temperature, $\\Delta {\\rm v}_{\\rm res}$ denotes the velocity resolution, and $\\Delta T_{\\rm rms}$ the average rms noise temperature per channel on a $T_{\\rm mb}$ scale.} \\begin{center} \\begin{tabular}{lcccccccc} \\hline \\hline & $\\nu$ & \\small{HPBW} & $\\eta_{mb}$ & T$_{sys}$ & $\\Delta {\\rm v}_{\\rm res}$ & $\\Delta T_{\\rm rms}$ \\\\ & [GHz] & & & [K] & [km s$^{-1}$] & [K] \\\\ \\hline {\\bf FCRAO} & & & & & & \\\\ \\hline $^{13}$CO 1$\\to$0 & 110.2 & 45$''$ & 0.48 & 210 & 0.067 & 0.48 \\\\ CS 2$\\to$1 & 98.0 & 48$''$ & 0.48 & 200 & 0.075 & 0.38 \\\\ N$_2$H$^+$ 1$\\to$0 & 93.2 & 48$''$ & 0.48 & 200 & 0.079 & 0.38 \\\\ \\hline {\\bf IRAM} & & & & & & \\\\ \\hline $^{12}$CO 2$\\to$1 & 230.8 & 11$''$ & 0.52 & 601 & 0.05 & 0.8 \\\\ H$_2$CO \\tiny{3(1,2)-2(1,1)} & 225.7 & 11$''$ & 0.55 & 549 & 0.05 & 0.65 \\\\ $^{13}$CO 2$\\to$1 & 220.4 & 11$''$ & 0.57 & 580 & 0.03 & 0.9 \\\\ C$^{34}$S 2$\\to$1 & 96.4 & 26$''$ & 0.77 & 210 & 0.025 & 0.5 \\\\ N$_2$H$^+$ 1$\\to$0 & 93.2 & 26$''$ & 0.77 & 158 & 0.03 & 0.1 \\\\ HCO$^+$ 1$\\to$0 & 89.2 & 28$''$ & 0.77 & 138 & 0.05 & 0.2 \\\\ H$^{13}$CO$^+$ 1$\\to$0 & 86.8 & 29$''$ & 0.78 & 114 & 0.025 & 0.25 \\\\ \\hline \\end{tabular} \\end{center} \\end{table} \\section {Cygnus X} \\label{cygnusx} \\noindent {\\sl The Cygnus X region} \\\\ Cygnus X is one of the richest star-formation regions in the Galaxy, covering an area of about 30 square degrees in the Galactic plane around Galactic longitude 80$^\\circ$ (Reipurth \\& Schneider \\cite{reipurth2008}). It contains a prominent OB-association, Cyg OB2, with $\\sim$100 O-stars and a total stellar mass of up to 10$^5$ M$_\\odot$ (Kn\\\"odlseder \\cite{knoedl2000}). From large-scale $^{13}$CO 2$\\to$1 (KOSMA\\footnote{Cologne Observatory for Submm-Astronomy}, Schneider et al. \\cite{schneider2006}) and $^{13}$CO 1$\\to$0 (FCRAO\\footnote{Five College Radio Astronomy Observatory}, Schneider et al. \\cite{schneider2007}, Simon et al., in prep.) surveys, we obtained a mass of a few 10$^6$ M$_\\odot$ for the whole molecular cloud complex that is divided into two parts -- Cygnus North and South. These studies also showed that the majority of molecular clouds is located at a common distance of about 1.7 kpc, i.e. the distance of Cyg OB2. Its proximity makes Cygnus X one of the rare laboratories where different phases of massive star formation can be studied in detail. More than 800 distinct HII regions, a number of Wolf-Rayet stars, several OB associations, and at least one star of spectral type O4If are known in Cygnus X, reflecting its record of high-mass star formation in the past. Ongoing massive star formation is revealed by wide-field (3 deg$^2$) 1.2 mm continuum imaging of the densest regions in the Cygnus X molecular clouds (Motte et al. \\cite{motte2007}). More than 100 protostellar dense-cores were detected of which 40 are likely to be the precursors of massive OB stars. High-angular resolution observations with the Plateau de Bure Interferometer (Bontemps et al. \\cite{bontemps2010}) have indeed confirmed that the most massive dense cores of this sample are the sites of massive star formation, mainly in the form of clusters. \\\\ \\noindent {\\sl The DR21 filament} \\\\ From our 1.2 mm continuum imaging and the $^{13}$CO surveys, we concluded that the molecular ridge containing DR21 and DR21(OH) (Dickel, Dickel \\& Wilson \\cite{dickeldickel1978}, Wilson \\& Mauersberger \\cite{wilson1990}, Jakob et al. \\cite{jakob2007}) is the most active, dense (average density$\\sim$10$^4$ cm$^{-3}$), and massive (34 000 M$_\\odot$) cloud in Cygnus X. DR21 (Downes \\& Rinehart \\cite{downes1966}) itself is a group of several compact HII regions (Harris et al. \\cite{harris1973}) with an associated, very energetic outflow (e.g. Garden et al. \\cite{garden1991a}, Russell et al. \\cite{russell1992}), located in the southern part of the ridge. Three arcminutes further north lies DR21(OH), famous for its maser emission (H$_2$O, Genzel \\& Downes \\cite{genzel1977}; OH, Norris et al. \\cite{norris1982}; CH$_3$OH, Batrla \\& Menten \\cite{batrla1988}). Even further north lies the massive star-forming region W75N (see Shepherd et al. \\cite{shepherd2004} for a review). Very recently, the DR21 filament was again the target for a wealth of studies in various wavelengths. Infrared images at 3.6, 4.5, and 8 $\\mu$m from the Spitzer satellite (Marston et al. \\cite{marston2004}, Hora et al. \\cite{hora2009}) show the complexity of the region with many IR-filaments perpendicular to the ridge that seem to be sites of star formation (Kumar et al. \\cite{kumar2007}). Dust continuum studies (Vall\\'ee \\& Fiege \\cite{vallee2006}, Motte et al. \\cite{motte2007}) confirmed the detection of 8 compact, dense cores in the DR21 filament (Chandler et al. \\cite{chandler1993}) and found additional ones. Some of them are high mass protostars and drive outflows detected in SiO (Motte et al. \\cite{motte2007}). An H$_2$ 1-0 S(1) image at 2.121 $\\mu$m (Davis et al. \\cite{davis2007}) shows at least 50 individual outflows driven by embedded low-mass stars. \\begin{figure*}[ht] \\includegraphics[angle=-90,width=150mm]{channels_filaments_paper.jpg} \\caption [] {Channel maps of $^{13}$CO 1$\\to$0 emission in Cygnus X North in the velocity range --3.2 to --0.9 km s$^{-1}$ observed with the FCRAO. The red triangle shows the position of DR21(OH). The three major subfilaments (F1, F2, F3) linked with the large DR21 filament are indicated in the plot of velocity --2.6 km s$^{-1}$.} \\label{fcrao-channel} \\end{figure*} We studied this particular filament in high-angular resolution observations with the IRAM\\footnote{Institut de Radio-Astronomie Millimetrique} 30 m telescope in different molecular tracers (isotopomeric CO lines, CS, C$^{34}$S, HCO$^+$, H$^{13}$CO$^+$, N$_2$H$^+$, H$_2$CO) to investigate the distribution of the different phases of molecular gas (cold dense cores, warm envelopes) and to uncover infall/outflow signatures. The lower angular resolution FCRAO data delineate the large-scale structure in which the DR21 filament is embedded. A more detailed analysis of the FCRAO $^{13}$CO data set is presented in Schneider et al. (\\cite{schneider2007}, \\cite{schneider2010}) and the CS and N$_2$H$^+$ data will be discussed in a forthcoming paper. All molecular line data we obtained serve as a basis for a physical model of the DR21 filament, including detailed non-LTE line modelling. ", "conclusions": "\\label{summary} We have presented a detailed molecular line study of the molecular ridge containing the star-forming regions DR21 and DR21(OH). This ridge is embedded in a large-scale network of filamentary structures, revealed by our maps of $^{13}$CO 1$\\to$0, CS 2$\\to$1, and N$_2$H$^+$ 1$\\to$0 emission obtained with the FCRAO. It is the most massive (around 30 000 M$_\\odot$) and dense (average density $\\sim$10$^4$ cm$^{-3}$) filament within the region and is labeled by us the 'DR21 filament'. Several sub-filaments are linked to the DR21 filament, the most massive one runs orthogonal to the North-South oriented ridge and has a mass (determined from $^{13}$CO 1$\\to$0) of 2600 M$_\\odot$ and an average density of 690 cm$^{-3}$. Its inferred dynamical time is $\\sim$2$\\times$10$^6$ yr. Higher angular resolution IRAM molecular line observations in HCO$^+$, H$^{13}$CO$^+$, $^{12}$CO/$^{13}$CO 2$\\to$1, C$^{34}$S, N$_2$H$^+$, and H$_2$CO resolve the detailed structure of the DR21 filalment. The H$^{13}$CO$^+$ 1$\\to$0 data show how the sub-filament seen in $^{13}$CO connects directly to the DR21(OH) clump. From the $^{12}$CO 2$\\to$1 line mapping, we confirmed the known outflow sources DR21 and DR21(OH) and detected three new ones, correlated with the mm-continuum sources N53, N44, and N45 (Motte et al. \\cite{motte2007}). The HCO$^+$ 1$\\to$0 line shows self-absorbed lines across the whole filament. Since optically thin lines peak in the gap and the blue wing of HCO$^+$ is more intense than the red one, we conclude that this emission feature is due to infalling gas. The typical infall speed, determined with a simple method described in Myers et al. (\\cite{myers1996}), is 0.6--0.8 km s$^{-1}$. A more sophisicated non-LTE modelling of the HCO$^+$ and H$^{13}$CO$^+$ lines using the {\\sl Simline} radiative transfer code yields an infall speed of 0.5 and 0.6 km s$^{-1}$ for the northern and southern part of the DR21 filament (but excluding the DR21 region itself), respectively. The kinematic structure of the DR21 filament is remarkable. We measured, using the N$_2$H$^+$, H$^{13}$CO$^+$, and HCO$^+$ maps, the highest values of the velocity dispersion in a vertical column of low column density gas that is offset from the dense clumps seen in N$_2$H$^+$. These results can be explained if the filament was produced by turbulent flows. In this case, the densest gas has been shocked and slowed down so that the largest velocity dispersions do occur not in the densest regions but in the outskirts (Klessen et al.~\\cite{klessen2005}, Vazquez-Semadeni et al. ~\\cite{vaz2008}). We also observe three horizontal (west-east) velocity gradients (--0.8, --2.2, +2.3 km s$^{-1}$ pc$^{-1}$) in the filament in which the one with the positive value marks the location where the filament is 'falling' onto the DR21(OH) clump. This velocity pattern cannot be explained by a single rotation of the filament along a vertical axis. An alternative explanation would be that at least part of the observed motions is due to convergent flows. By comparing our observations of the DR21 filament with a hydrodynamic (Federrath et al. \\cite{fed2009}) and a magneto-hydrodynamic turbulence model (Hennebelle et al., Teyssier \\cite{teyssier2002}, Fromang et al. \\cite{fromang2006}), we infer that a very dynamic and fast mode of star formation occurs in the filament, in which gas is continuously replenished by subfilaments attached to the main filament. These subfilaments are aligned with the magnetic field direction that is perpendicular to the DR21 filament (Vallee \\& Fiege \\cite{vallee2006}). The DR21 filament is globally collapsing. All our observational findings are incompatible with the view of a quasi-static, pressure-bounded clump scenario." }, "1003/1003.0896_arXiv.txt": { "abstract": "Although redshift-space distortions only affect inferred distances and not angles, they still distort the projected angular clustering of galaxy samples selected using redshift dependent quantities. From an Eulerian view-point, this effect is caused by the apparent movement of galaxies into or out of the sample. From a Lagrangian view-point, we find that projecting the redshift-space overdensity field over a finite radial distance does not remove all the anisotropic distortions. We investigate this effect, showing that it strongly boosts the amplitude of clustering for narrow samples and can also reduce the significance of baryonic features in the correlation function. We argue that the effect can be mitigated by binning in apparent galaxy pair-centre rather than galaxy position, and applying an upper limit to the radial galaxy separation. We demonstrate this approach, contrasting against standard top-hat binning in galaxy distance, using sub-samples taken from the Hubble Volume simulations. Using a simple model for the radial distribution expected for galaxies from a survey such as the Dark Energy Survey (DES), we show that this binning scheme will simplify analyses that will measure baryon acoustic oscillations within such galaxy samples. Comparing results from different binning schemes has the potential to provide measurements of the amplitude of the redshift-space distortions. Our analysis is relevant for other photometric redshift surveys, including those made by the Panoramic Survey Telescope \\& Rapid Response System (Pan-Starrs) and the Large Synoptic Survey Telescope (LSST). ", "introduction": "\\label{sec:intro} The late-time acceleration of the expansion of the Universe has been one of the most exciting cosmological discoveries in recent years \\citep{riess98,perlmutter99}. Understanding the nature of this acceleration is one of the main challenges facing cosmologists. One of the key observational methods that will be used to help meet this challenge involves using Baryonic Acoustic Oscillations (BAO) in the 2-point galaxy clustering signal as a standard ruler to make precise measurements of cosmological expansion. The acoustic signature has now been convincingly detected \\citep{percival01,cole05,eisenstein05} using the 2dF Galaxy Redshift Survey (2dFGRS; \\citealt{colless03}) and the Sloan Digital Sky Survey (SDSS; \\citealt{york00}). The detection has subsequently been refined using more data and better techniques, and is now producing interesting constraints on cosmological models \\citep{percival07a,percival07b,gaztanaga08,sanchez09,percival09}. Some of the next generation of sky surveys, including the Dark Energy Survey (DES {\\tt www.darkenergysurvey.org}), the Panoramic Survey Telescope and Rapid Response System (PanStarrs {\\tt pan-starrs.ifa.hawaii.edu}), and the Large Synoptic Survey Telescope (LSST {\\tt www.lsst.org}), will use photometric techniques to estimate galaxy redshifts, rather than more precise estimates from spectroscopic emission lines. The larger uncertainties on galaxy redshifts induce errors on inferred distances in the radial direction. The amplitude of the power spectrum and correlation function is reduced in the radial direction by this smoothing, removing information. In this scenario, where little information remains from fluctuations in the radial direction, it makes sense to use the projected 2-pt functions in photometric-redshift slices as the statistics to compare with models \\citep{padmanabhan07,blake07}. The projection does not completely remove problems caused by inferring distances from velocity data (i.e. working in redshift-space). The distribution of galaxies that we observe in sky surveys, where we measure radial distances from spectroscopic or photometric redshifts, is not a true 3D picture. We observe an apparent clustering pattern in {\\it redshift-space}, which is systematically different from the true distribution in {\\it real-space} because redshifts of galaxies are altered from their Hubble flow values by peculiar velocities. For example, on large scales, the infall of galaxies onto collapsed objects leads to an apparent enhancement of clustering in the radial direction as galaxies are projected along their velocity vectors \\citep{kaiser87,hamilton98}. When we infer galaxy distances assuming that the total velocity relative to the observer comes from the Hubble expansion flow, the result is that we see a distorted (redshift-space) density field. For angular measurements, these redshift-space distortions can alter the angular clustering in a redshift slice because the distortions are correlated across the direction of projection. Although redshift-space distortions are sub-dominant compared with photometric redshift uncertainties, they give rise to a systematic effect, which needs to be included when photometric redshift surveys are analysed \\citep{padmanabhan07,blake07}. This can complicate the analysis as the size of the redshift-space distortions, and therefore of this effect, is dependent on the cosmological model. Consequently, for every model to be tested against the data, we need to make a revised estimate of the redshift-space effect. \\begin{figure} \\centering \\resizebox{0.9\\columnwidth}{!}{\\includegraphics{photo-z-slice.ps}} \\caption{The radial distribution of galaxies selected in a bin of width $400\\mpcoh$, calculated using photometric redshifts to estimate distances (solid line). This is compared against the distribution of true distances to these galaxies (dashed line) assuming a photometric redshift error of $\\sigma_z=0.03(1+z)$. If the photometric redshifts of different galaxies are independent, then the expected projected correlation function of the photo-z selected sample, and a sample selected applying the dashed line as a selection function based on the true distances, are the same.\\label{fig:photo-z-slice}} \\end{figure} In this paper, we consider the simplified problem in the plane-parallel approximation, and only consider linear redshift-space distortions. Both photometric redshift errors and the random motion of galaxies in clusters provide an additional convolution of the overdensity field along the radial direction. While these effects need to be corrected in any analysis, the required correction is easily modelled and can be separated from the linear redshift-space distortion effects. For a measurement of the projected clustering, including such effects is equivalent to simply broadening the radial window function with which the galaxies were selected. This is demonstrated in Fig.~\\ref{fig:photo-z-slice}. A top-hat bin in photometric redshift gives the same expected projected correlation function as simply applying the convolved version of the bin as a selection function for the true distances. As we have to include a window function anyway, we simply assume in this paper that this window already includes the effects of both photometric redshift errors and the random motion of galaxies in clusters. In the following analysis, we therefore assume that there are no redshift errors without loss of generality. The layout of our paper is as follows. In Section~\\ref{sec:2pt} we analyse the projected overdensity field and redshift-space effects upon it, both analytically (Section~\\ref{sec:xi_p} \\&~\\ref{sec:pk_p}) and using Monte-Carlo simulations (Section~\\ref{sec:mc_sim}). We then consider how the recovered correlation function depends on galaxy binning (Section~\\ref{sec:dep_win}). Mock catalogues drawn from the Hubble Volume simulation are constructed and analysed in Section~\\ref{sec:hv_sim} in order to validate this analytic work. We incorporate hybrid selection functions based on both real and redshift-space boundaries into the analysis in Section~\\ref{sec:hybrid}. In Section~\\ref{sec:des} we consider a non-uniform redshift distribution similar to that of future sky survey DES, and the realistic implementation of our work is discussed in Section~\\ref{sec:discussion}. ", "conclusions": "\\label{sec:discussion} Redshift distortions produce a strong effect on projected clustering measurements --- one that is far stronger than the redshift-space distortion effect on the 3D clustering signal for galaxy samples with low bias and a narrow radial window. It is clear that redshift distortion effects must be included when modelling the projected galaxy clustering in redshift slices. If we consider the apparent motion of galaxies as we move from real- to redshift-space, then redshift-space distortions cause an apparent coherent motion of galaxies into and out of samples. This is true whether samples have sharp boundaries, or if the selection function changes more gradually with distance. In fact, we have argued that such motion does not in itself alter the projected correlation function --- we would recover the real-space projected correlation function if we could correct for the movement of the boundary (i.e. allow for the depth of the survey to change with the distortions). However, this is not easy to do, although it is theoretically possible and is an interesting alternative approach. The effect of redshift-space distortions is due to the redshift-space boundaries themselves having an angular clustering signal, and their correlation with the overdensity field. We can alternatively view the effect from a Lagrangian standpoint, where we have to consider that the projection does not remove redshift-space effects from the anisotropic correlation function. We have used Hubble Volume simulations to show that the projected correlation function can be modelled most easily by integrating the redshift-space correlation function over the radial selection function. Galaxy selection will often be a mix of real and redshift-space constraints, and we have shown that this can be modelled by splitting the population into samples that can be considered to have top-hat windows in either real-space, redshift-space or a hybrid of the two. In the hybrid situation, the projected correlation function can be modelled using both the real-space and redshift-space correlation function over the radial selection function, and that more complicated selection functions can be effectively modelled in a similar manner. Prior to this publication, no-one has considered how these hybrid selection functions affect the recovered projected clustering signal. \\subsection{Pair-Centre Binning} We have presented a new measurement technique, {\\it pair-centre} binning, and shown that it minimises the effects of redshift space distortions. In this new scheme, we only include galaxies where their apparent {\\it pair-centres} lie within a given radial bin, whereas traditional methods select pairs where both galaxies lie within the bin. The new scheme includes individual galaxies that lie {\\it outside} the traditionally applied top-hat boundaries. This simple modification acts to reduce the effect of the coherent movement of galaxies between slice boundaries on projected correlation function clustering analyses. It is important to note that this new technique does not {\\it prevent} the movement of galaxies between slices; redshift-space distortions due to peculiar velocities will always exist in the radial direction. It simply makes sure that they do not produce a coherent effect on the measurements. There are two potential disadvantages of the pair-centre binning scheme. One is the fact that the same galaxy may be included in multiple radial bins --- thus introducing a correlation between radial bins. Another is the fact that such a scheme results in necessarily wider radial bins, which causes the clustering signal to be diluted. We do not feel that either is a large problem. Applying the more traditional top-hat binning scheme to photometric surveys necessarily results in overlapping radial bins (due to photometric redshift errors) and there will always be considerable covariance between radial bins selected with photometric redshifts --- we do not think that pair-centre binning will make this problem considerably worse. The dilution effect can be mitigated by imposing a maximum separation between the pairs included in a pair-centre bin: we call this constrained pair-centre binning. As can be seen by comparing the middle and right-hand panels of Figs.~\\ref{fig:xi-des} and~\\ref{fig:xi-dess}, imposing such a constraint increases the expected signal while not causing a significant change in the effects of redshift-space distortions. More detailed studies of these effects are warranted, but we are confident that the reduction in the redshift distortion effect we observe when utilising pair-centre binning will make this scheme considerably preferable to a top-hat binning scheme. Pair-centre binning completely removes the effect of redshift distortions when given a uniform galaxy distribution. Such perfect distributions do not exist --- most galaxy samples selections are based on an apparent magnitude limit --- and thus realistic radial distributions of galaxies are more complicated. However, we have argued that if galaxy samples selected based on an apparent magnitude limit are cut back so that no galaxies $k$-corrected galaxies are missing from the sample, then this does not matter: the boundaries of the bins are either in real-space, or based on pair-centres, neither of which introduces redshift distortion effects. We have argued, and it is clear from previous work, that any interpretation of projected clustering measurements must account for redshift space distortions. In fact, comparing correlation functions calculated using different binning schemes might actually prove to provide a mechanism for measuring the amplitude of the redshift-space distortions. This is beyond the scope of our current draft, and we leave this for subsequent work. \\subsection{Future Surveys} To quantify the effect of redshift-space distortions for future surveys, we have used the expected radial selection function and photometric redshift distribution for the Dark Energy Survey to predict the effect of redshift-space distortions on projected clustering measurements. This analysis is also relevant to other planned surveys such as PanStarrs and the LSST, which will have similar radial selection functions. We have contrasted two different types of binning: top-hat --- in which we only allow galaxies between a given radial bound to enter our sample--- and pair-centre --- in which we only count galaxy pairs with an average radial position that lies within our bounds. For typical bin widths that will be applied to these surveys, we find that top-hat binning in the radial direction leaves a strong signal from redshift-space distortions. Using a pair-centre binning scheme reduces the redshift-space distortion signal, by as much as 80$\\%$ in realistic situations (see Fig.~\\ref{fig:xi-des}) and should therefore allow the measurements to be more sensitive to the cosmological parameters one wishes to constrain. In this analysis, we have only considered the simplified situation where the redshift-space distortions act along one axis of a Cartesian basis. However, the arguments we have put forward in favour of pair-centre binning do not rely on this assumption, and will remain valid even when wide-angle effects are included in any analysis." }, "1003/1003.2631_arXiv.txt": { "abstract": "We present results of an extensive morphological, spectroscopic, and photometric study of the galaxy population of MACS\\,J0025.4$-$1225 ($z=0.586$), a major cluster merger with clear segregation of dark and luminous matter, to examine the impact of mergers on galaxy evolution. Based on 436 galaxy spectra obtained with Keck DEIMOS, we identified 212 cluster members within 4 Mpc of the cluster centre, and classified them using three spectroscopic types; we find 111 absorption-line, 90 emission-line (including 23 e(a) and 11 e(b)), and 6 E+A galaxies. The fraction of absorption(emission)-line galaxies is a monotonically increasing(decreasing) function of both projected galaxy density and radial distance to the cluster center. More importantly, the 6 observed E+A cluster members are all located between the dark-matter peaks of the cluster and within $\\sim 0.3$Mpc radius of the X-ray flux peak, unlike the E+A galaxies in other intermediate-redshift clusters which are usually found to avoid the core region. In addition, we use Hubble Space Telescope imaging to classify cluster members according to morphological type. We find the global fraction of spiral and lenticular galaxies in MACS\\,J0025 to be among the highest observed to date in clusters at z$>0.5$. The observed E+A galaxies are found to be of lenticular type with Sersic indices of $\\sim2$, boosting the local fraction of S0 to 70 per cent between the dark-matter peaks. Combing the results of our analysis of the spatial distribution, morphology, and spectroscopic features of the galaxy population, we propose that the starburst phase of these E+A galaxies was both initiated and terminated during the first core-passage about 0.5--1Gyr ago, and that their morphology has already been transformed into S0 due to ram pressure and/or tidal forces near the cluster core. By contrast, ongoing starbursts are observed predominantly in infalling galaxies, and thus appears to be unrelated to the cluster merger. ", "introduction": "It is well known that the galaxy population of clusters is dominated by passively evolving, red, early-type galaxies \\citep[e.g.][]{dressler80, bower92, vandokkum98}, unlike the population of galaxies in less dense environments \\citep[e.g.][]{gomez03, lewis02, cooper07}. Many cluster-specific environmental processes, such as ram-pressure stripping \\citep{gunn72,bekki02}, strangulation\\citep{larson80}, galaxy-galaxy harassment \\citep{moore96}, or tidal disruption \\citep{merritt83,byrd90} have been proposed to explain the effects of cluster assembly on the galaxy population. However, the picture of galaxies being accreted in a static cluster environment is overly simplistic. Many optical and X-ray studies \\citep[e.g.][]{jones99, depropris04} reveal that a significant fraction, if not all, of the clusters even at recent epochs are still growing through mergers of galaxy groups and sub-clusters. To get a correct picture of the effects of environment on galaxy evolution, the dynamic processes governing the formation of these large-scale structures must be taken into account. For instance, the predominance of passively evolving galaxies in clusters was suggested to have its roots in the preprocessing of galaxies in the group environment \\citep[e.g.][]{zabludoff98,kodama01, tran09,wilman09}, thus not being directly related to the much denser environment in clusters. Although environment has thus been identified to play a critical role in the evolution of galaxies, the effect of events that cause the most dramatic change in environment, cluster mergers, is still poorly understood \\citep{girardi02}. For instance, it is still controversial whether star-formation activity would be enhanced during cluster mergers, as favored by numerical simulations \\citep[e.g.][]{bekki02,kronberger08}. Observationally, this picture is supported by the classical studies of \\citet{caldwell93} and \\citet{caldwell97}, which find evidence of recent star bursts in merging clusters. More recently, \\citet{owen05} suggested that the excess of radio-loud galaxies around the cluster Abell~2125 may be related to the ongoing cluster merger, and \\citet{ferrari05} report on actively star-forming galaxies in the region between the two sub-clusters of the merging cluster Abell~3921 ($z=0.094$). In addition, \\citet{hwang09} compare the two merging clusters Abell~168 and Abell~1750, and attribute the different amount of galaxy activity, including star formation and nuclear activity, in the region between the subclusters to the differences in the merger history of the two clusters. On the other hand, the ages of numerous post-starburst galaxies found in Abell~3921 are too old for the starbursts to have been triggered by the ongoing merger \\citep{ferrari05}. Also, a 24$\\mu$m observation of ClJ0152.7$-$1357 \\citep[$z\\sim0.83$, ][]{ebeling00} by \\citet{marcillac07} suggests that most of the mid-infrared luminous galaxies are infalling and thus not related to the merging process. The cause of the conflicting observational evidence may lie in the complexity of the dynamics of cluster mergers. For instance, the specific phase in which a merger is observed, the viewing angle, and the relative masses of the subclusters can all be expected to have a direct effect on the triggering of star formation and the resulting spatial and spectral observational data \\citep[see e.g.][]{hwang09}. Finding merging systems whose geometry and merger history are immediately evident is thus of paramount importance. To date, only three such systems have been identified through a clear segregation between dark matter and intra-cluster gas: the Bullet Cluster \\citep[1E0657-56, $z=0.296$;][]{clowe04,clowe06}, Abell~520 \\citep[$z=0.201$;][]{mahdavi07}, and MACS\\,J0025.4$-$1225 \\citep[$z=0.586$, hereafter MACSJ0025;][]{bradac08}. Among these three systems, MACSJ0025 stands out as the best target for a study of the effects of a major cluster merger because the masses of the two sub-clusters are similar\\footnote{Note, however, that the supersonic head-on collision of two systems of very different mass in the Bullet Cluster enabled a detailed study of the effects of ram pressure from supersonic gas during a cluster merger \\citep[see][]{chung09}.}. In this paper, we will report spectroscopic and morphological results for galaxies around MACSJ0025, a cluster discovered in the MAssive Cluster Survey \\citep{macs,ebeling07}, and, at L$_{X,Chandra} = 8.8\\pm 0.2\\times~10^{44}$~ergs~s$^{-1}$, one of the most X-ray luminous clusters at z$>0.5$ \\citep{ebeling07}. Our paper is structured as follows. \\S\\ref{sec:dynamics} summarizes the dynamical structure of MACSJ0025 as studied by \\citet{bradac08}. \\S\\ref{sec:data} gives an overview of the observation and data reduction procedures. \\S\\ref{sec:analysis_technique} describes the analysis of photometric, spectroscopic and morphological data. In \\S\\ref{sec:result}, we present our results. Finally, we discuss a simple merger model in \\S\\ref{sec:discussion} and provide a brief summary in \\S\\ref{sec:summary}. Throughout this paper, we adopt the concordance $\\Lambda$CDM cosmology with $h_0=0.7$, $\\Omega_{\\Lambda}=0.7$, $\\Omega_m = 0.3$. Magnitudes are quoted in the AB system. \\begin{figure*} \\epsfxsize=0.8\\textwidth \\epsffile{paper_plane_combine.eps} \\caption[f1.eps]{Spatial distribution of galaxies observed spectroscopically: small dots mark the location of all sources with measured spectrum, large dots mark confirmed cluster members. The overlaid contours show the projected galaxy density as computed by \\textit{asmooth} \\citep[ ][ see \\S\\ref{sec:phot_galdens} for details]{asmooth}. The red boundary indicates the outline of all DEIMOS masks and defines the study region for this paper. The blue dash square indicate the HST ACS image mentioned in \\S\\ref{sec:data_image_hst}, and defines the area for the morphology analysis in \\S\\ref{sec:morph}. The two plus signs show the location of gravitational lensing peaks \\citep{bradac08}, and the cross sign shows the location of the X-ray flux peak.\\label{plane}} \\end{figure*} ", "conclusions": "\\label{sec:dynamics} MACSJ0025 is the second cluster that shows a significant segregation between the hot gas traced by the X-ray surface brightness, and the mass traced by the gravitational shear \\citep{jeyhan08,bradac08} in a simple two-body merger configuration. We here briefly summarize the dynamics of the system as presented by \\citet{bradac08}. Like many other merging clusters, MACSJ0025 features a galaxy distribution in redshift space as well as in projection on the sky which, taken alone, does not show obvious evidence of substructure. Its galaxy redshift distribution is consistent with a Gaussian centered at z$_{cl}=0.5857$ with $\\sigma_{cl} = 835^{+58}_{-59}$ km~s$^{-1}$ with only mild evidence of a bimodal distribution\\footnote{In \\S\\ref{sec:data_spect_cat}, we discuss updates to the spectroscopic sample; the results are similar. } . In addition, the Dressler-Shectman test \\citep{dstest} detects only marginal cluster substructure. The redshift difference between the two sub-clusters is small ($\\Delta z=0.0005\\pm0.0004$), which implies that the two subclusters are colliding almost in the plane of the sky (within $5\\degr$). The X-ray temperature of the intracluster medium (ICM) within a circular region of radius $660$~kpc centered on the X-ray peak is $6.26^{+0.50}_{-0.41}$~keV; the metallicity of the ICM is $Z=0.37\\pm0.10$ solar. The corresponding X-ray mass of the gas is $0.55\\pm0.06 \\times 10^{14}$~M$_{\\sun}$. The virial radius for a cluster with the same gas temperature is then $\\sim1.3$~Mpc, calculated following the prescription provided by \\citet{arnaud02}. Finally, the radius at which a Milky-way-like galaxy would lose all its gas due to ram-pressure stripping is $\\sim0.7$~Mpc (see \\cite{ma08} for details). The gravitational lensing analysis of \\citet{bradac08} shows the two peaks of the mass distribution being offset from the X-ray peak by $0.82\\arcmin$ for the southeastern peak and $0.50\\arcmin$ for the northwestern peak. Since the X-ray centre lies close to the projected line connecting the gravitational lensing peaks, the separation of the two lensing peaks is thus $1.32\\arcmin$ ($0.52$~Mpc). The masses estimated from the lensing model are $2.5^{+1.0}_{-1.7}\\times10^{14}$~M$_{\\sun}$ and $2.6^{+0.5}_{-1.4}\\times10^{14}$~M$_{\\sun}$ for the southeastern and northwestern peak, respectively. \\subsection{A simple model}\\label{sec:toymodel} Because of its simple merger geometry (major equal-mass merger, in the plane of the sky, after first core passage), MACSJ0025 provides us with an excellent opportunities to study the effect of cluster mergers on fundamental galaxy properties, specifically their morphology and star formation activity. Our understanding of the merger history and dynamics of MACSJ0025 is depicted schematically in Fig.~\\ref{toymodel}. The key phases and physical processes at work can be described as follows: \\begin{description} \\item[Phase A:] pre-contact; no significant galaxy-gas interaction yet; relative galaxy velocities too high for galaxy mergers, but increasing probability of harassment events \\item[Phase B:] first contact; increased ICM interaction induces shocks; onset of segregation between dark and luminous matter; ram-pressure stripping and tidal interactions begin as galaxies encounter steep gradients in gravitational potential and ICM density; morphological tranformation and triggered star formation \\item[Phase C:] core passage; significant ICM interaction and shock heating; non-collisional components (dark matter, galaxies) proceed as gas is left behind; galaxies on the leading edge of either cluster now fully stripped of their gas content; galaxies at the trailing edge experience Phase B effects \\item[Phase D:] present stage; dark matter cores fully separated but single ICM component concentrated at global centre of mass; gas-galaxy interactions (triggered star formation, stripping) continue as trailing galaxy populations pass through core. \\end{description} Qualitatively, this picture is supported by the result of numerical simulation of ram-pressure stripping \\citep[e.g.][]{fujita99b,kronberger08,bekki09} and tidal interaction \\citep{byrd90, bekki99, gnedin03a, gnedin03b}. Although the environments encountered in this process by galaxies at the leading and trailing edges of the individual clusters are roughly symmetrical, galaxies on the far side of the merger experience a hostile environment for longer since they encounter strong tidal fields (passage through the dark matter core) and ram-pressure stripping sequentially; in addition, they pass through the more massive and denser post-collision gaseous core. We note that the velocities of galaxies relative to their respective cluster core are not negligible, and hence their observed relative locations may differ substantially from the ones prior to collision. Exceptions are, however, the E+A and the e(a) populations, both of which will remain reasonably faithful tracers of environment thanks to the relatively short duration of the star-formation and quenching processes. In the following, we will examine whether the features observed in the galaxy population of MACSJ0025 support the picture described above. \\begin{figure} \\epsfxsize=0.47\\textwidth \\epsffile{paper_CMD.ps} \\caption[f15]{Color-magnitude diagram for galaxies of different spectral type. The the symbols are the same as in Fig.~\\ref{galxdensity}. The E+A galaxies (green filled circles) are fainter than $\\rm m_{R_c}>22.0$ and are located mostly in the ``green valley\" below the red-sequence marked by the absorption-line galaxies (red open circles). Compare to the E+A galaxies, the e(a) galaxies (cyan-squares) are brighter and exhibit a broad range of color in ``blue cloud\". \\label{cc}} \\end{figure} \\begin{figure} \\epsfxsize=0.47\\textwidth \\epsffile{toymodel.eps} \\caption[f16]{A schematic illustration of the collision process. Red circles mark the centres of the dark-matter distributions of the two clusters, while light blue and black circles represent the distribution of the ICM and the cluster galaxy population, respectively. See text for discussion. \\label{toymodel}} \\end{figure} \\subsection{Distribution of E+A galaxies} As discussed in \\S\\ref{sec:spect_densdistr} and \\S\\ref{sec:spect_spatdistr}, we find all but one of the E+A galaxies in MACSJ0025 to be located near the peak of the X-ray surface brightness, and thus between the dark-matter cores of the individual clusters. This is exactly what would be expected in the scenario outlined in \\S\\ref{sec:toymodel} and depicted in Fig.~\\ref{toymodel}. The merging process in MACSJ0025 thus appears to cause dramatic and rapid galaxy evolution (the triggering and termination of startbusts) in an environment where it is rarely observed, namely the region of highest gas density in the very cluster core. We should mention in this context that E+A galaxies have in fact been detected in cluster cores before. In all cases, however, projection effects seem to play a major role. Specifically, the group of starburst galaxies found in Abell~851 \\citep[see][and references there in]{oemler09} forms a group of high peculiar velocity falling into the cluster along the line of sight. By contrast, no evidence of projection effects is found for MACSJ0025 (see \\S\\ref{sec:spect_zdistr}). \\subsection{Morphological transformation} In the merger scenario discussed above, the galaxy populations of both clusters should have been exposed to conditions conductive to trigger star formation and ram-pressure stripping for a large fraction of the duration of the cluster collision. As a result, one would expect the morphological transformation from late-type to S0 galaxies to have been particularly efficient in this system in recent times. Indeed, we measure a high global S0 fraction of about $0.3$ (Fig.~\\ref{morph_fraction} and Table~\\ref{table_morph}), among the highest found to date in any cluster at $z>0.5$, but still within the observed range of 0.05--0.33 \\citep{poggianti09b}. Furthermore, the high global S fraction of about $0.39$ may reflect that cluster merger help the mix of late-type galaxies at outskirt. A more dramatic increase in the global S0 and S fraction would have been surprising, given that the morphological mix of galaxies in clusters is constantly driven to a global average as S0s continue to evolve into ellipticals, and late-type galaxies are being accreted from the field. A snapshot of the merger-driven transformation process is, however, provided by our view of the core of MACSJ0025 where we find all E+A galaxies to be lenticulars, leading to a local S0 fraction of about 70\\%. Our hypothesis that interactions with the gravitational potential and the intracluster medium are primarily responsible for the accelerated evolution of galaxies in this massive merger is also supported by a comparison with the morphology of E+A galaxies in the field. Signature of tidal destruction or galaxy merging can be removed efficiently and quickly by both ram pressure \\citep{kapferer08} and tidal forces \\citep{bekki01c} in the cluster core. \\begin{figure} \\epsfxsize=0.47\\textwidth \\epsffile{paper_d4000hd2.eps} \\caption[f17]{The D$_n$(4000) and EW of H$\\delta$ of galaxies in the cluster; cyan for e(a) galaxies, green for E+A galaxies, magenta for absorption-line galaxies in the regions farther from the cluster cores along merger axis, and red for absorption-line galaxies in region 3 of Fig.~\\ref{hstcentre}. The four curves represent stellar-population synthesis models from \\citet{BC03}, assuming a delta-function starburst in a 5~Gyr-year-old galaxy with exponentially decaying SFR ($\\tau=1Gyr$). From top to bottom, the curves are for models with starburst fractions of $100\\%$, $30\\%$, $10\\%$, and $5\\%$. The labels along the curves denote the age of the burst in Gyr. \\label{D4000}} \\end{figure} \\subsection{Temporal evolution} We can use estimates of the duration of characteristic phases of galaxy evolution derived from spectral diagnostics for a consistency check of our simple merger model (Fig.~\\ref{toymodel}). Specifically, stellar synthesis models \\citep{BC03} predict how the strength of the D$_n$(4000) break and the EW of H$_{\\delta}$ vary along an evolutionary track from the beginning of the starburst to several Gyrs after. In Fig.~\\ref{D4000}, the data for galaxies of the various spectral types observed in MACSJ0025 are overlaid on the theoretically expected tracks\\footnote{The models assume a delta-function starburst in a 5~Gyr-year-old galaxy with exponentially decaying SFR ($\\tau=1Gyr$), Solar metalicity, and a Salpeter initial mass function.}. We note that these models can only be used to describe the absorption features in the stellar spectra; emission-lines from the interstellar medium are not included. While this makes the model predictions less meaningful for emission-line galaxies, they should be applicable to the E+A galaxies that we are primarily interested in. Along the track, the magnitude-weighted stellar age \\citep{kauffmann03} from the models agrees qualitatively with the location of the various galaxy subpopulations. If the E+A population observed near the peak of the X-ray emission was created by the equal-mass merger in MACSJ0025, these galaxies' approximate age of 0.5--1 Gyr (Fig.~\\ref{D4000}) would provide a constraint on the time elapsed since core passage (Phase C). As interestingly, in the context of our merger model, is that the regions farther from the cluster core along the merger axis are dominated by young absorption-line galaxies in which the star formation ended about $1-2$~Gyrs ago. A high fraction of young passively evolving galaxies along the merger path is consistent with our model in which the transformation of these galaxies would have begun in phase A, i.e.\\ at the onset of the cluster collision. These time scales are consistent with the predictions from numerical simulations \\citep{barnes01} which show that the relative velocity of the two free-falling dark-matter halos is significantly reduced by dynamical friction during a head-on merger. For reference, Fig.~\\ref{D4000} also shows the location of e(a) galaxies to demonstrate, qualitatively, that their stellar ages derived from this H$_{\\delta}$-D$_n$(4000) method are younger than those of the E+A galaxies. A more quantitative analysis of the e(a) galaxies' locus in this graph would require more complex models which are beyond the scope of this paper. Finally, and not surprisingly, we find the absorption-line galaxies in the two sub-cluster core (red symbols in Fig.~\\ref{D4000}) to be the oldest population, consistent with the expectations that the star-formation epoch in elliptical galaxies at center of ``proto''-clusters ends as early as z$\\sim 2$\\citep[e.g. ][]{blakeslee06} . \\subsection{Relation between starburst and cluster merger} Simulations suggest increased star-formation activity during cluster mergers \\citep[e.g.][]{bekki99}. However, most of the observational results \\citep[e.g.][]{marcillac07} for starburst galaxies in clusters are interpreted as being related to infall. The spatial distribution of e(a) galaxies in MACSJ0025 (\\S\\ref{sec:spect_spatdistr} and \\S\\ref{sec:spect_densdistr}) seems to support the infall scenario, in that most of the starburst galaxies are located at the cluster outskirts. This is consistent with the view that starbursts in e(a) galaxies are caused by galaxy-galaxy interactions which are highly improbable at the high relative velocities of galaxies closer to the cluster core. Nevertheless, evidence of starbursts near the X-ray centre is observed in the form of the quenched starbursts in E+A galaxies. The physical mechanism that caused these bursts is likely to be the same that ultimately quenches them, creating the E+A phase of galaxy evolution, namely the violent interaction with the dense intracluster medium encountered during the merger (b and c, in Fig.~\\ref{toymodel}). The fact that no starburst or post-starburst galaxies are found in the regions farther from the cluster core, but still well within the viral radius, suggests that extreme gas densities are required to trigger bursts, which is consistent with the few instances in which such dramatic gas-galaxy encounters have actually been observed (Cortese et al.\\ 2007; Richard et al.\\ 2010). This still leaves us with the question, however, of why no active starbursts are presently observed near or within the newly formed gaseous core (phase d of Fig.~\\ref{toymodel}). We argue that galaxies from either cluster approaching the core along the merger axis, and thus encountering a similarly hostile environment as the galaxies at the leading edge of either system during first approach, do in fact experience bursts of star formation and will, ultimately, also go through an E+A phase. Their density, however, and thus the number of such ICM-triggered starbursts decreases rapidly as the merger proceeds and only galaxies from the increasingly sparsely populated outer cluster regions continue to approach. \\begin{itemize} \\item We present results of a extensive morphological, spectroscopic, and photometric study of the galaxy population of MACS\\,J0025.4$-$1225, an equal-mass cluster merger, based on 212 galaxy spectra obtained with Keck DEIMOS as well as colour imaging of the cluster core with HST/ACS and wide-field imaging in seven passbands from Subaru SuprimeCam and CFHT/Megacam and CFHT/WIRCAM observations. \\item Six E+A galaxies are found in MACSJ0025, all of them near the cluster center defined by the X-ray surface brightness peak, in contrast to other studies which find the E+A population to avoid the cluster core. The color of these E+A galaxies fall in between those of red-sequence galaxies and emission-line galaxies. Compared to E+A galaxies in other clusters at intermediate redshift, the luminosity function of E+A galaxies in MACSJ0025 appears truncated (no galaxies with M$_R < -20.7$). Their morphology is broadly lenticular, i.e., characterized by a dominant bulge and relatively flat surface-brightness profiles (Sersic index$\\sim 2$). In addition, no evidence of tidal features or other asymmetric features is found. \\item The fraction of absorption(emission)-line galaxies is a smooth increasing (decreasing) function of projected galaxy density. A similar correlation is found with distance to the cluster center. The fraction of dusty-starburst galaxies (e(a)) is dramatically lower inside the virial radius than outside. \\item The global fractions of morphological types (E, S, S0) of galaxies in MACSJ0025 are consistent with those in clusters at $z>0.5$ compiled by \\citet{poggianti09b}, although the fraction of S0 and S types observed in MACSJ0025 are the highest among the sample. While the global S0 fraction is thus not dramatically enhanced, we find the local value boosted around the X-ray surface brightness, coinciding with the excess of E+A galaxies. \\item We propose a simple model in which star-formation activity and morphological transformations are caused by dynamic interactions enhanced and accelerated by the cluster merger. In this picture, the time since the termination of the starburst phase in the observed E+A galaxies naturally matches the time elapsed since the core passage of the merger; about 0.5--1.0 Gyrs ago in the case of MACSJ0025. Additional evidence for the causal role of the merger is provided by the relatively young age of the absorption-line cluster members near the leading edges of the two clusters engaged in the collision. \\end{itemize}" }, "1003/1003.3807_arXiv.txt": { "abstract": "{The transition from the L to the T spectral type of brown dwarfs is marked by a very rapid transition phase, remarkable brightening in the \\emph{J}-band and a higher binary frequency. Despite being an active area of inquiry, this transition regime still remains one of the most poorly understood phases of brown dwarf evolution.} { We resolved the L dwarf 2MASS\\,J03105986+1648155 for the first time into two almost equally bright components straddling the L/T transition. Since such a co-eval system with common age and composition provides crucial information of this special transition phase, we monitored the system over $\\sim$\\,3 years to derive first orbital parameters and dynamical mass estimates, as well as a spectral type determination.} {We obtained resolved high angular resolution, near-IR images with HST and the adaptive optics instrument NACO at the VLT including the laser guide star system PARSEC.} {Based on two epochs of astrometric data we derive a minimum semi-major axis of 5.2\\,$\\pm$\\,0.8 AU. The assumption of a face-on circular orbit yields an orbital period of 72\\,$\\pm$\\,4 years and a total system mass of $\\sim$ 30\\,-\\,60\\,$M_\\mathrm{Jup}$. This places the masses of the individual components of the system at the lower end of the mass regime of brown dwarfs. The achieved photometry allowed a first spectral type determination of L9\\,$\\pm$\\,1 for each component. In addition, this seems to be only the fifth resolved L/T transition binary with a flux reversal. } {While ultimate explanations for this effect are still owing, the 2MASS\\,J03105986+1648155 system adds an important benchmark object for improving our understanding of this remarkable evolutionary phase of brown dwarfs. Additionally, the observational results of 2MASS\\,J03105986+1648155\\,AB derived with the new PARSEC AO system at the VLT show the importance of this technical capability. The updated AO system allows us to significantly extend the sample of brown dwarfs observable with high-resolution from the ground and hence to reveal more of their physical properties.} ", "introduction": "\\label{LT_transition_intro} The transition from the L to the T spectral types of brown dwarfs is marked by a dramatic change in their near-IR spectral energy distribution (SED) and atmospheric properties. While this has already been an active area of inquiry, it still remains one of the most poorly understood phases of brown dwarf evolution. As discussed by e.\\,g. \\citet{Geballe02}, the late-type L dwarfs are characterized by very red near-IR colors, caused by condensate dust in their photospheres and metal hydrides, as well as CO absorption bands. In contrast, the T dwarfs are characterized by again bluer near-IR colors, due to the appearance of CH$_4$ absorption at 1.65 and 2.2 $\\mu$m, stronger H$_2$O absorption and the increasing importance of collision-induced H$_2$ absorption (CIA), as well as relatively dust-free photospheres \\citep{Geballe02}. This change occurs over a comparatively narrow effective temperature range ($\\Delta$\\,T$_{\\mathrm{eff}} \\approx$ 200\\,K) around 1500\\,-\\,1300\\,K for near-IR L7\\,-\\,T3 dwarfs \\citep{Goli04_2}, implying a very rapid transition phase. Taking into account the interaction between temperature, gravity, metallicity and the physics of atmospheric dust clouds, this area remains a challenge to theoretical models (for different possible explanations see e.\\,g. \\citealp{Knapp, Tsuji05, Tsuji99, Marley02, Burrows06, Ackerman01, Burgasser02, Folkes07}). \\begin{table*}[t!] \\centering \\caption{Observation log of high-angular resolution imaging of 2MASS\\,031059+164815\\,AB} \\label{ObsLog_2M0310} \\centering \\begin{tabular}{c c c c c c} % \\noalign{\\smallskip} \\hline\\hline \\noalign{\\smallskip} Date & Telescope/Instrument & Filter & Exp.\\,time & Seeing\\,$^{\\mathit{a}}$ & Strehl ratio\\\\ & & & [sec] & [ $\\arcsec$ ] & [ \\% ] \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 24\\,/\\,09\\,/\\,2004 & HST/NIC1 & F108N & 2560 & \\\\ & & F113N & 2816 & \\\\ \\noalign{\\smallskip} 05\\,/\\,11\\,/\\,2007 & VLT/NACO& H & 14 x 60 & 1.02\\,-\\,1.19 & 21.9\\,-\\,43.8\\\\ & & K$_\\mathrm{S}$ & 14 x 60 & 0.99\\,-\\,1.07 & 42.2\\,-\\,68.0 \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{\\mathit{a}}$] Site seeing measured by the differential image motion method (DIMM) in \\emph{V}-\\,band at zenith \\end{list} \\end{table*} Another very peculiar, yet unexplained observational feature is the remarkable brightening in the \\emph{Z}/\\emph{Y} ($\\sim$ 0.9\\,-\\,1.1 $\\mu$m) and \\emph{J} ($\\sim$ 1.2\\,-\\,1.3 $\\mu$m) bands of up to $\\Delta\\,M_{J} \\sim$\\,1\\,mag for the early- to mid-type T dwarfs. This so-called \\emph{J}\\,-\\,band ``bump\" (\\citealt{Dahn, Tinney03, Vrba04}) indicates a significant flux redistribution at almost constant luminosity. In the following, high-resolution imaging surveys revealed a binary frequency among the L\\,/\\,T transition objects almost twice as high as in earlier or later type brown dwarfs \\citep{Burgasser06_1}. In a first attempt, it was suggested that the ``bump\" might be artificially enhanced by systems appearing overluminous due to binarity (``crypto-binarity\") and that the integrated light of an L\\,+\\,T dwarf system could mimic the spectral characteristics of an early\\,-\\,type T dwarf (\\citealt{Burrows06, Liu06, Burgasser06_1}). However, recent discoveries suggest that at least a fraction of the observed \\emph{J}\\,-\\,band brightening is intrinsic to the atmospheres of early- to mid\\,-\\,type T dwarfs as they cool. Resolved high-resolution photometry revealed in four L\\,/\\,T dwarf binary systems a 1.0\\,-\\,1.3 $\\mu$m flux reversal, with the T dwarf secondary being brighter than the late L or early T dwarf primary in this wavelength regime (\\object{2MASS J17281150+3948593}\\,AB, \\citealt{Gizis03}; \\object{SDSS J102109.69-030420.1}\\,AB, \\citealt{Burgasser06_1}; \\object{SDSS J153417.05+161546.1}\\,AB, \\citealt{Liu06}; \\object{2MASS J14044941-3159329}\\,AB, \\citealt{Looper08_2}). In addition, a comparison with absolute \\emph{J}\\,-\\,band magnitudes of other resolved binary components having a spectral type of T1\\,-\\,T5 (e.\\,g. $\\epsilon$\\,Indi\\,Ba or SDSS\\,J042348-041403\\,B) shows that they are still $\\sim$ 0.5 mag brighter than the latest L dwarfs (\\citealt{Burgasser06_1, Looper08_2}). These findings imply that the brightening across the L\\,/\\,T transition is a real effect, since it also affects binaries which are assumed to be coeval systems with common age and metallicity. Therefore, further discoveries and high-resolution observations of L\\,/\\,T transition binaries will play an important role. An extended sample of L\\,+\\,T dwarf binaries should provide independent crucial information on these issues. The comparison of the observed properties with theoretical models (e.\\,g.\\,\\citealt{Baraffe03, Burrows03, Saumon08}) will then help to reveal the physical mechanism that drives the transition from dusty L dwarfs to dust-free T dwarfs. One such important newly-resolved binary is \\object{2MASS J03105986+1648155} (hereafter 2M0310+1648), whose likely coeval components straddle the L\\,/\\,T transition. 2M0310+1648 was originally discovered by \\citet{Kirk00} in the \\emph{Two Micron All Sky Survey} (2MASS) database. It was classified with a spectral type L8 in the optical and the presence of lithium implied a mass $M \\le$ 60\\,$M_\\mathrm{Jup}$, confirming its brown dwarf nature. The first near-IR spectroscopic observations revealed some spectral discrepancy compared to other late-type L dwarfs in terms of a significantly depressed \\emph{K} -\\,band spectrum starting around 2.2\\,$\\mu$m, interpreted to be caused by collision-induced H$_2$ absorption \\citep{Reid01_2}. By contrast, \\citet{Nakaj01} attributed this flux suppression to methane rather than H$_2$, which would also explain the very weak CO band head at 2.3\\,$\\mu$m and indicate that much of the carbon is in CH$_4$. Finally, \\citet{Geballe02}, with their new classification scheme slightly revised the spectral type of 2M0310+1648 to L9 in the near-IR. In September 2004, 2M0310+1648 was resolved as an almost equally bright binary system during our own \\emph{Hubble Space Telescope} (HST) NICMOS survey (\\citealt{Stumpf05}; Stumpf et al. submitted). As an L\\,/\\,T transition binary it added up to the already apparently higher binary fraction in the transition regime and provides an important testbed to derive information about the underlying physical and chemical processes in this transition regime. Therefore a monitoring program including resolved photometry and spectroscopy was started and the first results are presented in this paper. ", "conclusions": "HST/NICMOS imaging in the F108N and F113N filters revealed the binary nature of another very interesting L/T transition brown dwarf: 2M0310+1648\\,AB. In the following, second epoch astrometry and first resolved high-resolution photometry in the \\emph{H}\\,- and \\emph{K$_\\mathrm{S}$}\\,-\\,bands were obtained with VLT/NACO and its new LGS AO system PARSEC. The two epochs of astrometric measurements spanning $\\sim$ 3 years allowed for first rough orbital parameter estimations. Due to a non-significant change in the separation, a face-on circular orbit was assumed, yielding an orbital period of 72\\,$\\pm$\\,4 years. Depending on the assumed semi-major axis, Kepler's third law yielded a first total system mass estimate of $\\sim$ 30\\,-\\,60\\,$M_\\mathrm{Jup}$, placing the individual component masses at the lower end of the brown dwarf regime. The first orbital period estimate of $\\sim$ 72 years does not suggest the possibility for a meaningful dynamical mass determination on a short time scale. Nevertheless, follow-up observations in the next years will allow us to derive more accurate information on the orbital elements and hence the true orientation of the system in space, as well as the true orbital period. This will finally enable us to better constrain the total system mass. The derived photometry revealed a very intriguing property of 2M0310+1648\\,AB. The component fluxes show an unexpected decrease in brightness difference with increasing wavelength, resulting in a marginal flux reversal in the \\emph{K$_\\mathrm{S}$}\\,-\\,band. An additional comparison of the component colors obtained reveals a redder color of the B component. These results indicate that the designated primary component 2M0310+1648\\,A might actually be of slightly later spectral type than 2M0310+1648\\,B. This could at least partly explain the observed flux reversal as part of the \\emph{J}\\,-\\,band brightening of early- to mid-type T dwarfs, but a full explanation for the true nature of the reversal is still owing. Upcoming spatially resolved spectroscopic observations with VLT/SINFONI and the PARSEC AO system will allow a precise spectral type determination and an investigation of the underlying spectral morphologies. If it turns out that 2M0310+1648\\,A is really of later spectral type than 2M0310+1648\\,B, the system would add up to the currently small sample of flux reversal binaries. Additionally, 2M0310+1648\\,AB would be the first binary with a secondary showing the \\emph{J}\\,-\\,band brightening already at the very late\\,-\\,L (L9) or very early-T (T0) dwarf stage rather than at a T1.5 spectral type or later. This would challenge the existing theoretical models even more. In future work, the likely coeval system 2M0310+1648\\,AB will serve as a very important benchmark object in the L\\,/\\,T transition. Further high-resolution observations will provide an improved understanding of and new insights into the challenging picture of this still poorly understood, yet remarkable evolutionary phase of brown dwarfs." }, "1003/1003.1496_arXiv.txt": { "abstract": "Ultrahigh energy cosmic ray (UHECR) protons interacting with the cosmic microwave background (CMB) produce UHE electrons and gamma-rays that in turn initiate electromagnetic cascades on CMB and infrared photons. As a result, a background of diffuse isotropic gamma radiation is accumulated in the energy range $E\\lsim 100$\\,GeV. The Fermi-LAT collaboration has recently reported a measurement of the extragalactic diffuse background finding it less intense and softer than previously measured by EGRET. We show that this new result constrains UHECR models and the flux of cosmogenic neutrinos. In particular, it excludes models with cosmogenic neutrino fluxes detectable by existing neutrino experiments, while next-generation detectors as e.g.\\ JEM-EUSO can observe neutrinos only for extreme parameters. ", "introduction": "% The origin of ultrahigh energy cosmic rays (UHECRs) is not yet established despite more than 50~years of research. Natural candidates as UHECR primaries are extragalactic protons from astrophysical sources. In this case, interactions of UHE protons with the cosmic microwave background (CMB) leave their imprint on the UHECR energy spectrum in the form of the Greisen-Zatsepin-Kuzmin (GZK) cutoff and a pair-production dip~\\cite{GZK}. The GZK cutoff is a steepening of the proton spectrum at the energy $E_{\\rm GZK} \\approx (4 - 5)\\times 10^{19}$~eV, caused by photo-pion production on the CMB. Such a steepening has been observed by the HiRes~\\cite{Abbasi:2007sv} and the Auger collaboration~\\cite{Abraham:2008ru}, but its real cause is still unclear. An immediate consequence of the dominance of extragalactic protons in the CR flux and their interaction with CMB photons is the existence of ultrahigh energy (``cosmogenic'') neutrinos produced by charged pion decays, as suggested first in Ref.~\\cite{BZ}. Another signature for extragalactic protons is a pair-production dip \\cite{BG88,BGG-dip} in the CR flux around $5\\times 10^{18}$\\,eV, which is clearly seen in the experimental data. Photons and positrons from pion decay and $p + \\gamma_{\\rm CMB} \\to p + e^+ + e^-$ pair-production initiate electromagnetic cascades on photons from the CMB and the extragalactic background light (EBL), dumping all the energy injected into cascade particles below the pair-production threshold at $\\sim 100$\\,GeV. Clearly, the production of neutrinos by UHE protons is thus intimately tied to the one of photons and electrons, and both depend in turn on the flux of primary cosmic rays. While UHE photons and electrons start electromagnetic cascades by scattering on photons from the EBL, neutrinos reach us suffering no collisions. Therefore, the measurement of the diffuse extragalactic gamma-ray background (EGRB) can be used to impose a strict upper limit on the possible diffuse high energy neutrino flux, as suggested first in Ref.~\\cite{cascade}. We derive in this work an upper limit on the flux of cosmogenic neutrinos assuming that the primary UHECR particles are protons. In case that all primaries or part of them are nuclei, the cosmogenic neutrino flux is lower than for a pure proton composition~\\cite{nuclei}. Thus our assumption of a pure proton composition is justified, since we aim at deriving an {\\em upper limit\\/} on the cosmogenic neutrino flux. Note that the HiRes data~\\cite{Abbasi:2007sv} agree with a pure proton composition at $E \\gsim 1\\times 10^{18}$~eV, while the mass composition deduced from the Auger data indicates the presence of heavier nuclei in the primary UHECR flux~\\cite{Abraham:2008ru}. In the latter case, the maximally allowed cosmogenic neutrino flux would be below the upper limit derived for a pure proton primary flux in this paper. In the present work, we use a recently reported measurement~\\cite{fermi} of the EGRB by Fermi-LAT to constrain UHECR models. We show that the observed fast decrease of the EGRB with energy, $J(E)\\propto E^{-2.41}$, already constrains such models. In particular, versions of the dip model with strong redshift evolution contradict the Fermi data, while this model without or with weak redshift evolution remains viable. Moreover, the Fermi data allows us to derive a strong upper limit on the diffuse UHE neutrino flux. As a result, we conclude that the detection of cosmogenic neutrinos requires to increase the sensitivity of UHE neutrino experiments compared to current levels. As it is demonstrated below, the maximal energy density of cascade radiation $\\omega_{\\rm cas}^{\\max} \\approx 5.8\\times 10^{-7}$~eV/cm$^3$ allowed by the Fermi-LAT data can be used to select viable UHECR models without explicitly calculating electromagnetic cascade processes. ", "conclusions": "% We have used a recent measurement of the EGRB by Fermi-LAT to constrain models for UHECR and cosmogenic UHE neutrinos and to demonstrate that the latter are not detectable with the present experimental sensitivity. Both the dip and ankle model without or with weak evolution are consistent with the Fermi-LAT measurement of the EGRB. The cosmogenic neutrino flux is strongly limited by the new upper cascade bound and undetectable for a conservative choice of parameters by Auger-North and JEM-EUSO. Only for an extreme set of parameters, $E_{\\max} \\gsim 1\\times 10^{22}$~eV and $\\omega_{\\rm cas} \\sim \\omega_{\\rm cas}^{\\max}$, the cosmogenic flux is marginally detectable by JEM-EUSO. To achieve the observation of cosmogenic neutrinos for less extreme parameters, the detection threshold of JEM-EUSO (in the tilted mode) must be lowered down to $1\\times 10^{19}$~eV and the sensitivity of Auger-North should be increased by factor $\\sim 20$ in comparison with Auger-South. The further development of radio-detection methods gives another hope for detection of small fluxes of cosmogenic neutrinos. The results of our paper emphasize the necessity to develop more sensitive methods for the detection of cosmogenic neutrinos." }, "1003/1003.1893.txt": { "abstract": "The atmospheric water vapor content above the Roque de los Muchachos Observatory (ORM) obtained from Global Positioning Systems (GPS) is presented. GPS measurements have been evaluated by comparison with 940nm-radiometer observations. Statistical analysis of GPS measurements points to ORM as an observing site with suitable conditions for infrared (IR) observations, with a median column of precipitable water vapor (PWV) of 3.8 mm. PWV presents a clear seasonal behavior, being Winter and Spring the best seasons for IR observations. The percentage of nighttime showing PWV values smaller than 3 mm is over 60\\% in February, March and April. We have also estimated the temporal variability of water vapor content at the ORM. A summary of PWV statistical results at different astronomical sites is presented, recalling that these values are not directly comparable as a result of the differences in the techniques used to recorded the data. ", "introduction": "The triatomic molecules H$_2$0, CO$_2$ and O$_3$ are the main responsible for reducing the few transparent windows on the infrared atmospheric transmission spectrum, and producing absorption bands that are difficult to correct during the processing of astronomical data. Out of these windows, the atmosphere is opaque and observations at those wavelengths have to be carried out from space. Sites with large atmospheric water vapor levels have narrow and unstable transparent windows. High levels of water vapor reduce the atmospheric transparency, but also increase the thermal infrared background. It is assumed among the astronomical community that sites placed at lower altitudes are less suitable for high-quality IR observations; however, other parameters such the troposphere thickness might be also playing an important role \\citep{2004SPIE.5572..384G}. Indeed, the evaluation of the IR quality at the Observatorio del Teide in the Canary Islands, Spain \\citep{1985A&A...150..281M} and the comparison with prediction models \\citep{1992AJ....104.1650C} indicate that El Teide site, at $\\sim$2400 m above sea level, might be as good as Mauna Kea (4200 m) for IR astronomy in the 1-5 $\\mu$m window \\citep{1998NewAR..42..533H}. Observational results pointing to a similar conclusion were reported for Roque de los Muchachos observatory (ORM) on the island of La Palma (Canary Islands, Spain), also at $\\sim$2400 m above sea level, when comparing infrared observations from the 3.58m Telecopio Nazionale Galileo with data obtained at the 10m Keck telescope on Mauna Kea astronomical observatory (Hawaii, USA) (http://www.tng.iac.es/news/2003/03/21/nics$\\_$refurbish). There are many parameters accounting for the quality of an astronomical site, namely seeing, cloud cover, ground winds, high-altitude winds, etc (see \\cite{2007RMxAC..31...36M} for a review on the characterization of these parameters at ORM). The water vapor content is an important parameter affecting the IR quality of astronomical sites. Conditions for IR astronomical observations were classified in terms of precipitable water vapor (PWV hereafter) in four divisions \\citep{1998NewAR..42..537K}: 1) Good or excellent $\\rightarrow$ PWV $\\leq$ 3 mm; 2) Fair or mediocre $\\rightarrow$ 3 $<$ PWV $\\leq$ 6 mm; 3) Poor $\\rightarrow$ 6 $<$ PWV $\\leq$ 10 mm; and (4) Extremely poor $\\rightarrow$ PWV $\\geq$ 10 mm . The development of IR instrumentation and the requirements for current large and future extremely large telescopes demand a proper characterization of PWV and statistical studies of large temporal databases (covering years). The fraction of nights with good IR conditions (small column of water vapor) as a function of the epoch of the year will allow an optimal scheduling of telescope observing time. The total atmospheric water vapor contained in a vertical column of unit cross-sectional area extending between any two specified levels is known as precipitable water vapor and it is commonly expressed in terms of the height to which that water substance would stand if completely condensed and collected in a vessel of the same unit cross section \\citep{AMS00}. PWV is also referred to as the total column water vapor \\citep{ferrare02}. Measurements of PWV can be obtained in a number of ways, from {\\it in situ} measurements (radiosondes) to remote sensing techniques (photometers, radiometers, GPS, Imaging Spectroradiometers on satellites, etc). Radiosondes have been the primary {\\it in situ} observing system for monitoring water vapor; however, the operability of radiosondes is limited due to running costs and decreasing sensor performance in cold dry conditions \\citep{2003JGRD..108.4651L}. Usually, radiosondes are expected to measure PWV with an uncertainty of a few millimetres, which is considered to be the standard accuracy of PWV for meteorologists \\citep{2001JAtOT..18..830N}. Near infrared (NIR) radiometers measuring methodology assumes plane-parallel atmosphere, hence it is only satisfactory for measurements near the Zenith. Moreover, their uses are limited to photometric and bright nights (from first quarter to last quarter Moon). Errors of 20\\% or below are estimated for PWV measurements using NIR radiometers in the range 3-10 $\\mu$m, whereas errors in the 0.5-1$\\mu$m range might be as high as 40\\% \\citep{1989PASP..101..441Q}. Microwave water vapor radiometers (WVR) observations are not reliable when liquid water is present on the WVR frequency window \\citep{2001JApMe..40....5L} or when there is significant scattering from liquid water droplets and ice crystals in the field-of-view \\citep{Zhang99}. GPS is an increasingly useful tool for measuring PWV, which has gained a lot of attention in the meteorological community. The GPS procedure to estimate the PWV is based on the fact that the propagation of electromagnetic waves through the atmosphere is drastically affected by variations on the refraction index of the troposphere, which depends on the water vapor pressure, air pressure, and temperature. The consequence is an induced delay in any signal crossing the atmosphere, being due to a combination of a hydrostatic and a water vapor delay. The hydrostatic delay is very stable and has a direct relationship with local atmospheric pressure. The second component, the wet delay, is directly related to the water vapor content in the atmosphere above the site where the measurements are taken. The rapid temporal variations of the water vapor affects its prediction and proper measurement. The overall tropospheric delay at a GPS station allows the estimation of the PWV with a high-degree of accurancy \\citep{1994JApMe..33..379B, Boco06, 2009JGeod..83..537J}, although under very dry environments GPS might underestimate the PWV content \\citep{schneider09}. GPS system provides a better spatial coverage and continuous PWV estimations in comparison with other techniques \\citep{2000GeoRL..27.1915G}. The comparison of GPS, radiosondes and WVR \\citep{2000ITGRS..38..324E, 2001JAtOT..18..830N, 2003JGRD..108.4651L} shows agreement generally at the level of 1-2 mm of PWV, corresponding to 7-13 mm of zenith wet delay (error $<<$ 10\\%). In this paper, we present statistical results on the PWV above the Roque de los Muchachos Observatory (La Palma, Spain) derived from GPS measurements for the period spanning from June 2001 to December 2008. For comparison purposes, we also present PWV statistical results for Mauna Kea site derived also from GPS data. ", "conclusions": "We have analised the water vapor content for the period from June 2001 to December 2008 above the ORM using PWV estimations from GPS data. We have verified the consistency of 940nm-radiometer and GPS estimation of PWV, removing the offset between both techniques. We have also presented statitical results for Mauna Kea for the same period, analysing the GPS close to Mauna Kea site with $\\tau_{225GHz}$ measurements. Our main results and conclusions may be summarized as follows: \\begin{itemize} \\item[(i)] The nighttime median PWV above ORM is 3.79 mm, with slightly differences between day and nighttime statistics. \\item[(ii)] The PWV presents a clear seasonal variation at the ORM. Winter and spring nights present the lower PWV statistics, with a median value of 2.82 mm and 2.86 mm, respectively. \\item[(iii)] More than 60\\% of the nighttime during February, March and April present good or excellent conditions in terms of water vapor (PWV $\\leq$ 3 mm) at the ORM. \\item[(iv)] Comparing PWV statistical results for ORM and Mauna Kea, we deduce that for 10 hours of good conditions for IR observations at Mauna Kea during winter, the ORM will present 7.9 hours showing excellent conditions. \\item[(v)] The average temporal range presenting good or excellent conditions for IR observations (PWV $\\leq$ 3mm) is comparable at ORM and Mauna Kea (16.9 hours and 23.4 hours, respectively). \\item[(vi)] The comparison of PWV at different astronomical sites is difficult because there is not an unique defined technique/procedure to estimate the PWV. \\end{itemize} GPS is a promising technique to unify the PWV estimations at many astronomical sites." }, "1003/1003.0345_arXiv.txt": { "abstract": "{We obtained new Fabry-Perot data cubes and derived velocity fields, monochromatic and velocity dispersion maps for 28 galaxies in the Hickson compact groups 37, 40, 47, 49, 54, 56, 68, 79 and 93. We also derived rotation curves for 9 of the studied galaxies, 6 of which are strongly asymmetric. Combining these new data with previously published 2D kinematic maps of compact group galaxies, we investigate the differences between the kinematic and morphological position angles for a sample of 46 galaxies. We find that one third of the non-barred compact group galaxies have position angle misalignments between the stellar and gaseous components. This and the asymmetric rotation curves are clear signatures of kinematic perturbations, probably due to interactions among compact group galaxies. A comparison between the B-band Tully-Fisher relation for compact group galaxies and that for the GHASP field-galaxy sample shows that, despite the high fraction of compact group galaxies with asymmetric rotation curves, these lie on the Tully-Fisher relation defined by galaxies in less dense environments, although with more scatter. This is in agreement with previous results, but now confirmed for a larger sample of 41 galaxies. We confirm the tendency for compact group galaxies at the low-mass end of the Tully-Fisher relation (HCG 49b, 89d, 96c, 96d and 100c) to have either a magnitude that is too bright for its mass (suggesting brightening by star formation) and/or a low maximum rotational velocity for its luminosity (suggesting tidal stripping). These galaxies are outside the Tully Fisher relation, at the 1$\\sigma$ level, even when the minimum acceptable values of inclinations are used to compute their maximum velocities. The inclusion of such galaxies with v$<$100 km s$^{-1}$ in the determination of the zero point and slope of the compact group B-band Tully-Fisher relation would strongly change the fit, making it different from the relation for field galaxies, a fact that has to be kept in mind when studying scaling relations of interacting galaxies, specially at high redshifts.} {}{}{}{} ", "introduction": "Compact groups of galaxies are environments in which tidal encounters are thought to be common. It is, therefore, expected, that interactions have stripped or disturbed such systems at some level. Mendes de Oliveira et al. (2003) studied the Tully-Fisher (TF) relation for 25 compact group galaxies and highlighted the importance of having 2D velocity fields, derived from the H$\\alpha$ line, in the study of TF of interacting galaxies. They find similar TF relations for compact group and field galaxies and stress that a fine tuning of the kinematic parameters is needed in order to have meaningful rotation curves for interacting galaxies. This is highly relevant for the study of the evolution of the TF relation as a function of redshift, given that at high redshifts the numbers of interacting galaxies raise considerably. We are gathering a large dataset of Fabry-Perot velocity maps for interacting galaxies at low redshifts which allows the study of the effects of the environment on galaxy evolution in high density structures like compact groups of galaxies. With this goal in mind, we studied, in previous papers, the properties of the galaxies in these systems, pointing out interaction indicators for individual galaxies (e.g. Mendes de Oliveira et al. 1998); the evolution of the B-band TF relation (Mendes de Oliveira et al. 2003); the distribution of luminous and dark mass profiles of compact group galaxies (Plana et al. 2010). The present study comes to complement the sample of compact group galaxies with measured velocity maps presented by Plana et al. (1998, 2000, 2003), Mendes de Oliveira et al. (1998), Amram et al. (2003, 2004, 2007) and Torres-Flores et al. (2009). We present, for the first time, velocity maps for galaxies in nine Hickson compact groups (HCGs 37, 40, 47, 49, 54, 56, 68, 79 and 93.) and we revisit the B-band Tully-Fisher relation for a sample of 41 Hickson compact group galaxies, including a comparison with the GHASP sample of field galaxies, for which 2D velocity fields, derived from the H$\\alpha$ line, are also available (Epinat et al. 2008b). ", "conclusions": "We have presented new velocity fields, monochromatic H$\\alpha$ images and velocity dispersion maps for galaxies in nine compact groups. We have explored the B-band Tully-Fisher relation for a sample of 41 galaxies in compact groups and we compared it with the reanalyzed Tully-Fisher relation for the GHASP sample. We found that galaxies in dense environments lie on the Tully-Fisher relation defined by field galaxies, however, some of the low-mass galaxies have either a high luminosity for their mass or a high rotational velocity for their luminosities (e.g. HCG 49b). We also found that one third of the non-barred galaxies in compact groups have a misalignment between the kinematic and morphological position angle higher than 20 degrees." }, "1003/1003.0459_arXiv.txt": { "abstract": "The lensing data of the galaxy cluster Abell 1689 can be explained by an isothermal fermion model with a mass of 1-2 eV. The best candidate is the 1.5 eV neutrino; its mass will be searched down to 0.2 eV in KATRIN 2015. If its righthanded (sterile) modes were created too, there is 20\\% neutrino hot dark matter. Their condensation on clusters explains the reionization of the intercluster gas without Pop. III stars. Baryonic structure formation is achieved by gravitional hydrodynamics alone, without dark matter trigger. ", "introduction": "It is presently understood that the mass density of the universe is (nearly) equal to the critical density, with a fraction $\\Omega_B\\approx 4.5\\%$ in baryons, $\\Omega_D=20-25\\%$ in dark matter and the rest in dark energy. In the standard view, dark matter is cold, that is, a mass in the TeV regime, or at least warm, keV or more. Well over 40 searches have been performed, which failed to detect the dark matter particle, or came with claims that are not broadly accepted: the annual variation of the signal recorded in DAMA or the ``two hints or background events'' of CDMS-II. In Ref. \\citen{Nneutrino09} the present author takes a ``blind'' approach to describe the lensing data of the galaxy cluster Abell 1689. The assumptions are: the dark matter is a non-interaction (quantum) gas; the cluster is stationary; it may be approximated as spherically symmetrical; the mass distribution is isothermal for each of its components: Galaxies, X-ray gas and dark matter; all three components are subject to the common gravitational potential. While galaxies and gas are so dilute that a classical isothermal model applies, the dark matter may in principle be quantum degenerate, having a Bose-Einstein or Fermi-Dirac distribution. If the mass comes out large and thus the density low, either case will reduce to a Maxwell-Boltzmann. ", "conclusions": "" }, "1003/1003.5560_arXiv.txt": { "abstract": "Using a high resolution $N$-body simulation of a two-component dwarf galaxy orbiting in the potential of the Milky Way, we study two effects that lead to significant biases in mass estimates of dwarf spheroidal galaxies. Both are due to the strong tidal interaction of initially disky dwarfs with their host. The tidal stripping of dwarf stars leads to the formation of strong tidal tails that are typically aligned with the line of sight of an observer positioned close to the host center. The stars from the tails contaminate the kinematic samples leading to a velocity dispersion profile increasing with the projected radius and resulting in an overestimate of mass. The tidal stirring of the dwarf leads to the morphological transformation of the initial stellar disk into a bar and then a spheroid. The distribution of stars in the dwarf remains non-spherical for a long time leading to an overestimate of its mass if it is observed along the major axis and an underestimate if it seen in the perpendicular direction. ", "introduction": "\\begin{figure}[t] \\begin{center} \\includegraphics[width=16cm]{001elokas_fig1.ps} \\end{center} \\caption{An example of a kinematic sample of a thousand stars (left panel), the velocity dispersion profiles measured from it (middle panel) and the contours showing the inferred mass and anisotropy (right panel). In each panel the gray (black) color refers to the sample contaminated by (cleaned of) interloper stars from the tails.} \\label{fig1} \\end{figure} For the purpose of this study we used a high-resolution simulation of a two-component dwarf galaxy orbiting in the gravitational potential of the Milky Way (Klimentowski et al. 2007, 2009a). The dwarf progenitor of total mass $M = 4.3 \\times 10^9 M_{\\odot}$ consisted of a baryonic disk embedded in a dark matter halo. The initial mass of the disk was $1.5 \\times 10^8 M_{\\odot}$ and the mass of the NFW dark matter halo was $4.1 \\times 10^9 M_{\\odot}$. The dwarf galaxy was placed on an eccentric orbit around the host galaxy with apocenter to pericenter ratio of $r_{\\rm a}/r_{\\rm p}=120/25$ kpc and the disk initially inclined by 45$^\\circ$ to the orbital plane. The evolution was followed for 10 Gyr corresponding to five orbital times. The host galaxy was modelled by a static gravitational potential assumed to have the present-day properties of the Milky Way as described by mass model A1 of Klypin et al. (2002). During the evolution the dwarf galaxy is strongly affected by the tidal field of the host which results in a significant mass loss, the morphological transformation from a disk to a bar and then a spheroid and the transition from the streaming to the random motion of the stars. At all times, including the final stage, the core of the dwarf galaxy remains gravitationally bound but is surrounded by pronounced tidal tails and its shape departs from spherical. If dwarf spheroidal galaxies of the Local Group indeed formed as envisioned by this so-called tidal stirring scenario (Mayer et al. 2001, 2007), these effects must be taken into account when modeling their masses. Here we provide quantitative estimates of the biases imposed by the contamination of kinematic samples by unbound stars from the tails and by the departures from sphericity of the stellar component. ", "conclusions": "" }, "1003/1003.0942_arXiv.txt": { "abstract": "We study the evolution of linear density fluctuations of free-streaming massive neutrinos at redshift of $z<1000$, with an explicit justification on the use of a fluid approximation. We solve the collisionless Boltzmann equation in an Einstein de-Sitter (EdS) universe, truncating the Boltzmann hierarchy at $l_{\\rm max}=1$ and $2$, and compare the resulting density contrast of neutrinos, $\\delta_{\\rm\\nu}^{\\rm fluid}$, with that of the exact solutions of the Boltzmann equation that we derive in this paper. Roughly speaking, the fluid approximation is accurate if neutrinos were already non-relativistic when the neutrino density fluctuation of a given wavenumber entered the horizon. We find that the fluid approximation is accurate at few to 25\\% for massive neutrinos with $0.05a_{DE}$, and the scale around $k\\gtrsim k_{\\rm FS}(a_{DE})\\gg k_{\\rm FS}(a_{\\rm nr})$. We found that, as long as the term proportional to the gravitational potential, $\\psi(k,x)$, dominates the right hand side of Eq.(\\ref{eq:psi1}), the fluid approximation is valid. Therefore, unless the effect of the dark energy suppresses $\\psi(k,x)$ much faster than the growth of $\\epsilon^2(q,x)\\propto a^2$, $k_{max}$ at $a>a_{DE}$ will not change significantly. Since we have studied the evolution of the distribution function solving the collision-less Boltzmann equation, one can apply these results to other collision-less particles in general. Now, why is fluid approximation useful? Future and on-going dark energy missions aim at the accurate measurement of the galaxy/matter power spectrum with an accuracy better than 1\\%. One might think that the cosmological linear perturbation theory has already been well established, and the numerical codes such as CMBfast and CAMB can calculate the linear matter power spectrum with an accuracy better than 1\\%. However, the linear perturbation theory breaks down at small-scale and low redshift, where the density contrast becomes non-linear ($k\\gtrsim0.1~h~{\\rm Mpc^{-1}}$ at $z\\sim1$) \\cite{jeong/komatsu:2006,carlson/white/padmanabhan:2009}. Therefore, in order to exploit the cosmological information contained in a given survey, one needs to understand the non-linearities on the galaxy/matter power spectrum \\cite{yamamoto/bassett/nishioka:2005,rassat/etal:2008,shoji/jeong/komatsu:2009}. Among the non-linearities, the matter clustering has been well understood in the mildly non-linear regime (see \\cite{bernardeau/etal:2002}, for a review), but the theories have been limited to CDM dominated universe. The pressure gradient term in the Euler equation was completely ignored. In our previous work, we developed the 3rd-order perturbation theory with the pressure gradient terms explicitly included \\cite{shoji/komatsu:2009} (also see \\cite{saito/takada/taruya:2008,wong:2008,lesgourgues/etal:2009,mcdonald:2009}). With this extension to the higher order perturbation theory as well as within the limitation on the accuracy of $\\delta_{\\rm \\nu}$ calculated from the fluid approximation, we can now calculate the next-to-linear order matter power spectrum with massive neutrino free-streaming effect, properly included. Since the structure formation is mostly affected by the most massive species of neutrinos, and the current constraints on the total mass of neutrinos indicate that at least one of the neutrino species has a mass of order a tenth of ${\\rm eV}$, the use of fluid approximation is limited with an accuracy of few to 25\\% over $k\\lesssim0.4~h~{\\rm Mpc^{-1}}$ for $z<10$. As a result, for a small fraction of massive neutrino, $f_{\\rm\\nu}\\lesssim0.04$ for $\\sum_im_{\\rm\\nu,i}\\lesssim0.5~{\\rm eV}$, the fractional error on the matter density contrast, $\\delta_m=(1-f_{\\rm\\nu})\\delta_{\\rm c}+f_{\\rm\\nu}\\delta_{\\rm\\nu}$, calculated with the fluid approximation is accurate to sub-percent level. This material is based in part upon work supported by the Texas Advanced Research Program under Grant No. 003658-0005-2006, by NASA grants NNX08AM29G and NNX08AL43G, and by NSF grant AST-0807649. M.~S. thanks for warm hospitality of Astronomical Institute at Tohoku University where part of this work was done. \\appendix \\begin{widetext}" }, "1003/1003.1339_arXiv.txt": { "abstract": "The total infrared (TIR) luminosity from galaxies can be used to examine both star formation and dust physics. We provide here new relations to estimate the TIR luminosity from various {\\em Spitzer} bands, in particular from the 8~$\\mu$m and 24~$\\mu$m bands. To do so, we use 45\\arcsec\\ subregions within a subsample of nearby face-on spiral galaxies from the Spitzer Infrared Nearby Galaxies Survey (SINGS) that have known oxygen abundances as well as integrated galaxy data from the SINGS, the Local Volume Legacy Survey (LVL) and \\cite{engelbracht2008a} samples. Taking into account the oxygen abundances of the subregions, the star formation rate intensity, and the relative emission of the polycyclic aromatic hydrocarbons at 8~$\\mu$m, the warm dust at 24~$\\mu$m and the cold dust at 70~$\\mu$m and 160~$\\mu$m we derive new relations to estimate the TIR luminosity from just one or two of the {\\em Spitzer} bands. We also show that the metallicity and the star formation intensity must be taken into account when estimating the TIR luminosity from two wave bands, especially when data longward of 24~$\\mu$m are not available. ", "introduction": "With the end of the {\\em Spitzer} cold phase and the widespread availability of 8~$\\mu$m and 24~$\\mu$m bands observations in the archives, the availability of relations to determine the total infrared (TIR) emission from these wave bands by themselves is crucial to efficiently exploit the archives. In addition, the {\\em Herschel} Space Observatory observes dust emission at rest-frame 24~$\\mu$m and longer wavelengths for galaxies redshifted to $z\\simeq1.5$ in the PACS 60~$\\mu$m and longward bands with a spatial resolution as good as the {\\em Spitzer} 24~$\\mu$m band, for instance. Thus, estimating the total infrared flux from these bands becomes crucial for measuring total infrared fluxes using {\\em Herschel} data. Finally, the advent of new instrumentation in the coming years, such as the James Webb Space Telescope (JWST) or the Atacama Large Millimeter Array (ALMA) will also open a new window on the infrared emission of nearby and distant galaxies. For instance, the 18~$\\mu$m JWST/MIRI band and the 350~$\\mu$m ALMA band will probe the rest-frame 8~$\\mu$m and the 160~$\\mu$m emission of $z\\simeq1.2$ galaxies. Determining the total infrared emission using the 8-160~$\\mu$m {\\em Spitzer} bands with equation 4 from \\cite{dale2002a} or equation 22 from \\cite{draine2007a} yields a better estimate of the total infrared flux than using a single wave band as a proxy for the total infrared flux, but they necessitate 3 and 4 infrared bands respectively. Indeed, these relations necessitate using the much lower resolution 70~$\\mu$m and 160~$\\mu$m data, which have resolutions of 18\\arcsec\\ and 40\\arcsec, respectively. The poorer resolution of the 70~$\\mu$m and 160~$\\mu$m bands strongly constrains the scales on which total infrared fluxes can be measured, even for local galaxies. Attempts to derive a relation to estimate the total infrared emission from the 8~$\\mu$m and 24~$\\mu$m bands have been made by \\cite{calzetti2005a} using NGC~5194 (M51) data, \\cite{perez2006a} using NGC~3031 (M81) data, and \\cite{thilker2007a} using M33 data. For each galaxy, the scatter around the relation is about 40\\%. \\cite{rieke2009a} recently showed that using the 24~$\\mu$m band only provided good results. Using the 24~$\\mu$m and the 70~$\\mu$m bands \\cite{papovich2002a} have also provided an estimate of the TIR emission with a similar uncertainty. However, these relations may be applicable only to galaxies with similar metallicities and star formation rate intensities (Calzetti et al. 2010, submitted; Li et al. 2010, in preparation). Metallicity variations have been associated with variations in mid- and far-infrared colors, as has been observed in metal-poor galaxies \\citep[e.g.][]{engelbracht2005a, engelbracht2008a}. Indeed, the \\cite{calzetti2005a} relation, which was derived using data from a very metal-rich galaxy, underestimates the total infrared emission by a factor of a few in metal-poor dwarf galaxies \\citep{cannon2005a,cannon2006a,cannon2006b}. In this article we derive relations to estimate the total infrared emission using just one or two of the {\\em Spitzer} bands, with a strong focus on deriving the total infrared flux from just the higher resolution, shorter wavelength 8~$\\mu$m and 24~$\\mu$m bands. To do so, we use regions within a subset of nearby face-on spiral galaxies from the Spitzer Nearby Galaxies Survey \\citep[SINGS,][]{kennicutt2003a} that are resolved in the {\\em Spitzer} 160~$\\mu$m band as well as integrated galaxy luminosities for galaxies in the SINGS, the Local Volume Legacy survey (LVL) and the \\cite{engelbracht2008a}, hereafter E08, samples. In section \\ref{sec:sample}, we describe the samples of galaxies and the data processing. In section \\ref{sec:results}, we present the results, and we discuss them in section \\ref{sec:discussion}. Finally we summarize our results and conclude in section \\ref{sec:conclusion}. ", "conclusions": "\\label{sec:conclusion} Using data for spatially resolved subregions within 13 face-on spiral galaxies as well as integrated luminosities for the SINGS, LVL, and E08 galaxies, we have derived new relations to estimate the total infrared luminosity based on using only one or two {\\em Spitzer} bands, particularly the 8.0~$\\mu$m and 24~$\\mu$m bands. Relations incorporating 8.0~$\\mu$m data vary significantly with oxygen abundance and especially with $\\Sigma\\left(TIR\\right)$. However, TIR estimates that do not include 8~$\\mu$m data are less dependent on oxygen abundances. In particular, the relations between the TIR emission and the 70~$\\mu$m or 160~$\\mu$m bands are relatively independent of oxygen abundances compared to the 8~$\\mu$m and 24~$\\mu$m ones. \\appendix" }, "1003/1003.1878_arXiv.txt": { "abstract": "It is shown that the first law of thermodynamics and the holographic principle applied to a cosmic causal horizon demand the zero cosmological constant and non-zero dynamical dark energy in the form of the holographic dark energy. This dark energy has a parameter $d=1$ and an equation of state $w_0\\simeq -0.903$ comparable to current observational data, if entropy of the horizon saturates the Bekenstein-Hawking bound. ", "introduction": " ", "conclusions": "" }, "1003/1003.2815.txt": { "abstract": "\\noindent We recently presented a series of dark energy theorems that place constraints on the equation of state of dark energy ($\\wdark$), the time-variation of Newton's constant ($\\dot G$), and the violation of energy conditions in theories with extra dimensions. In this paper, we explore how current and future measurements of $\\wdark$ and $\\dot G$ can be used to place tight limits on large classes of these theories (including some of the most well-motivated examples) independent of the size of the extra dimensions. As an example, we show that models with conformally Ricci-flat metrics obeying the null energy condition (a common ansatz for Kaluza-Klein and string constructions) are highly constrained by current data and may be ruled out entirely by future dark energy and pulsar observations. ", "introduction": "% Beginning with the work of Kaluza and Klein \\cite{KK1,KK2,KK3} and continuing today with string theory and M-theory, extra dimensions have been a common feature of unified theories. The basic notion is that the observed 3+1-dimensional universe is actually described by a general relativistic theory in a space-time with one or more extra compactified dimensions. If the compactification scale is much greater than 1~TeV (or $<<10^{-16}$~cm), laboratory experiments, even at the Large Hadron Collider, are unable to uncover direct evidence %detect evidence of extra dimensions. In this paper, though, we show how measurements of the equation of state of dark energy ($\\wdark$) and the time variation of Newton's constant $\\dot G/G$ can be used to test or rule out the existence of extra dimensions for large classes of models {\\it independent of the compactification scale.} This surprising power to discriminate among extra-dimensional models derives from a set of ``dark energy theorems,\" first described in \\cite{Wesley:2008de,Wesley:2008fg, Steinhardt:2008nk}. The theorems are based on the observation that the expansion of the usual three large dimensions tends to cause extra dimensions to vary with time, which, in turn, causes a change of $\\wdark$ and $G$ in the corresponding 4d effective theory. These changes can be avoided in a decelerating universe by introducing conventional interactions strong enough to keep the sizes of the extra dimensions fixed. However, the dark energy theorems show that, once the universe starts to accelerate, conventional interactions satisfying the classical (strong, weak and null) energy conditions no longer suffice no matter the size of the extra dimensions. Many well-motivated extra-dimensional models satisfy one or more of the energy conditions. For these large classes, the dark energy theorems, combined with the observed acceleration rate, can be used to compute the predicted time-variation of $\\wdark$ and $G$ for given model parameters. As illustrated below, current measurements are already strong enough to rule out a substantial range of model parameters. The more exciting prospect is anticipated improvements in the measurement of $\\wdark$, as described by the Dark Energy Task Force (DETF), and of $\\dot G/G$, as constrained by pulsar timing, that can test or rule out whole classes of extra-dimensional models. The dark energy theorems were derived for the general case of $k$ extra spatial dimensions, but the predictions depend on $k$. For the purposes of illustration, we focus in this paper on the well-motivated class of 9+1-dimensional theories ($k=6$) with a conformally flat Ricci (CRF) metric and satisfying the null energy condition (NEC) -- theories commonly used in string- and M-theoretic phenomenological and cosmological models. In Ref~\\cite{Steinhardt:2008nk}, we showed that this class of models is inconsistent with the standard $\\Lambda$CDM model which has $\\wdark=-1$. In fact, the dark energy theorems show that, the closer $\\wdark$ is to $-1$, the more rapidly the extra-dimensional volume and, hence, $G$ must vary. Furthermore, if the theory contains no mechanism for violating the null energy condition, $\\wdark$ cannot remain close to $-1$ for an extended period. This means that improved limits on $\\dot G$, combined with ever-tightening bounds on the time-variation of $\\wdark$ from future experiments, can progressively constrain or rule out this entire class of extra-dimensional models. A corollary of this analysis is that a dark energy mission, even if it fails to find any time-variation of $\\wdark$ and is consistent with $\\wdark=-1$, can still be highly informative because it would eliminate well-motivated extra-dimensional models. A second corollary is that a coordinated effort is needed. The ambitious improvements in the measurements of $\\wdark$ alone, as projected by the DETF, or of $\\dot G/G$ alone, as estimated from planned pulsar timing surveys, are not sufficient. Each constrains some range of parameter space but leaves some substantial untested range. The two approaches are complementary, though: by pursuing both to some degree, entire classes of extra-dimensional models can be tested and ruled out. This paper is organized as follows. In Section \\ref{s:basiceqs}, we introduce the conformally Ricci-flat (CRF) class of extra-dimensional models used to exemplify our approach and review the constraints imposed by the dark energy theorems on the {\\it total} equation of state $\\wtot$ and Newton's constant $G$. {\\it Here and throughout this paper, the symbol $\\wtot$ refers to all contributions to the energy density of the universe (matter, radiation, etc.), not just dark energy}. The current value is $\\wtot =-0.74$ based on observations \\cite{Komatsu:2010fb}. We use the symbol $\\wdark$ to refer to the dark energy component alone. In Section \\ref{s:method}, we show how to translate these constraints into predictions for dark energy and pulsar timing experiments. The key results are in Section \\ref{s:results}, where we compare the predictions with current measurements and near-future experiments. In particular, we show how current measurements of dark energy and pulsar timing and anticipated improvements can be used to test and perhaps rule out the entire class of models. In Section \\ref{s:conclusions}, we conclude with a discussion of the generalization to other classes of extra-dimensional models. More details on the constraints on the time-variation of $G$ and the dark energy equation of state in extra-dimensional models are given in the Appendices. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ", "conclusions": "Discussion} The results of the previous section demonstrate that it is possible to test and possibly rule out an entire class of extra-dimensional models in a way that does not depend on the size of the extra dimensions. The approach relies on the fact that the expansion of the universe is accelerating today and that acceleration causes the volume of the compactified dimensions to change in cases where the theory satisfies a classical energy condition. Instead of probing the compactified dimensions directly, the approach is to test for the effect of time-variation of the compactified dimensions on $G$ and $\\wdark$. A key advantage of this approach is that the time variation required by dark energy theorems does not depend on the size of the extra dimensions. For the purposes of illustration, we have focused on CRF models that obey the NEC because they are common to many string- and M-theoretic constructions. We have demonstrated that current observations allow a significant range of these models, but that improvements in the measurement of $\\dot G/G$ and $\\wdark$ anticipated over the next few years can rule them out entirely. To accomplish the task by improved dark energy constraints alone would require the DETF's most ambitious Stage IV plan combined with current measurements of $\\dot G/G$. {\\it The most cost-effective approach,, though, is to combine a modest improvement in $\\dot G/G$ and $\\wdark$ limits, which can impose constraints tighter than those obtained from a very high-precision dark energy measurement alone.} Fig.~\\ref{f:requiredPulsar} summarizes how different combinations of improved measurements of $\\dot{G}/{G}$ from pulsar timing measurements and dark energy experiments described by the DETF can rule out the entire range of $k=6$-dimensional CRF models obeying the NEC at the $3 \\sigma$ level or higher. \\begin{figure} \\begin{center} \\includegraphics[width=6in]{Figure06.pdf} \\caption{\\label{f:requiredPulsar} The improvements in $\\dot G/G$ constraints required to rule out NEC-satisfying models with CRF-type metrics and $k=6$ extra dimensions, as a function of the available dark energy constraints. For each DETF stage, the bars show the improvement in pulsar constraints that would lead to exclusion of this family of models at the 3$\\sigma$, 4$\\sigma$, and 5$\\sigma$ levels. By contrast, this family cannot be ruled out (even at the 1$\\sigma$ level) from dark energy information alone, even with constraints from extremely ambitious measurements. } \\end{center} \\end{figure} A similar approach can be applied to test other metrics and/or other energy conditions. For example, we have carried out a similar analysis for the family of models in which the extra-dimensional theory satisfies the {\\it strong energy condition} (SEC) and is described by curved metrics of the form: \\begin{equation}\\label{e:TDmetric} \\text{d} s^2 = e^{2\\Omega(t,y)} g^{\\rm FRW}_{\\mu\\nu}(t,x) \\, \\text{d} x^\\mu \\text{d} x^\\nu + h_{\\alpha\\beta}(t,y) \\, \\text{d} y^\\alpha \\text{d} y^\\beta \\end{equation} where $g^{\\rm FRW}_{\\mu\\nu}$ is a flat Friedmann-Robertson-Walker metric; and the extra-dimensional metric $h_{\\alpha\\beta}$ and warp factor $\\Omega$ can be time-dependent. Unlike the CRF family, in (\\ref{e:TDmetric}) we allow the extra-dimensional metric $h_{\\alpha\\beta}$ to have arbitrary Ricci curvature. This family of models satisfies a set of dark energy theorems that is different from the theorems for NEC/CRF family. We find that, as in the case of the NEC/CRF metrics, the SEC/Curved family of models is not ruled out by current constraints on $\\wdark$ and $\\dot{G}/G$, but DETF Stage II limits combined with current constraints on $\\dot{G}/G$ is sufficient to rule out the entire family of models. We expect to be able to extend our approach to yet more combinations of metrics and energy conditions. Even for the NEC/CRF family of models, it may be possible to derive additional dark energy theorems. The current theorems specify conditions that are necessary to have cosmic acceleration and still satisfy the energy conditions, but they are not sufficient. The current theorems were selected because they are the simplest to prove, as discussed in Ref.~\\cite{Steinhardt:2008nk}. However, there may be stronger theorems that combine with observational constraints of $\\wdark$ and $\\dot G/G$ to produce much more stringent constraints on extra-dimensional models. \\appendix" }, "1003/1003.5340_arXiv.txt": { "abstract": "We present the observational results of an $L'$ and $M$ band Adaptive Optics (AO) imaging survey of 54 nearby, sunlike stars for extrasolar planets, carried out using the Clio camera on the MMT. We have concentrated more strongly than all other planet imaging surveys to date on very nearby F, G, and K stars, prioritizing stellar proximity higher than youth. Ours is also the first survey to include extensive observations in the $M$ band, which supplement the primary $L'$ observations. Models predict much better planet/star flux ratios at the $L'$ and $M$ bands than at more commonly used shorter wavelengths (i.e. the $H$ band). We have carried out extensive blind simulations with fake planets inserted into the raw data to verify our sensitivity, and to establish a definitive relationship between source significance in $\\sigma$ and survey completeness. We find 97\\% confident-detection completeness for 10$\\sigma$ sources, but only 46\\% for 7$\\sigma$ sources -- raising concerns about the standard procedure of assuming high completeness at 5$\\sigma$, and demonstrating that blind sensitivity tests to establish the significance-completeness relation are an important analysis step for all planet-imaging surveys. We discovered a previously unknown $\\sim0.15$M$_{\\sun}$ stellar companion to the F9 star GJ 3876, at a projected separation of about 80 AU. Twelve additional candidate faint companions are detected around other stars. Of these, eleven are confirmed to be background stars, and one is a previously known brown dwarf. We obtained sensitivity to planetary-mass objects around almost all of our target stars, with sensitivity to objects below 3 \\mjup~in the best cases. Constraints on planet populations based on this null result are presented in our Modeling Results paper, \\citet{modeling}. ", "introduction": "Nearly 400 extrasolar planets have now been discovered using the radial velocity (RV) method. RV surveys currently have good statistical completeness only for planets with periods of less than ten years \\citep{cumming,butlercat}, due to the limited temporal baseline of the observations, and the need to observe for a complete orbital period to confirm the properties of a planet with confidence. The masses of discovered planets range from just a few Earth masses \\citep{hotNep} up to around 20 Jupiter masses (\\mjup). We note that a 20 \\mjup~object would be considered by many to be a brown dwarf rather than a planet, but that there is no broad consensus on how to define the upper mass limit for planets. For a good overview of RV planets to date, see \\citet{butlercat} or \\url{http://exoplanet.eu/catalog-RV.php}. The large number of RV planets has enabled several good statistical analyses of planet populations \\citep{carnp, butlercat,cumming}. However, these apply only to the short-period planets accessible to RV surveys. We cannot obtain a good understanding of planets in general without information on long period extrasolar planets; nor can we see how our own solar system fits into the big picture of planet formation in the galaxy without a good census of planets in Jupiter- and Saturn-like long-period orbits around other stars. Several methods (transit detection, RV variations, astrometry, and direct imaging) have yielded repeatable detections of extrasolar planets so far. While RV and astrometric surveys may eventually deliver important information about long-period extrasolar planets, direct imaging is the only method that allows us to characterize them on a timescale of months rather than years or decades. Direct imaging of extrasolar planets is technologically possible at present only in the infrared, based on the planets' own thermal luminosity, not on reflected starlight. The enabling technology is adaptive optics (AO), which allows 6-10m ground-based telescopes to obtain diffraction limited IR images several times sharper than those from HST, despite Earth's turbulent atmosphere. Theoretical models of giant planets indicate that such telescopes should be capable of detecting self-luminous giant planets in large orbits around young, nearby stars. The stars should be young because the glow of giant planets comes from gravitational potential energy converted to heat in their formation and subsequent contraction: lacking any internal fusion, they cool and become fainter as they age. Several groups have published the results of AO imaging surveys for extrasolar planets around F, G, K, or M stars in the last five years (see for example \\citet{masciadri,kasper,biller1,GDPS,chauvin}). Of these, most have used wavelengths in the 1.5-2.2 $\\mu$m range, corresponding to the astronomical $H$ and $K_S$ filters \\citep{masciadri,biller1,GDPS,chauvin}. They have targeted mainly very young stars. Because young stars are rare, the median distance to stars in each of these surveys has been more than 20 pc. In contrast to those above, our survey concentrates on very nearby F, G, and K stars, with proximity prioritized more than youth in the sample selection. The median distance to our survey targets is only 11.2 pc. Ours is also the first survey to include extensive observations in the $M$ band, and only the second to search solar-type stars in the $L'$ band (the first was \\citet{kasper}). The distinctive focus on older, very nearby stars for a survey using longer wavelengths is natural: longer wavelengths are optimal for detecting objects with very red IR colors -- that is, low temperature planets. These are most likely to be found in older systems, since planets cool and redden with age \\citep{bar,bur}. However old, low-temperature planets also have low luminosities, rendering them undetectable around all but the nearest stars. In Section \\ref{sec:samp} we describe the criteria used in choosing our sample, and present the characteristics of our stars. In Section \\ref{sec:Observe}, we briefly describe our instrument, our observing strategy, and our image processing pipeline. In Section \\ref{sec:sensanal} we detail our sensitivity estimation methods, and show how we characterized them using blind tests in which simulated planets were inserted into our raw data -- a practice that should be standard for planet imaging surveys. In Section \\ref{sec:faintreal} we give astrometric and photometric data for all the faint companions detected in our survey, as well as precise astrometry of the bright known binary stars in our sample. We present our conclusions in Section \\ref{sec:concl}. Constraints on planet populations based on our survey null result are presented in \\citet{modeling}. ", "conclusions": "\\label{sec:concl} We have surveyed unusually nearby, mature star systems for extrasolar planets in the $L'$ and $M$ bands using the Clio camera with the MMT AO system. We have developed a sophisticated image processing pipeline for data from this camera, including some interesting innovations. We have carefully and rigorously analyzed our sensitivity. Accurately determining the sensitivity of AO planet-search images is a more complex task than, perhaps, has been widely appreciated. Our data support the conclusion of \\citet{mspeckle} that $5 \\sigma$ limits can substantially overestimate the meaningful sensitivity of an image. Blind tests involving fake planets inserted in raw data are the best way to confirm the validity of any sensitivity estimator, and should be included in all future planet-search publications. By extensive use of such tests, we established a definitive significance vs. completeness relation for planets in our data. This relation is important for use in Monte Carlo simulations to constrain planet distributions. We have discovered a physically orbiting $\\sim 0.15$M$_{\\sun}$ binary companion at a projected separation of 80 AU from the star GJ 3876. We have detected twelve additional candidate faint companions, one of which is the binary brown dwarf companion of GJ 564 discovered prior to our observations by \\citet{potter}. The remaining eleven are confirmed to be background stars. We note that shorter wavelength surveys, such as that of \\citet{GDPS} in the $H$ band regime, have typically found a much larger number of background stars, necessitating extensive follow-up observations. A long wavelength survey such as ours can obtain good sensitivity to planets, with their very red IR colors, while remaining blind to all but the brightest stars. This reduces the amount of telescope time spent following up planet candidates that turn out to be background stars. We did not detect any planets, but have set interesting limits on the masses of planets or other substellar objects that may exist in the star sytems we surveyed. In \\citet{modeling}, we use extensive Monte Carlo simulations to show how our null result constrains the mass and semimajor axis distributions of extrasolar planets orbiting sun-like stars." }, "1003/1003.0753_arXiv.txt": { "abstract": "We present a study of X-ray AGN overdensities in 16 Abell clusters, within the redshift range $0.0730.995$ level, as indicated by an appropriate $t$-student test. We conclude that the triggering of luminous X-ray AGN in rich clusters is strongly suppressed. % Furthermore, searching for optical {\\em Sloan Digital Sky Survey} (SDSS) counterparts of all the X-ray sources, associated with our clusters, we found that about half appear to be background QSOs, while others are background and foreground AGN or stars. The true overdensity of X-ray point sources, associated to the clusters, is therefore even smaller than what our statistical approach revealed. ", "introduction": "There is a growing body of studies investigating the effect of the environment on the nuclear activity of galaxies and on the possible triggering mechanisms of the AGN phenomenon (e.g., Dultzin-Hacyan et al. 1999; Koulouridis et al. 2006; 2009; Sorrentino et al. 2006; Gonz\\'alez et al. 2008; Choi et al. 2009; Silverman et al. 2009; Lee et al. 2009; von der Linden et al. 2009; Padilla, Lambas \\& Gonz\\'alez 2009 and references therein). One particular research direction is the study of X-ray AGN, as a function of environment% , since undoubtedly, one of the best AGN identification methods is through X-ray observations. Clusters of galaxies offer an ideal target for this type of study and indeed most X-ray based studies report overdensities of X-ray sources in clusters, with respect to the field % (e.g., Cappi et al. 2001; Molnar et al. 2002; D'Elia et al. 2004; Branchesi, et al. 2007; Galametz et al. 2009; Gilmour et al. 2009). There are attempts to substantiate such results with spectroscopic data and indeed various studies have verified the existence of a large population of X-ray AGN in clusters of galaxies and its probable evolution with redshift (eg. Martini et al. 2002; Johnson et al. 2003; Martini et al. 2007, 2009; van Breukelen et al. 2009). Some of the previous studies have also shown that only a small fraction of the X-ray sources, associated with clusters of galaxies, possess optical emission line ratios characteristic of AGN (see also Davis et al. 2003, Finoguenov et al. 2004). Similarly, starting from optical spectroscopy, % Arnold et al. (2009) found that 14 out of 144 galaxy members of 11 clusters and groups are AGN, but with only one being detected in X-rays. A similar trend of disconnection between X-ray and optically selected AGN, in eight poor groups of galaxies, was also found by Shen et al. (2007). We believe that there is a major question still not adequately answered, which is: {\\em Is there an enhancement of AGN activity and/or of its X-ray manifestation in clusters of galaxies with respect to what expected from the obvious optical galaxy overdensity?} We attempt to address this question by searching not only for the existence or not of a cluster X-ray AGN excess, with respect to the field, since such could possibly be expected on the basis of the known excess of optical galaxies in clusters, but rather by investigating whether the X-ray AGN overdensity shows a relative enhancement or suppression with respect to the corresponding optical galaxy overdensity. ", "conclusions": "We find that there is a relatively significant X-ray source overdensity in about half of the clusters in our sample, the rest showing the excepted background value of $\\delta_x$. This can be realized by inspecting Table 1, where we list the X-ray and optical overdensity values and their uncertainties, as well as from Figure 1 where we plot the cluster X-ray point-source overdensities versus the corresponding optical SDSS galaxy overdensities, within the indicated $r$-band magnitude range. Note that the 6 clusters with effective X-ray flux-limit corresponding to a minimum luminosity of $L_x \\magcir 2 \\times 10^{42}$ erg s$^{-1}$ (see specific values in Table 1), at the redshift of the clusters, are indicated by a triangular point type.% Although the number of clusters in our sample is quite low to provide stringent population statistics, a secure and important conclusion of our analysis is that the cluster X-ray point-source overdensities are always lower than the corresponding optical SDSS overdensities (with only two clusters having $\\delta_x\\sim\\delta_o$). In Table 2 we present such population statistics but separately for the two cluster subsamples having different limiting $L_x$ (as discussed previously). Note also that in order to take into account the variable uncertainty of the overdensity values we present in Table 2 the Poisson uncertainty-weighted mean optical and X-ray overdensities. A $t$-student test comparing the two means, assuming unknown and unequal variances, shows that they are different at a high significance level (see ${\\cal P}$ in Table 2). Inspecting Table 2 it becomes evident that {\\em the luminous X-ray AGN overdensity is suppressed by a factor of 3 - 4, on the mean, with what would have been expected from a constant fraction, independent of environment, of X-ray AGN to bright optical galaxies.} This result is in the same direction with the optical SDSS analysis of Lee et al. (2009). % In order to investigate in more detail the nature of the X-ray overdensities in our cluster sample, we have also cross-identified all detected X-ray sources with the SDSS database, finding in total only six out of the 88 detected X-ray point-sources (with $L_x\\ge 10^{42}$ erg s$^{-1}$ at the redshift of the cluster) being clearly associated with the clusters; among which one spectroscopically confirmed AGN (Sy1), 2 galaxies with no apparent emission lines % and 3 more galaxies, based on their photometric redshifts (one of which, in A1689, is indeed confirmed by the spectroscopic analysis of Martini et al. 2007). In Table 1 we also present the optical characterization of the X-ray point-sources, the total number of which for each cluster field, $N_x$, is listed in column 5. The 6$^{th}$ column indicates the number of probable background QSO, $N_{\\rm QSO}$, (mostly determined as such from their point-like images and their $u-g$ versus $g-r$ colors, while $\\sim 10\\%$ are also spectroscopically verified), the 7$^{th}$ column indicates the number of X-ray sources clearly associated with cluster galaxies, $N_{\\rm cgal}$, and the 8$^{th}$ column indicates the number of apparently irrelevant associations, $N_{\\rm other}$, like stars, foreground/background galaxies, no-counterparts, smudges, etc. About $\\sim 50\\%$ of all our AGN candidates % appear to be related to background QSO, which in many clusters represent the expected background provided by the $\\log N-\\log S$ of Kim et al. (2007). Therefore, the real overdensity of X-ray AGN associated with our cluster sample appears to be significantly smaller than what is listed in column 3 of Table 1, a fact which further strengthens our main result that the rich cluster environment (within 1 $h^{-1}_{72}$ Mpc) strongly suppresses the luminous X-ray AGN activity. In order to provide also a visual example of the source categorization, we present in Figure 2 the XMM field of A2065, the lowest $z$ cluster of our sample, together with the SDSS composite $gri$-band images of all the X-ray point-source counterparts.% Finally, we stress that the main conclusion of our present analysis is that in the intermediate intracluster distances (ie., between $3r_c$ and 1$h^{-1}_{72}$ Mpc), the relatively dense and hot ICM environment not only does not enhance AGN activity, but it rather strongly suppresses it (at least its X-ray luminous manifestation). {\\small" }, "1003/1003.0938_arXiv.txt": { "abstract": "The Australia Telescope Compact Array (ATCA) has been used to make the first extensive search for the class~I methanol masers at 9.9~GHz. In total, 48 regions of high-mass star formation were observed. In addition to masers in W33-Met (G12.80$-$0.19) and G343.12$-$0.06 (IRAS 16547$-$4247) which have already been reported in the literature, two new 9.9-GHz masers have been found towards G331.13$-$0.24 and G19.61$-$0.23. We have determined absolute positions (accurate to roughly a second of arc) for all the detected masers and suggest that some class~I masers may be associated with shocks driven into molecular clouds by expanding H{\\sc ii} regions. Our observations also imply that the evolutionary stage of a high-mass star forming region when the class~I masers are present can outlast the stage when the class~II masers at 6.7-GHz are detectable, and overlaps significantly with the stage when OH masers are active. ", "introduction": "Methanol masers are well established probes of high-mass star formation. They are divided into two categories first defined by \\citet{bat87}: class~II masers (of which the 6.7~GHz is the best known and usually strongest) are associated with millimetre and infrared sources \\citep[e.g.,][]{hil05} and reside in the close environment of high-mass young stellar objects (YSOs). Class~I masers (of which the 44 and 95~GHz are commonly observed) are often found apart from the strong continuum and infrared sources and can be separated by up to a parsec from the YSO responsible for excitation \\citep[e.g.,][]{kur04, vor06, cyg09}. Theoretical calculations are able to explain this empirical classification and strongly suggest that the pumping process of class~I masers is dominated by collisions with the molecular hydrogen \\citep[see, e.g.,][]{lee73,sob83,vor99} in contrast to class~II masers, which are pumped by radiative excitation \\citep[see, e.g.,][]{wil85, sob94, sut01}. The two pumping mechanisms were shown to be competitive \\citep[see, e.g.,][]{men91,cra92,vor99,vor05b}. For example, strong radiation from a nearby infrared source quenches class~I masers and increases the strength of class~II masers \\citep{vor05b}. Although it has been demonstrated by theoretical calculations that a weak class~II maser at 6.7~GHz can coexist with a bright class~I emission under special conditions \\citep{vor05b}, bright masers of different classes residing in the same volume of gas are widely accepted as mutually exclusive. However, on larger scales, they are often observed to coexist in the same star forming region within less than a parsec of each other. Class~I methanol masers are relatively poorly studied. The common consensus is that the majority of class~I masers trace interface regions between outflows and molecular gas, although direct observational evidence of this has been obtained for a limited number of sources only \\citep[e.g.,][]{pla90,kur04,vor06}. An alternative scenario, which involves cloud-cloud collisions, may be realised in some sources \\citep{sob92,meh96,sal02}. The common ingredient of these two scenarios is the presence of shocks. There is observational evidence that the methanol abundance is significantly increased in the shock processed regions \\citep[e.g.][]{gib98,sut04}. The gas in such regions is heated and compressed increasing the frequency of collisions with molecular hydrogen and, therefore, providing a more efficient pumping \\citep[e.g.][]{sut04}. \\citet{che09} demonstrated statistically the association of class~I masers with the shocks traced by extended features showing a prominent excess of the 4.5-$\\mu$m emission in the images obtained with the Spitzer Space Telescope's Infrared Array Camera (IRAC) also known as extended green objects \\citep[EGOs;][]{cyg08}. To date, there is observational evidence that more than 20 methanol transitions can produce class~I masers, more than half of which belong to the J$_2-$J$_1$~E series in the 25$-$30~GHz frequency range \\citep{vor99,mul04}. The observational properties of different class~I masers are not the same. Some masers, e.g. at 44 and 95~GHz, are known to be quite ubiquitous with about 100 sources known to date \\citep[see, e.g.,][]{has90,sly94,val00,kur04,cyg09}. On the other hand, there are other transitions with very few known maser sources. An example of such a transition showing rare maser activity is the $9_{-1}-8_{-2}$~E methanol transition at 9.9~GHz. The rarity of these masers is likely due to the strong dependence of the maser brightness on the physical conditions, especially in the requirement of higher typical temperatures and densities for these masers to form \\citep{sob05}. \\citet{sly93} conducted the only search for the 9.9-GHz masers reported in the literature so far. However, they observed just 11 targets which were largely well known regions of high-mass star-formation where methanol masers of either class had previously been reported and discovered a single 9.9-GHz maser (W33-Met also known as G12.80$-$0.19). In this paper we report the results of the search for the 9.9-GHz masers towards 48 independent positions. The majority of targets, 46 sources in total, were known class~I masers at 44 and/or 95~GHz drawn from \\citet{sly94}, \\citet{val00}, \\citet{kur04}, and \\citet{ell05}. The remaining two sources were included in order to cover all known star forming regions located south of declination of $-$20$^\\circ$, which have a periodically variable class~II methanol maser at 6.7-GHz according to \\citet{goe04}. Little is known about the physics of periodic variability of such masers. Therefore, a detection of the 9.9-GHz maser (which is very sensitive to the physical conditions and has a different pumping to the 6.7-GHz maser) in the same source might shed light on this enigmatic phenomenon. ", "conclusions": "\\begin{enumerate} \\item Two new class~I methanol masers at 9.9-GHz were found in addition to two other sources already reported in the literature. The absolute positions with arcsecond accuracy are summarised in Table~\\ref{positive} for all four 9.9-GHz masers known to date. Based on the trial observations, we also suspect that another 9.9-GHz maser may exist in G305.21$+$0.21 at a flux density below the detection threshold of our survey. \\item We suggest that some class~I methanol masers may be associated with shocks driven into molecular cloud by an expanding H{\\sc ii} region. This is an alternative scenario to the association with outflows which has been proved unambiguously in a number of other cases \\citep[e.g.,][]{kur04,vor06}. This new scenario applies also to the class~I methanol masers in other transitions, e.g. to the widespread masers at 36 and 44~GHz. It does not appear to be an exclusive mechanism to generate the 9.9-GHz masers, although could be responsible for 3 out of 4 known 9.9-GHz masers. \\item The evolutionary stage with the class~I maser activity is likely to outlast the stage when the 6.7-GHz methanol masers are present and overlap significantly in time with the stage when the OH masers are active. We expect that the class~I masers in one of the widespread maser transitions (e.g. at 44~GHz) will be detected towards a significant number of OH maser sources which do not have a class~II methanol maser at 6.7~GHz. Such sources were rarely observed in the class~I maser surveys currently available in the literature. \\item A 9.9-GHz methanol maser was found in G331.13$-$0.24 which also hosts one of the few known periodically variable class~II methanol masers at 6.7~GHz. This rare combination makes this source an attractive target for a long-term monitoring program, which has a high potential to shed light on the periodic variability of masers. \\end{enumerate}" }, "1003/1003.2268_arXiv.txt": { "abstract": "We report new radial velocities of the TrES-4 transiting planetary system, including observations of a full transit, with the High Dispersion Spectrograph of the Subaru 8.2m telescope. Modeling of the Rossiter-McLaughlin effect indicates that TrES-4b has closely aligned orbital and stellar spin axes, with $\\lambda = 6.3^{\\circ} \\pm 4.7^{\\circ}$. The close spin-orbit alignment angle of TrES-4b seems to argue against a migration history involving planet-planet scattering or Kozai cycles, although there are two nearby faint stars that could be binary companion candidates. Comparison of our out-of-transit data from 4 different runs suggest that the star exhibits radial velocity variability of $\\sim$20~m~s$^{-1}$ in excess of a single Keplerian orbit. Although the cause of the excess radial velocity variability is unknown, we discuss various possibilities including systematic measurement errors, starspots or other intrinsic motions, and additional companions besides the transiting planet. ", "introduction": "Transiting planets provide us with valuable opportunities to characterize exoplanetary systems. One such opportunity is to measure the Rossiter-McLaughlin effect (hereafter the RM effect: \\cite{1924ApJ....60...15R}, \\cite{1924ApJ....60...22M}), an apparent radial velocity anomaly during a planetary transit, which is caused by a partial eclipse of the rotating surface of the host star. By measuring and modeling this effect, one can measure the sky-projected angle between the stellar spin axis and the planetary orbital axis. Many theoretical investigations of the RM effect have been presented (e.g., \\cite{2005ApJ...622.1118O, 2006ApJ...650..408G, 2007ApJ...655..550G, 2010ApJ...709..458H}), and observations of the RM effect have been reported for about 20 transiting planetary systems (for the most recent compilation, see \\cite{2010arXiv1001.0416J} and references therein). One of the main theoretical motivations to observe the RM effect is that the observed degree of spin-orbit alignment is thought to be connected with the migration history of the transiting planet. The most frequently discussed planetary migration mechanisms are (1) gravitational interaction between a protoplanetary disk and a growing planet (disk-planet interaction models, e.g., \\cite{1985prpl.conf..981L, 1996Natur.380..606L, 2004ApJ...616..567I}), (2) gravitational planet-planet scattering and subsequent tidal evolution (planet-planet scattering models, e.g., \\cite{1996Sci...274..954R, 2002Icar..156..570M, 2008ApJ...678..498N, 2008ApJ...686..580C}), or (3) the Kozai mechanism caused by a distant massive companion and subsequent tidal evolution (Kozai migration models, e.g., \\cite{2003ApJ...589..605W, 2005ApJ...627.1001T, 2007ApJ...669.1298F, 2007ApJ...670..820W}). These scenarios are not necessarily mutually exclusive, but to the extent that they are, disk-planet interaction models would predict small orbital eccentricities and good spin-orbit alignments, while planet-planet scattering models and Kozai migration models predict a broader range of eccentricities and spin-orbit alignment angles. Until about a year ago, all of the measurements indicated close alignments, but recently 6 transiting planets have been reported with significant misalignments: XO-3b \\citep{2008A&A...488..763H, 2009ApJ...700..302W}, HD~80606b \\citep{2009A&A...498L...5M, 2009A&A...502..695P, 2009ApJ...703.2091W}, WASP-14b \\citep{2009PASP..121.1104J, 2009MNRAS.392.1532J}, HAT-P-7b \\citep{2009PASJ...61L..35N, 2009ApJ...703L..99W}, CoRoT-1b \\citep{2009MNRAS.tmpL.360P}, and WASP-17b \\citep{2010ApJ...709..159A}. With this increase in the number of measurements, and the diversity of results, we are approaching the time when we may test the validity and applicability of the different planetary migration models. The main target of this paper is TrES-4b, which is a transiting exoplanet discovered by \\citet{2007ApJ...667L.195M} (hereafter M07) in the course of the TrES survey, supplemented by Keck radial velocity measurements. The planet orbits an F8 \\citep{2009A&A...498..567D} host star every 3.55 days and is one of the most ``inflated'' hot Jupiters with a radius of about 1.8 $R_{\\rm Jup}$, which places this planet to be one of the least density exoplanets ever discovered. Refined spectroscopic and photometric characteristics of the host star TrES-4 were presented by \\citet{2009ApJ...691.1145S} (hereafter S09). The amplitude of the RM effect for TrES-4b was expected to be large, because of the large projected equatorial rotational velocity of the TrES-4 star ($V \\sin I_s = 8.5$~km~s$^{-1}$; S09). In addition, \\citet{2009A&A...498..567D} have recently reported a possible companion star around the TrES-4 system. Although it has not yet been confirmed that the companion is a true physical binary as opposed to a chance alignment, a companion star would raise the possibility of migration via Kozai cycles, lending additional motivation to the study of the RM effect in this system. We note that we adopt ``TrES-4'' as the host star name and ``TrES-4b'' as the planet name in this paper, although the planet was originally named ``TrES-4'' by the discoverers (see M07). The reason of our choice is because recent papers on this system (e.g., Daemgen et al. 2009, which we referred in this paper) often describe the host star as ``TrES-4''. Thus we consider that it would be confusing for readers if we describe the planet as ``TrES-4'' in our paper, and we hope to avoid such confusions. In this paper, we present new measurements of the radial velocity of TrES-4 made with the Subaru 8.2m telescope. Although TrES-4 is relatively faint ($V = 11.6$), the large aperture of the Subaru telescope has enabled us to measure radial velocities of TrES-4 with high precision. Our radial velocity dataset consists of 23 samples covering a full transit, and 8 samples obtained outside of transits on 3 different nights. In addition to reporting the spin-orbit alignment angle of TrES-4b, we also report the observation of radial velocity variation in excess of the previously observed single Keplerian orbit, and we confirm through direct imaging with the HDS slit viewer that there are candidate companion stars. The rest of this paper is organized as follows. Section~2 summarizes our Subaru observations, and section~3 describes analysis procedures of the RM effect of TrES-4b. Section~4 presents our main result on the spin-orbit alignment angle of TrES-4b, and section~5 discusses possible causes of the additional radial velocity variation in this system. Finally, section~6 summarizes the findings of this paper. ", "conclusions": "\\subsection{Possible Causes of the Additional Radial Velocity Variation} Since we could not find an appropriate model for the all observed radial velocities at this point (see figure~2), we here discuss possibilities of systematic effects as well as real sources of excess radial velocity variation. \\subsection*{-- Instrumental Instability of the Subaru HDS} Since the Subaru radial velocities were gathered on a few days in clusters spanning about 1 year, it is of utmost importance to know the instrumental stability of the HDS. For observations within a single night, \\citet{2007PASJ...59..763N} studied the radial velocity standard star HD~185144 and found that the Subaru HDS is stable within a few~m~s$^{-1}$. In addition, \\citet{2009ApJ...703L..99W} reported that the Subaru HDS is stable within a few~m~s$^{-1}$ over two weeks based on HAT-P-7 observations. Likewise, \\citet{2008ApJ...686..649J} did not find systematic offsets for the HAT-P-1 system over approximately 1 month, using the same setup (and even some of the same nights) as the TrES-4 observations presented here. However, specific stability over 1 year has not yet been confirmed through monitoring of radial velocity standard stars, although studies for such long-term stability of the Subaru HDS are in progress. Thus we note that the instrumental instability of the Subaru HDS is one of the prime possibilities of a cause of the additional radial velocity variation at this point. \\subsection*{-- Starspots} One possibility involves starspots on the photosphere of TrES-4. Since the rotational velocity of TrES-4 is relatively fast ($V \\sin I_s = 8.5$~km~s$^{-1}$: S09), stellar spots of similar size to a planet would cause an apparent radial velocity shift like the RM effect, on a timescale of the stellar rotation period ($P_{\\rm rot} \\approx 11$~days, assuming $\\sin I_s \\approx 1$). For example, a dark spot of approximately the same size as the planet would lead to a maximum shift of 85~m~s$^{-1}$, while smaller spots with less contrast would contribute smaller velocities. If this is the case, all the RV data are affected. However, one would not expect a hot F8 star to have large spots. It is because M07 did not report such possibility of stellar spots from the TrES transit survey, S09 reported no active $\\textrm{Ca}$ HK line emission ($\\log R'_{\\tiny{\\textrm{HK}}} = -5.11 \\pm 0.15$), and \\citet{2009ApJ...691..866K} concluded that spot activity is unlikely (but not impossible) based on Spitzer observations. Thus the spot explanation is doubtful, although further long-term photometric monitoring would be useful to constrain this hypothesis still further. \\subsection*{-- Other Sources of Stellar Jitter} Another possible explanation of the systematic radial velocity variation is an intrinsic stellar jitter (see e.g., \\cite{2005PASP..117..657W}), i.e., motions of the stellar photosphere due to pulsations or other flows. Although an empirical relation reported by \\citet{2005PASP..117..657W} predicted a typical stellar jitter of $4.4$~m~s$^{-1}$ for stars like TrES-4, it is conceivable that TrES-4 has an unusually unstable photosphere. To explain the observed radial velocities, we would need to invoke a stellar jitter of about $20$~m~s$^{-1}$ for TrES-4 based on the rms residuals of the Subaru and Keck datasets. The jitter would need to have a time scale longer than a few days, in order to explain why the M07 observations (conducted on consecutive three nights) do not exhibit such a large scatter. In this light the hypothesis that the stellar jitter explains all the excess variability seems too contrived. \\subsection*{-- Contamination of Companion Star's Lights or Sky Backgrounds} As described by \\citet{1996PASP..108..500B}, a slight change in the instrumental profile would result in a systematic shift of the apparent radial velocity. Thus, any contaminating light from the nearby candidate companion star of TrES-4 or sky backgrounds (e.g., moonlight or twilight) may have affected the radial velocities. We limited the aperture width of echelle orders in order to avoid contamination of lights from the north companion star. Moonlight contamination is unlikely since our TrES-4 observations were conducted in clear and moonless time, and sky background levels were still low although 2 exposures (HJD of 2454317.74338 and 2454535.14401) were conducted during twilight. Note that M07 did not report the existence of the companion star, and therefore it is possible that the companion star might have been on the slit during the M07 observations. Such an effect might have caused some systematic shifts in the M07 data. Although we were not able to estimate the systematic effect in the M07 data, it could be a small effect since the companion is very faint. \\subsection*{-- Eccentricity of TrES-4b} As M07 reported only 4 radial velocity samples, they did not include the eccentricity $e$ and the argument of periastron $\\omega$ in their radial velocity model. Instead they assumed the orbit to be circular as we have done, and more recently \\citet{2009ApJ...691..866K} found $e \\cos \\omega < 0.0058$ ($3\\sigma$) based on Spitzer secondary eclipse observations. With the Subaru data we now have a sufficient number of radial velocity samples to allow the eccentricity and argument of pericenter to be free parameters. However, allowing $e$ and $\\omega$ to vary does not improve the model fit, and the eccentricity of the best-fit model is nearly zero. This is consistent with the constraint by \\citet{2009ApJ...691..866K}. Thus a large eccentricity of TrES-4b could not be the explanation for the observed excess RV variability. \\subsection*{-- Additional Planets} If the preceding explanations for the observed radial velocity variation could be ruled out, we would consider a possibility of presence of additional planets. This is the most interesting case, however, at this point our Subaru observations were too sparse to find and confirm another periodicity in the radial velocities. We can only conclude that the radial velocity semiamplitude of about 20~m~s$^{-1}$ over a year is possible for hypothetical planets in the TrES-4 system. Obviously, further continuous radial velocity monitoring would be necessary to check on this possibility." }, "1003/1003.3744_arXiv.txt": { "abstract": "{The FU Orionis candidate V733 Cep was discovered by Roger Persson in 2004. The star is located in the dark cloud L1216 close to the Cepheus OB3 association. Because only a small number of FU Orionis stars have been detected to date, photometric and spectral studies of V733 Cep are of great interest.} {The studies of the photometrical variability of PMS stars are very important to the understanding of stellar evolution. The main purpose of our study is to construct a long-time light curve of V733 Cep. On the basis of $BVRI$ monitoring we also study the photometric behavior of the star.} {We gather data from CCD photometry and archival photographic plates. The photometric $BVRI$ data (Johnson-Cousins system) that we present were collected from June 2008 to October 2009. To facilitate transformation from instrumental measurements to the standard system, fifteen comparison stars in the field of V733 Cep were calibrated in $BVRI$ bands. To construct a historical light curve of V733 Cep, a search for archival photographic observations in the Wide-Field Plate Database was performed. As a result, 192 photographic plates containing the field of V733 Cep were found. Some plates were analyzed at our request to estimate the magnitude of V733 Cep.} {Our photometric study confirms the affiliation of V733 Cep to the group of FU Orionis objects. An outburst in the optical and a slow rise in brightness during the period 1971-1993 are well documented. During the period 1993-2004, V733 Cep exhibited its maximum brightness and the amplitude of the observed outburst exceeded 4$\\fm$5 (R). The $BVRI$ photometric data imply that from February 2007 to October 2009, a slow decrease in brightness of V733 Cep was observed. The observed color evolution of $V-I$ index also suggest that V733 Cep is currently fading. The long-time light curve of V733 Cep is similar to the light curves of other FU Orionis objects.} {} ", "introduction": "During the Pre-Main Sequence (PMS) stage of evolution, young stellar objects exhibit different types of photometric variability (Herbst et al. 1994). One of the most dramatic of these events, with very high amplitude variations, is the FU Orionis (FUor) outburst (Ambartsumian 1971). The flare-up of FU Orionis itself was documented by Wachmann (1939) and for several decades it was the only known object of that type. Herbig (1977) defined FUors as a class of young variables after the discovery of two new FUor objects, V1057 Cyg and V1515 Cyg. The main characteristics of FUors are an increase in optical brightness of about 4-5 mag, a F-G supergiant spectrum with broad blue-shifted Balmer lines, strong infrared excess, connection with reflection nebulae, and location in star-forming regions (Reipurth 1990, Bell et al. 1995, Clarke et al. 2005). Typical spectroscopic properties of FUors include a gradual change in the spectrum from earlier to later spectral type from the blue to the infrared, a strong Li I 6707 line, P Cygni profiles of H$\\alpha$ and Na I 5890/5896 lines, and the presence of CO bands in the near infrared spectra (Herbig 1977, Bastian \\& Mundt 1985). The prototypes of FUors seem to be low-mass PMS objects (T Tauri stars) with massive circumstellar disks. According to a commonly accepted view, the FUor outburst is produced by a sizable increase in accretion from a circumstellar disk on the stellar surface (Hartmann \\& Kenyon 1985). The cause of this increase in accretion from $\\sim$10$^{-7}$$M_{\\sun}$$/$yr to $\\sim$10$^{-4}$$M_{\\sun}$$/$yr appears to be thermal or gravitational instability in the circumstellar disk. This accretion disk model can account for most of the main properties of FUors. An alternative explanation of the FUor phenomenon is the rapid-rotator hypothesis (Herbig at al. 2003). Among all objects associated with the group of FUors, only three (FU Ori, V1057 Cyg, and V1515 Cyg) have detailed photometric observations taken during the outburst and during the fading period (Clarke et al. 2005). For a few objects, V1735 Cyg (Elias 1978, Peneva et al. 2009), V346 Nor (Graham \\& Frogel 1985), and V733 Cep (Reipurth et al. 2007), the presence of an optical outburst is also documented and they are labeled classical FUors. Two new suspected FUor objects with observed outbursts at optical wavelengths are reported: V582 Aurigae (Samus 2009, Munari et al. 2009) and the object CSS091110, coincident with the infrared source IRAS 06068-0641 (Wils et al. 2009). About a dozen objects have spectroscopic properties similar to the classical FUors, but there is no evidence of an outburst at optical wavelengths. These objects are termed FUor-like (Reipurth et al. 2002, Greene et al. 2008) and only partial photometric observations have been published for them. The PMS star V733 Cep (Persson's star) is located in the dark cloud L1216 close to Cepheus OB3 association. The variability of V733 Cep was discovered by Swedish amateur astronomer Roger Persson in 2004 (Persson 2004). Persson compared the plate scans from the first and the second Palomar Sky Survey. He noted the star of red magnitude 15$\\fm$7 on the POSS-II image (UT 1991 September 3) and its absence on the corresponding POSS-I image (UT 1953 October 31). The star is also visible on a Palomar Quick-V plate from 1984. On an image taken with the 2.2-m telescope in Mauna Kea, Hawaii on UT 2004 October 9, Reipurth et al. (2007) measured a red magnitude of V733 Cep at about 17$\\fm$3. Comparing this value with that reported by Persson (2004), Reipurth et al. (2007) concluded that the star had faded by 1$\\fm$6 (R) over a time period of about 13 yr (from 1991 to 2004). The authors propose that an outburst occurred in the period 1953-1984 and identify great spectral similarities to FU Ori itself. In our first paper (Semkov \\& Peneva 2008; hereafter SP08), data from photometric monitoring of V733 Cep in the period February 2007 - February 2008 were presented, and it was noted that no significant changes in the star brightness were registered during the period of observations. Because of the short time period of observations, a trend toward a decrease in brightness was not detected. ", "conclusions": "The analysis of the available photometric data has allowed us to precisely classify V733 Cep as a FUor variable. The long-time light curve of the star is similar to the light curves of others FUor objects. The observed photometric variations in the period of fading are also typical of some FUor stars. The time of rise in brightness and the star magnitude in the maximum light remain unclear. Therefore, the collection of new photometric data (from photographic plate archives and from ongoing photometric monitoring) will be of great importance for a precise determination of the outburst parameters. We conclude that V733 Cep is presently the FUor object with the longest time of increase in brightness and probably the first found to have an approximately symmetrical light curve." }, "1003/1003.6027_arXiv.txt": { "abstract": "HD~189733 is a K2 dwarf, orbited by a giant planet at 8.8 stellar radii. In order to study magnetospheric interactions between the star and the planet, we explore the large-scale magnetic field and activity of the host star. We collected spectra using the ESPaDOnS and the NARVAL spectropolarimeters, installed at the 3.6-m Canada-France-Hawaii telescope and the 2-m Telescope Bernard Lyot at Pic du Midi, during two monitoring campaigns (June~2007 and July~2008). HD~189733 has a mainly toroidal surface magnetic field, having a strength that reaches up to 40 G. The star is differentially rotating, with latitudinal angular velocity shear of $\\dom = 0.146 \\pm 0.049$~\\rpd, corresponding to equatorial and polar periods of $11.94 \\pm 0.16 $~d and $16.53 \\pm 2.43$~d respectively. The study of the stellar activity shows that it is modulated mainly by the stellar rotation (rather than by the orbital period or the beat period between the stellar rotation and the orbital periods). We report no clear evidence of magnetospheric interactions between the star and the planet. We also extrapolated the field in the stellar corona and calculated the planetary radio emission expected for HD~189733b given the reconstructed field topology. The radio flux we predict in the framework of this model is time variable and potentially detectable with LOFAR. ", "introduction": "\\label{sec:intro} Magnetic fields are present at different scales in the universe, from planets, to stars, galaxies and galaxy clusters. Thanks to new high-resolution spectropolarimeters, we are able to study the large scale magnetic field of stars, trying to understand its origin, but also its implication in different stellar phenomena (stellar wind, magnetic braking, stellar cycles, ...). In Hot Jupiter (HJ) systems (giant planets orbiting close to their parent stars, i.e. with semi-major axis of the planet lower than 0.1 AU), the study of the stellar magnetic field is important to understand Star-Planet interactions (SPI), and their effects on the evolution and properties of the system. SPI can be of two types: magnetospheric (e.g. caused by reconnections between the stellar and the planetary magnetic fields), or tidal (resulting from the proximity and masses of the two bodies). \\cite{cuntz00} suggested that such interactions may enhance the stellar activity. Studying the activity of HJ hosting stars, \\cite{shk03,shk05,shk08} concluded that not all the observed systems show hints of interactions and that for a single system with interactions, those interactions may be not observable during some observing epochs, yet observable at other times. Different theoretical scenarios of magnetospheric interactions were proposed. \\cite{preusse06} described SPI by adopting the Alfv\\'en wind model, \\cite{lanza08} considered a non-potential magnetic field configuration for the closed corona of the star. To explain the 'on-off' nature of SPI, \\cite{cranmer07} studied their signatures over many orbital cycles considering a cyclic stellar magnetic field. These signatures do not repeat exactly from orbit to orbit, nor from epoch to epoch. SPI is an intermittent phenomenon depending strongly on the configuration of the stellar field. The study of the stellar magnetic field is therefore important to understand SPI. We have started an observing program aimed at detecting and modelling the magnetic field of HJ hosting stars. Eleven systems are observed, having different stellar and planetary parameters. In this paper, we present the results for HD~189733. This system is interesting to study SPI, with a short orbital period different from the stellar rotation period. The properties of HD~189733 are listed in section \\ref{sec:star}. In sections \\ref{sec:obs} and \\ref{sec:mod}, we will present our data, data analysis and magnetic modelling of the star. Results of its magnetic topology and differential rotation will also be presented. Stellar activity will be studied in section \\ref{sec:Activity}. From magnetic maps of the stars obtained in section \\ref{sec:mod}, we will extrapolate the magnetic field in the stellar corona (section \\ref{sec:extrapolation}), deduce the expected planetary radio emission (section \\ref{sec:radio}), and end with our conclusions (section \\ref{sec:conclusions}). ", "conclusions": "\\label{sec:conclusions} In this paper, we present a detailed spectropolarimetric study of the star HD~189733, host of a transiting giant planet. The star was observed at two epochs (June~2007 and July~2008). Using Zeeman Doppler Imaging, we reconstructed the magnetic maps of the star. With a strength up to 40~G, the magnetic field is dominated by the toroidal component at both epochs. This component contributes 57\\%~and 77\\%~to the total energy respectively and is mainly axisymmetric. In contrast, the poloidal component is mainly non-axisymmetric. Its contribution to the total energy drops from 2006 to 2007 and 2008. We will continue monitoring this system to study the magnetic evolution on time-scales longer than 2 years and look for a potential magnetic cycle. HD~189733 rotates differentially and has a latitudinal angular rotation shear of $\\dom = 0.146 \\pm 0.049$~\\rpd~; the star has an equatorial period of $11.94\\pm0.16$ d and a polar period of $16.53\\pm2.43$~d. These values of the equatorial and polar periods bracket all published photometric periods within the error bars. The star is an active star, variable on small time-scales. We analyzed the activity residuals in the \\caii~H\\&K\\ and H$\\alpha$. These residuals are periodic and modulated with the rotational period of the star. Active regions apparently concentrate around the equator, given the modulation period of 12-13~d. A rotational modulation was also found by \\citealt{moutou07}, \\citealt{boisse09}, and \\citealt{shk08}. We looked for lower amplitude periodic fluctuations, the periods we found are different than the orbital and beat period (2.5-2.7~d) except in one occurrence (H$\\alpha$, July 2008) where one of the detected roughly matches the beat period. To enlarge our knowledge of this system, we studied the magnetic field in the stellar atmosphere using the extrapolation technique applied to the reconstructed surface magnetic maps. We find that HD~189733 has a complex magnetic topology for both epochs. Depending on where the SS (the surface beyond which the field is purely radial) is, we find that the magnetic field at the distance of the planet ($8.8~R_{\\star}$) is variable throughout the orbit, of the order of $4-23~\\rm mG$~in average. We also estimated the radio flux expected from SPI assuming the magnetic scenario model of \\cite{Zarka01} and find it to be of the order of $10-220~\\rm{mJy}$ on average. We also predict it to be variable with time on a time scale equal to the beat period (contrary to previous published predictions \\citealt{Griessmeier07AA}). The radio flux we predict at 0-6~MHz is potentially detectable with LOFAR (see fig. 1-3 in \\citealt{Griessmeier07AA}) in the coming years. Published radio observations only report upper limits on the planetary flux at higher frequencies \\citep[307-347~MHz,][244~MHz and 614~MHz]{smith09,lecavalier09}~-~providing no constraint on the model discussed in the present paper. The variability of the planetary radio flux with the subplanetary stellar phase will make the distinction between the planetary and stellar radio flux in the observational data more challenging; however, this effect can be used to distinguish between the magnetic energy model and other models of interactions. Our result confirms that a single observation of a star-planet system is not sufficient. Rather, it is important to have multiple observations densely sampling stellar rotation, planetary orbit and beat periods. The study of SPI is an ongoing effort. Monitoring stars at different epochs and through multi-wavelength campaigns will help us identify the nature of SPI and the origin of their apparent on-off behavior. Studying stellar magnetic cycles and comparing results for HJ hosting stars with different stellar and planetary parameters will enlarge our understanding of SPI, as well as stellar magnetism and activity in general." }, "1003/1003.5118_arXiv.txt": { "abstract": "In this work, we present a detailed study of the rotational properties of magnetized and self-gravitating dense molecular cloud cores formed in a set of two very high resolution three-dimensional molecular cloud simulations with decaying turbulence. The simulations have been performed using the adaptative mesh refinement code RAMSES with an effective resolution of 4096$^3$ grid cells. One simulation represents a mildly magnetically-supercritical cloud and the other a strongly magnetically-supercritical cloud. We identify dense cores at a number of selected epochs in the simulations at two density thresholds which roughly mimick the excitation densities of the NH$_{3}$ ($J-K$)=(1,1) transition and the N$_{2}$H$^{+}$ (1-0) emission line. A noticeable global difference between the two simulations is the core formation efficiency (CFE) of the high density cores. In the strongly supercritical simulations the CFE is $~33$ percent per unit free-fall time of the cloud ($t_{ff,cl}$), whereas in the mildly supercritical simulations this value goes down to $\\sim 6$ percent per unit $t_{ff,cl}$. A comparison of the intrinsic specific angular momentum ($j_{3D}$) distributions of the cores with the specific angular momentum derived using synthetic two-dimensional velocity maps of the cores ($j_{2D}$), shows that the synthetic observations tend to overestimate the true value of the specific angular momentum by a factor of $\\sim 8-10$. We find that the distribution of the ratio $j_{3D}/j_{2D}$ of the cores peaks at around $\\sim 0.1$. The origin of this discrepancy lies in the fact that contrary to the intrinsic determination of $j$ which sums up the individual gas parcels contributions to the angular momentum, the determination of the specific angular momentum using the standard observational procedure which is based on a measurement on the global velocity gradient under the hypothesis of uniform rotation smoothes out the complex fluctuations present in the three-dimensional velocity field. Our results may well provide a natural explanation for the discrepancy by a factor $\\sim 10$ observed between the intrinsic three-dimensional distributions of the specific angular momentum and the corresponding distributions derived in real observations. We suggest that previous and future measurements of the specific angular momentum of dense cores which are based on the measurement of the observed global velocity gradients may need to be reduced by a factor of $\\sim 10$ in order to derive a more accurate estimate of the true specific angular momentum in the cores. We also show that the exponent of the size-specific angular momentum relation are smaller ($\\sim 1.4$) in the synthetic observations than their values derived in the three-dimensional space ($\\sim 1.8$). ", "introduction": "An important issue for current ideas of star formation is whether most stars of a given mass are born single, in a binary system or in multiple systems (e.g., Bodenheimer 1995; Bodenheimer et al. 2000; Larson 2010). Among other several physical processes that affect the fragmentation of cores prone to star formation such as their geometry (e.g, Boss 2009), their metallicity (e.g., Hocuk \\& Spaans 2010), and degree of magnetization (e.g., Boss 1999; Price \\& Bate 2007; Hennebelle \\& Teyssier 2008) the rotational properties of the cores will also strongly affect their ability to fragment (e.g., Nelson 1998; Sigalotti \\& Klapp 2001; Matsumoto \\& Hanawa 2003; Hennebelle et al. 2004; Machida et al. 2009; Walch et al. 2009). This may have important consequences on the fraction of binary and multiple systems. The determination of the angular momentum of the cores and the fraction of their energy stored in rotational motions is complicated by the fact that the velocity fields in the cores usually exhibit complex supersonic motions when mapped with tracers such as CO, CS, and C$^{18}$O with a transition to subsonic, coherent motions when tracers with a higher excitation density are used such as NH$_{3}$ and N$_{2}$H$^{+}$ lines (e.g. Barranco \\& Goodman 1998; Goodman et al. 1998; Caselli et al. 2002a). Usually, the assumption is often made in the observations that cores have a uniform rotation and follow a rigid-body rotation law. Their angular velocity $\\Omega$ is deduced from measuring global velocity gradients in velocity maps built using velocity measurements along the line of sight. Such measurements for a large number of molecular clouds, clumps or cores were performed by Myers \\& Benson (1983), Goldsmith \\& Arquilla (1986), Kane \\& Clemens (1997), Pound \\& Goodman (1997), and Rosolowsky (2007) using CO lines, and by Goodman et al. (1993), Barranco \\& Godman (1998) using observations in the ($J-K$)=(1-1) transition of NH$_{3}$, and Caselli et al. (2002a), Pirogov et al. (2003), Olmi et al. (2005) and Chen et al. (2007) using N$_{2}$H$^{+}$ (1-0) line observations. Several other authors have also obtained rotation measurements for individual cores or small numbers of cores (e.g., Harris et al. 1983; Menten et al. 1984; Armstrong et al. 1985; Wadiak et al. 1985; Zheng et al. 1985; Ho \\& Haschick 1986; Jackson et al. 1988; Ho et al. 1994; Ohashi et al. 1997; Belloche et al. 2002; Caselli et al. 2002b; Tafalla et al. 2004; Shinnaga et al. 2004; Redman et al. 2004, Schnee et al. 2007; Chen et al. 2008, Chen et al. 2009; Csengeri et al. 2010). It is important to investigate whether the distributions of the rotational properties of cores show significant variations as a function of the adopted density tracer/threshold and as a function of the environment (e.g., for example for different magnetization levels of the parent cloud), and most importantly whether the distributions that are determined from observations are a faithful reproduction of the distributions in the intrinsic three-dimensional space, would the entire dynamical and structural properties of the cores be accessible. Furthermore, knowing the statistical distributions of the rotational properties of dense cores in molecular clouds is crucial in order to asses the statistical relevance of rotational parameters assigned to individual rotating core collapse simulations. In previous works, Burkert \\& Bodenheimer (2000) superimposed a random velocity field of power spectrum $P(k) \\propto k^{n}$ (with $n$ between $-3$ and $-4$) onto the density field of turbulent molecular cloud cores. They found that the projected velocity maps display velocity gradients that can be interpreted as rotation and measured the average specific angular momentum of the cores to be $7 \\times10^{20} (R_{c}/0.1)^{1.5}$ cm$^{2}$~s$^{-1}$, where $R_{c}$ is the size of the core, and the average value of the rotational parameter (i.e., ratio of rotational energy to gravitational energy) of $\\beta_{rot} \\sim 0.03$. Gammie et al. (2003) examined the distribution of specific angular momentum for cores formed in a set of magnetized, self-gravitating molecular clouds with decaying turbulence at the resolution of $256^{3}$. They showed the distribution of specific angular momentum at an intermediate epoch in one of their simulation (the one with beta plasma value of $\\beta_{p}=0.1$ after 0.09 sound crossing time). The specific angular momentum they found follows a distribution with values ranging between $10^{21}$ cm$^{2}$ s$^{-1}$ to $7 \\times 10^{23}$ cm$^{2}$ s$^{-1}$ with the peak of the distribution at $j=4 \\times 10^{22}$ cm$^{2}$ s$^{-1}$. Jappsen \\& Klessen (2004) found that there might be a dependence of the average specific angular momentum of the cores as a function of the Mach number of the cloud. Li et al. (2004) and Offner et al. (2008) also evaluated the specific angular momentum $j$ and rotational parameter $\\beta_{rot}$ for cores formed in their molecular cloud simulations, both with resolutions of $512^{3}$ grid cells, and found values $\\it {of~the~order}$ of the observed ones. We revisit the issue of rotation in molecular clouds cores using very high resolution simulations of magnetized, turbulent, and self-gravitating isothermal molecular clouds (MCs) performed with the Adaptative Mesh Refinment code RAMSES. Our primary aim is to quantify any systematic difference between the intrinsic angular momentum of the cores and the values of the specific angular momentum derived from synthetic observations built using an approach that mimicks the procedure commonly used in the observations. In \\S~\\ref{simul}, we describe the molecular cloud simulations and the clump finding algorithm used to identify dense cores in the simulations at different density thresholds. In \\S~\\ref{general} we describe some of the basic properties of the simulations. The rotational properties of dense cores are presented and discussed in \\S~\\ref{comp} and in \\S~\\ref{conc} we summarize our results and conclude. ", "conclusions": "\\label{conc} In this paper, we have analyzed the rotational properties of dense molecular cloud cores formed in two magnetized, self-gravitating molecular cloud simulations with a decaying turbulence. The two simulations differ by the strength of the magnetic field in the clouds with one cloud being mildly magnetically supercritical and the other being strongly magnetically supercritical. Our results show that the formation efficiency of dense cores is strongly reduced with increasing importance of the magnetic field in the cloud (going down from 33 percent per free-fall time in the strongly supercritical cloud to 6 percent for the mildly supercritical cloud). We also observe that the median value of the specific angular momentum of the high density cores in the mildly supercritical simulation is smaller than the values derived for cores in the strongly supercritical simulation. This result is consistent with the fact that magnetic braking which leads to angular momentum loss is playing a more important role in the cloud where the magnetic field is stronger. We have focused our attention on the discrepancies that may arise between estimates of the specific angular momentum of the cores derived from the global velocity gradient method commonly used in the observations and its the true value measured in the intrinsic three-dimensional space. In order to derive the specific angular momentum of the cores following the observational procedure, we generate synthetic velocity maps of the cores along three different projections. The global velocity gradient of the cores is measured from the velocity map using the VFIT routine employed initially by Goodman et al. (1993). The specific angular momentum is then calculated using the global velocity gradient value under the assumption of uniform rotation of the cores. We find, in the two simulations, that the distributions of the ratio of the specific angular momentum determined in the intrinsic 3D space to the one derived from projected velocity maps peaks at values around $\\sim 0.1$. This may well explain the difference by a factor $\\sim 10$ that is observed between the distribution of specific angular momentum derived from the intrinsic data in our simulations and the corresponding real observations using the NH$_{3}$ and N$_{2}$H$^{+}$ molecules of roughly similar excitation density than the density thresholds used to identify the cores in the simulations. We suggest that the origin of this discrepency (between 2D and 3D) lies in the fact that contrary to the intrinsic determination of $j$ which sums up the individual gas parcels contributions to the angular momentum, the observational determination of $j$ is based on a measurement on the global velocity gradient under the hypothesis of uniform rotation which smoothes out the complex fluctuations present in the three-dimensional velocity field. We therefore suggest that previous measurements of the specific angular momentum of the cores overestimate its true value and that a correction factor of $\\sim 10$ should be applied to these measurements as well as to new determinations of the specific angular momentum when using the global gradient method adopted so far in the observations. As already stressed by other groups (e.g., Padoan et al. 1998,2000; Pichardo et al. 2000; Ostriker et al. 2001; Ballesteros-Paredes \\& Mac Low 2002; Pineda et al. 2009; Federrath et al. 2010; Shetty et al. 2010) our work further highlights the importance of generating synthetic observations from three-dimensional numerical simulations that can be compared to real observations." }, "1003/1003.0810_arXiv.txt": { "abstract": "{ We analyse CMB data in a manner which is as independent as possible of the model of late-time cosmology. We encode the effects of late-time cosmology into a single parameter which determines the distance to the last scattering surface. We exclude low multipoles $\\ell<40$ from the analysis. We consider the WMAP5 and ACBAR data. We obtain the cosmological parameters $100\\ob =2.13\\pm 0.05$, $\\oc=0.124\\pm 0.007$, $n_s=0.93\\pm 0.02$ and $\\theta_A=0.593^{\\circ}\\pm 0.001^{\\circ}$ (68\\% C.L.). The last number is the angular scale subtended by the sound horizon at decoupling. There is a systematic shift in the parameters as more low $\\ell$ data are omitted, towards smaller values of $\\ob$ and $n_s$ and larger values of $\\oc$. The scale $\\theta_A$ remains stable and very well determined. } \\preprint{CERN-PH-TH/2010-053} \\begin{document} \\setcounter{tocdepth}{2} \\setcounter{secnumdepth}{3} ", "introduction": "\\label{sec:intro} The cosmic microwave background (CMB) is one of the most important cosmological probes. The pattern of acoustic oscillations of the baryon-photon plasma is imprinted on the CMB at the time of decoupling, and then rescaled (and on large scales modified) as the CMB photons propagate from the last scattering surface to the observer. The CMB is thus sensitive to cosmological parameters in two ways, via the physics at decoupling and via the evolution of the universe after that. While the physics at decoupling --essentially atomic physics and general relativity of a linearly perturbed Friedmann-Lema\\^\\i tre (FL) universe-- is well understood, the evolution at late times deviates from the predictions of linearly perturbed FL models with radiation and matter. The difference may be due to an exotic matter component with negative pressure such as vacuum energy, deviation of gravity from general relativity~\\cite{RR,Linder,Koyama,CapoFranca}, or a breakdown of the homogeneous and isotropic approximation \\cite{Rasanen:2006b,Buchert:2007,Enqvist:2007,Rasanen:2008a,Zibin:2008b,Celerier:2009}. It is not known which of these possibilities is correct, and there are large differences between the various models. It is therefore worthwhile to analyse the CMB in a manner which is as independent of the details of late-time cosmology as possible. On the one hand, this clarifies the minimal constraints that all models of late-time cosmology, whatever their details, have to satisfy in order to agree with CMB observations. On the other hand, our analysis provides limits on the physical parameters at decoupling that are independent of the details of what happens at later times. This is particularly important for cosmological parameters such as the density of baryons, density of dark matter and the spectral index, which are used to constrain particle physics models of baryogenesis, supersymmetry and inflation, which are independent of late-time cosmology. Such a separation of constraints is possible because the physics after decoupling affects the CMB in a rather limited manner (except at low multipoles), by simply changing the angular scale and modifying the overall amplitude of the CMB pattern. We encode the change in the angular scale in a single parameter related to the angular diameter distance to the last scattering surface and treat the amplitude as a nuisance parameter. We aim to be transparent about how the different cosmological parameters enter the calculation and the assumptions that go into the analysis. In section 2 we discuss how the physics at early and late times affects the CMB and explain our assumptions. In section 3 we present the results of the analysis of the WMAP 5-year data \\cite{Komatsu:2008, Dunkley:2008, Hinshaw:2008} and the ACBAR data \\cite{acbar} and give the constraints on cosmological parameters. In section 4 we summarise our results. Some details are collected in two appendices. ", "conclusions": "We have analysed the CMB data in a way which is independent of the details of late-time cosmology, i.e. the cosmology at redshifts $z\\lsim 60$. The results we have obtained are therefore valid for most models of late-time cosmology, whether they include dark energy, modified gravity, a local void or backreaction. We have presented model-independent limits on $\\ob$, $\\oc$, $n_s$ and the angular diameter distance to the last scattering surface $D_A(z_*)$, or its ratio with the sound horizon at last scattering, $\\theta_A=r_s(z_*)/D_A(z_*)$. The present CMB data give an extraordinarily precise measurement of $\\theta_A$, which every realistic model of the late universe must agree with. We can summarize the final result by \\bea\\label{e:fin} 100\\ob &=& 2.13 \\pm 0.05 \\ , \\qquad \\oc = 0.124 \\pm 0.007 \\el n_s &=& 0.93 \\pm 0.02 \\ , \\qquad \\theta_A = 0.593^\\circ\\pm 0.001^\\circ \\ . \\eea \\noindent Note that the values of $\\oc$ and $\\ob$ actually determine the matter and baryon density at last scattering via the relation $\\rho_x(z_*) = (1+z_*)^3(H_0/h)^2\\om_x$. The values of the densities today may be different e.g. if dark matter decays at late times \\cite{DMdecay}. In summary, every model which satisfies equations \\re{e:fin} will automatically be in agreement with the present CMB data for $\\ell\\geq40$. Only lower $\\ell$ CMB data, large scale structure, lensing and other observations can distinguish between models which have the above values for $\\ob, \\oc, n_s$ and $\\theta_A$. We have also found that there is a systematic shift in the cosmological parameters as more low $\\ell$ data are cut. As more data from low multipoles is removed, $\\ob$ and $n_s$ decrease, while $\\oc$ becomes larger. These changes keep the power spectrum at small scales fixed, but tend to increase the amplitude on large scales. These changes are not reflected in the statistical error bars: the small angle data prefer different parameter values than the full set of CMB data. This trend is visible to at least $\\lmin=100$. Whether this behaviour has any connection with the various directional features at low multipoles~\\cite{anom, Hansen:2008, Francis:2009, Bennett:2010}, is not clear." }, "1003/1003.2587.txt": { "abstract": "High-dispersion spectra of 89 potential members of the old, super-metal-rich open cluster, NGC 6253, have been obtained with the HYDRA multi-object spectrograph. Based upon radial-velocity measurements alone, 47 stars at the turnoff of the cluster color-magnitude diagram (CMD) and 18 giants are identified as potential members. Five turnoff stars exhibit evidence of binarity while proper-motion data eliminates two of the dwarfs as members. The mean cluster radial velocity from probable single-star members is -29.4 $\\pm$ 1.3 km/sec (sd). A discussion of the current estimates for the cluster reddening, derived independently of potential issues with the $BV$ cluster photometry, lead to an adopted reddening of E$(B-V)$ = 0.22 $\\pm$ 0.04. From equivalent width analyses of 38 probable single-star members near the CMD turnoff, the weighted average abundances are found to be [Fe/H] = $+0.43 \\pm 0.01$, [Ni/H] = $+0.53 \\pm 0.02$ and [Si/H] = $+0.43 ^{0.03}_{0.04}$, where the errors refer to the standard errors of the weighted mean. Weak evidence is found for a possible decline in metallicity with increasing luminosity among stars at the turnoff. We discuss the possibility that our turnoff stars have been affected by microscopic diffusion. For 15 probable single-star members among the giants, spectrum synthesis leads to abundances of $+0.46^{0.02}_{0.03}$ for [Fe/H]. While less than half the age of NGC 6791, NGC 6253 is at least as metal-rich and, within the uncertainties, exhibits the same general abundance pattern as that typified by super-metal-rich dwarfs of the galactic bulge. ", "introduction": "Given the amount and level of detailed information potentially available for the local region of the Galaxy, the chemical and dynamical evolution of the Galactic disk offers unrivalled insight into the process of galaxy formation and evolution, at least in the context of generating a well-defined spiral galaxy as the end product. As the database for the population components comprising the Galaxy within the neighborhood of the Sun and beyond has grown in size and precision, the simple picture of \\citet{els} has given way to a complex evolutionary history increasingly dominated by multiple substructures from within and multiple mergers from without (see, e.g., \\citet{pra08, fre08, wys08} among many others). To make sense of these complications, it is valuable to view the Sun and the solar neighborhood within the context of the temporal and spatial evolution of the entire disk. The delineation of that global context requires age, metallicity, and kinematic measures for systems that inhabit zones of the Galaxy well within and well beyond the solar circle. While they can be challenging observationally and their number is modest, open clusters increasingly have been selected as the objects of choice in attempting to probe the history of the Galaxy at large distance and over long periods of time \\citep{bra08, fri08}. The primary driver behind this trend is the often-noted benefit of being able to derive the age, composition, and distance from collective analysis of the homogeneous cluster stellar sample, rather than from an individual star. The challenges in interpreting the results from the cluster population are often coupled to small samples over a wide range in parameter space \\citep{car98}, a lack of homogeneity in merging the data from different observers applying different approaches \\citep{che03}, and the real possibility that subgroups of clusters reflect different origins indicative of the complex structure noted above for field stars \\citep{fri06, car06, war09}. Among the various properties of the disk that can be probed using open clusters, two primary questions of interest have been the existence of an age-metallicity relation (AMR) \\citep{pia95, car98} and the delineation of a radial metallicity gradient \\citep{fri95, taat, car98, fri02, che03}. Once one removes the effect of the latter from the sample, the conclusions regarding the absence of the former have changed little over the last 30 years \\citep{hir78, mcc78, fri95, fri02}. However, the failure to discern an AMR among open clusters may be partly due to the limited age range contained within most cluster samples. Numerous studies over the last few decades have concluded that there is little evidence for a significant change in the mean [Fe/H] among stars formed within the solar neighborhood over the last 5-6 Gyrs \\citep{twa80, meu91, edv, gar00, fel01, nor04}, though exceptions to this pattern occasionally arise \\citep{roc00, sou08}. Since few clusters survive to this age due to tidal disruption, most cluster-based determinations of the AMR are dominated by samples where no trend is expected, assuming clusters reflect the same history as the field stars. By contrast, if the absence of an AMR over the last 5-6 Gyrs holds throughout the disk, the abundance gradient of the open cluster sample should be approximately constant with time and, within the statistical uncertainty, should not be dependent upon the mix of ages included within the cluster analysis. As cluster studies extend the temporal and spatial baseline of the sample, the relevance of the objects that populate the extreme ends of the scales becomes an issue of greater concern. For example, the observation that the open cluster sample beyond a galactocentric distance of 10 kpc, on a scale where the Sun is positioned at 8.5 kpc, exhibits, at best, a shallow gradient in metallicity \\citep{taat} has now been confirmed through the addition of clusters beyond 20 kpc from the galactic center \\citep{yon05, ses08}. This immediately raises the question of whether or not the very distant members of this sample are, in fact, representative of the internal development and evolution of the outer Galactic disk or are interlopers accreted via interactions with nearby galactic companions. It should be apparent, however, that this dichotomy may be artificial if the origin and growth of the outer Galactic disk were the products of infalling systems/material from the extended environment of the Galaxy over Gyr timescales \\citep{mag09}. Looking in the opposite direction, cluster studies toward the Galactic center have been hampered by the observational limits imposed by field crowding, high reddening, and a real deficit of surviving clusters interior to a galactocentric radius of 7 kpc. Resulting interpretations of the interior galactic gradient often have been affected by the presence within the sample of perhaps the most anomalous open cluster studied to date, NGC 6791. The extreme nature of the cluster derives from the paradoxical combination of high metallicity and high age, though the degree of the anomaly has remained a point of controversy for decades. A consensus from a variety of studies is that the cluster has an overall metallicity of [Fe/H] = +0.4 $\\pm$ 0.1 based upon intermediate-band photometry \\citep{atm07}, high resolution spectroscopy of stars in various parts of the CMD, including a horizontal branch star \\citep{pg98}, four red giant clump stars \\citep{gr06, car07}, ten red giant stars \\citep{car06}, and two turn-off stars \\citep{boe09}, high resolution infrared spectroscopy of six M giants \\citep{or06}, and medium resolution spectroscopy of 24 giants \\citep{wo03}. Using three different isochrone sets \\citep{de04, van06, pie04}, an age of 8 $\\pm$ 1 Gyr has been determined from isochrone fitting of the cluster turnoff and from the mass-radius diagram of an eclipsing binary \\citep{bed08, gru08}, though \\citet{bed08} derive the puzzling age of 4-6 Gyr from the cooling curve for the cluster white dwarfs. In most models of galactic chemical evolution, attainment of such a high degree of chemical enrichment within such a short time interval after the formation of the disk presents a significant challenge. Some solutions to the problem have been based upon the unusual kinematics of the cluster, implying that the cluster is either an interloper from a region much closer to the galactic center where a steep linear gradient would indicate the existence of clusters of unusually high metallicity or a cluster captured in a merger event \\citep{car06}. Both interpretations are built upon questionable assumptions tied to the chemical composition of the cluster parent population, particularly the latter solution requiring a super-metal-rich cluster from a dwarf system; in fact, dwarf galaxies are metal-poor \\citep{sk,vz}. As long as NGC 6791 represented a unique case, both viewpoints retained some plausibility, but with the discovery of NGC 6253, the picture is significantly altered. The first comprehensive broad-band photometric study of NGC 6253 was carried out by \\citet{bra97}. Based upon multicolor comparisons to the morphology and luminosity functions of theoretical isochrones, \\citet{bra97} concluded that the cluster was $\\sim$ 3 Gyr old with a metallicity nearly twice solar. From broad-band photometry and integrated spectra, \\citet{pia98} found [Fe/H] = +0.2 and an age of 5 $\\pm$ 1.5 Gyr. High-dispersion spectra of four giants by \\citet{car00} indicated [Fe/H] = +0.36 $\\pm$ 0.15, a result used by \\citet{sag01} to obtain an age of 2.5 $\\pm$ 0.6 Gyr. A weakness in all these discussions was the exact value of the reddening, which ranged from E$(B-V)$ = 0.20 to 0.32. The first direct measurement of the cluster reddening and metallicity tied to intermediate-band photometry of a large sample of cluster turnoff stars \\citep{tatd} generated E$(B-V)$ = 0.26 $\\pm$ 0.003, where the error cited is the standard error of the mean (sem), and identified NGC 6253 as the most metal-rich open cluster known, with a metallicity at least as large as NGC 6791 and potentially much larger. A best-fit approach to a match with theoretical isochrones indicated that the cluster CMD morphology was ideally reproduced with an age of 3 $\\pm$ 0.5 Gyrs and alpha-enhanced isochrones. The weak point in the analysis was the extreme nature of the photometric indices which placed the cluster well beyond the limits of the photometric calibration based upon stars in the solar neighborhood. High-dispersion spectroscopic analysis of 4 clump giants by \\citet{car07} has generated [Fe/H] = +0.46 $\\pm$ 0.03, virtually identical to NGC 6791, while \\citet{ses07} find [Fe/H] = +0.36 $\\pm$ 0.07 from 4 probable members. A reanalysis of the intermediate-band photometry of \\citet{tatd} using an improved field star calibration \\citep{atm07} demonstrated that, independent of the adopted reddening, [Fe/H] for NGC 6253 is at least as high as that for NGC 6791 ([Fe/H] = +0.45) and could be 0.1 dex higher. The most recent addition to the cluster database has been supplied by \\citet{mon09}, who used a CCD mosaic to measure proper-motion memberships for stars in a large area around the cluster, as well as supplying a new set of $BVRIJHK$ photometry. Adopting a low reddening (E$(B-V)$ = 0.15) from multicolor fits to theoretical isochrones, \\citet{mon09} find a slightly older age (3.5 Gyr) for the cluster than commonly derived, a possible byproduct of comparisons with isochrones that may be too metal-poor ([Fe/H] = +0.22 and +0.36). While it is clear that, from a chemical composition standpoint, NGC 6253 is a younger cousin of NGC 6791, its extreme characteristics in relation to the rest of the open cluster population and a desire to place it within the context of Galactic chemical evolution make it an important object for study at higher resolution, beyond the handful of clump giants and turnoff stars that have been discussed to date. For this reason, HYDRA spectra of 89 potential cluster members from the tip of the giant branch to below the main sequence turnoff were obtained in multiple wavelength regions for detailed elemental analysis. In this first of three papers we present a discussion of the results for elements other than Li based upon absorption lines in the region of the Li 6708 \\AA\\ line. Future papers will discuss other wavelength regions \\citep{mad10}, as well as the pattern found for the Li abundance as a function of mass and evolutionary state within the cluster \\citep{cum10}. The layout of the paper is as follows: Sec. 2 presents the observations and their reduction, identification of probable members based upon the derived radial velocities coupled with proper-motion membership, and discussion of the photometric system used to define the stellar colors. Sec. 3 details the determination of the fundamental stellar parameters used in the derivation of the abundances from the absorption lines, Sec. 4 lays out the elemental abundance analysis, and Sec. 5 summarizes our conclusions. ", "conclusions": "High-dispersion spectra centered on the Li line of 89 potential members of the old open cluster NGC 6253 have been processed and analyzed, generating 65 probable radial-velocity members, including 47 dwarfs and 18 red giants. When coupled with recently published proper-motion memberships, the member sample is reduced to 45 dwarfs and 18 red giants. Excluding potential binaries and spectra with inadequate S/N per pixel, abundances have been derived from line analysis for 38 stars at the turnoff and from spectrum synthesis for 15 red giants, with impressive agreement between the two samples for the key element, Fe. For the adopted reddening (E$(B-V)$ = 0.22) and effective temperature scale, the mean cluster metallicity lies at [Fe/H] = +0.45, confirming previous work from intermediate-band photometry of turnoff stars and high-dispersion analyses of a handful of mostly red giants that placed NGC 6253 among the most metal-rich systems known in the Galaxy, if not the most metal-rich. The issues dogging any discussion of the metallicity of this object bear a strong similarity to those surrounding NGC 6791 - the exact value of the reddening, the appropriate color and effective temperature scale, and an embarrassing richness of lines that can confuse the interpretation of spectra more than enlighten. Despite these caveats, the choices made in setting the cluster parameters have been selected such that probable changes would raise our metallicity determination for NGC 6253. On an absolute scale, if the effective temperature and microturbulence scales had been set using the approaches found in previous high-dispersion work, our NGC 6253 metallicity would be increased by $\\sim$0.1 dex. We have presented results for three elements other than iron from the wavelength region around the lithium line. In spite of a very limited linelist, we have recovered essentially solar values for [A/Fe] for nickel and silicon. The implied abundance for [Ca/Fe] is surprising and at variance with the results of \\citet{car07} and \\citet{ses07}, both of whom recover an essentially scaled-solar abundance for calcium. Our Ca abundance results rest entirely on a single line at 6717 \\AA. In contrast, both studies from 2007 employ a wider spectral region including 6 lines for the \\citet{car07} study and 22 for the \\citet{ses07} work. A direct comparison of measured equivalent widths for MS/TO stars in our study to published equivalent widths for the 6717 \\AA\\ line indicates that our equivalent width measurements are consistent with those of \\citet{ses07}, implying that our measurement procedures are not faulty but that the line itself might be seriously contaminated. To test this possibility, equivalent widths were synthesized using MOOG for a solar model and a model typical of the MS/TO region's parameters in NGC 6253. For the solar model, the 6717 \\AA\\ line includes a blended feature from Ti II that contributes less than 8$\\%$ to the effective equivalent width of the line. For the NGC 6253 turnoff star, however, the Ti II line accounts for a 20$\\%$ contamination of the line feature. A quick correction to the equivalent width to ``remove\" the contamination from the Ti II line suggests an estimate for [Ca/H] that is nearly 0.3 dex lower, consistent with an essentially scaled-solar abundance for calcium relative to iron. While additional elements, including O, will be the focus of an analysis of a different wavelength region \\citep{mad10}, the pattern emerging from all high-dispersion work done to date is that NGC 6253, like NGC 6791 \\citep{car06, or06, car07, boe09}, exhibits approximately scaled-solar abundances for a majority of the elements. Note that for NGC 6791, the only element that shows statistically significant deviations from a scaled-solar ratio in multiple analyses is C \\citep{or06, car07}; other elements such as Ba and Ni may show deviations in one study, but not in another, or worse, deviations in the opposite sense. As noted by \\citet{boe09}, the pattern of roughly scaled-solar abundance ratios is consistent within the errors with what is observed among very metal-rich field dwarfs found in the solar neighborhood with kinematics that imply a potential origin near the galactic bulge \\citep{pom02}, as well as among probable bulge dwarfs studied {\\it in situ} through spectroscopy obtained during gravitational lensing events \\citep{coh09}. If a link exists between the bulge and NGC 6253, it will need to await additional observational evidence before becoming apparent." }, "1003/1003.0351_arXiv.txt": { "abstract": "We present {\\it Spitzer} MIPS observations at 24 $\\mu$m of 37 solar-type stars in the Pleiades and combine them with previous observations to obtain a sample of 71 stars. We report that 23 stars, or 32\\% $\\pm$ 6.8\\%, have excesses at 24 $\\mu$m at least 10\\% above their photospheric emission. We compare our results with studies of debris disks in other open clusters and with a study of A stars to show that debris disks around solar-type stars at 115 Myr occur at nearly the same rate as around A-type stars. We analyze the effects of binarity and X-ray activity on the excess flux. Stars with warm excesses tend not to be in equal-mass binary systems, possibly due to clearing of planetesimals by binary companions in similar orbits. We find that the apparent anti-correlations in the incidence of excess and both the rate of stellar rotation and also the level of activity as judged by X-ray emission are statistically weak. ", "introduction": "Understanding the evolution of planetary systems has long been central to our speculations about our place in the universe and is a challenging and rapidly advancing topic in astronomy and planetary science. Several hundred extrasolar planets have been discovered in the last 15 years, establishing that planetary systems form frequently. We have also significantly expanded our understanding of the planet formation process, through measurements of protoplanetary disks in the critical mid-infrared through submillimeter regimes \\citep[e.g.,][]{andrews2005,silverstone2006,hernandez2006} and development of detailed two- and three-dimensional numerical models to simulate and interpret these results \\citep[e.g.,][]{alexander2008,dullemond2007}. However, after protoplanetary disks dissipate at about 5 million years, the processes leading to planet formation and its remaining signatures become almost invisible to most types of observation. Planetary debris disks, resulting from collisional activity among planetesimals \\citep{wyatt2008}, are a notable exception and thus are our best current means to characterize planet system evolution. Debris disks are readily detected through the infrared emission from dust particles, which are produced in collisional cascades initiated among larger bodies, e.g., through gravitational stirring by planets. The generation of dust can be further enhanced when the particles become small enough that non-gravitational forces become significant, leading in extreme cases to avalanches of particle generation \\citep{grigorieva2007}. The dust is cleared from the disk on time scales of a thousand to a million years, so it must be replenished. Thus, debris disks are an effective indirect means to probe the current level of collisions in planetary systems. One application of debris disk studies is to characterize the collisional activity as planet systems age \\citep[e.g.,][]{habing2001,spangler2001}. Previous work has shown that debris disks tend to decay with time \\citep{rieke2005,siegler2007}. The excellent sensitivity of the Multiband Imaging Photometer on {\\it Spitzer} \\citep[MIPS;][]{rieke2004} and, more importantly, the accurate photometry it can deliver, have expanded such studies substantially \\citep[e.g.,][]{rieke2005,su2006,gorlova2006,siegler2007,meyer2008,trilling2008,hillenbrand2008,carpenter2009a,balog2009}. In addition to accurate infrared photometry, such studies depend on the accurate determination of stellar ages. Substantial uncertainties remain in even the best age determinations for field stars \\citep[e.g.,][]{mamajek2008}. Therefore, measurements of excesses in stellar clusters, where ages are better determined, can refine the estimates of decay trends. However, the number of clusters close enough for the measurements to probe to solar masses and below is limited, and characterizing the decay trend and other disk parameters is therefore limited by poor statistics \\citep{gaspar2009}. For these reasons, it is important to observe thoroughly the small number of nearby rich clusters available. Only through such measurements can we test whether there are differences in disk behavior around different stellar types, with their accompanying differences in luminosities, masses, and presumably protoplanetary disk masses. For example, initial studies failed to find significant differences in infrared excess decay with respect to spectral type \\citep[e.g.,][]{gorlova2006,siegler2007}, but with improved statistics they seem to be emerging \\citep{gaspar2009}. Removing age as a parameter also allows us to study other parameters affecting debris disk evolution, such as the various mechanisms for grain removal. This paper completes our analysis of {\\it Spitzer} debris disk surveys in the Pleiades. We report 24 $\\mu$m measurements of a sample of 37 stars of late F to early K type in this cluster. We combine these measurements with previous surveys by \\citet{stauffer2005} and \\citet{gorlova2006} to assemble a combined sample of 71 such stars. Our final excess rate provides the highest-weight determination available for debris disk behavior at 115 $\\pm$ 10 Myr, the age of the Pleiades \\citep{meynet1993,stauffer1998,martin2001}. In Section 2, we describe our observations and sample selection. In Section 3, we describe how we obtained excess ratios for our sample using both color-color plots and Kurucz model-fitting. In Section 4, we interpret these excesses through comparisons with similar studies of Praesepe, Blanco 1, and NGC 2547. We also present an analysis of binarity, rotation, and stellar wind drag, as possible ways for the dust to be removed. Finally, in Section 5 we conclude with a summary and some possibilities for future work. ", "conclusions": "We report the results of a new survey of solar-type stars in the Pleiades for 24 $\\mu$m excess indicative of circumstellar debris disks. We combine this work with previous surveys \\citep{stauffer2005,gorlova2006} to build a sample of 71 solar-type stars in this cluster with sufficiently accurate data to identify excesses as small as 10\\% at 24 $\\mu$m. Twenty-three of these stars have excesses at this level or above; on statistical grounds, it is likely that about six additional members have excesses in the 6\\% to 10\\% range, and the remaining 42 stars must have little or no 24 micron excess. The incidence of excesses at 24$\\mu$m and at the age of the Pleiades is high, $\\sim$ 31\\% $\\pm$ 6\\%. We find that the incidence of 24 $\\mu$m excesses for solar-type stars in the Pleiades is slightly smaller than for A stars \\citep[from a general sample mostly in the field:][]{su2006}. It appears that by an age of $\\sim$ 750 Myr, the excesses around solar-type stars have decayed faster than they decay around A stars \\citep{gaspar2009}. The effect probably arises through a mechanism that operates relatively slowly, such as systematic velocity differences in 24$\\mu$m-emitting zones of the debris disks around the two stellar types and the resulting difference in the speed of debris disk evolution. Our study of the Pleiades, plus similar work on other clusters, lets us test aspects of debris disk behavior independently of the evolution of these systems with age. We confirm the results of \\citet{stauffer2009} that close, high-mass binary systems tend not to harbor debris disks. This behavior is probably associated with binary companions that orbit close to the zone where debris disks tend to lie \\citep{trilling2007}. There appear to be anticorrelations between infrared excesses and both rotation ($v$sin$i$) and X-ray luminosity, as also indicated by some previous works. However, we find that these results are not statistically significant and may arise instead from selection effects within the debris disk sample. The excesses around stars with indicated strong winds (from X-ray surface brightnesses) suggest that the wind strengths may be overestimated." }, "1003/1003.0421.txt": { "abstract": "{\\it Hubble Space Telescope} (\\HST) Fine Guidance Sensor astrometric observations of the G4 IV star \\HD ~are combined with the results of the analysis of extensive ground-based radial velocity data to determine the mass of the outermost of two previously known companions. Our new radial velocities obtained with the Hobby-Eberly Telescope and velocities from the Carnegie-California group now span over eleven years. With these data we obtain improved RV orbital elements for both the inner companion, \\HD b and the outer companion, \\HD c. %We also find weak evidence for a tertiary, \\HD d. We identify a rotational period of \\HD~(P$_{rot}=31.65 \\pm 0.17^ d$) with FGS photometry. The inferred star spot fraction is consistent with the remaining scatter in velocities being caused by spot-related stellar activity. We then model the combined astrometric and RV measurements to obtain the parallax, proper motion, perturbation period, perturbation inclination, and perturbation size due to \\HD c. For \\HD c we find P = 2136.1 $\\pm$ 0.3 d, perturbation semi-major axis $\\alpha =1.05 \\pm 0.06$ mas, and inclination $i$ = 48.3\\fdg0 $\\pm$ 4\\fdg0. Assuming a primary mass $M_* = 1.48 M_{\\sun}$, we obtain a companion mass ${\\it M}_c = 17.6 ^{ +1.5}_{-1.2} {\\it M}_{Jup}$, $3\\sigma$ above a 13 \\mjup deuterium burning, brown dwarf lower limit. Dynamical simulations incorporating this accurate mass for \\HD c indicate that a near-Saturn mass planet could exist between the two known companions. %This analysis in the presence of a spot-induced noise floor provide only weak evidence for the identification of a P=197$^d$ signal in the RV residuals as a planetary-mass ($\\sim0.16$\\mjupe) companion. Until confirming, additional observations (radial velocities and/or {\\it Gaia} astrometry) are obtained, we treat this signal simply as a nuisance to be removed in order to obtain an accurate mass for \\HD c. We find weak evidence of an additional low amplitude signal that can be modeled as a planetary-mass ($\\sim$0.17{\\it M}$_{Jup}$) companion at P$\\sim194$ days. Including this component in our modeling lowers the error of the mass determined for \\HD c. Additional observations (radial velocities and/or {\\it Gaia} astrometry) are required to validate an interpretation of \\HD d as a planetary-mass companion. If confirmed, the resulting \\HD~planetary system may be an example of a ``Packed Planetary System\". ", "introduction": "\\HD (= HIP 27253 = HR 1988 = PLX 1320) hosts two known companions discovered by high-precision radial velocity (RV) monitoring \\citep{Fis01,Fis03,Wri09}. Previously published periods were P$_b$=14.31$^d$ and P$_c = 2146^d$ with minimum masses $M_b sini = 0.85$\\mjup and $M_c sini = 13.1$\\mjupe, the latter right above the currently accepted brown dwarf mass limit. A predicted minimum perturbation for the outermost companion, \\HD c, $\\alpha_c = 0.8$ millisecond of arc (mas), motivated us to obtain millisecond of arc per-observation precision astrometry with {\\it HST} with which to determine its true mass (not the minimum mass, $M_c sini$). These astrometric data now span 3.25 years. In the early phases of our project \\cite{Ref06} derived an estimate of the mass of \\HD c from \\HIP, obtaining ${\\it M}_c = 38 ^{ +36}_{-19} $ \\mjupe, well within the brown dwarf 'desert'. Recent comparisons of FGS astrometry with \\HIP, e.g. \\cite{lee07b}, suggest that we should obtain a more precise and accurate mass for \\HD c. Our mass is derived from combined astrometric and RV data, continuing a series presenting accurate masses of planetary, brown dwarf, and non-planetary companions to nearby stars. Previous results include the mass of Gl 876b \\citep{Ben02c}, of $\\rho^1$ Cancri d \\citep{McA04}, $\\epsilon$ Eri b \\citep{Ben06}, HD 33636B \\citep{Bea07}, and HD 136118 b \\citep{Mar10}. \\HD~is a metal-rich G4 IV star at a distance of about 40 pc. The star lies in the 'Hertzsprung Gap' (Murray \\& Chaboyer 2002), a region typically traversed very quickly as a star evolves from dwarf to giant. \\cite{Bai08a} have measured a radius. \\HD~also has a small IR excess found by \\citet{Mor07} with {\\it Spitzer} and interpreted as a Kuiper Belt at 20--50 AU from the primary. Stellar parameters are summarized in Table~\\ref{tbl-STAR}. In Section 2 we model RV data from four sources, obtaining orbital parameters for both \\HD b and \\HD c. We also discuss and identify RV noise sources. In Section 3 we present the results of our combined astrometry/RV modeling, concentrating on \\HD c. We briefly discuss the quality of our astrometric results as determined by residuals, and derive an absolute parallax and relative proper motion for \\HD, those nuisance parameters that must be removed to determine the perturbation parameters for the perturbation due to component c. Simultaneously we derive the astrometric orbital parameters. These, combined with an estimate of the mass of \\HD, provide a mass for \\HD c. Section 4 contains the results of searches for additional components, limiting the possible masses and periods of such companions. In Section 5 we discuss possible identification of an RV signal that remained after modeling components b and c. We discuss our results and summarize our conclusions in Section 6. ", "conclusions": "\\subsection{Discussion} Given the adopted Table~\\ref{tab:allorb3} errors, \\HD c is either one of the most massive exoplanets or one of the least massive brown dwarfs. We can compare our true mass to the (as of January 2010) 69 transiting exoplanetary systems, each also characterized by true mass, not \\msini. As shown in the useful Exoplanetary Encyclopedia \\citep{Sch09} only one companion (CoRoT-Exo-3 b, Deleuil \\etal 2008) has a mass in excess of 13\\mjupe. Whether this 'brown dwarf desert' (Grether \\& Lineweaver, 2006) \\nocite{Gre06} in the transiting sample is due to the difficulty of migrating high-mass companions (bringing them in close enough to increase the probability of transit), or to inefficiencies in gravitational instability formation is unknown. The age of the host star, $\\sim3.3$ By, would suggest that \\HD c has not yet cooled to an equilibrium temperature. An estimated temperature and self-luminosity for a 17\\mjup object that is 3.3 By old can be found from the models of \\cite{Hub02}. Those models predict that \\HD c has an effective temperature, T$_{eff}\\simeq400$K, and L=2.5e-7L$_{\\sun}$, about 20 times brighter than what we estimated using these same models for \\eps~b \\citep{Ben06}. Unfortunately \\HD~has about 16 times the intrinsic brightness of \\eps, erasing any gain in contrast. We note that due to the eccentricity of the orbit, \\HD c is actually within the present day habitable zone for a fraction of its orbit. As \\HD ~continues to evolve and brighten, the habitable zone will move outward and \\HD c will be in that zone for some period of time. If the inner known companion \\HD b is a minimum mass exoplanetary object (assuming $M = M sin i=0.8$\\mjup), our 1 mas astrometric per-observation precision precludes detecting that 2 microsecond of arc signal. Invoking (with no good reason) coplanarity with \\HD c similarly leaves us unable to detect \\HD b. However, with the motivation of our previous result for HD 33636 \\citep{Bea07}, we can test whether or not \\HD b is also stellar by establishing an upper limit from our astrometry. To produce a perturbation, $\\alpha_b > 0.2$ mas (a 3-$\\sigma$ detection, given $\\sigma_{\\alpha}=0.06$ mas from Table~\\ref{tab:allorb3}), and the observed RV amplitude, K$_b$=59\\ms, requires $M_b\\sim0.1M_{\\sun}$ in an orbit inclined by less than 0\\fdg5. Our limit is lower than that established with the CHARA interferometer \\citep{Bai08b}, who established a photometric upper limit of G5 V for the b component. While it might be possible to use 2MASS and SDSS \\citep{Ofe08} photometry of this object to either confirm or eliminate a low-mass stellar companion by backing out a possible contribution from an M, L, or T dwarf, using their known photometric signatures \\citep{Haw02, Cov07}, we lack precise (1\\%) knowledge of the intrinsic photometric properties of a sub-giant star in the Hertzsprung gap with which to compare. %Given that the amplitude of the P$\\sim197^d$ variation is exactly that predicted by stellar activity (Section~\\ref{SR}), we are reluctant to claim the detection of a tertiary in the \\HD~system. The interpretation of this signal as \\HD d is consistent with the results of dynamical simulations and the Packed Planet hypothesis \\citep{Ray09}. While these simulations are necessary to establish companion reality, they are not by themselves sufficient. Confirmation will require additional high-cadence RV observations, or future astrometry. A minimum mass component d would generate a peak to peak astrometric signature of 52 microarcseconds, likely detectable by Gaia \\citep{Cas08}. \\subsection{Conclusions} In summary, radial velocities from four sources, Lick and Keck \\citep{Wri09}, HJS/McDonald \\citep{Wit09b}, and our new high-cadence series from the HET, were combined with \\HST~astrometry to provide improved orbital parameters for \\HD b and \\HD c. Rotational modulation of star spots with a period P=31.66$\\pm$0.17$^d$ produces 0.15\\% photometric variations, spot coverage sufficient to produce the observed residual RV variations. Our simultaneous modeling of radial velocities and over three years of {\\it HST} FGS astrometry yields the signature of a perturbation due to the outermost known companion, \\HD c. Applying the Pourbaix \\& Jorrisen constraint between astrometry and radial velocities, we obtain for the perturbing object \\HD c a period, P = 2136.1 $\\pm$ 0.3 d, inclination, $i$=48.3\\fdg2 $\\pm$ 3.7\\arcdeg, and perturbation semimajor axis, $\\alpha_c = 1.05\\pm 0.06$ mas. Assuming for \\HD~a stellar mass $M_*$ = 1.48$\\pm$0.05$M_{\\sun}$, we obtain a mass for \\HD~c, $M_b$ = 17.6$^{+1.5}_{-1.2}$\\mjup, within the brown dwarf domain. Our independently determined parallax agrees within the errors with \\HIP, and we find a close match in proper motion. Our HET radial velocities combined with others establish an upper limit of about one Saturn mass for possible companions in a dynamically stable range of companion-star separations, $0.2 \\le a \\le 1.2$ AU. RV residuals to a model incorporating components b and c contain a signal with an amplitude equal to the RMS variation with a period, P$\\sim194^d$ and an inferred $a\\sim0.75$ AU. While dynamical simulations do not rule out interpretation as a planetary mass companion, the low S/N of the signal argues for confirmation." }, "1003/1003.3678_arXiv.txt": { "abstract": "{The young active star BD\\,+20\\,1790 is believed to host a substellar companion, revealed by radial-velocity measurements that detected the reflex motion induced on the parent star. }{A complete characterisation of the radial-velocity signal is necessary in order to assess its nature.}{We used CORALIE spectrograph to obtain precise ($\\sim$10\\,m/s) velocity measurements on this active star, while characterizing the bisector span variations. Particular attention was given to correctly sample both the proposed planetary orbital period, of 7.8\\,days, and the stellar rotation period, of 2.4\\,days. }{A smaller radial-velocity signal (with peak-to-peak variations $<$500\\,m/s) than had been reported previously was detected, with different amplitude on two different campaigns. A periodicity similar to the rotational period is found on the data, as well as a clear correlation between radial-velocities and bisector span. This evidence points towards a stellar origin of the radial-velocity variations of the star instead of a baricentric movement of the star, and repudiates the reported detection of a hot-Jupiter.}{} ", "introduction": "Since the discovery of the first exoplanet around 51\\,Peg by \\cite{1995Natur.378..355M}, the radial-velocity method (RV) has established itself as the workhorse for exoplanet detections. It allowed the detection of more than 3/4 of all planets we know and was instrumental in shaping the body of knowledge we gathered on the subject \\citep{2007ARA&A..45..397U}. Still, many open questions remain. For instance the planetary formation process is a subject of great debate. As a consequence, a positive detection of a planet around a young star would be highly valuable; it would put a stringent upper limit on the time scale of planet formation, providing a strong observational constraint for the modeling. In order to address this question, several RV surveys started targeting young objects, only to find that extrasolar planets around these hosts were very rare \\citep{2007ApJ...660L.145S}. On top of that, since young stars exhibit high photospheric and chromospheric activity, these surveys are plagued by stellar activity effects. The detection of an RV signature rooted on atmospheric phenomena is very common and false planetary detections around very young stars were provided by \\cite{2008ApJ...687L.103P} and \\cite{2008A&A...489L...9H}, among several others. Very recently \\cite{2009arXiv0912.2773H} provided compelling evidence for a planet around the young active star BD\\,+20\\,1790. RV observations were associated with a massive hot-Jupiter orbiting the star. Aiming primarily at activity characterisation, several high-frequency RV data sets had been obtained, scattered, spanning six years. The main limitations of the data used were the inadequacy of the time sampling to a planetary search campaign and the low RV precision of the measurements. Still, an extensive activity analysis and careful discussion were presented. The authors showed that some spectral/activity indicators exhibited variation with a time scale of the rotation period (2.8\\,days) but none on a time scale of the reported RV variation (7.8\\,days). This backed the planetary hypothesis, which was considered as the best explanation for the RV variations. Following the announcement of the putative planet we started an intensive RV campaign on BD\\,+20\\,1790 with the CORALIE spectrograph. We analyze our measurements in this paper. In Sect. 2 we present an overview of the parameters both of the star and of the putative planet. In Sect 3 we describe the results of our campaign and discuss these in Sect. 4. We finish by stating our conclusions in Sect. 5. ", "conclusions": "We present strong evidence that the RV variation on the star BD\\,+20\\,1790 is has its origin in the stellar atmosphere rather than being induced by the presence of an unseen companion. These conclusions were drawn from high-cadence, precise RV (at $\\sim$10\\,m/s level) that correctly sampled the previously proposed orbit. Instead of reproducing this orbit, we detected a lower RV variation with variable amplitude. The RV signal is correlated with BIS and shows a periodicity very similar to the reported photometric period. This work shows the importance of correctly sampling the phase of a candidate orbit. Otherwise the conjugated effect of starspots and stellar jitter can be mistaken for a planetary signature." }, "1003/1003.3955_arXiv.txt": { "abstract": "Using new and archival observations made with the \\swift\\ satellite and other facilities, we examine 147 \\xray\\ sources selected from the \\rosat\\ All-Sky-Survey Bright Source Catalog (RASS/BSC) to produce a new limit on the number of isolated neutron stars (INS) in the RASS/BSC, the most constraining such limit to-date. Independent of \\xray\\ spectrum and variability, the number of INSs is $\\leq$48 (90\\% confidence). Restricting attention to soft ($kT_{\\rm eff}<200$\\,eV), non-variable \\xray\\ sources -- as in a previous study -- yields an all-sky limit of $\\leq$31 INSs. In the course of our analysis, we identify five new high-quality INS candidates for targeted follow-up observations. A future all-sky \\xray\\ survey with \\erosita, or another mission with similar capabilities, can be expected to increase the detected population of \\xray-discovered INSs from the 8 to 50 in the BSC, to (for a disk population) 240 to 1500, which will enable a more detailed study of neutron star population models. ", "introduction": "\\label{sec:intro} The \\xray\\ source class of isolated neutron stars (INSs) are observationally defined by their high \\xray\\ to optical flux ratio, $\\fxfopt\\simgt 10^4$ \\citep{treves00}. At discovery, these objects do not exhibit radio pulsations and are not associated with supernova remnants. The study of this phenomenological class ultimately promises a more complete picture of neutron star properties and evolution. Following formation in a supernova explosion (SN), the ``standard'' neutron star (that is, one with a core composed of beta-equilibrium nuclear matter) evolves thermally, cooling to a temperature $T\\sim\\ee{6}$\\, after $t\\sim\\ee{6}$\\,yr \\citep{page04}. The cooling is modestly affected by the presence of strong ($B\\sim\\ee{15}$\\,G) magnetic fields \\citep{arras04}. During this period, if the neutron star is emitting isotropically in the \\xray\\ band, detection and study of the source will be a straightforward issue of improving the flux sensitivity of \\xray\\ instrumentation. This simple picture could be altered if neutron stars, at birth, have strong toroidal magnetic fields, which may significantly modify the isotropy of their early emission \\citep{page07}. Over 1500 radio pulsars have been discovered \\citep{manchester05}; and a few hundred low-mass-X-ray binaries (LMXBs) and high mass \\xray binaries (HMXBs) are catalogued \\citep{liu06,liu07}. However, star formation models predict $N_{\\rm MW}\\sim\\ee{9}$ neutron stars have formed within the Milky Way \\citep{timmes96}, implying that the vast majority of stellar evolution remnants have not yet been observed. While one can attempt to infer the natal properties of neutron stars observed as radio pulsars and \\xray\\ binaries, selection effects are present in both populations. Radio pulsars exhibit significant surface dipolar magnetic fields ($B\\sim\\ee{8}-\\ee{13}$\\,G; \\citealt{camilo94,camilo00}); the distribution of magnetic field strength and geometry at birth, and their subsequent time evolution, are not strongly observationally constrained; it may be that only a fraction of neutron stars are born with such magnetic fields, implying that relatively few NSs are ever observable as radio pulsars. The uncertain emission mechanism(s) of radio pulsars have made it challenging to quantify the observational selection effects which impact their detection in radio surveys. Nonetheless, recent analyses modelling this class concludes that the properties of known pulsars and pulsar surveys, including best estimates of the above selection effects, can account for all neutron stars observed \\citep{faucher06}. Observation and characterization of the INS population (which, in the present work, will include all neutron stars which are first identified via their high \\fxfopt) provides a means to more fully characterize the neutron star population of the Galaxy, without selections related to magnetic field strength, geometry, evolution, and the various phases (and consequences) of binary stellar evolution. Modelling of the INS population is, however, susceptible to theoretical uncertainties in neutron star cooling mechanisms, including the possibility of enhanced neutrino emissivity due to hyperon condensates \\citep{page04}. Determining the ages (to $\\simlt$50\\%) of individual isolated thermally-emitting neutron stars, independent of the observed surface temperature, is currently only possible for objects with precision parallax and proper motion measurements (e.g., \\citealt{kaplan02b,kaplan07}), which allow their space velocity to be back-extrapolated to a likely birthplace in an OB-association. In the absence of accurate ages for individual INSs, comparing individual INS properties with cooling models is effectively an interpretive act rather than a confrontation of theory with data. To test these models, therefore, will ultimately require a sufficient number of observed INSs such that models of their birth properties, evolution, and dynamics can be challenged to reproduce the observed population, as is now usefully done for the radio pulsar population \\citep{cordes98,arzoumanian02,gonthier04,faucher06}. INSs evolve in the absence of influence of a companion (such as accretion of $\\sim$0.1-0.3 \\msun over the lifetime of a low-mass X-ray binary), and so offer an opportunity to study the physics of compact objects directly -- their formation, dynamics, thermal evolution and magnetic field evolution. Initial estimates \\citep{blaes93} for the number of INSs which would be detected in the \\rosat\\ All-Sky-Survey Bright Source Catalog (RASS/BSC \\citep{voges99}) overestimated the number which were eventually observed, by a factor of $\\sim$100 (as discussed previously; \\cite[R03 hereafter]{rutledge03}). This overestimation was due to the fact that the observable population's luminosity was thought to be dominated by simple Bondi-Hoyle accretion, which scales approximately with the neutron star velocity $\\propto v^{-3}$. However, the velocity distribution for radio pulsars was later found to underestimate the typical velocity by a significant factor, which subsequently altered the conclusions by decreasing both the luminosity and the number of detectable INSs. Later estimates, which employ results of MHD simulations of Bondi-Hoyle accretion, found that the modified Bondi formula dramatically suppresses the observable population even further, to the extent that none are expected detected above RASS/BSC sensitivity \\citep{perna03}. Thus, the interpretation of INSs observed from the RASS/BSC has been as post-natal cooling from neutron stars \\citep{popov00b,popov01,popov03,popov06,popov06b,posselt07,posselt08}. The present work continues our search for INSs -- \\xray-bright sources with no observed off-band emission at discovery. This follows our previous work (R03), using essentially the same selection methods for INS candidates detected by the RASS/BSC. We have expanded the selection of INS candidates based on \\pnoid\\ -- the probability that the RASS/BSC X-ray source is not associated with any off-band counterparts in the catalogs considered (USNO-A2; NVSS and IRAS) -- from $\\pnoid\\geq0.90$ to $\\pnoid\\geq 0.8$. Doing so increases the fraction of all INSs detected within the analyzed region of the RASS/BSC ($\\delta\\geq -39\\deg$, due to the sky-coverage of NVSS) which pass our selection, from $\\sim$20\\% (R03) to 30\\% (in the present work); this insures a greater number of INSs detected in the RASS/BSC end up in our selected INS candidates list. To find 100\\% of the INSs within the RASS/BSC, an analysis such as the one performed here would be required for all 18806 RASS/BSC X-ray sources. The present analysis makes use of a statistical selection which reduces the number of sources to be analyzed by a factor of x100 (from 18806 RASS/BSC sources, to 147 sources), and will permit identification of almost 1/3 of all INSs among RASS/BSC sources. To identify 100\\% of the INSs in the RASS/BSC sources, thus will require approximately 100$\\times$ the amount of observations analyzed here, to obtain a factor of 3 more INSs. Thus, searches for INSs and classification of their populations using the techniques described here involve a balance between observational resources invested (in \\swift\\ and \\chandra\\ observations), and the number of INSs which will be identified. The size of the sample analyzed here is a compromise between these two goals, providing a significant improvement over previous characterization of the INS population in the RASS/BSC \\cite{rutledge03}. We also report follow-up observations with \\swift, as well as archival observations with \\chandra, \\xmm, and \\rosat/HRI, which permit localization and association with off-band (optical, IR and UV) sources. As previously, we note that the present analysis produces an upper-limit not only on objects that follow the strict definition of an INS - non-radio pulsar, non-magnetar, and non-SNR associated objects - but of high X-ray/optical flux ratio-selected populations of {\\em all types} (such as anomalous X-ray pulsars (AXPs)). It should be noted that in our analysis, we may refer to objects as ``INS'' based on their X-ray/optical flux ratio and not on other properties that may technically exclude it from the phenomenological class. As an example, one object in our selection sample, 1RXS~J141256.0+792204 (also known as Calvera) was identified as an isolated compact object of some kind based on its X-ray/optical flux ratio \\citep{rutledge08}. Although more observations are necessary to determine its classification as either an INS, AXP, or radio pulsar, we classify it in our analysis as an ``INS''. ", "conclusions": "\\label{sec:conclude} We have placed upper-limits on the number of INSs -- indepdendent of spectral and variability properties -- in the RASS/BSC, of $\\leq$48 (90\\% confidence) and $\\leq$70 (99\\% confidence). When we limit ourselves to spectrally soft, non-variable INSs, these limits are $\\leq$31 (90\\% confidence) and $\\leq$46 (99\\% confidence) -- a factor of 2 lower than limits derived previously (R03). Using the same analysis as in previous work ($\\S$ 5.2, R03), we can place a limit on supernova rate of progenitors to these objects, using a naive and optimistic model for the resulting sources. The following assumptions are made for this model: \\begin{enumerate} \\item An INS production rate of $\\gamma_{-2}$ per 100 yr \\item INSs are produced in a flat disk with constant SNe rate per unit area, out to a radius of $R_{\\rm disk}=15\\,{\\rm kpc}$ \\item INSs are produced velocities of $<10\\, {\\rm kpc}\\, {\\rm Myr^{-1}}$ perpendicular to this disk \\item The resulting INSs maintain a luminosity 2\\tee{32} \\cgslum for a period of $\\tau=\\ee{6}\\, {\\rm yr}$ in the \\rosat/PSPC passband of 0.1-2.4 keV \\item The INSs exhibit an $\\alpha=3$ spectrum \\item The limit on all INSs (assuming $\\leq$50) corresponds to a limit \\item The effects of absorption (which are likely to be important) are neglected (permitting INSs to be detected to a distance of 10.3 kpc) \\end{enumerate} We can then place a limit of $\\gamma_{-2}<0.019$ (90\\% confidence), corresponding to an INS birthrate of one per 5300 years. If we assume a spatial average absorption of $N_H=3\\tee{21}\\, {\\rm cm^{-2}}$ (limiting detection to sources within a distance of 1.5 kpc), the limit is $\\gamma_{\\rm -2}<$0.8. In the process of obtaining these limits, we have identified one new confirmed INS \\cite[Calvera]{rutledge08}. After second-epoch \\xray\\ observations with \\swift, we find 9 INS candidates. We are presently following-up observationally on these candidates, with better X-ray localizations and deep optical imaging, to place limits of $F_X/F_{\\rm opt}>1000$, on these sources." }, "1003/1003.4869_arXiv.txt": { "abstract": "We present a catalogue of 179 hyperluminous infrared galaxies (HLIRGs) from the Imperial IRAS-FSS Redshift (IIFSCz) Catalogue. Of the 92 with detections in at least two far infrared bands, 62 are dominated by an M82-like starburst, 22 by an Arp220-like starburst and 8 by an AGN dust torus. On the basis of previous gravitational lensing studies and an examination of HST archive images for a further 5 objects, we estimate the fraction of HLIRGs that are significantly lensed to be 10-30$\\%$. We show simple infrared template fits to the SEDs of 23 HLIRGs with spectroscopic redshifts and at least 5 photometric bands. Most can be fitted with a combination of two simple templates: an AGN dust torus and an M82-like starburst. In the optical, 17 of the objects are fitted with QSO templates, several with quite strong extinction. There are 5 objects fitted with galaxy templates in the optical, two of which show evidence for AGN dust tori and so presumably contain Type 2 (edge-on) QSOs. The remaining object is fitted with a galaxy template in the optical, but is of such high luminosity that this classification would be plausible only if there were very strong lensing. 20 of the 23 objects (87$\\%$) show evidence of an AGN either from the optical continuum or from the signature of an AGN dust torus, but the starburst component is the dominant contribution to bolometric luminosity in 14 out of 23 objects (61 $\\%$). The implied star-formation rates, even after correcting for lensing magnification, are in excess of 1000 $M_{\\odot} yr^{-1}$. We use infrared template-fitting models to predict fluxes for all HLIRGs at submillimetre wavelengths, and show predictions at 350 and 850 $\\mu$m. Most would have 850 $\\mu$m fluxes brighter than 5 mJy so should be easily detectable with current submillimetre telescopes. At least 15 $\\%$ should be detectable in the {\\it Planck} all-sky survey at 350 $\\mu$m and all {\\it Planck} all-sky survey sources with z $<$ 0.9 should be IIFSCz sources. From the luminosity-volume test we find that HLIRGs show strong evolution. A simple exponential luminosity evolution applied to all HLIRGs would be consistent with the luminosity functions found in redshift bins 0.3-0.5, 0.5-1 and 1-2. The evolution-corrected luminosity function flattens towards higher luminosities perhaps indicating a different physical mechanism is at work compared to lower luminosity starbursts. In principle this could be gravitational lensing though previous searches with HST have, perhaps surprisingly, not shown lensing to be widely prevalent in HLIRGs. ", "introduction": "Rowan-Robinson (2000, hereafter RR2000) defined hyperluminous infrared galaxies to be those with bolometric infrared luminosities $L_{ir} > 10 ^{13} L_{\\odot}$. He discussed a sample of 39 such galaxies found either by follow-up of the IRAS survey or through submillimetre observations, and modelled the infrared and submillimetre spectral energy distributions (SEDs) for those that were well-observed. Submillimetre observations and more detailed SED models for 13 hyperluminous infrared galaxies (HLIRGs) were reported by Farrah et al (2002a) and ISO data for 4 HLIRGs was discussed by Verma et al (2002). A more detailed SED model for the prototypical HLIRG IRAS F10214+4724 was give by Efstathiou (2006). Recently Ruiz et al (2010) have modelled the SEDs of 13 hyperluminous galaxies which are X-ray sources. Many of these galaxies show evidence of AGN dust torus emission at rest-frame wavelengths 3-30 $\\mu$m, but the submillimetre emission is almost always due to star-formation, an interpretation that is supported by the large gas-masses found in these galaxies (see RR2000 for a summary). 7 of the 39 objects in RR2000 were already known to be gravitationally lensed. A further 9 hyperluminous infrared galaxies were imaged with HST by Farrah et al (2002b), but none were found to be lensed. Even after correction for magnification by gravitational lensing, most of these galaxies imply star formation rates $> 1000 M_{\\odot} yr^{-1}$. Since many of the hyperluminous infrared galaxies in RR2000 were discovered through their submillimetre radiation, it is clearly of interest to consider whether all or most submillimetre galaxies are hyperluminous. Many discussions of the nature of 'submillimetre galaxies' assume the latter to be the case. However there is growing evidence that the submillimetre population is quite diverse. In their models for the SEDs of a complete sample of 850 $\\mu$m sources from the SHADES survey, Clements et al (2008) find a diverse range of SED types, including M82 starbursts, 'cirrus' (ie quiescent) galaxies, and AGN dust tori, as well as Arp220-type starbursts. Efstathiou and Rowan-Robinson (2003) have pointed out that the SEDs of as many as half of submillimetre galaxies could be interpreted in terms of a cirrus model. In modelling source-counts from 8-1100 $\\mu$m, and the far infrared and submillimetre background radiation, Rowan-Robinson (2008) concludes that at S(850) = 8 mJy there are approximately equal numbers of M82 and Arp220 starbursts. However cirrus galaxies, in these models, dominate the counts only at S(850) $>$ 100 mJy and at S(850) $<$ 1 mJy. The creation of the Imperial IRAS FSC Redshift Catalogue (IIFSCz, Wang and Rowan-Robinson 2009a) gives us the opportunity to define an almost all-sky ($|b| > 20^0$) catalogue of hyperluminous infrared galaxies, sampled at 60 $\\mu$m. This allows us to study the evolution and luminosity function of this population and to predict what will be seen in submillimetre surveys to be undertaken by {\\it Planck} and {\\it Herschel}. A spatially flat cosmological model with $\\Lambda$ = 0.7, $h_0$=0.72 has been used throughout. ", "conclusions": "We present a catalogue of 179 hyperluminous infrared galaxies from the Imperial IRAS-FSS Redshift (IIFSCz) Catalogue (Wang and Rowan-Robinson 2009a). 14 of these were included in the analysis of Rowan-Robinson (2000). We use the infrared template-fitting models of Wang and Rowan-Robinson (2009a) to predict fluxes at 350 and 850 $\\mu$m. Most would have 850 $\\mu$m fluxes brighter than 5 mJy so should be readily detectable by current submillimetre telescopes. At least 15 $\\%$ should be detectable in the {\\it Planck} all-sky survey at 350 $\\mu$m and all {\\it Planck} all-sky survey sources with z $<$ 0.9 should be IIFSCz sources. From the luminosity-volume test we find that HLIRGs show strong evolution. A simple exponential luminosity evolution applied to all HLIRGs would be consistent with the luminosity functions found in redshift bins 0.3-0.5, 0.5-1 and 1-2. The evolution-corrected luminosity function flattens towards higher luminosities perhaps indicating a different physical mechanism is at work compared to lower luminosity starbursts. In principle this could be gravitational lensing though previous searches with HST have not shown lensing to be widely prevalent in HLIRGs. Although for an interesting minority of HLIRGs the dominant contribution to the bolometric infrared luminosity is due to an AGN dust torus, star formation appears to be the dominant factor in the overwhelming majority of cases, with star-formation rates $>$ 1000 $M_{\\odot} yr^{-1}$ being implied. Even if the duration of these extreme starbursts is shorter than that for less luminous starbursts (Rowan-Robinson 2000), they still represent an extremely interesting and dramatic phase of galaxy evolution." }, "1003/1003.3214_arXiv.txt": { "abstract": "We analyze the quasar two-point correlation function (2pCF) within the redshift interval $0.810\\,h^{-1}$~Mpc the parameter describing the large-scale infall to density inhomogeneities is $\\beta=0.63\\pm0.10$ with the linear bias $b=1.44\\pm0.22$ that marginally (within 2$\\sigma$) agrees with the linear theory of cosmological perturbations. We discuss possibilities to obtain a statistical estimate of the random component of quasars velocities (different from the large-scale infall). We note rather slight dependence of quasars velocity dispersion upon the 2pCF parameters in the region $r<2$~Mpc. ", "introduction": "\\label{sec:1} \\indent\\indent Spatial distribution of quasars is one of the main sources of observational information about the largescale structure of the Universe; this is at present almost the only source of statistical information about matter inhomogeneity and clustering at cosmological redshifts. The most important results on quasars clustering were obtained using two largest quasar surveys up to date: the 2-degree Field Quasar Survey with 2QZ catalogue as a result (http://www.2dfquasar.org, \\citet{Croom_1998}) and the Sloan Digital Sky Survey (SDSS; http://www.sdss.org), which has been completed with the 7th release \\citep{Abazadjian_2009}. The two-point correlation functions (2pCF) \\citep{peebles_book} of quasars $\\xi(r)$ are important characteristics of matter spatial inhomogeneity that may be compared to cosmological theories of structure formation. Development of 2dF and SDSS surveys gave a powerful incentive to investigation of various aspects of correlation functions of quasars (see, e.g., \\citet{croom_2005,myers_2006} and references therein). In the present paper we use the 7-th data release of SDSS to study 2pCF of quasars. The main problems concerning reconstruction of the cosmological mass distribution from redshift surveys are as follows. The first problem is that the surveys of extragalactic objects give us an information only about distribution of the luminous matter which is biased relative to the dark matter \\citep{dekel_1987}. Biasing may depend on the physical peculiarities of extragalactic objects, e.g., on morphological type \\citep{einasto_2007,ross_2007}, luminosity \\citep{beisbart_2000,sorrentino_2006}, color-index \\citep{coil_2007}, star formation rate \\citep{owners_2007} etc. and evolves with redshift \\citep{croom_2005,myers_2006,myers_2006_1,weinstein_2004,porciani_2004,daAngela_2008,Mountrichas_2009}. Commonly used supposition of linear biasing means that the density variations of certain kind of objects (in our case quasars) is proportional to that of the whole matter; therefore \\begin{equation}\\label{bias1} \\xi(r)=b^2\\,\\xi_{m}(r), \\end{equation} where $\\xi_{m}(r)$ is 2pCF of the whole matter, $\\xi(r)$ is 2pCF of quasars and $b$ is the bias parameter. The second problem of the 3-dimensional analysis of the matter distribution is due to the fact that the observed redshifts of extragalactic objects are `contaminated' by measurement errors and non-Hubble motions. The distances calculated from these redshifts without taking into account the unknown peculiar velocities are called distances in redshift-space (in contrast to the real space). As our Universe is isotropic, the correlation function must be spherically symmetric in the real-space. But in the redshift-space it appears to be distorted. On smaller scales the profile of galaxies 2pCF is stretched along the line of sight (`Finger of God' effect) due to virial velocities of objects inside the galaxy clusters (this can be neglected for quasars if we consider $z>1$), their random velocities and redshift errors. This effect is especially noticeable for quasar pairs with the projected linear separations $\\lesssim2$~Mpc, where the SDSS data are expected to be incomplete \\citep{hennawi_06}. On larger scales the 2pCF profile is flattened along the line of sight (`The Bull's eye' or \\citet{kaiser_1987} effect) due to the gravitational infall to density inhomogeneities; this kind of redshift-space distortion dominates on the linear scales. These effects are parameterized by line-of-sight pairwise velocity dispersion $\\langle w^{2}\\rangle^{1/2}$ and infall parameter $\\beta$. It is worth to note that the non-Hubble motions and redshift errors are not the only sources of redshift-space distortions. One more effect of geometric flattening that can lead to distortion of 2pCF can be due to a wrong choice of cosmological parameters $\\Omega_{M}$, $\\Omega_{\\Lambda}$; this provides an additional tool for estimation of these parameters by means of a geometrical test of \\citet{alcock_paczynski_1979}. The redshift-space 2pCF of quasars can be in principle used for estimation of all parameters ($\\beta$, $\\langle w^{2}\\rangle^{1/2}$ and cosmological ones) simultaneously. But due to a degeneracy between the geometric distortions and the redshift-space distortions (see, e.g., \\citet{hoyle_2002}) this problem is complicated. \\citet{hoyle_2002} and \\citet{daAngela_2005} proposed a method that allows to break this degeneracy by means of combination of the \\citet{alcock_paczynski_1979} test with that based on evolution of the quasar clustering amplitude, which has a different dependence on $\\beta(\\bar{z})$ and $\\Omega_{M}(0)$. As we cannot estimate the non-Hubble motion of each quasar independently, the effects of the redshift-space distortion are at present the only source of statistical information about proper velocities of quasars. On the other hand, these effects prevent direct determination of 3D 2pCF that involves line-of-sight distances between objects determined from $z$-measurements. This urges us to use only the projected distances $\\sigma$ (i.e. orthogonal to the line of sight, which may be considered as independent of proper motions) to determine the projected 2pCF and then to restore the real-space 2pCF. It is well known that such reconstruction of the real-space 2pCF from projected one is mathematically ill-posed problem. However, most authors avoid this difficulty using a concrete functional form of 2pCF. Typically 2pCF is represented in a power-law form $\\xi(r)=\\left(r_0/r\\right)^{\\gamma}$, though it is clear that 2pCF slope and correlation length may be different on different interval. For example, \\citet{daAngela_2005} showed that double power-low model gives a good fit to 2pCF of quasars distribution with $\\xi(r)=\\left(6.0/r\\right)^{1.45}$ over scales of $10.01$ were also excluded. \\begin{figure} \\centering \\epsfig{figure=fig_1_sky.eps,width=8cm} \\caption{Sky coverage of our \\textit{full} quasar sample in equatorial coordinates.} \\label{fig_1} \\end{figure} \\begin{figure}\\centering \\epsfig{figure=fig_2_zdistr.eps,width=8cm} \\caption{Redshift distribution of our quasar samples: \\textit{full} (solid line), \\textit{uniform} (dashed), \\textit{low-reddening} (dash-dot), \\textit{high-reddening} (dash-dot-dot-dot), \\textit{good} (long dashes)} \\label{fig_2} \\end{figure} \\begin{figure}\\centering \\epsfig{figure=fig_2a_mdistr.eps,width=8cm} \\caption{Luminosity $-$ redshift distribution for \\textit{uniform} sample.} \\label{fig_2a} \\end{figure} The third step was examination of all objects left after the previous steps by eye (photometry and spectroscopy data) using the SDSS data base. During these examination 466 objects were excluded as the ones having too faint and noisy spectra, or being fault `double' objects (i.e. single objects processed twice and thus appeared twice in the list with similar coordinates), or being artifacts of observations, or being stars (we have found 5 stars in this sample). Our final sample, which we call \\textit{full}, contains 52303 objects and has the mean redshift $\\bar{z}=1.47$. The sky coverage of the sample and its redshift distribution are shown in Figs.~\\ref{fig_1} and \\ref{fig_2} (solid line). The redshift-luminosity distribution of this sample is shown in Fig.~\\ref{fig_2a}, where the effect of $m_{i}=19.1$ magnitude limit can be seen. Here the values of absolute magnitudes in $i$-band, $M_{i}$, are calculated using $K$-correction from \\citet{richards_2006}. \\subsection{Sample inhomogeneity}\\label{sec:2.2} \\indent\\indent As one can see from Fig.~\\ref{fig_1} the sky coverage by our sample in not heterogeneous: some parts seem to be significantly denser then others. This is due the fact that these parts were twice spectroscopically covered and can influence our results. First of all we note that the algorithms we use to estimate the 2pCF parameters (sections 3.1 and 4.1) are just worked out in order to reduce the selection effects due to unphysical inhomogeneity \\citep{ivzh_10,zhdan_ivashch2008}. Nevertheless, to test the influence of inhomogeneity we constructed 50 \\textit{homogeneous} samples in the following way. The whole sky area covered by our sample was divided into $2.5^{\\circ}\\times2.5^{\\circ}$ patches and the density of each patch, mean density ($\\bar{n}=7.25$ objects per square degree) and its rms ($\\sigma_{n}=2.55$ objects per square degree) were calculated. All the objects from the patches with density exceeding the mean one more than 0.5rms were excluded (see Fig.~\\ref{fig_1a}). Then we generated 50 new \\textit{homogeneous} samples by putting the excluded objects in their places, but choosing (in a random way) only those making the density of the given patch equal to the mean density. The sky coveredge of one of these samples is presented in Fig.~\\ref{fig_1b}. Each new test sample has 35643 objects and a density $\\bar{n}=6.35\\pm1.19$ objects per deg$^{2}$. Apparent decrease of the density with increasing declination is the result of projection of the celestial sphere onto the plane. \\subsection{Tilable targets}\\label{sec:2.3} \\indent\\indent Most quasar candidates in the SDSS are selected based on their locations in multidimentional SDSS colour space and their cross-identification with radio sources from the Faint Images of the Radio Sky at Twenty-cm (FIRST) survey \\citep{richards_2002}. Supplementing this primary quasar sample are quasars targeted by the GALAXY, X-RAY, STAR and SERENDIPITY SDSS software packages; no attempt at completeness was made for the last three categories. Thus, the sample we use is not complete and the possible effect of this incompleteness on our results has to be investigated. For this purpose we constructed the sample of the so-called `tilable' objects (\\textit{uniform}) from our sample. According to SDSS glossary web-page, tilable targets are supposed to have as closed to uniform completeness as possible. These are targets with primTarget flag `QSO\\_HIZ', `QSO\\_CAP', `QSO\\_SKIRT', `QSO\\_FIRST\\_CAP', `QSO\\_FIRST\\_SKIRT', `GALAXY\\_RED', `GALAXY', `GALAXY\\_BIG', `GALAXY\\_BRIGHT\\_CORE', `STAR\\_BROWN\\_DWARF', and second target flag `HOT\\_STD'. The number of objects in the subsample constructed following this criterion is 37290. One can see that the sky coverage by the \\textit{uniform} sample presented in Fig.~\\ref{fig_1c} is close to uniform and close to the sky coverage by our \\textit{homogeneous} samples (Fig.~\\ref{fig_1b}). The redshift distribution of this \\textit{uniform} sample is presented in Fig.~\\ref{fig_2} (dashed line). \\begin{figure} \\centering \\epsfig{figure=holes_0.5sigma_2.5deg.eps,width=8cm} \\caption{Sky coverage of the \\textit{full} sample with excluded patches with density higher than 0.5rms.} \\label{fig_1a} \\end{figure} \\begin{figure}\\centering \\epsfig{figure=new_0.5sigma_2.5deg.eps,width=8cm} \\caption{Sky coverage of one of the \\textit{homogeneous} samples.} \\label{fig_1b} \\end{figure} \\begin{figure}\\centering \\epsfig{figure=tiled.eps,width=8cm} \\caption{Sky coverage of the \\textit{uniform} sample.} \\label{fig_1c} \\end{figure} \\subsection{Reddening}\\label{sec:2.4} \\indent\\indent As the main criterion of quasar candidates selection in SDSS is based on their magnitudes and colors, that have been corrected for Galactic extinction using the maps of \\citet{schlegel_1998}, any systematic errors in the reddening model can induce additional effects on clustering results (see \\citet{ross_2009} for discussion of possible effects). That is why following \\citet{ross_2009} we devided our sample into low reddening ($0.0028L_0$, where $L$ is some characteristic comoving scale where 2pCF approaches to zero. We also make the following assumptions. First, we neglect variations of $n_{i}$ on the scales $ 10$~Mpc, where the linear theory is applicable. Here the parameter $\\beta=f(\\Omega_{M},z)/b$ appears as some consequence of the large-scale infall onto matter overdensities, $b$ is the bias parameter in Eq.~(\\ref{bias1}). The function $f(\\Omega_{M},z)$ is expressed by means of the growing mode of the density contrast \\citep{peebles_book}. There is a simple analytical approximation \\citep{carroll_1992} for this function in the case of the spatially-flat $\\Lambda$CDM cosmological model: \\begin{equation}\\label{eq:fom} f(\\Omega_{M},z)\\approx\\left[\\frac{\\Omega_{M}(1+z)^{3}}{\\Omega_{M}(1+z)^{3}+1-\\Omega_{M}}\\right]^{4/7} \\end{equation} Therefore we may proceed in two ways: either (i) by taking the slope $\\gamma_s=\\gamma$ from the results of Sec.~\\ref{sec:3} dealing with fitting of $s_0$ only, or (ii) by independent determination of $\\gamma_s,\\, s_0$. Like in Sec.~\\ref{sec:3} one can determine the total number of neighbours of the $i$-th quasar from the whole sample in a spherical layer $[s,s+\\Delta s]$ as \\[ DD(s,\\Delta s)=\\sum_{i}n_{i}\\left[1+\\frac{4\\pi}{\\Delta V}\\int_{s}^{s+\\Delta s}\\xi(s')s'^{2}ds'\\right]\\Delta V, \\] \\noindent where $\\Delta V=4\\pi \\Delta(s^2+s\\Delta s+(\\Delta s)^2/3)$. The similar estimation for random catalogue, which is considered to represent random spatial distribution of objects with no clustering, is \\[ RR(s,\\Delta s)=4\\pi \\sum_{i} n'_{i}\\Delta V. \\] \\noindent Assuming $\\sum\\limits_{i}n'_{i}\\approx\\sum\\limits_{i}n_{i}$ and taking into account the power-low form of the monopole part of 2pCF $\\xi_0(s)=\\left(s_{0}/s\\right)^{\\gamma_s}$, we have the relation for fitting: \\begin{equation}\\label{fit} \\frac{DD(s,\\Delta s)}{RR(s,\\Delta s)}-1=\\frac{s_{0}^{\\gamma_s}}{3-\\gamma_s} \\cdot \\frac{(s+\\Delta s)^{3-\\gamma_s}-s^{3-\\gamma_s}}{\\Delta s\\,(s^{2}+s\\Delta s+\\Delta s^{2}/3)}. \\end{equation} In order to determine the infall parameter $\\beta$ on the way (i) we use Eq.~(\\ref{fit}) to find $s_0$, with the slope $\\gamma_s=\\gamma$ being known from the results of the Sec.~\\ref{sec:3}. Then we derive $\\beta$ by solving the Eq.~(\\ref{eq_beta}) with known $s_0,\\,r_0$, and then we estimate bias $b$ for our redshift interval using Eq.~(\\ref{eq:fom}). These operations deal with the region $s>10$~Mpc. In addition, mainly for comparison with the results of the other authors, we also considered way (ii) with independent determination of $\\gamma$ and $s_0$ for different separation intervals. Note that the method we use for estimation of the infall parameter is similar to the so-called $J_{3}(s)/J_{3}(r)$-method (see e.g. \\citet{ratcliffe_IV_1998}), which uses the volume integrals $J_{3}(x)=\\int\\limits_{0}^{x}\\xi(y)y^{2}dy$. Using of these integrals allows to smooth the influence of small sizes of the samples. Our approach has the same advantage as it includes the 2pCF parameters which are the averaged characteristics of 2pCFs within the whole distance range. \\subsection{Results and discussion}\\label{sec:4.2} \\begin{figure} \\centering \\epsfig{figure=zspace_in.eps,width=8cm} \\caption{The quasar redshift-space 2pCF, $\\xi(s)$, for \\textit{full} sample. The quoted errorbars are jackknifes. The best fit single power-low over our full range $1-35\\,h^{-1}$~Mpc and over ranges $1-10\\,h^{-1}$~Mpc and $10-35\\,h^{-1}$~Mpc with parameters presented in Table~2 are shown with solid, dotted and dashed lines correspondingly.} \\label{fig_zsp_in} \\end{figure} \\indent\\indent We estimated the 2pCF parameters within the whole range $s=1-35$~$h^{-1}$~Mpc and separately for ranges $1-10$~$h^{-1}$~Mpc and $10-35$~$h^{-1}$~Mpc. The results are presented in Table~2 and Fig.~13. The errors are calculated using the jackknife technique. We also found the same parameters for distance range $1-25$~$h^{-1}$~Mpc for comparison with the results of \\citet{croom_2005} and \\citet{ross_2009} presented in the last 3 rows of this table. As one can see our parameters $s_{0}$ and $\\gamma_{s}$ within the distance range $1-25$~$h^{-1}$~Mpc agree well with the results of \\citet{croom_2005} and \\citet{ross_2009} for 2QZ catalogue and the 4th edition of the SDSS quasar catalogue \\citep{Schneider_2007} correspondingly. And for distance range $1-10$~$h^{-1}$~Mpc our results agree with \\citet{croom_2005} within 2$\\sigma$ for the slope and 1$\\sigma$ for the correlation length. \\begin{table*}\\label{tab:s-space} \\centering \\begin{minipage}{170mm} \\caption{The parameters of the redshift-space 2pCF and the results of other authors for comparision. Here DR5Q stands for quasar catalogue by Schneider et al. (2007).} \\centering \\begin{tabular}{|c|c|c|c|c|c|c|c|} \\hline dist. range, $h^{-1}$Mpc & $z$ range & mean $z$ & $s_{0}$, h$^{-1}$ Mpc & $\\gamma_{s}$ & $\\chi^{2}/d.o.f.$ & sample & authors\\\\ \\hline & & & $5.83\\pm0.47$ & $1.49\\pm0.12$ & 1.440 & \\textit{initial} & \\\\ & & & $5.95\\pm0.86$ & $1.44\\pm0.21$ & -- & \\textit{homogeneous} & \\\\ & & & $5.33\\pm0.51$ & $1.35\\pm0.08$ & 1.473 & \\textit{uniform} & \\\\ $s=1-35$ & $0.8-2.2$ & 1.47 & $5.76\\pm0.54$ & $1.46\\pm0.12$ & 1.363 & \\textit{good} & present paper \\\\ & & & $5.45\\pm0.96$ & $1.48\\pm0.24$ & 0.779 & \\textit{low reddening} & \\\\ & & & $4.56\\pm0.55$ & $1.24\\pm0.11$ & 1.892 & \\textit{high reddening} & \\\\ \\hline $s=1-25$ & & & $5.97\\pm0.51$ & $1.37\\pm0.13$ & 1.060 & & \\\\ $s=1-10$ & $0.8-2.2$ & 1.47 & $6.43\\pm0.63$ & $1.21\\pm0.24$ & 0.456 & \\textit{initial} & present paper \\\\ $s=10-35$ & & & $7.37\\pm0.81$ & $1.90\\pm0.24$ & 1.636 & & \\\\ \\hline $s=1-25$ & $0.3-2.2$ & 1.35 & $5.48^{+0.42}_{-0.48}$ & $1.20\\pm0.10$ & -- & 2QZ & \\cite{croom_2005}\\\\ $s=1-10$ & $0.3-2.2$ & 1.35 & $3.88^{+0.43}_{-0.53}$ & $0.86^{+0.16}_{-0.17}$ & -- & 2QZ & \\cite{croom_2005}\\\\ $s=1-25$ & $0.3-2.2$ & 1.27 & $5.95\\pm0.45$ & $1.16^{+0.11}_{-0.16}$ & -- & SDSS DR5Q & \\cite{ross_2009} \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*} \\begin{table*}\\label{tab:b-and-beta} \\centering \\begin{minipage}{170mm} \\caption{The values of infall parameter $\\beta$ and bias parameter $b_{Q}$ and the results of other authors for comparison. Here DR5Q stands for quasar catalogue by Schneider et al. (2007), DR4phot stands for a sample of photometrically classified quasars from the 4th Data Release of SDSS. The numbers in brackets correspond to methods proposed by \\protect\\citet{hoyle_2002} -- [1], \\protect\\citet{daAngela_2005} -- [2], \\protect\\citet{croom_2005} -- [3], and \\protect\\citet{porciani_2004} -- [4]; see explanations in the text.} \\centering \\begin{tabular}{|c|c|c|c|c|c|c|} \\hline mean $z$ & $b_{Q}$ & $\\beta$ & method & sample & authors\\\\ \\hline 0.55 & $1.4\\pm0.2$ & $0.55\\pm0.10$ & mean of [1],[2] & 2SLAQ+2QZ+SDSS DR5 & \\cite{Mountrichas_2009} \\\\ 1.27 & $2.06\\pm0.03$ & $0.43$ & [3] & SDSS DR5Q & \\cite{ross_2009} \\\\ 1.35 & $2.02\\pm0.07$ & $0.44\\pm0.02$ & [3] & 2QZ & \\cite{croom_2005}\\\\ 1.40 & $1.50\\pm0.20$ & $0.60^{+0.14}_{-0.11}$ & [2] & 2QZ+2SLAQ & \\cite{daAngela_2008} \\\\ 1.40 & $2.84^{+1.49}_{-0.57}$ & $0.32^{+0.09}_{-0.11}$ & $\\xi(s)/\\xi(r)$ & 2QZ & \\cite{daAngela_2005} \\\\ 1.40 & -- & $0.50^{+0.13}_{-0.15}$ & [2] & 2QZ & \\cite{daAngela_2005} \\\\ 1.40 & $2.41\\pm0.08$ & -- & [3] & SDSS DR4phot & \\cite{myers_2006_1} \\\\ 1.47 & $2.42^{+0.20}_{-0.21}$ & -- & [4] & 2QZ & \\cite{porciani_2004} \\\\ \\hline & $1.44\\pm0.22$ & $0.63\\pm0.10$ & & \\textit{initial} & \\\\ & $1.37\\pm0.22$ & $0.64\\pm0.10$ & & \\textit{homogeneous} & \\\\ 1.47 & $1.50\\pm0.37$ & $0.61\\pm0.15$ & decribed in text & \\textit{uniform} & present paper \\\\ & $1.35\\pm0.23$ & $0.67\\pm0.12$ & & \\textit{good} & \\\\ & $1.30\\pm0.44$ & $0.70\\pm0.24$ & & \\textit{low reddening} & \\\\ & $1.54\\pm0.51$ & $0.67\\pm0.12$ & & \\textit{high reddening} & \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*} Comparing the parameters for different distance ranges we can conclude that on the large scales the redshift-space 2pCF is steeper than on the small scales while the correlation lengthes do not differ significantly within 1$\\sigma$. This difference in the slope of the redshift-space 2pCF of quasars on different scales was pointed out earlier \\citep{croom_2005}. It is important to note that the results for the slope on $10-35\\,h^{-1}$~Mpc interval fairly well agree with the results for $\\gamma$ in the real space, as it must be. \\begin{figure} \\centering \\epsfig{figure=zspace_res.eps,width=8cm} \\caption{Correlation lengthes and slopes of the redshift-space 2pCF for 50 \\textit{homogeneous} samples.} \\label{fig_sp} \\end{figure} \\begin{figure} \\centering \\epsfig{figure=beta_params_35.eps,width=8cm} \\caption{Correlation lengthes (for fixed slope value) and infall parameters for 50 \\textit{homogeneous} samples.} \\label{fig_bs} \\end{figure} As discussed in Section~4.1, we have used the correlation length $r_{0}$ and the slope $\\gamma$ of the real-space 2pCF obtained in Section~\\ref{sec:3}. We have then obtained the redshift-space correlation length $s_0=7.27\\pm0.37\\,h^{-1}$~Mpc for the fixed $\\gamma$ on the scales $>10\\,h^{-1}$~Mpc and we have found $\\beta$ and $b$ using Eqs.~(\\ref{eq_beta}) and (\\ref{eq:fom}). The results are given in Table~3 that contains the obtained values of infall and bias parameters and results of other authors with indication of the method, used for their estimation. One of them (denoted as [1]), containing simultaneous determination of $\\beta$ and cosmological model, was proposed by \\citet{hoyle_2002} and includes fitting of the observed $\\xi(\\sigma,\\pi)$ by theoretical model from \\citet{peebles_book} (see also \\citet{ratcliffe_IV_1998}) with the relation for infall velocity from \\citet{hale-sutton_1990}. The improvement of this method proposed in \\citet{daAngela_2005} (denoted as [2]) allows to break the degeneracy between $\\Omega_{M}$ and $\\beta$ and is based in the consideration of the quasar clustering evolution. Another approach proposed in \\citet{croom_2005} (denoted as [3]) suggests the solution of the quadratic equation in $b_{Q}$ with the shape of $\\xi(r)$ governed by the underlying dark matter distribution described by the analytic equation from \\citet{hamilton_1991,hamilton_1995}. In \\citet{daAngela_2008} and \\citet{ross_2009} the last approach was applied with theoretical dark matter power spectrum from \\citet{smith_2003}. One more method which was used in \\citet{porciani_2004} (denoted as [4]) includes using the relation between projected 2pCFs of the quasars and that of the matter. As distinct from a number of authors (aimed, in particular, at estimations of the dark matter halo masses) we concentrate to the comparison with the linear theory; i.e. our results on $b_{Q}$ are completely based on the interval $r>10\\,h^{-1}$~Mpc, where there is no problems with the linear theory. For test we also provided the same calculations for redshift-space 2pCF, $\\beta$ and $b$ for \\textit{homogeneous}, \\textit{uniform}, \\textit{good}, \\textit{low reddening} and \\textit{high reddening} samples. The scattering of the resulting redshift-space 2pCF parameters and infall parameters with the corresponding values of $s_{0}$ for 50 \\textit{homogeneous} samples are presented in Figs.~\\ref{fig_sp} and \\ref{fig_bs}. Their mean values are presented in Tables~2 and 3 together with results for other samples. For all our samples errors of the parameters are jackknifes, except the \\textit{homogeneous} where errors are rms. All these results agrees well (within 1.3$\\sigma$) with the results for initial sample. ", "conclusions": "\\label{sec:7} \\indent\\indent We have performed a detailed analysis of 2pCF using the quasar sample from the SDSS DR7 data for objects with $0.810\\,h^{-1}$~Mpc the parameter describing the large-scale infall to density inhomogeneities is $\\beta=0.63\\pm0.10$ with the linear bias to be $b=1.44\\pm0.22$ that marginally (within 2$\\sigma$) agrees with the linear theory of cosmological perturbations. \\item We point out that the value of the quasars velocity dispersion varies almost within errors for a wide range of model parameters that determine 2pCF for $r<2$~Mpc. \\end{enumerate} We also performed investigations of our sample and showed that we can neglect boundary effects and all inhomogeneities of the sample caused by non-uniformity of the sky coverage by the spectroscopic survey, reddening and conditions of observations. It is worth to note, that the values of $\\langle w^2 \\rangle^{1/2}$ from Tables~4,\\,5 have been calculated on account of SDSS DR7 data on pairs with the projected separations $\\sigma<2$~Mpc, which are not complete and the number of pairs involved in treatment of the velocity dispersion is not too large. In view of these circumstances we consider our results on $\\langle w^2\\rangle^{1/2}$ rather as motivation to further studies to observe larger sample of $\\sim 1$~Mpc quasar pairs with accurate redshift measurements." }, "1003/1003.1301_arXiv.txt": { "abstract": "Wasp\\,-\\,10\\,b is a very interesting transiting extrasolar planet. Although its transit is very deep, about 40\\,mmag, there are very different estimates of its radius in the literature. We present new photometric observations of four complete transits of this planet. The whole event was detected for each transit and the final light curve consists of more than 1500 individual CCD exposures. We determine the following system parameters: planet to star radius ratio $R_{\\mathrm{p}}/R_{*}=0.168 \\pm 0.001$, star radius to semimajor axis ratio $R_{*}/a = 0.094 \\pm 0.001$ and inclination $i=87.3 \\pm 0.1$\\,deg. Assuming that the semimajor axis is $0.036\\,9 \\pm ^{0.0012}_{0.0014}$\\,AU (Christian et al. 2009), we obtain the following radius of the planet $R_{\\mathrm{p}}=1.22 \\pm 0.05\\,R_{\\mathrm{J}}$, and radius of the star $R_{*} = 0.75 \\pm 0.03$. The errors include the uncertainty in the stellar mass and the semimajor axis of the planet. Surprisingly, our estimate of the planet radius is significantly higher (by about 12\\,percent) than the most recent value of Johnson et al. (2009). We also improve the orbital period $P_{\\mathrm{orb}}=3.092\\,731 \\pm 1\\times 10^{-6}$\\,days and estimate the average transit duration $T_{\\mathrm{D}}=0.097\\,4\\pm 8\\times 10^{-4}$\\,days. ", "introduction": "\\begin{figure} \\centerline{ \\includegraphics[width=6.0cm,clip=]{prvni.eps} \\includegraphics[width=6.0cm,clip=]{druhy.eps} } \\centerline{ \\includegraphics[width=6.0cm,clip=]{treti.eps} \\includegraphics[width=6.0cm,clip=]{ctvrty.eps} } \\caption{The observations from the following nights from top left: 4/11/08, 31/8/09, 4/10/09 and 7/10/09 (dd/mm/yy).} \\label{f1} \\end{figure} Transiting extrasolar planets are the VIP's (very important planets) of the extrasolar planet community (exoplanets), being the key to the understanding of the physics and various processes in the interior and atmospheres of substellar objects. The orientation of their orbits in space with respect to the observer is so fortunate that they allow us to observe the transit of the planet in front of the parent star as a dip in the light curve. Photometric observations of these transits (in combination with spectroscopy) provide us with important parameters of the system: orbital period, inclination, mass and radius of the planet. The planet radius is a complicated function of the planet mass, age, chemical composition as well as properties of its orbit and the parent star (Guillot \\& Showman 2002, Burrows et al. 2007, Fortney et al. 2007, Baraffe et al. 2008, Leconte et al. 2010). Precise knowledge of planetary radii is essential for further progress in the field. Wasp\\,-\\,10\\,b is one of the 80 so far known transiting extrasolar planets (Schneider, 1995). It was discovered in 2008 by the Wide Angle Search for Planets (WASP) Consortium (Christian et al. 2009). The exoplanet orbits the parent star GSC 2752-00114 (spectral type K5) at the distance of 0.036\\,9\\,AU and has the mass of $2.96\\,M_{\\mathrm{J}}$ (Christian et al. 2009). Wasp\\,-\\,10\\,b has one of the deepest transit depth, cca 39\\,mmag (Poddany et al. 2010). This results in the planet radius of about $1.28\\,R_{\\mathrm{J}}$ (Christian et al. 2009). This radius is quite large and an additional internal heat source is required to understand such large radii. Unfortunately, their most precise observations do not cover the whole transit events. Soon after the discovery the properties of the exoplanet were revisited by Johnson et al. (2009). They obtained superb observations with the UH 2.2m telescope which covered the whole transit and significantly improved the precision of the radius determination. However, contrary to the discovery paper, they determined a very different radius of the planet: $R_{\\mathrm{p}}=1.08\\,R_{\\mathrm{J}}$, which is 16 percent smaller than that given in the discovery paper. The origin of the difference was not entirely clear. Nevertheless, the new radius is consistent with previously published theoretical radii for irradiated Jovian planets. Miller et al. (2009) could explain its new radius without any additional heat sources. ", "conclusions": "We have presented the observation and analysis of four light curves of exoplanetary system Wasp\\,-\\,10\\,b. This enabled us to improve the orbital period and revisit the properties of the system. The resulting values of the parameters together with their uncertainties are given in Table \\ref{tabulka}. For comparison, the parameters from previous works are added there, too. Figure \\ref{f2} shows the resulting best fit light curve together with measured data from all four nights. Our results indicate that the radius of the planet is significantly larger than the latest estimate of Johnson et al. (2009). Our observations are in agreement with the original observations of Christian et al. (2009), but are more precise. This is quite surprising since Johnson et al. (2009) had seemingly superior observations\\footnote{Recently, after the submission of our paper, new independent observations and analysis of Wasp\\,-\\,10\\,b with similar conclusions were carried out by Dittmann et al. (2010).}. Nevertheless, our value of the planet radius indicates that an additional alternative mechanism is required to inflate the planet size. Enhanced opacities in the atmosphere (Burrows et al. 2007) or a tidal heating mechanism studied recently by Miller et al. (2009), Ibgui et al. (2010) and Leconte et al. (2010) might be invoked to explain the inflated radius of Wasp\\,-\\,10\\,b." }, "1003/1003.4181_arXiv.txt": { "abstract": "We show that cold clumps in the intra--cluster medium (ICM) efficiently lose their angular momentum as they fall in, such that they can rapidly feed the central AGN and maintain a heating feedback process. Such cold clumps are predicted by the cold feedback model, a model for maintaining the ICM in cooling flows hot by a feedback process. The clumps very effectively lose their angular momentum in two channels: the drag force exerted by the ICM and the random collisions between clumps when they are close to the central black hole. We conclude that the angular momentum cannot prevent the accretion of the cold clumps, and the cold feedback mechanism is a viable model for a feedback mechanism in cooling flows. Cold feedback does not suffer from the severe problems of models that are based on the Bondi accretion. ", "introduction": "\\label{s-intro} For more than a decade now it is clear that the intra-cluster medium (ICM) in cooling flow (CF) clusters of galaxies and CF galaxies must be heated, and the heating process should be stabilized by a feedback mechanism (see reviews by \\citealp{Pet06} and \\citealp{McN07}). However, in many cases the heating cannot completely offset cooling (e.g., \\citealp{Wis04, McN04, Cla04, Hic05, Bre06, Sal08, Wil09}) and some gas cools to low temperatures and flows inward (e.g., \\citealp{Pet06}). The mass inflow rate is much below the one that would occur without heating, and the flow is termed a moderate cooling flow \\citep{Sok01, Sok03, Sok04}. Two basic modes have been proposed to feed the super massive black hole (SMBH) at the heart of the active galactic nucleus (AGN) in CF clusters (e.g., \\citealp{Com08}). They can be termed the hot and the cold feedback mechanisms. In the hot-feedback mode the hot gas from the vicinity of the SMBH is accreted, such as in the Bondi accretion process (e.g., \\citealp{Omm04, All06, Rus10}). More sophisticated schemes that use direct calculations of the inflow of the hot gas, but are basically similar to the Bondi accretion process, also exist (e.g., \\citealp{Cio09, Cio10}). However, the Bondi accretion process suffers from two severe problems \\citep{Sok06, Sok09}, to the point that it fails. First, there is no time for the feedback to work. The accretion of gas from the very inner region does not have time to respond to gas cooling in the outer regions ($r \\ga 10 \\kpc$). Second, although in galaxies the accretion rate can in principle account for the AGN power \\citep{All06, Cio09, Cio10}, in CF clusters the Bondi accretion rate is much too low \\citep{Raf06}. In a recent paper it was argued that the Bondi accretion process fails to maintain feedback in the process of galaxy formation as well \\citep{Sok09b}. In the cold feedback model \\citep{Piz05, Sok06, Piz07} the mass accreted by the central black hole originates in non-linear over-dense blobs of gas residing in an extended region of $r \\sim 5 - 30 \\kpc$; these blobs are originally hot, but then cool faster than their environment and sink toward the centre (see also \\citealp{Rev08}). The mass accretion rate by the central black hole is determined by the cooling time of the ICM, the entropy profile, and the presence of inhomogeneities \\citep{Sok06}. In the cold feedback model if gas cools at large distances, so do the overdense blobs. In a relatively short time they feed the SMBH, and the ICM is heated on a time scale shorter than its cooling time scale. This overcomes the main problem of models based on the Bondi accretion process \\citep{Sok09b}. \\citet{Wil09} suggest that the behaviour and properties of the cold clumps they observe in the cluster A~1664 support the cold feedback mechanism. In general, the presence of large quantities of cold ($T \\la 10^4 \\K$) gas in CF clusters (e.g., \\citealp{Edg02, Edw07}) suggests that the non-linear perturbations that are required to form the dense clumps do exist in CF clusters. The main arguments against the cold feedback mechanism was that the cold clumps supposed to feed the SMBH have too much angular momentum, and they cannot fall directly to the SMBH: for this reason, their accretion rate will be too slow and the response time too long (e.g., \\citealp{Rus10}). In this paper we show that dense clumps lose their angular momentum rapidly enough, such that there is no angular momentum problem in the cold feedback mechanism. The equations are presented in Sections~\\ref{s-eqns} and~\\ref{s-angular}, and the cluster properties used for the calculations are given in Section~\\ref{s-prof}. In Section~\\ref{s-numerical} we present the solution for falling clumps, both without (Sec.~\\ref{s-lzero}) and with (Sec~\\ref{s-l}) angular momentum. In Section~\\ref{s-collision} the loss of angular momentum via collision is estimated. Our main conclusions are in Section~\\ref{s-summary}. ", "conclusions": "\\label{s-summary} We have studied the role of the cold clumps' angular momentum in the framework of the cold feedback model, generalizing the results of Paper~I that did not include angular momentum. We first studied falling clumps with zero angular momentum. The efficiency of accretion increases (i.e., clumps are less stable) as the size of the clump decreases (for a given density contrast) and its cross section increases (for a given volume and density contrast), as evident from Figs.~\\ref{f-kdelta}-\\ref{f-atracks}. The reason is that a relatively large cross section implies slower infall, such that the clump has time to increase its overdensity due to radiative cooling. We then include angular momentum. As it is apparent from Figs.~\\ref{f-jtracks} and~\\ref{f-jtimes}, at large radii the drag force is the main agent cause for clumps to lose their angular momentum, and it is efficient enough to remove the centrifugal barrier of accreted clumps. The clumps can easily reach distances of $r \\la 100 \\psec$. At this stage, the clumps may be very dense and cool ($T \\la 10^4 \\K$), and they are no more in pressure equilibrium with the surrounding ICM. As discussed in Section~\\ref{s-collision}, the collisions between clumps become important, and even for small volume filling factors ($\\epsilon_{V} \\simeq 10^{-4}-10^{-2}$) the collision time scale is comparable or shorter than the infall time. The residual orbital angular momentum is lost, and the clumps are free to accrete on the central galaxy and ultimately feed the SMBH sitting at at the clusters' centre. This accretion powers the AGN and activate an efficient feedback with the ICM cooling. In our calculations the ICM was static. However, along the jets axis there is an outflow, and clumps will not fall there freely. We expect most of the clumps to be accreted from the equatorial plane. The interaction of the jets and rising bubbles with dense clumps will be the subject of a future study. We also note that although the AGN is the main heating source, the falling clumps do release gravitational energy \\citep{Fab03}. The solution of the angular momentum problem has removed the main theoretical objection to the cold feedback mechanism as a viable model for ICM heating in CF clusters. Observational support to the cold feedback mechanism (e.g., \\citealp{Wil09}) stand now on a more solid ground." }, "1003/1003.2558_arXiv.txt": { "abstract": "Multiple outbursts of type Ia SNe in one galaxy present a unique opportunity to study the homogeneity of these objects. NGC 3147 is only the second known galaxy with three SNe Ia, another one is NGC 1316. We present CCD $UBVRI$ photometry for SN Ia 2008fv and compare the light and color curves of this object with those for SNe Ia discovered earlier in NGC 3147: 1972H and 1997bq. The photometric properties of SNe 1997bq and 2008fv are nearly identical, while SN 1972H exhibits faster declining light curve. ", "introduction": " ", "conclusions": "" }, "1003/1003.0900_arXiv.txt": { "abstract": "Angular momentum transport owing to hydrodynamic turbulent convection is studied using local three dimensional numerical simulations employing the shearing box approximation. We determine the turbulent viscosity from non-rotating runs over a range of values of the shear parameter and use a simple analytical model in order to extract the non-diffusive contribution ($\\Lambda$-effect) to the stress in runs where rotation is included. Our results suggest that the turbulent viscosity is of the order of the mixing length estimate and weakly affected by rotation. The $\\Lambda$-effect is non-zero and a factor of 2--4 smaller than the turbulent viscosity in the slow rotation regime. We demonstrate that for Keplerian shear, the angular momentum transport can change sign and be outward when the rotation period is greater than the turnover time, i.e.\\ when the Coriolis number is below unity. This result seems to be relatively independent of the value of the Rayleigh number. ", "introduction": "Turbulence due to the convective instability is thought to account for much of the angular momentum transport in the outer layers of the Sun and other stars with convection zones (e.g.\\ R\\\"udiger 1989; R\\\"udiger \\& Hollerbach 2004). In the presence of turbulence the fluid mixes efficiently and diffusion processes occur much faster than in its absence. This effect is usually parameterized by a turbulent viscosity $\\nut$ that is much larger than the molecular viscosity $\\nu$. Often the value of $\\nut$ is estimated using simple mixing length arguments with $\\nut=\\urms l/3$, where $\\urms$ is the rms velocity of the turbulence and $l=\\alpha_{\\rm MLT} H$ where $\\alpha_{\\rm MLT}$ is a parameter of the order unity and $H$ is the vertical pressure scale height. Numerical results from simpler fully periodic isotropically forced systems suggest that the mixing length estimate gives the correct order of magnitude of turbulent viscosity (e.g.\\ Yousef et al.\\ 2003; K\\\"apyl\\\"a et al.\\ 2009a; Snellman et al.\\ 2009). However, it is important to compute $\\nut$ from convection simulations in order to see whether the results of the simpler systems carry over to convection. Furthermore, it is of interest to study whether the small-scale turbulent transport can be understood in the light of simple analytical closure models that can be used in subgrid-scale modeling. Measuring $\\nut$ and its relation to averaged quantities, such as correlations of turbulent velocities, is one of the main purposes of our study. In addition to enhanced viscosity, turbulence can also lead to non-diffusive transport. The $\\alpha$-effect (e.g.\\ Krause \\& R\\\"adler 1980), responsible for the generation of large-scale magnetic fields by helical turbulence, is one of the most well-known non-diffusive effects of turbulence. An analogous effect exists in the hydrodynamical regime and is known as the $\\Lambda$-effect (Krause \\& R\\\"udiger 1974). The $\\Lambda$-effect is proportional to the local angular velocity and occurs if the turbulence is anisotropic in the plane perpendicular to the rotation vector (R\\\"udiger 1989). The existence of the $\\Lambda$-effect has been established numerically from convection simulations (e.g.\\ Pulkkinen et al.\\ 1993; Chan 2001; K\\\"apyl\\\"a et al.\\ 2004; R\\\"udiger et al.\\ 2005) and simpler homogeneous systems (K\\\"apyl\\\"a \\& Brandenburg 2008). If, however, both shear and rotation are present it is difficult to disentangle the diffusive and non-diffusive contributions. This is particularly important in the case of accretion disks where the sign of the stress determines whether angular momentum is transported inward or outward. Convection is commonly not considered as a viable angular momentum transport mechanism in accretion disks since several studies have indicated that the transport owing to convection occurs inward (e.g.\\ Cabot \\& Pollack 1992; Ryu \\& Goodman 1992; Stone \\& Balbus 1996; Cabot 1996; R\\\"udiger et al.\\ 2002). Furthermore, in an influential paper, Stone \\& Balbus (1996, hereafter SB96) presented numerical simulations of hydrodynamic convection where the transport was indeed found to be small and directed inward on average. This result was used to provide additional evidence for the importance of the magneto-rotational instability (Balbus \\& Hawley 1991) as the main mechanism providing angular momentum transport in accretion disks. Although we agree with the conclusion that hydrodynamic turbulence is ineffective in providing angular momentum transport, we are concerned about the generality of the result of SB96. There are now some indications that hydrodynamic turbulence may not always transport angular momentum inward (cf.\\ Lesur \\& Ogilvie 2010). In order to approach the problem from a more general perspective, Snellman et al.\\ (2009) studied isotropically forced turbulence under the influence of shear and rotation and found that the total stress, corresponding to the radial angular momentum transport in an accretion disk, can change sign as rotation and shear of the system are varied in such a way that their ratio remains constant. They found that the stress is positive for small Coriolis numbers, corresponding to slow rotation. In what follows we show that outward transport is possible also for convection in a certain range of Coriolis numbers. When rotation is slow, the Reynolds stress is positive, corresponding to outward transport. In the regime of large Coriolis numbers the Reynolds stress changes sign and is directed inward, as in the study of SB96. The remainder of the paper is organized as follows: our numerical model is described in Section~\\ref{sec:model} and the relevant mean-field description in Section~\\ref{sec:mf}. The results and the conclusions are given in Sections~\\ref{sec:results} and \\ref{sec:conc}, respectively. ", "conclusions": "\\label{sec:conc} The present results have shown that hydrodynamic convection is able to transport angular momentum both inward and outward, depending essentially on the value of the Coriolis number, in accordance with earlier results from homogeneous isotropically forced turbulence (Snellman et al.\\ 2009). This underlines the importance of considering comprehensive parameter surveys and not to rely on demonstrative results from individual case studies. For given value of the Coriolis number, the stress is found to be relatively independent of the value of the Rayleigh number (Section~\\ref{DependenceReynolds}). By varying shear and rotation rates separately, we have been able to quantify the relative importance of diffusive and non-diffusive contributions to the Reynolds stress tensor. In agreement with earlier work, it turns out that the turbulent kinematic viscosity is of the order of the mixing length estimate and has roughly the same value as the turbulent magnetic diffusivity found earlier for similar runs (K\\\"apyl\\\"a et al.\\ 2009b). In other words, the turbulent magnetic Prandtl number is around unity, again in accordance with results from simpler systems (e.g.\\ Yousef et al.\\ 2003). The other important turbulent transport mechanism in rotating turbulent bodies is the $\\Lambda$ effect. Although the importance of this effect is well recognized in solar and stellar physics (e.g.\\ R\\\"udiger \\& Hollerbach 2004), it is not normally considered in the context of accretion disks. In the present paper we have been able to quantify its importance for a range of Coriolis numbers by means of a simple analytical model making use of the minimal $\\tau$-approximation. For slow rotation the coefficient $\\Lambda_{\\rm H}$ is of the order of $\\nutz$ and independent of the Coriolis number. However, once the Coriolis number exceeds a value around unity, our method of separating the turbulent viscosity and the $\\Lambda$-effect breaks down, which reinforces the need for a truly independent determination not only of diffusive and nondiffusive contributions to the Reynolds stress, but also of all the components of the full stress tensor." }, "1003/1003.2134_arXiv.txt": { "abstract": "{Though there is increasing evidence linking the moat flow and the Evershed flow along the penumbral filaments, there is not a clear consensus regarding the existence of a moat flow around umbral cores and pores, and the debate is still open. Solar pores appear to be a suitable scenario to test the moat-penumbra relation as evidencing the direct interaction between the umbra and the convective plasma in the surrounding photosphere, without any intermediate structure in between. } {The present work studies solar pores based on high resolution ground-based and satellite observations.} {Local correlation tracking techniques have been applied to different-duration time series to analyze the horizontal flows around several solar pores.} {Our results establish that the flows calculated from different solar pore observations are coherent among each other and show the determinant and overall influence of exploding events in the granulation around the pores. We do not find any sign of moat-like flows surrounding solar pores but a clearly defined region of inflows surrounding them.} {The connection between moat flows and flows associated to penumbral filaments is hereby reinforced by this work.} ", "introduction": "The solar photosphere displays a wide variety of magnetic features at different spatial scales being sunspots and pores the more conspicuous ones. The morphology and evolution of sunspots and pores have been extensively studied and both are thought to be ruled by the mutual effect of magnetic field that inhibits convection (though not completely) and plasma motions. Nevertheless, there is not clear consensus for a model explaining the transition from pores into sunspots (i.e.\\ development of penumbra, see the monograph by \\cite{thomasweiss08}), the fade of the Evershed flow once leaving the penumbra and entering the region dominated by the large outflows (the so-called moat flows), the flow patterns surrounding pores which are, alternatively, dominated by downflows surrounding them \\citep{giordano08} or moat-like flows \\citep{zuccarello09} and, also, the effect of the penumbra in the granular pattern surrounding sunspots.\\\\ By using local correlation tracking techniques \\cite{vargas2007} found a direct correlation between the presence of penumbrae and the appeareance of moat flows in a complex $\\delta$-configuration active region. A more extensive sample taking into account different penumbral configurations was analyzed by \\cite{vargas2008}, establishing a systematic moat-penumbra relation in all the sunspots under study. According to these studies, no moat flow was detected in the granulation next to umbral boundaries lacking penumbrae. Moat flows were always detected as a prolongation of the penumbral filaments once crossing the penumbral boundary. \\\\ Though there is increasing evidence linking the moat flow and the Evershed flow along the penumbral filaments \\cite[e.g.][]{sainz2005,cabrera2006}, the debate regarding the existence of a moat flow around umbral boundaries without penumbra and individual pores is still ongoing. In a recent work, \\cite{deng2007}, found that the dividing line between radial inward and outward proper motions in the inner and outer penumbra, respectively, survived the decay phase, suggesting that the moat flow is still detectable after the penumbra disappeared. However, previous works \\citep{sobotka1999, roudier2002, hirzberger2003} have measured horizontal proper motions in and around pores and have observed some penetrating flows at the umbral boundaries and a ringlike arrangement of positive divergences (\\emph{rosettas}) around the pores which is related to a continuous activity of exploding granules in the granulation around them. \\cite{roudier2002} identified a very clear inflow around pores which corresponds to the penetration of small granules and granular fragments from the photosphere into the pores, pushed by granular motions originated in the divergence centres around them. These authors conclude that the motions at the periphery of the pore are substantially and continuously influenced by the external plasma flows deposited by the exploding granules. It is important to note in this context that the annular area surrounding pores and filled with exploding granules generates an outward directed flow annulus that can give the impression of a persistent outflow. \\\\ Converging flows around pores have also been observationally reported previously by \\cite{wang1992} as well as downflows at their periphery \\citep{sankara2003}. More recently, 3D magnetohydrodynamic simulations found horizontal flows towards the pores that contribute to mantain their magnetic structure together \\citep{cameron2007}.\\\\ \\begin{table*} \\centering \\caption{Characteristics of the time series of solar pores observed from ground-based and satellite facilities.} \\begin{tabular}{cccccccc}\\\\ Telescope & Date 2007 & Series & Time, UT& Duration \\footnotesize[min] & N. images & Cadence \\footnotesize[sec] & FOV \\footnotesize[$\\arcsec$]\\\\\\hline\\hline\\\\ \\multirow{2}{*}{\\emph{SST}} & \\multirow{2}{*}{30 Sep} & 1 & 08:43-09:31 & 48 & 286 & 10 & 64.3$\\times$65.0 \\\\ & & 2 & 09:36-09:56 & 20 & 118 & 10 & 64.3$\\times$65.0 \\\\\\hline\\\\ \\multirow{3}{*}{\\emph{Hinode}} & \\multirow{2}{*}{1 Jun} & 1 & 21:35-21:55 & 20 & 40 & 30 & 27.9$\\times$55.7 \\\\ & & 2 & 22:26-23:33 & 67 & 134 & 30 & 27.9$\\times$55.7 \\\\\\cline{2-8}\\\\ & 30 Sep& 1 & 00:14-17:59 & 960 & 1030 & 60 & 55.7$\\times$111.4 \\\\\\hline \\end{tabular} \\label{poroseries} \\end{table*} \\begin{figure} \\vspace{-1.5cm} \\includegraphics[width=1.\\linewidth]{./imagenes.ps} \\vspace{-2cm} \\caption{SST co-temporal and co-spatial set of images from the emerging flux region on 30 September 2007. Dark regions in the Dopplergram denote upflows and the map ranges from approx. -1.2 (\\emph{white}) to 1.6 (\\emph{black}) km s$^{-1}$. The LOS Magnetogram shows negative/positive (\\emph{black/white}) polarities with field strengths up to$~\\sim$2300 G. The coordinates are expressed in arc sec.} \\label{images_red} \\end{figure} Pores are interesting to analyze since, as they do not display penumbrae \\citep{keil1999}, what we actually observe is the direct interaction between the umbra (with a strong vertical magnetic field that inhibits convection inside it) and the convective plasma in the surrounding photosphere, without any intermediate structure in between. Many observed features such as bright granules moving in the border of a pore \\citep{sobotka1999} show the complex exchange taking place between the pore and its surrounding granulation.\\\\ Our main interest is the characterization of the horizontal flows around a variety of solar pores on the basis of high-resolution time series of images. Observations from ground-based and space telescopes are analyzed by means of the local correlation tracking technique. In Sect.~2, the paper concentrates on the description of the images acquisition and data processing separately for ground-based and satellite data. The analysis of the data and the presentation of results are treated in Sect.~3. A general summary and final discussion are presented in Sect.~4.\\\\ ", "conclusions": "\\label{sec:dis} The proper motions in solar active regions displaying pores are analyzed from high-resolution time series of images. The observing material stems from coordinated ground-based and space observations. Thus, part of this material was acquired with the Swedish 1-m Solar Telescope, and reconstructed by employing the novel MFBD and MOMFBD techniques to achieve image resolutions near the diffraction limit. The other part of the data stems from the solar telescope on board the \\emph{Hinode} satellite. The long duration, stability and high-resolution of the time series achieved by \\emph{Hinode} enable us to study dynamical properties of the photospheric horizontal flows along periods of time much longer than those typically reachable from ground-based observations which are restricted by varying seeing conditions.\\\\ The local correlation technique applied to the time series allowed us to track the proper motions of structures in solar active regions and particularly in the areas nearby solar pores. Proper motions have been tracked in a variety of active regions for periods of typically 20-60 min but also one for several hours. We conclude that the flow patterns derived from different observational sets are consistent among each other in the sense that they show the determinant and overall influence of exploding events in the granulation around the pores and in the whole FOV. Motions toward the pores in their nearest vicinity are the dominant characteristic we claim to observe systematically. Thus, we do not find any trace of moat flow in the wide sample of pores studied. The motions at the periphery of the pores are basically influenced by the external plasma flows deposited by the exploding events, as suggested by other authors in previous works \\citep{sobotka1999, roudier2002, sankara2003}. In addition, the horizontal velocity magnitudes are clearly lower ($<$ 0.3 km s$^{-1}$) in the nearest locations surrounding the pores and, in general, in the more magnetized regions in the FOV, as expected due to the inhibition of convection taking place.\\\\ Our results are also in agreement with recently developed 3D radiative magnetohydrodynamic simulations of pore-like magnetic structures that report downflows surrounding them, maintained by horizontal flows towards the simulated pore \\citep{cameron2007}. Moreover, we interpret the dividing line between radial inward and outward motions, found by \\cite{deng2007} outside the residual pore in the last stage of a decaying sunspot, as corresponding to the location of the centres of divergence of the exploding events around the pore. The outward motions these authors describe, which are not in the immediate surroundings of the pore but separated by the annular inward motion, would then correspond not to moat flows but to the outward flows originated in the regular mesh of divergence centers around the pore. \\\\" }, "1003/1003.1717_arXiv.txt": { "abstract": "We present a detailed analysis of high resolution near-infrared imaging and spectroscopy of the potential star cluster IRS13E very close to the massive black hole in the Galactic Center. We detect 19 objects in IRS13E from Ks-band images, 15 of which are also detected reliably in H-band. We derive consistent proper motions for these objects from the two bands. Most objects share a similar westward proper motion. We characterize the objects using spectroscopy (1.45 to $2.45\\,\\mu$m) and (narrow-band) imaging from H- ($1.66\\,\\mu$m) to L'-band ($3.80\\,\\mu$m). Nine of the objects detected in both Ks- and H-band are very red, and we find that they are all consistent with being warm dust clumps. The dust emission may be caused by the colliding winds of the two Wolf-Rayet stars in the cluster. Three of the six detected stars do not share the motion or spectral properties of the three bright stars. This leaves only the three bright, early-type stars as potential cluster members. It is unlikely that these stars are a chance configuration. Assuming the presence of an IMBH, a mass of about $14000\\,M_{\\odot}$ follows from the velocities and positions of these three stars. However, our acceleration limits make such an IMBH nearly as unlikely as a chance occurrence of such a star association. Furthermore, there is no variable X-ray source in IRS13E despite the high density of dust and gas. Therefore, we conclude that is unlikely that IRS13E hosts a black hole massive enough to bind the three stars. ", "introduction": "The innermost parsec of the Galaxy hosts the supermassive black hole (SMBH) Sgr~A* \\citep {Schoedel_02, Ghez_03}, accompanied by a population of young WR/O-stars \\citep {Forrest_87, Allen_90, Krabbe_91,Genzel_96}. Most of these stars reside in one or two disk-like structures \\citep{Genzel_03, Paumard_06, Lu_09, Bartko_09}. A group of at least three such bright stars called IRS13E at a distance of $3.5'' =0.13\\,$pc from Sgr~A* is of special interest. This group has a diameter of about $0.5''$. \\citet {Maillard_04} identified two of the stars, E2 and E4, from their emission lines as early-type stars. In addition they identified four other early-type stars in IRS13E from broad-band SED fitting. The four brightest stars share a common proper motion \\citep {Ott_02}. \\citet {Maillard_04} concluded without using a statistical test that such an association of young stars cannot be a coincidence. From the radial velocities of two of the stars they estimated a cluster mass of at least $750\\,M_{\\odot}$ if the two stars are bound. This mass is higher than the stellar mass seen in IRS13E. They explained the additional mass by the presence of an intermediate mass black hole (IMBH). In the simulations of stellar clusters of \\citet{Zwart_02} the core of a dense cluster collapses and forms an IMBH. Such an IMBH would be necessary, if IRS13E were an inspiraling cluster that survives the infall into the Galactic Center (GC) and reaches the central parsec before disintegrating \\citep{Hansen_03}. A cluster without central mass would be disrupted \\citep{Gerhard_01} by the tidal forces of the SMBH. \\citet{Schoedel_05} measured the proper motions of the four brightest sources more accurately and estimated that the cluster has a mass of about $10000$ to $ 50000\\,M_{\\odot}$ if it is gravitationally bound. According to these authors, an IMBH of this mass is unlikely, mainly because of the lack of radio and X-ray emission in IRS13E, despite the presence of a lot of dust. They suggested that IRS13E could also be a cluster in the process of dissolution or a chance association. \\citet{Paumard_06} identified the spectral types of the three brightest early-type stars. In addition, they measured the stellar surface density around IRS13E on a deconvolved H-band image. Inside a radius of $0.30''$, the core of IRS13E, they found at least twelve stars. Furthermore, they found an over-density out to $0.68''$ which has a total significance of $4.5\\,\\sigma$. They concluded that IRS13E is a cluster. Given that velocities had been measured for only four of the stars, \\citet{Paumard_06} argued that the total velocity dispersion could be small. Then no dark mass is needed to explain IRS13E. \\citet{Trippe_08} measured the velocities of the stars between $0.30''$ and $0.68''$. The velocities of most of these stars are different from the stars in the center. Therefore most of them are not cluster members and thus the overdensity is less significant compared to \\citet{Paumard_06}. IRS13E2 and IRS13E4 are candidate members of the face-on counterclockwise disk \\citep{Paumard_06,Bartko_09}. In this case the physical distance of these stars to Sgr~A* is identical with the projected one. According to \\citet{Paumard_04, Paumard_06} IRS13E is embedded in the bar of the minispiral at a distance of z$=7-20$'' of Sgr~A* \\citep{Liszt_03}. In this case IRS13E is rather far away from Sgr~A* and IRS13E2 and IRS13E4 would not be candidates members of the counterclockwise disc \\citep{Bartko_09}. Accordingly, the nature of IRS13E is still a matter of debate. This is largely because the proper motions and the nature of the member objects are known only for very few sources in the core with a radius of about $0.3''$. In particular, it is not clear whether all objects identified by previous works actually are stars. In this paper we analyze seven years of imaging data (presented in section 2). We identify 19 objects in the core of IRS13E, for all of which we can measure proper motions (section 3). For the brightest sources E1 and E2, we obtain acceleration limits which start to to constrain the nature of IRS13E (section 3). From H- to L'-band photometry and H+K-band spectra we constrain the nature of the objects (section 4). In Section 5 we calculate the probability that our acceleration limits occur for the two cases that a) IRS13E is bound by an IMBH and b) that it is a chance association. Section 6 considers other data in a more qualitative fashion and theoretical models. Finally we summarize and conclude in Section 7. We assume a distance to the GC of $R_0=8\\,$kpc \\citep{Reid_93} and a mass of the SMBH of $M=4\\times 10^6 M_{\\odot}$ \\citep{Ghez_08,Gillessen_09}. ", "conclusions": "We present a detailed analysis of the potential cluster IRS13E in the GC. We use AO-based images in H- and Ks-band from 2002 to 2009 for identifying objects in IRS13E and for performing astrometry. We detect 19 objects on most high Strehl ratio images in Ks-band and 15 of them also in H-band. The proper motions of the objects are well determined. In addition, we characterize the SED of the objects in IRS13E using AO-based integral field spectroscopy (1.45 to $2.45\\,\\mu$m) for the brighter objects and (narrow-band) imaging from H- ($1.66\\,\\mu$m) to L'-band ($3.80\\,\\mu$m) for the fainter objects. We fit the SED of the objects by a model consisting of two blackbodies and an extinction parameter. From this we conclude that 13 objects are dust clumps without embedded stars and that the extinction appears to be constant in the IRS13E region. Three objects are fainter stars: One is a spectroscopically identified late-type star. The two fainter stars of the magnitude of the red clump are consistent with the expected number of stars per area at this magnitude bin. Thus, these three stars are likely background stars. Therefore, IRS13E is only a concentration of the remaining three early-type stars E1, E2 and E4.0. They have the following properties: \\begin{itemize} \\item E2 and E4.0 are probable located close in 3D, because of the bright hot dust clumps between them. \\item A chance association of three early-type stars has a probability of around 0.2\\%. Depending on the assumptions on dynamical properties of the stellar population this value varies between 0.06\\% and 0.6\\%. \\item In the case of a binding mass, we use Monte Carlo Simulations of a cluster of three stars around all reasonable masses of IMBHs (300 to $3\\times 10^5 M_{\\odot}$). From those simulations we select the ones that agree with the constraints from the observed positions and velocities of the stars. The most likely mass is about $20000\\,M_{\\odot}$. This mass exceeds the stellar mass (inclusive faint, invisible stars) in IRS13E. \\item We find that at any given mass less than 0.8\\% of the previously selected simulations agree with the observed acceleration limits of the three stars and the fact that Sgr~A* appears to be at rest. This value is uncertain by a factor of $\\approx 2$. \\item There are roughly 15 times too few young stars in the GC for the formation of an IMBH as massive as $20000\\,M_\\odot$. \\item The weak and non-variable X-ray source in IRS13E makes the presence of an IMBH unlikely. \\end{itemize} Overall, we conclude that it is more likely that IRS13E does not host an IMBH." }, "1003/1003.3238_arXiv.txt": { "abstract": "Comparing clustering of differently biased tracers of the dark matter distribution offers the opportunity to reduce the sample or cosmic variance error in the measurement of certain cosmological parameters. We develop a formalism that includes bias non-linearities and stochasticity. Our formalism is general enough that can be used to optimise survey design and tracers selection and optimally split (or combine) tracers to minimise the error on the cosmologically interesting quantities. Our approach generalises the one presented by \\cite{McDSel} of circumventing sample variance in the measurement of $f\\equiv d \\ln D/d\\ln a$. We analyse how the bias, the noise, the non-linearity and stochasticity affect the measurements of $Df$ and explore in which signal-to-noise regime it is significantly advantageous to split a galaxy sample in two differently-biased tracers. We use N-body simulations to find realistic values for the parameters describing the bias properties of dark matter haloes of different masses and their number density. We find that, even if dark matter haloes could be used as tracers and selected in an idealised way, for realistic haloes, the sample variance limit can be reduced only by up to a factor $\\sigma_{2tr}/\\sigma_{1tr}\\simeq 0.6$. This would still correspond to the gain from a three times larger survey volume if the two tracers were not to be split. Before any practical application one should bear in mind that these findings apply to dark matter haloes as tracers, while realistic surveys would select galaxies: the galaxy-host halo relation is likely to introduce extra stochasticity, which may reduce the gain further. { ", "introduction": "One of the active topics of current research is the formation and growth of large-scale structure in the universe. Knowledge of the physical origin of the growth of structure will allow us to know about the origin of dark matter and also provide a useful way to discriminate between different theories for the origin and evolution of dark energy. In particular, comparing and combining measurements of the Universe expansion history (as given by e.g., Baryon Acoustic Oscillations, cosmic chronometers, Supernovae) with measurements of the linear growth of structure, can provide a tool to test whether dark energy is an extra component with negative pressure or a manifestation of the breakdown of general relativity on large scales. To this end, usually, the goal is to measure the $f(z)$ parameter defined as $f(z)\\equiv d\\ln D(z)/d\\ln a(z)$, where $D(z)$ is the linear growth factor and $a(z)$ the scale factor. The two main approaches to measure the growth of structure are weak gravitational lensing (e.g., \\cite{cite6,Bacon05}) and galaxy clustering, which is the technique we consider here. Galaxy clustering is a relatively simple, high signal-to-noise measurement: the angular position of galaxies can be measured using photometry, and the radial position using spectroscopy. With this information the three-dimensional power spectrum of galaxies can be computed as a function of redshift. However, at the two-point level, the galaxy field can only trace the dark matter field up to a bias factor $b$, which may depend on scale and redshift. Thus, only once the bias is known, the galaxy power spectrum $P_{gg}$ can be related to that of the dark matter $P_{mm}$ and thus yield the growth of structure. The main drawback associated to this technique is that the value of this bias cannot readily be predicted from theoretical models of galaxy formation. There are several observational techniques to measure the bias parameter \\citep{Fry94,Feldmanbias,Verde2df02, WLbias, cite7a,cite7b}. In this work we use the approach that takes advantage of the redshift-space distortions. Peculiar velocities of dark matter tracers are set by the gravitational field: using the measured redshift as a distance indicator distorts clustering, enhancing it along the line-of-sight, and the redshift space distortion parameter is $\\beta\\equiv f/b$ \\citep{DavisPeebles83, kaiser, hamilton}. Using measurements in different directions (different Fourier modes), one can compute $\\beta$. In combination with the galaxy power spectrum measurement, this approach yields the divergence power spectrum $P_{\\theta \\theta}=f^2 P_{mm}$, which can be directly compared with theory predictions and encloses the desired dependence on the growth of structure. There are two sources of errors in the measurement of the galaxy power spectrum: the shot noise and the sample variance (or cosmic variance). On one hand, the shot noise is due to the fact that we use a discrete set of objects to characterise the matter field. If this noise is Poisson, it is scale independent and equal to the reciprocal of the number density of objects \\citep{cite8}. On the other hand, the sample variance effect is due to the fact that the matter field has its origin in a random realisation of the underlying cosmology. In a finite survey volume there are only a finite number of modes present, especially on large scales we only have a few modes to perform the averaging. Thus, the total error on the power spectrum $P$ at a given scale $k$, is $\\sigma_P/P=\\left(2/N\\right)^{1/2}\\left(1+\\sigma_n/P\\right)$. Here, $N$ is the number of modes measured (at the scale given by $k$) and $\\sigma_n$ is the shot noise contribution. We see that just reducing the shot noise (increasing the number density of objects) does not help, as there is a natural limitation on our capacity to measure $P$ (and consequently $f(z)$), due to sample variance. In order to reduce this limitation, a multi-tracer technique has been advocated recently \\citep{Seljakfnl}. It is based on the usage of two differently biased tracers of the dark matter field. With this method the sample variance limit can be reduced. The effectiveness of this method depends on a number of factors: the ratio of these different biases; the signal-to-noise regime and on the non-linearity of the biases. With the exception of gravitational lensing, one cannot see the dark matter directly nor the dark matter haloes, so in most practical applications, tracers need to be used such as galaxies, quasars, clusters, or 21 cm emission. The goal of this paper is to study the possibility of measuring the parameter $x(z)\\equiv f^2(z)D^2(z)$ using the single- and the multi-tracer formalism and see whether the reduction of the sample variance is significant. Here, we present a new formalism of how to estimate the error on $x$ using the multi-tracer formalism, taking into account that the bias may be scale dependent, non-linear and stochastic. This formalism may be useful for galaxy surveys, because it has been observed that the galaxy biasing is significantly non-linear and stochastic. N-body simulations and theoretical models, allow us to estimate which are the bias characteristics for dark matter tracers, and therefore which precision can be reached with this model. In \\S 2 we begin by introducing the formalism of our method and analyse how the different parameters affect the reduction of the sample variance effect. In \\S 3 we use both analytical approximations and N-body simulations to obtain physically motivated parameters for our model and compute realistic expected errors using the single- and multi-tracer formalism. In \\S 4 we conclude with a summary and a discussion of the results. ", "conclusions": "\\label{sec:discussion} We have revisited the method of circumventing sample variance in the measurement of $f=d\\ln D /d\\ln a$ ($D$ being the linear growth factor), based on comparing the clustering properties of two differently-biased tracers of the dark matter distribution. This method was recently investigated by \\citet{McDSel}, although a similar technique in a different context was presented in \\citet{pen04}. Along the same lines, \\citet{SlosarCV} and \\citet{Seljakfnl} propose to compare clustering of differently biased tracers to circumvent sample variance in the measurement of primordial non-Gaussianity. Most of the statistical power of these measurements comes from very large scales, where cosmic variance is the dominant contribution to the statistical error-bars. By suppressing cosmic variance, this approach promises to reduce drastically error-bars on cosmologically very important quantities; for example it would allow for a high-precision determination of growth of structure as a function of redshift, as encoded in $f D$, and an improvement of dark energy figures of merit by large factors. All these approaches assume that the observed objects (i.e., galaxies) trace the dark matter deterministically; the galaxy density field is assumed to be proportional to the dark matter field with the constant of proportionality given by a single parameter, the bias. This goes under the name of the linear bias model. An important underlying assumption is therefore made, that there is no stochasticity between the tracer field and the dark matter field on the scales of interest, which is expected to breakdown at some level, at least on small scales. While the linear bias model has been extremely successful in cosmology (e.g., \\cite{Reidetal09} and references therein), it is well known that the linear bias model might provide a good description for the galaxy power spectrum even if the relation between the galaxy and dark matter overdensities is not that of a linear bias (e.g., \\cite{heavensetal98}). Galaxies are believed to form inside dark matter haloes, but their formation probability as function of halo mass and their exact radial distribution are still the subject of active research. The process of halo occupation by galaxies is expected to be stochastic to some extent, but the details of the galaxy distribution within haloes is expected to become increasingly unimportant on large scales. Here we simplify the issue by assuming (possibly with an over-simplification) that dark matter haloes can be used as tracers. The linear deterministic bias model however is known not to be a perfect description of halo clustering and that the relation between dark matter and haloes and between haloes of different masses is stochastic. For example, \\citet{SW} point out that ``the fluctuations between haloes and the initial or final matter fields are never below 10-20 per cent\" and that ``the scatter between the fields in individual modes is significant and one cannot assume that the fields are simply proportional to one another\". This was further explored and quantified by \\citet{Bonoli}. Note that the halo overdensity field is expected to have a stochastic component even if it was a perfect Poisson sampling of a linearly biased dark matter field, but the above references and N-body simulations show that there are additional sources of stochasticity beyond shot noise. We have thus set out to generalise the approach of \\citet{McDSel}, by assuming that the bias of haloes may not be perfectly linear and allowing for some stochasticity. We have computed the expected error on the quantity $fD$ achievable by comparing clustering of differently biased tracers (thus suppressing cosmic variance) and by combining the different tracers in a single sample (thus reducing shot noise and stochasticity but carrying along sample variance in full). We have analysed how the bias, the noise, the non-linearity and stochasticity affect the measurements of $fD$ and explored in which signal-to-noise regime it is significantly advantageous to split a galaxy sample in two differently-biased tracers. We used results from simulations to set plausible values for these parameters to see how great the gains may be in practice. We find that even small amount of stochasticity (either in the form of Poisson noise or in more general form) and of non-linearity can limit significantly the performance of the two-tracers approach. In our analysis we also have assumed a scale-independent bias. This may be enough, for the mass range studied, if we consider only dark matter haloes as tracers. This is indeed what we have seen in simulations. On the other hand, it is also true that more realistic approaches, which account for galaxies as tracers instead of haloes, should include a scale-dependent bias. However, including this in our formalism can only reduce the gain achievable by splitting the sample. We expect the ratio of errors increases from the current value of 0.9 to even closer values to 1 if the bias is strongly scale dependent. We have shown that only in the very high signal-to-noise regime it is significantly advantageous to split the sample and that, even though the gain is maximised by increasing the {\\it ratio} of the biases of the two tracers, both tracers should be well sampled. We have explored different ways of selecting and splitting dark matter haloes obeying a $\\Lambda CDM$ mass function and found that one can achieve up to a 40\\% reduction of the error on $f D$. While this would correspond to the gain from a three times larger survey volume if the two tracers were not to be split, it is much smaller that the improvement forecasted in the absence of stochasticity and bias non-linearity. In addition we should note that these findings apply to dark matter haloes as tracers, while realistic surveys would select galaxies: the galaxy-host halo relation is likely to introduce extra stochasticity which would reduce the gain further. The formalism we have developed, however, is general enough that can be used to optimise survey design and tracers selection and optimally split (or combine) tracers to minimise the error on the cosmologically interesting quantities." }, "1003/1003.1521_arXiv.txt": { "abstract": "We present measurements of the Type Ia supernova (SN) rate in galaxy clusters based on data from the Sloan Digital Sky Survey-II (SDSS-II) Supernova Survey. The cluster SN Ia rate is determined from 9 SN events in a set of 71 C4 clusters at $z \\le 0.17$ and 27 SN events in 492 maxBCG clusters at $0.1 \\le z \\le 0.3$. We find values for the cluster SN Ia rate of $({\\cfsnur}^{+\\cfsnurhi+\\cfsnurhisyst}_{-\\cfsnurlo-\\cfsnurlosyst})$ $\\mathrm{SNu}r$ $h^{2}$ and $({\\bcgsnur}^{+\\bcgsnurhi+\\bcgsnurhisyst}_{-\\bcgsnurlo-\\bcgsnurlosyst})$ $\\mathrm{SNu}r$ $h^{2}$ ($\\mathrm{SNu}x = 10^{-12} L_{x\\sun}^{-1} \\mathrm{yr}^{-1}$) in C4 and maxBCG clusters, respectively, where the quoted errors are statistical and systematic, respectively. The SN rate for early-type galaxies is found to be $({\\cfsnurearly}^{+\\cfsnurhiearly+\\cfsnurhisystearly}_{-\\cfsnurloearly-\\cfsnurlosystearly})$ $\\mathrm{SNu}r$ $h^{2}$ and $({\\bcgsnurearly}^{+\\bcgsnurhiearly+\\bcgsnurhisystearly}_{-\\bcgsnurloearly-\\bcgsnurlosystearly})$ $\\mathrm{SNu}r$ $h^{2}$ in C4 and maxBCG clusters, respectively. The SN rate for the brightest cluster galaxies (BCG) is found to be $({\\cfsnurbcg}^{+\\cfsnurhibcg+\\cfsnurhisystbcg}_{-\\cfsnurlobcg-\\cfsnurlosystbcg})$ $\\mathrm{SNu}r$ $h^{2}$ and $({\\bcgsnurbcg}^{+\\bcgsnurhibcg+\\bcgsnurhisystbcg}_{-\\bcgsnurlobcg-\\bcgsnurlosystbcg})$ $\\mathrm{SNu}r$ $h^{2}$ in C4 and maxBCG clusters, respectively. The ratio of the SN Ia rate in cluster early-type galaxies to that of the SN Ia rate in field early-type galaxies is ${\\cfrat}^{+\\cfrathi+\\cfratsysthi}_{-\\cfratlo-\\cfratsystlo}$ and ${\\bcgrat}^{+\\bcgrathi+\\bcgratsysthi}_{-\\bcgratlo-\\bcgratsystlo}$, for C4 and maxBCG clusters, respectively. The SN rate in galaxy clusters as a function of redshift, which probes the late time SN Ia delay distribution, shows only weak dependence on redshift. Combining our current measurements with previous measurements, we fit the cluster SN Ia rate data to a linear function of redshift, and find $r_{L} = $ $[(0.49^{+0.15}_{-0.14}) +$ $(0.91^{+0.85}_{-0.81}) \\times z]$ $\\mathrm{SNu}B$ $h^{2}$. A comparison of the radial distribution of SNe in cluster to field early-type galaxies shows possible evidence for an enhancement of the SN rate in the cores of cluster early-type galaxies. With an observation of at most 3 hostless, intra-cluster SNe Ia, we estimate the fraction of cluster SNe that are hostless to be $(9.4^{+8.3}_{-5.1})\\%$. ", "introduction": " ", "conclusions": "" }, "1003/1003.5843_arXiv.txt": { "abstract": "To understand the formation of a magnetically dominated molecular cloud from an atomic cloud, we study the interaction of a weak, radiative shock with a magnetised cloud. The thermally stable warm atomic cloud is initially in static equilibrium with the surrounding hot ionised gas. A shock propagating through the hot medium then interacts with the cloud. We follow the dynamical evolution of the shocked cloud with a time-dependent ideal magnetohydrodynamic code. By performing the simulations in 3D, we investigate the effect of different magnetic field orientations including parallel, perpendicular and oblique to the shock normal. We find that the angle between the shock normal and the magnetic field must be small to produce clouds with properties similar to observed molecular clouds. ", "introduction": "Molecular clouds exhibit a hierarchical density structure \\citep[e.g.][]{BS86}. Stars form in the densest regions, or dense cores, which become gravitationally unstable. In fact, molecular clouds in the Solar neighbourhood that don't harbour any stars are rare. While most of the stars within molecular clouds are young ($\\approx $ 1-2 Myr), stellar associations older than 5 Myr are devoid of molecular gas \\citep[e.g.][]{BH07}. This suggests that molecular clouds are short-lived, transient objects and that the time-lag between cloud formation and stellar birth is short. The rapid onset of star formation requires that large density contrasts arise while the parental cloud forms. It has been suggested that this fragmentation results from thermal processes in the interstellar medium (ISM) \\citep[see e.g.][and references therein]{KKH09}. For a range of pressures and the heating and cooling rates appropriate for diffuse atomic gas, two thermally stable phases exist, i.e. a rarefied, warm phase and a cold, dense phase, which can co-exist in pressure equilibrium \\citep{FGH69,Wetal95}. Atomic gas at intermediate temperatures is subject to a thermal instability. A sufficient rise in pressure causes the cold phase to be the only stable one \\citep{F65}. Previous studies of molecular cloud formation due to thermal instability mainly focus on collisions of warm gas streams in the context of expanding supernovae shells or spiral arm shocks \\citep[e.g.][]{Hetal08, HSH09, II09}. Beside being thermally unstable in some circumstances, the collision region is prone to numerous dynamical instabilities such as the Kelvin-Helmholtz, Rayleigh-Taylor and nonlinear thin shell or Vishniac instabilities. In the turbulent shocked layer cold gas clumps then arise on short timescales. Although the derived density and velocity structures depend strongly on the magnitude and orientation of the magnetic field, they resemble the ones observed in molecular clouds and diffuse HI clouds \\citep{HSH09}. While this model of flow-driven structure generation shows rapid onset of star formation while the parental cloud is forming, it cannot account for the observed low star formation rates \\citep{KKH09}. As there is a continuous instream of gas into the collision region, too much of the accumulated gas will be converted into stars. For the same reason, it also cannot explain cloud lifetimes. Some of these limitations do not necessarily occur when studying the same processes in cloud-cloud interactions \\citep{KW98,Metal99} or shock-triggered models where shocks overrun warm, diffuse density perturbations \\citep[][hereafter Paper I]{IK04,PaperI}. Indeed, numerical models of shocks interacting with clouds show that the clouds fragment due to dynamical instabilities \\citep{Metal94}, thereby setting a limit on both the cloud lifetime and the star formation rate. Colliding flow-driven models only form thin filamentary clouds, while different cloud morphologies such as a cometary cloud structure with a massive head and long-spread tail, are also observed \\citep{Tetal02}. Such a morphology is characteristic for clouds harbouring cluster-forming cores. The W3 Giant Molecular Cloud (GMC) is an example of such a cloud \\citep{Metal07}. As there is no constraint on the geometry of density perturbations in the shock-triggered models several morphologies can be reproduced. The interaction of a shock with a magnetised cloud has been studied extensively in a two-dimensional (2D) axisymmetric geometry, both for adiabatic \\citep[e.g.][]{Metal94,Netal06} and radiative shocks \\citep[e.g.][Paper I]{Fetal05}. However, these simulations are limited to a single configuration of the magnetic field, i.e. the magnetic field is parallel to the symmetry axis and the shock normal. To study the effect of the magnetic field orientation, it is necessary to model the shock-cloud interaction in 3D. Although such simulations have been around for some time \\citep{SN92,Getal00}, only recently high enough resolution has been achieved for general cloud properties to converge \\citep{SSS08}. However, the Shin et al. study focuses on strong, adiabatic shocks, while the results of Paper I show that molecular clouds most likely form from interactions involving weak, radiative shocks. The first 3D simulations of the interaction of a radiative shock with a magnetised cloud were performed by \\citet{Letal09}, although they use a nearly isothermal equation of state to simulate strong radiative cooling. In this paper we study the interaction of a weak, radiative shock with a magnetised cloud. We focus on the early stages of the evolution before the cloud re-expands and fragments. In Sect.~\\ref{sect:model} we describe the numerical code and the initial conditions. The results are presented in Sect.~\\ref{sect:results}, while we discuss the cloud properties in \\ref{sect:properties}. Finally, we finish the paper with a summary (Sect.~\\ref{sect:conclusions}). ", "conclusions": "\\label{sect:properties} \\subsection{Size and density} In order to study the properties of the different shock-cloud interaction models, we use diagnostics similar to those used by previous authors, \\citep[e.g.][]{Metal94,SSS08}. Specifically, we use the density-weighted average of variables such as the plasma $\\beta$ and the density using Eq.~\\ref{eq:average}. Furthermore, we define \\[ a = \\left[5 \\left( - ^2 \\right)\\right]^{1/2}. \\] along the $x$-axis, with similar expressions for analogous quantities defined along the $y$ and $z$-axes. This gives us the axes for an ellipsoid with a similar mass distribution as the cloud's which are used to follow geometrical changes of the cloud. Figure~\\ref{fig:volume} shows the temporal evolution of the volume of the cloud. The parallel shock compresses the cloud the most, while the perpendicular shock is the least compressive. At $t_{cc}$ the volume of the cloud is only a tenth of its original volume. The ram pressure exerted by the post-shock flow of the perpendicular shock is smaller than that for the parallel shock, as both the post-shock density and gas velocity (in the lab frame) are lower. (A parallel shock compresses gas more than a perpendicular shock with the same Mach number.) Furthermore, work needs to be done to compress the transverse magnetic field. Similarly, it is not surprising to find that the 45$^\\circ$ oblique shock squeezes the cloud faster than the perpendicular shock, but slower than the parallel shock. For smaller angles we expect the confinement of the cloud to converge to the parallel shock model. Indeed, the volume of the cloud in the 15$^\\circ$ oblique shock model changes as for the parallel shock model with only a small deviation at later times. \\begin{figure} \\includegraphics[width = 8.4 cm]{fig7} \\caption{Mean cloud volume derived from the mass-weighted moments along the different axes for the parallel (solid), perpendicular (dotted) and oblique shock (thick dashed for 45$^o$ and dashed for 15$^o$) cases. The volume is normalised to the initial cloud size.} \\label{fig:volume} \\end{figure} Figure~\\ref{fig:volume} seems to show a minimum in the volume near 9~Myr. From the 2D axisymmetric simulations, we know that a minimum arises as the fast-mode shock is reflected at the centre of the cloud at around $t_{cc} (\\approx$ 8.3 Myr) and propagates back to the boundary. This shock reflection initiates a re-expansion phase in the direction perpendicular to the post-shock flow direction \\citep[see][Paper~I]{Metal94} Along the flow direction, the compression of the cloud continues and the cloud ends up as a thin disc after 1.5-2 $t_{cc}$. This means that the geometrical shape of the cloud resembles more closely an oblate spheroid than a sphere. The ellipticity of the cloud changes from 0 initially (i.e. a sphere) to roughly 0.5 at $t_{cc}$. As mentioned earlier, we do not follow this expansion phase and stop the simulation shortly after the cloud-crushing time $t_{cc}$. \\begin{figure} \\includegraphics[width = 8.4 cm]{fig8} \\caption{rms number density of the cloud for the parallel (solid), perpendicular (dotted) and oblique shock (thick dashed for 45$^o$ and thin dashed for 15$^o$) cases.} \\label{fig:rhoav} \\end{figure} Figure~\\ref{fig:rhoav} gives the evolution of the root-mean-square (rms) number density. Although the volumes of the cloud are only different by a factor of 2 as can be seen from Fig.~\\ref{fig:volume}, the rms density changes more significantly. At $t_{cc}$, the rms number density for the parallel model is $\\approx 40$~cm$^{-3}$, while it is only 10~cm$^{-3}$ for the perpendicular model. This large variation is due to the large difference in the density of the boundary layer between the models (see Sect.~\\ref{sect:results}). Only models with small angles between the shock normal and the magnetic field have mean densities similar to those of observed GMCs. Large angle models only reach mean densities similar to diffuse HI clouds. Such a dependence of the density structure on the magnetic field orientation is also observed by \\citet{HSH09}. \\begin{figure} \\includegraphics[width = 8.4 cm]{fig9} \\caption{Mass-weighted plasma beta of the cloud for the parallel (solid), perpendicular (dotted) and oblique shock model (thick dashed for 45$^o$ and thin dashed for 15$^o$). } \\label{fig:betaav} \\end{figure} \\begin{figure} \\includegraphics[width = 8.4 cm]{fig10} \\caption{Mass fraction of the cloud with $\\beta < 0.1$ for the parallel (solid), perpendicular (dotted) and oblique shock model (thick dashed for 45$^o$ and thin dashed for 15$^o$).} \\label{fig:massbelow} \\end{figure} \\subsection{Thermal and magnetic pressure} Observations show that molecular clouds are magnetically dominated with plasma $\\beta$ of the order 0.04 - 0.6 \\citep{CHT03}. Figure~\\ref{fig:betaav} shows the mass-weighed mean of $\\beta$. This weighted mean value of $\\beta$ for the parallel shock model does not lie within the observed range. During the early stages of the evolution the gas behind the slow-mode shock has a large thermal pressure, while the magnetic pressure is small. Hence, the plasma $\\beta$ is large for the high-density gas in the boundary layer (which dominates the mean value of $\\beta$), i.e. $\\beta \\approx 5$. After about 2.5~Myr, the weighted mean value of $\\beta$ decreases as the thermally unstable gas behind the fast-mode shock becomes magnetically dominated. Although the weighted mean value of $\\beta$ is now around unity, Fig.~\\ref{fig:massbelow} shows that a significant mass fraction of the cloud has $\\beta < 0.1$. At $t_{cc}$, about 75\\% of the gas is significantly magnetically dominated. The other models do produce a cloud with a weighted mean $\\beta$ of the order of 0.4 (see Fig.~\\ref{fig:betaav}). The main reason for this has to do with the role played by the transverse component of the magnetic field. While the thermal pressures behind the fast-mode and slow-mode shock are not as high as in the parallel shock model (i.e. the transition from the warm phase to the cold phase is smoother), the magnetic field is significantly compressed behind the fast-mode shock. The combined result of these effects produces a lower $\\beta$ in both the boundary layer and within the cloud. Figure~\\ref{fig:massbelow} indeed shows that magnetically-dominated gas appears much earlier for shock models with a transverse component to the magnetic field. After only 3~Myr, already 10\\% of the total mass in these models is magnetically dominated. Our simulations show that, to produce clouds with magnetically-dominated high-density gas, the angle between the shock normal and the magnetic field must be small. While perpendicular shocks produce magnetically-dominated gas at low densities, high-density clumps with $\\beta \\gg 1$ arise in the parallel shock models. \\subsection{Fragmentation} Our models show that large fractions of the cloud are magnetically dominated. This provides the ideal conditions for MHD waves to generate high-density clumps and cores within the cloud \\citep{FH02,VFH06,VFH08}. For shock models with a transverse component of the magnetic field, this process initiates earlier suggesting a higher degree of fragmentation. On top of that, another process is effective in these models. As the transition from the thermally-stable warm phase to the cold one follows the unstable part of the equilibrium curve, small perturbations can initiate the formation of dense, cold clumps embedded in warm, diffuse gas \\citep{KI06}. Unfortunately, we cannot follow these clump and core formation processes as our resolution is insufficient. While a low resolution is partly to blame for the low amount of fragmentation, the uniform initial conditions of the cloud also play an important role. In the colliding flow-driven models of \\citet{Hetal08} and \\citet{HSH09}, the generation of cold dense cores and clumps relies on seeded perturbations in either the incoming flow or at the collision front. Without these perturbations, the collision region remains roughly uniform and fragmentation occurs on very long timescales. Therefore, it can be expected that the introduction of perturbations within the cloud and at the edge of the cloud would produce much more fragmentation. Furthermore, our simulations do not include the effect of self-gravity. While self-gravity is dynamically unimportant for the global evolution of the cloud, i.e. the mass of the cloud is much smaller than its Jeans' mass, its effect will become important locally. For example, the high-mass boundary layer clumps of the parallel shock model have masses that exceed their Jeans mass. We will investigate the effect of self-gravity on the cloud evolution in a later paper. \\begin{figure} \\includegraphics[width = 8.4 cm]{fig11} \\caption{Mass fraction of cold (solid) and molecular (dashed) gas within the cloud for the parallel (thick) and perpendicular (thin) shock model. Molecular gas is assumed to be present when the thermal gas pressure is higher than 2500$k$ and the gas temperature lower than 100K.} \\label{fig:molecular} \\end{figure} \\subsection{Molecular clouds} Our simulations only describe the dynamical evolution of a cloud from warm atomic gas to cold atomic gas. Figure~\\ref{fig:molecular} shows the temporal evolution of the cold gas mass fraction for the parallel and perpendicular shock model. Cold gas arises earlier in the parallel shock model than in the perpendicular one, but after 9 Myr more than 50\\% of the initial cloud mass is in the thermally stable cold gas phase. Although we have not included a description of molecular cooling, nor do we follow the cloud chemistry, we can roughly estimate how much gas is converted from the cold atomic gas to molecular gas. \\citet{WBS95} find that the average number density of H$_2$ in the CO clumps of the Rosette Molecular Cloud is $\\approx 220~{\\rm cm^{-3}}$. With typical excitation temperatures between 10 and 20 K, the thermal gas pressure of these clumps is roughly 2500k. Therefore, we assume that any gas parcel in the shocked cloud with a thermal pressure higher than 2500k and a temperature below 100K will become molecular given enough time (see below). Using this criterion, we find that the parallel shock model generates a molecular cloud as half of the cold gas becomes molecular (see Fig.~\\ref{fig:molecular}). In the perpendicular cloud model only a small fraction of the cold atomic gas is converted into molecular gas. This suggests that the perpendicular shock model only produces a diffuse HI cloud. The above result is only valid if the time scale for the formation of molecules is short and if the physical parameters of the formation process are met. \\citet{Getal09} use high-resolution 3D simulations of turbulent interstellar gas to follow the formation and destruction of molecular hydrogen and CO. They find that most CO forms within 2-3~Myr for dense, turbulent gas, while the formation of H$_2$ is even faster, i.e. within 1-2~Myr. Their results indicate that once large enough spatial and column densities are reached, the conversion from atomic to molecular gas is rapid. A good indictor for the formation of molecules is the visual extinction $A_V$, which can be expressed as \\citep[e.g.][]{CC04} \\[ A_V = \\frac{N_{H}}{1.80 \\times 10^{21} {\\rm cm^{-2}}}. \\] For regions with $A_V \\gtrsim 0.5$ and high local densities, we can then expect that molecules are present. Figure~\\ref{fig:Av_pa} shows that the visual extinction is already high early on in the parallel shock model. These regions also correspond to high-density regions, as can be seen in Fig.~\\ref{fig:paboundary}. This model thus most likely produces a molecular cloud. Also, note the similarity of the column density plot for the parallel shock model with the emission map of the W3 GMC (see Fig.~6 of Paper~I). A more structured inner cloud can be expected with a non-uniform initial condition and a higher resolution. While the 15$^\\circ$ oblique shock model also produces high column densities which coincide with high density regions, the perpendicular and the 45$^\\circ$ models do not. Models with large transverse components of the magnetic field produce diffuse HI clouds instead of molecular clouds. However, this conclusion only holds for our current simulations. A higher resolution and inclusion of small-scale perturbations potentially would produce higher density clumps for these models. \\begin{figure*} \\includegraphics[width = 8.4 cm]{fig12a} \\includegraphics[width = 8.4 cm]{fig12b} \\caption{Visual extinction along the $y$-axis for a parallel shock interacting with a diffuse cloud. The left panel show the visual extinction at 2.5 Myr and the right panel at 9 Myr.} \\label{fig:Av_pa} \\end{figure*}" }, "1003/1003.5244_arXiv.txt": { "abstract": "We present the first results of a new Keck spectroscopic survey of UV faint Lyman break galaxies in the redshift range $33$ {\\it Lyman break galaxies} (LBGs), it is now established from various independent surveys that the star formation density, deduced from rest-frame UV luminosities, declines monotonically with redshift (e.g. \\citealt{Stanway03,Bunker04}) largely as a result of a corresponding fading of the characteristic UV luminosity (e.g., \\citealt{Ouchi04a, Yoshida06,Bouwens06a,Bouwens07,McLure10}) The associated stellar mass density in Lyman break galaxies, deduced from near-infrared Spitzer photometry, increases by $\\simeq$1 dex from $z\\simeq 6$ to 4 \\citep{Eyles07,Stark09}. As the rate of change of stellar mass is governed by ongoing star formation, it is useful to relate the two measures and such a comparison indicates a rapid duty-cycle of star formation activity at this time, unlike the more continuous modes seen for equivalent sources at $z\\simeq 2-3$ \\citep{Stark09}. By contrast, the redshift-dependent luminosity function of narrow-band selected {\\it Lyman alpha emitters}, shows no equivalent decline with redshift over $3 < z < 6$ \\citep{Ouchi08}, suggesting an increasing fraction of line emitters amongst the star forming population at early times. Moreover, detailed studies of the slope of the UV continuum in $z>3$ Lyman break galaxies indicates a decreasing dust content at earlier times \\citep[e.g.][]{Stanway05,Bouwens06a,Bouwens09b} as well as a luminosity dependence at $z\\simeq 3$ \\citep[e.g.][]{Reddy09}. Conceivably the combination of a reduced dust content and a shift to more intense, shorter-term star formation at high redshift, can explain these various redshift-dependent trends. Notwithstanding this considerable progress, a major concern is that the above conclusions rest largely on deductions made with {\\it photometric data}, particularly for the Lyman break population. Quite apart from the possibility of low redshift interlopers lying within the photometric samples (a problem that increases for drop-out selected samples at redder wavelengths), the wanted physical measures of the star formation rate, dust content and stellar mass are all rendered uncertain by the absence of precise spectroscopic redshifts. While considerable effort has been invested in the spectroscopic study of $z\\simeq 3$ Lyman break galaxies \\citep[e.g.][]{Shapley03,Steidel03,Quider10}, comparably little spectroscopy has been achieved for higher redshift samples. \\citet{Steidel99} obtained spectroscopic redshifts for nearly 50 bright ($I<25$) $z\\simeq 4$ LBGs. Most surveys of $z\\simeq 5-6$ $V$ and $i'$-drops have generally involved relatively small samples, typically comprising fewer than 10 sources \\citep[e.g.][]{Stanway04,Stanway07,Ando07,Dow-Hygelund07}. Recent deep HST {\\it ACS} Grism observations of the Hubble Ultra Deep Field have allowed the spectra of faint $z\\simeq 5$ LBGs to be characterised \\citep{Rhoads09}, albeit at very low spectral resolution, resulting in 39 redshift confirmations. The largest spectroscopic sample of $47$, it seems unlikely that quasar spectroscopy will constrain the epoch when the bulk of the IGM was reionised in the near future. Ly$\\alpha$ emitting galaxies offer a valuable additional probe of the IGM (e.g., \\citealt{Rhoads01,Malhotra04,Kashikawa06}). In principle, the test is straightforward to apply. Young galaxies emit copious amounts of Ly$\\alpha$ photons, which are resonantly scattered by neutral hydrogen. Hence as we probe the regime when the IGM becomes significantly neutral, the fraction of star-forming galaxies showing strong Ly$\\alpha$ emission should decrease \\citep[e.g.][]{Haiman99,Santos04a,Furlanetto06,McQuinn07,Mesinger08,Iliev08, Dayal10}. Recent measurements of the luminosity function of Ly$\\alpha$ emitters (LAEs) selected via narrowband imaging have revealed a tantalising decline between $z=5.7$ and $z=7.0$ \\citep[e.g.][]{Kashikawa06,Iye06,Ota08}, offering possible evidence that the ionisation state of the IGM evolves over $60.8$. The apparent tension between $\\sigma_8$ from primary cosmic microwave background power and from analytic SZ spectra inferred using ACT and SPT data is lessened with our AGN feedback spectra. ", "introduction": "When CMB photons are Compton-scattered by hot electrons, they gain energy, giving a spectral decrement in thermodynamic temperature below $\\nu \\approx 220$~GHz, and an excess above \\citep{1970Ap&SS...7....3S}. The high electron pressures in the intracluster medium (ICM) result in cluster gas dominating the effect. The integrated signal is proportional to the cluster thermal energy and the differential signal probes the pressure profile. The SZ sky is therefore an effective tool for constraining the internal physics of clusters and cosmic parameters associated with the growth of structure, in particular the {\\it rms} amplitude of the (linear) density power spectrum on cluster-mass scales $\\sigma_8$ \\citep[e.g.,][]{1999PhR...310...97B, 2002ARA&A..40..643C}. Identifying clusters through blind SZ surveys and measuring the SZ power spectrum have been long term goals in CMB research, and are reaching fruition through the South Pole Telescope, SPT \\citep{2009arXiv0912.4317L} and Atacama Cosmology Telescope, ACT \\citep{2010arXiv1001.2934T} experiments. The ability to determine cosmological parameters from these SZ measurements is limited by the systematic uncertainty in theoretical modelling of the underlying cluster physics and hence of the SZ power spectrum. The power contribution due to the kinetic SZ (kSZ) effect that arises from ionized gas motions with respect to the CMB rest frame adds additional uncertainty. There are two main approaches to theoretical computations of the thermal SZ (tSZ) power spectrum: from hydrodynamical simulations of SZ sky maps or from semi-analytical estimates \\citep[][B0205]{2002ASPC..257...15B,2005ApJ...626...12B}. Large cosmological simulations providing a gastrophysical solution to the pressure distribution should include effects of non-virialized motions, accretion shocks, and deviations from spherical symmetry. Averaging over many realizations of synthetic SZ sky projections yields the power spectrum and its variance \\citep[e.g., B0205;] []{2000MNRAS.317...37D,2001ApJ...549..681S,2001PhRvD..63f3001S,2006MNRAS.370.1309S}. In conjunction with primary anisotropy signals and extragalactic source models, the SZ power spectrum has been used as a template with variable amplitude $A_{\\rmn{SZ}}$ for extracting cosmological parameters by the Cosmic Background Imager (CBI) team \\citep[B0205;][]{2009arXiv0901.4540S} and the ACBAR team \\citep{2003ApJ...599..773G,2009ApJ...694.1200R}. $A_{\\rmn{SZ}}$ was used to estimate a $\\sigma_{\\mathrm{8, SZ}}\\propto A_{\\rmn{SZ}}^{1/7}$ as a way to encode tension between the SZ-determined value and the (lower) $\\sigma_8$ obtained from the primary anisotropy signal. The CBI team also has included an analytic model \\citep[][KS]{2002MNRAS.336.1256K} which was also the one adopted by the WMAP team \\citep{2007ApJS..170..377S}. The KS template yielded a lower value for $\\sigma_{\\mathrm{8, SZ}}$ than that obtained with the simulation template, by $\\sim 10\\%$. The KS model assumes a universal ICM pressure profile in hydrostatic equilibrium with a polytropic (constant $\\Gamma$) equation of state. The power spectrum is then obtained using an analytic fit to `halo model' abundances. So far the SPT and ACT have only used the KS template and a related semi-analytic one \\citep{2005ApJ...634..964O, 2009ApJ...700..989B}. This model \\citep[][S10]{2010ApJ...709..920S} allows map generation by painting dark matter halos in N-body simulations with gas. It expands on KS by calculating the gravitational potential from the DM particles, includes an effective infall pressure, adds simplified models for star formation, non-thermal pressure support and energy feedback which are calibrated to observations. Using these templates, the SPT team derived a $\\sigma_{\\mathrm{8, SZ}}$ lower than the primary anisotropy $\\sigma_8$ \\citep[e.g., WMAP7,][]{2010arXiv1001.4635L}. Current simulations with {\\em only} radiative cooling and supernova feedback excessively over-cool cluster centers \\citep[e.g.][]{2000ApJ...536..623L}, leading to too many stars in the core, an unphysical rearrangement of the thermal and hydrodynamic structure, and problems when compared to observations, in particular for the entropy and pressure profiles. The average ICM pressure profile found through X-ray observations of a sample of nearby galaxy clusters \\citep{2009arXiv0910.1234A} is inconsistent with adaptive-mesh cluster simulations \\citep{2007ApJ...668....1N}, as well as the KS analytic model \\citep{2010arXiv1001.4538K}. Pre-heating \\citep[e.g.][]{2001ApJ...555..597B} and AGN feedback \\citep[e.g.][]{2007MNRAS.380..877S, 2008MNRAS.387.1403S, 2008ApJ...687L..53P} help solve the over-cooling problem and improve agreement with observed cluster properties. Previously, an analytical model by \\citet{2005ApJ...634...90R} has explored the effects of effervescent heating on the SZ power spectrum and \\citet{2007MNRAS.382.1697H} use a semi-analytical model to calculate how an entropy floor affects the SZ power spectrum. There have been several simulations on galaxy and group scales that have studied how `quasar' feedback impacts the total SZ decrement \\citep{2006ApJ...653...86T, 2008ApJ...678..674S, 2008MNRAS.389...34B, 2008MNRAS.390..535C}. In this work we explore whether AGN feedback incorporated into hydrodynamical simulations of structure formation can suppress the over-cooling problem and resolve the current inconsistency between theoretical predictions and observations of the SZ power spectrum and X-ray pressure profile. ", "conclusions": "Without hydrodynamical simulations in a cosmological framework similar to the ones presented in this paper it is hard to come up with a consistent model of the gas distribution in clusters and the infall regions which both contribute significantly to the SZ power spectrum. In this paper, we identify three main points that a future semi-analytic model of such a pressure distribution has to provide. (1) In order to arrive at a consistent gas distribution that matches not only the integrated stellar mass fraction but also the X-ray derived pressure profiles within $R_{500}$, we need self-regulating AGN-type feedback. We emphasize that we tuned our parameters to match a previous single-cluster model that successfully suppressed the over-cooling by means of AGN feedback \\citep{2008MNRAS.387.1403S}. The excellent agreement with current data was a pleasant byproduct: our simulated pressure profiles agree with recently obtained observational ones that have been constructed from X-ray data; the scaling relations between the cluster mass and X-ray based Compton-$Y$ \\citep{2009arXiv0910.1234A} also agree; as do the integrated stellar and gas mass fractions \\citep{2007ApJ...666..147G,2007MNRAS.378..293A}. (2) The amount of non-gravitational energy injection into proto-clusters and groups by AGN and starburst galaxies at intermediate-to-high redshifts $z\\gtrsim0.8$ is poorly understood. Other observables are needed to constrain the physics and to answer this question which seems to be essential in understanding the resulting gas profiles. Our simulations suggest that AGN-type feedback lowers the central pressure values as a hydrodynamic response of the gas distribution to the non-gravitational feedback of energy. This effect inhibits gas from falling into the core regions which causes a flatter and more extended pressure profile and a noticeably reduced power of the SZ power spectrum at small angular scales for $\\ell\\gtrsim2000$. (3) For the SZ flux to be converged, an integration of the pressure profile out to $4 R_{200}$ is necessary; half of the SZ flux is contributed from regions outside $R_{200}$. To compute a reliable SZ power spectrum, it is essential to precisely characterize the state of the gas in these infall regions. In particular, we find that: (i) the pressure support from kinetic energy strongly increases as a function of radius to reach on average equipartition with the thermal energy at $\\sim 2 R_{200}$ in our AGN model with the exact dependence on cluster mass to be determined by future work; (ii) the effective adiabatic index $\\Gamma =\\dd\\ln p/\\dd\\ln \\rho \\sim 1.2$ in the interior, but upturns towards $\\Gamma \\sim 5/3$ beyond the virial radius; (iii) the inclusion of cluster asphericity at large radii may also become important. Hence a successful semi-analytic model of the spherical cluster pressure, if that is indeed a viable goal, at the least needs careful calibration using numerical simulations which accurately treat all of the effects. The variance of the average profiles also encodes important information that is manifested in the power spectrum. Our studies also show that simplified analytic models that employ hydrostatic gas models with a constant $\\Gamma$ necessarily overpredict the SZ power on large scales by up to a factor of two and predict an inconsistent shape of the SZ power spectrum. The alternative that we explore in a subsequent paper is to use stacked scaled simulational clusters which are rotated to principal axes to provide the pressure form factors for the semi-analytic approach. The tSZ power spectrum of our $512^3$ simulation agrees well with the average of our ten $256^3$ simulations. A large number of simulations are needed to properly sample the high-mass end of the cluster mass function and hence accurately deal with sample (cosmic) variance. Alternatively, larger cosmological volumes can compensate since they contain enough statistics on the large scale modes that are responsible in part for forming the highest-mass clusters which are also the rarest events. This, however, is quite challenging as we require the same (high-)resolution to accurately follow the physics in the cluster cores which is needed to obtain profiles that match current X-ray data. Our $256^3$ simulations do not quite sample large enough scales to provide a fully converged kSZ power spectrum at low $\\ell$ since we miss the long-wavelength tail of the velocity power spectrum. We also have ignored the patchy re-ionization kSZ which could be a significant contributor, up to $50\\%$ of the total kSZ \\citep[e.g.,][]{2008MNRAS.384..863I, 2007ApJ...660..933I}. We have found the $\\ell <2000$ multipole range to be relatively insensitive to cooling and feedback, at least for the range constrained by the X-ray data. We did find the higher multipole range ($\\ell\\sim2000-10000$) probed by the high-resolution ACT and SPT CMB telescopes is sensitive to the feedback prescription; hence the high-$\\ell$ SZ power spectrum can be used to constrain the theory of intracluster gas, in particular for the highly uncertain redshifts $>0.8$. In addition to the SZ power spectrum probe, our simulations can be used to address the cosmological significance of cluster counts as derived from the SZ effect. Counts provide complementary constraints on parameters that help to break some degeneracies that are present in the power spectrum method. By employing inhomogeneous, localized and self-regulated feedback we are not only able to match recent X-ray reconstructions of cluster core regions, but also decrease the tension in $\\sigma_8$ estimated from SZ power with $\\sigma_8$ from other cosmological probes. However, only a detailed confrontation between simulations exploring the vast terrain of feedback options with the rapidly improving high resolution observations of cluster interiors can move the theory of cluster gas physics and its use for precision cosmology forward. We thank Norm~Murray, Volker~Springel, Hy~Trac, Jerry~Ostriker, Gil~Holder, Niayesh Afshordi and Diasuke Nagai for useful discussions. Research in Canada is supported by NSERC and CIFAR. Simulations were run on SCINET and CITA's Sunnyvale HPC clusters." }, "1003/1003.5700_arXiv.txt": { "abstract": "An increasing number of Active Galactic Nuclei (AGNs) exhibit broad, double-peaked Balmer emission lines,which represent some of the best evidence for the existence of relatively large-scale accretion disks in AGNs. A set of 20 double-peaked emitters have been monitored for nearly a decade in order to observe long-term variations in the profiles of the double-peaked Balmer lines. Variations generally occur on timescales of years, and are attributed to physical changes in the accretion disk. Here we characterize the variability of a subset of seven double-peaked emitters in a model independent way. We find that variability is caused primarily by the presence of one or more discrete ``lumps'' of excess emission; over a timescale of a year (and sometimes less) these lumps change in amplitude and shape, but the projected velocity of these lumps changes over much longer timescales (several years). We also find that all of the objects exhibit red peaks that are stronger than the blue peak at some epochs and/or blueshifts in the overall profile, contrary to the expectations for a simple, circular accretion disk model, thus emphasizing the need for asymmetries in the accretion disk. Comparisons with two simple models, an elliptical accretion disk and a circular disk with a spiral arm, are unable to reproduce all aspects of the observed variability, although both account for some of the observed behaviors. Three of the seven objects have robust estimates of the black hole masses. For these objects the observed variability timescale is consistent with the expected precession timescale for a spiral arm, but incompatible with that of an elliptical accretion disk. We suggest that with the simple modification of allowing the spiral arm to be fragmented, many of the observed variability patterns could be reproduced. ", "introduction": "It is now generally accepted that at the heart of an Active Galactic Nucleus (AGN) lies a supermassive black hole with a mass in excess of $10^{6}\\;\\msol$ that is accreting matter from its host galaxy \\citep[see, e.g.,][]{salpeter64,rees84,petersonagn}. As matter spirals inward, it forms an equatorial accretion disk around the black hole (such as that described by Shakura \\& Sunyaev 1973), whose inner portions are heated to UV-emitting temperatures by ``viscous'' stresses which tap into the potential energy of the accreting gas. The most direct kinematic evidence for {\\it large-scale} accretion disks in AGNs comes from the broad, double-peaked Balmer emission lines detected in some AGNs. Double-peaked emission lines were first observed in the broad-line radio galaxies (BLRGs) Arp~102B \\citep*{ssk83,ch89,chf89}, 3C~390.3 \\citep{o87,p88}, and 3C~332 \\citep{h90}. Given the similarity between these lines and those observed in cataclysmic variables (CVs), these authors suggested that the lines originated in an accretion disk surrounding the black hole at distances of hundreds to thousands of gravitational radii ($\\rg$~=~$G\\;m_{\\rm BH}/c^{2}$, where {\\rm $m_{\\rm BH}$} is the mass of the black hole). As discussed in great detail in \\citet{eh03} and \\citet{gezari07}, alternative scenarios for the origin of the double-peaked lines are not consistent with the observed variability and multi-wavelength properties of double-peaked emitters. Thus we adopt this interpretation hereafter. A later survey of (mostly broad-lined) radio-loud AGN showed that $\\sim$~20\\% of the observed objects exhibited double-peaked Balmer emission lines \\citep{eh94,eh03}. \\citet{s03} found 116 double-peaked emitters in the Sloan Digital Sky Survey \\citep[SDSS;][]{y00}, representing 3\\% of the $z<0.332$ AGN population. In addition, several Low Ionization Nuclear Emission-line Regions \\citep[LINERs,][]{h80} were found to have double-peaked Balmer lines, including NGC~1097 \\citep*{sb93}, M81 \\citep{b96}, NGC~4203 \\citep{s00}, NGC~4450 \\citep{h00}, and NGC~4579 \\citep{b01}. Most recently, \\citet{mbek07} discovered that the obscured AGN NGC~2110 shows a double-peaked H$\\alpha$ emission line in polarized light. The prototypical double-peaked emitters (Arp~102B, 3C~390.3, and 3C~332) exhibited variations in their line profiles that took place on timescales of years \\citep*[for examples, see][]{vz91,zvg91,g99,shap01,sps00,serg02,gezari07}, with the most striking variability being occasional reversals in the relative strength of the red and blue peaks. This long-term variability is intriguing because it is un-related to more rapid changes in luminosity but might be related to physical changes in the accretion disk. This profile variability can be exploited to test various models for physical phenomena in the outer accretion disk. The observed {\\it long-term} variability of the double-peaked line profiles (on time scales of several months to a few years) and the fact that in many objects the red peak is occasionally stronger than the blue peak (in contrast to the expectations from relativistic Doppler boosting) mandates that the accretion disk is non-axisymmetric. A few of the best-monitored objects (Arp~102B, 3C~390.3, 3C~332, and NGC~1097) have been studied in detail using models that are simple extensions to a circular accretion disk. These include circular disks with orbiting bright spots or ``patches'' \\citep*[Arp~102B and 3C~390.3,][]{ne97,sps00,zvg91}, a disk which is composed of randomly distributed clouds that rotate in the same plane \\citep[Arp~102B,][]{sps00}, spiral emissivity perturbations \\citep[3C~390.3, 3C~332, and NGC~1097,][]{g99,sb03} and precessing elliptical disks \\citep[NGC~1097,][]{e95,sb95,sb97,sb03}. All of these non-axisymmetric models naturally lead to modulations in the ratio of the red to blue peak flux; in the case of the first two models, modulations occur on the dynamical timescale whereas for the latter two, the variations occur on the longer precession timescale of the spiral density wave or the elliptical disk. However, in all cases, these simple models do not reproduce the fine details of the profile variability. It is extremely important to determine whether the profile variability observed in the small number of well-monitored objects is common to all double-peaked emitters. With a larger sample, one can determine if there are {\\it universal} variability patterns that point to a global phenomenon that occurs in AGN accretion disks. Thus, a campaign was undertaken to observe a larger sample of $\\sim$20 radio-loud double-peaked emitters from the \\citet{eh94} sample, in addition to a few more recently discovered objects. This campaign spanned nearly a decade and many objects were observed 2--3 times per year, although some objects, particularly those in the southern hemisphere, were only observed once per year. In this paper we present results for seven objects whose long-term variability has not been studied in any detail: 3C~59, 1E~0450.3--1817, Pictor A, CBS~74, PKS~0921--213, PKS~1020--103, and PKS~1739+18C. The properties of these objects are given in Table~\\ref{variability_objects} and a plot of the temporal coverage of the observing campaign is provided in Figure~\\ref{obslog_fig}. \\citet{gezari07} present results for Arp~102B, 3C~390.3, 3C~332, PKS~0235+023, Mkn~668, 3C~227, and 3C~382. The most recent profile variability results for NGC~1097 are presented by \\citet{sb03}, while the remaining objects from this monitoring program will be presented in Flohic et al. (2010, in preparation). The primary goal of this paper is to characterize the observed variability patterns in these seven objects in a {\\it model-independent} way. Currently, the models which have been used to fit the line profiles represent the most simple extensions to a circular disk. As mentioned above, these models fail to reproduce the fine details of the profile variability. A detailed characterization of the variability will help guide the refinement of these models and more importantly inspire ideas for new families of models. In essence, these results will serve as benchmark that models can be easily tested against. This paper is organized as follows. In \\S\\ref{variability_reductions}, the observations and data reductions are described. \\S\\ref{variability_analysis} is devoted to describing the methods for characterizing the profile variability in a model-independent way and noting the common trends in the data. In \\S\\ref{variability_models} we detail the two simple models we will compare with the observations, including their physical motivation, the model profile calculation, and the common variability trends. In \\S\\ref{variability_discussion} we assess the viability of these two models through a comparison of the observed and predicted variability trends and timescales. We also suggest refinements to the models which might afford a better description of the observed variability trends. Finally, in \\S\\ref{variability_conclusions} we summarize the primary findings. ", "conclusions": "In this paper we have characterized, in a model-independent way, the variability of the broad, double-peaked H$\\alpha$ emission lines in seven objects (Pictor A, PKS 0921--213, 1E 0450.3--1817, CBS 74, 3C 59, PKS 1739+18, and PKS 1020--103) which have been monitored over the past decade. As indicated by the rms profiles, the greatest variability occurs at larger projected velocities, and in particular many objects show a peak in the rms profile at large negative projected velocities. Difference spectra showed that the variability is caused primarily by discrete lumps of excess emission that change in morphology and amplitude on timescales of a few years and drift in projected velocity on longer timescales. There are often multiple lumps of emission observed at a single epoch and they are generally located at both positive and negative projected velocities. For some objects the dynamical timescale is known to be only a few months, thus these lumps cannot be orbiting bright spots such as those used to model Arp~102B \\citep{ne97,sps00}. The most striking profile variations observed are changes in the ratio of the red to blue peak flux, thus this trend which has long been observed in some of the better-studied double-peaked emitters is very common indeed. In fact, with the exception of CBS~74, all of the objects in this study have a red peak that is stronger than the blue in at least 50\\% of the observations. Some objects in this sample are very extreme, most notably PKS 0921--213 and Pictor~A, in that the blue peaks, which are supposed to be boosted due to relativistic effects, are rarely observed to be stronger than the red peak. We also noted that many objects had profiles that were blue-shifted by up to 1000$\\;\\kms $, again contrary to the expectations of the simple circular disk model. Further observation of these objects is important; should they continue to show strong red peaks and overall blueshifts this will place important constraints upon models for the broad-line region in these objects. We compared these variability trends with those expected from two simple models, an elliptical accretion disk and a circular disk with a single-armed spiral emissivity perturbation. In general, neither of these models reproduces the observed variability trends in detail; most importantly spiral-arm models do not predict the presence of multiple lumps of emission at a single epoch and the elliptical disk model predicts profile parameters variations that are extremely smooth, uniform, and symmetric and also rms profiles with four peaks, neither of which is observed. From a consideration of physical timescales, at least for the three objects with a known black hole mass (Pictor A, PKS 0921--213, and 1E 0450.3--1817), the spiral arm models is able to produce variability on a reasonable timescale, while the eccentric disk model appears to be untenable. Thus we propose an extension to the spiral-arm model in which one or more {\\it fragmented} spiral arms are present in the accretion disk. This model retains many of the general theoretical characteristics of the simple, uniform spiral arm but is likely to produce the observed variability more successfully. This model is inherently stochastic and it will be necessary to perform simulations to determine whether, statistically, such a model can reproduce the types of behavior observed. Such work is already being undertaken, in the context of a clumpy disk by \\citet{fe08}, but has not yet been extended to the scenario proposed here. We reiterate that the models considered here are the simplest extensions to a circular disk and more sophisticated models will be required. In particular it may be important to consider an accretion disk wind or other outflow to explain the full range of behaviors that are observed. Finally, we note that significant variability appears to be occurring on timescales of less than a year in some objects. In particular lumps of emission were observed to change significantly in shape and/or amplitude within one year, and it is quite possible that some rapid variations are being missed by the current observing strategy. It would be very useful to monitor all objects at least twice per year, and to intersperse periods of intense monitoring (perhaps as often as every few weeks) for a few of the more variable objects such as 1E~0450.3--1817 or PKS~0921--213. In particular, to test the fragmented spiral-arm model, it is necessary to determine the lower limit for the timescale over which the individual lumps of emission change in amplitude and morphology, which would provide an estimate of the fragmentation timescale." }, "1003/1003.2911_arXiv.txt": { "abstract": "{ CMB polarization signal may be decomposed into gradient-like (E) and curl-like (B) mode. We have investigated E/B decomposition in pixel space. We find E/B mixing due to incomplete sky is localized in pixel-space, and negligible in the regions far away from the masked area. By estimating the expected local leakage power, we have diagnosed ambiguous pixels. Our criteria for ambiguous pixels (i.e. $r_c$) is associated with the tensor-to-scalar ratio of B mode power spectrum, which the leakage power is comparable to. By setting $r_c$ to a lower value, we may reduce leakage level, but reduce sky fraction at the same time. Therefore, we have solved $\\partial \\Delta C_l/\\partial r_c=0$, and obtained the optimal $r_c$, which minimizes the estimation uncertainty, given a foreground mask and noise level. We have applied our method to a simulated map blocked by a foreground (diffuse + point source) mask. Our simulation shows leakage power is smaller than primordial (i.e. unlensed) B mode power spectrum of tensor-to-scalar ratio $r\\sim 1\\times10^{-3}$ at wide range of multipoles ($50\\lesssim l \\lesssim 2000$), while allowing us to retain sky fraction $\\sim 0.48$. } ", "introduction": "Over the past years, CMB polarization has been measured by several experiments and is being measured by the Planck surveyor \\citep{DASI:data,DASI:instrument,DASI:I,DASI:II,DASI:III,DASI:3yr,QUaD1,QUaD2,QUaD:instrument,QUaD_improved, Planck_bluebook}. CMB polarization pattern may be considered as the sum of gradient-like E mode and curl-like B mode \\citep{Seljak-Zaldarriaga:Polarization,Kamionkowski:Flm}. In the standard model, B mode polarization is not produced by scalar perturbation, but solely by tensor perturbation. Therefore, measurement of B mode polarization makes it possible to probe the universe on the energy scale at inflationary period \\citep{Kamionkowski:Flm,Seljak-Zaldarriaga:Polarization,Modern_Cosmology,Inflation,Foundations_Cosmology}. In most inflationary models, tensor-to-scalar ratio $r$ is much smaller than one, and the WMAP 7 year data imposes an upper bound on $r<0.36$ at $95\\%$ confidence level \\citep{WMAP7:powerspectra,WMAP7:Cosmology}. Besides instrument noise, there are complications, which limits detectability of tensor perturbation. Imperfection in removing foreground and gravitational lensing imposes observational limit on tensor-scalar-ratio: $r\\sim10^{-4}$ and $r\\sim 3\\times10^{-5}$ respectively \\citep{B_lensing,B_mode_limit_foreground}. Due to the nature of the observation or heavy foreground contamination, reliable of estimation on CMB polarization signal is not available over a whole sky. Incomplete sky coverage leads to E/B mixing, and very significantly limit our capacity to measure tensor perturbation as well \\citep{Bunn:EB-Separation}. Therefore, there have been various efforts to understand and reduce E/B mixing \\citep{Kim:optimization,Kim:measuring_a2lm,Bunn:EB-Separation,EB_incomplete_sky,EB_harmonic,Smith:pseudo_EB}. It is best to implement E/B decomposition in map space, since diffuse foregrounds and point sources are well-localized in map space, and their spatial information are known relatively better than other properties. In this paper, we investigate E/B decomposition in pixel space. Our investigation shows that E/B mixing is highly localized in pixel space. Therefore, we may reduce E/B mixing effectively by excluding the ambiguous pixels. We have applied our method to simulated maps partially blocked by a foreground (diffuse + point source) mask. After excluding ambiguous pixels, we find that leakage power in retained pixels (sky fraction $\\sim 0.48$) is smaller than primordial (i.e. unlensed) B mode power spectrum of tensor-to-scalar ratio $r\\sim 1\\times10^{-3}$ at wide range of multipoles ($50\\la l \\la 2000$). The outline of this paper is as follows. In Sec. \\ref{Stokes}, we discuss all-sky analysis of CMB polarization. In Sec. \\ref{pixel_filter}, we derive E/B decomposition in pixel space. In Sec. \\ref{cutsky}, we discuss the application to cut sky, and the method to diagnose ambiguous pixels. In Sec. \\ref{simulation} and \\ref{scale}, we present our simulation result. In Section \\ref{Discussion}, we summarize our investigation. In Appendix \\ref{noise}, we discuss error analysis of pseudo $C_l$ estimation, and show interpixel noise correlation may be neglected. ", "conclusions": "\\label{Discussion} We have investigated E/B decomposition in pixel space, and shown that we may produce E/B decomposed maps by convolving polarization maps with certain filter functions of a sharp peak. We find that E/B mixing due to incomplete sky is localized in pixel-space, and negligible in the regions far away from masked area. By estimating the expected local leakage power and comparing it with the expected pure mode power, we have diagnosed ambiguous pixels and excluded them. Our criteria for ambiguous pixels (i.e. $r_c$) is associated with the tensor-to-scalar ratio of B mode power spectrum, which the leakage power is comparable to. The estimation error $\\Delta C_l$ may increases with lower $r_c$, because sky fraction decreases. Therefore, we have solved $\\partial \\Delta C_l/\\partial r_c=0$ and obtained the optimal $r_c$, which minimizes the estimation error, given a foreground mask and noise level. We have applied our method to simulated maps blocked by a foreground mask. Simulation shows that leakage power is subdominant in comparison with unlensed B mode power spectrum of $r\\sim 1\\times10^{-3}$ at wide range of multipoles ($50\\la l \\la 2000$), while pixels of sky fraction $0.48$ are retained. We may apply our method equally to small sky patch observation, by treating unobserved sky as masked region. From simulation with sky patch of simple symmetric shape, we have confirmed our method reduce E/B mixing very effectively. A rigorous investigation is deferred to a separate publication. Noise is slightly correlated from pixel to pixel in E and B maps, even when interpixel correlation is absent in Q and U maps. However, this interpixel noise correlation induced by E/B decomposition is not confined to our method, but E/B decomposition in general. Besides that, we find interpixel noise correlation may be neglected without sacrificing the accuracy of error analysis (refer to Appendix \\ref{noise} for details). Therefore, it does not limit the applicability of our method. Current observations such as WMAP were unable to detect B mode polarization. Therefore, we did not attempt to apply our method to observation data. When Planck polarization data of high Signal-to-Noise-Ratio (SNR) are available in near future, we may apply our method to the data, and be able to detect B mode polarization." }, "1003/1003.5536_arXiv.txt": { "abstract": "We consider the problem of estimating filamentary structure from planar point process data. We make some connections with computational geometry and we develop nonparametric methods for estimating the filaments. We show that, under weak conditions, the filaments have a simple geometric representation as the medial axis of the data distribution's support. Our methods convert an estimator of the support's boundary into an estimator of the filaments. We also find the rates of convergence of our estimators. ", "introduction": "Filaments are one-dimensional curves embedded in $\\mathbb{R}^d$ where $d>1$. Filament estimation has important applications in many fields including astronomy, geology, and medicine. Our basic filament model is \\begin{equation}\\label{eq::first} Y_i = f(U_i) + \\epsilon_i \\end{equation} where $f:[0,1] \\to \\mathbb{R}^d$. The unobserved variables $U_1,\\ldots, U_n$ are drawn from a distribution $H$ on $[0,1]$ and $\\epsilon_1, \\ldots, \\epsilon_n$ are drawn from a mean zero noise distribution $F$. The goal is to estimate \\begin{equation} \\Gamma\\equiv \\Gamma_f = \\{ f(u): 0 \\leq u \\leq 1\\}. \\end{equation} Later, we extend the model to include background clutter, other $Y_i$'s drawn uniformly from a compact set containing the filaments. See Figure \\ref{fig::model}. Estimating $f$ is an example of one-dimensional manifold learning. It may also be regarded as a type of principal curve estimation. There is a plethora of available statistical methods that can, in principle, be used for estimating filaments. These include: principal curves (\\cite{hs::1989}, \\cite{kegl::2000}, \\cite{sandilya::2002}, and \\cite{smola}); nonparametric, penalized, maximum likelihood \\citep{tibs::1992}; beamlets (\\cite{Donoho01beamletsand}, and \\cite{arias::2006}); parametric models (\\cite{stoica-etal:2007}); manifold learning techniques (\\cite{isomap}, \\cite{lle}, and \\cite{Huo02locallinear}); gradient based methods (\\cite{novikov-etal:2006}, and \\cite{us::2009} and methods from computational geometry (\\cite{Dey}, \\cite{Lee:1999}, and \\cite{Cheng:2005}). In this paper, we make some connections between the statistical problem and some ideas from computational geometry. We propose new, simple, nonparametric estimators for $\\Gamma_f$, and we find their rates of convergence. To the best of our knowledge, our methods are the first that are computationally simple, consistent, and have given rates of convergence with the exception of \\cite{Cheng:2005}. However, our methods are simpler than those in \\cite{Cheng:2005}, our assumptions are weaker, our loss function is more stringent and our estimators have faster rates of convergence. The optimal rates of convergence for this problem appear to be unknown. In related work (\\cite{us}) we derived the minimax rate under stringent conditions. In ongoing work, we are finding the minimax rate under more general conditions. These rates depends critically on various features of the noise distribution $F$. The methods in this paper are unlikely to be minimax optimal. Nonethless, they achieve reasonable rates of convergence and are simple to compute. Our basic strategy involves two steps: \\begin{enumerate} \\item Construct a set of fitted values that are close in Hausdorff distance to the filament. \\item Extract a curve from this set of fitted values. \\end{enumerate} \\begin{figure} \\vspace{-.5in} \\begin{center} \\includegraphics[width=4in,height=4in]{model} \\end{center} \\vspace{-.5in} \\caption{These plots illustrate the filament model. Top left: some points $U_i$ on $[0,1]$ are mapped to $\\Gamma_f$ by $f$. Top right: noise is added to the points. Bottom left: a larger sample. Bottom right: background clutter has been added.} \\label{fig::model} \\end{figure} {\\em Motivation.} The need to identify filamentary structures arises in a wide variety of applications. In medical imaging, for instance, filaments arise as networks of blood vessels in tissue and need to be identified and mapped. In remote sensing, river systems and road networks are common filamentary structures of critical importance (\\cite{lacoste-etal:2005,stoica-etal:2004}). In seismology, the concentration of earthquake epicenters traces the filamentary network of fault lines. Filaments are of particular interest in astronomy because the distribution of galaxies in the universe is concentrated on a network of filaments that is often called the ``cosmic web.'' Indeed, astronomers have substantial literature on the problem of estimating filaments; see \\cite{luo-vishniac:1995}, \\cite{weygaert}, \\cite{martinez-saar:2002}, \\cite{barrow::1985}, \\cite{stoica-etal:2005}, \\cite{eriksen-etal:2004}, \\cite{novikov-etal:2006}, \\cite{sousbie-etal:2006} and \\cite{stoica-etal:2007}. \\vspace{.5cm} {\\em Summary of Results.} Two key geometric ideas underlie our results -- the medial axis of a set and the thickness $\\Delta(f)$ of a curve $f$ -- both of which are defined in Section 3. The medial axis is like the median of a set. The thickness of a curve measures both the curvature and how close the curve comes to being self-intersecting. Our main results are the following: \\begin{enumerate} \\item If the noise level $\\sigma$ of $F$ is less than the thickness $\\Delta(f)$, the filament equals the medial axis of the support of $Y$'s distribution (Theorem \\ref{thm::filamentismedial}). \\item Any estimate of the boundary of the support of the distribution can be converted into an estimate of the filament that is close in Hausdorff distance to the true filament (Theorems \\ref{thm::edtestimator} and \\ref{thm::methodII}). If the rate of convergence of the boundary estimator is $r_n$ then the rate of convergence of the filament estimator is also $r_n$. \\item Our estimators produce a set of fitted values that contain the filament and are close to it in Hausdorff distance. In Section 5, we show how to extract curves from the set estimators that are Hausdorff close to the true filament. \\end{enumerate} Proofs of all results are given in Section \\ref{sec::proofs}. \\medskip {\\em Notation.} The boundary of a set $S$ is denoted by $\\partial S$. The Hausdorff distance between two sets $A$ and $B$ is \\begin{equation} d_H(A_1,A_2) =\\min \\left\\{\\delta:\\strut\\ A_1 \\subset A_2\\oplus\\delta \\;\\mathand A_2 \\subset A_1\\oplus\\delta\\right\\} \\end{equation} where \\begin{equation} A\\oplus\\delta = \\Union_{x\\in A} B(x,\\delta) \\end{equation} denotes the $\\delta$-\\emph{enlargement} of the set $A$, and $B(x,\\delta) = \\{ y:\\ ||y-x|| \\le \\delta\\}$ denotes a closed ball centered at $x$ with radius $\\delta$. If $A$ is a set and $x$ is a point then we write $d(x,A) = \\inf_{y\\in A}||x-y||$. The closure of $A$ is denoted by $\\overline{A}$ and the complement of $A$ by $A^c$. A curve is a map $f:[0,1]\\to \\mathbb{R}^d$. Throughout, we use symbols like $C, c_0, c_1\\ldots $ to denote generic positive constants whose value may be different in different expressions. ", "conclusions": "In recent work (\\cite{us}) we found the minimax rate for this problem under restrictive conditions (but in general dimensions). In current work, we are finding the minimax rates in general. This is a difficult problem because the rate depends critically on features of the noise distribution $F$. Moreover, the problem is essentially a deconvolution problem since the variables $\\xi_i = f(U_i)$ are unobserved and corrupted by noise. We will report on these results elsewhere. The estimators presented here are not minimax but are appealing because of their simplicity. Finding a practical estimator that achieves the minimax rate is an open question. Our approach, instead, consists of two steps: producing a set of fitted values $\\hat\\Gamma$ and then extracting a curve from $\\hat\\Gamma$. We gave two specific methods for obtaining the fitted values and a curve extraction method for each of the two approaches. The resulting estimators have reasonably fast rates of convergence. The noise model is critical. We assumed compact support which is reasonable for many applications. Without compact support, the behavior of the methods changes substantially as it does in nonparametric measurement error problems. It is interesting to compare our results to those in \\cite{Cheng:2005}. They show that each of their fitted values is $O_P((\\log n/n)^{1/8})$ from the filament. Under weaker conditions than they assumed, we get a rate which is faster as long as $\\alpha$ is not too large. (They implicitly assume that $\\alpha=0$.) Also, our rate is in Hausdorff distance which is a stronger notion of closeness than used in their paper. Currently, we are pursing several extensions of our results. These include: the aforemetioned extensions to higher dimensions (manifold learning), relaxing the smoothness condition, relaxing the constant $\\sigma$ condition, noise distributions with non-compact support and comparisons with beamlets. We are also investigating data-driven methods for choosing the tuning parameter $\\epsilon$ and we are studying the theoretical properties of the decluttering technique. \\newpage" }, "1003/1003.0557.txt": { "abstract": "Janus and Epimetheus are famously known for their distinctive horseshoe-shaped orbits resulting from a 1:1 orbital resonance. Every four years these two satellites swap their orbits by a few tens of kilometers as a result of their close encounter. Recently \\citet{TiThBu2009} have proposed a model of rotation based on images from the Cassini orbiter. These authors inferred the amplitude of rotational librational motion in longitude at the orbital period by fitting a shape model to the recent Cassini ISS images. By a quasiperiodic approximation of the orbital motion, we describe how the orbital swap impacts the rotation of the satellites. To that purpose, we have developed a formalism based on quasi-periodic series with long and short-period librations. In this framework, the amplitude of the libration at the orbital period is found proportional to a term accounting for the orbital swap. We checked the analytical quasi-periodic development by performing a numerical simulation and find both results in good agreement. To complete this study, the results regarding the short-period librations are studied with the help of an adiabatic-like approach. ", "introduction": "\\label{sec:intro} The orbital motion of Janus and Epimetheus presents a peculiar horseshoe-shaped orbit resulting from a 1:1 orbital resonance (e.g. \\citealt{DeMu1981b}; \\citealt{YoCoSyYo1983} ; \\citealt{MuDe1999}; \\citealt{JaSpPoBeCoEvMu2008} and references therein). Every four years the two satellites swap their orbits by a few tens of kilometers as a result of their close encounter. As the mass of Janus is 3.6 times greater than the mass of Epimetheus, the dynamical motion of the latter is more sensitive to the swap than the dynamical motion of Janus. The rotational motion of the satellites depends mainly on the gravitational torque of Saturn acting on the dynamical figure of each moon. The expression of the gravitational torque is: \\begin{equation} \\vec{T} = \\frac{3GM_S}{r^3} \\vec{u} \\times [I] \\vec{u} \\label{torquepointless} \\end{equation} with $G$ the gravitational constant, $M_S$ the mass of Saturn, $[I]$ the inertia tensor of the moon, $r$ the distance between Saturn and the moon, and $\\vec{u}$ the unit vector toward Saturn in the moon's reference frame. The gravitational torque $\\vec T$ depends on the relative Saturn-moon distance, hence the swap also yields his signature on the rotational motion of the satellites. First estimates of the rotational motion of the two coorbital satellites Janus and Epimetheus have been obtained by \\citet{TiThBu2009}. From images provided by the Cassini orbiter, they fitted a numerical shape model of the moons, which included the amplitude of the libration in longitude. The libration in longitude corresponds to the oscillation of the body along its equatorial plane. \\citet{TiThBu2009} obtained an amplitude of $5.9^{\\circ} \\pm 1.2^{\\circ}$ for Epimetheus. For Janus, the uncertainty on the fit of the libration determination is too large to yield an accurate librational amplitude. However, based on their shape model, Tiscareno et al. suggested a value of $0.33^{\\circ} \\pm 0.06^{\\circ}$ for the amplitude of the libration in longitude. In addition, they identified an unexplained constant phase of $5.3^{\\circ} \\pm 1^{\\circ}$ for Janus, whereas for Epimetheus such offset is in the error bar. A recent numerical study by \\cite{No2010} explored the three-dimensional rotational motion of these satellites based on the numerical shape deduced by \\citet{TiThBu2009}. Noyelles' study suggests a strong influence of the swap on the rotational motion of Janus and Epimetheus, which we propose to explore in the present study. The orbital motion of Janus and Epimetheus appears to be very regular, at least over a timescale of several thousands of years. Thus, the trajectories of these satellites can be considered as quasi-periodic. Schematically these trajectories evolve on three different timescales. The shortest corresponds to the mean motion of the satellites with a period of about $0.7$ days. The second component has a period of $8$ years and is associated with the close encounters of the satellites. The long-period component is the secular variations of the satellites eccentricities and inclinations over periods of a few thousands years. Since the rotation of Janus and Epimetheus is synchronous with their orbital motion, it reflects these different timescales. One of the goals of this paper is to understand the influence of the $8$-year horseshoe motion on the rotational librations of Janus and Epimetheus. To this purpose, we develop an analytical solution of the rotation of these bodies that can simulate the main features of their spin. Then, in Section \\ref{sec:orbit}, we model the orbits of the two satellites through quasi-periodic expansions. The fourth section is dedicated to the description of the librational motion of each satellite. We develop three approaches to describe the librational motion in details: (1) a quasi-periodic development that highlight the fundamental frequencies involved in this problem, (2) an adiabatic invariant approach that focuses on the short-period librations, and (3) a numerical approach to reach high accuracy. Then we discuss the offset of Janus's orientation and investigate the effect of high spherical harmonics (order 3) and tidal coupling. We also discuss the influence of triaxiality on libration amplitudes. ", "conclusions": "In this paper we have investigated the librational motion of the co-orbital satellites Janus and Epimetheus by using three methods: (1) a perturbative technique based on quasi-periodic expansions, (2) an adiabatic invariant approach by expanding in power series of the small parameter $\\nu/\\nb$, and (3) a numerical integration. With the perturbative technique, we have detailed the librational behavior. For both satellites the solutions are composed of long-period librations linked to the orbital swap and short-period librations related to the orbital period. We found that the amplitudes of the short-period librations depend on the magnitude of the forcing and the proximity to the resonance, as for the librations analyzed in a Keplerian framework, but also on Bessel functions of the amplitude of orbital libration of the moons mean longitudes along their horseshoe orbit. These amplitudes bear the signature of the mass distributions in the satellites and are crucial to investigate the internal structure signature of these objects. On the other hand, the amplitudes related to the long-period librations do not contain any information on the distribution of mass. The numerical integration allows us to assess the accuracy of the perturbative development. The accuracy of the analytical solution is good for Janus but poor for Epimetheus because its mass is smaller than Janus and therefore its dynamics is more perturbed. In addition, in order to obtain a compact analytical solution easy to manipulate, we have developed an adiabatic approach yielding directly the amplitude of the short periods. The analytical approaches have been developed in the most general formalism and may be applied for co-orbitals like Telesto, Calypso, Helene and Polydeuces. The adiabatic approach seems a convenient approach to fit the short-period librations to the observations. %" }, "1003/1003.0063_arXiv.txt": { "abstract": "Spontaneous rapid growth of strong magnetic fields is rather ubiquitous in high-energy density environments ranging from astrophysical sources (e.g., gamma-ray bursts and relativistic shocks), to reconnection, to laser-plasma interaction laboratory experiments, where they are produced by kinetic streaming instabilities of the Weibel type. Relativistic electrons propagating through these sub-Larmor-scale magnetic fields radiate in the jitter regime, in which the anisotropy of the magnetic fields and the particle distribution have a strong effect on the produced radiation. Here we develop the general theory of jitter radiation, which includes (i) anisotropic magnetic fields and electron velocity distributions, (ii) the effects of trapped electrons and (iii) extends the description to large deflection angles of radiating particles thus establishing a cross-over between the classical jitter and synchrotron regimes. Our results are in remarkable agreement with the radiation spectra obtained from particle-in-cell simulations of the classical Weibel instability. Particularly interesting is the onset of the field growth, when the transient hard synchrotron-violating spectra are common as a result of the dominant role of the trapped population. This effect can serve as a distinct observational signature of the violent field growth in astrophysical sources and lab experiments. It is also interesting that a system with small-scale fields tends to evolve toward the small-angle jitter regime, which can, under certain conditions, dominate the overall emission of a source. ", "introduction": "There is a lore that a relativistic particle of charge $e$ and a Lorentz factor $\\gamma$ moving through a magnetic field $B$ produces synchrotron radiation, whose spectrum peaks at $\\omega_s\\sim (eB/mc)\\gamma^2$, has an asymptotic $\\omega^{1/3}$ dependence below the peak and falls off exponentially at higher frequencies (it makes a second power-law for an isotropic ensemble of particles having a power-law distribution in energy). This is often true, but not always. If the field is inhomogeneous on scales comparable or smaller then the particle Larmor radius, $\\lambda_B\\lesssim R_L\\sim\\gamma m c^2/e\\langle B\\rangle$, the produced radiation spectrum may be far different from synchrotron. Gamma-ray bursts, supernovae shocks, relativistic pulsar winds and shocks, relativistic jets from quasars and active galactic nuclei, magnetic reconnection sites, plasmas produced by high-intensity lasers --- they all are the high-energy density environments where conditions are favorable for the spontaneous magnetic field production. The field generation via the Weibel instability \\citep{Weibel59,Fried59,Silva+03,WA} or its modifications \\citep{D+06,Bret+08,Bret09,Jacob+08} has been predicted to occur in astrophysical shocks with low ambient magnetic field and rare particle collisions, e.g., in gamma-ray burst and large-scale structure shocks \\citep{ML99,MSK06,MZ09}. It has been observed in numerical simulations of relativistic non-magnetized shocks \\citep{Nish+04,Jacob+04,Spitk08,Keshet+09,Nish+09}, nonrelativistic shocks unmagnetized and weakly magnetized shocks \\citep{KT08,KT10}, cosmic rays interacting with a pre-shock medium \\citep{Jacek+10}, magnetic reconnection in electron-positron relativistic and non-relativistic plasmas \\citep{ZH08,Swisdak+08,Liu+09}, as well as in simulations of and even real laser plasma experiments \\citep{Ren+04,T+03}. Given such a ubiquity of the process at hand, a natural question to ask is: Are there any observational signatures, which can benchmark the process in astrophysical sources and, if any, what can we learn about the physical conditions there? With the radiation techniques being developed and implemented into numerical codes \\citep{Hededal05,Nish+08,SS09,Jacob+2010}, we will soon be able to answer this question in detail \\citep[for instance,][showed that PIC simulations can realistically model some astrophysical sources]{MS09}. In order to correctly interpret the results of simulations and observational data, a comprehensive theory of radiation processes in a strong small-scale magnetic turbulence is, therefore, of great demand. The effects of small-scale inhomogeneities on radiation emission have been of long-standing theoretical interest \\citep[see, for instance,][and many more]{LP53,Migdal54,Migdal56,GS65,LL}. Techniques developed in these papers have further been applied to synchrotron radiation from large-scale homogeneous magnetic fields with a small-scale random field component, as a model of radio emission by cosmic rays in the interstellar medium \\citep[see, e.g.,][and references therein]{NT79,BNT80,Toptygin,TF87}. The so-called perturbative approach of radiation emission from random small-scale magnetic fields without a large-scale component has first been discussed as a model of radiation from Weibel-mediated relativistic colisionless shocks of gamma-ray bursts \\citep{M00}, where it was referred to as jitter radiation. This approach was generalized in \\citep{F06} and further corrected% \\footnote{It is important to note here some problems in \\citep{F06}, which are relevant to the present paper. First, it was argued that the spectrum $F_\\nu\\propto \\nu^1$ below the peak, ``valid in the presence of ordered small-scale magnetic field fluctuations, does not occur in the general case of small-scale random magnetic field fluctuations.\" This statement was shown \\citep{M05,M06} to be flawed in that jitter radiation from {\\em random} magnetic fluctuations with a fairly general distribution function (not just ``ordered small-scale'' fields) does allow for $\\propto \\nu^1$ spectra. Moreover, the entire range from $\\propto \\nu^1$ to $\\propto\\nu^0$ is allowed for a single electron emission, softer spectra can be expected for some electron energy distributions. Second, another confusing issue is related to the absorption-like $\\propto\\nu^2$ jitter spectrum. Such a spectrum is due to plasma dispersion and does not occur in the absence thereof, as the reader might incorrectly infer \\citep[see][for more discussion]{M05}. Finally, we should note that contrary to the claims, none of their papers treat the jitter regime properly. The approximations made in their analysis apply to the systems with large-scale plus small-scale isotropic turbulent fields, which do not hold for the generic Weibel-like magnetic turbulence.} % in \\citep{M05,M06,MPR09,RPM10}. Because of this and also because \\citep{M06} was the first to consider anisotropic magnetic turbulence (e.g., Weibel- or filamentation-instability-generated magnetic field turbulence), we will refer to this paper in the following discussion. We also would like to mention here a recent paper by \\citep{RK10}, who developed a new algorithm to compute radiation from small-scale turbulent fields. In this paper we develop a theory of jitter radiation that accounts for anisotropies of the magnetic field and particle velocity distributions, including a trapped population, and further extend the theory to the large angle jitter regime. Our theoretical findings are tested with dedicated particle-in-cell simulations. Interesting conclusions are presented in the final section. ", "conclusions": "The primary results of this paper are as follows. First, we present a general expression for the spectral energy per solid angle emitted by an ensemble of particles in the small-angle jitter approximation, Eq. (\\ref{dWmain}). Second, we have found that the electrons streaming through the filaments and being trapped in them produce a transient hard spectrum, Eqs. (\\ref{Omegab},\\ref{trapjitt}). Third, we analyzed the large-angle deflection regime and showed that the spectrum starts to resemble the synchrotron spectrum near the peak, but a new spectral break at a lower frequency appears, Eq. (\\ref{largejitt}). From the positions of the spectral peak and break, one can deduce the field correlation length. Fourth, PIC simulations show that the radiation spectrum produced at the onset of and during the phase of the exponential growth of the magnetic field is grossly inconsistent with synchrotron, Fig. \\ref{spec}. The appearance of such a spectrum in the beginning of an emission episode can be used as a benchmark signal of the onset of the magnetic field generation in astrophysical sources and laboratory experiments. Among possible astrophysical systems where such emission can be or could have already been observed are gamma-ray bursts. The data show hard synchrotron-violating spectra in some bursts and the majority of spectra are flat \\citep{Preece+98,Preece+00,Kaneko+06}, which are difficult to explain within the synchrotron model. Such spectra have recently been interpreted in the jitter emission paradigm \\citep{MPR09}. We should caution the reader that ``blind'' application of the jitter spectrum template for interpretation of observational data, without checking the physical conditions at and the validity of the jitter approximation for the source in hand (galactic and quasar jets, supernovae remnants, etc.), can yield incorrect results. Possible laboratory experiments include laser-plasma interactions in which a beam (e.g., a probe electron beam) propagating through turbulent fields can emit jitter radiation \\citep{RM11}. Fifth, although the spectra after saturation are consistent with the synchrotron `template', an overall trend of the system toward the small-angle jitter regime (i.e., toward $\\delta_{\\rm jitt}<1$) is observed in Fig. \\ref{delta}, which suggest the scaling: $\\delta_{\\rm jitt}\\propto\\sqrt{\\epsilon_B}\\,\\Gamma (\\omega_p t)^{-0.5}$. Although the magnetic field decays, it does so rather slowly, $B(t)\\propto (\\omega_p t)^{-1}$, see Fig. \\ref{epsB} (for example, for $\\Gamma=10$ the systen shall return to the small-angle jitter regime at times $t\\ga100$). So, if the field is continuously produced (as in the case of the propagating shock, for instance) the field decay can be compensated by the increase of the emitting volume, so that the total spectral emissivity will increase logarithmically, $P_{\\rm tot}(\\omega)\\propto\\int B\\,dV\\propto \\ln(\\omega_p t)\\propto \\ln(\\omega_p L_{ps}/c)$, assuming these scalings hold at asymptotically late times, where $L_{ps}$ is the size (longitutinal extent) of the post-shock medium. Since the plasma time $\\sim\\omega_p^{-1}$ is generally very short in astrophysical sources, the overall time-integrated (and even time resolved, but with a coarse temporal resolution) spectrum can be dominated by the small-angle jitter spectrum, which is expected to be flat (unless the strong anisotropy of the fields and/or particles is somehow maintained). Here we also comment on the relation of our results and the simulations of \\citep{SS09}. First of all, one should understand that the physical set-ups are entirely different. In this paper we study the evolution of radiation during the filamentation instability: the plasma in the simulation box is initially homogeneous, the particle distribution is unstable and the entire system is intrinsically non-stationary. In contrast, \\citet{SS09} simulated a well-developed collisionless shock: a steady state system with very slow, if any, evolution of conditions and radiation spectra. Thus, the two simulations are complementary and the direct comparison of them is not well-posed. Since, however, the shock is moving through a medium with a constant speed, the temporal evolution of the filamentation instability and its subsequent saturation and further nonlinear evolution studied here is, to a certain degree, represented by the spatial profiles in the pre-shock and post-shock domains. The filamentation instability and its saturation occur far in front of the shock (hundreds or thousands skin lengths, in typical simulations) by particles escaping from the shock, hence this region would roughly correspond to the early and saturation times in our simulations. After saturation, mergers of magnetic filaments increase their sizes and radiation spectrum now mimics synchrotron, as our simulations show. In shock simulations, this merging stage occurs in a large region in front of the shock and this is where the radiation is collected in \\citep{SS09}. Their results are, thus, in agreement with the ones presented here. The shock itself and the medium just behind the shock do not correspond to our simulations. However, the magnetic field strength decreases behind the shock and in a few hundred skin lengths radiation should enter the jitter regime. Unfortunately, radiation from neither the early pre-shock, nor from the far downstream regions have been shown in \\citep{SS09}. We stress that radiation from the far downstream can be of great importance and dominate the entire shock emission, provided the scalings presented in the previous paragraph hold." }, "1003/1003.3583_arXiv.txt": { "abstract": "{The winter seeing at Concordia is essentially bimodal, excellent or quite poor, with relative proportions that depend on altitude above the snow surface. This paper studies the temporal behavior of the good seeing sequences.} {An efficient exploitation of extremely good seeing with an adaptive optics system needs long integrations. It is then important to explore the temporal distribution of the fraction of time providing excellent seeing.} {Temporal windows of good seeing are created by a simple binary process. Good or bad. Their autocorrelations are corrected for those of the existing data sets, since these are not continuous, being often interrupted by technical problems in addition to the adverse weather gaps. At the end these \tcorrected autocorrelations provide the typical duration of good seeing sequences. This study has to be a little detailed as its results depend on the season, summer or winter.} {Using a threshold of 0.5 arcsec to define the ``good seeing'', three characteristic numbers are found to describe the temporal evolution of the good seeing windows. The first number is the mean duration of an uninterrupted good seeing sequence: it is $\\tau_0=7.5$~hours at 8~m above the ground (15 hours at 20~m). These sequences are randomly distributed in time, with a negative exponential law of damping time $\\tau_1=29$~hours (at elevation 8~m and 20~m). The third number is the mean time between two 29~hours episodes. It is $T=10$~days at 8~m high (5 days at 20~m). } {There is certainly no other site on Earth, except for the few other high altitude Domes on the Antarctic plateau, that can get close to these really peculiar seeing conditions.} ", "introduction": "Regarded as astronomical sites, the highest points of the Antarctica plateau present many obvious advantages due to the local climate and the remoteness from any polluting civilization. They also benefit from an interestingly unique distribution of turbulence. This has been extensively measured at Dome C since the first winter-over permitted in 2005 by the French - Italian Concordia station operation (Aristidi et al, 2009). Winter and Summer display very different but both unusual vertical distributions of the turbulent energy. Generally speaking, the situation is dominated by the presence of a surface inversion layer that becomes very turbulent when the temperature gradient is strong in winter, and can completely vanish in summer when this gradient becomes flat. In summer it depends on the Sun's elevation, and is then strongly time dependent, with an optimum period of a few hours of excellent seeing every day in the middle of local afternoon (Aristidi et al, 2005). In the other 3 seasons, the mean seeing is essentially altitude dependent above the snow surface. Above this surface turbulent layer that contains, statistically, 95 percent of the total C$n^2$ along the line of sight, it has proved to be statistically independent of the season, within the measurements accuracy. Its mean value is between 0.3 and 0.4 arcsec as soon as the telescope is located above a sharply defined altitude threshold, that fluctuates around a mean value of the order of 25 m. The consequence is that the non summer seeing displays a nearly bimodal statistical distribution. It is indeed as good as 0.3 to 0.4 arcsec 50 percent of the time at 25m above the surface, this fraction of time decreasing to about 40 percent at 20m and slightly less than 20 percent at 8m. But it is obviously not equivalent to have 40 percent of good seeing spread in many short sequences of seconds to minutes rather than distributed in extended long sequences measured in hours or days. This paper addresses the temporal distribution of this good seeing percentage. It goes beyond the first analysis made in Aristidi et al. (2009), using a method to compensate for the gaps in the data. ", "conclusions": "Thanks to the very long data sets exploited here, the validity of the autocorrelation division for correcting the data interruptions and the final statistical robustness are well validated by the summer season analysis and results. At least at 8-m high. The reduced amount of data available at 20-m high makes the statistical robustness of the numbers visibly weaker, but the general tendency tends to reinforce the interpretation of the 8-m data. The 29-hour exponential decay of the distribution of the good runs starting times, assumed to depend on the meteorological situation, is found to be the same at both altitudes. The second characteristic time $\\tau_0$, that can be regarded as the minimum duration of nearly uninterrupted good runs, is found twice longer at 20 m, i.e. 15 hours versus 7.5 hours at 8 m. Finally, the number of episodes of excellent seeing is estimated to be twice more frequent at 20m, 20 times versus 10 times per winter, and that is also easy to understand with a reasonable assumption on the vertical motions of the boundary layer upper limit." }, "1003/1003.1779_arXiv.txt": { "abstract": "We study Q-ball formation in the expanding universe on 1D, 2D and 3D lattice simulations. We obtain detailed Q-ball charge distributions, and find that the distribution is peaked at $Q^{3D}_{\\rm peak} \\simeq 1.9\\times 10^{-2}(|\\Phi_{\\rm in}|/m)^2$, which is greater than the existing result by about 60$\\%$. Based on the numerical simulations, we discuss how the Q-ball formation proceeds. Also we make a comment on possible deviation of the charge distributions from what was conjectured in the past. ", "introduction": "\\label{sec:intro} If scalar fields are ubiquitous in nature, some of them may remain light in the early universe. Such scalar fields could be deviated from the low-energy minimum due to quantum fluctuations during inflation and/or modification of the scalar potential through gravitationally suppressed interactions with the inflaton. At a later time when the expansion rate becomes smaller than the mass, the scalar will start to roll down and oscillate about the potential minimum. The dynamics of such a scalar can exhibit rich phenomena in cosmology. We focus on a complex scalar field with a global U(1) symmetry at the potential minimum. The U(1) symmetry is not necessarily exact, and could be violated at high energy scales; our arguments apply if it is a good symmetry at low energy. If supersymmetry (SUSY) is realized in nature, there may be many such scalar fields. In Ref.~\\cite{Coleman:1985ki} Coleman found that, if the scalar potential is shallower than a quadratic potential, there is a non-trivial field configuration that minimizes the energy for a fixed U(1) charge $Q$. Since the solution has a spherical symmetry, the solution is named Q-ball. A Q-ball is a non-topological soliton whose stability is supported by the U(1) symmetry. It was noticed in Refs.~\\cite{Kusenko:1997zq,Dvali:1997qv,Kusenko:1997si,Enqvist:1997si} that Q-balls play an important role in a context of the Affleck-Dine (AD) mechanism~\\cite{Affleck:1984fy}. The mechanism utilizes a flat direction of the supersymmetric standard model (SSM), which possesses a non-zero baryon (or lepton) number. A flat direction responsible for the AD mechanism is referred to as the AD field (denoted by $\\Phi$ in the following). The AD field $\\Phi$ develops a large expectation value during inflation, and it starts to oscillate after inflation when the cosmic expansion rate becomes comparable to its mass. The baryon number is effectively created at the onset of the oscillations. Finally, $\\Phi$ decays into the ordinary quarks, leaving the universe with a right amount of the baryon asymmetry. It was realized however that, soon after the onset of oscillations, the AD field experiences spatial instabilities and deforms into clumpy Q-balls~\\cite{Kusenko:1997si}. Later it was shown in \\cite{Kasuya:2000wx} that most of the baryon asymmetry is absorbed into Q-balls. Its non-linear property necessitates numerical approach to the formation and the subsequent evolution. The properties of Q-balls crucially depend on the size of the U(1) charge $Q$. For instance, the mass per unit charge is a decreasing function of $Q$ in case of gauge-mediation type Q-balls; for a large enough $Q$, the Q-ball can be absolutely stable against the decay into nucleons and therefore contribute to dark matter~\\cite{Kusenko:1997si}. In the case of unstable Q-balls, the decay rate depends on the charge, and Q-balls can be very long-lived in a cosmological time scale for a sufficiently large $Q$~\\cite{Enqvist:1998xd}, since the Q-ball decay process takes place only in the vicinity of the surface~\\cite{Cohen:1986ct}. Thus it is of utmost importance to determine the charge distribution of the Q-balls at the formation. There is a number of numerical studies on the Q-balls. The Q-ball formation has been first examined by Kasuya and one of the authors (MK) on 3D lattice simulations in a context of gauge mediation~\\cite{Kasuya:1999wu,Kasuya:2001hg} and gravity mediation~\\cite{Kasuya:2000wx}, taking account of the cosmic expansion. They found that most of the baryon number is absorbed into the largest Q-balls in the lattices and estimated the typical Q-ball charge. However, the estimated charge might have been affected by relatively coarse spatial resolution due to the limited computational power at that time. The Q-ball charge distribution was studied in 2D~\\cite{Enqvist:2000cq} and 3D~\\cite{Multamaki:2002hv} lattices, with an emphasis on negative Q-balls formed for an orbit with a large ellipticity. In gravity mediation, however, the AD field naturally acquires a large kick into the phase space, leading to an orbit with a comparatively small ellipticity. As we will see later, the formation of negative Q-balls are also observed in our simulations for a large ellipticity. However, the final charge distribution does not fit well the empirical formula conjectured in Ref.~\\cite{Enqvist:2000cq} based on the entropy argument. The difference may be attributed to the fact that the charge re-distribution through exchanges of small secondary Q-balls seems less efficient in our simulations. Also the finite grid size could have affected the Q-ball evolution especially at a very late time when one Q-ball was represented by a single grid point. In our analysis we stop following the evolution once the spatial resolution becomes insufficient to describe the Q-ball solution, in order to avoid any artificial effects due to the finite grid size. Recently Tsumagari studied the Q-ball formation and thermalization processes in a great detail in a Minkowski space neglecting the cosmic expansion \\cite{Tsumagari:2009na}. In an expanding universe, however, the thermalization processes through Q-ball interactions may be weakened, and the resultant Q-ball configuration could be quite different. We will see in the following that the cosmic expansion indeed plays an essential role in the Q-ball formation. In this paper we study the Q-ball formation in a scalar potential motivated by the AD mechanism in gravity mediation and obtain the Q-ball charge distribution on 1D, 2D and 3D lattice simulations with the cosmic expansion for different values of ellipticity and strength of instabilities. We develop a sophisticated numerical code based on the 6th-order symplectic integrator equipped with a new Q-ball identification algorithm. The paper is organized as follows. In Sec.~\\ref{sec:dynamics}, we summarize the governing equation of Q-balls in the gravity-mediation model. In Sec.~\\ref{sec:setup}, we briefly mention the numerical scheme used in the simulations and, in Sec.~\\ref{sec:numerical}, we show numerical results obtained from simulations in 1D, 2D and 3D configurations. After that, Sec.~\\ref{sec:applications} is devoted to some applications of our numerical results, and we conclude in Sec.~\\ref{sec:conclusion}. ", "conclusions": "\\label{sec:conclusion} In this paper, we have studied physical properties of the gravity-mediation type Q-balls in the AD mechanism on 1D, 2D and 3D lattice simulations {\\it with} the cosmic expansion. We have found that including cosmic expansion has a crucial effect on the Q-ball formation and its subsequent evolution, using the 1D lattice simulations. In the non-expanding background, the Q-ball interactions through collisions never decouple until one large Q-ball is formed. In fact, we confirmed the finite box effect in our simulations without cosmic expansion. On the other hand, in the expanding background, the major collisions among large Q-balls do not occur frequently. Even in the cases of the 2D and 3D simulations, we expect that, without the cosmic expansion, the Q-ball evolution would be quite different. In various setups with the cosmic expansion, we have performed simulations in 2D and 3D systems to obtain the charge distributions of the Q-balls. For the initial circular orbit with $\\epsilon = 1$, the charge at the peak of the distribution is $Q^{2D}_{\\rm peak}\\simeq0.1\\unitQQ$ and $Q^{3D}_{\\rm peak}\\simeq1.9\\times 10^{-2}\\unitQQQ$. We have found that filamentary structure is formed when the initial fluctuation becomes of order unity, as pointed out in Ref.~\\cite{Enqvist:2000cq}. One of our new observations is that this filamentary structure plays an important role in the Q-ball formation and affects the final charge distribution. In particular, the Q-balls forming the peak of the distribution are mainly produced at the intersections of the filaments, while Q-balls produced from the side and void region have generically small charges and do not give main contribution to the charge distribution. Note however that these small Q-balls overwhelm the large Q-balls in the number. This argument holds true for a small value of $\\epsilon$, in which large positive and negative Q-balls are mostly produced from the excited Q-balls at the intersections. (We call the latter as the first generation Q-balls and the former as the second generation Q-balls in the text.) We have compared the empirical formula derived in Ref.~\\cite{Enqvist:2000cq} to our numerical results on the charge (not number) distribution on the 2D lattice and found that, if we fit the formula at the peak of the charge distribution, it underestimates the charge distribution at smaller $Q$. In contrast to the claim made by Ref.~\\cite{Enqvist:2000cq}, therefore, we conclude that these small Q-balls do not play any significant role to determine the Q-ball distribution. We also mention here that we stopped following the evolution once the Q-balls become no longer resolved compared with our simulation grid size, in order to avoid mis-identification of background fluctuations with the Q-balls as well as unphysical effects on the Q-ball evolution. This also may be the reason for the difference of ours from the results in the past. Our findings, especially detailed charge distribution of the Q-balls, will be very important for the AD baryogenesis and related topics, because the Q-balls are an essential ingredient of the mechanism. For instance, the Q-ball decay temperature is modified accordingly, which affects the resultant dark matter density in a scenario that the lightest SUSY particles are produced from the Q-ball decay. Also the charge distribution gives a lower bound on the evaporated charge, which in turn makes it difficult to realize a large hierarchy between the baryon and lepton asymmetries in the scenario using the L-balls. Lastly we mention the limitation of our numerical simulations. Strictly speaking, the Q-balls we identified in the simulations are in a excited state, in a sense that the $E/mQ$ ratio is slightly larger than $1$. We expect that these Q-balls will in the end relax to the Q-ball solution, emitting the excessive energy. Also we have neglected the interactions of the AD field with gauge fields and other matter fields. We leave these issues as well as further exploration of the application of our results for future work." }, "1003/1003.4706_arXiv.txt": { "abstract": "{We present new spectral (FPI and long-slit) data on the Eastern optical filament of the well known radionebula W50 associated with SS433. We find that on sub-parsec scales different emission lines are emitted by different regions with evidently different physical conditions. Kinematical properties of the ionized gas show evidence for moderately high ($V \\sim 100\\,\\kms$) supersonic motions. \\oiii$\\lambda$5007 emission is found to be multi-component and differs from lower-excitation \\sii$\\lambda$6717 line both in spatial and kinematical properties. Indirect evidence for very low characteristic densities of the gas ($n \\sim 0.1\\cmc$) is found. We propose radiative (possibly incomplete) shock waves in low-density, moderately high metallicity gas as the most probable candidate for the power source of the optical filament. Apparent nitrogen over-abundance is better understood if the location of W50 in the Galaxy is taken into account. } ", "introduction": "W50 radionebula was first catalogued in the work by \\citet{westerhout} and classified as a radiative-stage supernova remnant by \\citet{holcas69}. However, the nebula is strongly affected by the activity of the central source SS433 (see \\citet{dubner} for review). It is not evident whether the contribution from the initial supernova explosion, that probably precursed formation of the compact accretor in SS433 system, may be distinguished from the impact of the jet and wind activity of the central system. Most of the peculiarities of W50 may be attributed to the jet activity of SS433. The total power of the jets is of the order $10^{39}\\,\\ergl$. The jets are an immense source of energy capable to provide the energy of $\\sim 10^{51}\\,\\rm erg$ characteristic for a usual supernova remnant (SNR) in about $3\\times 10^4\\,\\yr$, that is at least several times lower than the observed dynamical age of W50 \\citep{yama}. The morphology and radio properties of the nebula are fairly reproduced in numerical simulations with the preceeding SNR \\citep{zavala,velazquez}. However, it is not clear whether the central, quasi-spherical part of the radionebula requires a preceeding SNR cavity or not. The simulations also do not consider the possible emission of the warm gas that may appear in some parts of the nebula where the density is high enough for the shocked material to cool. Warm gas in W50 was detected by \\citet{vdberg} as the Eastern and the Western filaments. Recent investigation by \\citep{Boumis} proved that warm gas emission is present in different parts of the nebula but peaks in several bright optical filaments offset from the areas bright in X-ray and radio ranges. The W50 optical filaments are distinguished among the known shell-like and SNR-like nebulae by the strong \\nii\\ lines that were attributed by \\citet{zealey} to low velocity shocks and by \\citet{shuder} to nitrogen over-abundance. In this article, we analyse the reasons for the apparent nitrogen over-abundance, involving Galactic abundance gradients (see \\S \\ref{sec:disc:abund} for details). In the next section we describe the observational material and the methods used for data reduction. In \\S~\\ref{sec:res} main results are presented. The results are discussed in \\S~\\ref{sec:disc}. ", "conclusions": "\\label{sec:conc} We come to the conclusion that the optical filaments of W50 are mainly powered by shocks with velocities of the order 100$\\div$200\\kms\\ and less. Flows with line-of-sight velocities differing by $\\pm$100\\kms\\ are detected in the \\oiii$\\lambda$5007 emission line produced in the recently shocked, relatively hot gas. For lower-ionisation lines kinematics is different: we detect only general broadening of the \\sii\\ line. Probably, cooler gas shows both lower turbulent velocities and larger number of velocity components. We also find at least two spatially resolved shock fronts in the FPI field of view propagating in different directions. Their line-of-sight velocities differ by about 50\\kms. Our results are compolementary to the results reported by \\citet{Boumis} and provide more detailed picture of a smaller part of the nebula. The optical emission-line spectrum of W50 Eastern filament bears the main signatures of a radiative (possibly, incomplete) shock such as enhanced collisionly excited lines of different ionization potentials. The observed gas shows a broad range of physical properties. In general, the gas is rarefied ($n \\sim 10^{-2}\\div 1\\,\\cmc$) and moderately over-abundant in oxygen and nitrogen. This apparent over-abundance is likely to be attributed to the high metallicity of the Galactic gas at the comparatively low ($\\sim 4\\div 5\\, \\kpc$) distance from the Galactic center. Nitrogen lines are primarily affected because of stronger nitrogen gradient in the Galactic disc." }, "1003/1003.0724_arXiv.txt": { "abstract": "A new hydrodynamics code aimed at astrophysical applications has been developed. The new code and algorithms are presented along with a comprehensive suite of test problems in one, two, and three dimensions. The new code is shown to be robust and accurate, equalling or improving upon a set of comparison codes. \\Fyris\\ will be made freely available to the scientific community. ", "introduction": "Multi-dimensional computational fluid dynamics is a mature field, with roots in the earliest days of computation. In the astrophysical setting, with compressible, trans--, and super--sonic flows, the hyperbolic Eulerian system of ideal gas equations are typically used, allowing for the formation of shock discontinuities. The three dimensional non-viscous Euler equations may be written in terms of the conserved variables (density, $\\rho$, momentum $m$, and total energy $E$) at time $t$; $U_t$, and flux vectors, $F(U)$, $G(U)$, $H(U)$: \\begin{equation} U_t + F(U)_x + G(U)_y + H(U)_z = 0, \\end{equation} with \\begin{eqnarray} U_t &=& \\left[ \\begin{array}{c} \\rho \\\\ m_x \\\\ m_y \\\\ m_z \\\\ E \\end{array} \\right]\\, , \\\\ F(U)_x = \\left[ \\begin{array}{c} v_x m_x + P \\\\ v_x m_y \\\\ v_x m_z \\\\ v_x ( E + P ) \\end{array} \\right] \\, , \\, G(U)_y &=& \\left[ \\begin{array}{c} v_y m_x \\\\ v_y m_y + P \\\\ v_y m_z \\\\ v_y ( E + P ) \\end{array} \\right] \\, , \\, H(U)_z = \\left[ \\begin{array}{c} v_z m_x \\\\ v_z m_y \\\\ v_z m_z + P \\\\ v_z ( E + P ) \\end{array} \\right] \\, , \\end{eqnarray} where the momentum terms are; $m_x = \\rho v_x$, $m_y = \\rho v_y$, $m_z = \\rho v_z$. The total energy per unit volume is $E = \\rho ( \\frac {1}{2} {\\bf v}^2 + e)$, where ${\\bf v}$ is the total velocity and $e$ is the specific internal energy. $v_x$, $v_y$ and $v_z$ are the velocity components. $P$ is the gas pressure. With a baric equation of state $P = EOS_{\\rm B}( \\rho, e, \\gamma)$, relating density, $\\rho$, pressure, $P$, and internal energy $e$ ( using the polytropic index $\\gamma$), the system can be closed. When the temperature is required, such as when cooling is used, a caloric equation of state, $T = EOS_{\\rm C}( \\rho, P, \\mu )$ is invoked, using, $\\mu$, the mean molecular weight. This system has been solved on a grid of cells by considering the cell average values and the flow, or fluxes, between cells. The methods used have often been founded on the solution of the Riemann problem, a simple non-linear interaction between two adjacent regions of ideal gasses, as first outlined by Godunov in the 1950s, \\citep{godunov59}, and \\cite{glimm65} in the 1960s. With the formation of shockwaves in a computed flow, a fundamental problem arises where high order methods inevitably lead to non-physical oscillations of the flow properties in cells near a discontinuity, and the pioneering work of Godunov showed, in the Godunov Theorem, that no monotonic shock capturing method could be higher than first order accuracy. Since then, the development of hybrid high order methods that give rapid convergence in smooth flows, while also computing locally first order treatments of shock discontinuities, have been the goal of many. The first practical methods to localise and manage Godunov's theorem were developed in a series of papers by van Leer in the 1970s ( \\cite{vanleerII, vanleerIII, vanleerIV, vanleerV} ). Godunov's method of solving the Riemann gas problem was combined with the development of adaptive piecewise grid reconstructions, were the order of interpolation was reduced in the vicinity of shock discontinuities, but allowed to be higher than first order in areas of smooth flow. Conservative fluxes were computed from the solution of inter--cell Riemann problems and used to update the cell quantities. A key development in 1984 as the use of a conservative 4th order integral polynomial intercell boundary reconstruction, which resulted in the widely used PPM, piecewise parabolic method of Collella \\& Woodward ( \\citet{cw84, cw84b}). The PPM method is extremely efficient for the order of accuracy achieved, up to 3rd order in spatial reconstruction in smooth flows on a regular grid. Apart from higher order inter--cell reconstruction, work progressed on solving the non-linear Riemann problem more quickly, which was especially slow before hardware support for transcendental functions became widespread. One approach was to linearise the Riemann problem, where an approximate Riemann problem is solved exactly -- resulting in the Roe--Pike approximate solvers in the mid 1980s (see for example \\cite{roe1981, roe1986}), which became very widely used. Other less common approaches included simplifying the iterations when solving the full non-linear problem, giving approximate solutions to the exact Riemann problem \\citep{toro99}. Full, exact solvers of the full non-linear Riemann problem were generally abandoned. In the 1980s the issues surrounding the control of `post-shock' ringing, or generally non-monotonic behaviour of any method higher than first order were explored and refined with the development of TVD (Total Variation Diminishing) \\citep{harten83, harten84} and later ENO/WENO (Essentially Non-Oscillating and Weighted Essentially Non-Oscillating) methods \\citep{liu1994,jiang1996}. The PPM method continues to be effective, despite its somewhat ad-hoc monotonicity constraints compared to the more formal TVD flux limiting methods, due to it's great speed and high order interpolation which -- has a particularly simple form on a regular grid due to exploited symmetries in the polynomial construction. Since the 1990s, interest in the Astrophysical regime has been to extend the microphysics ( Cooling, dust, ionisation, magnetic fields etc) and there are a large number of CFD codes targeting different aspects of astrophysical fluids. With this diversity of applications, the key issues have become similarly diverse, and applying a single code to all problems is not practical. Here a new code has been developed, to provide a means of computing astrophysical flow problems with more sophisticated microphysics than previous codes, aimed at providing more reliable observational quantities from the simulations, such as H$\\alpha$ surface brightness, for example. In this work, we present the astrophysical hydrodynamics code, \\Fyris, with a focus on those aspects which have been developed specifically for the new code, along with a comprehensive suite of existing and new test problems to verify the new code. Verification by such test suites is essential, as has been well described in the Lenska \\& Wendroff review \\citep{liska03}, hereafter LW03, where they showed that even well established codes can have particular weaknesses that are not apparent from a simple description of their algorithms. LW03 presented a set of 1D and 2D test problems, many with known analytical solutions, and derived robust L1-norm measures of the deviations from the known solutions, providing an objective measure of the complete code performance. Here we apply the LW03 tests to \\Fyris, and present additional 2D and 3D tests to cover aspects either not fully covered in LW03, or not covered at all. In particular there is a need for analytical 3D test problems for testing code, as large scale 3D simulations are becoming widely used. Paper II will cover the non-fluid dynamics aspects of the code, such the time-dependent ionisation, along with aspects of the magneto-hydrodynamic treatment, in the code {\\em Fyris Beta}. ", "conclusions": "\\Fyris\\ is a new implementation of an established hydrodynamics methodology. The key new features are; \\begin{itemize} \\item The significant performance improvement, up to a factor of four times faster for the same adiabatic 2D calculations. \\item A new two-shock Riemann solver, based on the \\cite{gg88} solver, which is accurate and faster than the equivalent two shock solver of \\cite{cw84}. \\item Robust and accurate solutions to test problems demonstrate that \\Fyris\\ is competitive with other hydrodynamic codes, and in many cases produces better results with smaller errors. \\item Fast and accurate microphysics, including cooling and generalised equation of state allowing for a variable polytropic index, $\\gamma$, and variable mean molecular weight $\\mu$. \\item Active floating point `noise' control, controlling roundoff errors and floating point representation bit errors that would otherwise affect coordinates and allow non-physical changes to the fluid quantities. \\item Fully parallelized, both in shared memory using pre-emptive threading and in distributed mode using the MPI library ({\\small \\tt http://www.mpi-forum.org/docs/docs.html}), or a combination of both. \\end{itemize} \\subsection { Code Availability } Source code, animations and figures for all the problems presented here, along with some additional qualitative problems are available at the website. {\\small \\tt http://www.mso.anu.edu.au/fyris/} \\subsection{ Acknowledgements} \\Fyris\\ is approximately 84000 lines of C code, and was begun during a study visit to UAO, supported by the Centre for Complex Dynamics, at the \\AA ngstr\\\"om Labotoriet, Uppsala University, Sweden, and the code is named after the Fyris river that flows through that town. The author would like to thank the Centre for the financial support during the visit. \\clearpage" }, "1003/1003.5876_arXiv.txt": { "abstract": "We present spectrophotometric data from 0.4 to 4.2 $\\mu$m for bright, northern sky, Be stars and several other types of massive stars. Our goal is to use these data with ongoing, high angular resolution, interferometric observations to model the density structure and sky orientation of the gas surrounding these stars. We also present a montage of the H$\\alpha$ and near-infrared emission lines that form in Be star disks. We find that a simplified measurement of the IR excess flux appears to be correlated with the strength of emission lines from high level transitions of hydrogen. This suggests that the near-IR continuum and upper level line fluxes both form in the inner part of the disk, close to the star. ", "introduction": "% The observed absolute flux from an astronomical source (after correction for telluric and interstellar extinction) is directly related to its emitted flux and angular size in the sky. As we enter the era of optical long-baseline interferometry, it will become easier to measure the angular dimensions of many objects and, consequently, to explore the relationship between the observed and emitted flux distributions. This effort is especially important to determine effective temperatures of stars, but it also plays a key role in the interpretation of circumstellar environments, in particular the disks surrounding Be stars and the winds and outflows of massive stars. Be stars are rapidly rotating B-type stars that manage to eject gas into a circumstellar disk (observed in H emission lines, an infrared flux excess, and linear polarization; \\citealt{por03}). The IR flux excess from the disk results from bound-free and free-free emission from ionized gas, and this emission increases with wavelength, so that in the near and mid-IR the disk flux will dominate over the stellar flux. Models of the IR excess can relate the observations to the disk radial density function \\citep{wat86,dou94,por99}. Such models are also required to interpret recent near-IR interferometric observations of Be stars where the ratio of disk to stellar flux is a key parameter \\citep{ste01,gie07,mei07,car09}. However, Be star disks are intrinsically variable on timescales of months to years \\citep{hub98,por03,mcs08}, so it is necessary to obtain contemporaneous spectrophotometry in order to model both the total flux and its angular distribution in the sky. There are many emission lines of H and He in the near-IR spectra of Be stars \\citep{cla00,stl01,len2a,men09}, and they offer additional diagnostics of the disk density, temperature, and geometry \\citep{hon00,len2b,jon09}. We have embarked on a number of programs of interferometry with the Georgia State University Center for High Angular Resolution Astronomy (CHARA) Array, a six-telescope, optical/IR interferometer with baselines up to 330 m \\citep{ten05}. Here we describe a program of complementary optical and near-IR spectrophotometry of our targets that we will use in detailed modeling of the source angular flux distribution. The observations and their calibration are described in \\S2, and we present figures of the target spectral energy distributions and emission line strengths in \\S3 and \\S4, respectively. In \\S5 we discuss the relationship between the disk continuum and line emission of Be stars, and we present a summary of the work in \\S6. ", "conclusions": "% Our spectrophotometric observations of nearby Be stars show that all the stars with strong H$\\alpha$ emission also display an IR excess relative to the expected photospheric flux distribution. The size of the IR excess is correlated with the H$\\alpha$ equivalent-width but the relation shows the largest scatter among those stars with the densest and largest circumstellar disks. On the other hand, the IR excess shows a better correlation with the equivalent-widths (corrected for disk continuum emission) of high excitation transitions like Hu14. Since only a trace number of H atoms populate these excited states, transitions like Hu14 have a low opacity except in the densest parts of the disk. We argue that these results can be understood in terms of the spatial range in radius over which any emission mechanism is optically thick. The good correlation between the IR continuum emission and the high excitation line emission suggests that both form in the inner, dense part of the disk, while the less marked correlation between the IR continuum and H$\\alpha$ emission results from changes in the density distribution in the outer part of the disk (perhaps due to the temporal evolution of the disk and/or the tidal influence of a binary companion). We are currently making near-IR interferometric observations of some 20 northern Be stars with the CHARA Array. We will combine the SED data presented here with the interferometric visibility data to develop consistent models of the disk density structure and orientation in the sky (see our first examples in \\citealt{gie07})." }, "1003/1003.4794_arXiv.txt": { "abstract": "This paper presents a detailed analysis of temporal and spectral variability of the low energy peaked BL Lac object S5~0716+714 with a long ($\\sim 74$~ks) X-ray observation performed by \\xmm\\ on 2007 September 24--25. The source experiences recurrent flares on timescales of hours. The soft X-ray variations, up to a factor of $\\sim 4$, are much stronger than the hard X-ray variations. With higher energy, the variability amplitude increases in the soft X-rays but decreases in the hard X-rays. The hard X-ray variability amplitude, however, is effectively large. For the first time, we detect a soft lag of $\\sim1000$~s between the soft and hard X-ray variations. The soft lags might become larger with larger energy differences. The overall X-ray spectra exhibit a softer-when-brighter trend, whereas the soft X-ray spectra appear to show a harder-when-brighter trend. The concave X-ray spectra of the source can be interpreted as the sum of the high energy tail of the synchrotron emission, dominating in the soft X-rays, and the low energy end of the inverse Compton (IC) emission, contributing more in the hard X-rays. The synchrotron spectra are steep ($\\Gamma\\sim2.6$), while the IC spectra are flat ($\\Gamma\\sim1.2$). The synchrotron spectra appear to harden with larger synchrotron fluxes, while the IC spectra seem to soften with larger IC fluxes. When the source brightens, the synchrotron fluxes increase but the IC fluxes decrease. The synchrotron tail exhibits larger flux variations but smaller spectral changes than the IC component does. The crossing energies between the two components and the trough energies of spectral energy distributions (SEDs) increase when the source brightens. The X-ray spectral variability demonstrates that the synchrotron and IC SED peaks of S5~0716+714 shift to higher energies when it brightens. The temporal variability also elucidates that the hard X-ray variations of the source might be dominated by the synchrotron tail. The simultaneous optical and UV data obtained with \\xmm\\ are compared with the X-ray observations. ", "introduction": "\\label{sec:intro} Blazars, an assembly of BL Lac objects and flat spectrum radio quasars (FSRQs), are the most extreme class of Active Galactic Nuclei (AGN). They are remarkably characterized by variability of different amplitude on various timescales across most of the electromagnetic wavelengths (e.g., Ulrich et al. 1997). The radiation from blazars is thought to originate in a relativistic jet closely aligned with the line of sight, which is thus relativistically beamed (e.g., Urry \\& Padovani 1995). In the $\\log(\\nu F_{\\nu})-\\log(\\nu)$ representation, the spectral energy distributions (SEDs) of blazars comprise of two broad bumps. The low energy bump peaks at the frequencies ranging from sub-millimeter to X-ray bands, while the peak of the high energy bump is thought to be at the MeV-TeV gamma-ray bands though it has not been well-known for most of blazars yet. The low energy component is widely believed to be the synchrotron emission of a relativistic electron population residing in the jet tangled with magnetic field, whereas the high energy component is thought to be the Inverse Compton (IC) radiation of the same electron population, scattering the low energy photons of either its own synchrotron photons (the synchrotron self-Compton [SSC, e.g., Maraschi et al. 1992] model mostly adopted for BL Lac objects) or the external photons of surrounding environment (the External Compton [EC, e.g., Sikora et al. 1994] model mostly used for FSRQs). The current classification of blazars is preferably based on the peak energy of the synchrotron emission component. BL Lac objects are differentiated into high energy peaked BL Lac objects (HBLs) and low energy-peaked BL Lac objects (LBLs) (Giommi \\& Padovani 1994; Padovani \\& Giommi 1995). The synchrotron emission of HBLs and LBLs peaks at the UV/X-ray and the IR-optical wavelengths, respectively. The synchrotron peak of FSRQs might shift down to lower (e.g., sub-millimeter) frequencies. The location of the synchrotron peak is thought to be related with the bolometric luminosity of the source: the higher the luminosity, the lower the peak energy (Fossati et al. 1998; Ghisellini et al. 1998). This is the so-called blazar SED sequence widely cited in the blazar studies. Historically, LBLs and HBLs are best observed and studied in the optical and X-ray bands, respectively, because LBLs are usually the brightest and most variable objects in the optical wavelengths, whereas HBLs are the brightest and most variable ones in the X-rays. The optical variability properties of the LBL BL Lacertae are analogous to the X-ray variability properties of bright HBLs (Papadakis et al. 2003; Bian et al. 2007). This similarity is anticipated, since the optical emission of LBLs and the X-ray emission of HBLs correspond to the synchrotron peak of their own, respectively. Because it originates in a different dominance of the synchrotron over the IC radiation, the X-ray emission of different classes of blazars shows very different characteristics of temporal and spectral variability (see Pian 2002 and Zhang 2003 for reviews). HBLs are bright and best studied X-ray sources. Their X-ray emission is commonly interpreted as the synchrotron radiation from the high-energy tail of an electron distribution, which is sensitive to the particle acceleration and cooling and thus shows rapid and strong variability (e.g., Kirk et al. 1998). Repeated X-ray observations have been performed for a few X-ray bright HBLs with various X-ray telescopes. Although it is complicated, the temporal and spectral variability of HBLs might be characterized as follows. The X-ray fluxes of HBLs exhibit large amplitude variability on different timescales. Rapid variations are common, e.g., the fluxes of PKS~2155--304 changed by a factor of $\\sim2$ on timescales of the order of a few hours (Sembay et al. 1993; Zhang et al. 1999). The variability amplitude increases with higher energy (Mrk~421: Ravasio et al. 2004; PKS~2155--304: Zhang et al. 2005, 2006b; Mrk~510: Gliozzi et al. 2006), in accordance with the fact that the higher energy electrons have shorter cooling timescales. The X-ray spectra of HBLs are steep ($\\Gamma > 2$) and show a convex shape. More precisely, the X-ray spectra continuously steepen with energies (e.g., Perman et al. 2005; Tramacere et al. 2007). Flux changes are frequently accompanied by spectral variability. The spectra typically harden with higher fluxes (e.g., Mrk~421: Brinkmann et al. 2005; PKS~2155-304: Sembay et al. 2002), i.e., the so-called harder-when-brighter phenomenon. The synchrotron peaks have been determined to locate in the X-ray band for a few bright HBLs, which shift to higher energies as the source brightens (e.g., Mrk~421: Fossati et al. 2000; Massaro et al. 2004a; Tanihata et al. 2004; Tramacere et al. 2009; Mrk~501: Massaro et al. 2004b). The variations between different energies are well correlated within the X-ray band, often with time lags. The lags, when detected, appear to change from flare to flare (e.g., Mrk~421: Takahashi et al. 2000; PKS~2155--304: Zhang et al. 2002; 2006a). The variations at lower energies often lag behind those at higher energies (soft lag), although the opposite behavior (hard lag) has been also claimed (Mrk~421: Zhang 2002; Ravasio et al. 2004; Brinkmann et al. 2003; PKS 2155--305: Zhang et al. 2006a; 1ES~1218+304: Sato et al. 2008). With the \\xmm\\ PN timing mode X-ray observations of Mrk~421, Brinkmann et al. (2005) claimed that both the sign and the length of lags may change on timescales of a few thousand seconds. The lags are also energy dependent, increasing with larger energy differences (e.g., Mrk~421: Takahashi et al. 1996; Zhang 2002; Zhang et al. 2010; PKS~2155--304: Kataoka et al. 2000; Zhang et al. 2002). Spectral variability may correspond to the existence of the inter-band time lags. Spectral evolution with flux has been systematically resolved for some well-defined flares. On the spectral index--flux ($\\alpha-F$) plot, clockwise and counter-clockwise loops have been effectively observed to correspond to the soft and hard lag (e.g., Mrk~421: Takahashi et al. 1996; Zhang 2002, 2010; Ravasio et al. 2004; PKS~2155--304: Kataoka et al. 2000; Zhang et al. 2002), respectively. The sign of the lags and the pattern of the loops have been thought to be the results of the balance between the acceleration and cooling timescales of the electron population responsible for the observed X-ray emission (Kirk et al. 1998; Zhang et al. 2002). This in turn has been used to constrain some of the physical parameters (the magnetic field and the Doppler factor) of the emitting region. The X-ray spectra of LBLs have a concave shape with a break energy of $\\sim2$~keV, which has been clearly detected in a few bright objects with the synchrotron emission peaking in the optical band (e.g., BL Lacertae: Tanihata et al. 2000; Ravasio et al. 2002, 2003; S5~0716+714: Wagner et al. 1996; ON~231: Tagliaferri et al. 2000). The spectra are steep ($\\Gamma \\sim 2.3$) in the soft X-rays, while they are flat ($\\Gamma \\sim 1.7$) in the hard X-rays. On timescale of hours, the soft X-ray fluxes are highly variable, whereas the hard X-ray fluxes tend to be stable (e.g., Tanihata et al. 2000; Ravasio et al. 2003; Giommi et al. 1999; Tagliaferri et al. 2000). Therefore, LBLs shows different spectral and temporal properties in the soft and hard X-rays, which is thought to be due to different emission components. The soft X-rays are dominated by the synchrotron emission from the high-energy tail of an electron distribution, whereas the IC scattering of the low-energy side of the same electron distribution off the synchrotron (and/or external in some cases) photons contributes more in the hard X-rays (e.g., Giommi et al. 1999). The high energy tail of the synchrotron emission is variable and has a steep spectrum. The low energy side of the IC emission is stable and its spectrum is flat. Accordingly, the studies of LBLs in the X-rays are able to understand the physical conditions of the low- and high-energy electrons simultaneously. By fitting the X-ray spectra with the double power law model, Tanihata et al. (2000) and Ferrero et al. (2006) disentangled the two emission components for BL Lacertae and S5~0716+714. The synchrotron emission is variable on short timescales (e.g., $\\sim$~hours), whereas the IC emission appears to vary on longer timescales (e.g., $\\sim$~days). The breaking energies become higher with higher fluxes (e.g., BL Lacertae: Ravasio et al. 2003). The overall X-ray spectra show the softer-when-brighter behaviour (e.g., Giommi et al. 1999. These phenomena are ascribed to the upshift of the synchrotron peak to higher energy when the sources brighten. No inter-band time lags or the spectral index--flux loops have been clearly reported so far for LBLs, although they are theoretically expected (B${\\rm \\ddot{o}}$ttcher \\& Chiang 2002). S5~0716+714 is identified as a prototype of LBLs since its synchrotron emission peaks in the optical band (e.g., Nieppola et al. 2006). The photometric detection of its host galaxy suggested a redshift of $0.31\\pm0.08$ (Nilsson et al. 2008). It has been intensively observed and studied in many wavelengths, especially in the optical band (e.g., Wu et al. 2007 and reference therein). It is strongly variable from radio to X-ray bands on different timescales (e.g., Wagner et al. 1996). The EGRET onboard the Compton Gamma-ray Observatory detected its high energy gamma-ray ($>100$~MeV) emission several times from 1991 to 1996 (Hartman et al. 1999). The gamma-ray fluxes varied by a factor of two on timescale of years, but the spectral index--flux correlation was not found (Nandikotkur et al. 2007). In 2008, AGILE detected a variable gamma-ray flux, with a peak flux above the maximum obtained by EGRET (Chen et al. 2008). Observations by HEGRA resulted in an upper limit of flux at very high energy (VHE) gamma-ray energies ($>1.6$~TeV) (Aharonian et al. 2004). It is also in the first Fermi-LAT bright source list (Abdo et al. 2009). Recently, the MAGIC collaboration reported the first detection of VHE gamma-rays from the source at $5.8\\sigma$ level (Anderhub et al. 2009). The discovery of S5~0716+714 as a VHE gamma-ray LBL was triggered by its very high optical state, suggesting a possible correlation between the VHE gamma-ray and the optical emissions. S5~0716+714 is also a preferred target to perform simultaneous multi-wavelength observations (e.g., Villata et al. 2008; Giommi et al. 2008). Wu et al. (2009) reported the first detection of time lags between the variations in different optical wavelengths. S5~0716+71 has been observed by various X-ray telescopes. The observations in 1991 March with the PSPC onboard ROSAT revealed a behavior of flux-related spectral variations in the 0.1--2.4~keV band. Two distinct spectral components are needed to describe the concave X-ray spectra (Cappi et al. 1994). ASCA observed the source in 1994 March and confirmed the spectral shape found by ROSAT in higher energies: the spectra flatten with increasing energies (Kubo et al. 1998). The three \\sax\\ observations in 1996, 1998 and 2000 revealed the spectral and temporal variability of the two spectral components. It was in faint states in the 1996 and 1998 observation (Giommi et al. 1999). The spectral fits with a broken power law model resulted in concave spectra in the 0.1--10~keV band, breaking at $\\sim 2-3$~keV. The 2000 observation caught the source in its high state (Tagliaferri et al. 2003). The concave spectra can be disentangled into two power law components, crossing at $\\sim1.5$~keV. The steeper power law component dominates the soft X-ray emission, whereas the flatter one contributes more to the hard X-rays. The soft X-ray spectral index is the largest out of the three \\sax\\ observations, indicating a softer-when-brighter behaviour for the soft X-ray variability. The \\sax\\ observations also revealed large and rapid variations in the soft X-ray band and the lack of variations above $\\sim3$~keV. In the 1996 observation, a flare with duration of $\\sim20$~hours was detected in the soft but not in the hard X-ray band (Giommi et al. 1999). Variations by a factor of $\\sim3$ on a timescale of one hour was detected below $\\sim3$~keV and no variations above (Tagliaferri et al. 2003). Nevertheless, the comparisons between the three observations show that the hard X-ray fluxes were also variable over the timescales of years. The temporal and spectral variability of S5~0716+714 is in good agreement with the interpretation that the soft and hard X-ray emission, separating at $\\sim2-3$~keV, are dominated by the strongly variable high energy tail of the synchrotron radiation and the \"stable\" low energy side of the IC radiation, respectively. The softer-when-brighter behaviour, opposite to the harder-when-brighter one frequently observed in HBLs, could be ascribed by the upshift of the synchrotron peak to higher energy as the source brightens, causing the synchrotron tail to enter into the soft X-ray band more (e.g., Giommi et al. 1999). However, previous X-ray observations were not able to disentangle the two spectral components and related variations on short timescales. \\xmm\\ pointed to S5~0716+714 twice, on 2004 April 4--5 and 2007 September 24--25 (Orbit 791 and 1427, PI: G. Tagliaferri), lasting for $\\sim59$~ks and $\\sim74$~ks, respectively. The first \\xmm\\ observation was analyzed by Ferrero et al. (2006, hereafter FE06; see also Foschini et al. [2006]), who studied the temporal and spectral variability of the synchrotron and IC component on timescales of hours. S5~0716+714 was in a higher state compared to previous observations. Strong variations with flux changes by a factor of more than three were detected on timescale of hours. The variability amplitude is significantly higher in the soft than in the hard X-ray band. The soft X-ray variations correlate with the harder X-ray ones, but no pronounced time lags were found. The time-resolved spectral analysis with a double power law model showed that the crossing points of the synchrotron and IC emission move to higher energies with increasing fluxes. Both of the two components vary on timescales of hours. The synchrotron emission becomes more dominant when the source brightens, following a harder-when-brighter trend as HBLs. The IC emission exhibits a more complicated variability behaviour. In this paper, we present in detail the temporal and spectral variability of the synchrotron and IC emission for the second \\xmm\\ observation of S5~0716+714. Some new variability properties are found. Our results will be compared with those of the first \\xmm\\ observation and the earlier observations by previous X-ray telescopes. The X-ray observations are also compared with the simultaneous optical and UV data obtained with \\xmm. ", "conclusions": "\\label{sec:disc} We perform a detailed spectral and temporal analysis for the second \\xmm\\ observation of S5~0716+714. Most of our results are in agreement with previous results obtained from the first \\xmm\\ observation (FE06) and other X-ray observations (e.g., Giommi et al. 1999; Tagliaferri et al. 2003). Nevertheless, we also find some new phenomena, adding new clues to better understand the underlying physical processes taking place in the source. The concave X-ray spectra of S5~0716+714 can be disentangled into two power law components. The steep power law ($\\Gamma \\sim 2.6$) component is interpreted as the high energy tail of the synchrotron emission, whereas the flat power law ($\\Gamma \\sim1.2$) component is ascribed to the low energy side of the IC emission. FE06 obtained similar results with the first \\xmm\\ observation. It is worth noting that the photon indices of the steep power law component are similar to those of HBLs in the hard X-ray band (e.g., PKS~2155--304: Zhang 2008), supporting the interpretation of the steep power law component as the synchrotron tail. The X-ray variability amplitude of HBLs monotonically increases with higher energy (e.g., PKS~2155--304: Zhang et al. 2005; Mrk~421: Zhang et al. 2010), which is thought to be the signature of synchrotron emission. However, LBLs are highly variable in the soft X-rays, whereas they show little variability in the hard X-rays (e.g., Giommi et al. 1999). Our results demonstrate that S5~0716+714 is indeed strongly variable in the soft X-rays, showing the maximum variability by a factor of $\\sim4$ throughout the whole observation and several episodes of rapid variations on timescales of hours. In a sharp-cut contrast, the hard X-ray fluxes of the source are much less variable, exhibiting only $\\sim50\\%$ change between the minimum and maximum count rates and no rapid events. For the first time, we quantify the energy dependence of the variability amplitude for S5~0716+714. The variability amplitude increases from the 0.3--0.5 to 0.5--7.5~keV band, but from the 0.75--1~keV band, it decreases with higher energy. Moreover, the OM data suggest lower variability amplitude in the UV band than in the soft X-ray band. The energy dependence of the variability amplitude of S5~0716+714 is thus clearly different from those of HBLs. During the first \\xmm\\ observation, the \\fvar\\ of S5~0716+714 amounts to $40\\pm3\\%$ in the 0.5--0.75~keV and $27\\pm1\\%$ in the 3--10~keV (FE06). Though longer duration of the second \\xmm\\ observation, the \\fvar\\ is only $32.03\\pm2.12\\%$ and $7.62\\pm1.5\\%$ in the corresponding energy bands. It is important to understand whether the synchrotron or IC emissions or both are responsible for the significant variations of S5~0716+714 in different energy bands. It is usually argued that the variable high energy tail of the synchrotron emission is responsible for the highly variable soft X-ray emission, whereas the \"stable\" low energy side of the IC emission accounts for the lack of variability in the hard X-ray band (e.g., Giommi et al. 1999; Tagliaferri et al. 2003). This argument is somewhat ambiguous, since the two \\xmm\\ observations clearly demonstrate that the hard X-ray fluxes are also highly variable, although their variability amplitudes are significantly smaller than the soft X-ray ones. The hard X-ray variations could be due to either the IC or synchrotron variations. It is certain that the synchrotron component is strongly variable on short timescales, whereas it is unclear whether the IC component is variable on similar timescales. The spectral analysis by FE06 showed that the IC component might be variable on timescales of hours, which is, however, not supported by our results. Although the IC component contributes more to the total hard X-ray fluxes, we still assume that the hard X-ray variations might be controlled by the synchrotron tail. With increasing energies, the increasing dilutions of the \"stable\" IC component to the increasing variations of the synchrotron tail, might result in the observed overturn of the energy-dependent variability amplitude for S5~0716+714. The model-independent hardness-ratio analysis shows that the 0.5--10~keV spectra of S5~0716+714 soften when it brightens. The phenomenon, also noticeable in the first \\xmm\\ observation (FE06; Foschini et al. 2006), was already found in the \\sax\\ observations of the source (Giommi et al. 1999). The softer-when-brighter phenomenon is interpreted in terms of the high energy tail of the synchrotron emission extending more to higher energies when the source brightens (e.g., Giommi et al. 1999). Nevertheless, the soft X-ray spectra appear to harden when it brightens, though the trend is not significant. The variability property found in the soft X-ray band of the source with ROSAT observations (Cappi et al. 1994) is similar to what we found here. The harder-when-brighter trend is analogous to those of HBLs, presenting a strong evidence that the soft X-ray emission of S5~0716+714 is dominated by the synchrotron tail. Although the inter-band time lags and the related energy dependence have been detected in a few X-ray bright HBLs, they have not be firmly detected in LBLs yet. FE06 claimed that the lags of $\\gtrsim100$~s were not present in the first \\xmm\\ observation of S5~0716+714. However, we found that the 0.3--0.5~keV variations lag the 3--10~keV ones by $\\sim 1000$~s in the second \\xmm\\ observation. We also found a weak evidence that the lags increase with larger energy differences. In at least one episode of the optical-UV light curves, the U band variations might lag the X-ray variations by $\\sim 2000$~s. As far as we know, it is the first evidence for a definite detection of soft lag in the X-ray variability of S5~0716+714 and possibly LBLs. Interestingly, the soft lags and the related energy dependence of S5~0716+714 are similar to what have been detected in HBLs (e.g., Kataoka et al. 2000; Zhang et al. 2010). The similarity suggests that the X-ray variations of the source have the same origin as those of HBLs, i.e., the variations of the synchrotron tail. Therefore, the hard X-ray fluxes of the source are dominated by the IC component, but the hard X-ray variations might be still controlled by the synchrotron tail, as already suggested by the energy dependent variability amplitude. If we assume that both the soft and hard X-ray variations are caused by the synchrotron tail, the observed soft lags provide a way to constrain the physical parameters of the emitting region (e.g., Zhang \\et 2002), \\begin{equation} B\\delta^{1/3} =209.91 \\times \\left (\\frac{1+z}{E_{\\rm l}}\\right)^{1/3} \\left [\\frac{1 - (E_{\\rm l}/E_{\\rm h})^{1/2}} {\\tau_{\\rm soft}} \\right ]^{2/3} \\quad {\\rm G} \\,, \\label{eq:soft} \\end{equation} where $\\tau_{\\rm soft}$ is the observed soft lag (in second) between the low ($E_{\\rm l}$) and high ($E_{\\rm h}$) energy (in keV), $z$ the source's redshift, $B$ the magnetic field (in G) and $\\delta$ the bulk Doppler factor of the emitting region. If adopting $\\tau_{\\rm soft} \\sim 1000$~s between the 0.3--0.5 and 3--10~keV, one gets $B\\delta^{1/3} \\sim 2.56$~G. During a model fit to the SED involving the first \\xmm\\ observation of S5~0716+714, Foschini et al. (2006) assumed $B=3$~G and $\\delta =16.7$, i.e., $B\\delta^{1/3} = 7.67$~G. A larger $B\\delta^{1/3}$ implies a smaller $\\tau_{\\rm soft}$, qualitatively consistent with the soft lags of no larger than 100~s claimed by FE06. The unprecedented high signal-to-noise ratio PN observation allows us to disentangle the synchrotron and IC components in the X-ray band of S5~0716+714 and to synchronously study the variations of the two components on timescales of hours. The results, obtained with the divisions of the observation by individual episodes and by count rate levels, are consistent with each other. The synchrotron photon indices are constrained in a limited range of $\\Gamma \\sim 2.5-2.7$, which are typical of HBLs in the hard X-rays (e.g., PKS~2155--304: Zhang 2008). The IC photon indices show relatively large changes ($\\Gamma \\sim 0.9-1.4$). Due to large uncertainties, the synchrotron photon indices might be consistent with each other within the error bars, which could be also true for the IC photon indices. Although the synchrotron and IC photon indices do not correlate with the total fluxes, the synchrotron spectra appear to harden with higher synchrotron fluxes, and the IC spectra seem to soften with higher IC fluxes. When the source brightens, the synchrotron fluxes increase, while the IC fluxes decrease. The synchrotron fluxes also show larger variations than the IC fluxes. The X-ray SEDs convolved with the synchrotron and IC components exhibit significant concave shapes. The crossing energies and the SED trough energies increase with the increasing total fluxes. The flux dependence of the crossing energies and the SED trough energies is consistent with that of the first \\xmm\\ observation (FE06). We further notice that the SED trough energies are smaller than the crossing energies in most cases, indicating that they are not the exact energies at which the synchrotron component transits to the IC component in terms of the equal contributions of the two components to the total fluxes. The SED evolution, characterized by the shifts of the SED troughs to higher energies with higher fluxes, might be mainly caused by the changes of the synchrotron normalization. The changes of the synchrotron and IC photon indices and of the IC normalization may affect the SED evolution in a weaker way. The spectral and temporal behaviors of S5~0716+714 and its synchrotron and IC emission are just the consequence of the peaks of both the synchrotron and IC SEDs shifting to higher energies with increasing fluxes. When the source brightens, the synchrotron peak moves to higher energy. In turn, the high energy tail of the synchrotron emission extends to higher energy. The synchrotron peak is therefore more close to the observed X-ray band, bring on that the synchrotron flux increases and the synchrotron spectrum hardens. At the same time, the IC peak also shifts to higher energy, incurring that the low energy end of the IC emission recedes from lower energy to the observed X-ray band. The IC peak is thus more far from the observed X-ray band. As a result, the IC flux decreases and the IC spectrum hardens. Accordingly, the synchrotron flux anti-correlates with the IC flux when the source brightens. The series of changes make the SED trough move to higher energy with higher total flux. The high energy tail of the synchrotron emission originates from the high energy electrons, showing strong and rapid variations. The low energy end of the IC emission comes from the low energy electrons, exhibiting small and slow variations. Soft lag is expected if both the soft and hard X-ray variations are caused by the cooling of the high energy electrons. In conclusions, S5~0716+714 exhibits different X-ray variability properties between the two \\xmm\\ observations. During the second observation, it shows harder synchrotron and IC spectra and lower variability amplitude. Even though the low energy end of the IC emission contributes more to the hard X-ray fluxes than the high energy tail of the synchrotron emission does, the energy dependence of the variability amplitude suggests that the hard X-ray variations might be dominated by the synchrotron tail. The large hard X-ray variability amplitude suggests its synchrotron origin as well. More importantly, soft lags are detected for the first time, also supporting that the soft and hard X-ray variations are caused by the same mechanism. When the the source brightens, the synchrotron fluxes increase but the IC fluxes decrease. The synchrotron spectra might harden with higher synchrotron fluxes, while the IC spectra might soften with larger IC fluxes. The synchrotron tail shows larger flux variations but smaller spectral variations than the IC emission does. With higher fluxes, the crossing points between the two components and the SED troughs move to higher energies. The X-ray variability of S5~0716+714 is in accordance with the results of its synchrotron and IC SED peaks shifting to higher energies with higher fluxes. The decompositions of its X-ray emission into the synchrotron and IC components are helpful to understand the behaviours of both the low and high energy part of the electron population." }, "1003/1003.1949_arXiv.txt": { "abstract": "{} {To separate stars and galaxies in the far infrared AKARI All-Sky Survey data, we have selected a sample with the complete color information available in the low extinction regions of the sky and constructed color-color plots for these data. We looked for the method to separate stars and galaxies using the color information. } {We performed an extensive search for the counterparts of these selected All-Sky Survey objects in the NED and SIMBAD databases. Among 5176 selected objects, we found {4272} galaxies, {382 other extragalactic objects,} {349 Milky Way stars, 50 other Galactic objects,} and {101} sources detected before in various wavelengths but of an unknown origin. {22} sources were left unidentified. Then, we checked colors of stars and galaxies in the far-infrared flux-color and color-color plots. } {In the resulting diagrams, stars form two clearly separated clouds. One of them is easy to be distinguished from galaxies and allows for a simple method of excluding a large part of stars using the far-infrared data. The other {smaller} branch, overplotting galaxies, {consists of stars known to have an infrared excess, like Vega and some fainter stars discovered by IRAS or 2MASS. The color properties of these objects in any case make them very difficult to distinguish from galaxies.} } {We conclude that the FIR color-color diagrams allow for a high-quality star-galaxy separation. With the proposed simple method we can select more that 95~\\% of galaxies rejecting {at least} 80~\\% of stars.} ", "introduction": "Emission from galaxies at wavelengths beyond a few micrometers is produced mainly by dust. As widely known, star formation activity is always accompanied by dust production, { though the actual mechanism of dust supply is not very well understood yet. } The energy emitted by massive stars at ultraviolet (UV) is scattered efficiently and finally absorbed by dust grains. The heated grains re-emit the energy at far-infrared (FIR). Thus, radiative processes in FIR are related to the composition and the amount of dust in galaxies as well as to the properties of their stellar population, especially newly formed massive stars. Many studies have shown that significant amount of star formation in galaxies is obscured by dust \\citep{lefloch05,caputi07,buat07a,buat07b,buat08,reddy08}. Especially, such ``hidden'' star formation is revealed to become more and more important with increasing redshift from $z = 0$ to 1 \\citep{takeuchi05a}. Therefore, FIR observations have crucial importance to understand the {\\it true} star formation activity in the Universe. Any FIR surveys would first provide us with a point source catalog, and extragalactic astrophysicists may want to have a reliable list of galaxy candidates from the FIR catalog. However, usually it is not a trivial task to classify detected sources and extract a specific class of objects only from FIR information. Obviously not only galaxies would be detected at FIR, but also many other kinds of Galactic objects would be included in { a purely flux-limited catalog which is obtained as a first product of the survey. } Very crudely one may cut out the Galactic plane to avoid Galactic star-forming regions or H{\\sc ii} regions, but still there are numerous stars with dust, e.g., Vega-like stars and AGBs. An effective and convenient method for star-galaxy separation which makes use of FIR information only would be desired. Such a method would also be very useful to construct a good candidate list of dusty stars. The Infrared Astronomical Satellite \\citep[IRAS; ][]{neugebauer84} has brought a vast amount of statistics and very efficient methods have been invented from IRAS Point Source Catalog (PSC). The four bands of IRAS enabled us to perform even a very detailed classification of extragalactic and various galactic objects, like blue and red galaxies, Seyferts and QSOs, carbon stars, H{\\sc ii} regions, reflection nebulae, planetary nebulae, T Tauri stars etc. A {thorough} description of the method from IRAS can be found in \\citep[e.g., ][]{walker89}. After many years since IRAS, the advent of AKARI (ASTRO-F) opened a new window to explore the Universe, as a survey-oriented space telescope {at MIR and FIR} \\citep{murakami07}. The primary purpose of the mission is to provide second-generation infrared (IR) catalogs to obtain a better spatial resolution and a wider spectral coverage than the IRAS catalog. AKARI is equipped with a cryogenically cooled telescope of 68.5~cm aperture diameter and two scientific instruments, the Far-Infrared Surveyor \\citep[FIS; ][]{kawada07} and the Infrared Camera \\citep[IRC; ][]{onaka07}. Among various astronomical observations performed by AKARI, an all sky survey with FIS {and IRC} has been carried out (AKARI All-Sky Survey). Since FIS is an instrument dedicated to FIR $\\lambda = 50 \\mbox{--} 180\\;\\mu$m, all the AKARI FIS bands are in the FIR wavelengths: {\\it N60} ($65\\;\\mu$m), {\\it WIDE-S} ($90\\;\\mu$m), {\\it WIDE-L} ($140\\;\\mu$m), and {\\it N160} ($160\\; \\mu$m) \\citep{kawada07}. Hereafter, we use a notation $S_{65}$, $S_{90}$, $S_{140}$ and $S_{160}$ for flux densities in these bands. Especially, since FIS has sensitivity at longer wavelengths than IRAS, a new method of classification scheme is needed if we try to select a list of a certain class of objects. Such a scheme is not merely an empirical technique but also will provide us with a new understanding of objects with cool dust which were difficult to detect by IRAS bands. Since our central interest is on the physics of IR galaxies, we set our main aim to select galaxies by fluxes at four AKARI FIS bands. In this paper, we present the first color-color and color-flux diagrams obtained by the All-Sky Survey data only based on FIS bands, and show the method of star-galaxy separation with these diagrams. The paper is organized as follows: in Section~\\ref{sec:data}, we present the data: the sample with the complete color information, In Sections~\\ref{sec:flux_color} and \\ref{sec:color_color}, we present the FIR flux-color and color-color diagrams of the data, respectively, and we show how galaxies and stars can be separated in them. We show our results and conclusions in Section~\\ref{sec:conclusion}. ", "conclusions": "\\label{sec:conclusion} \\begin{enumerate} \\item The far infrared flux-color diagrams show two clear branches of sources. One of them is populated almost only by stars, the second one is dominated by galaxies. \\item The far infrared color-color diagrams allow for a high quality star-galaxy separation. In all the combinations of color-color diagrams we can distinguish two separate clouds. It was shown that one of them contains in all cases more than 9{5}~\\% of galaxies and the other one {in most cases} consists in more than {8}0~\\% of stars. \\item At a first glance the population of stars is divided into (roughly) two categories: a {larger group} occupies a separate area on the diagrams while the {remaining smaller part} resides in the ``galaxy cloud''. A more careful examination has revealed that stars occupying a separate area on the color-color diagram are {mostly} known bright stars. { This group is dominated by pulsating variable stars, among them many Mira-type variables, and contains also a number of evolved stars.} {Branch II of stars, overlapping the galaxy cloud, includes stars discovered by the infrared surveys, bright stars known to have an infrared excess, like Vega, pulsating stars, a few T Tauri stars and faint red stars.} We estimate that the percentage of galaxies intervening the stellar part of the diagram is in each case less than 6~\\%, and the percentage of stars in the galaxy part of the diagram - {in most cases} less than 20~\\%. {A few exceptions can be probably attributed to some systematic error in the flux estimation, in particular in the $N160$ band.} { Obviously, since this analysis was made in the sky regions with low-Galaxtic cirrus emission, all the detailed numbers shown above are applicable for the regions except low-Galactic latitude zone. } \\item We conclude that the FIR color-color diagrams are a powerful tool to separate stars from galaxies in AKARI FIS All-Sky Survey, and further, similar large FIR catalogs. With the presented simple method we can select more than 95~\\% of galaxies, rejecting 80~\\% of stars. A more detailed methodology will be a subject of future works. \\end{enumerate}" }, "1003/1003.2369.txt": { "abstract": "We present CO observations of nine ULIRGs at $z$\\,$\\sim$\\,2 with $f_\\nu(24\\mu m)$\\,$\\simgt$\\,1\\,mJy, previously confirmed with the mid-IR spectra in the {\\sl Spitzer} First Look Survey. All targets are required to have accurate redshifts from Keck/GEMINI near-IR spectra. %The analyses of mid-IR spectra and SEDs suggest that dust obscured AGNs dominate the mid-IR luminosities at (3\\,--\\,6)\\um\\ for most of our objects. Using the Plateau de Bure millimeter-wave Interferometer (PdBI) at IRAM, we detect CO J(3-2) [7 objects] or J(2-1) [1 object] line emission from eight sources with integrated intensities $I_{c}$\\,$\\sim$\\,(5\\,--\\,9)$\\sigma$. %The CO detectability is not strongly correlated with mid-IR spectral features such as PAH emission and silicate absorption. The CO detected sources have a variety of mid-IR spectra, including strong PAH, deep silicate absorption and power-law continuum, implying that these molecular gas rich objects at $z$\\,$\\sim$\\,2 could be either starbursts or dust obscured AGNs. The measured line luminosity $\\rm L^{'}_{CO(3-2)}$ is (1.28\\,--\\,3.77)$\\times$10$^{10}$\\,K\\,km/s\\,pc$^2$. The averaged molecular gas mass $\\rm M_{H_2}$ is $1.7\\times$10$^{10}M_\\odot$, assuming CO-to-H$_2$ conversion factor of 0.8\\,$M_\\odot$/[K\\,km/s\\,pc$^2$]. %Although the statistics is small, 24\\um\\ ULIRGs may have on average less cold molecular gas than SMGs and gas rich QSOs, and have probably narrower CO velocity width than that of SMGs, but similar to that of QSOs. %the averaged cold molecular mass may be %Assuming CO-to-H$_2$ conversion factor of 0.8\\,$M_\\odot$/[K\\,km/s\\,pc$^2$] and $T_{CO(3-2)}/T_{CO(1-0)}$ of 1, as what used for SMGs and QSOs, we infer the cold molecular gas masses of (1\\,--\\,3)$\\times$10$^{10}M_\\odot$ with $\\rm \\langle M_{gas}\\rangle$\\,=\\,1.7$\\times$10$^{10}M_\\odot$, suggesting probably a factor of $2$ smaller than that of sub-millimeter galaxies (SMGs) and QSOs with CO detections. Our sample has a CO velocity width distribution similar to that of QSOs, but its median velocity width (406\\,km/s) is a factor of 1.5 smaller than that of SMGs, although the statistics is small. %Two sources in our sample are spatially resolved, with MIPS506 having two discrete components separated by 45\\,kpc and MIPS16144 extending across 50\\,kpc (beam $\\sim$\\,20\\,kpc$\\times$18\\,kpc). Three sources (33\\%) -- MIPS506, MIPS16144 \\&\\ MIPS8342 -- have double peak velocity profiles. The CO double peaks in MIPS506 and MIPS16144 show spatial separations of 45\\,kpc and 10.9\\,kpc, allowing the estimates of the dynamical masses of %MIPS506 and MIPS16144 have two velocity peaks separated by $\\sim$\\,200\\,km/s and in space by 10.9\\,kpc and 45\\,kpc respectively. The inferred, corresponding dynamical masses are 3.2$\\times$10$^{11}$\\,$\\rm{sin}^{-2}(i)\\,M_\\odot$ and 5.4$\\times$10$^{11}$\\,$\\rm{sin}^{-2}(i)\\,M_\\odot$ respectively. The implied gas fraction, $\\rm M_{gas}/M_{dyn}$, is 3\\%\\ and 4\\%, assuming an average inclination angle. %implying the gas fraction, $\\rm M_{gas}/M_{dyn}$, of 5\\%\\ and 6\\%\\ respectively assuming $\\rm\\langle sin^{2}(i)\\rangle$\\,=\\,2/3. %the gas fraction, $\\rm M_{gas}/M_{dyn}$, is 12\\%\\ for both systems, much smaller than the average values of $\\sim$(30\\,--\\,40)\\%\\ observed for SMGs and gas rich, high-$z$ QSOs. Finally, the analysis of the HST/NICMOS images, mid-IR spectra and IR SED revealed that most of our sources are mergers, containing dust obscured AGNs dominating the luminosities at (3\\,--\\,6)\\um. Together, these results provide some evidence suggesting SMGs, bright 24\\um\\ ULIRGs and QSOs could represent three different stages of a single evolutionary sequence, however, a complete physical model would require much more data, especially high spatial resolution spectroscopy. ", "introduction": " ", "conclusions": "" }, "1003/1003.1722_arXiv.txt": { "abstract": "We present a unified framework for the study of late time cosmic acceleration. Using methods of effective field theory, we show that existing proposals for late time acceleration can be subsumed in a single framework, rather than many compartmentalized theories. We construct the most general action consistent with symmetry principles, derive the background and perturbation evolution equations, and demonstrate that for special choices of our parameters we can reproduce results already existing in the literature. Lastly, we lay the foundation for future work placing phenomenological constraints on the parameters of the effective theory. Although in this paper we focus on late time acceleration, our construction also generalizes the effective field theory of inflation to the scalar-tensor and multi-field case for perturbatively constructed backgrounds. ", "introduction": "The observation of cosmic acceleration has provided a fundamental quandary for particle physics. The acceleration could be due to a tiny but non-zero vacuum energy density (a cosmological constant), but then the fine-tuning problem between the tiny cosmological constant and the Planck scale must be addressed, either via anthropic arguments or some new dynamics. It is also interesting to consider, however, that the observed cosmic acceleration may be arising from new dynamics in the gravity sector. Experimental results suggest that at solar system scales our world is accurately described by General Relativity (GR) \\cite{Damour:2008zz}, which is the unique Lorentz invariant theory of spin-2 gravitons at low energies \\cite{Weinberg:1965nx}. In order to explain the observation that the universe is accelerating on much larger scales, without invoking a cosmological constant, requires that we either add new degrees of freedom in the stress energy tensor, or that we alter the structure of general relativity (more precisely the graviton propagator). In the former case, the new dynamics enters through an additional scalar degree of freedom, for which a new hierarchy problem between the Planck scale and the scalar mass arises that must be addressed. In the latter case, consistent modifications to GR lead to an additional scalar degree of freedom as well. As we will elaborate on below, one finds that the dynamics has an effective description as a scalar-tensor theory for the range of scales relevant for observations. To date, most of the theoretical effort on dynamical dark energy models have been phenomenological in nature. This includes Quintessence~\\cite{Quintessence}, K-essence~\\cite{ArmendarizPicon:2000dh}, Brans-Dicke theories~\\cite{bransdicke}, modifications of the Ricci scalar ($f(R)$ models)~\\cite{fofr}, first order phase transitions~\\cite{Stojkovic:2007dw}, spontaneous violation of Lorentz invariance (ghost condensation)~\\cite{ArkaniHamed:2003uy} and the Dvali-Gabadadze-Porrati (DGP) model~\\cite{dgp}. Though on the surface, each appears to be a modification or addition to a different part of GR, all of these models can be re-written as scalar-tensor theories in a different regime. For example, Quintessence is a scalar-tensor theory where the scalar enters the action only through its potential and interactions (and so reduces to ordinary GR plus a scalar sector), while modified gravity theories come with an additional coupling to matter in the Einstein frame (or curvature in the Jordan frame). Thus far, the dominant approach in the literature has been to choose one model and analyze its effects on the expansion history and the matter power spectrum. Little attempt has been made to unify these models into one fundamental, theoretical framework. With a single framework in place, one can analyze the constraints from data on the coefficients of the terms in the effective action, and more generally determine which classes of models are consistent with the observations and more clearly see for which reasons. Such an approach should simplify and unify the phenomenological analysis of models of dark energy. Given the generality of our approach, we will see that many limiting cases of our action have already appeared in the literature in one form or another. At first it might seem hopeless to construct a single action that could account for so many different ideas to address cosmic acceleration. However, as long as we restrict our attention phenomenologically to the long distance (low energy) regime where the theory is valid, we are able to subsume many theories with different ultraviolet behavior into one theory. An analogous situation arises in particle physics, where many different theories for electroweak symmetry breaking in the ultraviolet describe the same low energy phenomenology of the weak and electromagnetic interactions. To obtain the effective theory describing our low energy world, one simply writes down all the lowest order operators consistent with the symmetries of the theory. Our goal here is to carry out this procedure to describe the low energy phenomenon of cosmic acceleration. In doing this, we unify all cosmic acceleration models (and modified gravity alternatives) using the approach of effective field theory (EFT). We write down the lowest order corrections to the scalar-tensor theory and show that with these corrections all the models of dark energy described above can be reproduced. After some work and re-writing of the lowest order terms in the action (which we discuss in more detail later), the result is remarkably simple\\footnote{Throughout this paper we will work with the metric signature mostly plus ($-,+,+,+$), and with natural units $\\hbar = c =1$. For a full list of our notation and conventions we refer the reader to the appendix.}: \\bea \\label{masteraction} S &=& \\int\\d^4 x\\sqrt{-g} \\Big\\{ \\frac{1}{16\\pi G_N}\\Omega^2(\\vp)R-\\frac{1}{2} Z(\\varphi) g^{\\mu\\nu}\\pt_\\mu\\vp\\pt_\\nu\\vp-U(\\vp) + \\frac{\\alpha(\\vp)}{\\Lambda^4} (g^{\\mu\\nu}\\pt_\\mu\\vp\\pt_\\nu\\vp)^2 \\Big\\} + S_m . \\;\\; \\eea We then proceed to do a perturbation analysis to determine cosmologically relevant parameters for the evolution of structure formation, such as the anisotropic stress. We show that our analysis reproduces the existing results in the literature. We leave for a second paper a more concrete application to phenomenological analysis of dark energy models, which will be useful for constraining them. The outline of our paper is as follows. In the next section we lay out details for arriving at the effective action (\\ref{masteraction}) and make some comparisons and connections with models already existing in the literature. In the following section, we connect (\\ref{masteraction}) to observations by deriving observable quantities, such as the expansion history, the anisotropic stress and the effective Newton constant. We also discuss how our action again reduces in special cases to existing ideas and models of dark energy. We then conclude and outline future directions for the application of our framework to data analysis and constraints on dark energy models. To avoid obscuring the presentation, we leave a number of technical results to the appendices. These include the full derivation of the action (\\ref{masteraction}), as well as the resulting equations of motion at both the background level and for cosmological perturbations. Our choice of conventions are summarized in Appendix \\ref{appendix0}. ", "conclusions": "In this paper we have derived a formalism for combining existing proposals for dark energy and modified gravity into a single framework. We find that the effective field theory analysis of scalar-tensor theory provides a compelling framework to accomplish such a task. We have constructed the most general, local and unitary action, and calculated the leading corrections to fourth order in derivatives for all of the fields. We find that we can reproduce the results of many different models of dark energy with this framework. We have seen that the effective theory (\\ref{masteraction}) can capture the crucial physics of modified gravity theories for intermediate scales between the horizon and the solar system. However, there remain several challenges to obtaining an even more complete discussion. It would be interesting to see if our approach could be pushed to super-horizon scales, where measurements like those of the cosmic microwave background could allow further constraints. This has been done through a more phenomenological approach -- the so-called PPF formalism of \\cite{Hu:2009ua}, however it is not clear whether the approach we have taken here can be pushed to such scales in a similar fashion. The PPF formalism makes the crucial assumption that in models like DGP, above the horizon where the model becomes intrinsically five dimensional, if four dimensional energy/momentum conservation holds one can still proceed. If one takes {\\em all} fields into consideration, this seems a reasonable assumption, even though a given sector (like that of the scalar $\\varphi$) will not be separately conserved. We leave extending our methods to super-Horizon scales to a future publication \\cite{us}. A more challenging problem is that of small scales, as perturbations evolve into the non-linear regime. This offers an important test of these theories, especially in the case of modified gravity, for it is in this regime that one expects the strong coupling of the scalar to reduce to GR. Moreover, we have focused on cosmological applications, but in the presence of clustered and dense objects, we expect that something like the Vainshtein effect \\cite{Vainshtein:1972sx} could become relevant and our effective theory will break down. This breakdown is not due to the the strong coupling at the scale $\\Lambda_{UV}$, but is instead a purely classical effect and could be crucial in understanding how effective field theories in the case of modified gravity can be connected to GR in the appropriate limit. Within the regime of validity of our EFT, below the horizon size but above the length scale of the solar system, we have argued that these theories are appropriately described by GR plus a non-minimally coupled scalar field. This fact was already well known in the case of DGP and $f(R)$ theories, but we argued here that this can be extended to all other consistent modifications as well. Within our EFT framework, we have seen that strong observational constraints can be placed on these theories. These constraints come from observations that restrict the expansion history and the growth of structure. We saw that the choice of parameters in the effective theory can determine whether growth can be either faster or slower depending on the value of the effective Newton constant. In addition, modified gravity theories result in a non-negligible anisotropic stress, which, by combining measures of structure growth with \\eg~weak lensing, can be used to distinguish these models from those of dark energy. Such an approach can be carried out in a similar way to that accomplished in the literature for the special cases of DGP and $f(R)$ (see \\eg~\\cite{Hu:2009ua} and references within). We plan to continue this approach in a future publication \\cite{us}, to demonstrate generally which are the features of a theory of modified gravity necessary to produce the observed acceleration and structure growth." }, "1003/1003.5721_arXiv.txt": { "abstract": "We report one of the most accurate measurements of the three-dimensional large-scale galaxy power spectrum achieved to date, using $56{,}159$ redshifts of bright emission-line galaxies at effective redshift $z \\approx 0.6$ from the WiggleZ Dark Energy Survey at the Anglo-Australian Telescope. We describe in detail how we construct the survey selection function allowing for the varying target completeness and redshift completeness. We measure the total power with an accuracy of approximately $5\\%$ in wavenumber bands of $\\Delta k = 0.01 \\, h$ Mpc$^{-1}$. A model power spectrum including non-linear corrections, combined with a linear galaxy bias factor and a simple model for redshift-space distortions, provides a good fit to our data for scales $k < 0.4 \\, h$ Mpc$^{-1}$. The large-scale shape of the power spectrum is consistent with the best-fitting matter and baryon densities determined by observations of the Cosmic Microwave Background radiation. By splitting the power spectrum measurement as a function of tangential and radial wavenumbers we delineate the characteristic imprint of peculiar velocities. We use these to determine the growth rate of structure as a function of redshift in the range $0.4 < z < 0.8$, including a data point at $z=0.78$ with an accuracy of $20\\%$. Our growth rate measurements are a close match to the self-consistent prediction of the $\\Lambda$CDM model. The WiggleZ Survey data will allow a wide range of investigations into the cosmological model, cosmic expansion and growth history, topology of cosmic structure, and Gaussianity of the initial conditions. Our calculation of the survey selection function will be released at a future date via our website {\\tt wigglez.swin.edu.au}. ", "introduction": "\\renewcommand{\\thefootnote}{\\fnsymbol{footnote}} \\setcounter{footnote}{1} \\footnotetext{E-mail: cblake@astro.swin.edu.au} The pattern of density fluctuations in the low-redshift Universe results from the physical processes which govern the evolution of matter perturbations after the Big Bang. In the early Universe, the primordial spectrum of fluctuations created by inflation is processed before recombination in a manner depending on the physical matter density, baryon fraction and massive neutrino fraction (e.g.\\ Bond \\& Efstathiou 1984; Bardeen et al.\\ 1986; Holtzman 1989; Hu \\& Sugiyama 1996; Eisenstein \\& Hu 1998). After recombination, perturbations of all scales are amplified by gravity at an identical rate whilst linear theory applies. This growth rate depends on the matter and dark energy components which drive the cosmic expansion (e.g.\\ Heath 1977; Hamilton 2001; Linder \\& Jenkins 2003; Percival 2005). The growth of fluctuations enters a non-linear regime at progressively larger scales at lower redshifts: in today's Universe, only perturbations with Fourier wavescales $k < 0.1 \\, h$ Mpc$^{-1}$ evolve linearly to a good approximation (e.g.\\ Smith et al.\\ 2003; Jeong \\& Komatsu 2006; McDonald 2007). The clustering pattern of galaxies at different redshifts is related to the underlying density fluctuations and may be used to test this model of structure formation. The shape of the clustering power spectrum -- the relative amplitudes of large-scale and small-scale modes -- depends on the composition of the early Universe and may be used to extract information about the matter and baryon fractions (e.g.\\ Tegmark et al.\\ 2004a; Cole et al.\\ 2005; Percival et al.\\ 2007b). The amplitude of the clustering power spectrum as a function of redshift, together with the pattern of redshift-space distortions induced by galaxy peculiar velocities, can be used to measure the growth rate of structure (e.g.\\ Hamilton 1992; Hawkins et al.\\ 2003; Guzzo et al.\\ 2008; Percival \\& White 2009). Higher-order or topological descriptors of the density field, such as the bispectrum or genus, can be applied to test whether the initial conditions are consistent with scale-invariant Gaussian random perturbations generated by inflation (e.g.\\ Gott, Dickinson \\& Melott 1986; Fry \\& Scherrer 1994; Sefusatti \\& Komatsu 2007; James, Lewis \\& Colless 2007). The interpretation of the shape and amplitude of the galaxy power spectrum is complicated by several factors. Firstly, the manner in which galaxies trace the density field -- the ``galaxy bias'' -- is in general a complex function of scale, dark matter halo mass, galaxy type and redshift (Dekel \\& Lahav 1999; Tegmark \\& Bromley 1999; Wild et al.\\ 2005; Conway et al.\\ 2005; Percival et al.\\ 2007a; Smith, Scoccimarro \\& Sheth 2007; Cresswell \\& Percival 2009). However, the bias of galaxy fluctuations on sufficiently large scales ($k < 0.1 \\, h$ Mpc$^{-1}$ at $z=0$) appears to be well-described by a simple constant of proportionality whose value depends on galaxy type and luminosity, or more fundamentally dark matter halo mass (Peacock \\& Dodds 1994; Scherrer \\& Weinberg 1998; Verde et al.\\ 2002). Secondly, small-scale density perturbations eventually begin to evolve in a non-linear fashion requiring more complex modelling techniques such as higher-order perturbation theory or numerical $N$-body simulations (Smith et al.\\ 2003; Jeong \\& Komatsu 2006; McDonald 2007). Thirdly, there is a practical challenge of acquiring galaxy survey data across a ``fair sample'' of the Universe (Tegmark 1997). For the large-scale linear modes of clustering, which provide the most robust link to underlying theory, this sample must map a volume of the order 1 Gpc$^3$ using of the order $10^5$ galaxies. These demands require multi-year campaigns with ground-based telescopes utilizing hundreds of clear nights (Glazebrook \\& Blake 2005). Despite these challenges, a series of galaxy redshift surveys have been undertaken to provide such datasets at redshifts $z < 0.5$. The state-of-the-art projects which have mapped the ``local'' ($z \\approx 0.1$) Universe are the 2-degree Field Galaxy Redshift Survey (2dFGRS; Colless et al.\\ 2001) and the Sloan Digital Sky Survey (SDSS; York et al.\\ 2000). The 2dFGRS obtained redshifts for $2 \\times 10^5$ galaxies covering 1500 deg$^2$ in the period between 1997 and 2002. The ``main'' spectroscopic survey of the SDSS gathered $8 \\times 10^5$ galaxy redshifts over 8000 deg$^2$ between the years 2000 and 2005. The SDSS project also included observations of $1 \\times 10^5$ Luminous Red Galaxies (LRGs) reaching up to a redshift $z = 0.5$ (Eisenstein et al.\\ 2001). These datasets have provided a rich source of information about the clustering of galaxies. For example, power spectra have been extracted for the 2dFGRS by Percival et al.\\ (2001) and Cole et al.\\ (2005); for the SDSS ``main'' galaxy sample by Pope et al.\\ (2004), Tegmark et al.\\ (2004a) and Percival et al.\\ (2007a); and for the LRGs by Eisenstein et al.\\ (2005), Huetsi (2006), Tegmark et al.\\ (2006) and Percival et al.\\ (2007b). Analysis of these surveys, in combination with the Cosmic Microwave Background fluctuations, has confirmed that we inhabit a low-density Universe where matter today provides only $25-30\\%$ of the total energy governing the large-scale dynamics, with the rest located in a mysterious ``dark energy'' component. In addition the baryonic fraction of the matter is only $15-20\\%$, with the remainder composed of non-baryonic, cold particles whose nature is currently unknown (e.g.\\ Percival et al.\\ 2002; Tegmark et al.\\ 2004b; Tegmark et al.\\ 2006; Komatsu et al.\\ 2009). The clustering pattern is also sensitive to the presence of hot dark matter such as massive neutrinos, which comprise a small fraction of the energy budget (Elgaroy et al.\\ 2002; Seljak et al.\\ 2005). These galaxy surveys also describe how the underlying density fluctuations are modulated by galaxy bias (Verde et al.\\ 2002; Wild et al.\\ 2005; Conway et al.\\ 2005; Percival et al.\\ 2007a; Cresswell \\& Percival 2009). In this context the comparison of power spectrum measurements from the 2dFGRS and SDSS, which targeted galaxy populations selected in blue and red optical wavebands respectively, is of particular interest. When the differing galaxy types in these surveys are assigned linear bias factors, the resulting model fits to the linear-regime power spectra produce best-fitting matter densities which are inconsistent at the statistical level of $2\\sigma$. Careful treatment of scale-dependent and luminosity-dependent galaxy bias can potentially explain this discrepancy (Percival et al.\\ 2007a; Sanchez \\& Cole 2008). There are strong motivations for extending these large-scale structure measurements to higher redshifts ($z > 0.5$). Firstly, the growth of structure implies that the linear regime of evolving perturbations extends to smaller scales at higher redshifts, enabling cleaner and more accurate model fits. Secondly, the shape of the survey cone allows access to significantly greater cosmic volumes at higher redshift, enabling more accurate determinations of the large-scale power spectrum amplitude. Thirdly, baryon oscillations in galaxy power spectra at different redshifts may be used as a standard ruler to extract the cosmic distance-redshift relation and infer the properties of dark energy (Blake \\& Glazebrook 2003; Seo \\& Eisenstein 2003; Hu \\& Haiman 2003). Fourthly, measurements of the growth of cosmic structure as a function of redshift increases our ability to discriminate between dark energy models including modifications to Einstein's theory of gravity (Guzzo et al.\\ 2008; Wang 2008; White, Song \\& Percival 2009). Our current tools for probing the matter power spectrum at redshifts $z > 0.5$ are limited. The clustering of high-redshift quasars has been studied by the 2dF Quasar Survey (Outram et al.\\ 2003) and the SDSS (Ross et al.\\ 2009) but the scarcity of QSOs implies that the large-scale clustering measurements are strongly limited by shot noise. Photometric redshifts from imaging surveys have been used to study the projected clustering pattern in redshift slices (Blake et al.\\ 2007; Padmanabhan et al.\\ 2007). However, this approach loses the information from small-scale radial clustering modes (Blake \\& Bridle 2005) and in particular prevents the extraction of the patterns of peculiar velocities, which are swamped by photometric redshift errors. Alternatively, fluctuations in the Lyman-$\\alpha$ forest absorption spectrum on the sight lines to bright quasars have been used to infer the amplitude of small-scale clustering fluctuations in the high-redshift Universe (Croft et al.\\ 2002; McDonald et al.\\ 2005; McDonald et al.\\ 2006). However, this method is potentially susceptible to systematic modelling errors and is only applicable at redshifts $z > 2.3$ where the Lyman-$\\alpha$ absorption lines pass into optical wavebands. The WiggleZ Dark Energy Survey at the Anglo-Australian Telescope (Drinkwater et al.\\ 2010) will provide the next step forwards in large-scale spectroscopic galaxy redshift surveys, mapping a cosmic volume of the order 1 Gpc$^3$ over the redshift range $z < 1$. The survey, which began in August 2006 and is scheduled to finish in July 2010, is obtaining of the order $200{,}000$ redshifts for UV-selected emission-line galaxies covering of the order 1000 deg$^2$ of equatorial sky. The principal scientific goal is to measure baryon oscillations in the galaxy power spectrum in redshift bins up to $z = 1$ and provide a robust measurement of the dark energy model. The dataset will also trace the density field over unprecedented cosmic volumes at $z > 0.5$, providing a sample comparable to the SDSS LRG catalogue at $z < 0.5$. Moreover, the spatial overlap between the WiggleZ and LRG catalogues in the redshift range $0.3 < z < 0.5$ will allow careful studies of the systematic effects of galaxy bias on power spectrum estimation. This paper presents a determination of the current WiggleZ survey selection function and galaxy power spectrum, using a dataset comprising of the order $25\\%$ of the final survey observations. The selection function, which describes the angular and radial survey coverage in the absence of clustering, is complicated by the relatively high level of incompleteness in the survey affecting both the parent target catalogues and the spectroscopic follow-up observations (although this latter type of incompleteness will decrease as the survey progresses). However, we demonstrate that despite these complications the galaxy clustering power spectrum may be successfully extracted and already provides accurate tests of the cosmological model that rival lower-redshift surveys. The structure of this paper is as follows: in Section \\ref{secsel} we present a detailed account of the survey selection function including the coverage masks, completeness of our UV imaging catalogues, variations in redshift completeness, redshift distribution as a function of sky position, and redshift blunder rate. In Section \\ref{secpowerspec} we describe our power spectrum calculation and its correction for redshift blunders. We compare the predictions of cosmological models to the resulting power spectra in Section \\ref{secparfit} and itemize our conclusions in Section \\ref{secconc}. We note that we construct our selection function for a fiducial flat $\\Lambda$CDM cosmological model with matter density $\\Omega_{\\rm m} = 0.3$. ", "conclusions": "\\label{secconc} In this paper we have described our method of determining the selection function of the WiggleZ Dark Energy Survey, and have presented the current measurement of the large-scale galaxy power spectrum using $56{,}159$ redshifts of bright emission-line galaxies spanning redshifts $0.3 < z < 0.9$. This sample constitutes approximately $25\\%$ of the final WiggleZ survey. We have quantified and categorized the redshift blunder rate and determined its effect on the power spectrum measurement via analytical calculations and detailed simulations. We conclude that: \\begin{itemize} \\item The selection function of the WiggleZ survey is complicated by the proximity of the faint magnitude threshold to the completeness limit of the input catalogues, in particular for the GALEX UV data. We quantified the incompleteness in the parent target catalogue as a function of GALEX exposure time and Galactic extinction via fitting formulae. \\item We adopted a Monte Carlo technique to determine the relative completeness of the spectroscopic follow-up at any position. This technique allows for the complex overlapping of survey pointings and for the systematic variation of redshift completeness across the 2-degree field-of-view of the instrument. We also allowed for the magnitude prioritization of the spectroscopic follow-up which results in a position-dependent galaxy redshift distribution. \\item The WiggleZ survey contains redshift blunders resulting from emission-line confusion (most significantly, [OIII], H$\\beta$ and H$\\alpha$ mis-identified as [OII]) and from sky emission lines mis-identified as [OII]. The overall blunder rate is about $5\\%$. The effect of the redshift blunders on the power spectrum measurement is well-approximated as a constant reduction in amplitude for scales $k > 0.05 \\, h$ Mpc$^{-1}$ combined with an enhanced level of reduction for large scales $k < 0.05 \\, h$ Mpc$^{-1}$. \\item We measured 1D (angle-averaged) and 2D (binned in tangential and radial modes) galaxy power spectra for nine independent survey regions and redshift slices using the method of Feldman, Kaiser \\& Peacock (1994). The 1D power spectrum for the whole sample, combining these measurements, has a fractional accuracy of about $5\\%$ in Fourier bins of width $\\Delta k = 0.01 \\, h$ Mpc$^{-1}$. The 2D power spectra show the expected anisotropic signatures of redshift-space distortions due to large-scale coherent infall and small-scale virialized motions. \\item The power spectrum data are well-described by a model power spectrum with matter and baryon densities consistent with those determined from observations of the Cosmic Microwave Background radiation. The model includes non-linear corrections, redshift-space distortions and a linear galaxy bias factor. \\item The 2D power spectra allow us to measure the growth rate of structure across the redshift range $0.4 < z < 0.8$. We obtain results similar in precision to previous determinations at $z < 0.4$, including a measurement at $z=0.78$ with $20\\%$ accuracy. \\end{itemize} Future studies will present full cosmological parameter fits and combinations of these results with other datasets, including implications for the growth of cosmic structure and Gaussianity of the initial conditions, and extend these analyses to the final WiggleZ survey catalogues." }, "1003/1003.0618_arXiv.txt": { "abstract": "In order to understand the formation mechanism of the disks around Be stars it is imperative to have a good overview of both the differences and similarities between normal B stars and the Be stars. In this paper we address binarity of the stars. In particular, we investigate a previous report that there may be a large population of sub-arcsecond companions to Be stars. These had hitherto not been detected due to a combination of a limited dynamic range and spatial resolution in previous studies. We present the first systematic, comparative imaging study of the binary properties of matched samples of B and Be stars observed using the same equipment. We obtained high angular resolution (0.07-0.1 arcsec) {\\it K} band Adaptive Optics data of 40 B stars and 39 Be stars. The separations that can be probed range from 0.1 to 8 arcsec (corresponding to 20-1000 AU), and magnitude differences up to 10 magnitudes can in principle be covered. After excluding a few visual binaries that are located in regions of high stellar density, we detect 11 binaries out of 37 Be targets (corresponding to a binary fraction of 30 $\\pm$ 8 \\%) and 10 binaries out of 36 B targets (29 $\\pm$ 8\\%). Further tests demonstrate that the B and Be binary systems are not only similar in frequency but also remarkably similar in terms of binary separations, flux differences and mass ratios. The minimum physical separations probed in this study are of order 20 AU, which, combined with the similar binary fractions, indicates that any hypotheses invoking binary companions as responsible for the formation of a disk need the companions to be closer than 20 AU. Close companions are known to affect the circumstellar disks of Be stars, but as not all Be stars have been found to be close binaries, the data suggest that binarity can not be responsible for the Be phenomenon in all Be stars. Finally, the similarities of the binary parameters themselves also shed light on the Be formation mechanism. They strongly suggest that the initial conditions that gave rise to B and Be star formation must, to all intents and purposes, be similar. This in turn indicates that the Be phenomenon is not the result of a different star formation mechanism. ", "introduction": "Be stars are rapidly rotating stars surrounded by an ionized circumstellar disk, whose emission is best known for the distinct doubly peaked H$\\alpha$ lines. About 15\\% of all B stars belong to the Be category \\citep[chap. 9]{jaschek_book}, making them a significant component of the hot, massive star population. However, the origin of the circumstellar disks around Be stars is still an unsolved mystery (see e.g. the review by \\citealt{porter_review}). Binarity has sometimes been proposed to be responsible for the Be phenomenon. For example, \\citet{harmanec_2002} investigate the effects of a detached, secondary component's gravitational field on the mass ejection into the equatorial plane from the primary. \\citet{gies_1998} discuss the rotational history of $\\phi$ Per, and find the (compact) companion has had a significant tidal effect on the central Be star. In addition, it should be noted that companions are also known to affect the disk's properties such as its density law or even size (e.g. \\citealt{jones_2008} on $\\kappa$ Dra). Most models based on binarity have in common that close binaries with separations of order AU, are responsible for the formation of a disk (for an overview, see \\citealt{Neg_2007}). When one considers the more general case of binaries, it is plausible that disks are affected during the star formation stages by the presence of a companion star, even in the case of a wide binary. However, to be able to address binarity as a crucial ingredient for the existence of Be stars, we need to know whether Be stars are found in binary systems more often than normal B stars. If so, this would indicate that binarity must have something to do with the production of the conditions required for disk formation -- either directly, where the tidal forces acting within a (close) binary system result in a disk around a mass losing object, or indirectly, where the influence of the binary companion allows the existing disk to survive. \\begin{figure} \\includegraphics[width=0.4\\textwidth,angle=-90]{./b1892.eps} \\caption{Example of the data. Shown in this 3$\\times$3 arcsec image is the known B-type binary HR 1892. The previously known and catalogued companion is located at a distance of 1.6 arcsec towards the South-West. The close companion at 0.15 arcsec (double the spatial resolution in this image) in the North-West direction is a new detection. \\label{h1892}} \\end{figure} The most straightforward manner to determine whether the Be binary fraction is significantly different from that of ``normal'' B-type stars and thereby to assess whether existing binary models are viable in the first place is to establish whether Be stars are more frequently found in binaries than normal B stars. The most recent comparative study on Be binarity was performed 25 years ago by \\citet{abt_1984}. In this study, incorporating the spectroscopic survey of Be and B stars by \\citealt{abt_1978}, and including both spectroscopic and visual binaries, they found the Be binary fraction to be 30\\%, similar to that of B stars. However, this paper was essentially a compilation of the then, inevitably incomplete, literature on visual binaries and was never repeated by a homogeneous dedicated investigation. Later studies, such as the Speckle interferometric study by \\citet{mason_1997} concentrated on Be stars alone. Despite their high angular resolution, Mason et al. did not discover many new binary systems. This is most likely due to their limited dynamic range -- only companions which were fainter by up to 3 magnitudes than the primary could be detected. The original motivation for the current study was the unexpected result that in an investigation into the spectro-astrometric signatures of the H$\\alpha$ emission lines of a random sample of 8 Be stars, 5 were found to be a binary with separations at the sub-arcsecond to arcsecond scale (\\citealt{baines_thesis}, see also \\citealt{oudmaijer_2008}). The parameter space probed by the spectro-astrometry (from 0.1 arcsec to 1-2 arcsec, up to 6 magnitudes difference, \\citealt{baines_2006}) was not covered in any previous study into the binarity of Be stars (see references above), and observationally this was a completely unconstrained problem. If true, the large fraction of (sub-)arcsecond Be star binaries, would have important implications for our understanding of Be stars. Obviously, if wide binarity could be linked to the Be star phenomenon it would require new physics or new models as the effect from distant companions will be subtle. Alternatively, low number statistics could have mimicked a high binary fraction, whereas, in reality, this may not be the case. A follow-on study with more objects matched by a sample of normal B stars should allow us to decide on the issue. Whereas the technique of spectro-astrometry is very powerful in detecting binaries amongst {\\it emission}-line stars, it is not as sensitive when considering normal, absorption line, B-type stars \\citep{baines_2006}. In order to do a comparative study, we need to employ a technique that is both powerful and applicable to both types of star. Here we present observations for a simple experiment: are Be stars more likely than normal B stars to be in a binary system? To this end, we obtained Adaptive Optics (AO) data of a matched sample of 40 B and Be stars using an 8 meter class telescope. \\begin{figure} \\begin{center} \\includegraphics[width=0.45\\textwidth]{./densfig.ps} \\caption{The number of 2MASS sources per square arcminute around the target objects. A density of 20 arcmin$^{-2}$ corresponds to an average of one background object per field of view of 13.6$\\times$13.6 arcsec$^2$. All targets in fields with source densities exceeding this value have indeed more than one object in the FoV and were initially flagged as having a binary companion, but are not included in the final sample. \\label{densfig}} \\end{center} \\end{figure} ", "conclusions": "We have observed a total of 79 B and Be stars at high spatial resolution to study their binary properties. This is the first systematic comparative and deepest imaging study into the binarity of Be stars published. The observations probe scales down to less than 0.1 arcsecond and allow companions up to 10 magnitudes fainter to be detected. In the process, we derived a simple relation between the {\\it K}-band magnitude difference and the mass-ratio of the binary components under the assumption that both components are on the Main Sequence. The latter is justified in our sample. After excluding a few visual binaries that are located in regions of high stellar density, we find that : \\begin{enumerate} \\item There are 10 detected binaries out of 36 B stars and 11 detected binaries out of 37 Be stars (corresponding to binary fractions of 29 $\\pm$ 8 \\% and 30 $\\pm$ 8\\% respectively). The binary fractions are nearly identical. \\item The separations of the B and Be binary systems follow the same cumulative distributions. \\item The mass ratios for the 9 B and 10 Be binaries for which we have differential photometry have similar distributions. Like for the separations, the samples appear to be drawn from the same parent distribution. \\end{enumerate} \\noindent This leads to our main conclusion : \\begin{itemize} \\item The binary fraction, and the binaries' properties, of Be stars in the separation range 20 - 1000 AU are very similar to those of a matched sample of B stars. This is a statistically robust result and as both samples were observed with the same observational equipment, and therefore were subject to the same observational biases, we can rule out binary companions at separations of 20 AU or more as being responsible for the Be phenomenon. \\end{itemize} Returning to the original motivation of the survey, the previously found large fraction of wide Be binaries has been reduced by a factor of 2 due to improved number statistics. More importantly, the comparison with normal B stars now allows us to firmly conclude that there is no need, based on this study, to explore new models involving wide binaries for the Be phenomenon. Closer binaries than 20 AU can only be found with methods that have even higher spatial resolution than here (diffraction limited imaging at an 8 meter telescope in the {\\it K} band), or with radial velocity studies. A systematic imaging study at even higher resolution (indeed at any resolution) such as this one has not been performed, but the systematic spectroscopic study by \\citep{abt_1978} does not reveal a large percentage of close binaries. Indeed, on the contrary, they find a lower close binary fraction among Be stars than B stars. Therefore, binarity can not be responsible for all Be stars. Finally, a word on the similar mass ratios and binary separations. The mass ratio and separation of a binary system are the result of the initial conditions governing the formation of both components and their subsequent evolution. The stars in our sample are field stars, which means that if they were formed in clusters, as most massive stars do, these clusters have dispersed long since. Given that the separations of the B and Be binaries under consideration are fairly wide, the objects were most likely formed in low density clusters or even in isolation. This is because large separation binaries will be destroyed early in the evolution of dense clusters, while close, more strongly bound binaries are most likely to survive throughout the entire evolution \\citep{kouwenhoven_2009,parker_2009}. It is beyond the scope of this paper to derive the precise cluster conditions, but it would appear that both the evolutionary history and the birthplaces of the B and Be stars are very much the same. It can be concluded from our data, therefore, that the Be phenomenon does not result from a substantially different star formation mechanism or conditions compared to those for normal B stars." }, "1003/1003.0873_arXiv.txt": { "abstract": "The generation of a large recoil velocity from the inspiral and merger of binary black holes represents one of the most exciting results of numerical-relativity calculations. While many aspects of this process have been investigated and explained, the ``antikick'', namely the sudden deceleration after the merger, has not yet found a simple explanation. We show that the antikick can be understood in terms of the radiation from a deformed black hole where the anisotropic curvature distribution on the horizon correlates with the direction and intensity of the recoil. Our analysis is focussed on Robinson-Trautman spacetimes and allows us to measure both the energies and momenta radiated in a gauge-invariant manner. At the same time, this simpler setup provides the qualitative and quantitative features of merging black holes, opening the way to a deeper understanding of the nonlinear dynamics of black-hole spacetimes. ", "introduction": " ", "conclusions": "" }, "1003/1003.4661_arXiv.txt": { "abstract": "Abundances relative to iron for carbon, nitrogen, strontium and barium are presented for 33 stars on the red giant branch of the globular cluster $\\omega$ Centauri. They are based on intermediate-resolution spectroscopic data covering the blue spectral region analyzed using spectrum synthesis techniques. The data reveal the existence of a broad range in the abundances of these elements, and a comparison with similar data for main sequence stars enables insight into the evolutionary history of the cluster. The majority of the red giant branch stars were found to be depleted in carbon, i.e.\\ ${\\rm [C/Fe]} < 0$, while ${\\rm [N/Fe]}$ for the same stars shows a range of $\\sim$1 dex, from ${\\rm [N/Fe]} \\approx 0.7$ to 1.7 dex. The strontium-to-iron abundance ratios varied from solar to mildly enhanced ($0.0 \\leq {\\rm [Sr/Fe]} \\leq 0.8$), with [Ba/Fe] generally equal to or greater than [Sr/Fe]. The carbon and nitrogen abundance ratios for the one known CH star in the sample, ROA 279, are ${\\rm [C/Fe]}=0.6$ and ${\\rm [N/Fe]}=0.5$ dex. Evidence for evolutionary mixing on the red giant branch is found from the fact that the relative carbon abundances on the main sequence are generally higher than those on the red giant branch. However, comparison of the red giant branch and main sequence samples shows that the upper level of nitrogen enhancement is similar in both sets at ${\\rm[N/Fe]} \\approx 2.0$ dex. This is most likely the result of primordial rather than evolutionary mixing processes. One red giant branch star, ROA~276, was found to have Sr and Ba abundance ratios similar to the anomalous Sr-rich main sequence star S2015448. High resolution spectra of ROA~276 were obtained with the Magellan Telescope/{\\sc mike} spectrograph combination to confirm this result, revealing that ROA~276 is indeed an unusual star. For this star calculations of the depletion effect, the potential change in surface abundance that results from the increased depth of the convective envelope as a star moves from the main sequence to the red giant branch, strongly suggest that the observed Sr enhancement in ROA~276 is of primordial origin, rather than originating from a surface accretion event. ", "introduction": "\\label{intro} \\defcitealias{nfm96}{NFM96} \\defcitealias{nd95b}{ND95} \\defcitealias{pan03}{PAN03} \\defcitealias{bw93}{BW93} \\defcitealias{sta06a}{STA06a} \\defcitealias{sta06b}{STA06b} The unusual variation in abundances of stars within the globular cluster $\\omega$ Centauri has been well studied over the last four decades. From the original photographic and photoelectric work of \\citet{woo66} and \\citet{cs73} respectively, and the more recent analysis of CCD photometry by \\citet{sol05}, the spread in color on the red giant branch is apparent with distinct branches visible in the color-magnitude diagram (CMD). The large range in metallicity of over 1 dex has been studied extensively on the red giant branch (RGB) (\\citealt[hereafter NFM96]{nfm96}; \\citealt{sk96,lee99, pan00, rey04, sol05, joh08, cal09}), and to a limited extent on the main sequence and turnoff (MSTO) area \\citep{pio05, sol05, sta06a, kay06, vil07}. Large ranges in abundance for all elements studied in the cluster have been found on the RGB (\\citealt[hereafter ND95]{nd95b}; \\citealt{scl95, smi00, joh08, cal09}). This contrasts with what is found in normal globular clusters where typically small ranges for the light elements (C, N, O, Mg, Al and Na) are seen. The light elements in $\\omega$ Cen stars show large variations for a given {\\feh} (\\citealt{per80}, \\citealt{bw93}, hereafter BW93; \\citetalias{nd95b}). The $\\alpha$ elemental (Mg, Si, Ca and Ti) abundance ratios with respect to iron are largely constant for ${\\rm [Fe/H]} < -1.0$ (\\citetalias{bw93}, \\citealt{scl95}; \\citetalias{nd95b}; \\citealt{smi00}) but [$\\alpha$/Fe] then decreases at metallicities greater than -1.0 \\citep{pan02}. Sodium ([Na/Fe]) and aluminum ([Al/Fe]) abundances are correlated, and both are anticorrelated with oxygen ([O/Fe]) (\\citetalias{bw93}; \\citetalias{nd95b}; \\citealt{nd95b, smi00}). For lower metallicities (below ${\\rm [Fe/H]} < -0.8$) \\citet{smi00} and \\citet{cun02} report constant [Cu/Fe]. Above a metallicity of ${\\rm [Fe/H]} = -1.2$, however, \\citet{pan02} find an increase in [Cu/Fe] as the metallicity increases. Variations in abundance for the neutron-capture elements have also been found. The s-process element abundance ratios, [s/Fe], increase with increasing {\\feh}, but then are constant above metallicities greater than $-1.2$ (\\citetalias{nd95b}, \\citealt{scl95}). The {\\it source} of these abundance variations can be attributed to several different possibilities. The decrease in $\\alpha$-elemental abundances at higher metallicities (${\\rm [Fe/H]} \\geq -1.0$) can be attributed to Type Ia supernovae enrichment. The origin of the s-process enrichment is possibly due to ejecta from asymptotic giant branch (AGB) stars, and the iron peak and metal abundance patterns are consistent with primordial enrichment from Type II supernovae. At least three different {\\it processes} are likely to be involved. The first are the processes that mix material within the stars themselves. These include the first dredge up on the RGB in which the surface layers of the star are mixed with lower processed material \\citep{ibe65}. This can, in principle, account for some of the variations seen within the RGB stars of C, N and O. These processes, however, are unlikely to account for the variations seen in elements heavier than oxygen. Secondly, an enriched object may have accreted matter onto its surface layers from either stellar winds from AGB stars or interstellar material. As the accretion events only affect the surface layers of the star, once it ascends the RGB and the convective envelope deepens, the accreted matter will be mixed with the internal layers, thereby reducing the level of enrichment. A comparison of MS and RGB stars may then show higher abundances for the unevolved MS stars compared to those on the RGB. \\citet{joh09} suggest large binary fractions, and hence accretion of enhanced material, as the source of the 25\\% of their sample with ${\\rm [La/Eu]} \\geq 1$. \\citet{cal09} also suggest the same scenario to explain the $\\sim20\\%$ of CN strong RGB stars identified in their sample. However, this is in contrast to that observed by \\citet{may96} who monitored $\\sim310$ giants in $\\omega$ Cen over a ten year period and found the binary fraction to be as low as 3 -- 4\\%. The third possibility is that the variations have a primordial origin; i.e. the stars formed from gas enriched by a combination of AGB stars, massive stars and/or supernovae. It is likely that more than one of these scenarios is at work. Nevertheless, a comparison between abundances obtained from MS and RGB stars may help to distinguish among the various processes involved in the enrichment of this unusual star cluster. In particular, in the absence of thermohaline mixing (discussed further in \\S5), it can distinguish whether enrichment of s-process and CNO elements is due to surface contamination (which would be diluted by the growing convective envelope as the stars move on to the giant branch) or is instead uniform throughout the stars. An abundance analysis of MSTO stars has already been performed by the present authors \\citep{sta07} for carbon, nitrogen, and strontium. In the present paper we investigate abundances on the RGB in a similar manner. The sample of RGB stars and their observations is described \\S\\ref{rgbs}. In \\S\\ref{sparam} we discuss the stellar parameters and the techniques used to obtain abundances for C, N, Sr and Ba in the stars. The results are presented in \\S\\ref{abund}. In \\S\\ref{disc} we present comparisons between the RGB and MSTO and discuss the consequences. ", "conclusions": "Abundances of C, N and the s-process elements Sr and Ba were determined for a biased sample of 33 RGB stars in {\\wcen}. Almost all objects show depletion of carbon, and solar or enhanced nitrogen. The abundances of Sr and Ba show enhancement as well. One of the known CH stars in {\\wcen} has been analyzed for the first time for carbon and nitrogen, resulting in ${\\rm [C/Fe]} = 0.6$, ${\\rm [N/Fe]} = 0.5$ and ${\\rm [Sr/Fe]} = {\\rm [Ba/Fe]} = 0.6$. This star is less enhanced in carbon compared with other CH stars in the cluster, but still has considerable enhancement in carbon compared with other RGB stars. The levels of N enhancement on both the MS and RGB reach similar relative abundances, ${\\rm [N/Fe]} \\sim 1.8$. This may indicate that mixing on the RGB occurs to different extents for individual stars. A RGB star with high enhancement of the light s-process element Sr was found, but with little enhancement in Ba. This star is similar to the strongly Sr-enhanced MS object, S2015448 \\citep{sta06b} from which the conclusion is reached that the Sr enhancement is likely to be primordial in origin, rather than the result of some accretion event. \\\\ {\\bf Acknowledgments} We wish to thank Don VandenBerg and Leo Girardi for providing unpublished details from their models. Australian access to the Magellan Telescopes was supported through the Major National Research Facilities II program of the Australian Government. We thank the referee for their comments which led to improvements in the manuscript. LMS wishes to thank K. Ward for advice on statistical calculations. {\\it Facilities:} {ANU:SSO2.3m(DBS); Magellan:Clay(MIKE)}" }, "1003/1003.4985_arXiv.txt": { "abstract": "We have discovered recent star formation in the outermost portion (1--4$\\times$ $R_{25}$) of the nearby lenticular (S0) galaxy NGC~404 using GALEX UV imaging. FUV-bright sources are strongly concentrated within the galaxy's HI ring (formed by a merger event according to del Rio et al.), even though the average gas density is dynamically subcritical. Archival HST imaging reveals resolved upper main sequence stars and conclusively demonstrates that the UV light originates from recent star formation activity. We present FUV, NUV radial surface brightness profiles and integrated magnitudes for NGC~404. Within the ring, the average star formation rate surface density ($\\Sigma_{\\rm SFR}$) is $\\sim2.2\\times10^{-5}$ M$_\\odot$~yr$^{-1}$~kpc$^{-2}$. Of the total FUV flux, 70\\% comes from the HI ring which is forming stars at a rate of $2.5\\times10^{-3}$ M$_\\odot$~yr$^{-1}$. The gas consumption timescale, assuming a constant SFR and no gas recycling, is several times the age of the Universe. In the context of the UV-optical galaxy CMD, the presence of the SF HI ring places NGC~404 in the green valley separating the red and blue sequences. The rejuvenated lenticular galaxy has experienced a merger-induced, disk-building excursion away from the red sequence toward bluer colors, where it may evolve quiescently or (if appropriately triggered) experience a burst capable of placing it on the blue/star-forming sequence for up to $\\sim$1 Gyr. The green valley galaxy population is heterogeneous, with most systems transitioning from blue to red but others evolving in the opposite sense due to acquisition of fresh gas through various channels. ", "introduction": "Strateva et al. (2001) used an early SDSS galaxy catalog to demonstrate that the distribution of galaxies in color space is bimodal, and correlated with galaxy morphology. Galaxies populate separate red and blue sequences in a ($u$ - $r$) versus $M_r$ color-magnitude diagram, CMD (Baldry et al. 2004). Evolution of galaxies from the blue sequence locus of actively star-forming systems to the red sequence occurs via transition through an intermediate CMD zone, christened the green valley (Martin et al. 2005). Bell et al. (2004) and Faber et al. (2007) have both shown that the red sequence has grown in mass over the period from $z\\sim1$ to $z\\sim0$. Given that the color bimodality is driven principally by star formation activity and UV bands are an excellent tracer of recent star formation, it is not surprising that the galaxy color sequences are especially well-separated in the UV-optical CMD (see Wyder et al. 2007, hereafter W07). GALEX proved instrumental in characterising the incidence of green valley transition galaxies (W07; Martin et al. 2007, hereafter M07) and understanding their evolution (M07, Schiminovich et al. 2007) from star-forming systems (Salim et al. 2007) to ``red and dead'' galaxies. The effect of gas removal, and the subsequent quenching of SF, has been modeled (M07) in the context of galaxies evolving from the blue sequence to the red sequence across the green valley. However, the influence of late addition of gas to red sequence galaxies has received little attention. Contrary to classical expectations, galaxies of the types predominately populating the red sequence (elliptical and lenticular, S0) sometimes are associated with gaseous reservoirs, especially if they are in a low density environment rather than within a cluster (Morganti et al. 2008). Oosterloo et al. (2007) showed the structure of such gas is varied, with some regularly-rotating disks and a complement of extended, offset, even tail-like morphologies -- suggesting a diversity in origin. Both galaxy mergers and IGM accretion are viable mechanisms. The presence of centralized star formation in these early-type galaxies (ETGs) populating the red sequence has been associated with low angular momentum sources (such as retrograde mergers) because gas can be efficiently transported to the remnant center and consumed (Serra et al. 2008), leaving behind minimal HI. On the other hand, prograde mergers of a red sequence galaxy with a gas rich object may be the dominant mechanism of forming massive, extended HI distributions around ETGs. \\object{NGC 404} is a prime target for studies addressing the structure and evolution of field S0 galaxies. It is particularly interesting because it is the nearest lenticular galaxy, lying just outside the Local Group at an estimated distance of 3.1--3.3 Mpc (Karachentsev et al. 2002, Tonry et al. 2001). \\object{NGC 404} is presently isolated in a zone of radius $\\sim$1.1 Mpc (Karachentsev \\& Markarov 1999). Although the galaxy is not affected by a cluster environment (as is true for many S0s), it may still have a complex history. Karachentsev \\& Markarov (1999) suggested \\object{NGC 404} represents the end product in the coalescence of a small galaxy group, having accreted with all other (less massive) group members. At least one merger may have occurred in the recent past ($<$1 Gyr) and be fundamentally responsible for the current UV morphology. del Rio et al. (2004) argue that a large, disk-like, neutral atomic hydrogen (HI) ring surrounding \\object{NGC 404} is the remnant of a merger with a dwarf irregular galaxy which took place some 900 Myr ago, according to their kinematic estimation. Our GALEX observations revealed that this HI ring is now forming stars. However surprising, this rather late discovery is understandable post facto since \\object{NGC 404} is nearly hidden in the glare of red supergiant Beta Andromedae (Mirach), 406$\\arcsec$ to the south east of the galaxy. This bright foreground star makes it difficult to image the faint surface brightness portions of \\object{NGC 404} at red wavelengths (including H$\\alpha$), and the S0 galaxy became colloquially known as \"Mirach's Ghost\". In this letter, we present evidence for star formation in \\object{NGC 404}'s outer disk-like HI ring. We refrain from detailed interpretation of UV properties of the inner disk and bulge. Section 2 describes our GALEX observations. Section 3 contains our data analysis. We conclude with discussion in Section 4. ", "conclusions": "Star formation and associated gas are no longer viewed as uncommon in lenticular and elliptical galaxies, owing to ever more sensitive UV and HI imaging surveys. Residual star formation is thought to be present in 10--30\\% of the early-type galaxies examined with GALEX (Yi et al. 2005, Donas et al. 2007, Schawinski et al. 2007, Kaviraj et al. 2007), even after excluding classical UV upturn candidate galaxies from the sample. An external source of gas (accretion or mergers) very likely supplements the SF fuel available via recycling from stellar mass loss. HI imaging for ETG samples supports this conclusion, frequently revealing extended gaseous distributions (Morganti et al. 2008). What remains to be determined is the long term effect of externally-fueled SF on the morphology of red-and-dead galaxies and on the appearance of the UV-optical galaxy CMD. Sometimes the gas is quickly consumed near the galaxy center (Serra et al. 2008), having little net effect on morphology and a very short-lived movement within the galaxy CMD. Activity of this sort is the ``frosting'' variety described by Trager et al. (2000), though Schawinski et al. (2009) also describe ETGs with rather high centralized SFRs ($>50$ M$_{\\odot}$ yr$^{-1}$). We have shown that rejuvenation of large-scale disk formation is another possible outcome, as discovered in NGC~404 at a very low level. NGC~404 is classified as a Type 1 XUV-disk using the criteria of Thilker et al. (2007). Similar cases have been reported by Cortese \\& Hughes (2009), Donovan et al. (2009) and Rich \\& Salim (2009). Rejuvenation events may effectively transplant red sequence galaxies to the green valley and blue sequence. We note that E/S0 galaxies have been detected already on the blue sequence (Kannappan et al. 2009, Schawinski et al. 2009). NGC~404 represents a possible example of this transition underway. We conclude that the galaxy CMD transition zone known as the green valley represents a heterogeneous population of objects, with many evolving from blue to red but others going in the opposite direction. Stochastic excursions into the green valley from the red-sequence, driven by acquisition of fresh gas for star formation, are observed. Traffic through the green valley is not one-way. NGC~404 confirms that the merger of ETGs and gas-rich dwarfs are one mechanism establishing this diversity. This implies UV imaging is fundamentally required to place individual ETGs into an accurate evolutionary context, assessing whether they are red and dead (at least for the moment), experiencing residual star formation in the galaxy center, or entering a rejuvenated phase of disk building fueled by a long-lasting gas reservior. In fact, UV imaging is an effective means of detecting such potentially transformational reservoirs." }, "1003/1003.1875_arXiv.txt": { "abstract": "In this work the phenomenology of models possessing a non-minimal coupling between matter and geometry is discussed, with a particular focus on the possibility of describing the flattening of the galactic rotation curves as a dynamically generated effect derived from this modification to General Relativity. Two possibilities are discussed: firstly, that the observed discrepancy between the measured rotation velocity and the classical prediction is due to a deviation from geodesic motion, due to a non-(covariant) conservation of the energy-momentum tensor; secondly, that even if the principle of energy conservation holds, the dynamical effects arising due to the non-trivial terms in the Einstein equations of motion can give rise to an extra density contribution that may be interpreted as dark matter. The mechanism of the latter alternative is detailed, and a numerical session ascertaining the order of magnitude of the relevant parameters is undertaken, with possible cosmological implications discussed. ", "introduction": "Many modifications of the theories of gravity are strongly motivated by the issue of an alternative solution to the problem of the flattening of the galaxy rotation curves --- one which does not require dark matter in the strictest sense, {\\it {\\it i.e.}} weakly or non-interacting matter fields \\cite{dm1,dm2,dm3,dm4}. With this goal in mind, standard $f(R)$ models \\cite{fR1,fR2,fR3} have been studied endowed with power-law curvature terms $f_1(R) \\propto R^n$ in the action, instead of the linear curvature depicted in the General Relativity (GR) action. For $n>1$, an addition to the Newtonian potential $\\Delta \\Phi(r) = -GM(r/r_c)^\\be / 2r$ is found, with $\\be$ a function of $n$ and $r_c$ being a parameter characteristic of each galaxy. A fit of several rotation curves indicates that $ n = 3.5$ (corresponding to $\\be = 0.817$) yields the best agreement with observations \\cite{CapoLSB}. In another approach, it has been shown that an asymptotically flat velocity dispersion curve can be obtained if the curvature term in the Einstein-Hilbert action includes a logarithmic factor, $f_1(R) = f_0 R(1 + v \\log R) $ \\cite{LoboLog}; given the non-relativistic condition $v \\ll 1$, this may be approached by a power-law with an exponent very close to unity, $f_1(R) \\approx f_0 R^{1 + v^2}$. Since this description does not take into consideration any galaxy-dependent parameters, an universal asymptotic velocity $v$ is obtained; this is in stark contrast with the Tully-Fisher and Faber-Jackson relations, which posit a dependence of the visible mass of a galaxy $M \\propto v^m$, for spiral ($m=4$) and elliptical galaxies ($m=6$), respectively. This work reviews the results obtained in Ref. \\cite{paramos}, to which the reader is directed for a more technical discussion and the full set of computations. It comprises a small introduction to the non-minimal coupling model under scrutiny, proceeds with the implied non-geodesic motion already hinted above, and then proposes the already mentioned mechanism for dynamically mimicking dark matter; a numerical session and discussion follows. ", "conclusions": "In this work the possibility of obtaining a solution to the dark matter puzzle, embodied by the flattening of galaxy rotation curves, was approached by resorting to the main phenomenological implications of models possessing a non-minimal coupling of matter to curvature, following results reported in Ref. \\cite{paramos}. As a first attempt, one first examined the non-conservation of the energy-momentum tensor of matter and the implied deviation from geodesic motion, and ascertained what coupling $f_2(R)$ should be so that the derived extra force would lead to the reported flattening of the rotation curves. This requires a logarithmic coupling of the form $\\la f_2(R) = -v^2/m \\log (R/ R_*)$ (where $m$ is the outer slope of the visible matter density $\\rho$), which can be approximated by a simpler power-law $\\la f_2(R) \\approx (R_*/R)^\\al$, with the asymptotic velocity given by $v_\\infty^2 = \\al m$. However, this solution yields an almost universal $v_\\infty$, prompting for another possible mechanism for the dynamical mimicking of ``dark matter''. This need led to the second, more evolved approach, which is rooted empirically in the phenomenological Tully-Fisher relation: one may instead assume that geodesical motion is preserved, $\\nabla^\\mu T_\\mn = 0$ (a condition proved to be self-consistent), but that the metric itself is perturbed: the mimicked dark matter density is then given by the difference $R/2\\ka - \\rho$. By resorting to a power-law coupling with matter $f_2(R) = (R/R_0)^n$ (with a negative exponent $n$ yielding the desired effect at low curvatures and long range). The proposed power-law directly leads to a ``dark matter'' density that depends on a power of the visible matter density, thus accounting for the Tully-Fisher law in a natural way. The obtained ``dark matter'' component has a negative pressure, as commonly found in cosmology in dark energy models: this hints the possibility that the non-minimal coupling model might unify dark matter and dark energy. Given their relevance in the literature, two different scenarios were considered: the NFW and isothermal sphere dark matter profiles ($n=-1/3$ and $n=-1$, respectively). Although separate fits of the selected galaxy rotation curves to each coupling do not yield satisfactory results, a composite coupling of both power-laws produced a much improved adjustment. The characteristic lengthscales $r_1$ and $r_3$ were taken as fitting parameters for each individual galaxy, yielding the order of magnitudes $r_1 \\sim 10~Gpc$ and $r_3 \\sim 10^5 ~Gpc$. This enabled the computation of the cosmological background matching distances for each galaxy (which depend on their characteristic lengthscale $a$) and the obtained astrophysical range showed that the $n=-1$ isothermal sphere ``dark matter'' halo dominates the $n=-1/3$ NFW component. Furthermore, the $n=-1$ scenario was shown to have possible relevance in a cosmological context, as $r_1 \\sim r_H$, the latter being the Hubble radius --- again hinting at a possible unification of the dark components of the Universe. The lack of the desirable universality in the model parameters $R_1$ and $R_3$ was duely noticed, although these quantities are characterized by the same order of magnitude for the best fits obtained; several possible causes for this variation were put forward, from deviation from spherical symmetry to a more complex form for the non-minimal coupling between curvature and matter $f_2(R)$ or the curvature term $f_1(R)$, or the need for a more evolved Lagrangian description of the latter. As a final remark, one concludes that the rich phenomenology that springs from the model with a non-minimal coupling between matter and curvature enables a direct and elegant alternative to standard dark matter scenarios and many modifications to GR, which usually resort to extensive use of additional fields and other {\\it ad-hoc} features. As an example of the latter, one points out the MOdified Newtonian Dynamics (MOND) hypothesis, which is by itself purely phenomenological, and whose underlying Tensor-Vector-Scalar theory is based upon an extensive paraphernalia of vector and scalar fields \\cite{MOND1,MOND2,MOND3,MOND4} (see Ref. \\cite{MONDcritic} for a critical assessment). As a final remark, we point out that the considered model may also be translated into a multi-scalar theory \\cite{scalar}, with two scalar fields given by \\beq \\varphi^1 = {\\sqrt{3}\\over2} \\log \\left[ 1 + n \\left({R \\over R_0}\\right)^n \\right] \\qquad \\varphi^2 = R, \\eeq \\noindent with dynamics driven by a potential \\beq U(\\varphi^1,\\varphi^2) = {1 \\over 4} \\exp \\left( -{2 \\sqrt{3}\\over 3} \\varphi^1 \\right) \\left[\\varphi^2 - {f_1(\\varphi^2 )\\over 2\\ka } \\exp \\left( -{2 \\sqrt{3}\\over 3} \\varphi^1 \\right) \\right]. \\eeq \\ack The authors thank E. Vagena for fruitful discussions and the hospitality shown during the {\\it First Mediterranean Conference on Classical and Quantum Gravity}, Crete, September 2009. \\par ~" }, "1003/1003.3041_arXiv.txt": { "abstract": "For the first time, {\\it Swift} is giving us the opportunity to study supergiant fast X--ray transients (SFXTs) throughout all phases of their life: outbursts, intermediate level, and quiescence. We present our intense monitoring of four SFXTs, observed 2--3 times per week since October 2007. We find that, unexpectedly, SFXTs spend most of their time in an intermediate level of accretion ($L_{X}\\sim 10^{33-34} $ erg s$^{-1}$), characterized by rich flaring activity. We present an overview of our investigation on SFXTs with {\\it Swift}, the key results of our Project. We highlight the unique contribution {\\it Swift} is giving to this field, both in terms of outburst observations and through a systematic monitoring. ", "introduction": " ", "conclusions": "" }, "1003/1003.3107_arXiv.txt": { "abstract": "We establish the jump conditions for the wavefunction and its derivatives through the formal solutions of the wave equation. These conditions respond to the requirement of continuity of the perturbations at the position of the particle and they are given for any mode at first order. Using these jump conditions, we then propose a new method for computing the radiated waveform without direct integration of the source term. We consider this approach potentially applicable to generic orbits. ", "introduction": "The complexity in assessing the continuity of the perturbations at the position of the particle in the Regge-Wheeler gauge has led, among other motivations, to work in the Lorenz gauge, at the price of loosing the availability of the wave equation. Nevertheless, it has been indicated by two different heuristic arguments \\cite{lo00, lona09} that the even metric perturbations for radial fall should belong to the $C^0$ continuity class at the position of the particle, in the Regge-Wheeler gauge\\footnote{The discontinuities of the wave function and its derivatives were also addressed elsewhere \\cite{sola06}.}. Herein, we require that the perturbations are $C^0$ by identifying the conditions that the wavefunction and its derivatives have to satisfy for allowing the perturbations to belong to such category. Our analysis is based on the solutions of the Zerilli equation, not on the equation themselves. The inverse relations for the perturbation functions $K$, $H_2$, $H_1$ are given by (having suppressed the $l$ index and being $m=0$): \\beq K=\\frac{6M^2+3M\\lambda r+\\lambda (\\lambda +1)r^2} {r^2(\\lambda r+3M)}\\Psi +\\left( 1-\\frac{2M}r\\right) \\,\\Psi_{,r} -\\frac{ \\kappa \\ U^0(r-2M)^2}{(\\lambda +1)(\\lambda r+3M)r}\\delta \\nn \\eeq \\[ H_2\\!=\\!-\\!\\frac{9M^3\\!+\\!9\\lambda M^2r\\!+\\!3\\lambda ^2Mr^2\\!+\\!\\lambda ^2(\\lambda \\!+\\!1)r^3} {r^2(\\lambda r\\!+\\!3M)^2}\\,\\Psi \\!+\\! \\frac{3M^2\\!-\\!\\lambda Mr\\!+\\!\\lambda r^2}{r(\\lambda r\\!+\\! 3M)}\\Psi_{,r} \\!+\\! (r\\!-\\!2M)\\Psi_{,rr} \\] \\beq + \\frac{\\kappa U^0(r - 2M)(\\lambda ^2r^2+2\\lambda Mr-3Mr+3M^2)}{r (\\lambda +1)(\\lambda r+3M)^2}\\delta - \\frac{\\kappa U^0(r-2M)^2}{ (\\lambda +1)(\\lambda r+3M)}\\delta' \\nn \\eeq \\beq H_1=\\frac{\\lambda r^2-3M\\lambda r-3M^2}{\\left( r-2M\\right) (\\lambda r+3M)}{\\Psi_{,t}}+r\\Psi_{,tr} - \\frac{\\kappa \\ U^0\\stackrel{.}{z}_u(\\lambda r+M)}{(\\lambda +1)(\\lambda r+3M)}\\delta + \\frac{\\kappa \\ U^0\\stackrel{.}{z}_u r(r-2M)}{(\\lambda +1)(\\lambda r+3M)}\\delta ' \\nn \\eeq where $ \\lambda \\!= \\!1/2(l\\! -\\! 1)(l\\! + \\!2) $, $\\kappa\\! = \\!4m\\sqrt{(2l\\!+\\!1)\\pi}$, $\\delta \\! = \\! \\delta\\left[r\\!-\\!z_u(t)\\right ]$ and $\\delta'\\! =\\! \\delta'\\left[r\\!-\\!z_u(t)\\right ]$; $U^0 \\!= \\!E/(1\\! - \\!2M/z_u)$ is the time component of the 4-velocity, $z_u$ the coordinate time dependent position, $E$ the energy of the particle. Since the wavefunction $\\Psi$ belongs to the $C^{-1}$ continuity class, it and its derivatives can be written as, \\beq \\Psi(t,r)=\\Psi^+(t,r)~\\Theta_1+\\Psi^-(t,r)~\\Theta_2 ~~~~~~~~~~~~~~~~ \\Psi_{,r} = \\Psi^+_{,r}\\Theta_1 + \\Psi^-_{,r} \\Theta_2 + \\left(\\Psi^+ -\\Psi^-\\right) \\delta \\nn \\eeq \\beq \\Psi_{,rr} = \\Psi^+_{,rr}\\Theta_1 + \\Psi^-_{,rr} \\Theta_2 + \\left( \\Psi^+_{,r} - \\Psi^-_{,r} \\right) \\delta + \\left(\\Psi^+ - \\Psi^-\\right)\\mid_{r = z_u} \\delta' \\nn \\eeq \\beq \\Psi_{,t} = \\Psi^+_{,t}\\Theta_1 + \\Psi^-_{,t} \\Theta_2 - \\left(\\Psi^+ -\\Psi^-\\right) \\dot{z}_u \\delta \\nn \\eeq \\beq \\Psi_{,tr} = \\Psi^+_{,tr}\\Theta_1 + \\Psi^-_{,tr} \\Theta_2 + \\left(\\Psi^+_{,t} -\\Psi^-_{,t}\\right) \\delta - \\left(\\Psi^+ -\\Psi^-\\right)\\mid_{r = z_u} \\dot{z}_u \\delta' \\nn \\eeq where $\\Theta_1 = \\Theta\\left[r-z_u(t)\\right ]$, and $\\Theta_2 = \\Theta\\left[z_u(t) - r \\right ]$ are two Heaviside step distributions. For the second derivatives, the property of the Dirac delta distribution, at the position of the particle: $f(r) \\delta'[r-z_u(t)]=f(z_u(t)) \\delta'[r-z_u(t)] - f'(z_u(t)) \\delta[r-z_u(t)]$, has been used. The discontinuities of $\\Psi$ and its derivatives must be such that they cancel when combined in $K$, $H_2$ and $H_1$ at the position of the particle. After replacing $\\Psi$ and its derivatives in the perturbations, continuity requires that the coefficients of $\\Theta_1$ must be equal to those of $\\Theta_2$, while the coefficients of $\\delta$ and $\\delta '$ must vanish separately. Finally, the jump conditions for $\\Psi$ and its derivatives are found (the jump conditions provided by $K$, $H_2$ and $H_1$ are equivalent): \\beq \\Psi^+ - \\Psi^- = \\frac{\\kappa E z_u}{(\\lambda +1) (3 M+\\lambda z_u)} ~~~~~~~~~~~~~~ \\Psi^+_{,r} - \\Psi^-_{,r} = \\frac{\\kappa E \\left[6 M^2+3 M \\lambda z_u+\\lambda (\\lambda +1) z_u^2\\right]}{(\\lambda +1) (2 M-z_u) (3 M+\\lambda z_u)^2} \\nn \\eeq \\beq \\Psi^+_{,rr} - \\Psi^-_{,rr}\\! = \\! -\\frac{\\kappa E \\left[3 M^3 (5 \\lambda -3)+6 M^2 \\lambda (\\lambda -3) z_u+3 M \\lambda ^2 (\\lambda -1) z_u^2-2 \\lambda ^2 (\\lambda +1) z_u^3\\right]}{(\\lambda +1) (2M-z_u)^2 (3 M+\\lambda z_u)^3} \\nn \\eeq \\beq \\Psi^+_{,t} - \\Psi^-_{,t}\\! = \\! -\\frac{\\kappa E z_u \\dot{z}_{u}}{(2 M - z_u) (3 M+\\lambda z_u)} \\nn ~~~~~~~~~~~~~ \\Psi^+_{,tr} - \\Psi^-_{,tr} \\!= \\! \\frac{\\kappa E \\left(3 M^2+3 M \\lambda z_u - \\lambda z_u^2\\right)\\dot{z}_{u}}{(2M - z_u)^2 (3 M+\\lambda z_u)^2} \\nn \\eeq ", "conclusions": "" }, "1003/1003.4511_arXiv.txt": { "abstract": "We present multi-wavelength observations of the centre of RXCJ1504.1-0248 -- the galaxy cluster with the most luminous and relatively nearby cool core at $z\\sim 0.2$. Although there are several galaxies within 100\\,kpc of the cluster core, only the brightest cluster galaxy (BCG), which lies at the peak of the X-ray emission, has blue colours and strong line-emission. Approximately 80\\Msunpyr\\ of intracluster gas is cooling below X-ray emitting temperatures, similar to the observed UV star formation rate of $\\sim$140\\Msunpyr. Most star formation occurs in the core of the BCG and in a 42\\,kpc long filament of blue continuum, line emission, and X-ray emission, that extends southwest of the galaxy. The surrounding filamentary nebula is the most luminous around any observed BCG. The number of ionizing stars in the BCG is barely sufficient to ionize and heat the nebula, and the line ratios indicate an additional heat source is needed. This heat source can contribute to the H$\\alpha$-deduced star formation rates (SFRs) in BCGs and therefore the derived SFRs should only be considered upper limits. AGN feedback can slow down the cooling flow to the observed mass deposition rate if the black hole accretion rate is of the order of 0.5 \\Msunpyr at 10\\% energy output efficiency. The average turbulent velocity of the nebula is $v_{\\rm turb}\\sim 325\\,\\kmps$ which, if shared by the hot gas, limits the ratio of turbulent to thermal energy of the intracluster medium to less than 6 percent. ", "introduction": "} Cool core clusters are characterised by bright central peaks in their X-ray surface brightness profiles, cool and radially-increasing core temperatures, and short central cooling times. Unless the gas is reheated, the large amount of energy radiated in X-rays implies several hundred solar masses of gas cools below $10^{7}$\\,K every year \\citep{CowieBinney77,FabianNulsen77}. However, X-ray spectra from \\emph{XMM-Newton} and \\emph{Chandra} show cool-core clusters lack strong cooling lines, so less gas is cooling below X-ray emitting temperatures than is expected from the X-ray luminosity \\citep{Peterson01,Peterson03,Boehringer01,Peterson06,McNamara07}. The implication is that the intracluster medium (ICM) must be heated. On the other hand, large quantities of cool gas and dust have been found in brightest cluster galaxies at the centre of cool-core clusters. Searches conducted at optical and infrared wavelengths reveal both warm and cool gas reservoirs (e.g. \\citealt{Edge01}; \\citealt{Salome03}; \\citealt{Hatch05}; \\citealt*{Jaffe05}; \\citealt{Johnstone07}). The BCGs in cool-core clusters generally have anomalous blue colours implying recent star formation, and emit in recombination and low-ionization lines \\citep{McNamara97,Peres1998,Cavagnolo2008}. These observations imply a relationship between the properties of the cluster and the central galaxy, but this interaction has not yet been understood or quantified precisely. RXCJ1504.1-0248 is one of the most massive cool core clusters, with a prominent central X-ray brightness peak, and a short central cooling time. A classical cooling flow model leads to a mass deposition rate of $1400-1900$\\,\\Msunpyr, and the brightest cluster galaxy emits strong low ionization emission lines \\citep{Boehringer1504}. These properties mark this cluster as a prime target to study the relationship between the ICM, the brightest cluster galaxy, and a possible central active galactic nucleus (AGN). Here we report results of a multi-wavelength study of the core of RXCJ1504.1-0248. It lies at $z=0.2153$, which allows for a detailed spatial analysis of the BCG. We present results of integral field unit (IFU) spectroscopic observations aimed at understanding the morphology, kinematics and ionization state of the line-emitting nebula surrounding the BCG. We also present an analysis of X-ray spectra obtained with the XMM-Newton Reflection Grating Spectrometer (RGS) and European Photon Imaging Cameras (EPIC) and compare the ICM mass deposition rate to the star formation rate of the BCG. Sections \\ref{sec:obs} and \\ref{sec:dr} describe the observations and the data reduction process, while in section \\ref{sec:results} we present the properties of the BCG, the nebula, and the cluster core cooling rate. In section \\ref{sec:disc} we combine the results from the multi-wavelength observations to quantify the relationship between the ionized gas, the BCG and the cluster core. We summarize our findings in section \\ref{sec:concl}. Unless otherwise stated, the following cosmological parameters have been used: $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, $H_0=70\\,{\\rm km\\, s^{-1}\\, Mpc^{-1}}$. For a redshift $z=0.2173$ (the redshift of the BCG), the scale is $3.519\\,\\,{\\rm kpc\\,arcsec^{-1}}$. ", "conclusions": "\\label{sec:concl} RXCJ1504.1-0248 is one of the most extreme cooling flows ever discovered. The observations presented here offer the first multi-wavelength view of the cluster core, examining both the hot and warm gas as well as the stars of the BCG. The core of RXCJ1504.1-0248 has an X-ray luminosity of at least $3\\times 10^{45}$ ${\\rm erg\\, s^{-1}}$, approximately 10 percent of which is supplied by gas cooling and forming stars. Accretion onto a supermassive black hole of $2\\times 10^9\\,\\, \\Msun$ is the likely mechanism needed to account for most of the remaining $2.7\\times 10^{45}$ ${\\rm erg\\,s^{-1}}$ heating. The corresponding Bondi accretion rate at 10 percent radiative efficiency is 0.5 \\Msunpyr. Using the correlation between 1.4 GHz radio luminosity and jet power \\citep{Birzan2008}, we infer a jet power of $1.5\\times 10^{44}$ ${\\rm erg\\, s^{-1}}$, still insufficient to heat the ICM. The intracluster medium mass deposition rate is in agreement with the star formation rate of the BCG derived through FUV observations ($\\sim$140\\Msunpyr). The IFU images show that most of the star formation occurs in the galaxy core and in a 42\\,kpc-long filament that stretches out from the core to the SW. These regions have blue continuum colours, and shallow 4000\\AA\\ breaks consistent with recent star formation. Additionally, the [O{\\sc ii}]/H$\\beta$ line ratio indicates stellar photoionization is the dominant source of ionization along the filament. The X-ray emission also extends along the filament indicating a direct connection between the intracluster medium and the recent star formation. The line-emitting nebula that surrounds the BCG is the most luminous observed to date, with extinction-corrected [O{\\sc ii}] and H$\\alpha$ luminosities of $4.7\\times 10^{43}\\,\\,{\\rm erg\\,s^{-1}}$ and $3.4\\times 10^{43}\\,\\,{\\rm erg\\,s^{-1}}$ respectively. The number of ionizing O and B type stars in the BCG is barely sufficient to power this bright nebula, and the BPT diagrams and line ratio maps show that an additional heating mechanism is required. Because this heating source produces H$\\alpha$ emission, BCG star formation rates that are derived from H$\\alpha$ luminosities should be considered upper limits. The nebula kinematics are ordered and the nebula comprises of a number of kinematically distinct regions. The velocities are generally low ($<250\\,\\,{\\rm km\\,s^{-1}}$), and there is no evidence for rotation or free-falling gas infall. The 42\\kpc\\ filament has a velocity shear of $\\sim$400\\kmps\\ resulting in a dynamical time of $10^{8}$\\,years, thus the nebula must be long-lived. The average velocity dispersion of the nebula is $v_{\\sigma} \\sim 200$ ${\\rm km\\,s^{-1}}$ and the dispersion of the velocity shear of the cold gas is $v_{\\Delta}\\sim 117\\,\\kmps$. These set an upper limit of $\\sim 325\\,\\kmps$ on the mean turbulent velocity in the ICM, meaning that the ratio of turbulent to thermal energy of the intracluster medium is limited to less than 6 percent. Our results suggest mass transfers from the intracluster medium to the brightest cluster galaxy at a rate of approximately $10^{2}$\\Msunpyr. The cooling of the intracluster medium results in a gas reservoir in the brightest cluster galaxy. The rate at which gas condenses into this reservoir is similar to the rate at which this reservoir forms stars, so the reservoir should be approximately stationary. Our observations cannot reveal the coolest gas components of this reservoir, but we find that the $10^{4}$\\,K gas is partly heated by a source other than the forming stars." }, "1003/1003.4219_arXiv.txt": { "abstract": "{The isocyanic acid (HNCO) presents an extended distribution in the centers of the Milky Way and the spiral galaxy IC342. Based on the morphology of the emission and the HNCO abundance with respect to H$_2$, several authors made the hypothesis that HNCO could be a good tracer of interstellar shocks.} {Here we test this hypothesis by observing a well-known Galactic source where the chemistry is dominated by shocks.} {We have observed several transitions of HNCO towards L1157-mm and two positions (B1 and B2) in the blue lobe of the molecular outflow.} {The HNCO line profiles exhibit the same characteristics of other well-known shock tracers like CH$_3$OH, H$_2$CO, SO or SO$_2$. HNCO, together with SO$_2$ and OCS, are the only three molecules detected so far whose emission is much more intense in B2 than in B1, making these species valuable probes of chemical differences along the outflow. The HNCO abundance with respect to H$_2$ is 0.4-1.8\\,10$^{-8}$ in B1 and 0.3-1\\,10$^{-7}$ in B2. These abundances are the highest ever measured, and imply an increment with respect to L1157-mm of a factor up to 83, demonstrating that this molecule is actually a good shock tracer.} {Our results probe that shocks can actually produce the HNCO abundance measured in galactic nuclei and even higher ones. We propose that the gas phase abundance of HNCO is due both to grain mantles erosion by the shock waves and by neutral-neutral reactions in gas phase involving CN and O$_2$. The observed anticorrelation of CN and HNCO fluxes supports this scenario. The observed similarities of the HNCO emission and the sulfured molecules may arise due to formation pathways involving also O$_2$. } ", "introduction": "Interstellar {\\it isocyanic acid} (HNCO) was first detected towards Sgr B2 \\citep{Snyder72, Churchwell86, Kuan96}. Since the first detection in this source, the molecule has been observed in other hot cores around massive \\citep{Blake87, MacDonald96} and low mass protostars \\citep{vanDishoeck95, Bisschop08}. It has also been detected in translucent clouds \\citep{Turner99} and in the dense regions of Galactic molecular clouds \\citep{Jackson84, Zinchenko00}, including those in the Galactic center \\citep{Huttemeister93, Lindqvist95, Dahmen97, Rizzo00, Minh05, Martin08}. HNCO has also been detected in some extragalactic sources \\citep{Nguyen91, Meier05, Martin09}. The isotopologue HCNO (fulminic acid) has recently been detected by \\cite{Marcelino09}. The emission in HNCO lines with $K_{-1} = 0$ \\footnote{ HNCO is a planar, nearly linear, slightly asymmetric prolate rotor. The notation for a given level is $J_{K_{-1}K_1}$. The $K_{-1} = 0$ are usually excited thermally, while the excited $K$ ladders ($K_{-1} > 0$) are probably excited by FIR radiation \\citep{Churchwell86}. } is clearly extended in the two Galactic nuclei where it has been mapped: IC342 \\citep{Meier05} and the Milky Way \\citep{Lindqvist95, Dahmen97}. In the Galactic center the spatial distribution of the HNCO emission is different from that of most other tracers. In particular, there is a HNCO peak without an associated CO peak at Galactic longitude $l=1.65^\\circ$ \\citep{Dahmen97}. This very special distribution suggests that HNCO can be an important tracer of some physical processes that are not well revealed by other molecules. It is possible that the molecule is tracing shocks, which are thought to take place at this Galactic longitude due to the gas dynamics in the barred potential of the Milky Way \\citep{Rodriguez08}. Indeed, in the Galactic center, the highest gas phase abundances of SiO (a well-known shock tracer) are measured in this region \\citep{Huttemeister98, Rodriguez06}. Also on the basis of a special morphology of the HNCO emission in Sgr B2, \\cite{Minh06} suggested that HNCO is enhanced by shocks. Regarding IC342, the HNCO emission resembles that of CH$_3$OH. In particular, HNCO is detected not only in the nuclear ring but also in the inner spiral arms or \\emph{dustlanes}. Since the dustlanes are thought to be the locus of strong shocks in barred galaxies \\citep[see for instance][and references therein]{Rodriguez08}, \\cite{Meier05} suggested that HNCO, as CH$_3$OH, could trace large scale shocks. Nevertheless, the hypothesis that HNCO is a good shock tracer at the scale of galaxies still needs to be probed since the HNCO emission has never been studied in well-known Galactic templates of interstellar shocks. In order to better understand the excitation and the chemistry of this promising molecule, we have observed the protostar L1157 and its associated molecular outflow. This outflow presents the morphological signature of shocks \\citep[][]{Gueth98, Codella09} and it is frequently used to benchmark numerical models of shocks \\citep{Gusdorf08a, Gusdorf08b}. L1157 is the best example of ``chemically active outflow\" \\citep[][]{Bachiller01} and a template of shock chemistry, since many species exhibit large abundance increments with respect to the protostar \\citep[][]{Bachiller97}. Therefore, it is a source of choice to characterize the emission of a given molecule in a shocked environment \\cite[see for instance][]{Bachiller97, Bachiller01, Benedettini07, Arce08, Codella09}. \\begin{figure} \\begin{center} \\includegraphics[angle=-90,width=8cm]{fig_l1157.eps} \\end{center} \\caption{ Integrated CO(2-1) emission from the L1157 outflow \\citep{Bachiller01}. We show the position of L1157-mm (star), B1 (empty square) and B2 (filled square). B1 and B2 are the positions of the shocks as traced, for instance, by SiO(3-2) \\cite[see Fig. 6 of][]{Bachiller01}. The CO emission peaks behind the actual positions of the shocks. The different circles represent the half maximum contours of the IRAM 30m beam at the frequencies of the lines discussed in this paper (see Table~\\ref{tab:obs}). } \\label{fig:l1157} \\end{figure} ", "conclusions": "\\subsection{HNCO formation in molecular clouds and hot cores} \\label{sect:chemistryhnco} The possible formation pathways of HNCO in shocks have never been modeled. In contrast, HNCO has been included in some models of dark clouds and hot core chemistry. In the models by \\cite{Iglesias77}, HNCO is produced by the ion-neutral reaction $\\mathrm{H}_2 + \\mathrm{HNCO}^+ \\rightarrow \\mathrm{H}_2 \\mathrm{NCO}^+ + \\mathrm{H} $ followed by $\\mathrm{H}_2\\mathrm{NCO}^+ + e^- \\rightarrow \\mathrm{HNCO} + \\mathrm{H}$. \\cite{Turner99} have considered three possible formation pathways, among them the only efficient one is the neutral-neutral reaction $\\mathrm{CN} + \\mathrm{O}_2 \\rightarrow \\mathrm{NCO} + \\mathrm{O}$ followed by $ \\mathrm{NCO} + \\mathrm{H}_2 \\rightarrow \\mathrm{HNCO}+\\mathrm{H}$. { The last reaction has an activation barrier that has been estimated to be $\\sim 1160~K$ by \\cite{Turner99} but that could be as high as 4465~K \\citep{Tideswell10}. In any case, this reaction is not efficient at the typical temperature of a hot core ($\\sim$ 200 K). Therefore, gas phase chemistry alone cannot explain the HNCO abundances in hot cores. This is also the case for other species as the complex organic molecules. Chemical models have consequently been developed to include reactions on the dust grains surface \\citep{Caselli93, Garrod08, Tideswell10}. In those works, hot cores are modeled in two stages. In the first phase, a dark cloud of $n_H \\sim 10^4$~\\cmmt \\ suffers an isothermal collapse. The collapse phase is halted once $n_H \\sim 10^7$ \\cmmt , which occurs after approximately $5\\,10^{5}$ years (free-fall time). The typical temperature at this phase is 10-20 K and surface chemistry is very important then. Surface chemistry models \\citep{Garrod08, Tideswell10} produce HNCO as a secondary radical formed via the reaction $\\mathrm{NH+CO} \\rightarrow \\mathrm{HNCO}$. \\footnote{HNCO formation on the grain surfaces could also occur by reactions of OCN$^-$ with NH$_4^+$ or H$_3$O$^+$. The inverse reactions (HNCO with NH$_3$ or H$_2$O) are invoked to explain the presence of OCN$^-$ ices in the dust grains, \\citep{Demyk98, vanBroekhuizen04}. To our knowledge, these reactions have never been included in chemical models.} On the grain surface, HNCO can reach a maximum abundance of 10$^{-5}$ \\citep{Tideswell10} before being destroyed by new reactions with primary radicals (H, CH$_3$, HCO, NH, ...) giving complex species as HNCHO, HNCOCHO, or CH$_3$CONH \\citep{Garrod08}. The cloud-collapse phase is followed by a warm-up phase in which the temperature increases to $\\sim 200$~K. The timescale of the warm-up phase can be as short as $\\sim 5\\,10^4$~yr in hot cores but it can reach 10$^6$ yr in their low mass equivalent ({\\it hot corinos}), giving a similar although slightly different chemistry \\citep{Garrod08}. During the warm-up phase molecules are evaporated from the grain mantles. With respect to HNCO, the models by \\cite{Garrod08} and \\cite{Tideswell10} show that its abundance increases monotonously with time in the gas phase, which implies that HNCO is not directly ejected from the grains. Instead, HNCO is formed by the destruction in the gas phase of the complex molecules formed from HNCO on the grain surface \\citep{Garrod08, Tideswell10}. This mechanism can explain the HNCO abundances measured in hot cores of $10^{-9}-10^{-8}$. As already mentioned, the contribution of the \\cite{Turner99} and \\cite{Iglesias77} formation pathways is negligible at the moderate temperature of a hot core. The abundance of HNCO in L1157-mm is similar to that measured in the hot corino of the only other low mass protostar where HNCO has been observed (IRAS16293-2422, see Table \\ref{tab:comp}). Therefore, the HNCO abundance measured in L1157-mm can also be explained in the context of the hot-cores models discussed above. This would imply that L1157-mm can be considered as a hot corino, which is indeed in agreement with the intense emission of water detected by Herschel towards mm (Nisini et al. {\\it in prep}). } However, more observations will be needed to clearly establish the hot corino character of L1157-mm, in particular observations of complex organic molecules. \\subsection{HNCO formation in shocks} \\label{sect:chemistry} What is the formation pathway of HNCO in the L1157 outflow shocks? The most likely scenario is a combination of grain surface and gas phase chemistry. { In the context of shocks, molecules formed on the grain surfaces can be ejected to the gas phase due to grain sputtering instead of thermal evaporation as in hot cores. Therefore, the chemistry will be sensitive to the grain mantles composition at the time of the shock arrival. Currently, there are no models that study the HNCO abundance in dark clouds including grain surface chemistry. Dark clouds models by \\cite{Tideswell10} only consider gas phase reactions. Therefore, the exact grain mantle composition before the arrival of the shock in L1157-B1 and L1157-B2 is not known. Nevertheless, one can compare with the cloud-collapse phase of hot core models (the main difference is that the density during the collapse phase reaches higher values than in dark clouds). As discussed in the previous section, during the collapse phase the HNCO abundance on the grain surfaces has a peak of 10$^{-5}$, grain sputtering at that time will immediately give rise to very high HNCO abundance in gas phase. Even the very high HNCO abundance in L1157-B2 ($\\sim 10^{-7}$) could be accounted for in this context. On the other hand, gas phase chemistry in shocks differs considerably from that in hot cores because the temperature in the shocked gas can be much higher than that in a hot core. As already pointed out by \\cite{Zinchenko00}, in shocks the reaction $\\mathrm{NCO} + \\mathrm{H}_2 $ can be efficient in spite of the activation barrier. In addition, the reaction $\\mathrm{CN} + \\mathrm{O}_2 \\rightarrow \\mathrm{NCO} + \\mathrm{O}$ will be favored by the high O$_2$ abundances that are predicted in post-shock gas \\citep{Bergin98, Gusdorf08a, Gusdorf08b}. \\footnote{This does not mean that the \\textit{current} $O_2$ abundance in the L 1157 shocks is high, but it may have been high just after the passage of a shock-wave as predicted by shock models. Since the chemistry in shocked regions is fast, O$_2$ can be rapidly converted into other species. Therefore, there is not any contradiction with the fact that the O$_2$ abundance in the interstellar medium has been found to be low \\citep[$< 10^{-7}$][]{Larsson07}. } Using \\cite{Tideswell10} results, it is possible to verify if gas phase chemistry alone could explain the HNCO abundance measured in the L1157 shocks. In one of their dark cloud models (DC1), they have studied the HNCO formation exclusively via the gas phase neutral-neutral reactions of \\cite{Turner99} as a function of the rate coefficient for the reaction $ \\mathrm{NCO} + \\mathrm{H}_2 \\rightarrow \\mathrm{HNCO}+\\mathrm{H}$. Their Fig. 1 shows that a rate coefficient $k$ higher than $10^{-14}$ cm$^{3}$\\,s$^{-1}$ is needed to explain the observed abundances in L1157 of $10^{-8}-10^{-7}$. However, to have high HNCO abundance (a few $10^{-8}$) {\\it in a short time} (several 10$^3$ yr) the rate coefficient must be higher than $10^{-12}$~cm$^{3}$\\,s$^{-1}$. Taking into account the formula $k(T)=1.5\\,10^{-11} \\, \\exp(-4465/T) \\, \\mathrm{cm}^{-3} \\, \\mathrm{s}^{-1}$ given by \\cite{Tideswell10}, $k=10^{-12}$~cm$^{3}$\\,s$^{-1}$ implies a temperature higher than 1600 K. In contrast, the minimum temperature would be 428 K if the activation barrier is only 1160~K \\citep{Turner99}. \\cite{Gusdorf08b} have recently computed shock models to compare the modelled SiO and H$_2$ emission to observations of L1157. The best agreement between models and observations is found for a preshock density of 10$^{4}$ cm$^{-3}$ and a shock speed of 20 km\\,s$^{-1}$. In such a shock, the gas temperature is higher than 430-1600 K for a few thousand years, which is consistent with the dynamical age of the B1 and B2 shocks \\citep[2000-3000 yr, ][]{Gueth98} and with the chemical time necessary to have an HNCO abundance of a few 10$^{-8}$ only with gas phase reactions \\citep[Figure 1 of ][]{Tideswell10}. Therefore, gas phase chemistry alone could explain the HNCO abundance in L1157-B1. However, it is unclear whether a HNCO abundance as high as that measured in L1157-B2 ($\\sim 10^{-7}$) can be attained in a few thousand years only with gas phase reactions. } Therefore, we reckon that the most likely explanation to the high HNCO abundances in L1157 is dust grain mantles processing by the shock waves followed by neutral-neutral reactions in gas phase. \\begin{table} \\caption{HNCO abundance with respect to H$_2$ in different astronomical environments.} \\label{tab:comp} \\centering \\begin{tabular}{lll} \\hline Source & X(HNCO) & References\\\\ & 10$^{-9}$ &\\\\ \\hline L1157-B2 & 25-96 &1\\\\ Translucent clouds & 0.2-5 &2 \\\\ Orion bar PDR & $<10^{-2}$ &3 \\\\ Dense cores & 0.2-8.7 &4 \\\\ Orion Hot core & 5 &5\\\\ SgrB2M & 1 &6\\\\ W3(H$_2$O) & 5.0-27 &7 \\\\ G34.3+0.15 & 1.4 &8 \\\\ IRAS 16293 & 0.17-4.0 & 9, 10 \\\\ GC clouds & 3-48 & 11,12\\\\ Starburst galaxies & 1-6.3 & 13\\\\ \\hline \\end{tabular}\\\\ \\tablebib{(1)~This work; (2)~\\citet{Turner99};(3)~\\citet{Jansen95}; (4)~\\citet{Zinchenko00}; (5)~\\citet{Blake87}; (6)~\\citet{Minh98}; (7)~\\citet{Helmich97}; (8)~\\citet{Bockelee-Morvan00}, using the HNCO column density given by \\cite{MacDonald96} and the H$_2$ column density given by \\cite{Hatchell98}; (9)~\\citet{vanDishoeck95}; (10)~\\citet{Bisschop08}; (11)~\\citet{Rizzo00}; (12)~\\citet{Martin08}, and references within; (13)~\\citet{Martin09} } \\end{table} \\subsection{Chemical differences between B1 and B2 and the role of the O$_2$ chemistry} \\label{sect:o2chemistry} We have discussed the line profiles and the line intensities of different species in Sect.~\\ref{sect:profiles}. The HNCO profiles are very similar to those of SO and SO$_2$. In addition, the lines of most of the molecules show similar intensities in B1 and B2 with the exception of CN, which is more intense in B1, and HNCO, SO$_2$ and OCS, whose lines are more intense emission in B2 than in B1. These similarities of HNCO and the sulfured molecules is somewhat puzzling. The actual reason of the chemical differences in B1 and B2 is not clear. The present gas density in B1 and B2 is similar \\citep[see Figs. \\ref{fig:radex1} and \\ref{fig:radex2} and][]{Bachiller97, Nisini07}. In contrast, based on the line profiles, the shock velocity can be higher in B1 than in B2. Taking into account the extreme SiO line wings, the difference could reach 10 \\kms, which after the \\cite{Gusdorf08a} models could be significant. Alternatively, \\cite{Bachiller97} have suggested that the chemical differences between B1 and B2 could be due to different shock ages. Indeed, in contrast to SO$_2$ and OCS, the H$_2$S abundance is lower in B2 than in B1. This is consistent since H$_2$S is a parent molecule for other sulfured species. H$_2$S is formed in the grain surfaces and released to the gas phase by effect of the shock waves. Other sulfur-bearing molecules like SO and SO$_2$ are produced in gas phase very quickly (few 10$^3$ yr) via reactions with H, OH and O$_2$ \\citep{Pineau93, Wakelam05}. One possible explanation of the differences in B1 and B2 is that the B2 shock could be older than that in B1 and that H$_2$S has been converted into SO and SO$_2$. {Even if there are uncertainties regarding the sulfur chemistry \\citep[see for instance][and references therein]{Codella05}, sulfur-bearing molecules like H$_2$S, SO, SO$_2$, H$_2$CS and OCS have been invoked to be potential valuable probes of chemical evolution \\citep{Codella05, Wakelam05, Herpin09}. In particular, \\cite{Wakelam05} have found that the SO$_2$/SO ratio increases with the shock age. } The higher SO$_2$/SO ratio in B2 compared to B1 also points to an older shock in B2. \\cite{Bachiller97} have also proposed that the different CN abundance in B1 and B2 can be accounted for in this scenario of an older shock in B2, since after an enhancement of CN in the shock, reactions like CN+O $\\rightarrow$ CO+N could be very efficient to destroy CN. Our data globally agree with this scenario, nevertheless the clear anticorrelation of the CN and HNCO intensities and the similar abundances of CN and HNCO in L1157-B1 suggest that $all$ the CN can indeed be transformed into HNCO once the shock has increased the temperature and the neutral-neutral reactions form NCO from CN and O$_2$ and HNCO from NCO (see Sect. \\ref{sect:chemistry}). Therefore, the chemical differences from B1 and B2 do support the scenario proposed in Sect. \\ref{sect:chemistry} to explain the HNCO abundances in the L1157 shocks. In addition, the scheme proposed in Sect. \\ref{sect:chemistry} also gives insight on the possible link of HNCO and the sulfured molecules, in particular with SO and SO$_2$. In shocks these molecules can be produced by the reactions $\\mathrm{S+O_2} \\rightarrow \\mathrm{SO+O}$ and $\\mathrm{SO+OH}\\rightarrow \\mathrm{SO_2+OH}$ \\citep{Pineau93, Charnley97, Wakelam05}. Therefore, the common link with HNCO would be formation pathways involving O$_2$. Another important shock tracer that could be linked to the O$_2$ chemistry is SiO. Nevertheless, the situation regarding this molecule is more complex. First, recent models can explain the SiO emission in shocks without a significant contribution of Si oxidation in gas phase, either by sputtering in gas-grain collisions if there is already SiO in the grain mantles \\citep[][]{Gusdorf08b} or by dust vaporization in grain-grain collisions \\citep{Guillet09}. Second, if one assumes that SiO is formed in gas phase, there is a threshold shock velocity of $\\sim 25$ \\kms \\ in order to eject Si from the grain cores by sputtering \\citep{Gusdorf08a}. In contrast, other molecules as the organics are in the grain mantles and there is not such a high velocity threshold for them to be ejected to gas phase. { Therefore, both if SiO comes directly from the grains \\citep{Gusdorf08b, Guillet09} or if it is formed in gas phase \\citep{Gusdorf08a}, due to the shock velocity threshold, it is not expected to find similarities (and not observed) in the SiO emission and that of SO, SO$_2$ and HNCO. It is interesting to note that the impressive SiO blue wing in B1 \\citep{Bachiller97} suggests that SiO formation is indeed favored, in comparison to other molecules, in the higher velocity gas. In this context, chemical differences in B1 and B2 could also be due to different shock velocities. If the shock velocity in B2 is actually lower than in B1 and the SiO formation less efficient, more O$_2$ will remain available in B2 to form HNCO, SO and SO$_2$.} \\subsection{On the origin of the HNCO emission in galactic nuclei} The HNCO abundance in the molecular clouds in the center of the Milky Way and in the nuclei of starburst galaxies is similar or a bit higher (at most by a factor of 2) than in Galactic hot cores (Table \\ref{tab:comp}). HNCO becomes another piece in the well known puzzle of the chemistry of the Galactic center molecular clouds. This puzzle can be summarized as follows: the abundance of SiO and complex organic molecules in the Galacic center is as high as in hot cores \\citep{Martin-Pintado97, Rodriguez06, Requena06}. In contrast, ({\\it i}) the emission in the Galactic center is extended over the central 300 pc and it does not resemble a collection of discrete sources with the size of a hot core ($\\sim$ 0.1 pc), ({\\it ii}) the gas density ($10^3-10^4$ \\cmmt) is much lower than in hot cores ({\\it iii}) and the dust temperatures ($<30$ K) are also much lower than in hot cores. The Galactic center clouds present a ``hot core chemistry without hot cores\" \\citep{Requena06}. The origin of this chemistry is not known although it is thought to be due to some type of mechanical processes as shock waves \\citep{Martin-Pintado97, Martin-Pintado01, Rodriguez04}. The origin of the shocks can be related to the complex dynamics in the inner regions of the Galaxy \\citep{Huttemeister98, Rodriguez06, Rodriguez08}. Based on the spatial distribution of the HNCO and the comparison with other species, as CH$_3$OH, \\cite{Meier05} have also suggested that the HNCO emission in IC342 is tracing shocks. Recently, \\cite{Martin08, Martin09} have also proposed that shocks could be the explanation of the high HNCO abundances measured in galactic nuclei. Discussing the precise origin of shocks in Galactic nuclei is out of the scope of this paper. Nevertheless, our observations support the scenario of HNCO tracing shocks in galactic nuclei since our L1157 results probe a medium of moderate H$_2$ density where the HNCO abundace is indeed high due to the grain processing and gas heating by shock waves." }, "1003/1003.4733_arXiv.txt": { "abstract": "We present a physical model for origin of the cosmic diffuse infrared background (CDIRB). By utilizing the observed stellar mass function and its evolution as input to a semi-empirical model of galaxy formation, we isolate the physics driving diffuse IR emission. The model includes contributions from three primary sources of IR emission: steady-state star formation owing to isolated disk galaxies, interaction-driven bursts of star formation owing to close encounters and mergers, and obscured active galactic nuclei (AGN). We find that most of the CDIRB is produced by equal contributions from objects at $z\\sim 0.5-1$ and $z\\gsim 1$, as suggested by recent observations. Of those sources, the vast majority of the emission originates in systems with low to moderate IR luminosities ($L_{IR} \\lsim 10^{12}$ $L_\\odot$); the most luminous objects contribute significant flux only at high-redshifts ($z\\gsim 2$). All star formation in ongoing mergers accounts for $\\lsim 10\\%$ of the total at all wavelengths and redshifts, while emission directly attributable to the interaction-driven burst itself accounts for $\\lsim 5\\%$. We furthermore find that obscured AGN contribute $\\lsim 1-2\\%$ of the CDIRB at all wavelengths and redshifts, with a strong upper limit of less than $4\\%$ of the total emission. Finally, since electron-positron pair production interactions with the CDIRB represent the primary source of opacity to very high energy (VHE: $E_\\gamma \\gsim 1$ TeV) $\\gamma$-rays, the model provides predictions for the optical depth of the Universe to the most energetic photons. We find that these predictions agree with observations of high-energy cutoffs at $\\sim $ TeV energies in nearby blazars, and suggest that while the Universe is extremely optically thick at $\\gsim 10$ TeV, the next generation of VHE $\\gamma$-ray telescopes can reasonably expect detections from out to $\\sim 50-150$ Mpc. ", "introduction": "\\label{sec:intro} Diffuse extragalactic background light (EBL) represents the sum total of all the photons produced by luminous matter over the lifetime of the Universe. In the ultraviolet, optical, and near-infrared (IR) the EBL is directly attributable to star formation and active galactic nucleus (AGN) activity \\citep[for reviews, see][and references therein]{tyson1990,tyson1995,henry1991,henry1999,leinert1998}. However, in regions with significant dust opacity radiation generated by stars and AGN is reprocessed by dust into the IR as thermal emission \\citep[see, e.g.][]{soifer1991,sanders1996}, leading to the expectation of a significant EBL component at $\\lambda \\sim 10-1000\\mu$m. Measurements of this Cosmic Diffuse IR Backround (CDIRB) were notoriously difficult to obtain, owing primarily to the presence of significant zodiacal and galactic foregrounds and the lack of access to this region of the spectrum from the ground. The launch of the {\\it Cosmic Background Explorer} \\citep[COBE: for an overview, see][]{boggess1992} finally revealed a CDIRB comparable in brightness to the optical EBL, providing a complete census of obscured star formation and AGN activity across cosmic time \\citep{hauser1998,kelsall1998,aredt1998,dwek1998,fixsen1998,hauser2001}. Because the emission mechanisms that generate the CDIRB are intimately connected to galaxy formation and evolution, it provides a powerful observational constraint on models \\citep{partridge1967}. A particularly popular technique for modeling the CDIRB has been backwards evolution \\citep[e.g.,][]{rowanrobinson2001,rowanrobinson2009,lagache2003,lagache2004,lagache2005,xu2003,franceschini2008,finke2009b}, in which a parameterized fit to the evolution of the IR luminosity function -- using low-redshift observations as a baseline -- to reproduce a number of observables, including number counts and the CDIRB itself. A complementary approach has been to use semi-analytic models \\citep[SAMs:][]{cole1994,cole2000,somerville1999}, in which the hierarchical growth of structure is tracked by N-body cosmological simulations and baryonic physics -- e.g., star formation, radiative cooling, gas accretion, AGN activity, etc. -- are implemented as simple analytic prescriptions, tuned to match a predetermined set of (usually local) observational constraints. Mock galaxy catalogs generated by SAMs can then be combined with template spectral energy distributions (SEDs) to construct predictions for number counts at a given wavelength and the CDIRB more generally \\citep[e.g.,][]{primack1999,primack2005,primack2008,devriendt2000,baugh2005,swinbank2008}. Finally, cosmic chemical evolution (CCE) techniques \\citep[e.g.,][]{pei1995,pei1999} model the EBL from the optical through IR by solving for the self-consistent evolution of globally averaged quantities -- such as gas depletion and star formation -- in an analogous manner to galactic chemical evolution models \\citep[e.g.,][]{tinsley1980}. Each of these approaches has clear advantages and disadvantages. Backwards evolution models, while robust and successful in reproducing the observations, must assume simple evolutionary scalings. Therefore, they contain no physics, and cannot separate out the relative contributions to the CDIRB from different emission mechanisms. SAMs, by contrast, can determine the relative importance of, e.g. AGN versus star formation, but rely on simple analytic implementations that introduce parameter degeneracies and may not include all of the complexities in the relevant physics -- though some have sought to address this shortcoming by incorporating the results of hydrodynamical simulations of merger driven starbursts and AGN activity into the SAM framework \\citep{somerville2008}. Furthermore, SAMs have also faced difficulty simultaneously matching the observed galaxy mass function and cosmic star formation history \\citep{somerville2008}; it is no generally possible to separate these well-known issues from other model elements in driving their predictions for the CDIRB. Finally, while CCEs do not require a detailed implementation of these complicated baryonic physics, their global nature means they have little to say about specific populations of objects. An alternative approach, outlined in detail by \\citet{hopkins2007a,hopkins2007b} attempts a semi-empirical model of galaxy formation in the context of a merger-driven cosmic cycle: while isolated galaxies dominate the overall cosmic energy budget \\citep[however, see discussion in][]{hopkins2010.isolated}, major mergers of gas-rich disks drive an evolutionary sequence through a period of intense star formation and AGN activity, to produce passive elliptical galaxies \\citep{sanders1988a,hopkins2006}. The model connects measurements of the stellar mass function to the dark matter halo population through a halo occupation distribution (HOD) approach \\citep[e.g.,][]{peacock2000,scoccimarro2001,berlind2002,berlind2003,kravtsov2004,zehavi2004,zheng2005,zheng2009,brown2008,conroy2009}, and implements baryonic physics by incorporating the results of high-resolution hydrodynamical simulations. This approach has successfully reproduced the observed quasar luminosity function \\citep{hopkins2007a}, the growth of the red sequence \\citep{hopkins2007b}, the evolution of massive elliptical galaxies with redshift \\citep{hopkins2009.scale.evolve}, bulge-to-disk rations \\citep{hopkins2009.gasmorph}, and the IR luminosity function \\citep{hopkins2009.ulirg}. Semi-empirical modeling also represents a fundamental advance over pervious techniques for modeling the CDIRB; it is predictive, incorporates the relevant physics, includes the results of high-resolution hydrodynamical simulations where appropriate, while at the same time is constructed to match the galaxy mass function at all redshifts. Therefore, we can isolate the physics driving the production of the CDIRB without significant degeneracies model components and parameters. In this work, we use this semi-empirical framework to construct a physical model for the origin of the CDIRB. In addition to providing an important verification of this particular model of galaxy formation more generally, this approach will allow us to unfold the contributions from objects as a function of luminosity and redshift, as well as by emission mechanism -- including star formation in isolated disks, merger-driven starbursts, and AGN activity. Furthermore, because electron-positron pair production interactions ($\\gamma+\\gamma \\rightarrow e^+ + e^-$) with CDIRB photons is the primary source of attenuation of extragalactic, $\\approx$ TeV photons \\citep[often referred to as Very High Energy, or VHE $\\gamma$-rays;][]{jelly1966,fazio1966,fazio1970,gould1967,stecker1969,stecker1992}, model predictions for the evolution of this background with redshift yields a prediction for the opacity of the Universe to the most energetic photons. This work is organized as follows: in \\S~\\ref{sec:methods} we outline our methodology, in \\S~\\ref{sec:uncertainty} we summarize the primary sources of uncertainty in the model predictions, in \\S~\\ref{sec:comparison} we verify that the model is consistent will the relevant observational constraints, in \\S~\\ref{sec:assumptions} we investigate the result of changing some of the model assumptions, in \\S~\\ref{sec:interpret} we present predictions for the relative important of different populations of objects in generating the CDIRB, in \\S~\\ref{sec:gray} we make predictions for the $\\gamma$-ray opacity of the Universe, and in \\S~\\ref{sec:conclusion} we conclude. Throughout this work, we will make frequent reference to the total IR luminosity which is defined according to convention as the integrated luminosity from 8-1000\\micron \\citep[e.g.,][]{sanders1996}. Furthermore, we assume the most recent cosmological parameters from \\citet{komatsu2010}: ($\\Omega_m,\\Omega_\\Lambda,h) = (0.26,0.74,0.71)$, the WMAP 7-year mean. However, it is important to note that varying these within a reasonable range has no noticeable effect to our results. ", "conclusions": "\\label{sec:conclusion} We present a physical model for the origin of the CDIRB utilizing the semi-empirical framework of \\citet{hopkins2009.ulirg}. The model tracks three distinct sources of IR emission: steady-state star formation, interaction-induced starbursts, and obscured AGN activity. We also include all the relevant systematic uncertainties, which are dominated by the dynamic range in observational estimates of the stellar mass function (and increasing with redshift). The IRLFs generated by this model, combined with a library of template SEDs for starbursts \\citep{dale2001,chary2001,lagache2003} and obscured AGN \\citep{siebenmorgen1991,siebenmorgen1993,siebenmorgen1992a,siebenmorgen1992b,siebenmorgen2001,siebenmorgen2004a,siebenmorgen2004b}, provide an excellent match to observations of the CDIRB from $\\lambda_{obs} \\approx 10-1000\\mu$m. In contrast of alternative techniques -- including backwards evolution \\citep[e.g.,][]{rowanrobinson2001,rowanrobinson2009,lagache2003,lagache2004,lagache2005,xu2003,franceschini2008,finke2009b}, semi-analytic \\citep{primack1999,primack2005,primack2008,devriendt2000}, and CCE \\citep{pei1995,pei1999} models -- our approach provides a robust context in which to examine the physics its production. By varying the model assumptions, we find that the normalization and peak of the CDIRB emission are determined in large part by the strong redshift evolution of the gas content of steady-state star-forming galaxies. Without this evolution, the model vastly underpredicts the total intensity of CDIRB and peaks at shorter wavelengths than is observed, owing to the lack of a significant contribution form high-luminosity and high-redshift disks. Therefore, the observed CDIRB can be taken to reflect this strong evolution in the gas content of disks -- an effect which has been seen in numerous samples of individual objects \\citep[e.g.,][]{bell2001,kannappan2004,mcgaugh2005,shapley2005,daddi2009.fg,tacconi2010}. We can also use the model predictions to determine the relative importance of IR-luminous galaxies in producing the CDIRB as a function of redshift, luminosity, and emission mechanism. The model results indicate the following: \\begin{enumerate} \\item The CDIRB is primarily produced by equal contributions form objects at $z\\sim 0.5-1$ and $z\\gsim 1$, in agreement with recent observations by the BLAST experiment. However, in the observer's frame wavelengths past the peak contain an larger contribution from $z\\gsim 1$ than those at shorter wavelengths. \\item Most of the CDIRB is contributed by normal galaxies and LIRGs, though the ULIRG contribution becomes significant at high-redshift $z\\gsim 2$. This also leads to a larger contribution from ULIRGs at wavelengths past the peak. \\item Ongoing mergers contribute less than 10\\% of the total emission at all wavelengths, with less than $\\sim 1-3\\%$ produced by the merger-driven burst itself. \\item Obscured AGN account for $\\lsim 1\\%$ of the CDIRB at all wavelengths, with a strong upper limit at 4\\% of the total emission. \\end{enumerate} In the future, these predictions can be tested via deep, wide surveys with the {\\it Herschel Space Telescope} (e.g., the Herschel Multi-tiered Extragalactic Survey, or HerMES)\\footnote{{\\tt http://astronomy.sussex.ac.uk/$\\sim$sjo/Hermes/}} combined with stacking analyses similar to those presented by \\citet{devlin2009} and \\citet{marsden2009}. The CDIRB also represents the primary source of opacity for VHE $\\gamma$-rays, owing to electron-positron pair production interactions \\citep[$\\gamma+\\gamma \\rightarrow e^+ + e^-$;][]{jelly1966,fazio1966,fazio1970,gould1967,stecker1969,stecker1992}. Thus, our results also provide predictions for the opacity of the Universe to the most energetic photons. We find that the model predictions are consistent with high-energy cutoffs for TeV sources at $z\\sim 0.03$ \\citep{aharonian1999,aharonian2002a}. They also indicate that while the Universe is highly opaque to $E_\\gamma \\gsim 10$ TeV photons ($\\tau_{\\gamma\\gamma} \\gsim 10$ at $z\\approx 0.06$), with the next generation of ICATs we can reasonably expect to detect sources out to $\\sim 50-150$ Mpc. \\subsection*{Acknowledgements} Thanks to Michael Kuhlen, Desika Narayanan, and Chris Hayward for helpful converstaions. JDY acknowledges support from NASA through Hubble Fellowship grant \\#HF-51266.01 awarded by the Space Telescope Science Institute,which is operated by the Association of Universities for Research in Astronomy, Inc., for NASA, under contract NAS 5-26555. Support for PFH was provided by the Miller Institute for Basic Research in Science, University of California Berkeley. The computations in this paper were run on the Odyssey cluster supported by the FAS Research Computing Group at Harvard University." }, "1003/1003.3055_arXiv.txt": { "abstract": "We show that the cosmological phase transition from the first accelerated expansion in the early universe to the second accelerated expansion over the intermediate decelerated expansion is possible in the HL gravity without the ``detailed balance'' condition if the \\emph{dark scalar} energy density is assumed to be negative. Moreover, we obtain various evolutions depending on the scale factor and the expansion rate. Finally, we discuss the existence of the minimum scale in connection with the singularity free condition. ", "introduction": "Recently, the Ho\\v{r}ava-Lifshitz (HL) gravity has been proposed as an ultraviolet (UV) completion of general relativity~\\cite{horava}, motivated by the Lifshitz theory in the condensed matter physics~\\cite{lifshitz}. The key of the UV completion is the anisotropic scaling between space and time, \\begin{equation} \\label{scaling} t \\to b^{z}\\, t, \\qquad x^i \\to b\\, x^i, \\end{equation} where the Lifshitz parameter $z$ becomes 1 in the infrared (IR) limit. In general, the HL theory is not invariant under the full diffeomorphism group of general relativity but under its subgroup, called the foliation-preserving diffeomorphism; however, in the IR limit, the full diffeomorphism is somehow recovered. A mechanism for recovering the full diffeomorphism or the renormalization group flow is yet unsolved issue, but the HL gravity has been intensively studied in the area of cosmology~\\cite{ts,kk,mukohyama,brandenberger,ks,ww,mipark,ls,sj} and black hole physics~\\cite{lmp,cco,cy,ysmyung,mann,gh,kk:bh,majhi}. Considering Arnowitt-Deser-Misner (ADM) decomposition~ of the metric with $ds^2 = - N^2 c^2 dt^2 + g_{ij} (dx^i + N^i dt) (dx^j + N^j dt)$~\\cite{adm}, the Einstein-Hilbert action can be rewritten as \\begin{equation} \\label{act:EH} \\begin{aligned} I_{EH} &= \\frac{c^3}{16\\pi G_N} \\int d^4x \\sqrt{-\\mathcal{G}} \\left[ \\mathcal{R} - 2\\Lambda \\right] \\\\ &= \\frac{c^2}{16\\pi G_N} \\int dt d^3x \\sqrt{g} N \\left[ K_{ij} K^{ij} - K^2 + c^2 \\left( R - 2\\Lambda \\right) \\right], \\end{aligned} \\end{equation} where $K_{ij} \\equiv \\frac{1}{2N} \\left[ \\dot{g}_{ij} - \\nabla_i N_j - \\nabla_j N_i \\right]$ is the extrinsic curvature of $t= \\text{constant}$ hyper-surface, and the dot denotes the derivative with respect to $t$. Here, $g_{ij}$, $R$, and $\\nabla_i$ are the metric, the intrinsic curvature, and the covariant derivative in the three-dimensional hyper-surface, respectively. Note that the action~\\eqref{act:EH} can be regarded as one of the $z=1$ HL theory and each terms are invariant under the foliation-preserving diffeomorphism. Then, the first two terms are referred to as kinetic terms, while the other two are potential terms. In particular, to study the HL gravity, higher-order potential terms, for instance, $R^3$, will be taken into account. But there are almost ten possible terms for $z=3$ case, and the ``detailed balance'' condition can reduce the ten coefficients to three effective ones~\\cite{horava}. Of course, it also implies that $(D+1)$-dimensional renormalization can be reduced to the simpler $D$-dimensional renormalization. By the way, it has been claimed that matter is not UV stable with this condition~\\cite{calcagni}, and the HL theory with the detailed balance condition does not have the Minkowski vacuum solution. On the other hand, there have been interesting studies on the cosmological evolution in the HL gravity based on the phase space analysis~\\cite{ls}. Authors showed that the cosmological phase changing from the decelerated expansion to the accelerated expansion is possible assuming the detailed balance condition and the constant equation-of-state parameter. However, the first accelerated expansion corresponding to the inflationary era has not been discussed. We believe that there should be such a structure if the HL theory is indeed the UV completion. By the way, the HL gravity itself may be problematic, for instance, counting its degrees of freedom~\\cite{lp,hkg}, description of the asymptotically flat spacetime~\\cite{ks,lmp}, and so on. However, even in spite of these problems, it deserves to study various aspects since it may give some insight to understand the quantum gravity. So, we would like to investigate whether the smooth cosmological phase transition from the first accelerated expansion in the early universe to the second accelerated expansion over the intermediate decelerated expansion is possible or not. In section~\\ref{sec:hl:cos}, we recapitulate the Ho\\v{r}ava-Lifshitz cosmology and define some relevant quantities without the ``detailed balance'' condition. Then, the nonlinear equations of motion will be solved using nonlinear methods in order to study the cosmological behavior in section~\\ref{sec:hl:nl}. By considering the observational data on density parameters, we shall plot the phase portrait in section~\\ref{sec:hl:pp}, and discuss the behavior in the early stage of our universe and its various destiny. If a dark scalar energy density is assumed to be negative, then the desired first accelerated expansion appears. Moreover, we will show that it can be smoothly connected with the second accelerated expansion. Finally, discussions will be given in section~\\ref{sec:dis}. ", "conclusions": "\\label{sec:dis} We have studied the HL gravity coupled to the matter without the detailed balance condition in order to show the possibility to get the smooth phase transition from the first accelerated expansion corresponding to the early stage of the universe to the second accelerated expansion throughout the intermediate decelerated expansion assuming the energy density for the dark scalar to be negative $\\rho_\\text{ds}<0$. In spite of this negative energy contribution, the total energy density $\\rho_\\text{tot}$ is positive at any cosmological scale. Note that there have been researches which provide concrete justifications for models with negative density, in particular, a brane universe moving in a curved higher dimensional bulk space~\\cite{kk:mirage} and a model of dark energy stemming from a fermionic condensate~\\cite{abc}. Of course, we can confirm that there is no such a phase transition unless the dark scalar density is assumed to be negative. For instance, if we take the detailed balance condition, then it simply reduces the additional energy contribution to $\\rho_\\text{vac} = \\Lambda c^4 / 8\\pi G_N$, $\\rho_k = -3c^4k / 8\\pi G_N a^2$, $\\rho_\\text{dr} = 27c^4k^2 / 32\\pi G_N \\Lambda a^4$, $\\rho_\\text{ds} = 0$ as seen from fundamental constants~\\eqref{fund:const} and redefinition of coefficients~\\eqref{coeff:dbc}. In this case, the first accelerated expansion does not appear, that is the reason why we did not take the detailed balance condition. In addition, it has been shown that our universe may be oscillatory even with the same density parameters of the current observation, depending on the expansion rate. Actually, the expansion rate is related to the total energy density by Eq.~\\eqref{rel:conserv}, so that small expansion rate $v$ corresponds to small $h\\sim\\rho_\\text{tot}$ for a given scale factor. On the other hand, it has been well known that if there were $\\sim60$ e-foldings of inflationary expansion, then the universe was driven to be nearly flat, removing any need for fine-tuned initial conditions. In addition, the inflationary expansion would have driven the density of magnetic monopoles to be negligible today, explaining their apparent absence. Also, our entire observable portion of the universe would have inflated from an initially small causally connected region, thereby ensuring a high degree of isotropy today. Although our analysis shows that the first accelerated expansion can be obtained by assuming $\\rho_\\text{ds}<0$, unfortunately, it should be pointed out that it does not last for 60 e-foldings. We hope this problem will be discussed elsewhere. It is interesting to note that there exists a minimum scale $a_\\text{min}$ obtained from $\\rho_c\\ge0$ and $\\rho_\\text{tot}\\ge0$ with the relation $\\rho_c=\\rho_\\text{tot}+\\rho_k$. Note that $\\rho_\\text{tot}(a_\\text{min})=-\\rho_k(a_\\text{min})>0$ and $H^2(a_\\text{min})\\sim\\rho_c(a_\\text{min})=0$ for $k=+1$, so the universe starts with a certain amount of energy and zero expansion rate, while $\\rho_\\text{tot}(a_\\text{min})=0$ and $H^2(a_\\text{min})\\sim \\rho_c(a_\\text{min})=\\rho_k(a_\\text{min})>0$ for $k=-1$, so the universe starts with zero energy density and a certain expansion rate. For any case, as far as we consider the positive total energy density and the critical energy density, there should be no initial singularity problem because of the existence of the minimum scale. The final comment is in order. The energy density and the pressure of the dark scalar which play a crucial role in this work depend on the three terms in the potential~\\eqref{act:pot:3+1}, $\\sigma_1 R^3$, $\\sigma_2 R R_{ij} R^{ij}$, $\\sigma_3 R_i^j R_j^k R_k^i$. One may think that three terms give rise to kinetic contributions so that they provide the leading spatial dependence in the gravition propagator in the UV region. What it means is that graviton modes may be unstable. However, this is not the case since they do not modify the gravition propagator. In general, the lowest metric order in curvature tensors and curvature scalars is linear, which means that the three terms are not quadratic in metric. Of course, such terms which are quadratic in curvature, $\\sigma_4 \\nabla_i R \\nabla^i R$, $\\sigma_5 \\nabla_i R_{jk} \\nabla^i R^{jk}$ in Eq.~\\eqref{act:pot:3+1}, will not only add interaction but also modify the propagator; however, these are irrelvant to our energy density formulae~\\eqref{add:en}." }, "1003/1003.4991_arXiv.txt": { "abstract": "In this brief report, we summarize our recent studies in brane cosmology in both string theory and M-Theory on $S^{1}/Z_{2}$. In such setups, we find that the radion is stable and its mass, with a very conservative estimation, can be of the order of $0. 1 \\sim 0.01$ GeV. The hierarchy problem can be addressed by combining the large extra dimension, warped factor, and tension coupling mechanisms. Gravity is localized on the visible brane, and the spectrum of the gravitational Kaluza-Klein (KK) modes is discrete and can have a mass gap of TeV. The corrections to the 4D Newtonian potential from the higher order gravitational KK modes are exponentially suppressed. Applying such setups to cosmology, we find that a late transient acceleration of the universe seems to be the generic feature of the theory, due to the interaction between branes and bulk. A bouncing early universe is also rather easily realized. ", "introduction": " ", "conclusions": "" }, "1003/1003.0056_arXiv.txt": { "abstract": "We study the statistical nature of primordial fluctuations from an anisotropic inflation which is realized by a vector field coupled to an inflaton. We find a suitable gauge, which we call the canonical gauge, for anisotropic inflation by generalizing the flat slicing gauge in conventional isotropic inflation. Using the canonical gauge, we reveal the structure of the couplings between curvature perturbations, vector waves, and gravitational waves. We identify two sources of anisotropy, i.e. the anisotropy due to the anisotropic expansion of the universe and that due to the anisotropic couplings among variables. It turns out that the latter effect is dominant. Since the coupling between the curvature perturbations and vector waves is the strongest one, the statistical anisotropy in the curvature perturbations is larger than that in gravitational waves. We find the cross correlation between the curvature perturbations and gravitational waves which never occurs in conventional inflation. We also find the linear polarization of gravitational waves. Finally, we discuss cosmological implication of our results. ", "introduction": "The primordial fluctuations from inflation is supposed to be statistically isotropic, Gaussian, and scale invariant. The nature of fluctuations is associated with the nature of de Sitter spacetime. However, since the expansion during inflation is not exactly de Sitter, the power spectrum is slightly tilted by the order of the slow roll parameter~\\cite{Komatsu:2010fb} which characterizes the deviation of the expansion from the exact de Sitter expansion. The deviation from the Gaussianity is also known to be related to the slow roll parameter~\\cite{Maldacena:2002vr}. On the other hand, the statistical isotropy has been regarded as a robust prediction so far because the cosmic no-hair conjecture is thought to be robust~\\cite{Wald:1983ky}. From an observational point of view, there are various indications that there exists statistical anisotropy in the cosmic microwave background radiation (CMB)~\\cite{Eriksen:2003db}. Although the statistical significance of these anomalies is still under debate, the possibility of the statistical anisotropy certainly deserves further theoretical investigation~\\cite{Gordon:2005ai}. Recently, breaking the statistical isotropy through the vector fields in an inflationary universe is proposed in the paper~\\cite{Yokoyama:2008xw} and extended in various ways~\\cite{Dimopoulos:2009vu, Dimastrogiovanni:2010sm,ValenzuelaToledo:2009af}. However, if the vector field is relevant to inflation, it may also produce anisotropy in an inflationary universe whatever small it is, which seems to contradict the cosmic no-hair conjecture. From the above perspective, it is interesting to ask if it is possible to have anisotropic inflationary universe~\\cite{Ford:1989me,Kaloper:1991rw, Kawai:1998bn,Barrow:2005qv,Barrow:2009gx,Campanelli:2009tk,Golovnev:2008cf, Kanno:2008gn}. If possible, it provides a simple mechanism to break the statistical isotropy by breaking the isotropy of the spacetime~\\cite{Ackerman:2007nb}. In the light of no-hair conjecture~\\cite{Wald:1983ky}, one may deny this possibility. In fact, many attempts to construct anisotropic inflationary models suffer from the instability~\\cite{Himmetoglu:2008zp}. However, recently, stable anisotropic inflationary models are found for the first time~\\cite{Watanabe:2009ct,Kanno:2009ei}. This can be regarded as a counter example to the cosmic no-hair conjecture. The interesting point is that the deviation from isotropy is related to the slow roll parameter, namely, the deviation from the exact deSitter expansion. Of course, that means the degree of the anisotropy is quite small. From the point of view of precision cosmology, however, it is worth exploring theoretical fine structure in an inflationary scenario. In this paper, we study cosmological perturbations in an anisotropic inflationary scenario we have found. The expected phenomenology of the anisotropic inflation is as follows: \\begin{itemize} \\item There should be statistical anisotropy in curvature perturbations. \\item There should be statistical anisotropy in gravitational waves. \\item There should exist the cross correlation between curvature perturbations and gravitational waves. \\item There should be linear polarization of gravitational waves. \\end{itemize} The first item will be tested by the PLANCK~\\cite{Pullen:2007tu}. The second one may be detected through B-mode polarization in the CMB~\\cite{Baumann:2008aq}. The third one will imply T-B correlation in CMB~\\cite{Gluscevic:2010vv}. The last one could be important for the future direct measurement of gravitational waves through the interferometer~\\cite{Seto:2001qf}. The purpose of this paper is to calculate the above quantities numerically and analytically and reveal the physics behind them. Since the spacetime is anisotropic, the formalism treating perturbations is non-standard. Although there are many works treating the cosmological perturbations in an anisotropic universe~\\cite{Tomita:1985me,Dunsby:1993fg,Noh:1987vk, Pereira:2007yy,Gumrukcuoglu:2007bx,Himmetoglu:2009mk}, there have been several obstructions in extracting concrete predictions for CMB. The main obstruction was the lack of the concrete anisotropic cosmological models. Now, since we have such models, we have succeeded in obtaining concrete results by utilizing the canonical gauge which is a generalization of the flat slicing in the conventional isotropic inflationary scenario. Recently, during our slow preparation of this paper, two papers have appeared on the archive~\\cite{Dulaney:2010sq,Gumrukcuoglu:2010yc}. The first one \\cite{Dulaney:2010sq} studied the primordial perturbations in an anisotropic inflationary universe using a perturbative method. The second one \\cite{Gumrukcuoglu:2010yc} investigated the same issue numerically. The conclusion is quite similar to ours. The main difference is the gauge used in analysis. Our canonical gauge allows us to reveal the nature of primordial fluctuations from anisotropic inflation in a transparent way. Interestingly, on the contrary to a naive expectation, all of these works including ours imply that even if the anisotropy of the universe is very small, a large statistical anisotropy in the spectrum of curvature perturbations could be created. The organization of this paper is as follows: In section II, we review an anisotropic inflation which is caused by the inflaton coupled to the vector field. Here, we will see the anisotropy is determined by the slow roll parameter. In section III, we choose the canonical gauge which is a generalization of the flat slicing in the conventional inflation and classify perturbations in anisotropic universe based on the 2-dimensional rotation symmetry. Then, we obtain the quadratic action for perturbed quantities. In section IV, we reduce the action to that for physical variables from which we can read off the structure of couplings between those variables. Based on the reduced action, we calculate various statistical quantities numerically and analytically to reveal the nature of primordial fluctuations in anisotropic inflation. In section V, we discuss cosmological implication of our results. The final section is devoted to the conclusion. In the Appendix A, we provide a detailed derivation of the action for 2-dimensional scalar sector perturbations. ", "conclusions": "We have studied the statistical nature of primordial fluctuations from an anisotropic inflation which is realized by a vector field coupled to an inflaton. First, we have classified metric fluctuations according to the 2-dimensional rotational symmetry. To choose a convenient gauge in an anisotropic universe, we have started from the flat slicing gauge in an isotropic universe and made an appropriate gauge transformation to get a canonical gauge in an anisotropic universe. This gauge choice has made the subsequent analysis and the interpretation of the variables easier. Using the canonical gauge, we have revealed the structure of the couplings between curvature perturbations, vector waves, and gravitational waves. We found that there are two sources for anisotropy, i.e. the anisotropy due to the anisotropic expansion of the universe and that due to the anisotropic couplings among variables. It turned out that the latter effect is dominant. We have numerically obtained power spectra. We also presented analytical formula using in-in formalism. Since the coupling between the curvature perturbations and vector waves is the strongest one, the anisotropy in the curvature perturbations is larger than that in gravitational waves. More interestingly, we found the cross correlation between curvature perturbations and gravitational waves which is peculiar to anisotropic inflation. We also found the linear polarization of gravitational waves. Although there are several mechanism to produce circular polarization in the primordial gravitational waves~\\cite{Lue:1998mq}, this is the first example which realized the linear polarization in the primordial gravitational waves. We have only considered power spectrum for simplicity. However, as is pointed out in the paper~\\cite{Yokoyama:2008xw}, the statistical anisotropy could appear in the non-Gaussianity strongly and modify the shape of the bispectrum and trispectrum. Hence, it is interesting to study non-Gaussianity in anisotropic inflation models. We can extend anisotropic inflation in various ways. Although we have investigated a chaotic inflation in this paper, it is easy to extend the analysis to other inflation models. It is possible to incorporate multi-vector fields. From the string theory point of view, it is intriguing to consider anti-symmetric tensor field. It is also interesting to consider other Bianchi type models~\\cite{Dechant:2008pb} in the context of anisotropic inflation." }, "1003/1003.4308.txt": { "abstract": "We present results from a wide-field imaging campaign at the Canada-France-Hawaii Telescope to study the spectacular outburst of comet 17P/Holmes in late 2007. Using image-processing techniques we probe inside the spherical dust coma and find sixteen fragments having both spatial distribution and kinematics consistent with isotropic ejection from the nucleus. Photometry of the fragments is inconsistent with scattering from monolithic, inert bodies. Instead, each detected fragment appears to be an active cometesimal producing its own dust coma. By scaling from the coma of the primary nucleus of 17P/Holmes, assumed to be 1.7 km in radius, we infer that the sixteen fragments have maximum effective radii between $\\sim$ 10 m and $\\sim$ 100 m on UT 2007 Nov.\\ 6. The fragments subsequently fade at a common rate of $\\sim$ 0.2 mag day$^{-1}$, consistent with steady depletion of ices from these bodies in the heat of the Sun. Our characterization of the fragments supports the hypothesis that a large piece of material broke away from the nucleus and crumbled, expelling smaller, icy shards into the larger dust coma around the nucleus. ", "introduction": "Comet 17P/Holmes is a dynamically and compositionally typical Jupiter Family Comet (Schleicher 2009) but it has exhibited three dramatic outbursts that caused an increase in brightness large enough to lift it from obscurity to naked-eye visibility (Holmes 1892; Palisa 1893; Buzzi et al. 2007). The first outburst led to its discovery on UT 1892 Nov.\\ 6 by Edwin Holmes (Holmes 1892) and was followed by a second outburst three months later in January 1893. The third outburst, first identified by J.\\ A.\\ Henriquez Santana on UT 2007 Oct.\\ 24 (Buzzi et al. 2007), caused the comet to reach a brightness of 2$^{nd}$ magnitude. %be more specific, magnitude changes for first outburst [ref]. could talk about observations from 1st explosion...? The nature of cometary mass loss varies widely between comets, ranging from gentle outgassing to violent outbursts as observed in the case of 17P/Holmes. Possible causes of large outbursts are numerous but in this case we are able to rule out several. The 1892 and 1893 outbursts of 17P/Holmes were attributed to impacts with a satellite (Whipple 1984) but this possibility is rendered extremely unlikely by a third, similar outburst 115 years later. Rotational breakup requires a rotation period of less than 5.2 hours (assuming a spherical, strengthless body with a density of 400 kg m$^{-3}$; Richardson \\& Melosh 2006). Work by Snodgrass et al. (2006), while not revealing a definitive rotation period, suggests a value several times longer. Tidal breakup is implausible given the position of 17P/Holmes (far from any planet or the Sun) at the time of outburst. A possible trigger for the outburst is a decrease in the perihelion distance from 2.16 AU to 2.05 AU caused by a close approach to Jupiter in January 2004, resulting in an increase in solar insolation (but only by $\\sim$ 10\\%) to greater depths in the comet's interior. However, the detailed mechanism by which an increase in insolation might lead to the observed outburst remains unknown. %can you explain the 5 month delay from perihelion to the time of outburst? ... time for heat to build up. In this paper we present a set of coordinated observations taken at the %University of Hawaii 2.2m telescope and Canada-France-Hawaii Telescope (CFHT) in a program designed to monitor the development of the coma in outburst. A major result is the discovery of multiple sub-nuclei ejected from Comet 17P/Holmes during the October 2007 outburst. We discuss their dynamical and physical characteristics and the constraints placed by their existence on the outburst mechanism. ", "conclusions": "\\subsection{\\label{subsec:phasespace} Phase Space Distribution of Fragments} The position angle of the fragments measured from the nucleus appears uniformly distributed (Figure \\ref{fig:lumpxy}). This suggests that the true three--dimensional distribution is either spherical, or a cone with its axis along the line of sight. %The latter appears less probable, but remains a possibility because %the line of sight direction is closely aligned with the anti--solar %direction. In a radially expanding system, we expect the true and projected positional radii and radial velocities to be perfectly correlated. Comparing both to the same theoretical projected distribution thereby provides an independent validation of the data. An important caveat is that such tests are sensitive to the completeness of the sample. For example, if the manual peak-finding procedure missed slow moving fragments near the nucleus, we may understate the level of central concentration. In Table \\ref{table:dismodels}, we compare the projected distribution of the radii and velocities of the fragments with various three dimensional models. We consider 1) a spherical distribution of fragments; 2) a model in which fragments lie on an infinitely thin shell; 3,4) finitely thick hollow shells (or hollow spheres) of fragments that have 20\\% and 50\\% of the thickness of the shell's radius; 5,6) radially symmetric space-filling distributions of fragments with $r^{-1}$ and $r^{-2}$ number density profiles; 7) a line of sight hollow cone of fragments; and 8) a line of sight solid cone of fragments. For the positions, we use the median radius and time of the first three nights, and for the velocities, we use only the radial component of the best--estimate median velocity, under the assumption that any transverse component is noise. We consider a set of model distributions consisting of hollow spherical shells of various thicknesses, filled spheres, and filled and hollow cones. We use a Kolmogorov-Smirnov (KS) test to compare our projected $R$ and $v$ values with the distribution predicted by each model. This is a slight misuse of the KS test, because we fix our outermost point to be at a cumulative probability of 1. However, our interest is in ruling out models, and this effect will tend to make all models agree better with our data. Table \\ref{table:dismodels} shows the KS agreement of the distribution of the data with the models. We can rule out the thin shell and the 20\\% shell (in which the fragments occupy a shell of thickness equal to 20\\% of the radius) on the basis of both the radial data and the velocity data. The 50\\% thick shell, solid sphere, and, to a lesser extent, $r^{-1}$ models are compatible with the data. If the eruption is conical rather than spherical, then an edge-enhanced cone is preferred over one that is solid. In cases in which the radial data disagree with the velocity data, like the 20\\% spherical shell and the solid cone, we are inclined to believe that the radial data are more robust. These results must be interpreted taking into account the completeness caveat mentioned above. If we assume that fragments closer than $0.3\\times R_{\\rm max}$, where $R_{\\rm max}$ is the projected radius of the largest fragment, are invisible to us, then the shell-like distributions become less likely, but the centrally concentrated $r^{-2}$ density model and the solid cone are no longer ruled out. At most, we can claim that the fragment distribution is not concentrated at the peripheries, but we cannot rule out a strongly centralized arrangement. Our conclusion is that the spatial and radial velocity distributions of the fragments in the sky-plane are consistent with isotropic ejection or conical ejection centered around the line of sight, or very close to it. It is difficult to envisage a plausible scenario in which 16 fragments are ejected isotropically from the nucleus, without catastrophic disruption of the nucleus. One mechanism has been suggested by Samarasinha (2001) who proposed that small nuclei ($\\sim$~1~km) could contain connected voids that allow sublimated supervolatiles to move rapidly through the nucleus. Assuming these voids have no outlet to the surface, internal gas pressure could build up until it exceeds the tensile strength of the mantle. At this point, an outburst could occur over a large fraction of the nucleus' surface. A difficulty with this model is the very small tensile strength of the cometary nucleus, which will prevent the build-up of high internal pressures. % == talk about cone versus isotropy. why cone might be better. isotropy suggested by samarasinha but results in catastrophic disruption of the nucleus - not observed here. %% {\\bf NOTE:} Isn't an edge enhanced cone (vs filled cone) the %% consequence of an impact, which is unlikely, leaving only a %% filled-spherical model as a likely remaining contender? Also should %% note that a hemispherical erupution is indistinguishable from a sphere, %% and represents the meeting point between a cone and an sphere. %% \\subsection{Velocities of lumps} %% \\begin{itemize} %% \\item lumps travel at 100 $m/s$ 5000 Joules/kg - latent heat of %% fusion is 384,000 Joules/kg, s amorphous to crystalline %% transformation has the energy (but no mechanism) %% \\item Suppose you have a 10x10 cm 1 kg chunk. To accelerate it to %% 100 m/s you meed momentum $dP=100 \\rm kg m/s$. A one second %% force of $100 \\rm kg/m/s^2$ would do it. This force could be %% generated by a pressure of $10^4$ pascal over 0.01 square %% meter. Such a pressure corresponds to 0.01 atomospheres. %% \\end{itemize} %photometry \\subsection{\\label{subsec:sizeestimates} Size Estimates} At the earliest detection in our data set, the fragments have magnitudes ranging from r' = 17.6 to 23.5 (Table \\ref{table:megafrag}). We consider two models to estimate the sizes of the fragments from the magnitudes. The apparent magnitude of a monolithic body is related to the viewing geometry and the body's physical characteristics according to the relation \\begin{equation} g_{\\lambda} \\Phi_{\\alpha} C = 2.25\\times10^{22} R^{2} \\Delta^{2} \\pi 10^{0.4(m_{\\sun}-m_{\\lambda})} \\end{equation} \\noindent where $g_{\\lambda}$ is the geometric albedo, $\\Phi_{\\alpha}$ is a function to account for the variation of brightness of the body with phase angle, $C$ [m$^{2}$] is the geometric cross-section of the body, $R$ [AU] and $\\Delta$ [AU] are the heliocentric and geocentric distances, and $m_{\\sun}$ and $m_{\\lambda}$ are the apparent magnitudes of the Sun and the body respectively (Jewitt 1991). We use a linear approximation for the phase function and set $\\Phi_{\\alpha}$ = 10$^{-0.4\\alpha\\beta}$, where $\\alpha$ [deg] is the phase angle and $\\beta$ [mag deg$^{-1}$] is the phase-coefficient. We assume a value of 0.035 mag deg$^{-1}$ for the phase-coefficient (Lamy et al. 2004) and $m_{\\sun}$ = -26.95 mag when using the SDSS r' filter (Ivezi\\'{c} et al. 2001). %magnitude taken from SDSS calibration (see Ivezic et al. 2001). obtained from Vsun=-26.75, (B-V)sun=0.65 (Allen 1973) w/ aid of photometric transform from Fukugita et al. 1996. Model A: If the detected fragments are monolithic, spherical bodies with geometric albedos of 0.1, we infer that the median radius of a fragment is 1.79 km, with sizes ranging from 0.8 km to 3.0 km (Table~\\ref{table:megafrag}). Since the radius of the parent nucleus is only $\\sim$ 1.7 km (Lamy et al. 2000;, Snodgrass et al. 2006) this interpretation can be rejected. Increasing the albedo to 0.15 yields a range of radii from 0.6 km m to 2.4 km with the median radius being $\\sim$ 1.5 km. We conclude that it is unlikely that the fragments are bare nuclei, and instead proceed to consider the possibility that they are actively-outgassing sub-nuclei. Model B: Using the complementary set of 5 second exposures obtained on the same nights at CFHT, we measured the brightness of the unsaturated nucleus, without coma-subtraction. We find that the apparent magnitude corresponds to an effective radius of $\\sim$ 330 km, demonstrating that the dust coma around the nucleus dominates the scattering cross-section, as it appears to do for the fragments we have discovered. Thus, the magnitude of the fragment ($m_{f}$) or nucleus ($m_{n}$) depends primarily on the amount of dust present in the aperture. Assuming that the nucleus and the fragments have material sublimating from active regions that cover similar fractions of their surfaces, then the difference of the observed magnitudes is proportional to the ratio of their surface areas, or their radii squared if we assume spherical bodies: \\begin{equation} 10^{-0.4(m_{\\rm f}-m_{\\rm n})} = \\frac{R^{2}_{\\rm f}}{R^{2}_{\\rm n}} \\end{equation} where $R_{f}$ and $R_{n}$ are the radii of the fragment and the nucleus, respectively. Using this scaling argument, and apparent magnitudes of the fragments and nucleus determined using $2^{\\prime\\prime}.22$ circular apertures as described in Section 3.1, we obtain fragment radii between 10 m and 110 m on the first night of detection (Table~\\ref{table:megafrag}). Cometary nuclei typically have active regions that cover only a small percentage of the surface (A'Hearn et al. 1995; Jewitt 2004), due to the gradual formation of an inert mantle by irradiation, micro-meteorite bombardment and loss of volatiles. It is possible that the fragments were more active than the nucleus. They may have been rotating rapidly and exposing much of their surface to sunlight, or composed mainly of volatile material with little or no mantle. Thus, the sizes derived here should be considered as upper limits. With a density of 400 kg m$^{-3}$ (Richardson \\& Melosh) and the scaled radii listed in Table \\ref{table:megafrag} for Model B we find that the sixteen fragments have a combined mass of 10$^{10}$ kg, corresponding to $\\sim$ 0.1 $\\%$ of the mass of a 1.7 km radius, spherical nucleus. Again, this is an upper limit to the mass in the fragments, and shows that the outburst of 17P/Holmes ejected only a tiny fraction of the total nucleus mass. %\\citep{2008ICQ....30....3S}. \\subsection{Acceleration} We do not detect any systematic acceleration of the fragments between 2007 Nov. 6 and 2007 Nov. 14 UT, since a single mean velocity over the observational dataset predicts a time of ejection that agrees with the published eruption time (Wilkening et al. 2007). This suggests two things: firstly, that radiation pressure does not significantly affect the motion of the fragments and, secondly, that the fragments are not self-propelled by directional sublimation of volatiles. The first point suggests that the fragments are macroscopic, as opposed to clusters of micron-sized particles, or smaller, that would be easily accelerated in the anti-solar direction by radiation pressure. The second constrains the nature of the fragments. Given that these fragments are volatile-rich and actively outgassing (as demonstrated in \\S\\ref{subsec:sizeestimates}), one may expect self-propulsion to accelerate the fragments in the anti-solar direction, as volatiles would be typically expelled in the sunward-direction. However, the fragments are likely to be spinning rapidly and may be isothermal, resulting in sublimation in all directions, and hence no net acceleration. Thus, the fragments are observed to continue in the directions in which they were ejected, with no noticeable increase in velocity. \\subsection{Correlation of Flux with Radius} If the observed surface brightness maxima originate from discrete solid fragments expelled by gas pressure, we expect lighter fragments to be launched faster, and so to appear at larger radii. Specifically, if we make the assumption that the observed fragments are icy fragments of size $\\ell_{\\rm f}$, density $\\rho_{\\rm f}$ and mass $m_{\\rm f}=\\rho_{\\rm f} \\ell_{\\rm f}^3$, and are ejected by gas pressure $P$ acting over a fixed acceleration distance $d$, then the energy transmitted to each fragment is ${1\\over2} m_{\\rm f} v_{\\rm f}^2= P \\ell_{\\rm f}^2 d$, and each fragment's distance to the nucleus $r_{\\rm f}$ is given by $r_{\\rm f} \\propto v_{\\rm f} \\propto \\ell_{\\rm f}^{-1/2}$, where $v_{\\rm f}$ is the fragment's velocity. Assuming that a fragment's brightness is given by its (sublimating) surface area $\\ell_{\\rm f}^2$, the expected relationship of $r_{\\rm f}$ to photon flux $f_{\\rm f}$ is then $r_{\\rm f}\\propto f_{\\rm f}^{-1/4}$. Alternatively, if we change our assumptions so that the gas pressure acts for a fixed time $t$ instead of a fixed distance $d$, then the imparted momentum is $P \\ell_{\\rm f}^2 t = m_{\\rm f} v_{\\rm f}$, and $r_{\\rm f}\\propto f_{\\rm f}^{-1/2}$. In both instances, fragment brightness should be weakly anti-correlated with non--projected radius $r$. In a two dimensional projection onto the sky radius $R$, the above inverse relation will be somewhat washed out. Nevertheless, we still expect to find brighter fragments at smaller radii. Figure \\ref{fig:lump-r-vs-flux} shows the observed relationship of sky radius, $R$, to flux. There is a statistically significant (Spearman rank-order $p_{\\rm SRO}=0.017$) positive correlation between $R$ and flux, in contrast to the expected anti-correlation. It is reassuring that a strong deficit of faint fragments at small $R$ was not observed, because this would be suggestive of a selection bias against finding fragments in the bright central coma near the nucleus. The fact that the shape of the flux distribution varies with $R$ is consistent with the space-filling distribution suggested by \\S\\ref{subsec:phasespace}, because a thin expanding shell of fragments would produce a distribution of fluxes that is invariant in $R$. In conclusion, the observation that flux increases rather than decreases with radius argues against a model in which the fragments consist of monolithic fragments with a sublimation rate and reflective flux proportional to their surface area. Otherwise, by simple gas pressure arguments, one would expect the largest and heaviest fragments to be closest to the nucleus. Instead, each detected fragment might in fact be a collection of active objects, rather than a single cohesive fragment. %In conclusion, the observation that flux increases rather than %decreases with radius argues against a model in which the fragments %consist of single fragments with a sublimation rate and reflective flux %proportional to their surface area. We use this observation to suggest that each detected fragment may in fact be a collection of active objects, rather than a single cohesive fragment. \\subsection{Fragment Fading} Figure~\\ref{fig:fading} shows the temporal fading of the nucleus and the median of the fragments' magnitudes during our observations. The nucleus fades at a rate of 0.15 mag day$^{-1}$ while the fragments, on average, fade at a similar rate of 0.19 mag day$^{-1}$. Figure \\ref{fig:lump-fading} rescales the flux of each fragment to a common baseline using an exponential fit with a shared time constant. Over the nine days plotted, the fragments fade by about 80\\%. % have removed discussion of fading rate vs brightness entirely. complex plot. %It is of interest to examine whether the fading rate $\\dot %F/F=d\\,\\log(F)/dt$ varies with the brightness of the fragments. %Figure \\ref{fig:lump-flux-vs-dlogfluxdt} shows the distribution of %this logarithmic fading rate as a function of fragment brightness. It %is evident, despite the very large photometric uncertainties at the %faint limit, that the fading rate is independent of brightness, as %shown by the straight lines denoting the median over two ranges of %flux. The assumption of monolithic pieces sublimating with constant %$\\dot\\ell$ implies that the flux of reflected light falls as $F(t) %\\propto (\\ell_0 - \\dot\\ell \\times t)^2$ and $\\dot F/F = -2 F^{-1/2}$. %This is inconsistent with the observed invariance of the %fading rate with flux. %To better constrain the temporal variation of the fading, we attempt a %robust fit of three fading models: linear, quadratic, and exponential. %We fit the fragment fading by minimizing $\\Delta$ in the following %expression, where $F_{i,j}$, $t_{i,j}$ is the flux and observation %time of fragment $i$ at observation date $j$, $b_i$ is an offset constant for %each fragment, and $c$ is a global fading parameter: %\\begin{equation} %\\label{eq:genlinfit} %\\Delta = %\\sum_{i}^{N_{\\rm Fragments}} \\sum_j^{N_{\\rm obs}} % \\left| % F_{i,j} - g(b_i + c t_{i,j}) % \\right| %\\end{equation} %In this expression, $g(x)$ is either an identity function %$g(x)=\\max(x,0)$ for a linear falloff; a square $g(x)=[\\max(x,0)]^2$ %for the monolithic sublimation model discussed above; or $g(x)=e^x$ %for the constant fractional falloff suggested by Figure %\\ref{fig:lump-flux-vs-dlogfluxdt}. The absolute value was used as a %robust fit metric because the more customary quadratic $\\chi^2$ is %excessively sensitive to the abundant outliers %\\citep[see][]{NumericalRecipesC}. The absence of a per--datum error %(like the individual $\\sigma$ in a conventional $\\chi^2$) corresponds %to our belief that the dominant noise is a constant, long-tailed, %additive error term arising from random fluctuations in the coma. We %perform two passes of optimization, the first on the entire data set, %and the second after rejecting points that are over five times more %deviant than the median per--datum deviation after the first fit. %This outlier rejection removes about 10\\% of points, but has little %effect on the results or quoted uncertainties, supporting our belief %in the robustness of the fit. % bootstrap results: % EXPr)\\propto r^{-D}$ with %% $D=2.9$ to 3.9, or differential of $N(r)=r^{-\\alpha}$ with %% $\\alpha=1.9$ to 2.9. Primary fragmentation leads to $D=2.5$ and more %% fragmentation increases $D$. So a power law distribution of particles %% is natural for eruptions. This is the same as Dohnanyi's asteroid mass %% result of $N(>m)\\propto m^{-5/6}$ giving $N(>r) \\propto r^{-2.51}$. %% %%\\end{itemize} %Jupiter family comet nuclei have been observed to fragment in several cases (e.g. 73P, see \\cite{Fernandez2009} for a review of recently observed splitting events). Typically, fragments are identified as mini-comets moving with small relative velocities. These objects are often difficult to identify given that they are usually a small fraction of the nucleus size, and their presence may be masked by the nucleus' coma. We have circumvented these problems by using a Laplacian filter to identify small-scale features in our images. This technique is sensitive to faint objects that would otherwise remain hidden by the dust coma ejected from the nucleus. %Much of this work supports the hypothesis that an unstable fragment broke off from the nucleus \\citep{2008ICQ....30....3S, 2009ICQ....31....5S}. This fragment probably consisted of a mixture of dirty ice and volatile-depeted regolith. As the fragment moved into sunlight, the volatiles that had previously been stored beneath the surface of the nucleus were exposed, heated and began to sublimate. This led to the release of large amounts of dust and, as cohesive chunks of material were released from the original secondary nucleus, the fragments that we discuss in this paper. The fragments appear to have an approximately isotropic distribution in space, suggesting their progenitor was unstable across much of its surface. %We find that the magnitudes of the observed fragments are inconsistent with inert monolithic %bodies. It is instead probable that they are active cometesimals with their own dust comae. This suggests %that the outburst carried a large amount of volatile material away from the nucleus, as well as dust. %Sekanina (2008) estimated that 10$^{11}$ kg of material was ejected from comet 17P/Holmes during the %outburst. We \u00dend that the mass of the fragments is $\\leq$10\\% of the total ejected mass, or 0.1\\% %of the nucleus mass. Assuming approximately 1\\% of the nucleus mass was ejected in total, and %that comet 17P/Holmes outbursts in a similar fashion every $\\sim$100 years, we calculate a lifetime of 10,000 years. This is comparable with the physical lifetime estimated for JFCs \\citep{2004come.book..659J}, %and much shorter than the median dynamical lifetime of JFCs of 450,000 years. %We suggest that comet 17P/Holmes will simply become unobservable towards the end of its life, rather %than undergoing tidal disruption by Jupiter or the Sun. We have identified and characterized fragments that were ejected from the nucleus of 17P/Holmes during its spectacular outburst in October 2007. Our findings are as follows: % don't underestimate importance of simple existence of fragments. \\begin{enumerate} \\item Sixteen fragments are detected in Laplacian-filtered images where the coma has been suppressed using an azimuthal average. \\item The motion of the fragments implies either isotropic or conical ejection from the nucleus on UT 2007 Oct.\\ 24.3 $\\pm$ 1.2. \\item Results from aperture photometry are inconsistent with inert, monolithic bodies. Modeling the fragments as sublimating cometesimals yields radii of 10 m to 110 m. Assuming a density of 400 kg m$^{-3}$, the fragments account for 10$^{10}$ kg of the total mass ejected, or $\\sim$ 0.1 \\% of the nucleus mass. \\item The fragments move unusually fast, with on-sky velocities of up to 125 m s$^{-1}$. Acceleration by CO (or other supervolatile) gas drag forces might be able to generate such large velocities given appropriate launch conditions at the nucleus. \\item We detect no systematic acceleration of the fragments and deduce that the bodies are not self-propelled by sublimation in a preferred direction. \\item The fragments fade at a rate of $\\sim$ 0.19 mag day$^{-1}$, consistent with the idea that they are active bodies, eventually becoming inert as surface volatiles are depleted. % lifetimes too (at least, when do they disappear), but make sure this is mentioned in discussion \\end{enumerate} The authors thank CFHT Director Christian Veillet for allocating time to this target of opportunity program and Pierre Martin, Jean-Charles Cuillandre, and the QSO team at CFHT for providing observational assistance. Pedro Lacerda and Bin Yang provided helpful comments. We appreciate support from a NASA Outer Planets Research grant to DJ. %have presented detection of 16 fragments moving radially outward from the Holmes parent nucleus one week after initial outburst %isotropic ejection %size/magnitudes/really detected! %motion consistent with having originated from the nucleus at time of outburst %observed fading \\newpage" }, "1003/1003.2323_arXiv.txt": { "abstract": "We have combined the thermo-chemical disc code \\ProDiMo\\ with the Monte Carlo radiative transfer code \\MCFOST\\ to calculate a grid of $\\sim$300\\,000 circumstellar disc models, systematically varying 11 stellar, disc and dust parameters including the total disc mass, several disc shape parameters and the dust-to-gas ratio. For each model, dust continuum and line radiative transfer calculations are carried out for {29} far IR, sub-mm and mm lines of [OI], [CII], $^{12}$CO and o/p-H$_2$O under 5 inclinations. {The grid allows to study the influence of the input parameters on the observables, to make statistical predictions for different types of circumstellar discs, and to find systematic trends and correlations between the parameters, the continuum fluxes, and the line fluxes. The model grid, comprising the calculated disc temperature and chemical structures, the computed SEDs, line fluxes and profiles, will be used in particular for the data interpretation of the {\\sc Herschel} open time key programme GASPS. The calculated line fluxes show a strong dependence on the assumed UV excess of the central star, and on the disc flaring. The fraction of models predicting [OI] and [CII] fine-structure lines fluxes above {\\sc Herschel/Pacs} and {\\sc Spica/Safari} detection limits are calculated as function of disc mass. The possibility of deriving the disc gas mass from line observations is discussed.} ", "introduction": "The {structure, composition and evolution of protoplanetary discs are important corner-stones to unravel the mystery of life, as they set the initial conditions for planet formation}. Spectral energy distributions (SEDs), {although their analysis is known to be degenerate}, probe the amount, temperature and overall geometry of the dust in the discs, such as disc flaring \\citep{Meeus2001}, puffed-up inner rims \\citep{Dullemond2001, Acke2009}, and indications of an average grain growth in discs as young as a few Myr \\citep{D'Alessio2001}. Most works in the past decade have focused on the analysis of SEDs of individual objects \\citep[e.g.][]{D'Alessio2006} or to study systematic trends in infrared colours of various types of discs, ranging from embedded young stellar objects to exposed T\\,Tauri stars \\citep{Robitaille2006}. Some ambiguities inherent in SED analysis can be resolved by images in scattered light \\citep{Stapelfeldt1998}, in mid-infrared thermal emission \\citep{McCabe2003} or in the mm-regime \\citep{Andrews2007}. The Spitzer observatory has enabled detailed studies on dust mineralogy, constraining dust properties in the upper layers of the inner disk regions by using solid-state features \\citep{Furlan2006, Olofsson2009}. % Multi-technique, panchromatic approaches, combining the aforementioned observations, are now becoming possible but remain limited to a few objects with complete data sets \\citep[e.g.][]{Wolf2003, Pinte2008, Duchene2009}. \\begin{figure*} \\centering {\\ }\\\\*[-2ex] \\hspace*{-1mm}\\includegraphics[width=16cm]{ModelPipeline2.eps}\\\\[-1mm] \\caption{Unified {\\sc Mcfost / ProDiMo} model pipeline to calculate one radiation thermo-chemical disc model with SED and line flux prediction.} \\label{fig:ModelPipeline} \\vspace*{-1mm} \\end{figure*} However, all these observational findings are related to the dust component. {Initially}, about 99\\% of the disc mass can be assumed to be present in form of gas, and the progress toward a better understanding of the gas component, such as chemical composition, gas temperature structure and vertical disc extension, is hampered by a {current} lack of observational data. With several {key programs of the {\\sc Herschel Space Observatory}, such as GASPS and WISH, a new body of gas emission line observations will be provided very soon, and there is a clear need to include the gas in such systematic studies.} Recent work has focused on the prediction of far IR line emissions from individual discs \\citep{Meijerink2008,Ercolano2008,Woitke2009b,Cernicharo2009}, or rather small parameter studies \\citep{Kamp2009}. \\citet{Goicoechea2009} have recently discussed the detection rates of the $\\rm[OI]\\,63\\,\\mu$m, $\\rm[SI]\\,56\\,\\mu$m and $\\rm [SiII]\\,34\\,\\mu$m fine-structure lines with {\\sc Herschel} and the proposed {\\sc Spica/Safari} mission on the basis of {\\sl one} low mass disc model ($M_{\\rm gas}\\!=\\!10^{-5}\\rm\\,M_\\odot$) by \\citet{Gorti2004}. {Their conclusions, however, depend on the choices of the other model parameters, and it is difficult to put them on a firm statistical basis.} To address these issues, {we have combined our state-of-the-art computer codes to calculate the dust and line radiative transfer with the \\MCFOST-code \\citep{Pinte2006}, and the gas thermal balance and chemistry with the \\ProDiMo-code \\citep{Woitke2009a,Kamp2009}}. We have computed a large grid of 300\\,000 disc models to simultaneously predict SEDs and gas emission lines from parametrised disc density distributions. The grid name DENT stands for \\underline{D}isc \\underline{E}volution with \\underline{N}eat \\underline{T}heory. The following sections describe the model pipeline (Sect.~\\ref{DENTgrid}) and the first results on fine structure emission line fluxes and detectability (Sect.~\\ref{results}). Section \\ref{summary} summarises the preliminary findings of the DENT grid and provides an outlook to future studies. ", "conclusions": "\\label{summary} In a concerted effort of the theory groups in Edinburgh, Grenoble and Groningen, we have computed a grid of 300\\,000 circumstellar disc models, {simultaneously solving} gas-phase, UV-photo and ice chemistry, detailed heating \\& cooling balance, and continuum \\& line radiative transfer. The first results of the DENT grid show a strong dependence of the calculated emission line fluxes on the assumed stellar UV excess and on the flaring of the disc. The stellar UV is essential for the heating of the upper disc layers. In combination with positive disc flaring, a strong stellar UV irradiation creates an extended warm surface layer with $\\Tg\\!>\\!\\Td$ responsible for the line emissions. However, if the disc is not flared (self-shadowed), discs with total mass $\\ga\\!10^{-5}-10^{-4}\\rm\\,M_{\\odot}$ increasingly shield the stellar UV by their inner parts, which causes much cooler surface layers, and a saturation of the line fluxes with increasing disc mass. Despite these complicated parameter dependencies, we have shown that the $\\rm[OI]\\,63.2\\mu$m line flux depends basically on two quantities, namely the total disc gas mass and the mean disc temperature. We will continue this work by two follow-up papers (Kamp et al.~2010, in prep., M{\\'e}nard et al.~2010, in prep.) that will provide more insight into the statistical behaviour of gas line and dust continuum predictions, respectively, to identify trends and robust correlations with disc mass. \\smallskip\\noindent In summary, the DENT grid allows to\\\\[-0.5ex] \\hspace*{-11mm}\\begin{tabular}{cp{80mm}} $\\bullet$\\hspace*{-2.5mm} & study the effects of stellar, disc, and dust parameters on continuum and line observations,\\\\ $\\bullet$\\hspace*{-2.5mm} & allow for a qualified interpretation of observational data,\\\\ $\\bullet$\\hspace*{-2.5mm} & quickly predict line and continuum fluxes for planning observations,\\\\ $\\bullet$\\hspace*{-2.5mm} & search for best-fitting models concerning a given set of observed line and continuum fluxes,\\\\ $\\bullet$\\hspace*{-2.5mm} & study the robustness of certain fit values against variation of the observational data.\\\\[1.5ex] \\end{tabular} \\noindent We intend to make the calculated DENT grid available to the scientific community. A graphical user interface called {\\tt xDENT} has been developed to allow researchers to visualise the DENT results, to make plots as presented in this letter, and to search for best-fitting models for a given set of continuum and line flux data. We emphasise, however, that the DENT grid has not been developed for detailed fitting of individual objects. The coarse sampling of the 12-dimensional parameter space can mostly be used to narrow down the parameter range for individual objects, for example to design a finer sub-grid, especially for not so well-known objects. With a comprehensive data set of far IR gas emission lines to be obtained by Herschel/GASPS very soon, we aim at breaking the degeneracy of SED fitting and make possible a more profound analysis of the physical, chemical and temperature structures of discs around young stars. \\appendix" }, "1003/1003.0110_arXiv.txt": { "abstract": "We evaluate two dominant nuclear reaction rates and their uncertainties that affect $^{44}$Ti production in explosive nucleosynthesis. Experimentally we develop thick-target yields for the $^{40}$Ca($\\alpha$,$\\gamma$)$^{44}$Ti reaction at $E_{\\rm \\alpha} = 4.13, 4.54$, and $5.36$ MeV using $\\gamma$-ray spectroscopy. At the highest beam energy, we also performed an activation measurement which agrees with the thick target result. From the measured yields a stellar reaction rate was developed that is smaller than current statistical-model calculations and recent experimental results, which would suggest lower $^{44}$Ti production in scenarios for the $\\alpha-$rich freeze out. Special attention has been paid to assessing realistic uncertainties of stellar reaction rates produced from a combination of experimental and theoretical cross sections. With such methods, we also develop a re-evaluation of the $^{44}$Ti($\\alpha$,$p$)$^{47}$V reaction rate. Using these two rates we carry out a sensitivity survey of $^{44}$Ti synthesis in eight expansions representing peak temperature and density conditions drawn from a suite of recent supernova explosion models. Our results suggest that the current uncertainty in these two reaction rates could lead to as large an uncertainty in $^{44}$Ti synthesis as that produced by different treatments of stellar physics. ", "introduction": "The dynamic synergy between observation, theory, and experiment developed over many years around the field of $\\gamma$-ray astronomy has as its ultimate goal observations of specifc radionuclides informing our understanding of stellar explosions and the theoretical models that predict nucleosynthesis. Of the radioactive species observed so far, the most long lived, $^{26}$Al and $^{60}$Fe (with half-lives $\\tau_{1/2}=7.17\\pm 0.24 \\times 10^5$ and $2.62\\pm 0.04 \\times 10^6$ yr, respectively), are in reasonably good agreement with theoretical predictions \\citep{timmes96,diehl06}. Observations of those in the iron group, $^{56,57}$Ni ($\\tau_{1/2}=6.075 \\pm 0.01$ d and $35.6 \\pm 0.06$ hr, respectively) and their decay products $^{56,57}$Co ($\\tau_{1/2}=77.233\\pm 0.027$ d and $271.74\\pm 0.06$d, respectively), are used in many ways to constrain our current models of the core collapse explosion mechanism. The radionuclide $^{44}$Ti ($\\tau_{1/2}=58.9\\pm 0.3$ yr), made in the same explosive environment but in much lower amounts compared to the very abundant nickle isotopes \\citep{wac73,woosley95}, is hoped to one day serve as an even more sensitive diagnostic and a valuable probe of the conditions extant in some of the deepest layers to be ejected. Observationally most of the attention has focused on the detection of the 68, 78, and 1157 keV $\\gamma$-rays from the $^{44}$Ti $\\rightarrow$ $^{44}$Sc $\\rightarrow$ $^{44}$Ca decay chain (see Figure 1). The 1157 keV $\\gamma$-ray has been observed directly from a point source in Cassiopeia A (Cas A; Iyudin et~al. 1994). This was later confirmed by observation of the low-energy $^{44}$Sc $\\gamma$-rays using the $\\textit{BeppoSax}$ \\citep{vink01} and $\\textit{INTEGRAL}$ \\citep{renaud06} observatories. Using values for the distance, age, and $\\gamma$-flux of Cas A, the amount of $^{44}$Ti ejected was found to be 1.6$^{+0.6}_{-0.3}$ $\\times$10$^{-4}$ M$_\\odot$ \\citep{renaud06}, in agreement with earlier observations using CGRO \\citep{timmes96}. Although the presence of $^{44}$Ti is currently below detection limits in SN1987A in the nearby Large Magellanic Cloud, its light curve is theorized to now be powered by the decay of $^{44}$Ti. The yield of $^{44}$Ti in SN1987A has been estimated from its light curve to be $1 - 2 \\times 10^{-4} M_\\odot$, a factor of 3 greater than predicted by models \\citep{diehl06}. A third $^{44}$Ti source (of lower significance compared to the one in CasA) has been observed in the Vela region \\citep{iyudin98} but the existence of a co-located young supernova (SN) remnant has not been confirmed. While the mass of $^{44}$Ti observed in SN remnants appears to be underproduced by past models, the number of observed sources of $^{44}$Ti in all-sky surveys appears to be less than expected from estimates of the Galactic SN rate and the known $^{44}$Ti half life, leading some to question whether $^{44}$Ti-producing SNe are exceptional \\citep{the06}. Theoretically, $^{44}$Ti production is traditionally ascribed to regions experiencing a strong \"$\\alpha-$rich freeze out\", where material initially in nuclear statistical equilibrium (NSE) at relatively low density is cooled so rapidly that free $\\alpha-$particles do not have time to reassemble, through inefficient $3-$body reactions that span the mass gaps at A=5 and A=8, back into the iron group. A likely scenario is material in or near the silicon shell as it experiences shock wave passage during a core collapse SN event. Assuming the material cools adiabatically over a hydrodynamic timescale ($\\tau_{\\rm HD} = 446\\chi/\\sqrt{\\rho_i}$), where $\\chi$ is a scaling parameter (here unity) and $\\rho_i$ is the initial (peak) density in g cm$^{-3}$, initial conditions that would result in a final $\\alpha$-particle mass fraction of $\\sim$ 1\\% would require temperatures high enough to ensure NSE ($T_{9i}\\gtrapprox 5$) and an initial density $\\rho_i < {\\rm min}(4.5\\times 10^5 T_{9i}^3, \\quad 2.5\\times 10^5 T_{9i}^4 \\chi^{-2/3})$ \\citep{wac73}. This translates to a radiation entropy greater than unity (Section 3.2). In models of massive stars the $\\alpha-$rich freeze out dominates the solar production of several species, including $^{44}$Ca (made as Ti), $^{45}$Sc, $^{57}$Fe (made as Ni), and $^{58,60}$Ni, while still others seem to require a sizable component to account for their solar abundances, including $^{50,52}$Cr, $^{59}$Co, $^{62}$Ni, and $^{64}$Zn \\citep{woosley95,tnh96}. Although $^{56}$Ni is the dominant species produced in both NSE and the $\\alpha-$rich freez out over a wide range ($0 \\leq \\eta \\leq 0.01$) of neutron excess \\citep{hwei85}, the solar abundance of $^{56}$Fe is dominated by production in explosive silicon burning in massive stars and by a large contribution from SNe Ia \\citep{twoosley95}. In order to address one aspect of the model uncertainties associated with the theoretical predictions of $^{44}$Ti in SNe, we focus on exploring the nuclear data uncertainties of two key reaction rates. Experimentally we address the $^{40}$Ca($\\alpha$,$\\gamma$)$^{44}$Ti cross section where the existing experimental results are inconsistent and theoretical estimates are complicated by the suppression of $E1$ $T = 0 \\rightarrow T = 0$ gamma transitions in self-conjugate ($N$ = $Z$) nuclei \\citep{rauscher00a}. This cross section has been measured in the past using several techniques. \\cite{cooperman} used in-beam $\\gamma$-ray spectroscopy to measure the capture of $\\alpha$-particles on a metallic calcium target in the center-of-mass energy range $E_{\\rm CM} = 2.5 - 3.65$ MeV. In that work the excitation function for the 1083 keV first excited state transition in $^{44}$Ti was determined and the resonance strengths were developed. \\cite{Nassar} determined the integral cross section in the range $E_{\\rm CM} = 2.2 - 4.17$ MeV by bombarding a He gas target with a $^{40}$Ca beam and collecting the recoiling $^{44}$Ti in a catcher foil. Accelerator mass spectroscopy was then used to determine the ratio of $^{44}$Ti/Ti from the known content of Ti in the catcher foil. Recently, a slightly broader energy range of $E_{\\rm CM} = 2.11 - 4.19$ MeV was explored using the DRAGON recoil mass spectrometer \\citep{vockenhuber07} in inverse kinematics. They developed individual resonances whose sum was used to determine a reaction rate for $1\\le T_9 \\le 5.5$. Yet despite these often heroic expenditures of time, toil, and treasure, in the temperature range of astrophysical interest there still exists a factor of 3 or more difference between the experimentally determined reaction rates. In this work, we develop the cross section for $^{40}$Ca($\\alpha$,$\\gamma$)$^{44}$Ti by two separate methods as a check on systematic uncertainties. First we used in-beam $\\gamma$-ray spectroscopy to measure a thick target yield of the $^{40}$Ca($\\alpha$,$\\gamma$)$^{44}$Ti reaction. We then determined the number of $^{44}$Ti nuclei produced by counting low-energy $\\gamma$-rays from the decay of $^{44}$Ti in an irradiated target. Special attention was devoted to checking the internal consistency of the measurements and to establishing realistic uncertainties in developing the stellar reaction rate from a combination of experimental and theoretical cross sections. We have made a similar evaluation of the stellar reaction rate for the dominant destruction reaction, $^{44}$Ti($\\alpha$,$p$)$^{47}$V, based on the original experimental work of \\cite{sonz00} and the theoretical cross section work of \\cite{rauscher01}. Our results are presented in four parts. In Section 2, we describe our experimental efforts. Section 3 then discusses the development of stellar reaction rates, and an expose of the past and present experimental and theory efforts for the two rates in question. In Section 4, we present nucleosynthesis results for the production of $^{44}$Ti and $^{56,57,58}$Ni reported in previous surveys of massive star evolution. We then carry out a sensitivity survey of $^{44}$Ti to variations in the principle production and destruction rates using more recent SN models. In Section 5 we provide a discussion and conclusions. ", "conclusions": "We have considered the sensitivity to $^{44}$Ti production in expansions that approximate freeze outs from NSE due to variations in the principle production and destruction reactions $^{40}$Ca($\\alpha,\\gamma$)$^{44}$Ti and $^{44}$Ti($\\alpha,p$)$^{47}$V, and contrast them to experimental and theory reaction rate developments over the past 20 years. Experimentally we have also measured a thick-target yield for the $^{40}$Ca($\\alpha$,$\\gamma$)$^{44}$Ti reaction. In-beam $\\gamma$-ray spectroscopy was used to determine the yield of the 1083 keV prompt $\\gamma$-ray from $^{44}$Ti at $E_\\alpha$ = 4.13, 4.54, and 5.36 MeV. In order to correct for those transitions which bypass the 1083 keV transition, the Monte Carlo code DICEBOX was used to estimate a correction of 20$\\%$ to the in-beam thick target yield. An off-line activation measurement using the target from the $E_\\alpha$ = 5.36 MeV irradiation showed good agreement with the in-beam measurement. We then derived a thermonuclear reaction rate by normalizing the NON-SMOKER cross section (down by a factor of 1.71) to agree with our measured off-line thick target yield. We derived an error bar for our recommended rate whose magnitude ($\\pm \\times 2$) was dominated by the {\\sl theoretical} cross section error. We also carry out a similar re-evaluation of the $^{44}$Ti($\\alpha,p$)$^{47}$V reaction rate whose cross section was measured by \\cite{sonz00}. For both reactions, we conclude that the experimental data were far above the Gamow window, and suggest that further measurements be attempted. We then carried out a sensitivity survey of $^{44}$Ti nucleosynthesis in adiabatic expansions from eight peak temperature and density combinations drawn from conditions in three recent stellar explosion models \\citep{mag08}. For each expansion we survey a range of initial compositions ($0.505 \\geq Y_e \\geq 0.485$). We also vary the principle production and destruction rates affecting $^{44}$Ti in these expansions (eight choices of production rate, four for destruction). Our results show $^{44}$Ti produced in proportion to solar iron for only one expansion drawn from a model for Cassiopia A, even though the final mass fractions for $^{44}$Ti in most of the expansions are typically of order $10^{-4}$. With one exception, the other expansions are consistent with $^{44}$Ti production seen in previous surveys of one-dimensional stellar evolution and with constraints imposed by solar abundances. Our results suggest that a strong $\\alpha$-rich freeze out ($X(\\alpha)_f\\sim 0.2-0.3$) is highly conducive to $^{44}$Ti synthesis. With respect to reaction rate sensitivity, our experimental results suggest a recommended $^{40}$Ca($\\alpha,\\gamma$)$^{44}$Ti reaction rate that is smaller than those predicted by the most recent experimental efforts \\citep{vockenhuber07,Nassar}, but with a fairly large error bar ($\\pm \\times 2$) that would in fact encompass the former one. Nucleosynthesis models using our recommended rate would suggest less $^{44}$Ti than these two recent experiments, although it would be higher than that suggested by an earlier measured rate \\citep{cooperman} and the current theory rate \\citep{rauscher00b}. The total range of sensitivity is a factor of 1.5 (36\\%) when considering all of the production rates available. We also find that the uncertainties associated with the dominant destruction rate, $^{44}$Ti($\\alpha,p$)$^{47}$V, have roughly double the impact on $^{44}$Ti synthesis ($\\times 3.2$, or 70\\%) than that exhibited by the entire range of production rates. We suggest the use of our re-evaluation of the only available experimental reaction rate \\citep{sonz00} and strongly suggest consideration of its attendant larger uncertainty ($\\pm \\times 3$). This is a slightly larger spread in $^{44}$Ti sensitivity than observed in two recent surveys of massive star nucleosynthesis \\citep{rauscher02,cl06} that used essentially the same nuclear data, suggesting that current uncertainties in reaction rates could lead to as large an uncertainty in $^{44}$Ti synthesis as that produced by different treatments of stellar physics. Since the seminal work of \\cite{wac73}, several new and novel theories of SN nucleosynthesis have featured regimes within the neutrino wind where the $\\alpha-$rich freeze out plays an important role, including scenarios for the $r$-process \\citep{woo94,hwq97}, and the $\\nu p$-process \\citep{frolich06,pruet06}. In each a combination of high entropy $(S_{\\rm rad}\\geq 50)$ and short expansion time-scale are required to produce the unique signatures of each process (the $r$-abundances, and light $p$-nuclei, respectively). However, as a consequence of extreme neutrino irradiation, the composition is forced away from neutron-proton equality and ultimately high entropy material enters regions above the iron group where local effects due to rapid changes in particle separation energies have a strong influence on the net nuclear flows. In these works no $^{44}$Ti is reported. Further out in the SN ejecta theory does not show $^{44}$Ti enhancement due to the $\\nu$-process in massive stars \\citep{whhh90,woosley95}, nor in recent low-mass (electron capture) core-collapse scenarios \\citep{hmj08,wan09}. The later has been shown to effectively synthesize the long elusive $\\alpha$-rich freeze-out candidate $^{64}$Zn, but not $^{44}$Ti. Our limited survey suggests that $^{44}$Ti synthesis requires modest entropy $(S_{\\rm rad}\\sim 35$) and expansion timescales ($\\sim \\tau_{\\rm HD} = 0.45$ s) over a fairly narrow range of $Y_e$ ($0.5 \\geq Y_e \\geq 0.4980$). To date surveys of massive star evolution and nucleosynthesis have typically underproduced species whose production is attributed to the $\\alpha$-rich freeze-out, in particular $^{44}$Ca (made as $^{44}$Ti) and $^{64}$Zn \\citep{woosley95,cl06}. This has largely been ascribed to variations in treatments of stellar physics, most notably the parameterization of the explosion in one-dimensional models. Incorporating stellar yields from these surveys into models for galactic chemical evolution indicate the degree of underproduction, roughly a factor of a $2-3$ \\citep{twoosley95}. Recent work has lead to models for the progenitor of Cassiopia A that suggest increased production, but the underlying physics is still quite uncertain and in need of improvement \\citep{young06}. Future attention may focus on issues related to {\\sl asymmetrical} explosions where a noted increase in $^{44}$Ti production compared to models with imposed spherical symmetry has been suggested for many years \\citep{nag97,nag98,hwang03,young06}. The community is on the threshold of three-dimensional calculations that should provide valuable insight into the core-collapse mechanism, including physics such as rotation and magnetic fields which will likely enforce an asymmetrical result. Such results should help guide future parameterizations of the explosion in our one-dimensional models which will continue to carry the burden in future surveys of massive star evolution and nucleosynthesis. We believe we have addressed the uncertainty in two key nuclear reaction rates affecting $^{44}$Ti synthesis, but ultimately, if Type II SNe are the dominant site of $^{44}$Ti production, future models will have to include a larger fraction of their ejecta that experience an $\\alpha$-rich freeze out than they have in the past. Another solution would be an additional source of $^{44}$Ti, such as rare SNe of type Ia \\citep{wtw86,woo97} or Ib \\citep{perets10}." }, "1003/1003.4113_arXiv.txt": { "abstract": "{ {\\it Aims.} The ESO SN Ia Progenitor Survey (SPY) aims at finding merging double degenerate binaries as candidates for supernova type Ia (SN Ia) explosions. A white dwarf merger has also been suggested to explain the formation of rare types of stars like R CrB, extreme helium or He sdO stars. Here we present the hot subdwarf B binary GD\\,687, which will merge in less than a Hubble time.\\\\ {\\it Methods.} The orbital parameters of the close binary have been determined from time resolved spectroscopy. Since GD\\,687 is a single-lined binary, the spectra contain only information about the subdwarf primary and its orbit. From high resolution spectra the projected rotational velocity was derived. Assuming orbital synchronisation, the inclination of the system and the mass of the unseen companion were constrained.\\\\ {\\it Results.} The derived inclination is $i=39.3^{+6.2}_{-5.6}\\,^{\\circ}$. The mass $M_{\\rm 2}=0.71_{-0.21}^{+0.22}\\,M_{\\rm \\odot}$ indicates that the companion must be a white dwarf, most likely of C/O composition. This is only the fourth case that an sdB companion has been proven to be a white dwarf unambiguously. Its mass is somewhat larger than the average white dwarf mass, but may be as high as $0.93\\,M_{\\rm \\odot}$ in which case the total mass of the system comes close to the Chandrasekhar limit. \\\\ {\\it Conclusions.} GD\\,687 will evolve into a double degenerate system and merge to form a rare supermassive white dwarf with a mass in excess of solar. A death in a sub-Chandrasekhar supernova is also conceivable. ", "introduction": "} Double degenerate (DD) binaries consisting of two white dwarf (WD) stars experience a shrinkage of their orbits caused by gravitational wave radiation. Sufficiently close binaries with orbital periods of less than half a day will eventually merge within less than a Hubble time. The outcome of such a merger can be a single compact object like a white dwarf or a neutron star. But the merger event may also lead to the ignition of nuclear burning and trigger the formation of rare objects such as extreme helium stars (Saio \\& Jeffery \\cite{saio}), R CrB stars (Webbink \\cite{webbink}) or He sdO stars (Heber \\cite{heber3}; Napiwotzki \\cite{napiwotzki10}). The merger of two sufficiently massive C/O white dwarfs may lead to a Supernova of type Ia (SN~Ia, Webbink, \\cite{webbink}). SN~Ia play a key role in the study of cosmic evolution. They are utilised as standard candles for determining the cosmological parameters (e.g. Riess et al. \\cite{riess}; Leibundgut \\cite{leibundgut}; Perlmutter et al. \\cite{perlmutter}). There is general consensus that the thermonuclear explosion of a white dwarf of Chandrasekhar mass causes a SN~Ia. The DD merger (Iben \\& Tutukov \\cite{iben}) is one of two main scenarios to feed a white dwarf to the Chandrasekhar mass. Progenitor candidates for the DD scenario have to merge in less than a Hubble time and the total binary mass must exceed the Chandrasekhar limit. In general, DD candidates are assumed to consist of two white dwarfs. The discovery of the sdB+WD binary KPD\\,1930$+$2752, however, has highlighted the importance of such systems as SN~Ia progenitor candidates (Bill\\`{e}res et al. \\cite{billeres}; Maxted et al. \\cite{maxted1}; Geier et al. \\cite{geier1}). Hot subdwarf stars (sdBs) are core helium-burning stars situated at the hot end of the horizontal branch with masses around $0.5\\,M_{\\rm \\odot}$ (Heber \\cite{heber1}). Although the formation of these objects is still under debate, it is general consensus that they must loose most of their envelope at the tip of the RGB and evolve directly from the EHB to the WD cooling tracks without ascending the AGB (see Heber \\cite{heber} for a review). Since a high fraction of sdBs resides in close binaries with unseen companions (Maxted et. al \\cite{maxted2}; Napiwotzki et al. \\cite{napiwotzki6}), such systems can qualify as SN~Ia progenitors, if the companion is a white dwarf and the lifetime of the sdB is shorter than the merging time of the binary. This is the case for the sdB+WD binary KPD\\,1930+2752, which is the best known DD progenitor candidate for SN~Ia (Maxted et al. \\cite{maxted1}; Geier et al. \\cite{geier1}). Systematic radial velocity (RV) searches for DDs have been undertaken (e.g. Napiwotzki \\cite{napiwotzki2} and references therein). The largest of these projects was the ESO SN~Ia Progenitor Survey (SPY). More than $1000$ WDs were checked for RV-variations (Napiwotzki et al. \\cite{napiwotzki9,napiwotzki2}). SPY detected $\\sim 100$ new DDs (only $18$ were known before). One of them may fulfil the criteria for SN~Ia progenitor candidates (Napiwotzki et al. \\cite{napiwotzki3}). This sample includes about two dozen radial velocity variable sdB stars with invisible companions. As the sdB stars are intrinsically brighter than the white dwarfs, the nature of the companion is not so easily revealed. Most sdBs in close binary systems are single-lined and no features of the companions are visible in their spectra. In this case the nature of the unseen companion is hard to constrain since only lower limits can be derived for the companion masses. Given the fact that sdBs are subluminous stars, main sequence companions earlier than K-type can easily be excluded. But if the derived lower limit for the companion mass is lower than about half a solar mass it is in general not possible to distinguish between a main sequence companion and a compact object like a white dwarf. Despite the fact that the catalogue of Ritter \\& Kolb (\\cite{ritter}) lists more than 80 sdB binaries, the nature of their companions can only be constrained in special cases (e.g. For et al. \\cite{for}; Geier et al. \\cite{geier2}), in particular from photometric variations due to eclipses, reflection effects or ellipsoidal deformations. Results for eight DD systems discovered in the SPY survey have been presented in papers I--IV (Napiwotzki et al. \\cite{napiwotzki5}; Napiwotzki et al. \\cite{napiwotzki7}; Karl et al. \\cite{karl2}; Nelemans et al. \\cite{nelemans}). Here we study the sdB binary GD\\,687, a DD progenitor system which will merge in less than a Hubble time. \\begin{figure}[t!] \\centering \\resizebox{\\hsize}{!}{\\includegraphics{power_gd687.eps}} \\caption{In this power spectrum $-\\log\\chi^{2}$ of the best sine fit is plotted against the orbital periods. The region around the best solution is shown in the inlet. Confidence limits ($1\\sigma$,$3\\sigma$,$6\\sigma$) are marked with horizontal lines (dotted, dashed, solid).} \\label{power} \\end{figure} ", "conclusions": "The derived companion mass was calculated under the assumption of orbital synchronisation. Since theoretical synchronisation timescales for hot stars with radiative envelopes are not consistent (Zahn \\cite{zahn}; Tassoul \\& Tassoul \\cite{tassoul}), empirical evidence for orbital sychronisation in sdB binaries is needed. Geier et al. (\\cite{geier2}) found such evidence by detecting a variation in the lightcurve of the sdB+WD binary PG\\,0101$+$039, which could be identified as ellipsoidal deformation of the sdB. Since the orbital period of PG\\,0101$+$039 is $0.57\\,{\\rm d}$, sdB binaries with shorter periods like GD\\,687 are very likely synchronised as well. Recently van Grootel et al. (\\cite{vangrootel}) performed an asteroseismic analysis of the pulsating sdB binary Feige 48 and for the first time proved orbital sychronisation in this way. The orbital period of Feige 48 ($0.36\\,{\\rm d}$) is very similar to the one of GD\\,687. Furthermore, the atmospheric parameters of GD\\,687 ($T_{\\rm eff}=24\\,300\\,{\\rm K}$, $\\log{g}=5.32$) indicate that it has already evolved away from the ZAEHB (Zero Age Extreme Horizontal Branch) and should therefore, according to evolutionary calculations (e.g. Dorman et al. \\cite{dorman}), have a similar age as PG\\,0101$+$039 and Feige\\,48 (see Fig.~\\ref{tefflogg}). We thus conclude that the assumption of orbital sychronisation is fully justified in the case of GD\\,687. In about $100\\,{\\rm Myr}$ the helium-burning in the core of the sdB will come to an end. After a short period of helium-shell-burning this star will eventually become a white dwarf consisting of C and O. GD\\,687 is one of only a few known DD progenitor systems, where both components are C/O white dwarfs and which will merge in less than a Hubble time. Compared to the sdB+WD binary KPD\\,1930$+$2752 ($P\\approx0.1\\,{\\rm d}$, Maxted et al. \\cite{maxted1}; Geier et al. \\cite{geier1}), which is the best known candidate for DD SN~Ia progenitor, the orbital period of GD\\,687 is rather long. This leads to a merging time of $11.1\\,{\\rm Gyr}$, which is just a little shorter than the Hubble time, compared to only $200\\,{\\rm Myr}$ for KPD\\,1930$+$2752. With a total mass of $1.18_{-0.21}^{+0.22}\\,M_{\\rm \\odot}$ for the most likely subdwarf mass it may come close to the Chandrasekhar limit of $1.4\\,M_{\\rm \\odot}$ and is therefore placed at the edge of the progenitor parameter space (see Fig. \\ref{progen}). In contrast to KPD\\,1930$+$2752, where the primary mass could be constrained by an additional analysis of the subdwarfs ellipsoidal deformation visible as variation in its lightcurve, no such constraint can be put on the primary mass of GD\\,687 yet. Instead of exploding as SN~Ia, the merger of the two white dwarfs will most likely lead to the formation of a supermassive white dwarf with O/Ne/Mg-core. Up to now four binaries with total masses between about $1.2$ and $1.4\\,{\\rm M_{\\odot}}$ have been discovered. In two of them the visible component is an sdB. Since the sdB close binary fraction is much higher than the one of white dwarfs, it may be easier to find double-degenerate binary progenitors in the hot subdwarf population. Two other massive DD systems, a central star of a planetary nebula (Tovmassian et al. \\cite{tovmassian}; Napiwotzki et al. \\cite{napiwotzki8}) and a white dwarf from the SDSS survey (Badenes et al. \\cite{badenes}; Marsh et al. \\cite{marsh}; Kulkarni \\& van Kerkwijk \\cite{kulkarni}), were discovered serendipitously. Even though GD\\,687 does not qualify as SN~Ia progenitor candidate, the discovery of a population of double degenerate binaries (and progenitor systems) with total masses close to the Chandrasekhar limit in the course of the SPY survey (see Fig. \\ref{progen}) provides evidence for a similar population exceeding this limit. The same binary evolution channel that produces the sub-Chandrasekhar systems will also produce super-Chandrasekhar systems with slight changes in the initial conditions only. \\begin{figure}[t!] \\resizebox{\\hsize}{!}{\\includegraphics{Progen.ps}} \\caption{Total mass plotted against logarithmic period of double degenerate systems from the SPY survey. GD\\,687 is marked with the filled circle. For KPD\\,1930$+$2752 a mass range is given (Geier et al. \\cite{geier1}). The filled rectangles mark double-lined WDs, for which the absolute masses can be derived. The filled circle with arrow marks the lower mass limit derived for HE\\,1047$-$0436 (Napiwotzki et al. \\cite{napiwotzki5}), the filled diamond the lower mass limit derived for PN\\,G\\,135.9$+$55.9 (Napiwotzki et al. \\cite{napiwotzki8}). The open symbols mark single-lined WDs, sdBs, and sdOs. The companion masses of the single-lined systems are derived for the expected average inclination angle ($i=52^{\\circ}$) (Napiwotzki et al. \\cite{napiwotzki7}; Karl et al. \\cite{karl2}; Karl \\cite{karl}; Nelemans et al. \\cite{nelemans}; Napiwotzki et al. \\cite{napiwotzki3}). } \\label{progen} \\end{figure}" }, "1003/1003.3505_arXiv.txt": { "abstract": "A tight linear correlation is established between the HCN line luminosity and the radio continuum (RC) luminosity for a sample of 65 galaxies (from Gao \\& Solomon's HCN survey), including normal spiral galaxies and luminous and ultraluminous infrared galaxies (LIRGs/ULIRGs). After analyzing the various correlations among the global far-infrared (FIR), RC, CO, and HCN luminosities and their various ratios, we conclude that the FIR-RC and FIR-HCN correlations appear to be linear and are the tightest among all correlations. The combination of these two correlations could result in the tight RC-HCN correlation we observed. Meanwhile, the non-linear RC-CO correlation shows slightly larger scatter as compared with the RC-HCN correlation, and there is no correlation between ratios of either RC/HCN-CO/HCN or RC/FIR-CO/FIR. In comparison, a meaningful correlation is still observed between ratios of RC/CO-HCN/CO. Nevertheless, the correlation between RC/FIR and HCN/FIR also disappears, reflecting again the two tightest FIR-RC and FIR-HCN correlations as well as suggesting that FIR seems to be the bridge that connects HCN with RC. Interestingly, despite obvious HCN-RC and RC-CO correlations, multi-parameter fits hint that while both RC and HCN contribute significantly (with no contribution from CO) to FIR, yet RC is primarily determined from FIR with a very small contribution from CO and essentially no contribution from HCN. These analyses confirm independently the former conclusions that it is practical to use RC luminosity instead of FIR luminosity, at least globally, as an indicator of star formation rate in galaxies including LIRGs/ULIRGs, and HCN is a much better tracer of star-forming molecular gas and correlates with FIR much better than that of CO. ", "introduction": "Ever since far-infrared (FIR) emission was thought to mainly originate as a result of star formation in giant molecular clouds (GMCs; \\eg, Mooney \\& Solomon 1988), it had been used as a standard indicator of star formation rate (SFR). But until now, the spatial resolution of FIR observation has limited further detailed research, particularly in external galaxies. As another result of the formation of young massive stars, though at the end of their rapid formation and evolution, the radio continuum (RC) emission at centimeter (cm) wavelength has a completely different emission mechanism compared with that of the FIR emission. RC offers a hope, with its spatially resolved capability, to study SFR in galaxies both near (e.g., Condon et al. 1996) and far (e.g., Schinnerer et al. 2007) . The FIR emission arises from the dust heated by new-born massive stars. The young OB stars are imbedded in very tiny, yet massive and dense regions and all of their UV/optical radiation is absorbed by dust which re-radiates mostly in mid/FIR. Though for normal disk galaxies the situation is more complex in that the dust heating may also be contributed by older stellar population, and the optical depth of dust may not be thick enough for the FIR luminosity to truthfully measure the bolometric luminosity of entire galaxies except for the nuclear starburst regions. However, the FIR emission should provide an excellent measure of the SFR in luminous and ultraluminous infrared galaxies (LIRGs/ULIRGs) where dusty circumnuclear starbursts dominate (Kennicutt 1998a). The RC is mainly non-thermal synchrotron emission that arises from the interaction of relativistic electrons with the ambient magnetic field in which they diffuse. The supernova remnants (SNRs) of Type II and Type Ib supernovae that are produced by massive stars (M$\\gtrsim$8\\Msun), which have lifetimes $\\lesssim 3 \\times 10^{7}$~yr, are thought to have accelerated most of the relativistic electrons. If we take the typical spiral disk field strength $B \\sim 5 \\mu G$ and comparable magnetic energy density with inverse-Compton radiation energy density, the synchrotron lifetime of the relativistic electrons is $\\lesssim 10^{8}$ yr while emitting at 1.5 GHz. Therefore, RC can probe very recent star formation activity in normal star-forming galaxies whose radio emission is not dominated by the active galactic nucleus (AGN; Condon 1992). On global scales, the FIR and RC emissions are linearly correlated over 5 orders of magnitude in luminosities for various galaxies from dwarf galaxies, normal spiral galaxies, irregular galaxies, to starbursts, Seyferts, radio-quiet quasars (Condon et al. 1991), and local LIRGs/ULIRGs (Yun, Reddy \\& Condon 2001), out to the most extreme star-forming galaxies in the early universe, at $z$ = 1 and beyond (Appleton et al. 2004). The FIR-RC correlation for the \\IRAS 2 Jy galaxy sample (Yun et al. 2001, including 1809 galaxies) is well described by a linear relation, and over 98\\% of the sample galaxies follow this tight FIR-RC correlation. This correlation offers a potential method of deriving the SFR using the measured RC luminosity (Yun \\etal 2001) even if the FIR luminosity is unknown, and the RC images could potentially be used to judge how the FIR distribution would look like on detailed small scales in galaxies (Condon \\etal 1996). With the high resolution capability of \\Spitzer \\ (and now {\\it Herschel}), detailed local FIR-RC correlation is also shown to be valid (\\eg, Murphy \\etal 2006, 2008). Many researchers have conducted plenty of observations of CO emission in galaxies, and the global correlations between FIR and CO (\\eg, Devereux \\& Young 1990; Young \\& Scoville 1991), and between RC and CO (Rickard, Turner \\& Palmer 1977; Israel \\& Rowan-Robinson 1984; Adler, Allen \\& Lo 1991; Murgia \\etal 2002) are also well established for different samples. Adler \\etal (1991) and Murgia \\etal (2002, 2005) found that the RC-CO correlation is linear from global scale down to $\\sim 100$ pc size scale, at which scale there is still no evidence that this correlation is going to break down, with a dispersion that is less than a factor of 2. Therefore, the mean star formation efficiency (SFE), which measures the SFR (deduced from RC) per unit mass of molecular gas available to form stars, is found to vary weakly with Hubble morphological type (among galaxies) and distances from galaxy centers (within individual galaxy disks; Murgia \\etal 2002). For the HCN survey sample (Gao \\& Solomon 2004a,b, hereafter GS04a,b) of 65 spiral galaxies, starbursts, LIRGs, and ULIRGs, we show here that this FIR and RC correlation also holds and appears to be the best correlation among all correlations. Before the HCN survey of Gao \\& Solomon (GS04b), the HCN emission in external galaxies has only been observed in less than 30 galaxies, and only a few of them are measured globally to derive the total HCN emission (\\eg, Sorai \\etal 2002; Shibatsuka \\etal 2003). Based on this systematic HCN survey, a tight linear FIR-HCN correlation has been established (GS04a). Most recent follow-up observations of dense molecular gas tracers, such as HCN, HCO$^+$, HNC, etc., in galaxies (Baan et al. 2008; Gracia-Carpio et al. 2008) mostly confirm this correlation. Some theoretical models also predict such linear correlations (\\eg, Krumholz \\& Thompson 2007). According to this linear FIR-HCN correlation, a Schmidt (1959) star formation law in terms of dense molecular gas is established with a power-law index of 1.0. Furthermore, compared with the FIR-CO correlation, the FIR-HCN is linear and tighter as the authors argued that the combination of the stronger correlations between FIR-HCN and between HCN-CO may account for the FIR-CO correlation. Based on these comparisons, they also argued that the amount of the dense molecular gas traced by HCN, but not the total amount of molecular gas traced by CO, could be a critical molecular parameter that measures SFR in star-forming galaxies (GS04a). These new results from the HCN survey obviously invoke a question on the comparison between the correlations of RC-HCN versus RC-CO: is the RC-HCN correlation significantly better than the RC-CO correlation? How do the FIR and RC as SFR indicators relate with the star formation materials, the molecular gas (CO), and particularly the dense molecular gas (HCN)? Stars are born mainly in GMCs. The total mass of molecular gas of GMCs can be determined from the CO luminosity. However, the excesses of the CO luminosity in GMC cores where active high-mass star formation occurs are not specific enough to reveal the star formation potential of the dense cores. The physical conditions of star-forming GMC cores (Evans 1999) are better revealed by emission from high dipole-moment molecules, like HCN, whose emission traces the dense molecular gas ($n(H_2) \\approxgt 3 \\times 10^4$cm$^{-3}$) associated with the cores of the star-forming GMCs (GS04b). FIR-HCN is shown to be linearly correlated (GS04a). Does HCN also strongly correlate with RC? What is the role of dense molecular gas in the FIR-RC relationship? Here, we utilize this HCN sample to analyze the various relationships among the global HCN, FIR, CO, and RC luminosities to answer some questions raised here, and to further demonstrate the possibility of using the RC luminosity instead of FIR luminosity as an indicator of the SFR. The HCN sample and the total RC luminosity of the sample are presented in Section 2. Section 3 presents the results and analysis of the various correlations of these global luminosities and their ratios and shows that the global correlation between RC and HCN appears to be a combinational result of two tightest correlations between FIR-RC and between FIR-HCN even though RC-HCN correlation is significantly better than the RC-CO correlation. Discussion on the possible physical relationship between the HCN (dense molecular gas tracer) and the RC emission (the indicator of the rate of high-mass star formation) in galaxies is presented in Section 4, followed by the main points of this study in Section 5. ", "conclusions": "\\subsection{The Most Fundamental Correlation: FIR-RC versus FIR-HCN} In the RC-HCN correlation shown in Figure~1(a), no significant differences can be seen between the AGN host galaxies and others (the galaxies with AGN embedded are represented by stars, a slight excess in RC appears to be present at high luminosity end). This indicates that these AGNs (except those in ULIRGs) do not show stronger excess in RC than the rest of the star-forming galaxies. Namely, AGN contribution to RC is, on average, presumably roughly less than half of the total RC emission, and not much systematic difference that can be seen globally in this HCN sample of star-forming galaxies and LIRGs/ULIRGs (GS04a). Similarly, there is also no prominent difference between the AGN host galaxies and the rest of the star-forming galaxies in the FIR-HCN correlation (GS04a). The correlation between global FIR and RC luminosities is found to be tight and linearly valid over 5 orders of magnitude in thousands of galaxies (\\eg, Yun \\etal 2001). This is also reinforced here by the HCN survey sample even though the ranges in luminosities are over only 3 orders of magnitude due to very limited sample size. Moreover, radio AGNs usually show distinct excess in RC over FIR as compared to the general populations of normal galaxies in the FIR-RC correlation. But, AGNs in the HCN sample do not particularly show any significant radio emission excess. As detailed above, it turns out that the FIR-RC is the tightest correlation among all. But the FIR-HCN correlation is essentially as tight as the FIR-RC correlation, and these two correlations are the tightest. It might even be possible that the FIR-HCN correlation could be tighter than the FIR-RC correlation, especially in view of the significantly larger uncertainties in the estimate of HCN than other parameters. This is also intuitively implied since both FIR and HCN are directly and physically related to active massive star formation. Several models were previously posed to interpret this correlation (\\eg, Voelk 1989; Helou \\& Bicay 1993; Niklas \\& Beck 1997; Murgia \\etal 2005). Most recently, Lacki, Thompson, \\& Quataert (2009) suggest that the FIR-RC correlation is a combinational result of the efficient cooling of cosmic-ray (CR) electrons in starbursts and a conspiracy of several factors: the decrease in the radio emission due to bremsstrahlung, ionization, and inverse Compton cooling in starbursts is countered by secondary electrons/positrons and the decreasing critical synchrotron frequency, which both increase the radio emission; low effective UV dust opacity leads to the decreasing FIR emission, which balances the decrease in radio emission caused by CR escape for lower surface density galaxies. In fact, this kind of calorimetry theory that galaxies act in a reasonable approximation as calorimeters for the stellar UV radiation and for the energy flux of the CR electrons was proposed 20 years ago (Voelk 1989). Although the molecular gas connection to the FIR-RC correlation in galaxies is proposed by Murgia \\etal (2005), we demonstrated here that the addition of the total dense molecular gas traced by HCN appears to be globally insignificant in providing the dense star-forming gas connection between the FIR (SFR) and RC (the RC(\\Lhcn, \\Lir, \\Lco \\ ) model, Equations (2), (5), and B(3)). Yet, the contribution of CO is non-negligible and appears to be more important than HCN. Nevertheless, the star formation materials, particularly the dense gas (HCN), definitely play an important role in the obvious connection of SFR to both the FIR and RC emission. Both HCN and RC contribute nearly equally to FIR and there is no additional contribution to FIR by adding CO (the IR(\\Lhcn, \\Lrc, \\Lco) model, eqs. (4), (7) and B(3)). Thus, CO seems to be the least important one among all the four parameters RC, FIR, CO, and HCN in relating to SFR even though they all show quite good correlations with each other. \\subsubsection{The Origin of the Tightest Correlation: FIR-RC and FIR-HCN} In conventional models, the RC emission produced from synchrotron radiation is related with massive star formation via several steps: star formation $\\rightarrow$supernovae and SNRs $\\rightarrow$relativistic electrons accelerated by SNRs. So, too, is the FIR emission related: star formation$\\rightarrow$UV radiation$\\rightarrow$absorption and re-emission by dust enshrouding new-born stars. These should all be expected to be the dominant FIR and RC emission in active star-forming galaxies. Nevertheless, CR electrons involve with other radiation processes, generating mostly the background RC emission that is not related to massive star formation (the general interstellar radiation field). Additionally, the old stars also heat up the dust and additionally contribute to the general interstellar radiation field. The contribution of the general interstellar radiation field to the total FIR emission in galaxies might be significant at the low-\\Lir\\ end, where the general infrared interstellar radiation field is comparable to or even more dominant than the FIR radiation from the active star formation. The temporal and spatial inconsistencies between these processes could be the obstacles to understand fully the physical origin of the FIR and RC correlation. Given the obvious star formation connection of both FIR and RC and the important role of the star formation materials, i.e., the molecular gas, especially the dense molecular gas, in actually giving births of massive stars, it can surely help us figure out which of these correlations are more fundamental, and which are possibly indirectly less fundamental correlations. Globally, we have seen the bigger role of HCN than CO in the FIR-RC correlation. Locally, the spatially resolved studies perhaps help us much better to make the whole picture clearer, by taking the molecular gas and dense molecular gas into account in the analysis of various comparisons of the correlations. Temporally, massive stars live $\\sim 10^{7}$ yr and the relativistic electrons probably have lifetimes of $\\sim 10^{8}$ yr in star-forming galaxies; the RC luminosity therefore probes star formation activity not much earlier than $10^{8}$ yr ago. Taking the molecular gas in the dense phase as the initial stage in initiating the massive star formation procedure, and the total molecular gas as the potential star formation gas reservoir as a necessity in eventually forming the dense cores of GMCs, we can sort the three probes of SFR into a time line: CO (GMCs, assembly of star formation material---a necessity for later core collapse and star formation in GMCs)$\\rightarrow$HCN (probes dense GMC cores, massive star formation sites)$\\rightarrow$FIR (probes new-born young massive stars)$\\rightarrow$RC (probes the end products of the short-lived massive stars). If the dense cores of GMCs are forming massive stars on a time-scale smaller than the depletion time of dense gas in galaxies ($\\leq 10^7$ yr; GS04a), followed by the life span of massive stars, then this is phenomenologically consistent with our result that only HCN-FIR and FIR-RC correlations are the tightest and have the least scatter from linearity, while HCN and RC luminosity appears to be the result of a combination of the former two correlations, not an immediate relation in time. The obvious contrast between Figure~3(a) (or FIR-HCN normalized by CO (Figure~5(a) in GS04a)) and Figure~5(b) might be reflected in the differences in these time-scales that it is much more dynamic and scattered in luminosity ratios involving HCN rather than CO. Spatially, CO traces large-scale molecular gas distribution (entire GMCs plus diffuse molecular clouds), where most of the molecular gas is not forming stars, whereas the high-density regions (dense GMC cores traced by high-dipole moment molecules such as HCN) of much smaller spatial scales are indeed the locations of active star formation in galaxies. Murgia \\etal (2005) showed that the spatially resolved tight CO-RC correlation holds down to $\\sim 100$pc size scale in galaxies which is close to the GMC size scale. Paladino et al. (2006) further probed the RC-FIR-CO correlations down to linear scales of a few hundred pc using new {\\it Spitzer} IR images of six BIMA CO Survey of Nearby Galaxies (BIMA SONG, Helfer et al. 2003) and observed local deviations from the correlations in regions with a high SFR where a low RC/FIR ratio is found. We showed here, however, that the global HCN-RC correlation is actually much tighter than the CO-RC correlation, but we need a similarly detailed, spatially resolved local comparison between RC and HCN, and FIR and HCN in galaxies in order to examine how the FIR-RC-HCN correlations extend to much smaller size scale. Murphy \\etal (2006, 2008) have taken an initial look at the FIR-RC correlation within the disks of 4, and later increased to 29, nearby face-on galaxies in the {\\it Spitzer} SINGS legacy program (Kennicutt \\etal 2003), and found the trend that the ratio of FIR to RC decreases with increasing radius, which is consistent with what Marsh \\& Helou (1995) found at intermediate spatial resolution. They also studied how the star formation activity affects the FIR-RC correlation within galaxies by testing a phenomenological model which smears the FIR images to match the radio images. They found that the mean distance traveled by the CR electrons is most sensitive to the dominant age of the CR electron population, rather than the interstellar medium (ISM) parameters, which may inhibit their propagation, such as the ISM density, radiation-field energy density, and magnetic field strength. Comparison of such detailed spatially resolved correlations in line with our findings in the global quantities could help us reach our final goal, i.e., understanding the FIR-RC correlation. Globally, HCN seems better than CO, the validity of correlations of RC-CO and FIR-RC on small scale has already been proven; we need high-fidelity HCN imaging and/or resolved HCN observations to truly compare all. Both FIR and RC emission involves physical processes of both small and large spatial scales even though most of the FIR emission is dominated by active star-forming regions of small scales. Mooney \\& Solomon (1988) showed that the FIR-CO correlation for GMCs improves, once the diffuse FIR emission that originated from the general interstellar radiation field of large spatial scale was subtracted. RC might be dominated by large-scale shocks/bubbles associated with SNRs as well as even larger scale of magnetic fields, where CR electrons pass along the field lines and experience efficient cooling, whereas HCN traces a smaller size scale than that of FIR and RC. Yet, they are associated with three different periods in the time sequence connected with the entire star formation processes. FIR emission originates from the dust that enshrouds the new-born stars while HCN emission outlines regions of dense molecular gas that eventually nurse new-born stars, whereas RC emission is produced from the significantly diffused CR electrons and large-scale shocks, which have traveled a long distance from the previous star-forming sites. The difference between these three emissions in time sequence is in de facto agreement with that of the corresponding locations in spatial scales: dense molecular cores further collapse to form massive stars that then quickly evolve and go through supernovae to become SNRs. In short, CO and FIR emission are almost entirely associated with the GMCs, but most FIR emission is probably associated with star-forming sites inside the dense cores of the much smaller scales traced by the HCN emission. On the other hand, RC is probably dominated by large-scale diffuse emission though some FIR emission is also associated with the large-scale general interstellar radiation field. These differences are possibly suggested by the multi-parameter fits that RC correlates much tighter with FIR than with HCN though CO seems more important than HCN in predicting RC (Equations (2) and (5)), and HCN correlates much tighter with FIR than with RC, though CO seems to be as equally important as FIR in predicting HCN (Equations (3) and (6)). Both FIR and HCN trace smaller scale star-forming regions, whereas RC traces the overall large-scale environment, which are affected by the feedback of massive star formation and previously hosted massive stars. \\subsection{Radio Continuum (RC, and FIR and HCN) as Star Formation Rate (SFR) Indicator} Gao \\& Solomon (GS04a) have extensively discussed the dense molecular gas in relationship to the total molecular gas and SFR indicated by the FIR emission. The HCN emission is associated with the high-density molecular gas which is the direct active star formation material. Compared with CO, the HCN luminosity is much better at predicting the FIR luminosity for all galaxies including ULIRGs (GS04a). This is represented by a much better \\Lir--\\Lhcn \\ correlation than the \\Lir--\\Lco \\ correlation (GS04a). A similar result is obtained here by the comparison between the \\Lrc--\\Lhcn \\ correlation and the \\Lrc--\\Lco \\ correlation. Although the extreme claim that globally, the RC-CO correlation is as good as the FIR-RC correlation are made by Murgia \\etal (2005), perhaps due to a small range of parameters and limited sample size, we find here that RC luminosity correlates with \\Lhcn \\ much more tightly than with \\Lco, and that the FIR-RC and FIR-HCN are the tightest correlations among all. The physical explanation for the tight correlation between the HCN and FIR is obviously that stars are formed in dense molecular gas, whereas the tight correlation between HCN and RC would further strongly support this interpretation and at least globally indicate that the RC could be used as a tracer of star formation. For a detailed discussion on SFR, dense molecular gas, total molecular gas, and their various ratios readers are referred to Section 4 in Gao \\& Solomon (GS04a). According to the tight FIR-RC correlation which is also valid down to kpc/sub-kpc size scale within galaxies (e.g., Marsh \\& Helou 1995; Lu et al. 1996; Murphy \\etal 2006, 2008), detailed SFE maps can be deduced by using local RC luminosity as a local SFR tracer and comparing them to the CO images. This application has been first used in studies of individual interacting galaxies, like Arp 244---pioneered by Gao \\etal (2001)--- and Taffy galaxy (Gao, Zhu, \\& Sequist 2003; Zhu \\etal 2007). These SFE maps allow us to identify the most active star-forming sites, and characterize and investigate the star formation properties in local regions of kpc/sub-kpc size scales. Nonetheless, the mean SFE, which measures the SFR (deduced from RC) per unit mass of molecular gas available to form stars, is found to vary weakly with Hubble morphological type (among galaxies) and distances from galaxy centers (within individual galaxy disks; Murgia \\etal 2002). Murgia \\etal (2005), however, do not measure any systematic trend in the CO/RC ratio as a function of radius in the nine BIMA SONG galaxies studied. A hydrostatic pressure regulation model was used to interpret the excellent correlation between the CO, RC, and FIR emissions in galaxies on both large and small scales to avoid invoking any explicit dependence on the star formation scenario (Murgia \\etal 2005). Given the assumption that CO surface brightness is proportional to molecular gas surface density, the model predicts $I_{\\rm RC} \\propto I_{\\rm CO}^{1.4}$, which is consistent with our global result and also, probably results from the linear correlation between FIR-RC, consistent with the 1.4 power in global FIR versus CO (GS04a) and FIR versus (\\htwo +HI) (Kennicutt 1998b) as well as recent local resolved studies (Calzetti et al. 2007; Kennicutt et al. 2007). We note that if the observational result that $\\Lhcn = 1.38 log \\Lco -4.79$ (GS04a) is used, the prediction of this model will become $I_{\\rm RC} \\propto I_{\\rm HCN}$, also perfectly consistent with the result of our global correlation fits. However, Paladino \\etal (2006) observed local deviations from the RC-IR-CO correlations in regions with a high SFR, such as the spiral arms, in six galaxies for which high-resolution {\\it Spitzer} 24$\\mu$m mid-IR data are available. Further studies in a variety of star-forming galaxies are necessary to enlarge the dynamical range of parameter spaces and to achieve better agreements in the local RC-FIR-CO and even RC-FIR-CO-HCN correlations. Nevertheless, Paladino \\etal (2006) concluded that, down to $\\sim 100$ pc scale, they have not yet probed the physical scales at which the correlations break down. Based on \\Spitzer mid-IR observations, Wu et al. (2005) find that the 8$\\micron$ and 24$\\micron$ luminosities of star-forming galaxies are both strongly correlated with their RC (1.4GHz) and H$\\alpha$ luminosities over a range in luminosities of nearly 3 orders of magnitude. This suggests that, alternatively, mid-IR emission of much better spatial resolution than FIR can be approximately used as local tracers of SFR as well. Yet, it is still unclear how quantitatively different the various correlations between each SFR indicator: from mid-IR to FIR, and from RC to molecular gas tracers CO and HCN, could be, on local sub-kpc scales. With the help of the high-resolution IR capability of \\Spitzer, and now {\\it Herschel}, as well as high-resolution and high-sensitivity HCN and CO imaging (with the upcoming ALMA), we can further analyze these correlations down to sub-kpc scales in galaxies and extremely smaller scale (sub-pc) inside the dense GMC cores in the Milky Way. Whether these relations hold or break down will help make our understanding of star formation in both external galaxies approaching microscopic scales and the active star-forming regions in our Galaxy more comprehensive." }, "1003/1003.4263_arXiv.txt": { "abstract": "We analyze optical and radio properties of radiogalaxies detected in the Sloan Digital Sky Survey (SDSS). The sample of radio sources are selected from the catalogue of Kimball \\& Ivezi\\'c (2008) with flux densities at 325, 1400 and 4850 MHz, using WENSS, NVSS and GB6 radio surveys and from flux measurements at 74 MHz taken from VLA Low-frequency Sky Survey \\citep{cohen}. We study radiogalaxy spectral properties using radio colour-colour diagrams and find that our sample follows a single power law from 74 to 4850 MHz. The spectral index vs. spectroscopic redshift relation ($\\alpha-z$) is not significant for our sample of radio sources. We analyze a subsample of radio sources associated with clusters of galaxies identified from the maxBCG catalogue and find that about 40\\% of radio sources with ultra steep spectra (USS, $\\alpha<-1$, where $S_\\nu \\propto \\nu^{\\alpha}$) are associated with galaxy clusters or groups of galaxies. We construct a Hubble diagram of USS radio sources in the optical $r$ band up to $z\\sim0$.8 and compare our results with those for normal galaxies selected from different optical surveys and find that USS radio sources are around as luminous as the central galaxies in the maxBCG cluster sample and typically more than 4 magnitudes brighter than normal galaxies at $z\\sim0$.3. We study correlations between spectral index, richness and luminosity of clusters associated with radio sources. We find that USS at low redshift are rare, most of them reside in regions of unusually high ambient density, such of those found in rich cluster of galaxies. Our results also suggest that clusters of galaxies associated with steeper than the average spectra have higher richness counts and are populated by luminous galaxies in comparison with those environments associated to radio sources with flatter than the average spectra. A plausible explanation for our results is that radio emission is more pressure confined in higher gas density environments such as those found in rich clusters of galaxies and as a consequence radio lobes in rich galaxy clusters will expand adiabatically and lose energy via synchrotron and inverse Compton losses, resulting in a steeper radio spectra. ", "introduction": "Radio sources are frequently associated with massive systems at low and high redshifts. In the Local Universe these objects are usually identified with evolved red ellipticals and with luminous cD galaxies located at the centres of clusters of galaxies \\citep{west}. Radio sources at high redshifts ($z\\sim4$) are identified with massive forming systems, i.e galaxies with diffuse UV morphologies and consistent with small substructures around a dominant bright galaxy \\citep{miley}. Several works show that some distant radio sources are embedded in highly spatially extended ionized gas nebulae of 100-200 kpc \\citep{vene02, villar}. These gas structures are observed in massive ellipticals or cD galaxies at the center of nearby clusters of galaxies. This evidence suggests that distant radio sources represent the progenitors of the most massive galaxies observed in the Local Universe and therefore important for the study of structure formation, such as clusters or groups of galaxies. A high fraction of radio sources associated with clusters of galaxies have steep radio spectra as measured by the slope between two fixed ranges of frequencies, i.e the spectral index $\\alpha$ (USS, $\\alpha<-1$, where $S_{\\nu} \\propto \\nu^{\\alpha}$, \\citet{rot,cha96,jarvis01}). Such a correlations between spectral index and redshift ($\\alpha-z$ relation), has been used with success to search for distant galaxies \\citep{DB00, jarvis04, cruz06, brode}. There are at least three main explanations given in the literature for this phenomenon: The first is a radio k-correction. The spectral energy distribution (SED) of radio sources is usually concave, i.e the spectral index increases with frequency. The observed spectral index determined from two fixed observed frequencies will therefore sample a steeper part of the radio spectrum for sources at higher redshifts. The second is related with the interaction between photons of the cosmic microwave background radiation (CMB) and the relativistic electron population observed in the plasma gas of radio sources. The energy density of the CMB increases as $(1+z)^4$ and hence Inverse Compton scattering (IC) of the CMB becomes increasingly important for sources at high redshifts \\citep{krolik}. The third is related to an intrinsic correlation between radio luminosity and {\\it rest-frame} spectral index \\citep{blun, cham} that due to a Malmquist like bias in flux density limited surveys translates into a correlation between spectral index and redshift. An alternative explanation is related with an ambient density effect. The presence of USS radio sources residing closest to cluster centres may be due to a manifestation of pressure-confined radio lobes which slow adiabatic expansion of the plasma. Radio lobes will be pressure-confined and lose energy primarily via synchrotron and IC losses \\citep{klamer}. Studies of radio sources at low and moderate redshifts ($z <0.5$) show different environments. Hill \\& Lilly (1991) find that only 50\\% of powerful radio sources at $z\\sim0.5$ are located in rich galaxy clusters, even though similar sources avoid such environments at low redshifts. However, Geach et al. (2007) analyzed radio sources in the Subaru/XMM-Newton Deep Field and found that low-power radio galaxies at $z\\sim0.5$ reside in moderately rich groups - intermediate environments between poor groups and rich clusters. \\citet{prestage} investigate the local galaxy density around powerful radio sources using the angular cross-correlation technique and find that compact radio sources appear to lie in regions of low galaxy density. Moreover, that complex Fanaroff-Riley class I sources are typically found in regions of significantly enhanced galaxy density. \\citet{alli} study the evolution of galaxies in radio-selected groups at $z<0.5$ with the same range of radio power and find that strong radio galaxies are located in a wide range of environments but not as wide as for groups in general. At low redshifts ($z<0.1$) radio-selected groups have the same richness and blue fraction as do optically-selected groups. At high redshifts ($z\\sim 0.4$) all groups have the same proportion of blue galaxies. Studying the prevalence of radio loud AGN activity in nearby groups and galaxy clusters, selected from the SDSS catalog, \\citet{best} find that brightest group and cluster galaxies are more likely to host a radio loud AGN than other galaxies of the same stellar mass. In this work we analyze optical and radio properties of radio sources with steep radio spectra, and we determine the main characteristics of clusters of galaxies associated with radio sources. The structure of this paper is organized as follows: Section 2 describes the optical and radio samples analyzed. We analyze radio spectral properties of our sample in Section 3. Section 4 compares the optical and the radio luminosities of sources associated with steep spectrum and central galaxy clusters. In section 5, we study the Hubble diagram in the optical $r$ band for USS sources. We study in Section 6 spectral index and richness properties of galaxies associated with clusters of galaxies and finally, in Section 7, we discuss our main results. Throughout this work we assume a standard $\\Lambda$CDM model Universe with cosmological parameters, $\\Omega_{M}$=0.3, $\\Omega_{\\Lambda}$=0.7 and a Hubble constant of $H_0=$100~h~km~s$^{-1}$Mpc$^{-1}$. ", "conclusions": "We study optical and radio properties of radiogalaxies detected in the Sloan Digital Sky Survey (SDSS) with flux densities of 74, 325, 1400 and 4850 MHz, using the VLSS, WENSS, NVSS and GB6 radio catalogues. We search for a possible empirical correlation between the spectral index and redshift, however we find no significant trend. We analyze the functional form of the SED using colour-colour diagrams at radio frequencies. We do not find a clear tendency of radio sources to show flattening towards low frequencies, as expected assuming concave curvature in the radio SED. It is well known that a narrow relation exists between $K-$band and redshift as observed in Hubble diagrams \\citep{debreuck02b, willott,jarvis01, eales}. In this work we construct a Hubble diagram of USS radio sources in the optical $r$ band to $z\\sim0$.8. Despite any k-correction and possible extinction effects, our $r$ band Hubble diagram (Figure 6) also clearly shows a tight correlation. We find that USS radio sources are as bright as central galaxies in the maxBCG cluster sample and are typically more than 4 magnitudes brighter than normal galaxies at $z\\sim0$.3. We note that this result is not entirely new, for example \\citet{debreuck02b} also find that at redshifts $<\\sim1$ radio-loud galaxies define the luminous envelope using near infrared $K-$band magnitudes. Regarding the possible dependence of radio luminosity on environment, we notice that the radio luminosity distribution of USS, radio sources in general, and radio sources in clusters are remarkably similar (Figure 4), indicating that USS prefer higher denstity environments, independent of radio luminosity.These results are consistent with those by \\citet[fig. 8,][]{hill} who find no significant correlation in cluster richness $N_{0.5}$ and the rest-frame 2.0 GHz radio power for a sample of radio sources with $z<0.$5. Similar results were obtained by \\citet{alli} who find no trend of richness with radio luminosity at 408 MHz for a sample of radiosources with $z<0.$5. We also analyze the richness and spectral index properties of clusters of galaxies associated with radio sources and find that 40\\% of USS sources identified in the SDSS spectroscopic catalogue are associated with cluster or groups of galaxies identified in the literature, such as in the maxBCG, Abell, or Zwicky catalogues. We analyze the local density of galaxies around the sample of USS sources without a know cluster association from the literature, using the $\\Sigma_5$ and N1 estimators and find that these USS sources have similar galaxy densities to clusters selected from the maxBCG catalogue. We also find that USS sources at low redshift are rare objects (99 from a total sample of 2885 radio sources detected in the SDSS spectroscopic catalogue). However a majority reside in regions of unusually high ambient density, such as those regions found in rich cluster of galaxies. Our results complement those found by \\citet{DB00}. These authors define a sample of 669 USS sources selected from the WENSS, TEXAS, MRC, NVSS and PMN radio surveys. They conclude that the majority of relative nearby ($z<\\sim 0.4$) USS objects are located in galaxy clusters. They find that at least 85\\% of the X-ray objects associates with USS sources are galaxy clusters or known groups from the literature. At lower redshifts, we find that radio sources with $\\alpha_{325}^{1400}< -0.65$, are preferentially located in galaxy cluster environments. This result contrast with \\citet[fig. 7,][]{prestage} where it is found no dependence of the spatial cross--correlation amplitude on spectral index. We note although that this statistical analysis concerns more the large-scale, rather than the local environment of our study. We also find that clusters hosting radio sources with spectra steeper than the average have a higher galaxy richness and are populated by brighter galaxies in comparison to clusters associated to radio sources with $\\alpha_{325}^{1400}> -0.65$. A natural explanation for these correlations is that radio emission in rich cluster of galaxies is pressure-confined in a high gas density environment. Radio lobes in galaxy cluster environment will expand adiabatically and lose energy via synchrotron and inverse Compton losses, resulting in a steeper radio spectra \\citep{klamer}. \\clearpage \\begin{figure*} \\includegraphics[width=15cm]{fig14.ps} \\caption{Cutout colour images of a subsample of galaxies associated with Ultra Steep radio sources in the SDSS catalogue. The open cross indicates the radio position taken from the FIRST survey. The ID and redshift (increasing from left to right) are indicated above each plot. Information for the complete sample is listed in Table~1.} \\label{overlays} \\end{figure*} \\small \\begin{table*} {\\bf Table 1.} Sample of USS radio sources in the SDSS catalogue with spectroscopic redshifts.\\\\ \\begin{center} \\begin{small} \\begin{tabular}{rrrrrcrc} \\hline ID & R.A$^{radio}_{J2000}$ & DEC$^{radio}_{J2000}$ & z & $\\alpha_{352}^{1400}$ & Log($L_{1.4}$)& $M_r$ & ID from literature\\\\ & $^h\\;\\; ^m\\;\\;\\;\\; ^s\\;\\;\\,$ & \\degr$\\;\\;\\;$ \\arcmin$\\;\\;\\;$ \\arcsec$\\;$ &&& W Hz$^{-1}$&&Designation\\\\ \\hline 608550 &07 25 57.08&$+$41 23 05.13& 0.1113& $-$1.40& 23.76 & $-$22.18 &{\\tt MaxBCG J111.48808+41.38519}\\\\%, d=0.7} \\\\ 664059 &07 51 31.86&$+$43 49 29.49& 0.4249& $-$1.00& 24.22 & $-$22.63 & \\nodata \\\\ 668611 &07 53 32.47&$+$38 57 52.71& 0.1484& $-$1.04& 23.37 & $-$21.71 & \\nodata\\\\ 672034 &07 54 57.66&$+$38 15 22.71& 0.3030& $-$1.10& 23.73 & $-$20.00 & \\nodata\\\\ 683214 &07 59 49.48&$+$35 32 33.82& 0.4823& $-$1.04& 25.19 & $-$22.54 & \\nodata\\\\ 709796 &08 10 54.66&$+$49 11 03.90& 0.1147& $-$1.24& 23.28 & $-$22.27 &{\\tt MaxBCG J122.72750+49.18436}\\\\ 717149 &08 13 50.80&$+$39 32 32.11& 0.2045& $-$1.03& 23.90 & $-$21.79 &{\\tt MaxBCG J123.46125+39.54183}\\\\ 734111 &08 20 32.39&$+$30 34 48.65& 0.3628& $-$1.04& 25.98 & $-$21.97 & \\nodata\\\\ 735423 &08 21 03.64&$+$52 44 35.82& 0.4441& $-$1.38& 25.08 & $-$22.79 & \\nodata\\\\ 748224 &08 26 00.38&$+$40 58 51.75& 0.0576& $-$1.11& 22.78 & $-$22.79 &{\\tt SDSS-C4-DR3 3247}\\\\ 780016 &08 38 23.27&$+$29 45 21.67& 0.1068& $-$1.06& 22.68 & $-$22.00 & {\\tt SHK 182 GGroup}\\\\%, d=0.9}\\\\ 794634 &08 43 59.18&$+$51 05 25.55& 0.1264& $-$1.12& 24.37 & $-$22.63 & \\nodata\\\\ 801800 &08 46 37.85&$+$51 27 16.56& 0.1800& $-$1.27& 23.88 & $-$21.71 &{\\tt MaxBCG J131.65796+51.45436}\\\\ 814626 &08 51 17.29&$+$37 04 29.00& 0.2207& $-$1.00& 24.11 & $-$20.32 &\\nodata\\\\ 837935 &08 59 57.32&$+$56 47 12.15& 0.1833& $-$1.06& 23.71 & $-$22.11 & \\nodata\\\\ 838957 &09 00 20.28&$+$52 29 39.73& 0.0302& $-$1.02& 22.01 & $-$24.49 & {\\tt CGCG 264-047}\\\\ 842071 &09 01 30.10&$+$55 39 16.42& 0.1155& $-$1.80& 23.81 & $-$22.75 & {\\tt MaxBCG J135.37558+55.65463 }\\\\ 846087 &09 03 00.14&$+$35 27 04.82& 0.3488& $-$1.01& 25.53 & $-$21.09 & \\nodata\\\\ 848084 &09 03 44.85&$+$41 38 19.31& 0.2189& $-$1.01& 24.42 & $-$21.42 & {\\tt MaxBCG J135.93682+41.63908}\\\\ 932725 &09 34 42.25&$+$35 14 16.48& 0.4618& $-$1.15& 24.31 & $-$22.62 & \\nodata\\\\ 940538 &09 37 37.11&$+$37 05 35.37& 0.4492& $-$1.01& 25.24 & $-$22.53 & \\nodata\\\\ 955678 &09 43 09.29&$+$29 50 18.31& 0.2969& $-$1.05& 24.42 & $-$18.71 & \\nodata \\\\ 981931 &09 52 49.14&$+$51 53 04.99& 0.2151& $-$1.77& 24.08 & $-$22.49 & {\\tt ZwCl 0949.6+5207}\\\\%, d=1 } \\\\ 989139 &09 55 29.87&$+$60 23 17.47& 0.1989& $-$1.55& 23.82 & $-$22.03 & {\\tt MaxBCG J148.87452+60.38814} \\\\ 1002970&10 00 31.01&$+$44 08 42.94& 0.1533& $-$1.44& 23.16 & $-$21.32 & {\\tt RBS 0819} \\\\ 1027752&10 09 28.28&$+$46 17 37.17& 0.3858& $-$1.04& 24.29 & $-$19.86 & \\\\ 1051000&10 17 58.13&$+$37 10 54.02& 0.0442& $-$1.16& 22.04 & $-$22.34 & {\\tt CGCG 183-009 } \\\\ 1075981&10 27 09.98&$+$39 08 06.04& 0.3378& $-$1.02& 24.93 & $-$23.20 & \\nodata \\\\ 1080596&10 28 54.68&$+$48 09 38.20& 0.4851& $-$1.45& 24.99 & $-$22.19 & \\nodata \\\\ 1109542&10 39 32.11&$+$46 12 05.54& 0.1864& $-$1.21& 24.85 & $-$21.92 & \\nodata \\\\ 1128337&10 46 25.51&$+$59 37 37.59& 0.2282& $-$1.07& 23.56 & $-$22.91 & {\\tt MaxBCG J161.60635+59.62690} \\\\ 1138526&10 50 10.03&$+$32 22 05.09& 0.1150& $-$1.33& 23.01 & $-$21.94 & {\\tt NSCS J105005+322256 /CLtr } \\\\ 1160615&10 58 19.46&$+$41 03 40.76& 0.1299& $-$1.07& 23.41 & $-$22.44 & {\\tt MaxBCG J164.58100+41.06140 } \\\\ 1177112&11 04 33.11&$+$46 42 25.96& 0.1410& $-$1.82& 23.45 & $-$21.08 & \\nodata \t\t\t\t\\\\ 1179743&11 05 30.73&$+$31 14 36.74& 0.4381& $-$1.43& 24.12 & $-$21.34 & \\nodata \t\t\t\t\\\\ 1216462&11 18 45.25&$+$52 16 00.95& 0.4309& $-$1.24& 24.60 & $-$22.18 & \\nodata \t\t\t\t\\\\ 1228632&11 23 22.90&$+$47 55 14.34& 0.1262& $-$1.03& 22.91 & $-$21.08 & \\nodata \t\t\t\t\\\\ 1233822&11 25 16.31&$+$42 29 10.97& 0.1882& $-$1.02& 24.12 & $-$23.00 & {\\tt ABELL 1253 } \\\\ 1239404&11 27 18.46&$+$53 02 21.12& 0.3236& $-$1.04& 24.49 & $-$23.56 & \\nodata \t\t\t\t\t\\\\ 1244341&11 29 01.60&$+$32 45 50.65& 0.5759& $-$1.10& 25.02 & $-$22.24 & \\nodata \t\t\t\t\\\\ 1250737&11 31 20.94&$+$33 34 46.95& 0.2219& $-$1.61& 23.72 & $-$22.26 & {\\tt MaxBCG J172.83707+33.57975}\\\\ 1260844&11 34 57.39&$+$53 46 24.20& 0.1695& $-$1.05& 23.24 & $-$21.52 & \\nodata \t\t\t\t \\\\ 1262111&11 35 26.68&$+$31 53 33.14& 0.2310& $-$1.05& 24.71 & $-$22.67 & {\\tt MACS J1135.4+3153}\\\\%, d=0.1 } \\\\ 1268493&11 37 50.23&$+$46 36 33.65& 0.3151& $-$1.03& 24.72 & $-$23.08 & \\nodata \\\\ 1286560&11 44 27.21&$+$37 08 32.42& 0.1148& $-$1.56& 23.55 & $-$21.11 & \\nodata \\\\ 1294316&11 47 12.35&$+$38 19 26.32& 0.5977& $-$1.01& 25.34 & $-$22.11 & \\nodata \\\\ 1304138&11 50 49.21&$+$62 19 49.04& 0.3453& $-$1.69& 24.37 & $-$23.04 & \\nodata \\\\ 1307368&11 51 58.63&$+$31 40 32.05& 0.5079& $-$1.04& 25.69 & $-$23.39 & \\nodata \\\\ 1307436&11 52 00.09&$+$33 13 42.49& 0.3573& $-$1.48& 23.97 & $-$23.15 & \\nodata \\\\ \\hline \\end{tabular} \\end{small} \\end{center} \\end{table*} \\small \\begin{table*} {\\bf Table 1.} \\\\ \\begin{center} \\begin{small} \\begin{tabular}{rrrrrcrc} \\hline ID & R.A$^{radio}_{J2000}$ & DEC$^{radio}_{J2000}$ & z & $\\alpha_{352}^{1400}$ & Log($L_{1.4}$) &$M_r$ & ID from literature\\\\ & $^h\\;\\; ^m\\;\\;\\;\\; ^s\\;\\;\\,$ & \\degr$\\;\\;\\;$ \\arcmin$\\;\\;\\;$ \\arcsec$\\;$ && &W Hz$^{-1}$ & &Designation\\\\ \\hline 1309157&11 52 36.33&$+$37 32 43.86& 0.2294& $-$1.19& 24.22 & $-$20.23 &{\\tt MaxBCG J178.15191+37.54548 } \\\\ 1319027&11 56 05.51&$+$34 33 05.33& 0.2536& $-$1.10& 25.05 & $-$21.83 &{\\tt [EAD2007] 200 Arcs } \\\\ 1326519&11 58 48.07&$+$57 17 19.11& 0.2598& $-$1.04& 25.12 & $-$24.49 & \\nodata \\\\ 1348701&12 06 47.88&$+$51 57 10.95& 0.3446& $-$1.28& 24.66 & $-$21.51 & \\nodata \\\\ 1351855&12 08 00.78&$+$43 39 19.12& 0.2657& $-$1.00& 24.77 & $-$22.57 &{\\tt MaxBCG J182.00318+43.65537 } \\\\ 1353531&12 08 37.16&$+$61 21 06.52& 0.2748& $-$1.48& 23.96 & $-$22.17 & \\nodata \\\\ 1354974&12 09 08.84&$+$44 00 11.30& 0.0376& $-$1.12& 22.12 & $-$22.98 &{\\tt NGC4135, (G. group)} \\\\ 1362044&12 11 46.22&$+$32 38 38.16& 0.6115& $-$1.07& 24.94 & $-$22.25 & \\nodata \t\t\t\t \\\\ 1406400&12 28 02.17&$+$34 40 40.12& 0.2775& $-$1.40& 23.81 & $-$22.53 &{\\tt MaxBCG J187.00902+34.67753} \\\\ 1439302&12 40 04.88&$+$37 44 15.46& 0.1879& $-$1.15& 23.75 & $-$21.53 & {\\tt NSC J124001+374544} \t\t\t\t \\\\ 1468909&12 51 07.51&$+$56 25 44.98& 0.2008& $-$1.22& 23.42 & $-$21.73 & \\nodata \t\t\t\t \\\\ 1505383&13 04 31.36&$+$51 43 42.64& 0.2757& $-$1.56& 24.43 & $-$22.28 &{\\tt MaxBCG J196.15441+51.71551}\\\\%, d=1.2 } \\\\ 1510127&13 06 12.17&$+$51 44 06.94& 0.2773& $-$1.16& 24.55 & $-$22.54 &{\\tt MaxBCG J196.55069+51.73530}\\\\%, d=0.4}\\\\ 1532575&13 14 18.32&$+$41 24 30.18& 0.1987& $-$1.05& 23.25 & $-$20.90 &{\\tt MaxBCG J198.57609+41.40825}\\\\%, d=1 }\\\\ 1547051&13 19 38.92&$+$61 39 11.68& 0.1333& $-$1.21& 23.49 & $-$22.47 & \\nodata \t\t\t\t \\\\ 1559285&13 24 12.38&$+$31 17 24.33& 0.4268& $-$1.19& 24.63 & $-$22.47 & \\nodata \t\t\t\t \\\\ 1604023&13 40 32.89&$+$40 17 38.79& 0.1719& $-$1.17& 23.32 & $-$22.35 &{\\tt RX J1340.5+4017 GGroup}\\\\%, d=0.2} \\\\ 1607910&13 41 59.68&$+$42 21 32.32& 0.4261& $-$1.22& 24.07 & $-$22.75 & \\nodata \t\t\t\t \\\\ 1627288&13 49 03.74&$+$30 52 27.51& 0.0814& $-$1.13& 22.53 & $-$19.43 & \\nodata \t\t\t\t \\\\ 1628350&13 49 27.88&$+$46 20 15.29& 0.4212& $-$1.14& 24.17 & $-$14.98 & \\nodata \t\t\t\t \\\\ 1673236&14 06 03.34&$+$52 09 51.98& 0.4823& $-$1.01& 24.75 & $-$23.38 & \\nodata \t\t\t\t \\\\ 1707796&14 18 37.62&$+$37 46 22.63& 0.1349& $-$1.34& 23.37 & $-$21.87 &{\\tt ABELL 1896}\\\\%, d=1} \\\\ 1714152&14 20 56.84&$+$53 13 07.25& 0.7430& $-$1.05& 25.29 & $-$23.57 & \\nodata \\\\ 1714768&14 21 10.18&$+$42 09 12.97& 0.3529& $-$1.12& 24.03 & $-$20.99 &{\\tt NSCS J142115+420743}\\\\%, d=1.7}\\\\ 1735168&14 28 41.23&$+$43 41 34.03& 0.2136& $-$1.04& 23.70 & $-$22.42 &{\\tt NSCS J142842+434009}\\\\%, d=1.4} \\\\ 1735918&14 28 57.67&$+$54 36 27.65& 0.3819& $-$1.50& 24.39 & $-$21.60 & \t\t\t\t \\\\ 1744373&14 32 04.05&$+$46 37 43.79& 0.0927& $-$1.03& 22.65 & $-$19.72 &{\\tt NSC J143143+463738}\\\\%, d=3.5} \t\t\t\t \\\\ 1755382&14 36 02.53&$+$33 07 53.79& 0.0939& $-$1.02& 22.80 & $-$20.51 & \\nodata \t\t\t\t \\\\ 1756038&14 36 19.44&$+$48 32 10.68& 0.1912& $-$1.07& 24.03 & $-$21.62 & \\nodata \t\t\t\t \\\\ 1759889&14 37 42.41&$+$39 27 45.12& 0.2455& $-$1.36& 24.35 & $-$21.94 &{\\tt MaxBCG J219.42655+39.46313} \\\\ 1768885&14 40 57.03&$+$46 36 46.91& 0.8395& $-$1.12& 25.35 & $-$19.00 & \\nodata \t\t\t\t\\\\ 1775379&14 43 17.07&$+$46 43 48.40& 0.2424& $-$1.34& 23.68 & $-$21.83 & \\nodata \t\t\t\t\\\\ 1793973&14 50 03.51&$+$31 30 15.02& 0.2746& $-$1.10& 24.25 & $-$21.95 & {\\tt MaxBCG J222.55394+31.49750}\\\\% d=2.1} \\\\ 1795282&14 50 31.54&$+$32 53 03.73& 0.1775& $-$1.38& 23.45 & $-$21.67 & \\nodata \t\t\t\t\\\\ 1830754&15 03 23.78&$+$46 06 16.28& 0.4269& $-$1.06& 24.19 & $-$22.74 & \\nodata \t\t\t\t\\\\ 1836180&15 05 23.43&$+$47 06 25.59& 0.2615& $-$1.63& 23.90 & $-$22.52 &{\\tt ABELL 2024} \\\\ 1837171&15 05 46.23&$+$54 54 01.56& 0.2824& $-$1.24& 24.49 & $-$22.69 & \\nodata \t\t\t\t\\\\ 1838173&15 06 08.41&$+$60 02 16.86& 0.5196& $-$1.03& 24.72 & $-$24.11 & \\nodata \t\t\t\t\\\\ 1862173&15 15 05.54&$+$43 09 01.38& 0.0177& $-$1.14& 21.11 & $-$23.69 & {\\tt CGCG 221-045} \t\t\t\t\\\\ 1929427&15 39 50.77&$+$30 43 03.90& 0.0971& $-$1.18& 23.10 & $-$22.36 &{\\tt MaxBCG J234.96158+30.71777} \\\\ 1932815&15 41 05.46&$+$32 04 50.85& 0.0529& $-$1.06& 22.51 & $-$20.79 & \\nodata \t\t\t\t\\\\ 1934628&15 41 46.53&$+$45 56 14.29& 0.2024& $-$1.08& 24.31 & $-$21.37 & \\nodata \t\t\t\t\\\\ 1958800&15 50 51.44&$+$42 02 30.47& 0.0334& $-$1.09& 21.91 & $-$20.61 & \\nodata \\\\%{\\tt WBL 588, d=11.7 } \t\t\t\t\\\\ 1963482&15 52 41.11&$+$37 24 34.16& 0.3710& $-$1.92& 24.23 & $-$23.69 & \\nodata \t\t\t\t\\\\ 1967858&15 54 23.54&$+$48 41 07.36& 0.2271& $-$1.04& 23.93 & $-$22.32 &{\\tt MaxBCG J238.59817+48.68496} \\\\ 2001242&16 07 25.43&$+$47 50 24.15& 0.3282& $-$1.02& 24.76 & $-$22.54 &{\\tt ABELL 2157}\\\\%, d=0.9} \\\\ 2064005&16 33 10.92&$+$36 07 35.15& 0.1648& $-$1.02& 23.93 & $-$21.42 &{\\tt MaxBCG J248.29532+36.12611} \\\\ 2088195&16 43 26.82&$+$39 30 39.92& 0.4119& $-$1.05& 24.72 & $-$22.15 & \\nodata \t\t\\\\ 2122675&16 59 01.01&$+$32 29 38.93& 0.0627& $-$1.26& 23.82 & $-$21.61 &{\\tt ABELL 2241}\\\\%, d=0.9} \\\\ 2134416&17 04 26.40&$+$39 10 12.25& 0.1283& $-$1.01& 23.23 & $-$21.34 &{\\tt NSC J170432+390956 } \\\\ \\hline \\end{tabular} \\end{small} \\end{center} \\label{table} \\end{table*}" }, "1003/1003.0676_arXiv.txt": { "abstract": "We include an energy term based on Dark Matter (DM) self-annihilation during the cooling and subsequent collapse of the metal-free gas, in halos hosting the formation of the first stars in the Universe. We have found that the feedback induced on the chemistry of the cloud {\\it does} modify the properties of the gas throughout the collapse. However, the modifications are not dramatic, and the typical Jeans mass within the halo is conserved throughout the collapse, for all the DM parameters we have considered. This result implies that the presence of Dark Matter annihilations does not substantially modify the Initial Mass Function of the First Stars, with respect to the standard case in which such additional energy term is not taken into account. We have also found that when the rate of energy produced by the DM annihilations and absorbed by the gas equals the chemical cooling (at densities yet far from the actual formation of a proto--stellar core) the structure does {\\it not} halt its collapse, although that proceeds more slowly by a factor smaller than few per cent of the total collapse time. ", "introduction": "In the currently favoured $\\Lambda$CDM cosmological model, the bulk of the matter component is believed to be made of (so far) electromagnetically undetected particles, commonly dubbed Dark Matter (DM). Although the evidence for the existence of DM is compelling on different scales, yet its nature is unknown, and many particle models beyond the standard one have been proposed in the literature as DM candidates. We address the reader to a recent review of observational evidence and particle candidates for DM (e.g. Bertone, Hooper \\& Silk 2005), and will concentrate in this paper on a particular class of candidates, i.e. Weakly Interacting Massive Particles (WIMPs). Many WIMP DM models are stable (under the conservation of the suitable symmetry, for each model) and hence do never decay into standard model particles. However, in many of these very same models the WIMPs are Majorana particles, thus carrying the remarkable property of being self-annihilating; the value of the self--annihilation cross section, arising naturally in many WIMP models, reproduces the dark matter relic abundance required by the $\\Lambda$CDM cosmology, if the mass scale of WIMPs is within the GeV/TeV scale and they are to be thermally produced in the early Universe. We adopt this as a benchmark scenario for our paper, and will often refer to it as a ``Vanilla WIMP''. The actual DM distribution in the local Universe is such that even in the densest regions (e.g. galactic nuclei and black hole surroundings) from which the annihilation signal could be in principle detected, the energy released by WIMP DM annihilations (hereafter, DMAs) is only a negligible fraction of the one associated with standard gas processes. This implies that, locally, DM affects the host system almost uniquely through its gravitational effects, perhaps with the only possible exception of peculiar locations, such as the central parsec of the Milky Way (Fairbairn, Scott \\& Edsjo 2008, Scott, Edsjo \\& Fairbairn 2009, Casanellas \\& Lopes 2009). The effects of annihilating (or decaying, a scenario we do not consider here) DM upon the evolution of the intergalactic medium (IGM) at high redshift have been thoroughly studied (e.g. Chen \\& Kamionkowski 2004; Padmanabhan \\& Finkbeiner 2005; Mapelli \\& Ferrara 2005; Mapelli, Ferrara \\& Pierpaoli 2006; Furlanetto, Oh \\& Pierpaoli 2006; Zhang et al. 2006; Ripamonti, Mapelli \\& Ferrara 2007a; Vald\\'es et al. 2007; Shchekinov \\& Vasiliev 2007), and are now believed to be small, except perhaps in the case of an extremely high clumping factor (Myers \\& Nusser 2008; Chuzhoy 2008; Natarajan \\& Schwarz 2008, Lattanzi \\& Silk 2009), if one takes into consideration a standard, Vanilla WIMP scenario. The effects of DMAs upon primordial star-formation might be more significant. As the IGM could be heated by the energy deposition from DMAs, its temperature might in principle exceed the virial temperature of the smallest halos with the result of quenching gas accretion onto them. This effect has been shown to be unimportant by Ripamonti, Mapelli \\& Ferrara (2007b, hereafter RMF07). However, DMAs are expected to become more important as the collapse proceeds to protostellar scales (Ascasibar 2007). Spolyar, Freese \\& Gondolo (2008; hereafter SFG08) found that during the protostellar collapse of the first (Pop III) stars, the energy released by DMAs and absorbed by the gas could compensate (or even overcome) the radiative cooling of the gas. The increasing importance of such process arises from the combined enhancement during the collapse of DM density (due to gravitational dragging) and gas optical depth, implying a higher annihilation luminosity and absorption by the gas. The final phases of the collapse, after the formation of a hydrostatic core for gas central densities $n_{\\rm c} \\equiv \\rho_{\\rm c}/m_{\\rm p} > 10^{18}\\,\\percmc$ (where $\\rho_{\\rm c}$ is the central baryonic density, and $m_{\\rm p}$ is the proton mass), have been investigated by Iocco et al. (2008, hereafter I08), Freese et al. (2008b, 2009) and Spolyar et al. (2009). Initially, the DM pile-up is purely driven by gravitational interactions, but as the protostar approaches the Zero Age Main Sequence, DM accretion becomes dominated by the capture of WIMPs located in the star host halo after they scatter stellar baryons. As a consequence of the peculiar formation process of Pop III, following the smooth collapse of the gas cloud at the very center of the DM halos hosting them, Iocco (2008) and Freese, Spolyar \\& Aguirre (2008) suggested that DM capture is relevant for primordial stars; however, it can be safely neglected once local star formation is concerned, as the latter takes place anywhere in galactic discs, and it doesn't follow from a single, centered gas collapse episode. Further studies (I08, Yoon, Iocco \\& Akiyama 2008; Taoso et al. 2008) have concluded that WIMP DM capture's most remarkable effect is the possible increase of the stellar lifetime. Quite surprisingly, the early phases of the collapse have received so far less attention with respect to the more advanced ones, i.e. after hydrostatic core formation. For example, it is still unclear if the energy injection following annihilations results in a nett heating or cooling of the gas. In fact, high energy photons and electrons heat the gas through ionizations; however, this heat input could be overwhelmed by the increased production of cooling species (as for example molecular hydrogen) stimulated by the larger abundances of free electrons, thus resulting in a net gas cooling. This, among others, is one of the aspects of the collapse of first stars in presence of WIMP annihilation that we would like to address here. We plan to do so by a set of sophisticated numerical simulations including all the relevant chemical reactions and cooling processes. A first attempt to model the effects of DMA energy input was presented in Ripamonti et al. (2009); this paper represents a substantial extension and improvement of that study. Throughout the paper we assume the following set of cosmological parameters: $\\Omega_{\\rm \\Lambda}=0.76$, $\\Omega_{\\rm m}=0.24$, $\\Omega_{\\rm b}=0.042$, $\\Omega_{\\rm DM} = \\Omega_{\\rm m} - \\Omega_{\\rm b} \\equiv \\Omega_{\\rm WIMP}= 0.198$, and $h=0.73$. ", "conclusions": "We have studied the effects of WIMP Dark Matter Annihilations (DMAs) on the evolution of primordial gas clouds hosting the first stars. We have followed the collapse of gas and DM within a $10^6{\\rm M}_{\\odot}$ halo virializing at redshift $z = 20$, from $z=1000$ to slightly before the formation of a hydrostatic core, properly including gas heating/cooling and chemistry processes induced by DMAs, and exploring the dependency of the results on different parameters (DM particle mass, self-annihilation cross section, gas opacity, feedback strength). Independently of such parameters, when the central baryon density, $n_{\\rm c}$, is lower than the critical density, $n_{\\rm crit} \\approx 10^{9-13}\\,\\percmc$, corresponding to a model-dependent balance between DMA energy input and gas cooling rate, DMA ionizations catalyze an increase in the \\HH abundance by a factor $\\sim 100$. The increased cooling moderately reduces the temperature (by $\\approx 30$ per cent) but does not significantly reduce the fragmentation mass scale. For $n_{\\rm c} \\ge n_{\\rm crit}$, the DMA energy injection exceeds the cooling, with the excess heat mainly going into \\HH dissociations. In the presence of DMA the transition to the continuum dominated cooling regime occurs earlier and generally is not associated with abrupt temperature variations. In conclusion, no significant differences are found with respect to the case without DMAs; in particular, and contrary to previous claims, the collapse does not stall and the cloud keeps contracting even when $n_{\\rm c}\\gg n_{\\rm crit}$. Our simulations stop at central densities $\\approx 10^{14}\\,\\percmc$, and cannot follow the hydrostatic core formation, nor its accretion. At the final simulation stage, the lower temperature/infall velocity of the layers enclosing a mass of $\\approx 10^2 M_{\\odot}$ suggest that DMAs might lead to slightly longer stellar formation timescales, with a possible $\\approx 20$ per cent increase over models without DMAs. The latter finding strengthens our conclusions, although the final answer will come from numerical simulations (hopefully also three-dimensional) able to address this very same problem in the yet unexplored density regime $10^{14}\\,\\percmc \\lesssim n_{\\rm c} \\lesssim 10^{18}\\,\\percmc$." }, "1003/1003.0440_arXiv.txt": { "abstract": "Chemically Peculiar (CP) stars have been subject of systematic research since more than $50$ years. With the discovery of pulsation of some of the cool CP stars, the availability of advanced spectropolarimetric instrumentation and high signal-to-noise, high resolution spectroscopy, a new era of CP star research emerged about $20$ years ago. Together with the success in ground-based observations, new space projects are developed that will greatly benefit for future investigations of these unique objects. In this contribution we will give an overview of some interesting results obtained recently from ground-based observations and discuss on future outstanding Gaia space mission and its impact on CP star research. ", "introduction": "Back in $1897$, more than one century ago, the first peculiar stars were found in the course of the Henry Draper Memorial classification work at Harvard by Antonia Maury and Annie Cannon. Maury used the designation ``peculiar'' for the first time to describe spectral features in the remarks to the spectrum of $\\alpha^2$~CVn \\citep{pickering1897}, making a first attempt for two-dimensional classification system considering the strength and the width of the spectral lines. In $1974$ Preston proposed the division of main-sequence CP stars into four groups according to their spectroscopic characteristics \\citep{preston}: CP1 (Am/Fm stars), CP2 (Si, SrCrEu stars), CP3 (HgMn stars), CP4 (He-weak stars). More detailed spectroscopic consideretion of CP stars required to introduce new subtypes of CP stars, such as He-rich and $\\lambda$~Boo stars. CP2 stars, including Bp/Ap, host strong surface magnetic fields \\citep{lantz} that are likely stable on large time-intervals \\citep{north}. Abundance peculiarities were measured using the curve-of-growth method based on simple assumptions about formation of absorption lines (models of Schuster-Schwarzschild, Uns\\\"old, Milne-Eddington). Since that time and with development of new high-resolution, high signal-to-noise CCD based spectrometers, big progress has been made in abundance analysis of CP stars, revealing the presence of vertical \\citep{str1,str2,str3,str4} and horizontal \\citep{di0,di1,di2} elements separation in their atmospheres caused by the processes of microscopic particle diffusion \\citep{michaud}. The discovery of strong stellar surface magnetic fields \\citep{babcock} opened a new research field in astrophysics~--~stellar magnetism. A $200$~Gauss accuracy of the magnetic field detection usually obtained with photographic plates has increased to $\\approx1$~Gauss with modern spectropolarimetry and new techniques (such as Least Square Deconvolution, or LSD, for example) \\citep{lsd,wade}. With the discovery by D.~Kurtz \\citep{kurtz1978} of a $12$~min pulsation period in HD\\,101065 a subgroup of the cool CP stars, the so-called rapidly oscillating Ap (roAp) stars, became extremely promising targets for asteroseismology, a most powerful tool for testing theories of stellar structure. Driving of the oscillations results from a subtle energy balance depending directly on the interaction between the magnetic field, convection, pulsations, and atomic diffusion. Amazing insights in the 3-D structure of stellar atmospheres became available \\citep[see, for example,][]{puls1,puls2}. ", "conclusions": "" }, "1003/1003.2735_arXiv.txt": { "abstract": "A quasi-Keplerian parameterisation for the solutions of second post-Newtonian (\\pN) accurate equations of motion for spinning compact binaries is obtained including leading order spin-spin and next-to-leading order spin-orbit interactions. Rotational deformation of the compact objects is incorporated. For arbitrary mass ratios the spin orientations are taken to be parallel or anti-parallel to the orbital angular momentum vector. The emitted gravitational wave forms are given in analytic form up to 2\\pN\\ point particle, 1.5\\pN\\ spin-orbit and 1\\pN\\ spin-spin contributions, whereby the spins are assumed to be of 0PN order. ", "introduction": " ", "conclusions": "" }, "1003/1003.5323_arXiv.txt": { "abstract": "We have carried out an $L'$ and $M$ band Adaptive Optics (AO) extrasolar planet imaging survey of 54 nearby, sunlike stars using the Clio camera at the MMT. Our survey concentrates more strongly than all others to date on very nearby F, G, and K stars, in that we have prioritized proximity higher than youth. Our survey is also the first to include extensive observations in the $M$ band, which supplemented the primary $L'$ observations. These longer wavelength bands are most useful for very nearby systems in which low temperature planets with red IR colors (i.e. $H - L'$, $H - M$) could be detected. The survey detected no planets, but set interesting limits on planets and brown dwarfs in the star systems we investigated. We have interpreted our null result by means of extensive Monte Carlo simulations, and constrained the distributions of extrasolar planets in mass $M$ and semimajor axis $a$. If planets are distributed according to a power law with $dN \\propto M^{\\alpha} a^{\\beta} dM da$, normalized to be consistent with radial velocity statistics, we find that a distribution with $\\alpha = -1.1$ and $\\beta = -0.46$, truncated at 110 AU, is ruled out at the 90\\% confidence level. These particular values of $\\alpha$ and $\\beta$ are significant because they represent the most planet-rich case consistent with current statistics from radial velocity observations. With 90\\% confidence no more than 8.1\\% of stars like those in our survey have systems with three widely spaced, massive planets like the A-star HR 8799. Our observations show that giant planets in long-period orbits around sun-like stars are rare, confirming the results of shorter-wavelength surveys, and increasing the robustness of the conclusion. ", "introduction": "Nearly 400 extrasolar planets have now been discovered using the radial velocity (RV) method. RV surveys currently have good statistical completeness only for planets with periods of less than ten years \\citep{cumming,butlercat,carnp}, due to the limited temporal baseline of the observations, and the need to observe for a complete orbital period to confirm the properties of a planet with confidence. The masses of discovered planets range from just a few Earth masses \\citep{hotNep} up to around 20 Jupiter masses (\\mjup). We note that a 20 \\mjup~object would be considered by many to be a brown dwarf rather than a planet, but that there is no broad consensus on how to define the upper mass limit for planets. For a good overview of RV planets to date, see \\citet{butlercat} or \\url{http://exoplanet.eu/catalog-RV.php}. The large number of RV planets makes it possible to examine the statistics of extrasolar planet populations. Several groups have fit approximate power law distributions in mass and semimajor axis to the set of known extrasolar planets (see for example \\citet{cumming}). Necessarily, however, these power laws are not subject to observational constraints at orbital periods longer than 10 years -- and it is at these orbital periods that we find giant planets in our own solar system. We cannot obtain a good understanding of planets in general without information on long period extrasolar planets. Nor can we see how our own solar system fits into the big picture of planet formation in the galaxy without a good census of planets in Jupiter- and Saturn-like orbits around other stars. Repeatable detections of extrasolar planets (as opposed to one-time microlensing detections) have so far been made by transit detection (e.g. \\citet{hd209458}), by RV variations \\citep{51peg}, by astrometric wobble \\citep{benedict}, or by direct imaging \\citep{hr8799}. Of these methods, transits are efficient only for detecting close-in planets. As noted above, precision RV observations have not been going on long enough to detect more than a few planets with periods longer than ten years, but even as RV temporal baselines increase, long period planets will remain harder to detect due to their slow orbital velocities. The amplitude of a star's astrometric wobble increases with the radius of its planet's orbit, but decades-long observing programs are still needed to find long-period planets. Direct imaging is the only method that allows us to characterize long-period extrasolar planets on a timescale of months rather than years or decades. Direct imaging of extrasolar planets is technologically possible at present only in the infrared, based on the planets' own thermal luminosity, not on reflected starlight. The enabling technology is adaptive optics (AO), which allows 6-10m ground-based telescopes to obtain diffraction limited IR images several times sharper than those from HST, despite Earth's turbulent atmosphere. Theoretical models of giant planets indicate that such telescopes should be capable of detecting self-luminous giant planets in large orbits around young, nearby stars. The stars should be young because the glow of giant planets comes from gravitational potential energy converted to heat in their formation and subsequent contraction: lacking any internal fusion, they cool and become fainter as they age. Several groups have published the results of AO imaging surveys for extrasolar planets around F, G, K, or M stars in the last five years (see for example \\citet{masciadri,kasper,biller1,GDPS}; and \\citet{chauvin}). Of these, most have used wavelengths in the 1.5-2.2 $\\mu$m range, corresponding to the astronomical $H$ and $K_S$ filters \\citep{masciadri,biller1,GDPS,chauvin}. They have targeted mainly very young stars. Because young stars are rare, the median distance to stars in each of these surveys has been more than 20 pc. In contrast to those above, our survey concentrates on very nearby F, G, and K stars, with proximity prioritized more than youth in the sample selection. The median distance to our survey targets is only 11.2 pc. Ours is also the first survey to include extensive observations in the $M$ band, and only the second to search solar-type stars in the $L'$ band (the first was \\citet{kasper}). The distinctive focus on older, very nearby stars for a survey using longer wavelengths is natural: longer wavelengths are optimal for lower temperature planets which are most likely to be found in older systems, but which would be undetectable around all but the nearest stars. More information on our sample selection, observations, and data analysis can be found in our Observations paper, \\citet{obspaper}, which also details our careful evaluation of our survey's sensitivity, including extensive tests in which fake planets were randomly placed in the raw data and then recovered by an experimenter who knew neither their positions nor their number. Such tests are essential for establishing the true relationship between source significance (i.e. 5$\\sigma$, 10$\\sigma$, etc.) and survey completeness. Our survey places constraints on a more mature population of planets than those that have focused on very young stars, and confirms that a paucity of giant planets at large separations from sun-like stars is robustly observed at a wide range of wavelengths. In Section \\ref{sec:rv}, we review power law fits to the distribution of known RV planets, including the normalization of the power laws. In Section \\ref{sec:tmod}, we present the constraints our survey places on the distribution of extrasolar giant planets, based on extensive Monte Carlo simulations. In Section \\ref{sec:long} we discuss the promising future of planet-search observations in the $L'$ and especially the $M$ band, and in Section \\ref{sec:concl} we conclude. ", "conclusions": "\\label{sec:concl} We have surveyed unusually nearby, mature star systems for extrasolar planets in the $L'$ and $M$ bands using the Clio camera with the MMT AO system. By extensive use of blind sensitivity tests involving fake planets inserted into our raw data (reported in detail in \\citet{obspaper}), we established a definitive significance vs. completeness relation for planets in our data, which we then used in Monte Carlo simulations to constrain planet distributions. We set interesting limits on the masses of planets and brown dwarfs in the star systems we surveyed, but we did not detect any planets. Based on this null result, we place constraints on the power laws that may describe the distribution of extrasolar planets in mass and semimajor axis. We also place constraints on planet abundances independent of the distributions. If the distribution of planets is a power law with $dN \\propto M^{\\alpha} a^{\\beta} dM da$, the work of \\citet{cumming} and \\citet{butlercat} indicates that the most optimistic (i.e. planet-rich) case permitted by the statistics of known RV planets correponds to about $\\alpha = -1.1$ and $\\beta = -0.46$. Normalizing the distribution to be consistent with RV statistics, we find that these values of $\\alpha$ and $\\beta$ are ruled out at the 90\\% confidence level, unless the semimajor axis distribution is truncated at a radius $R_{trunc}$ less than 110 AU. Though $\\beta=0.0$ is not physically plausible, previous work has sometimes used it an example: for $\\alpha=-1.31$, corresponding to the best-fit value from \\citet{cumming}, we rule out $\\beta=0.0$ unless $R_{trunc}$ is less than 38 AU. Independent of distribution models, with 90\\% confidence no more than 50\\% of stars like those in our survey have a 5 \\mjup~or more massive planet orbiting between 30 and 94 AU, no more than 15\\% have a 10 \\mjup~planet orbiting between 22 and 100 AU, and no more than 25\\% have a 20 \\mjup~object orbiting between 8 and 100 AU. Our constraints on planet abundances are similar to those placed by \\citet{kasper} and \\citet{biller1}, but less tight than those of \\citet{nielsen} and especially \\citet{GDPS}, The recent work of \\citet{nielsenclose} and \\citet{chauvin} also placed tighter constraints on exoplanet distributions than our survey. However, we have surveyed a more nearby, older set of stars than any previous survey, and have therefore placed constraints on a more mature population of planets. Also, we have confirmed that a paucity of giant planets at large separations from sun-like stars is robustly observed at a wide range of wavelengths. The best current $H$ regime observations, those of \\citet{GDPS}, would attain sensitivity to lower mass planets than did our $L'$ and $M$ band observations for all of our survey targets except those lying within 4 pc of the Sun. However, as larger telescopes are built and longer exposures are attempted, the sensitivity of $M$ band observations may be expected to increase at least as fast as that of $H$ band observations (in part because $M$ band detectors are currently a less mature technology). As shown in Figures \\ref{fig:HLM1} and \\ref{fig:HLM2}, a modest increase from current sensitivity levels, even if paralleled by an equal increase in $H$ band sensitivity, would render the $M$ band the wavelength of choice for extrasolar planet searches around a large number of nearby stars." }, "1003/1003.5924_arXiv.txt": { "abstract": "Light curves in the $B$, $V$, and $I_c$ passbands have been obtained for the type II Cepheids V154 in M3 and V42 and V84 in M5. Alternating cycle behavior, similar to that seen among RV Tauri variables, is confirmed for V84. Old and new observations, spanning more than a century, show that V154 has increased in period while V42 has decreased in period. V84, on the other hand, has shown large, erratic changes in period that do not appear to reflect the long term evolution of V84 through the HR diagram. ", "introduction": "Type II Cepheids were among the first variable stars discovered within the globular clusters M3 (NGC 5272) and M5 (NGC 5904) \\citep{p89, p90}. Their initial discovery came long before the realization that Cepheids were pulsating stars \\citep{sh14}, and even longer before Baade's \\citep{b56} discovery that Cepheids could be divided into two population groups. Until recently, the vast majority of observations available for globular cluster Cepheids were obtained photographically. This paper presents new $B$, $V$, and Cousins $I_c$ band CCD light curves of the type II Cepheids V154 in M3 and V42 and V84 in M5, as well as some new photographic data. These variable stars were first observed more than a century ago, raising the possibility that their long term period changes may indicate the direction and speed of their evolution through the instability strip. Our new light curves, together with earlier observations and, in the case of V42, observations from the All Sky Automated Survey (ASAS) \\citep{po02}, are used to rediscuss the long term period changes of these variables. ", "conclusions": "Theory predicts that long period type II Cepheids enter the instability strip either while undergoing blueward instability loops from the asymptotic red giant branch as a consequence of helium shell flashes, or during final blueward evolution as the hydrogen burning shell nears the surface of the star \\citep{sch70, m73, g85, cl88, bo97}. \\citet{bo97} found that, for the lower mass but brighter Cepheids, the instability strip could be crossed two or three times as a consequence of thermal pulses. The period of a pulsating star is linked to its density via Ritter's pulsation equation, $Q = P \\sqrt{\\rho}$, where Q is the pulsation constant, P is the period, and $\\rho$ is the mean stellar density. The pulsation period of a Cepheid is often its most accurately known property, and, as noted long ago by \\citet{ed18}, a small change in the structure of the Cepheid will reveal itself as a change in pulsation period before it can be recognized in any other measured quantity. Each of our three stars showed long term period change behavior, but the changes are different in each case. V154 showed a modest increase in period consistent with movement to the red in the instability strip. V42 showed a decrease in period, consistent with movement to the blue. If these period changes indicate the long term evolution of these stars, V154 could be interpreted as being on the redward evolving, and V42 on the blueward evolving portion of the instability loops predicted by theory during shell helium burning. Alternatively, blueward moving V42 might be in the final blueward evolutionary phase. In neither case, however, is a parabolic fit to the phase diagram, implying a constant rate of period change, significantly better than the assumption of abrupt period changes. Using the theoretical timescales of \\citet{g76}, \\citet{cl88} found that one might expect a rate of period decrease of $P^{-1}dP/dt = -0.0005$ to $-0.002$ cycles per 100 years during the final quiescent blueward evolving stage. The observed decrease in the period of V42 is about $\\Delta P/P = -0.0007$ over 120 years, consistent with that expectation. Although V42 and V84 are both near the luminosity dividing type II Cepheid and RV Tauri behavior in the HR diagram \\citep{wc84,bo97}, V84 showed period changes much more erratic than those of V42. V84 is not, however, the only type II Cepheid exhibiting period fluctuations. \\citet{cl88} found that period fluctuations were not unusual in type II Cepheids in globular clusters, although rarely do they seem to reach the extent exhibited by V84. Apparently random period fluctuations have also been observed in the O-C diagrams of other variable stars \\citep{be09, tu09}. V1, a 15.5~day period Cepheid in the globular cluster M12, does perhaps show jumps in the phase shift diagram on a scale similar to that of V84 \\citep{cl88}. V84 shows evidence of RV Tauri behavior, and strong cycle-to-cycle period fluctuations have been observed for RV Tauri stars in the field, e.g. \\citet{p05}. However, \\citet{cl88} do not report RV Tauri type behavior for V1 in M12, so that these irregular periods appear not to be limited only to long period variables near the RV Tauri domain in the HR diagram." }, "1003/1003.2818_arXiv.txt": { "abstract": "The measurements of the Doppler shifts of the Fraunhofer lines, scattered by the dust grains in the solar F-corona, provides the insight on the velocity field of the dust and hence on its origin. We report on such measurements obtained during the total eclipse of March 29, 2006. We used a Fabry-P\\'{e}rot interferometer with the FOV of $5.9\\degr$ and the spectral resolution of about 5000 to record Fraunhofer spectral lines scattered by the dust of the F-Corona. The spectral region was centered on the \\mbox{Mg\\,{\\sc i}} $5172.69 \\rmn{\\AA}$ line. The measured line-on-sight velocities with the amplitude in the range from -10 to 10~km$\\cdot \\rmn{s^{-1}}$ show that during our observations the dust grains were on the orbit with a retrograde motion in a plane nearly perpendicular to the ecliptics. This indicates their cometary origin. Indeed, at the end of March, 2006, SOHO recorded several sungrazing comets with the orbital elements close to what was deduced from our measurements. We conclude that the contribution of comets to the dust content in the region close to the Sun can be more important albeit variable in time. We also deduce that the size of the most of the dust grains during our observations was less than 0.1 $\\umu$m. ", "introduction": "The work of \\citet{hulst47} established, after an earlier suggestion by Grotrian \\citep[as cited by][]{hulst47}, that the emission of the solar outer corona, called F-corona after the Fraunhofer lines composing its spectrum, is due to the scattering of the solar light by the dust particles. In recent decennia it became more clear that the dust component near the Sun cannot be described as a mere extrapolation of the zodiacal light disc with the dust grains orbiting in the plane of the ecliptics and slowly drifting to the Sun under the Poynting-Robertson effect as one could accept in a first approximation. The measurements of the Doppler shifts of the Fraunhofer lines in the spectrum of the F-corona, first done during the total eclipse of July 31, 1981 by \\citet{scheglov87}, and the analysis of the derived velocities by \\citet{shest87}, indicated that, although most of the dust particles are orbiting in the ecliptic plane, there are also grains orbiting at high values of inclination angle $i$, implying that the region close to the Sun receives a dust contribution from long-period comets. Also, the \\textit{in situ} measurements of the dust particle velocity distribution on board of HELIOS-1 showed the existence of two distinct components in the interplanetary dust, cometary and asteroidal \\citep{grun80}. Among more recent results, the high frequency of the Sun-grazing comets recorded by SOHO \\citep{macqueen91} makes also to consider the cometary contribution as important. Furthermore, studies of extrasolar planetary and dust systems showed the existence of Falling Evaporating Bodies, possibly comets \\citep[e.g.][]{beust98}. This observational progress was followed by a thorough theoretical modelling taking into account complex composition of the dust in the solar vicinity (see \\citealp{mann00}, \\citealt{koba08}). It would have been important to obtain further measurements of the line-on-sight (thereafter LOS) velocities of the dust near the Sun to have a better idea of the relative contributions of the cometary and the ecliptic disc dust components and of their possible variation in time. After the first work of 1981, these measurements, to the best of our knowledge, were successfully attempted only once, during the eclipse of July 11, 1991 \\citep{aimanov95} confirming the first results. Here we present new higher quality measurements of the LOS velocity of the dust in the F-corona obtained during the eclipse on March 29, 2006 \\citep[see][for a preliminary report]{shest07}, and discuss their implications. \\begin{figure*} \\begin{center} \\includegraphics [width=18cm] {LayFig1.eps} \\end{center} \\caption{The Fabry-P\\'{e}rot interferograms of the daylight sky and of the circumsolar region up to elongation $\\epsilon \\simeq 11$. Here, both axes are in pixels, the platescale is 20.66 arcsec/pix, the FOV\\,$=5.9\\degr$. \\textit{From left to right:} a comparison daylight sky frame ``D'', the total eclipse frame ``E'', and the frame ``L'' taken at the end of the totality, but already including the green coronal emission line of [Fe XIV] at $\\lambda 5302.86 \\rmn{\\AA}$. The fiducial wire (black line) indicates the East-West direction, the West is to the left along the wire and the North is down. The frames are rotated so that the horizontal axis of the frames coincides with the ecliptic plane. The frames are dark current subtracted. For more details see the text.} \\label{fig:fig1} \\end{figure*} \\section[]{Observations} The site of the observations was the village of Mugalzhar in the Aktobe region of the Republic of Kazakhstan, at $\\phi = 48\\degr 35\\arcmin$ and $\\lambda = 58\\degr 27\\arcmin$, situated in the middle of the totality band. According to our calculations, the beginning of the total eclipse was at UT 11h 32m 40, its end at 11h 35m 30s, and the duration of the total phase was 170\\,s. The Sun was at $27\\fdg 5$ above the horizon. The weather was mostly rainy and windy during the days preceding the eclipse, but on the eclipse day, the sky was clear with no wind and excellent transparency. We used the same optical set-up as in our previous measurements in 1991 \\citep{scheglov87}, i.e. a Fabry-P\\'{e}rot (hereafter FP) spectrometer with a coronographic mask rejecting the light of the solar corona and thus reducing the background. The entrance lens of 10 cm diameter has its focal plane at the field lense, the latter is followed by a collimator. The FP etalon in series with an order separation filter is placed in the parallel beam close the exit pupil imaged by the field lense. It is followed by a photographic objective and a CCD detector. The latter was an Apogee Alta-10 CCD device with $2048^{2}$ pixels of 14 $\\umu$m size. During the data reduction, the frames were binned by $2\\times 2$ pixels, so that to avoid any confusion we will refer from now on to the detector format of $1024^{2}$ pixels with the 28 $\\umu$m pixel size. The solar angular radius at the time of the eclipse was $\\mathcal{R}_{\\sun} = 961$ arcsec, the measured on the CCD value was 46.5\\,pix. This gives the equivalent focal length of 279.4 mm, and the platescale of 20.66 arcsec/pixel. The field of view (FOV) is $5.9\\degr$, corresponding to the elongations $\\epsilon < 11$, however, to reduce the background light, the central region around the Sun, $\\epsilon < 2.6$, is hidden by the coronographic mask (hereafter, the elongation $\\epsilon$ is given in angular solar radii $\\mathcal{R}_{\\sun}$). \\begin{figure*} \\begin{center} \\includegraphics [width=14 cm] {LayFig2.eps} \\end{center} \\caption{The interferogram of the circumsolar region during the total eclipse in different states of data reduction procedure. \\textit{Left:} after median filtering and dividing on the flat-field as described in the text. \\textit{Right:} The same with the low frequency two-dimensional trend subtracted. The elongation $\\epsilon$ is in angular solar radii $\\mathcal{R}_{\\sun}$.} \\label{fig:fig2} \\end{figure*} \\subsection[]{Spectro-imaging} We used a mechanically adjustable FP etalon with the free space $\\Delta=70\\,\\umu$m. The order separation filter has the FWHM = $10\\,\\rmn{\\rmn{\\AA}}$ with the central wavelength close to that of the \\mbox{Mg\\,{\\sc i}} line at 5172.69 $\\rmn{\\AA}$. In the wavelength space, the set-up is such that a frame includes 4 FP orders. The highest, on-the-axis order $N=\\Delta/\\lambda$, is at the centre of the frame and is unseen due to the use of the coronographic mask; the orders lower than $N-4$ are outside of the FOV. The dark subtracted frames obtained during the eclipse are shown in Fig.\\,\\ref{fig:fig1}. The left frame, denoted as ``D'', is one of the calibration interferograms of the daylight sky, scattered on a white screen, recorded shortly before and after the total phase of the eclipse. The use of the white screen allows a homogenous brightness distribution over the field of view without changing the spectrometer position which remains at the same position as during the total phase. The frame shows concentric rings of the Fraunhofer absorption lines scattered in the terrestrial atmosphere. The dark circle at the centre of the frame corresponds to the coronographic mask. The dark line corresponds to a fiducial wire indicating the East-West axis, the West to the left along the wire and the North down. For convenience of interpretation, the frames are rotated so that the frames horizontal axis coincides with the ecliptic plane. The frames are dark current subtracted. The frame ``E''in the middle of the Fig.\\,\\ref{fig:fig1} is our main scientific exposure, of 130\\,s, taken on the circumsolar region during the eclipse. It was started a few seconds after the beginning of the totality and stopped well before its end, so that contamination by coronal or chromospheric lines was avoided. The concentric dark rings are the solar Fraunhofer lines scattered on the dust particles of the F-corona. Finally, the frame ``L''at the right of the Fig.\\,\\ref{fig:fig1}, with the exposure time of 20\\,s was started close to the end of the total phase and lasted a few seconds beyond, so that it caught the light of the solar corona which was just appearing from behind the mask (but not yet the photosphere). On this frame, one can see, additionally to the absorption spectral features of the F-corona, the bright emission rings of the green coronal [Fe XIV] line at $\\lambda 5302.86 \\rmn{\\AA}$. Its wavelength lies far from the centre of the our narrow-band filter, however the emission is so strong that its light passed in the filter transmission wings. The corresponding interference rings well visible in Fig.\\,\\ref{fig:fig1} traced the location of the 3 used Fabry-P\\'{e}rot orders and gave useful reference points for the data reduction. Its measured $\\rmn{FWHM} = 1.2\\pm 0.1\\rmn{\\AA}$ gives the spectral resolution of the instrument. The daylight sky interferograms ``D'' provided the wavelength standard and allowed to eventually measure the Doppler shift of the dust particles of the F-corona, while the interferogram of the coronal green line $\\lambda 5302.86\\rmn{\\AA}$ allowed to calibrate the spectral geometry of the frames, and in particular to accurately define the rings centre as it will be discussed in Section\\,\\ref{Section:Reduction}. \\begin{figure*} \\begin{center} \\includegraphics [width=12 cm] {LaySpe.eps} \\end{center} \\caption{Extracted spectrograms: The flux integrated over the position angle $PA$ in a function of the radius r. \\textit{Upper plot:} The daylight sky, frame ``D\"; \\textit{Middle plot:} The circumsolar region during the total eclipse, frame ``E\"; \\textit{Lower plot:} Same but also including the strong coronal emission [$\\mbox{Fe\\,{\\sc xiv}}]\\,\\lambda 5302.86 \\rmn{\\AA}$, frame ``L\". The relevant Fraunhofer lines and their wavelengths are indicated. Note that there are 3 overlapping spectral orders, they are traced by the [$\\mbox{Fe\\,{\\sc xiv}}$] emission pics.} \\label{fig:fig3} \\end{figure*} \\subsection[]{Brightness calibration} The weather conditions being excellent and stable, the Fabry-P\\'{e}rot data were calibrated in brightness relative to that of the solar disc, $B_{\\sun}$. Two 0.1\\,s exposures on the Sun through a combination of a neutral and green filters were taken to the East, at $\\epsilon = 4.0$, and to the West, at $\\epsilon = 4.4$ from the Sun position at the total phase of the eclipse. The Sun elevation for the calibration and for the eclipse frames was nearly the same, about $30\\degr$. The uncovered surface of the solar disc was 0.94 and 0.25 respectively for the two exposures. Taking this into account, the daylight sky brightness off the eclipse is estimated to has been $4\\cdot 10^{-5}B_{\\sun}$. The brightness of the F-corona continuum emission, $B_{F}$, was measured to be $4.1\\cdot 10^{-9}B_{\\sun}$ at the elongation $\\epsilon = 4.0$, it was $3.9\\cdot 10^{-9}B_{\\sun}$ at $\\epsilon = 4.4$, and $1.7\\cdot 10^{-9}B_{\\sun}$ for the $\\epsilon$ range from 9 to 10. The brightness decrement in the F-corona is in agreement with the results of \\citet{koutchmy78}. Compared to the daylight sky, the F-corona is fainter by a factor of $10^{4}$. \\section[]{Data reduction} \\label{Section:Reduction} The extraction of the Doppler shift of the solar absorption lines out of the Fabry-P\\'{e}rot interferograms needs a thorough data reduction process. It is worthwhile to be describe here in details. \\subsection[]{Noise and trends handling } After the bias and dark current subtraction, all the frames were $2\\times 2$ binned, which increased the signal-to-noise ratio without changing spectral or spatial resolution. The next step was defining all regions meaningless for further reduction, which includes the central dark region corresponding to the coronographic mask, that of the fiducial wire, the part lying out of the field-of-view of the detector and the regions of strong light induced by spurious reflections. The resulting numerically masked area was about 30\\%. The ``hot\", ``cold\" and ``dead\" pixels were cleaned out using a median filter. The visual analysis showed that the ``clouds\" of defective pixels did not exceed a region of 10-12 pixels a size. We used therefore a non-linear circular filter covering 31 neighbour pixels (i.e. 2n+1, where n is the maximum scale of ``clouds\"). The correction was applied only to the pixels with the count exceeding 20\\% of the median value. The number of corrected pixels was less than 4\\% on the eclipse frames and less than 0.5\\% on the calibration frames. For convenience of interpretation, all frames were then rotated so that the horizontal axis coincides with the ecliptic plane. The next step was the flat-fielding. Usually, it is done by a straightforward division by a averaged flat-field frame with a subsequent masking of the pixels resulting from the division by zero or close to zero values. We decided to optimize this operation by adding a small constant to the flat-field frame. This is similar in its spirit to the Tikhonov regularisation of ill-defined inverse problems. The value of the constant was defined by analysing the histogram of the counts of the flat-field frame. The aim of this operation is to keep moderate, or negligible, the change induced in the signal-to-noise ratio. Thus, the division on the flat-field does not change the structure of the frames and avoids a useless addition of a noise. Further, we measured the two-dimensional low frequency trend of the resulting frames in order to subtract it and to keep only useful spectro-spatial variations. It was done by using iteratively a linear circular moving average filter. The best filter diameter and the number of iterations were searched by trials in such a way that the frame with the subtracted 2D trend would still keep the structures on a 10-15 pixels scale, which is that of the recorded Fraunhofer lines. The best filter had the diameter of 13 pixels, covering 137 closest pixels, and the best number of iterations was 4. Higher the number of iterations, closer the iterative filter to a linear gaussian smoothing filter, $exp[-(x{^2} + y{^2}/(2 \\sigma_{g}^{2})]$, where $\\sigma_{g}$ defines the degree of the smoothing. The advantage of using an iterative filter is the possibility to control the achieved smoothness by limiting the number of iterations. We tried square and rectangular moving average filters, however they give rise to spur line-like features. We also applied the methods of high and low enveloping curves, but it did not give better results. The FP interferogram after median filtering and subtraction of the 2D low frequency trend is shown in Fig.\\,\\ref{fig:fig2}. Finally, the data were passed through the gaussian filter with $\\sigma_{g} = 1.3$ pixels, or $FWHM\\approx 3$ pixels which is close to the measured $FWHM$ of the point-spread function of the experiment. \\begin{table*} \\centering \\caption{Spectral lines identification.} \\begin{tabular}{| c | c |c |c |} \\hline Radius & Elongation $\\epsilon$ & Wavelength & Element \\\\ (pixels) & ($\\mathcal{R}_{\\sun}$) & $\\lambda (\\rmn{\\AA}$) & \\\\ \\hline 182 & 3.92 & 5168.91+5169.04 & \\mbox{Fe\\,{\\sc i}} + \\mbox{Fe\\,{\\sc ii}} \\\\ 196 & 4.22 & 5167.33 & \\mbox{Mg\\,{\\sc i}} \\\\ 220 & 4.72 & 5183.62 & \\mbox{Mg\\,{\\sc i}} \\\\ 238 & 5.12 & 5162.28 & \\mbox{Fe\\,{\\sc i}} \\\\ 266 & 5.72 & 5159.06 & \\mbox{Fe\\,{\\sc i}} \\\\ 298 & 6.40 & 5172.69 & \\mbox{Mg\\,{\\sc i}} \\\\ 328 & 7.06 & 5167.33 & \\mbox{Mg\\,{\\sc i}} \\\\ 344 & 7.40 & 5183.62 & \\mbox{Mg\\,{\\sc i}} \\\\ 358 & 7.68 & 5162.28 & \\mbox{Fe\\,{\\sc i}} \\\\ 374 & 8.04 & 5159.06 & \\mbox{Fe\\,{\\sc i}} \\\\ 402 & 8.64 & 5172.69 & \\mbox{Mg\\,{\\sc i}} \\\\ 425 & 9.14 & 5167.33 & \\mbox{Mg\\,{\\sc i}} \\\\ 448 & 9.64 & 5162.28 & \\mbox{Fe\\,{\\sc i}} \\\\ 460 & 9.88 & 5159.06 & \\mbox{Fe\\,{\\sc i}} \\\\ 485 & 10.42 & 5172.69 & \\mbox{Mg\\,{\\sc i}} \\\\ \\hline \\end{tabular} \\label{tab:tab_lines} \\end{table*} \\begin{table*} \\centering \\caption{Areas used for the Doppler shift measurements.} \\begin{tabular}{| c | c |c |} \\hline Elongation range & Mean $\\epsilon$ & Lines used \\\\ \\hline 3.12 - 4.09 & 3.66 & 5172.7, 5169.0 \\\\ 3.12 - 4.82 & 4.09 & 5169.0, 5167.3 \\\\ 4.09 - 4.82 & 4.52 & 5167.3, 5183.6 \\\\ 4.82 - 5.63 & 5.27 & 5183.6, 5167.3 \\\\ 5.63 - 6.37 & 6.02 & 5159.1, 5172.7 \\\\ 6.36 - 7.01 & 6.71 & 5172.7, 5167.3 \\\\ 7.01 - 7.74 & 7.40 & 5167.3, 5183.6, 5162.3 \\\\ 7.74 - 8.43 & 8.08 & 5162.3, 5159.1 \\\\ 8.43 - 9.12 & 8.77 & 5172.7, 5167.3 \\\\ 9.12 - 9.63 & 9.38 & 5167.3, 5162.3 \\\\ 9.63 - 10.25 & 9.98 & 5162.3, 5159.1, 5172.7 \\\\ \\hline \\end{tabular} \\label{tab:tab_range} \\end{table*} \\subsection[]{Defining the interferograms centre} In an ideal Fabry-P\\'{e}rot interferogram, the wavelength $\\lambda$ is constant on a circular ring, and varies with its radius as $r{^2}$. Let $X_{c}$ and $Y_{c}$ denote the coordinates of the centre of the interference rings. Their values depend on many optical parameters, which can vary with temperature and mechanical flexures, so that they must be carefully defined for each frame. After different trials, we found that the most reliable way to measure the centre position was to use a correlation method in the following way. Let us denote $z=r^{2}$. We adopt a first guess of the centre coordinates $X_{0}$ and $Y_{0}$, and divide the frame on two sub-frames, left and right, symmetrically with the respect to $X_{0}$. For each sub-frame, we compute counts \\textit{vs} $z$, which gives us two functions $L(z)$ and $R(z)$ respectively for the left and the right sub-frames. We compute then the correlation of $L$ and $R$ , and vary $X_{0}$. The maximum of the correlation gives the value of $X_{0}$, which is adopted as $X_{c}$. In a similar way, we define the best value of $Y_{c}$, correlating the upper and the lower sub-frames. Such a correlation measurement can give false and biased results if, for example, there is a strong asymmetry in the intensity of interference patterns. To have an additional check, we applied another method to smaller parts of the patterns, using only arcs of the rings. For a given ring, we take first guess values of $X_{c}$, $Y_{c}$, $r$, and the width of the ring $dr$. Varying the values of $X_{c}$, $Y_{c}$ and $r$, we find the values such that the sum of counts is the less (for absorption lines). To insure a good statistics, the value of $dr$ must be sufficiently large, but without covering neighbor ring patterns. We used values of $dr$ in the range from 2 to 12 pixels. The difference of the centre coordinates defined by this method and that of correlations did not exceed 0.3 pix, which is at the level of the expected uncertainty. The resulting values of the center coordinates are defined with an accuracy better than 0.3 pix. \\subsection[]{Reduced spectrograms} Once the centre of the interference rings was found, it is convenient to transform the data presentation from cartesian to polar coordinates $(r, \\phi)$. The frames were oriented so that the polar angle $\\phi$ and the position angle on the sky, $PA$ are the same for all of them. The extracted spectrograms in the form of the flux integrated over all values of the position angle $PA$ in a function of the radius counted from the interferogram centre are plotted in Fig.\\,\\ref{fig:fig3} for the daylight sky, the circumsolar region during the total eclipse and the frame ``L\" including the coronal $\\mbox{Fe\\,{\\sc xiv}}, \\lambda 5302.86 \\rmn{\\AA}$ line. The 3 pics of the coronal emission line trace the 3 Fabry-P\\'{e}rot spectral orders. The relevant scattered $\\mbox{Fe\\,{\\sc i}}$, $\\mbox{Fe\\,{\\sc ii}}$ and $\\mbox{Mg\\,{\\sc i}}$ Fraunhofer lines, in absorption, are indicated. Their wavelengths, the values of the ring radius $r$ and the corresponding value of the elongation $\\epsilon$ are given in the Table\\,\\ref{tab:tab_lines}. Albeit the used Fabry-P\\'{e}rot orders overlap, the spectral features, fortunately, are distinct and can be easily identified. \\begin{figure*} \\begin{center} \\includegraphics [width=12 cm] {LayFig4.eps} \\end{center} \\caption{The line-on-sight velocity of the dust $\\bar{V}_{\\epsilon}$ \\textit{vs} the elongation $\\epsilon$; $\\bar{V}_{\\epsilon}$ is the average over all range of the position angle $PA$.} \\label{fig:fig_vel_elong} \\end{figure*} \\begin{figure*} \\begin{center} \\includegraphics [width=12 cm] {LayVelAngles.eps} \\end{center} \\caption{The line-on-sight velocity of the dust averaged over the studied range of elongations, $\\bar{V}_{PA}$, \\textit{vs} the position angle $PA$ on March 29, 2006 (upper plot) and on July 31, 1981 (lower plot) together with the fitted $sinus$ curves. For 2006, the triangle marks indicate measurements using two different daylight sky interferograms. The values for the $PA$ range $360\\degr-450\\degr$ were added for convenience, they merely repeat those for $0-90\\degr$. The $PA$ values of the extrema of the fitted to $V_{d}$ $sinus$ curves are given. For the circular orbits lying strictly in the ecliptic plane, the extrema should be at $90\\degr$ and $270\\degr$.} \\label{fig:fig_vel_pa} \\end{figure*} \\begin{figure*} \\begin{center} \\includegraphics [width=12 cm] {LayVradKep.eps} \\end{center} \\caption{\\textit{Upper plot}: The difference $V_{orb} = V - \\bar{V}_{\\epsilon}$ \\textit{vs} the elongation $\\epsilon$ in the directions of $PA=15\\degr$ (rectangles) and $PA=195\\degr$ (triangles) and the corresponding linear fits. The values of $V$ are the average over $\\pm30\\degr$ wide sectors, they are given in absolute values. The actual sign is positive in the first direction and negative in the second one. \\textit{Lower plot}: The average of the two and the linear fit to it, $V_{orb} = -2.2 + 3.1 \\epsilon $} \\label{fig:fig_vel_orb} \\end{figure*} \\subsection[]{The line-on-sight velocity of the dust. } The Doppler shift between the eclipse and the daylight spectrograms was measured by cross-correlation. The corresponding dust LOS velocity $V$ writes: \\begin{equation} V = \\frac{c \\Delta_{r}^{2}}{2f{^2}} \\label{eq:eq1} \\end{equation} or, substituting the numerical values: \\begin{equation} V= \\frac{3\\times 10^{6}\\times 0.028{^2}\\times \\Delta_{r}^2}{2\\times 86.65{^2}} = 0.01566 \\times \\Delta_{r}^2 \\label{eq:eq2} \\end{equation} where $f = 86.55$ mm is the focal distance of the camera objective, 0.028 mm is the pixel size, $\\Delta_{r}^2 = r_{e}^{2} - r_{sky}^{2}$, where $r_{e}$ and $r_{sky}$ are the radii in pixels respectively for the eclipse and daylight sky frames. The pixels convert to the elongation $\\epsilon$ given that $\\mathcal{R}_{\\sun} = 46.5$ pixels. The sampling in elongation $\\epsilon$ and in position angle $PA$ was as follows. We used intervals of $\\epsilon$ as given in Tab.\\,\\ref{tab:tab_range} choosing their spanning as regular as possible in a function of the useful spectral lines. The full range of $PA$ was simply divided on 36 sectors $20\\degr$ wide each. For convenience of the analysis, we computed two kinds of average velocities, $\\bar{V}_{\\epsilon}$, which is the average over all values of $PA$, and $\\bar{V}_{PA}$, which is the average over all values of the elongation $\\epsilon$. The value $\\bar{V}_{\\epsilon}$ is the function of $\\epsilon$ only; it is plotted in Fig.\\,\\ref{fig:fig_vel_elong}. If the LOS velocities of the dust grains are distributed in a central symmetry with the respect to the Sun, as it is expected for a pure circular motion, then $\\bar{V}_{\\epsilon}$ is 0; if not, it reflects the presence of a radial motion, an infall or outflow, at this particular value of $\\epsilon$. The Fig.\\,\\ref{fig:fig_vel_elong} shows that radial motions with velocities about $10\\,\\rmn{km}\\cdot\\rmn{s}^{-1}$ or less might be present at $\\epsilon=4.5$ and $\\epsilon\\approx9$, and they are absent, or very small, in between these elongations. The value $\\bar{V}_{PA}$ is a function of the position angle $PA$ only, it is plotted in Fig.\\,\\ref{fig:fig_vel_pa}. It is so particular that we give also for comparison a similar plot from the eclipse on July 31, 1981. For 2006, the triangle marks indicate measurements using two different daylight sky interferograms; their scattering provides an estimate of the uncertainty $\\delta \\bar{V}_{PA}\\simeq 1.6\\,\\rmn{km}\\cdot\\rmn{s}^{-1}$. The values for the $PA$ range $360\\degr-450\\degr$ are added for convenience, they merely repeat those for $0-90\\degr$. Plotted are also the least squares $sinus$ fits in the form: $\\bar{V}_{PA} = k_{1} + k_{2}\\cdot sin(k_{3}\\cdot PA +k_{4})$. The fit coefficients, for 2006, are as follows: $k_{1} = 0.3\\pm 0.3$, $k_{2} = 18.7\\pm 0.3$, $k_{3} = 0.01773\\pm 0.00016$ and $k_{4} = 1.22\\pm 0.04$, and for 1981, $k_{1} = 8.8\\pm 2.5$, $k_{2} = 22.0\\pm 3.5$, $k_{3} = 0.019\\pm 0.001$ and $k_{4} = 2.4\\pm 0.3$. ", "conclusions": "The reported here measurements of the Doppler shifts of the Fraunhofer lines, scattered by the dust grains in the solar F-corona, show that at the date of our observations the dust grains were on the orbit with a retrograde motion in a plane at $i\\approx 105\\degr$, i.e. nearly perpendicular to the ecliptics. This points to their cometary origin. Indeed, at the end of March, 2006, SOHO recorded several sungrazing comets with the orbital elements close to what was deduced from our measurements. We conclude that the contribution of comets to the dust content in the region close to the Sun can be important albeit variable in time. This contribution can explain the already noticed change of the dust distribution from one in the axial symmetry far from the Sun to that in the central, or may be spherical, symmetry \\citet{mann00}. We also derive that the observed plane of the dust grains orbit is slightly different from that of the parent comet(s), which indicates that the size of grains is small, less than 0.1 $\\umu$m, so that they are deviated from the initial orbit by the Lorentz force. This also means that the observed dust grains were released by the comet(s) shortly before our observations. The importance of comets in a circumstellar environment is general, let us recall e.g. the ``Falling Evaporating Bodies'' recorded in the spectra of $\\beta$ Pic \\citep[e.g.][]{beust98}. It would be interesting to investigate whether they can provide a sufficient transport of the dust grains between far and close environments of the central star, and to contribute to the dissemination of crystallized material in the recently detected exozodiacal dust discs (\\citealp{absil06}, \\citeyear{absil09}). For a more detailed discussion on the possible role of comet see also \\cite{augereau09} and references therein." }, "1003/1003.0483_arXiv.txt": { "abstract": "Recent cosmological observations indicate that the present universe is flat and dark energy dominated. In such a universe, the calculation of the luminosity distance, $d_L$, involve repeated numerical calculations. In this paper, it is shown that a quite efficient approximate analytical expression, having very small uncertainties, can be obtained for $d_L$. The analytical calculation is shown to be exceedingly efficient, as compared to the traditional numerical methods and is potentially useful for Monte-Carlo simulations involving luminosity distances. ", "introduction": "The most recent cosmological observations indicate that the present universe is flat and vacuum dominated~\\citep{Komatsu2009}. In such a vacuum dominated space-time, the distance analysis requires computer intensive numerical calculations. Even though, computers today are very fast, efficient analytical calculation of distance scales would be very useful for various types of Monte Carlo simulations. The most fundamental distance scale in the universe is the luminosity distance, defined by $d_L = \\sqrt{L/(4 \\pi f)}$, where $f$ is the observed flux of an astronomical object and $L$ is its luminosity. Current astronomical observations indicate that the present density parameter of the universe satisfy $\\Omega_\\Lambda + \\Omega_M = 1$ with $\\Omega_\\Lambda \\sim 0.7$. Here $\\Omega_\\Lambda$ is the contribution from the vacuum and $\\Omega_M$ is the contribution from all other fields. The distance calculations in such a vacuum dominated universe involve repeated numerical calculations and elliptic functions~\\citep{Eisenstein1997}. In order to simplify the numerical calculations, \\cite{Pen1999} (hereafter Pen99) has developed quite an efficient analytical recipe. In this paper, we show another analytical method, similar in many respect to that of Pen99, that can be used to calculate the distances in a vacuum dominated flat universe. Our analytical calculation is shown to run faster than that of Pen99 and has smaller error variations with respect to redshift ($z$) and $\\Omega_\\Lambda$. Our recipe for calculating the luminosity distance is the following ($H_0$ is the present Hubble constant and $c$ is the speed of light): \\begin{equation}\\label{eq1} d_L = \\frac{c}{3 H_0} \\frac{1+z}{\\Omega_\\Lambda^{1/6} (1-\\Omega_\\Lambda)^{1/3}}[\\Psi(x(0, \\Omega_\\Lambda)) - \\Psi(x(z, \\Omega_\\Lambda))], \\end{equation} \\begin{equation}\\label{eq2} \\Psi(x) = 3 \\, x^{1/3} 2^{\\,2/3} \\bigg[1-\\frac{x^2}{252}+\\frac{x^4}{21060}\\bigg], \\end{equation} \\begin{equation}\\label{eq3} x=x(z, \\Omega_\\Lambda) = \\ln (\\alpha + \\sqrt{\\alpha^2-1}), \\end{equation} \\begin{equation}\\label{eq4} \\alpha=\\alpha(z, \\Omega_\\Lambda) = 1+2\\,\\frac{\\Omega_\\Lambda}{1-\\Omega_\\Lambda}\\frac{1}{(1+z)^3}. \\end{equation} \\section[]{Approximation} We first begin by analyzing how the scale factor, $a(t)$ varies as a function of time $t$ in a flat universe in which $\\Omega_\\Lambda \\neq 0$. In this case, $a(t)$ is given by~\\citep{Weinberg2008} \\begin{equation}\\label{eq:no3} \\dot a^2 = H_{0}^{2} \\Omega_{\\Lambda} a^{2} + H_{0}^{2} \\Omega_{m} \\frac{a^{3}_{0}}{a}, \\end{equation} where $a_0$ is the present value of the scale factor. The above equation is then immediately integrated into \\begin{equation}\\label{eq:no4} \\bigg(\\frac{a}{a_0}\\bigg)^3 = \\frac{1}{2} \\frac{\\Omega_m}{\\Omega_\\Lambda} \\bigg[\\cosh (3 H_{0} t \\sqrt{\\Omega_\\Lambda}) -1\\bigg]. \\end{equation} The scale factor is directly related to the $z$ as, \\begin{equation}\\label{eq:no4b} \\frac{a}{a_0} = \\frac{1}{1+z}. \\end{equation} Let us define $x = 3 H_{0} t \\sqrt{\\Omega_\\Lambda}$ and indicate its present value by $x_0=x(0, \\Omega_\\Lambda)$. Then, equations \\ref{eq:no4} and \\ref{eq:no4b} give \\begin{equation}\\label{eq:no5} x=x(z, \\Omega_\\Lambda) = \\cosh^{-1} \\bigg[1+2\\,\\frac{\\Omega_\\Lambda}{1-\\Omega_\\Lambda}\\frac{1}{(1+z)^3}\\bigg] \\end{equation} If we define $\\alpha$ as follows \\begin{equation}\\label{eq4} \\alpha=\\alpha(z, \\Omega_\\Lambda) = 1+2\\,\\frac{\\Omega_\\Lambda}{1-\\Omega_\\Lambda}\\frac{1}{(1+z)^3} \\end{equation} and since $\\alpha > 1$ we can write $x$ as \\begin{equation}\\label{eq3} x=x(z, \\Omega_\\Lambda) = \\ln (\\alpha + \\sqrt{\\alpha^2-1}). \\end{equation} We note that $x$ is a monotonically decreasing function beyond $x(0, 0.7) = 2.42$. We choose the standard Robertson-Walker metric~\\citep{Weinberg2008} as the metric of the background space-time. With usual notation, this is \\begin{equation}\\label{eq:no6} ds^2 = c^2 dt^2 - a^2 \\bigg[\\frac{dr^2}{1-kr^2} + r^2 (d\\theta^2 + \\sin^2 \\theta \\, d\\phi^2) \\bigg]. \\end{equation} In the above space-time, we can use equation \\ref{eq:no3} to obtain $r$. A straightforward integration for a flat universe ($k = 0$) yields, \\begin{equation}\\label{eq:no7} r = \\frac{c}{a_0 H_0} \\frac{1}{3 \\Omega_\\Lambda^{1/6} \\Omega_M^{1/3}} \\int_{x}^{x_0} \\frac{dx'}{[\\sinh \\frac{x'}{2}]^{2/3}}. \\end{equation} This integral can be evaluated in terms of hypergeometric functions and related elliptic integrals. But here we take an alternate, simple approach by defining a new function, \\begin{equation}\\label{eq:no8} \\Psi(x) = \\lim_{\\delta \\rightarrow 0} \\int_{\\delta}^{x} \\frac{dx'}{[\\sinh \\frac{x'}{2}]^{2/3}}. \\end{equation} In the standard model the luminosity distance is defined as $d_L = a_0 r (1+ z)$. Now we can use equation \\ref{eq:no8} to write the luminosity distance as \\begin{equation}\\label{eq:no9} d_L = \\frac{c}{3 H_0} \\frac{1+z}{\\Omega_\\Lambda^{1/6} \\Omega_M^{1/3}}[\\Psi(x_0) - \\Psi(x)]. \\end{equation} Expanding $\\Psi$ in a series expansion to the 4th order, we find that \\begin{equation}\\label{eq:no10} \\Psi(x) = 3 \\, x^{1/3} 2^{\\,2/3} \\bigg[1-\\frac{x^2}{252}+\\frac{x^4}{21060}\\bigg] + \\Psi(0), \\end{equation} where $\\Psi(0) = -2.210$. Now, equation \\ref{eq:no9} reduces to the required expression for the luminosity distance as \\begin{equation}\\label{eq:no11} d_L = \\frac{c}{3 H_0} \\frac{1+z}{\\Omega_\\Lambda^{1/6} (1-\\Omega_\\Lambda)^{1/3}}[\\Psi(x_0) - \\Psi(x)]. \\end{equation} \\begin{figure} \\includegraphics[width=84mm]{lum_dist03} \\caption{The absolute relative percentage error ($\\Delta E$) as a function of the redshift for $\\Omega_\\Lambda = 0.7$.}\\label{WickPenCompare} \\end{figure} \\begin{figure} \\includegraphics[width=84mm]{lum_dist01} \\caption{Contour plot of absolute relative percentage error ($\\Delta E$) for the method of Pen99 with various $z$ and $\\Omega_\\Lambda$.}\\label{ContourPen} \\end{figure} \\begin{figure} \\includegraphics[width=84mm]{lum_dist02} \\caption{Contour plot of absolute relative percentage error ($\\Delta E$) for our method with various $z$ and $\\Omega_\\Lambda$.}\\label{ContourWick} \\end{figure} ", "conclusions": "In order to compare the method of Pen99 to ours, lets define the absolute relative percentage error as follows. \\begin{equation}\\label{eq:no12} \\Delta E = \\frac{|d_{L}^{\\tiny \\textrm{approx}} - d_{L}^{\\tiny \\textrm{num}}|}{d_{L}^{\\tiny \\textrm{num}}} \\times 100 \\textrm{\\%}. \\end{equation} Here $d_{L}^{\\tiny \\textrm{approx}}$ and $d_{L}^{\\tiny \\textrm{num}}$ are luminosity distance values calculated from approximate analytical methods and numerical method respectively. A comparison of $\\Delta E$ for both analytical methods for $\\Omega_\\Lambda = 0.7$ is shown in figure~\\ref{WickPenCompare}. Our method has a better absolute relative percentage error value for $z < 1.0$, $1.6 < z < 5.5$ and $z > 8.0$ compared to that of Pen99. We note that the error in our method decreases steadily with redshift approaching $< 0.014$ \\% at $z=1100$. In comparison, for high redshifts, Pen99 error always stays $\\sim$ 0.09\\% and does not decrease appreciably. A contour plot of $\\Delta E$ based on the method of Pen99 with various $z$ and $\\Omega_\\Lambda$ is shown in figure~\\ref{ContourPen}. Relatively complicated distribution of variations in the $\\Delta E$ can be seen for the parameter space characterized by $z$ and $\\Omega_\\Lambda$. However, a contour plot of $\\Delta E$ for our method, which is shown in Figure~\\ref{ContourWick}, shows a smooth behavior over the same parameter space. \\begin{figure} \\includegraphics[width=80mm]{HistoPenWick} \\caption{Histogram of running times of both Pen99 and our methods.}\\label{HistoPenWick} \\end{figure} In order to investigate the running time of the two analytical methods we performed the following test. With $z=1$ and $\\Omega_\\Lambda = 0.7$, we calculated the running time for 1 million calculations on a typical personal computer (Intel Core 2 Processor, 2127 MHz, 1 GB RAM, IDL\\footnote{Interactive Data Language \\\\ \\texttt{http://www.ittvis.com/ProductServices/IDL.aspx}} Version 6.2 running on Windows XP Service Pack 3). Then we repeated the above process 100 times for the both methods. Histogram of both running time results are shown in figure~\\ref{HistoPenWick}. Our method is significantly faster than the method of Pen99. In addition, we performed the same test on the numerical method and found that our method is more than an order of magnitude faster. However, we note that the above test is hardware and compiler dependent and results may vary depending on the hardware and the compiler used. With less than 0.1\\% error, our analytical method becomes quite desirable as the most interesting astronomical phenomena happen at $z > 1$ ($\\Omega_\\Lambda \\sim 0.7$). Furthermore, the analytical computation is more elegant and faster compared to traditional numerical computations invoked in connection with calculations of distances in a vacuum dominated flat universe. Once we know the luminosity distance, it becomes a simple matter to evaluate the other distances such as the angular diameter distance or the proper distance." }, "1003/1003.6093_arXiv.txt": { "abstract": "{The {\\it Kepler} space mission, successfully launched in March 2009, is providing continuous and high-precision photometry of thousands of stars simultaneously. The uninterrupted time-series of stars of all known pulsation types are a precious source for asteroseismic studies. The {\\it Kepler} data do not provide information on the physical parameters, such as $T_{\\rm eff}$, $\\log g$, metallicity, and $v \\sin i$, which are crucial for successful asteroseismic modelling. Additional ground-based time-series data are needed to characterize mode parameters in several types of pulsating stars. Therefore, ground-based multi-colour photometry and mid/high-resolution spectroscopy are needed to complement the space data. We present ground-based activities within KASC on selected asteroseismic {\\it Kepler} targets of several pulsation types.} ", "introduction": "The {\\it Kepler} satellite (Borucki et al. 1997), launched in March 2009, is collecting light curves with an unprecedented long time span of 3.5 years and a precision at the level of several ppm (Gilliland et al. 2010) for thousands of stars simultaneously. Of all {\\it Kepler} targets, more than 5000 stars have been selected as potential targets for seismic studies by the {\\it Kepler} Asteroseismic Science Consortium, KASC\\footnote{http://astro.phys.au.dk/KASC}. The selection of asteroseismic targets is one of the three basic aims of the KAI ({\\it Kepler} Asteroseismic Investigation). The other two are the asteroseismic characterization of planet hosting stars, e.g. derivation of accurate ages, masses, and radii (e.g. Stello et al. 2009), and the comparison of general stellar properties of Main Sequence stars with those of evolved stars. To fully exploit the excellent {\\it Kepler} light curves for asteroseismic means and to reach the science goals of the KAI, additional information from ground-based multi-colour photometry and high-resolution spectroscopy is indispensable (see, e.g., Uytterhoeven et al. 2009; Uytterhoeven 2009). The KASC subWorking Groups on ground-based observations (GBOsWG) take care of the organisation of ground-based observations in support of the {\\it Kepler} space data. In this paper we outline the importance of these efforts and we present an overview of the ground-based activities carried out within KASC to date. ", "conclusions": "" }, "1003/1003.3480_arXiv.txt": { "abstract": "% {} {The aim of this work is to study the contribution of the \\lae emitters to the star formation rate density (SFRD) of the Universe in the interval $2 1.5\\times10^{-18}erg/s/cm^{2}$ are $33$ galaxies/arcmin$^2$ and $\\sim 4 \\times 10^{-2}$Mpc$^{-3}$, respectively. We find that the the observed luminosity function of LAE does not evolve from z=2 to z=6. This implies that, after correction for the redshift-dependent IGM absorption, the intrinsic LF must have evolved significantly over 3 Gyr. The SFRD from LAE is found to be contributing about 20\\% of the SFRD at $z=2-3$, while the LAE appear to be the dominant source of star formation producing ionizing photons in the early universe $z\\sim>5-6$, becoming equivalent to that of Lyman Break galaxies. } {} ", "introduction": "The \\lae line is the strongest hydrogen emission line in the Universe, and it is observed in the optical range for galaxies at $z>2$. It has thus naturally been used to search for high-$z$ galaxies (Partridge\\&Peebles~1967; Djorgovski~et~al.~1985; Cowie~\\&Hu~1998). The \\lae emission in galaxies is thought to be produced by star formation, as the AGN contribution to the \\lae population at $z<4$ is found to be less than 5\\% (Gawiser~et~al.~2006; Ouchi~et~al.~2008; Nilsson~et~al.~2009). However, the physical interpretation of the observed \\lae flux is not simple, because the \\lae photons are resonantly scattered by neutral hydrogen. \\lae photons can therefore be more attenuated than other UV photons, and they have an escape fraction that can depend on the spatial distribution of neutral and ionized gas, as well as on the velocity field of the neutral gas (Giavalisco~et~al.~1996; Kunth~et~al.~1998; Mas-Hesse~et~al.~2003; Deharveng~et~al.~2008 and references therein). \\lae emitters (LAE) observed up to now are forming stars at rates of $\\sim1\\div10 M_{\\odot}yr^{-1}$ (Cowie\\&Hu~1998; Gawiser~et~al.~2006; Pirzkal~et~al.~2007), and they have stellar masses as low as $10^8\\div10^9 M_{\\odot}$ and ages $<50Myr$ (Pirzkal~et~al.~2007; ; Gawiser~et~al.~2007; Nilsson~et~al.~2009). However, Nilsson~et~al.~(2007) find ages between 0.1 and 0.9 Gyr, and more recently Pentericci~et~al.~(2007) and Finkelstein~et~al.~(2009) point out that \\lae galaxies are a more heterogeneous family than young star-forming galaxies: they also find \\lae emitters with old stellar populations (ages of $\\sim1~Gyr$) and a wide range of stellar masses (up to $10^{10} M_{\\odot}$). Interestingly, a class of \\lae galaxies with rest-frame equivalent width $EW>240$\\AA~ has been found (Malhotra\\&Roads~2002; Shimasaku~et~al.~2006). Galaxies with such large EW cannot be explained by star formation with a Salpeter IMF, but must have a top heavy IMF, a very young age $<10^7 yr$ and/or a very low metallicity. Many of these large EW objects are spatially extended, and thus are good candidates to be cooling clouds or primeval galaxies (Yang~et~al.~2006; Schaerer~2002) and can give interesting clues about the first stages of star formation. Together with studying the properties of \\lae galaxies at different redshifts, it is important to study the evolution of their luminosity function, comparing large and complete samples of \\lae galaxies at different redshifts. The most common technique used so far has been to build large samples from imaging in narrow band filters tuned to detect \\lae emission at $z\\sim2\\div9$ (Hu~et~al.~2004; Cuby~et~al.~2003; Tapken~et~al.~2006; Kashikawa~et~al.~2006; Gronwall~et~al.~2007; Murayama~et~al.~2007; Ouchi~et~al.~2008; Nilsson~et~al.~2009; Guaita~et~al.~2010). Blank field spectroscopy has been used blindly to search for \\lae emitters in deep HST-ACS slitless spectroscopic observations (Malhotra~et~al., 2005) or slit spectroscopy (van~Breukelen,~Jarvis~\\&~Venemans~2005; Martin~et~al.~2008; Rauch~et~al.~2008; Sawicki~et~al.~2008), with Rauch~et~al. (2008) exploring the faintest emitters. The general consensus today is that the apparent luminosity function of \\lae galaxies, that is the non-IGM corrected luminosity function, does not evolve at $z\\sim3\\div6$ (Rhoads\\&Malhotra~2001; Ouchi~et~al.~2003; van~Breukelen,~Jarvis~\\&~Venemans~2005; Shimasaku~et~al.~2006; Murayama~et~al.~2007; Gronwall~et~al.~2007; Ouchi~et~al.~2008; Grove~et~al.~2009). However, this conclusion is drawn from small samples, as the narrow band imaging techniques sample only thin slices in redshift (typically $\\Delta z$=0.1) and needs to be spectroscopically confirmed. Extensive spectroscopic follow-ups of narrow band \\lae candidates have been carried out in recent years, gathering hundreds of spectroscopic confirmations between $z\\sim2$ and $z\\sim7$. However, the spectroscopic coverage rarely reaches $30\\div50$\\% of the photometric sample (Murayama~et~al.~2007; Gronwall~et~al.~2007; Ouchi~et~al.~2008), but usually is much lower (Kashikawa~et~al.~2006; Nilsson~et~al.~2007; Matsuda~et~al.~2005). Moreover, current blind spectroscopic surveys sample only small areas to relatively shallow fluxes. It is likely that the apparent lack of evolution is a coincidence of the evolving intrinsic \\lae LF combined with an evolution of the intergalactic medium absorption with redshift (e.g. Ouchi~et~al.~2008). Firm conclusions about the possible evolution of the luminosity function have not yet been secured. Moreover, existing spectroscopic and narrow band samples are not sufficiently deep to constrain, even at intermediate redshift ($z\\sim3$), the slope of the luminosity function. Measuring the luminosity function in turn enables to compute the luminosity density and star formation rate density evolution under a set of well constrained hypotheses. The contribution of \\lae to the total star formation rate is yet not robustly measured mainly because the faint end slope of the luminosity function remains poorly constrained. In this paper, we present the results from a very deep blind spectroscopic survey search of Ly$\\alpha$ emitters over an unprecedented large sky area. We looked for the serendipitous detection of \\lae emission in the slits of the VIMOS VLT Deep Survey, concentrating on the VVDS--Deep and VVDS--Ultra-Deep surveys reaching up to 18h of integration on the VLT-VIMOS. We describe the spectroscopic and photometric data and the associated selection function in Section 2. The search for \\lae emitters and the final sample are presented in Section 3, and we discuss its properties in Section 4. The luminosity function calculation is discussed in Section 5, and the star formation rate density is derived. We discuss these results and give a summary in Section 6. Throughout the paper, we use and AB magnitudes and a standard Cosmology with $\\Omega_M=0.3$, $\\Omega_{\\Lambda}=0.7$ and $h=0.7$. ", "conclusions": "In this paper we have reported the discovery of 217 faint LAE in the range $2 \\leq z \\leq 6.62$ from targeted and serendipitous very deep observations using the VIMOS multi-slit spectrograph on the VLT. Adding together the areas covered by each slitlet combined to the wide wavelength coverage 5500-9350\\AA~ of the Deep survey and 3600-9350\\AA~ for the UltraDeep, we surveyed effective sky areas of 22.2 $arcmin^2$ and 3.3 $arcmin^2$ respectively. This produces a survey volume of $\\sim2.5\\times10^5 Mpc^3$, observed to unprecedented depth F$\\sim 1.5\\times10^{-18}erg/s/cm^{2}$. This volume is about one order of magnitude bigger than all other spectroscopic surveys produced up to now at comparable fluxes: van~Breukelen,~Jarvis~\\&~Venemans~(2005) sampled $10^4 Mpc^3$ down to 1.4$\\times10^{-17} erg/cm^2/s$; Martin~et~al.~(2008) sampled $4.5\\times10^4 Mpc^3$ down to similar fluxes in a narrow redshift range. Narrow band imaging surveys sampled bigger volumes than ours, but to a shallower flux: Ouchi~et~al.~(2008) covered a volume of $\\sim10^6 Mpc^3$, down to fluxes $\\sim2\\times10^{-17} erg/cm^2/s$. Serendipitous surveys have been presenting low number of objects (Sawicki~et~al., 2008), even if somewhat deeper (Malhotra~et~al., 2005; Rauch~et~al., 2008). We are therefore sampling deeper into the LAE luminosity function as we discussed in section \\ref{sect:lf}. From an observational point of view, we demonstrate the efficiency of blind LAE searches with efficient multi-slit spectrographs. The success of our approach is the result of combining a broad wavelength coverage to a large effective sky area, with long integration times, made possible by the high multiplex of the VIMOS instrument. The broad wavelength coverage has been essential to secure the spectroscopic redshifts from one single observation, without the need for follow-up to confirm the \\lae nature of the emission lines detected. This observing efficiency compares favorably with the time needed to perform narrow band imaging searches followed by multi-slit spectroscopy, and comparing the wide range in redshift covered by the former versus a narrow range for the latter. When the density of faint LAE is high, of the order several LAE/arcmin$^{2}$, multi-slit spectrographs become more efficient to secure a large number of confirmed sources than narrow band imaging searches, while at bright fluxes covering a wide field is essential to find rarer sources and narrow band imaging is more efficient. The two approaches will therefore remain complementary. Our main findings are the following: \\begin{enumerate} \\item We found a total of 217 LAE with confirmed spectroscopic redshifts in the range $2 \\leq z \\leq 6.62$, 133 coming from the serendipitous discovery in the multi-object spectrograph slits of the VVDS (105 from the Ultra-Deep and 28 from the Deep), and 84 coming from targeted VVDS observations of galaxies with $17.5 \\leq i_{AB} \\leq 24.75$ (Le F\\`evre~et~al., 2010, in prep.). About 50\\% of the Ultra-Deep and 40\\% of the Deep serendipitous targets have a detected optical counterpart down to magnitude $AB\\sim28$ in deep CFHTLS images. \\item The observed projected density of LAE with a \\lae emission brighter than F$\\sim 1.5\\times10^{-18}erg/s/cm^{2}$ in the range $2 \\leq z \\leq 6.6$ is $33$ LAE per arcmin$^2$, with $25$ LAE per arcmin$^2$ with $2 \\leq z \\leq 4.5$ and $8$ LAE per arcmin$^2$ with $4.5 < z \\leq 6.62$. The corresponding volume density of faint LAE with $L(Ly\\alpha) \\geq 10^{41}$ergs.s$^{-1}$ is $\\sim4 \\times 10^{-2}$Mpc$^{-3}$, a high density not yet observed at these redshifts. \\item The mean rest-frame EW(\\lae) of LAE in our sample range from about 40\\AA~ at $z\\sim2-3$ to $\\sim300-400$ at z$\\sim5-6$, and the star formation rate covers a wide range $0.1-20$M$_{\\odot}$yr$^{-1}$, assuming no ISM or IGM extinction and a Salpeter IMF. The HeII-1640\\AA~ emission has EW$\\sim4-14$ at $z\\sim2-4$ indicating the presence of young stars of a few Myr. We therefore detected vigorously star forming galaxies as well as galaxies with star formation comparable to dwarf starburst galaxies at low redshifts. \\item The \\lae {\\it apparent} luminosity function does not evolve between z=2 and z=6, within the error bars of our survey. Taking into account the average differential evolution in the IGM absorption with redshift therefore translates into a positive evolution of the {\\it intrinsic} \\lae LF of about 0.5 magnitude from $z\\sim2-3$ to $z\\sim5-6$. \\item We obtain a robust estimate of the faint end slope of the LAE luminosity function from a large sample of spectroscopically confirmed LAE. It is very steep: we find $\\alpha\\simeq-1.6$ at $z\\sim$2.5 and $1.8$ at $z\\sim$4. \\item The SFRD contributed by \\lae galaxies is increasing from $5\\simeq10^{-3}$M$_{\\odot}$yr$^{-1}$Mpc$^{-1}$ at $z\\simeq2.5$ to $\\simeq 2\\times10^{-2}$M$_{\\odot}$yr$^{-1}$Mpc$^{-1}$ at $z\\simeq6$. The contribution of the \\lae galaxies to the total SFRD of the universe as inferred by UV luminosity functions reported in the literature increases from $\\sim20$\\% at z=2.5 to $\\sim100$\\% at z=6. This seems to imply that all the galaxies that are forming stars at $z=6$ must show \\lae emission, therefore are in a very low dust medium. At z=2.5 80\\% of the star forming galaxies must have the \\lae emission produced by star formation blocked by some mechanism so that only 20\\% of the star forming galaxies show \\lae emission. A direct consequence would be that the \\lae escape fraction varies from 0.2 at $z\\simeq2.5$ to about 1 at $z\\simeq6$. This result would remain robust only if the total SFRD estimates based on UV luminosity functions using Lyman break galaxies identifications are complete. \\end{enumerate} The new VVDS measurements reported here bring a new important constraint to the LAE LF with a steep faint end slope observed at $25-6$, becoming equivalent to that derived from Lyman Break Galaxies searches. The steep faint end slope further implies that during reionisation sub-L$_*$ galaxies may have played an important role in keeping the universe ionized. These results further demonstrate that efforts dedicated to constraining the evolution of the luminosity function of high redshift LAE will remain an important tool to probe into the reionisation period." }, "1003/1003.3449_arXiv.txt": { "abstract": "{ We present a detailed analysis of the gas conditions in the H$_2$ luminous radio galaxy 3C~326~N at z$\\sim$0.1, which has a low star-formation rate (SFR$\\sim$0.07 M$_{\\odot}$ yr$^{-1}$) in spite of a gas surface density similar to those in starburst galaxies. Its star-formation efficiency is likely a factor $\\sim$10-50 lower than those of ordinary star-forming galaxies. Combining new IRAM CO emission-line interferometry with existing Spitzer mid-infrared spectroscopy, we find that the luminosity ratio of CO and pure rotational H$_2$ line emission is factors 10-100 lower than what is usually found. This may suggest that most of the molecular gas is warm. The Na~D absorption-line profile of 3C~326~N in the optical suggests an outflow with a terminal velocity of $\\sim -$1800 km s$^{-1}$ and a mass outflow rate of 30-40 M$_{\\odot}$ yr$^{-1}$, which cannot be explained by star formation. The mechanical power implied by the wind, of order $10^{43}$ erg s$^{-1}$, is comparable to the bolometric luminosity of the emission lines of ionized and molecular gas. To explain these observations, we propose a scenario where a small fraction of the mechanical energy of the radio jet is deposited in the interstellar medium of 3C~326~N, which powers the outflow, and the line emission through a mass, momentum and energy exchange between the different gas phases of the ISM. Dissipation times are of order $10^{7-8}$ yrs, similar or greater than the typical jet lifetime. Small ratios of CO and PAH surface brightnesses in another 7 H$_2$ luminous radio galaxies suggest that a similar form of AGN feedback could be lowering star-formation efficiencies in these galaxies in a similar way. The local demographics of radio-loud AGN suggests that secular gas cooling in massive early-type galaxies of $\\ge 10^{11}$ M$_{\\odot}$ could generally be regulated through a fundamentally similar form of 'maintenance-phase' AGN feedback. } ", "introduction": "\\label{sec:introduction} Molecular gas plays a critical role for our growing understanding of galaxy evolution. It often dominates the mass budget of the interstellar medium in galaxies, and is most closely related to the intensity at which galaxies form stars \\citep[e.g.,][]{kennicutt98}. Being strongly dissipative, it is also particularly susceptible to the astrophysical processes that drive galaxy evolution -- interactions, or feedback from starbursts and AGN -- and therefore plays a key role for our understanding of how these processes regulate star formation and galaxy assembly. It has only recently been recognized that powerful AGN may play a significant role in regulating galaxy growth over cosmological timescales by suppressing gas accretion and star formation \\citep[e.g.,][]{silk98,friaca98,scannapieco04,springel05,bower06,croton06,ciotti07,merloni08}. Such AGN 'feedback' would help resolve some of the remaining discrepancies between hierarchical models of galaxy evolution -- implying a rather gradual assembly of massive galaxies -- and observations, which suggest that massive galaxies formed most of their stars at high redshift, whereas star formation at later epochs was strongly suppressed. Observationally, a picture is emerging where radio jets may play a large role in transforming the energy ejected by the AGN into kinetic and thermal energy of the interstellar medium of the host galaxy. Observations of radio-loud AGN \\citep[e.g.,][]{heckman91, heckman91b, morganti05, emonts05, best05, best06, nesvadba06, nesvadba07, mcnamara07, nesvadba08, holt08, baldi08, fu09, humphrey09} and a large number of hydrodynamical simulations \\citep[e.g.,][]{krause05, saxton05, heinz06, sutherland07, merloni07, antonuccio08} suggest that radio-loud AGN inject a few percent of their mechanical energy into the ambient gas, parts of which produce significant outflows of warm gas \\citep{morganti03, morganti05, emonts05, nesvadba06, nesvadba07, holt08, nesvadba08, fu09}. However, most previous studies focused on the warm and hot gas at temperatures $\\ge 10^4$ K, and did not address the impact on the molecular phase, which is a serious limitation if we want to understand how the radio-loud AGN may regulate star formation in the host galaxy. Observations with the Spitzer IRS spectrograph recently revealed a significant number of ``H$_2$-luminous'' galaxies, where the molecular gas does not appear associated with star formation \\citep[][see also \\citealt{haas05}]{appleton06, egami06, ogle07, ogle08, ogle09, demessieres09, sivandam09}\\footnote{\\citet{ogle07, ogle09} propose to introduce a new empirical classification based on this H$_2$ excess, and refer to such targets as ``Molecular Hydrogen Emission Galaxies'', MOHEGs}. The mid-infrared spectra of H$_2$-luminous galaxies are dominated by bright, pure rotational emission lines of warm molecular hydrogen, (${\\cal L}(H_{2}) =10^{40}$ -- $10^{43}\\, {\\rm erg \\, s^{-1}}$), while classical star-formation indicators like a bright infrared continuum, mid-infrared lines of [NeII] and [NeIII], and PAH bands are weak or absent. Interestingly, \\citet{ogle09} find that 30\\% of their radio-loud AGN taken from the 3CR are H$_2$ luminous, suggesting this may be a common phenomenon which could be related to interactions with the radio source. To test this hypothesis and to evaluate possible consequences for the gas properties and star formation in the host galaxy, we have started CO emission-line observations of H$_2$ luminous radio galaxies with the IRAM Plateau de Bure Interferometer. Our goal is to constrain the physical properties and masses of the multiphase warm and cold gas in these galaxies, and to measure the gas kinematics. Here we present a detailed analysis of the multiphase gas content, energetics, and dissipation times of the H$_2$-luminous radio galaxy 3C~326~N at z=0.09 \\citep{ogle07,ogle08,ogle09}, which has particularly high H$_2$/PAH ratios. This analysis is based on our new CO(1-0) observations, as well as existing mid-infrared Spitzer and SDSS optical spectroscopy. Specifically we address three questions: What powers the H$_2$ emission in this galaxy? What is the physical state of the molecular gas, and perhaps most importantly, why is 3C~326~N not forming stars? We find that the interstellar medium of 3C~326~N has very unusual physical properties, where the warm molecular gas may dominate the overall molecular gas budget (\\S\\ref{ssec:massbudget}) and where the emission-line diagnostics suggest that the molecular as well as the ionized gas may be mainly excited by shocks (\\S\\ref{sec:diagnostics}) giving rise to luminous line emission at UV to mid-infrared wavelengths. We also identify a significant outflow of neutral gas from Na~D absorption profiles, which cannot be explained by star formation (\\S\\ref{sec:outflow}). We propose a physical framework in which these observations can be understood as a natural consequence of the energy and momentum coupling between the gas phases, which is driven by the mechanical energy injection of the radio jet (\\S\\ref{sec:enmom}). This scenario is an extension of the classical 'cocoon' model \\citep[e.g.,][]{scheuer74,begelman89} explicitly taking into account the multiphase character of the gas, with an emphasis on the molecular gas. We use our observational results to quantify some of the parameters of this scenario, including, perhaps most importantly, the dissipation time of the turbulent kinetic energy of the gas, and find that it is self-consistent and in agreement with the general characteristics of radio-loud AGN. 3C~326~N shows evidence for a low star-formation efficiency leading to a significant offset from the Schmidt-Kennicutt relationship of ordinary star-forming galaxies by factors 10$-$50, which is similar to other H$_2$ luminous radio galaxies with CO observations in the literature, as would be expected if 3C~326~N was a particularly clear-cut example of a common, underlying physical mechanism that is throttling star formation (\\S\\ref{sec:SF}). Long dissipation times suggest that the gas may remain turbulent over timescales of $10^{7-8}$ years, of order of the lifetime of the radio source, or perhaps even longer (\\S\\ref{ssec:dissipationtime}), while the energy supplied by the radio source may be sufficient to keep much of the gas warm for a Hubble time (\\S\\ref{sec:implications}) as required during the maintenance phase of AGN feedback, assuming typical duty cycles of order $10^8$ yrs. Throughout the paper we adopt a H$_0 =$70 km s$^{-1}$, $\\Omega_{M}=0.3$, $\\Omega_{\\Lambda} =$0.7 cosmology. In this cosmology, the luminosity distance to 3C~326~N is D$_L^{326N}=$ 411 Mpc. One arcsecond corresponds to a projected distance of 1.6 kpc. ", "conclusions": "We presented a detailed analysis of the physical gas conditions in the nearby powerful, H$_2$-luminous radio galaxy 3C~326~N, which does not show the signatures of active star formation in spite of few $\\times 10^9$ M$_{\\odot}$ in molecular gas. Our main goals were to investigate the gas conditions in this galaxy, to unravel the astrophysical mechanism which powers the remarkably bright mid-infrared H$_2$ emission, and to elucidate why this galaxy is not forming stars in spite of a considerable reservoir of molecular gas. To this end, we combined newly obtained CO(1-0) millimeter spectroscopy of 3C~326~N obtained at the Plateau de Bure Interferometer with the Spitzer-IRS spectroscopy by \\citet{ogle07,ogle09} and optical spectra from the SDSS. Our main conclusions are as follows: \\noindent (1) We compare gas masses for warm and cold gas, estimated from Spitzer mid-infrared spectroscopy and IRAM millimeter imaging spectroscopy. We find that most of the molecular gas in 3C~326~N is warm ($\\sim 2\\times 10^9$ M$_{\\odot}$, compared to $1.5\\times 10^{9}$ M$_{\\odot}$ estimated from our CO(1-0) observations for a ''standard'' CO-to-H$_2$ conversion factor). This ratio of warm to cold gas mass is about 1 to 2 orders of magnitude larger than that found in star-forming galaxies, indicating that the gas is in a distinct physical state, which may account for the low star-formation efficiency. \\noindent (2) We introduce a new 'molecular' diagnostic diagram based on the pure-rotational H$_2$, CO and PAH emission to show that most of the line emission from molecular gas in 3C~326~N is not produced by UV heating. We argue that the gas is likely to be powered by the dissipation of mechanical energy through shocks. The same is found for the ionized gas. \\noindent (3) Interstellar Na~D absorption marks an outflow of neutral gas at velocities of up $\\sim -1800$ km s$^{-1}$, and most likely mass and energy loss rates of 30$-$40 M$_{\\odot}$ yr$^{-1}$ and $\\sim 10^{43}$ erg s$^{-1}$, respectively. These values are similar to those found in starburst-driven winds, but this is the first time that such an outflow is detected in the Na~D line of a galaxy which does not have strong star formation. The star-formation intensity in 3C~326~N is orders of magnitudes below what is necessary to drive a wind. Similarly, the line profile cannot be explained through the interaction between 3C~326~N and 3C~326~S. \\noindent (4) Based on these observations we propose a scenario where the outflow and H$_2$ line emission are intricately related. It represents an extension of the well-explored 'cocoon'-model of interactions between radio jets and the ambient gas, and includes the physics of the molecular gas. In this scenario,the H$_2$ line emission is powered by turbulence induced within dense clouds that are embedded in the expanding cocoon, with a dissipation time of $\\sim 10^{7-8}$ yrs. Dissipation times of $10^8$ yrs are in rough agreement with the duty cycle of jet activity estimated from rejuvenated radio sources. \\noindent (5) Comparing PAH and CO surface brightness (in analogy to the ``Schmidt-Kennicutt'' diagram), we find a significant offset between 3C~326~N and other H$_2$-luminous galaxies towards lower PAH surface brightnesses. This may suggest that star-formation efficiencies in these galaxies are lower by roughly a factor $\\sim 10-50$ than those in 'ordinary' star-forming galaxies. \\noindent (6) Generalizing our results for 3C~326~N we find that outflows during similar radio-loud episodes in massive galaxies may balance the secular supply in cold gas through accretion and the mass return from evolved stars over a Hubble time. If radio-activity is a common, but episodic property of most massive early-type galaxies as suggested by \\citet{best06}, then the long dissipation times of up to $\\sim 10^8$ yrs may have consequences for the population of massive galaxies as a whole. We emphasize that this analysis can only be the first step towards an understanding of how radio-loud AGN regulate the physical conditions of the molecular gas and ultimately star formation in massive galaxies. It illustrates the importance of following a multi-wavelength approach if we seek to relate the physics and astrophysics of AGN feedback with the physical conditions of the interstellar medium including molecular gas in the host galaxy and the regulation of star formation." }, "1003/1003.4119_arXiv.txt": { "abstract": "Cosmology provides a unique and very powerful laboratory for testing neutrino physics. Here, I review the current status of cosmological neutrino measurements. Future prospects are also discussed, with particular emphasis on the interplay with experimental neutrino physics. Finally I discuss the possibility of a direct detection of the cosmic neutrino background and its associated anisotropy. ", "introduction": "Neutrino physics provides one of the prime examples of the interplay between particle physics and cosmology. Because neutrinos are so abundant, the cosmic neutrino background contributes significantly to the cosmic energy density at all times, and therefore have a profound influence on the evolution of our Universe. At early times, around the epoch of neutrino decoupling at $T \\sim 1$ MeV, they influence the formation of light nuclei, and later they have an influence on cosmic structure formation. This also means that precision cosmology can be used to probe details of neutrino physics, such as the absolute value of neutrino masses, and the presence of light or heavy sterile neutrinos. Here I will focus on the use of cosmology to probe the mass of neutrinos. As it turns out the main influence of light neutrinos on structure formation comes via their contribution to the energy density. At late times when at least some of the mass eigenstates are non-relativistic, the energy density in neutrinos can be quantified via the sum of neutrino masses, $\\sum m_\\nu$. To a very good approximation the effect of neutrinos on structure formation can be described using just this one parameter, i.e.\\ it does not matter how the mass is distributed internally between the different states (see \\cite{Lesgourgues:2004ps,Lesgourgues:2006nd} for a thorough discussion of this). Interestingly, cosmology is sensitive to a different combination of mass eigenstates than other experimental probes. In neutrinoless double beta decay experiments, the important parameter is the coherent sum \\cite{Aalseth:2004hb,Bilenky:1987ty} \\begin{equation} m_{\\beta\\beta} = \\left|c_{13}^2 c_{12}^2 m_1 + c_{13}^2 s_{12}^2 m_2 e^{i \\phi_2} + s_{13}^2 m_3 e^{i \\phi_3} \\right|, \\end{equation} which allows for phase cancellation. The current best upper bound on $m_{\\beta \\beta}$ comes from the Heidelberg-Moscow experiment and is $m_{\\beta \\beta} < 0.27 \\, {\\rm eV}$ (90\\% C.L.) \\cite{KlapdorKleingrothaus:2000sn,Rodin:2007fz}. However, the most direct current upper bound on the neutrino mass comes from the final data analysis of the Mainz experiment, and yields $m_{\\nu_\\beta} \\leq 2.3$ eV at 95\\% C.L. \\cite{kraus}, where \\begin{equation} m_{\\nu_\\beta} = \\left(c_{13}^2 c_{12}^2 m_1^2 + c_{13}^2 s_{12}^2 m_2^2 + s_{13}^2 m_3^2 \\right)^{1/2} \\end{equation} is effective parameter measured in beta decay spectra. In this case the mass states are weighed with their mixing with $\\nu_e$. Assuming just the three active neutrino species this corresponds to an upper bound on the sum of neutrino masses of $\\sim 7$ eV. The current neutrino temperature is $T_{\\nu,0} \\sim 1.7 \\times 10^{-4}$ eV so that any mass eigenstate heavier than this is non-relativistic at present. The contribution of neutrinos of any mass eigenstate, $i$, to the current energy density is given by $\\Omega_{\\nu,i} h^2 = m_{\\nu,i}/93$ eV, where $\\Omega$ is the density parameter and $h$ the Hubble parameter in units of $100 \\, {\\rm km} \\, {\\rm s}^{-1}$. The total neutrino contribution to the energy density is therefore $\\Omega_{\\nu,i} h^2 = \\sum_i m_{\\nu,i}/93$ eV, where the sum is over all non-relativistic states. The current upper bound on the dark matter density is roughly $\\Omega_{\\rm dm} \\lwig 0.1$ so that for $\\sum m_\\nu$, neutrinos would make up a very large fraction of the dark matter density. However, this is strongly ruled out by observations because of the free streaming property of neutrinos. Being light particles, neutrinos are relativistic approximately until the epoch of matter radiation equality. This means that all neutrino structures inside the horizon at this epoch have been erased, and if neutrinos constituted all the dark matter structure formation would have been impossible. This possibility is clearly excluded, and the argument can be refined to set constraints on the neutrino mass using precision measurements of cosmic structure formation. ", "conclusions": "Cosmology remains one of the main tools for the study of neutrino physics. Currently cosmology provides the most stringent upper bound on the neutrino mass, and even though the exact number is model dependent a very conservative upper bound on the sum of neutrino masses can be put at $\\sum m_\\nu \\lwig 0.6-0.7$ eV at 95\\% C.L. More aggressive use of data leads to more stringent bounds but in that case it is also necessary to rely on less well controlled effects from non-linear structure formation. The other important parameter which can be probed using structure formation data is the energy density in neutrinos at early times, quantified by the parameter $N_\\nu$, the effective number of neutrino species. The current bound from the same set of data is $3.03 < N_\\nu < 7.59$ at 95\\% C.L. Intriguingly, the preferred value of $N_\\nu$ is consistently higher than the standard model value of 3.04, but not at more than approximately 2$\\sigma$. In the future a range of different experiments will improve the sensitivity to neutrino parameters. Most important for a precision determination of $\\sum m_\\nu$ will be measurements of the matter power spectrum using larger volumes and going to higher redshifts than current surveys. During the next decade the most precise data will probably come from weak lensing survey of the LSST telescope. Together with measurements of the CMB anisotropy by the Planck satellite it has been estimated that a 1$\\sigma$ uncertainty on $\\sum m_\\nu$ of $\\sim 0.04$ eV can be achieved. In the more distant future it may be possible to decrease this error bar significantly by large scale measurements of 21-cm fluctuations at very high redshift, using for example the proposed FFTT project \\cite{fftt}. This could in principle increase the sensitivity by another factor of a few, making a precise neutrino mass determination possible even for the normal hierarchy. However, it should be stressed that there are many currently unadressed systematics involves in this, and it remains unclear if 21-cm surveys can ever reach this sensitivity. In conclusion, cosmology is an important and complementary laboratory for probing neutrino physics. Some neutrino parameters, like the neutrino mass, are in principle much easier to measure using precision cosmological data than in direct laboratory experiments. Furthermore, since cosmology is measuring a different effective mass quantity than beta or double beta decay experiments it remains an intriguing possibility that they will yield different and seemingly incompatible results. For example it may be the case that cosmology provides a stringent upper limit while for example a beta decay experiment shows positive evidence for a non-zero $m_\\nu$. Such a possibility could point to non-standard physics such as right handed currents masking as a neutrino mass in the beta decay experiment." }, "1003/1003.5215_arXiv.txt": { "abstract": "The theoretical maximum time variation in the electronic charge permitted by the Generalized Second Law of Thermodynamics applied to black holes radiating and accreting in the cosmic microwave background matches the measured cosmological variation in the fine structure constant claimed by Webb et al.. Such black holes cannot respond adiabatically to a varying fine structure constant. ", "introduction": "\\label{aba:sec1} Measurements\\cite{W} of absorption in the spectra from distant quasars suggest that the electromagnetic fine-structure constant -- $\\alpha = e^2/\\hbar c$ where $e$ is the electron charge, $\\hbar$ is Planck's constant and $c$ is the speed of light -- may be increasing as the Universe ages, at least in astrophysical environments. These observations motivated Davies {\\it et al.} \\cite{DDL} to apply the Second Law of Thermodynamics to black holes to derive theoretical limits on $\\alpha$ variation. Ref.~\\refcite{DDL} and subsequent papers\\cite{CV,FT,DK}\\,, however, applied the Generalized Second Law\\cite{B1} (GSL) incorrectly by investigating the entropy change with respect to $\\Delta\\alpha$, instead of a time interval $\\Delta t >0$. Here we take as our starting point that the net generalized entropy of the black hole system can not decrease {\\it over any} $\\Delta t >0$, i.e. that $\\Delta S_{tot}=\\Delta S_{bh} +\\Delta S_{env} \\ge 0$ where $\\Delta S_{bh} $ and $\\Delta S_{env} $ are the change in entropy of the black hole and of the ambient radiation and matter, respectively. We ask the question does the GSL applied to a black hole in our present Universe rule out a time variation in $e$ corresponding to the $\\alpha$ variation claimed in the quasar measurements, i.e. $de/dt\\approx 10^{-23} e$ per second. Below we summarize and update our detailed calculation of Ref.~\\refcite{JHM1}. The tightest constraints are derived by considering a charged, non-rotating black hole. In the following we use standard General Relativity and standard QED but extension to theories with additional terms is straightforward. In Section 3 we also discuss adiabaticity. ", "conclusions": "" }, "1003/1003.3575_arXiv.txt": { "abstract": "The evolution of high order correlation functions of a test scalar field in arbitrary inflationary backgrounds is computed. Whenever possible, exact results are derived from quantum field theory calculations. Taking advantage of the fact that such calculations can be mapped, for super-horizon scales, into those of a classical system, we express the expected correlation functions in terms of classical quantities, power spectra, Green functions, that can be easily computed in the long-wavelength limit. Explicit results are presented that extend those already known for a de Sitter background. In particular the expressions of the late time amplitude of bispectrum and trispectrum, as well as the whole high-order correlation structure, are given in terms of the expansion factor behavior. When compared to the case of a de Sitter background, power law inflation and chaotic inflation induced by a massive field are found to induce high order correlation functions the amplitudes of which are amplified by almost one order of magnitude. These results indicate that the dependence of the related non-Gaussian parameters -- such as $f_{\\rm NL}$ -- on the wave-modes is at percent level. ", "introduction": "It is now clearly established that standard single field inflation cannot produce significant non-Gaussianities (NG) during or immediately after the inflationary phase. This has been explicitly shown by Maldacena in Ref. \\cite{2003JHEP...05..013M} where it clearly appears that standard single field inflation leads to no or very little primordial non-Gaussianities. Multiple-field inflation has now long been recognized as a possible mechanism for the generation of primordial metric NG fluctuations (see recent review papers in Refs \\cite{2010CQGra..27l4001K,2010CQGra..27l4007L,2010CQGra..27l4002W,2010CQGra..27l4004B}). The exploration of the various types of coupling terms that appear in the action then leads to distinguish gravity from non-gravity mediated couplings (see \\cite{2010CQGra..27l4004B} for recent review on the origin of this distinction). This decomposition comes from the various behaviors of the terms that are present in the third order action. In one-field inflation, because the field fluctuations and the metric fluctuation are locked together, only the former can be found. In multiple-field inflation however this is not necessarily so. In particular isocurvature degrees of freedom open the possibility of having a richer phenomenology. And whereas gravity mediated couplings are ubiquitous but induce only modest effects \\cite{2005JCAP...09..011S,2006JCAP...05..019V}, non-gravity mediated couplings can be very efficient, although nothing ensures that they are generically at play (this is at the heart of Refs. \\cite{2007PhRvD..76d3526B,2003PhRvD..67l1301B,2002PhRvD..66j3506B}) . To make such perturbations play a role, one indeed needs a mechanism to transfer isocurvature modes into adiabatic fluctuations. For instance the curvaton model is based on the survival of (massive) isocurvature modes until late after the end of inflation that can alter the subsequent expansion history of the universe \\cite{2002PhLB..524....5L}. This is a particular case of modulated inflation \\cite{2003astro.ph..3614K,PhysRevD.69.023505,2004PhRvD..70h3004B}. Other mechanisms assume that isocurvature modes can change the end-point of inflation or alter the (p)-reheating sequences (as in \\cite{2009PhRvL.103g1301B}). Such mechanisms can also happen in the context of hybrid inflation. This latter situation is in particular advocated in Refs. \\cite{1990PhRvD..42.3936S,2002PhRvD..65j3505B,2002PhRvD..66j3506B,2003PhRvD..67l1301B} where isocurvature modes are shown to be able to induce large NGs in the metric fluctuations. In this paper we assume that such a mechanism is at play and focus our analysis on the intrinsic statistical properties of isocurvature modes, that is of modes that do not participate in the metric fluctuations at the time we are interested in. Even though this is a simple setting, exact results are still difficult to obtain since they depend on the actual dynamics of the expansion of the universe. Some exact results have been obtained in case of a de Sitter background (in Ref. \\cite{1993ApJ...403L...1F} bispectrum of a field with a cubic potential is derived; in Ref. \\cite{2003JHEP...05..013M} more complicated types of cubic interaction are considered and in Ref. \\cite{2004PhRvD..69f3520B} the four-point function is derived for quartic interactions) that has been the favorite playground for such calculations. The evolved bispectra, trispectra, have in particular been found to grow like the number of efolds after horizon crossing. This is an important result that lays the foundation for the computation of the mode dependence of the coupling parameters (which are usually expressed in terms of $f_{\\rm NL}$ parameter as introduced in \\cite{2002astro.ph..6039K}). This possibility is now being explored (as in Ref. \\cite{2009arXiv0911.2780B} for instance for some classes of model.) However, though inflationary models lead to background evolution that are generically close to a de Sitter, this is only an approximation, and which at best can be valid for only a limited period of time. In this article we aim at being as much exhaustive as possible in presenting exact results on the evolution of the statistical properties of a test scalar field, with non-vanishing self interaction terms, in arbitrary inflationary background. We will make use of the relation between quantum and classical evolutions. However we will not use explicitly the $\\delta N$ formalism (as introduced in \\cite{1982PhLB..117..175S,1996PThPh..95...71S} and see Ref. \\cite{2005PhRvL..95l1302L} in the context of nonlinear expansions) as it does not necessarily give controlled approximations but try to be as much precise as possible to draw the line between results that are of sub-Hubble origin and those that can be accounted for in a super-Hubble classical calculation. In the ``tree-theorem'' in Ref. \\cite{2008PhRvD..78f3534W} Weinberg demonstrated that there exists a classical system that exhibits the same tree order correlations properties as the quantum system. Although we will not exactly use this solution, it will eventually lead us to a complete description of the super-Hubble tree order correlation functions of the field. The paper is divided as follows. In section 2 we recall the method and the known results in case of a de Sitter background and describe in some details the late time super-Hubble limit of bispectra and trispectra. In the following section we make explicit the connection between calculations of correlation functions in a quantum context and classical calculations of stochastic field evolution. It shows in particular that the late time evolution of correlation functions can be computed in a classical context with the introduction of the Green function of the classical evolution of the free fields. This observation is exploited to obtain new results for the computation of the late time behavior of correlation functions in arbitrarily backgrounds. These results are presented in the section 4. In section 5 the resulting correlation functions are presented in a systematic way with the description of the tree structure of the field correlation functions at leading order in perturbation calculations. ", "conclusions": "We have explored the dependence of the evolution of the high order correlation functions of a test scalar field with minimal coupling on the background dynamics. The exact resolution of this problem requires involved quantum field theory calculation and there are only a limited number of cases where the calculations can be fully completed. They correspond to a de Sitter background and to some specific cases of power law inflation. Some of these results were already known but a novel and effective method is presented here that allows to perform those calculations efficiently. The explicit $k$-dependence of the bispectrum and trispectrum is derived for such cases for respectively a cubic or quartic potential. It is furthermore explicitly shown how the super-horizon evolution is equivalent to the mode coupling evolution of a classical system. It is then possible to derive the super-horizon evolution of the amplitude of the higher order spectra for arbitrary background. The result is encapsulated in Eq. (\\ref{nueexpression}). This expression is computed for different classes of models such as general power law inflation or background of chaotic inflation. It is found that when the background evolution departs from a de Sitter behavior the coupling amplitude is amplified. For realistic cases, we found that this amplification is about a factor 7. Those results suggest that the evolution is actually minimal when the background is de Sitter. It is to be noticed however that the resolution of the related classical system does not permit, even within a perturbation scheme, to compute the exact $k_{i}$ dependence of the effective $\\tau_{\\rm NL}$ and $f_{\\rm NL}$ parameters of the iso-curvature perturbations. Only in specific cases, where the quantum evolution can be computed, are those results known. If the background is close enough to de Sitter we found that the shape variations can induce variation of $\\tau_{\\rm NL}$ or $f_{\\rm NL}$ parameters up to 7 percent when the number of efolds is about 50 from horizon crossing to the end of inflation. The scale dependence (e.g. dependence with respect to scale for a given configuration) is precisely given by $1/N_{e}$ where $N_{e}$ is the number of efolds. Results obtained for a different background suggest that these dependence are about 7 times smaller for the inflationary backgrounds investigated in this paper. Those dependences are expected to be transferred to those of the metric fluctuations for certain models of multiple-hybrid inflation. Note however that for models where the transfer of modes takes place sooner, that is after a much smaller number of efolds, the $k$-dependence can be made arbitrarily larger. Its precise value depends however on the details of the model and needs a full quantum calculation to be derived. We finally show that when one restricts the computation to tree order, once super-Hubble scales have been reached, the high-order correlation functions follow a genuine tree structure (see text for details) with $k$-independent vertices. Correlation functions are then entirely encoded in the amplitude of the two-point correlation function and in the generating function of the vertices that can be explicitly computed for any backgrounds. Examples strongly suggest that, to a very good approximation, the background effects are entirely encoded in the expression of the lowest order vertex so that higher order correlation functions, when estimated for super-horizon limits, can indeed be built from this vertex value only. \\ifcqg" }, "1003/1003.1108_arXiv.txt": { "abstract": "We investigate how environmental effects by gas stripping alter the growth of a super massive black hole (SMBH) and its host galaxy evolution, by means of 1D hydrodynamical simulations that include both mechanical and radiative AGN feedback effects. By changing the truncation radius of the gas distribution ($R_{t}$), beyond which gas stripping is assumed to be effective, we simulate possible environments for satellite and central galaxies in galaxy clusters and groups. The continuous escape of gas outside the truncation radius strongly suppresses star formation, while the growth of the SMBH is less affected by gas stripping because the SMBH accretion is primarily ruled by the density of the central region. As we allow for increasing environmental effects - the truncation radius decreasing from about 410 to 50 kpc - we find that the final SMBH mass declines from about ${\\rm 10^{9}}$ to ${\\rm 8 \\times 10^{8} ~ M_{\\odot}}$, but the outflowing mass is roughly constant at about ${\\rm 2 \\times 10^{10} ~ M_{\\odot}}$. There are larger changes in the mass of stars formed, which declines from about ${\\rm 2 \\times 10^{10}}$ to ${\\rm 2 \\times 10^{9} ~ M_{\\odot}}$, and the final thermal X-ray gas, which declines from about ${\\rm 10^{9}}$ to ${\\rm 5 \\times 10^{8} ~ M_{\\odot}}$, with increasing environmental stripping. Most dramatic is the decline in the total time that the objects would be seen as quasars, which declines from 52 Myr (for $R_{t} = 377$ kpc) to 7.9 Myr (for $R_{t} = 51$ kpc). The typical case might be interpreted as a red and dead galaxy having episodic cooling flows followed by AGN feedback effects resulting in temporary transitions of the overall galaxy color from red to green or to blue, with (cluster) central galaxies spending a much larger fraction of their time in the elevated state than do satellite galaxies. Our results imply that various scaling relations for elliptical galaxies, in particular, the mass ratio between the SMBH and its host galaxy, can have dispersions due to environmental effects such as gas stripping. In addition, the simulations also suggest that the increase in AGN fraction in high-redshift galaxy clusters might be related to environmental effects which shut down the SMBH mass accretion in satellite galaxies and reduce their AGN activity. ", "introduction": "The role of the environment in galaxy evolution has been suggested in various forms which strip out gas from a galaxy and goes all the way back to the early suggestion by \\citet{spitzer51}. For example, the ram pressure of the intracluster medium is a possible way to strip out gas from falling galaxies partially or completely in galaxy clusters and to stop the supply of cold gas for star formation \\citep{gunn72,larson80,takeda84,gaetz87,begelman90,abadi99, domainko06,tonnesen08,kapferer09}. Other possible processes include thermal evaporation and viscous stripping \\citep{cowie77,livio80,nulsen82,nepveu85,valluri90,roediger08}. The combined gas loss by these different types of destruction effects is expected in galaxy clusters or groups \\citep{stevens99,quilis00,toniazzo01,kawata08, mccarthy08,smith09}. Tidal stripping can also play an important role in changing the gas and stellar mass of cluster or group galaxies \\citep{merritt83,moore99,dercole00}. Recent multi-wavelength observations have proved the loss of gases in cluster or group elliptical galaxies. In some elliptical galaxies, diffuse X-ray emitting gases show a long tail structure which can be explained by ram pressure stripping \\citep[e.g][]{kim08a,randall08}. Infrared observations also revealed dust emissions that can be from gas stripped from elliptical galaxies \\citep[e.g.][]{white91}. A central galaxy which hosts satellite galaxies in groups and clusters \\footnote{A primary galaxy accreting satellites in group and cluster environment corresponds to the brightest cluster galaxy (BCG) in this paper. For isolated galaxies, the term {\\it central galaxy} is correspondent to a primary galaxy which hosts satellites \\citep[e.g.][]{ann08}. In this paper, the key feature to define satellites is that they experience gas stripping which is strong enough to affect their evolution.} exhibits different features compared to satellites which can be affected by the various stripping processes. Because the central galaxy sits near the bottom of a deep gravitational potential well, its hot gas halo is much larger than those of satellites \\citep[e.g.][]{sun05}, and it is not surprising that it shows signatures of cooling flows in some cases \\citep[see][for a review]{fabian94,reiprich04}. Quite obviously, the differences in the stellar populations between satellites and central galaxies also have been studied, being relevant to our understanding of the environmental effects. For example, red satellite galaxies are redder than central galaxies of the same stellar mass \\citep{vandenbosch08}. Many central galaxies in clusters display either recent star formation or ongoing star formation which may be related to the cooling flow from their hot gas halos \\citep{cardiel98,rafferty06,bildfell08,odea08,pipino09} and which would be of reduced significance for satellite galaxies. Satellite galaxies in galaxy clusters seem to gradually lose of their gas by environmental effects and to truncate star formation, comparing their current and past star formation with those of field galaxies \\citep[e.g.][]{balogh99}. Differences in the properties of a central super massive black hole (SMBH) and its activity are expected between central and satellite galaxies, when considering the strong correlation between spheroidal galaxies and their SMBHs. The ratio of the SMBH mass over the spheroidal mass is found to be about $10^{-3}$ with a small dispersion for a large mass range in local galaxies \\citep{kormendy95,magorrian98, ferrarese00,gebhardt00,ferrarese02,yu02,marconi03,haring04}. If the growth of stellar mass and SMBH mass are coupled to the same kind of environmental effects, it is natural to conclude that the properties of SMBHs in satellite galaxies must be somewhat different from those of central galaxies. As we emphasized in our previous papers (Ciotti et al. 2009a, hereafter Paper I; Shin et al. 2010, hereafter Paper II) the self-regulated growth of SMBHs is profoundly dependent on how the accretion energy is converted to heat the ambient interstellar medium and how frequently and quickly the feedback process is ignited in the right place: we called these two key issues {\\it the problem of energy conversion} and {\\it the timing problem}. One would expect that the environmental effects might induce differences in the frequency of AGN feedback and the growth of the SMBHs by varying the conditions of the self-regulation process. In this paper, we tackle the issue of environmental effects on the coevolution of the SMBH and its host galaxy by simulating the evolution of a galaxy in hydrodynamical models with a simplified setup of different environments. Although there is no current systematic investigation of the SMBH mass and AGN activity for separating satellite and central galaxies, the theoretical prediction from our simulations can be used to constrain the hypothesis of the coevolving SMBH and its host galaxy and the connection to environmental effects. This paper is organized as follows. In Section \\ref{sec:sim}, we describe the models and the simulations. The results are presented in Section \\ref{sec:result}. Discussion and conclusions follow in Section \\ref{sec:discussion}. ", "conclusions": "\\label{sec:discussion} In this paper, we have investigated the effects of AGN feedback (in the form of combined radiative and mechanical energy and momentum deposition) on galaxies in different environments corresponding satellite and central galaxies. Although not including fully realistic stripping effects, our simulations predict various observable features that can be compared to the current and future observations, and which should change in a systematic way between central and satellite galaxies. The simulation results need to be interpreted as a guide to understand a general consequence of environmental effects. Here, we discuss implications from our simulations. First of all, the local scaling relationship such as the mass ratio between the SMBH and its host galaxy can have a dispersion due to environmental effects. For example, if gas in a galaxy has experienced stripping, the growth of its SMBH is less disturbed than star formation which generally happens over the whole galaxy. Hence, the mass ratio between the SMBH and the host galaxy of fixed galaxy mass could be significantly higher in satellite ellipticals in galaxy clusters. This also implies that a bias effect on deriving the scaling relationship probably exists due to selecting satellite galaxies less likely as observation samples. Objects measured for their SMBH mass are generally bright objects which might not include galaxies having experienced gas stripping. Even though the recent measurements of the SMBH mass might have an intrinsic scatter of the ratio \\citep[e.g.][]{kim08b,gultekin09}, the current measurements of the SMBH mass in satellite elliptical galaxies have not been conducted systematically. Probably, the systematic investigation of early-type galaxies in the Virgo Cluster might be useful to test our results for satellites in future \\citep{cote06,decarli07,gallo08}. But we warn that the uncertainty and biases in measuring the SMBH mass might dominate the dispersions in the mass ratios even for galaxies in the Virgo cluster \\citep[e.g][]{bernardi07}. The results of our simulations also imply the less frequent and earlier termination of AGN activity in satellite galaxies compared to central galaxies (i.e. galaxies with a large $R_{t}$). This result can be tested against the AGN fraction in galaxy clusters and groups \\citep[e.g.][]{gilmour07}. The observed AGN fraction in galaxy clusters is generally higher than in field galaxies partially because of the high number density of galaxies (i.e. the effect from the cluster richness) and how AGN activities are selected in radio, X-ray, and IR observations \\citep[e.g.][]{hickox09}. When considering this effect of the high number density of galaxies, AGN activity actually seems suppressed in galaxy clusters \\citep[e.g.][]{koulouridis10}. The evolution of AGN fractions in galaxy clusters as a function of redshift also shows the increase in AGN activity at higher redshifts when satellite galaxies are still frequently active in their SMBH accretion, supporting our hypothesis of environmental effects on AGN activity \\citep{galametz09, martini09}. The frequent multiple intensive SMBH accretion events in simulated central galaxies is consistent with several distinctive observable features of cluster central galaxies. For example, many central galaxies show multiple outburst structures such as several gas clumps and surrounding ripples even though the origin of those observed structures and what kind of AGN activity causes those is still uncertain \\citep[e.g.][]{fabian05,graham08,blanton09, clarke09}. The very weak dependence of the ejected mass on $R_{t}$ implies that satellite galaxies can be important sources of metal enrichment in the intracluster medium. Even though our simulation does not follow the evolution of the metal contents in the outflowed gas, the mass stripped out of the galaxy, which is about $2 \\times 10^{10} {\\rm M_{\\odot}}$ (see Figure \\ref{fig:time_comp}), must be metal-rich being produced by gas recycling from the evolving stars and supernovae explosion. As explained in \\S\\ref{sec:sim}, our models are similar to the {\\it leaky box model} of galactic chemical evolution for satellites \\citep{pagel09} because of the outflowing boundary condition. Finally, the outflows from both satellites and central galaxies might cause mixing which is coupled to AGN activity and turbulence within the intracluster medium \\citep[e.g.][]{osullivan05}. The detailed abundance patterns in the cluster gas should thus reflect the outflow from the satellite cluster ellipticals in addition to the processes by a cluster central galaxy \\citep{domainko06,schindler08,rasmussen09,sivanandam09}. The recent star formation in cluster central galaxies \\citep[e.g.][]{rafferty08,kirkpatrick09} is also naturally explained in our simulations by continuous recurrence of cooling flows and consequently heating by AGN feedback in galaxies with a large $R_{t}$. In our simulations, star formation from a cooling flow is temporally permitted while some amount of the cooling flow is accreted onto the SMBH and triggers AGN feedback which finally heats up and reverses the cooling flow as the cool gas is consumed by star formation. Contrary to a classical solution to the cooling flow problem in the cluster galaxies which stops the cooling flow completely \\citep[see][for a review]{fabian94}, our solution to the cooling flow problem for the cluster central galaxies is simply permitting temporary rejuvenation of star formation and AGN activity which is self-regulated, leading to the termination of the cooling flow \\citep[see][for a further discussion]{ciotti07}. This kind of solution can be called the {\\it intermittent cooling flow} scenario \\citep{salome06,bildfell08,pipino09,wilman09}. A further test of our results will be the Eddington ratio of the rejuvenated SMBH accretion in local cluster central galaxies because the SMBH accretion rate in the simulated central galaxies is significantly lower than the Eddington accretion rate despite the change of $R_{t}$. To make our simulations more realistic, we need to improve our simulations by including two important physics: feedback by a radio jet from AGN activity and cosmological setup of environmental effects. In addition to the limits of our simulations already mentioned in our previous Paper II, for example, 1D calculations, and possible different initial conditions such as models presented in Appendix, the radio jet is an important feedback component for cluster central galaxies in environments of galaxy clusters and groups \\citep{begelman04,konigl06}. The observation of radio jets proves that the kinetic energy of jets is enough to be an effective feedback mode even though we do not have a clear explanation about how much of the relevant energy is transported to ISM and intracluster medium \\citep{mcnamara07,shin11}. We note that powerful narrow jets such as that seen in the giant elliptical M87 tend to drill through the ambient gas in the galaxy depositing some amount of energy with a relativistic fluid directly to the intergalactic medium \\citep{ferrari98,owen00}. Our implementation of stripping effect from different environments is simply parametrized by $R_{t}$ in one dimension, and is not close to dynamically changing environmental effects in cosmological evolution \\citep[e.g.][]{takeda84,vollmer01,toniazzo01}. Moreover, our simulations do not include effects from tidal gravitational field which is commonly expected in cluster or group environments \\citep{dercole00}. Although the simple implementation makes the interpretation of simulation results obvious, the direct test of simulations will require better but complicated models of ram-pressure stripping and other environmental processes such as tidal stripping in galaxies which have realistic orbits in galaxy clusters and groups. \\appendix" }, "1003/1003.3611_arXiv.txt": { "abstract": "{Since most of the star clusters do not survive the embedded phase, their early dissolution appears to be a major source of field stars. However, catching a cluster in the act of dissolving is somewhat elusive.} {We investigate the nature and possible evolution of the young Galactic star clusters Collinder\\,197 (Cr\\,197) and vdB\\,92.} {Photometric and structural properties are derived with near-infrared photometry and field-star decontamination. Kinematical properties are inferred from proper motions of the main sequence (MS) and pre-MS (PMS) member stars.} {The colour-magnitude diagrams (CMDs) are basically characterised by a poorly-populated MS and a dominant fraction ($\\ga75\\%$) of PMS stars, and the combined MS and PMS CMD morphology in both clusters consistently constrains the age to within $5\\pm4$\\,Myr, with a $\\sim10$\\,Myr spread in the star formation process. The MS$+$PMS stellar masses are $\\approx660^{+102}_{-59}\\,\\ms$ (Cr\\,197) and $\\approx750^{+101}_{-51}\\,\\ms$ (vdB\\,92). Cr\\,197 and vdB\\,92 appear to be abnormally large, when compared to clusters within the same age range. They have irregular stellar radial density distributions (RDPs) with a marked excess in the innermost region, a feature that, at less than 10\\,Myr, is more likely related to the star formation and/or molecular cloud fragmentation than to age-dependent dynamical effects. The velocity dispersion of both clusters, derived from proper motions, is in the range $\\sim15 - 22\\,\\kms$.} {Both clusters appear to be in a super-virial state, with velocity dispersions higher than those expected of nearly-virialised clusters of similar mass and size. A possible interpretation is that Cr\\,197 and vdB\\,92 deviate critically from dynamical equilibrium, and may dissolve into the field. We also conclude that early cluster dissolution leaves detectable imprints on RDPs of clusters as massive as several $10^2\\,\\ms$. Cr\\,197 and vdB\\,92 may be the link between embedded clusters and young stellar associations. } ", "introduction": "\\label{Intro} It is now well-established that a significant fraction of the embedded star clusters dissolve into the field on a time-scale of a few $10^7$\\,Myrs. Basically, dissolution occurs mainly because the gravitational potential can be rapidly reduced by internal processes, such as the impulsive gas removal by supernovae and massive star winds associated with this early period. As a consequence, an important fraction of the stars, especially of low mass, end up moving faster than the scaled-down escape velocity, and may be lost to the field (e.g. \\citealt{GoBa06}). This process can dissolve the very young star clusters on time-scale of $10-40$\\,Myr. From the stellar content perspective, the early cluster dissolution depends essentially on the effective star-formation efficiency, the total mass converted into stars, and the mass of the more massive stars (e.g. \\citealt{tutu78}; \\citealt{GoBa06}). However, there is also evidence indicating that the determining factor for cluster survival during gas expulsion is the virial state of the stars just before the onset of this phase, so that clusters formed with a dynamically cold stellar component are more likely to survive (\\citealt{GoodW09}). In any case, the early dissolution of embedded clusters may lead to the formation of OB stellar groups (e.g. \\citealt{Goul00}), the subsequent dispersion of which may be an important source of field stars (e.g. \\citealt{Massey95}). Indeed, \\citet{LL2003} suggest that only about 5\\% of the Galactic embedded clusters dynamically evolve into gravitationally bound open clusters (OCs). With such a high dissolution rate, the embedded clusters could be the major contributors of field stars in galaxies for generations. However, recent studies suggest that the early dissolution rate in the Magellanic Clouds is significantly lower ($\\la30\\%$ - \\citealt{deGG08}; \\citealt{deGG09}) than in the Milky Way (\\citealt{LL2003}), or in other galaxies such as the Antennae (e.g. \\citealt{Whit07}) or M\\,51 (e.g. \\citealt{Bast05}). Observationally, low-mass star clusters younger than $\\sim10$\\,Myr, in general, have Colour-Magnitude Diagrams (CMDs) with an under-populated and developing main sequence (MS) and a more conspicuous population of pre-MS (PMS) stars. Typical examples are NGC\\,6611, NGC\\,4755, NGC\\,2239, NGC\\,2244, Bochum\\,1, Pismis\\,5, NGC\\,1931, vdB\\,80, and BDBS\\,96 (\\citealt{Pi5} and references therein). In terms of CMD morphology, bound and un-bound young clusters are expected to present similar evolutionary sequences. On the other hand, the important early changes in the potential that affect the large-scale internal structure of clusters should be reflected on the stellar radial density profile (RDP). Bochum\\,1 (\\citealt{Bochum1}) and NGC\\,2244 (\\citealt{N2244}), for instance, appear to be representatives of this scenario (i.e. structures evolving towards dissolution in a few $10^7$\\,yr), in which an irregular RDP cannot be represented by a cluster-like (i.e. an approximately isothermal sphere) profile. In this context, irregular RDPs in young clusters - when coupled to an abnormally high velocity dispersion - may reflect significant profile erosion or dispersion of stars, and point to important deviations from dynamical equilibrium. The rather complex interplay among environment conditions, effective star-formation efficiency (as defined in \\citealt{GoBa06}), and total mass converted into stars, is probably what explains the difference between (dissolving) objects like Bochum\\,1 and NGC\\,2244, and bound young OCs (in which the MS$+$PMS stars distribute according to a cluster RDP as in NGC\\,6611 and NGC\\,4755). Consistent with the above mass-dependent scenario, the MS$+$PMS mass of Bochum\\,1 and NGC\\,2244 is a factor 2-3 lower than in NGC\\,6611 and NGC\\,4755. In this paper we investigate the nature of the poorly-studied, young (age $\\sim5$\\,Myr; Sect.~\\ref{DFP}), large (with radii within $8-12$\\,pc; Sect.~\\ref{struc}), and PMS rich (with stellar masses within $660 - 750\\,\\ms$; Sect.~\\ref{MF}) clusters Cr\\,197 and vdB\\,92. Our main goal is to determine whether such young stellar systems can be characterised as typical OCs or if they are heading towards dissolution. In addition, we will derive their fundamental and structural parameters, most of these for the first time. \\begin{figure} \\begin{minipage}[b]{0.50\\linewidth} \\includegraphics[width=\\textwidth]{Fig1a.eps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.50\\linewidth} \\includegraphics[width=\\textwidth]{Fig1c.eps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.50\\linewidth} \\includegraphics[width=\\textwidth]{Fig1b.eps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.50\\linewidth} \\includegraphics[width=\\textwidth]{Fig1d.eps} \\end{minipage}\\hfill \\caption[]{Top: $20\\arcmin\\times20\\arcmin$ DSS-II R images centred on Cr\\,197 (left) and vdB\\,92 (right). Gas emission, dust reflection and/or absorption are present in the fields in varying proportions. Bottom: 2MASS \\ks\\ images covering $5\\arcmin\\times5\\arcmin$ (Cr\\,197) and $6\\arcmin\\times6\\arcmin$ (vdB\\,92). Orientation: North to the top and East to the left.} \\label{fig1} \\end{figure} \\begin{table*} \\caption[]{Fundamental parameters} \\label{tab1} \\tiny \\renewcommand{\\tabcolsep}{0.73mm} \\renewcommand{\\arraystretch}{1.25} \\begin{tabular}{cccccccccccccccc} \\hline\\hline &\\multicolumn{6}{c}{WEBDA}&&\\multicolumn{8}{c}{This work}\\\\ \\cline{2-7}\\cline{9-16} Cluster&$\\alpha(2000)$&$\\delta(2000)$&Age&\\ebv&\\ds&D&&$\\alpha(2000)$&$\\delta(2000)$& $\\ell$&$b$&Age&\\ebv&\\ds&\\dSC\\\\ & (hms)&($\\degr\\,\\arcmin\\,\\arcsec$)&(Myr)&(mag)&(kpc)&(\\arcmin)&&(hms)&($\\degr\\,\\arcmin\\,\\arcsec$) &(\\degr)&(\\degr)&(Myr)&(mag)&(kpc)&(kpc)\\\\ (1)&(2)&(3)&(4)&(5)&(6)&(7)&&(8)&(9)&(10)&(11)&(12)&(13)&(14)&(15)\\\\ \\hline Cr\\,197&08:44:51&$-$41:14:00&13&0.55&0.84&17.0&&08:44:40.3&$-$41:16:48.4&261.51&$+$0.94 &$5\\pm4$&$0.34\\pm0.16$&$1.05\\pm0.20$&$0.23\\pm0.04$\\\\ vdB\\,92&07:03:54&$-$11:29:23&40&0.25&1.50&24.0&&07:03:56.4&$-$11:34:54.7&224.66&$-$2.52 &$5\\pm4$&$0.22\\pm0.13$&$1.38\\pm0.26$&$1.04\\pm0.19$ \\\\ \\hline \\end{tabular} \\begin{list}{Table Notes.} \\item Col.~7: Optical diameter; Col.~14: distance from the Sun; col.~15: distance from the Solar circle. WEBDA data for Cr\\,197 is based on photometry from \\citet{VM73}. \\end{list} \\end{table*} This paper is organised as follows. In Sect.~\\ref{RecAdd} we recall literature data on both objects. In Sect.~\\ref{2mass} we discuss the 2MASS photometry and build the field-star decontaminated CMDs. In Sect.~\\ref{DFP} we derive fundamental cluster parameters. In Sect.~\\ref{struc} we derive structural parameters. In Sect.~\\ref{MF} we estimate cluster mass. In Sect.~\\ref{propMot} we build the intrinsic proper motion distribution. In Sect.~\\ref{Discus} we compare structural parameters and dynamical state with those of a sample of template OCs. Concluding remarks are given in Sect.~\\ref{Conclu}. ", "conclusions": "\\label{Conclu} Given the complex interplay among environmental conditions, star-formation efficiency, initial dynamical state of the stellar content, and the total mass converted into stars or expelled, the majority of the embedded clusters do not survive the first few tens of Myr, especially the low-mass ones. In this context, it is important to investigate the structural and photometric properties of OCs (with a mass range and in different environments) that are undergoing this early phase. In the present paper we derive astrophysical parameters and investigate the nature of the young OCs Cr\\,197 and vdB\\,92. Their location in the $3^{rd}$ Galactic quadrant minimises contamination by field stars. We work with field-star decontaminated 2MASS photometry (with errors $\\la0.1$\\,mag), which enhances CMD evolutionary sequences and stellar RDPs, yielding more constrained fundamental and structural parameters. The decontaminated CMDs are characterised by similar properties, an under-populated and developing MS, a dominant fraction ($\\ga75\\%$) of PMS stars, and some differential reddening. In both cases, the MS and PMS CMD morphologies consistently imply a time-spread of $\\sim10$\\,Myr in the star formation. Thus, we set the age of Cr\\,197 and vdB\\,92 as constrained within $5\\pm4$\\,Myr. Their MS$+$PMS stellar masses are $\\approx660^{+102}_{-59}\\,\\ms$ (Cr\\,197) and $\\approx750^{+101}_{-51}\\,\\ms$ (vdB\\,92). By means of the proper motion of the member MS and PMS stars, we estimate their velocity dispersions to be in the range $\\approx15\\,\\kms$ to $\\approx22\\,\\kms$. Compared to a set of {\\em classical} bound OCs, both Cr\\,197 and vdB\\,92 appear to have core and cluster radii abnormally large, with $\\rc\\approx1.5,~1.9$\\,pc and $\\rl\\approx12,~8$\\,pc, respectively for Cr\\,197 and vdB\\,92. Structurally, the stellar RDPs follow a cluster-like profile for most of the radial range, except in the central region, where they have a pronounced cusp. At less than about 10\\,Myr, this cusp is probably related to the star formation and/or molecular cloud fragmentation, and not the product of dynamical evolution. A possible conclusion is that Cr\\,197 and vdB\\,92 deviate critically from dynamical equilibrium, and similarly to the equally young NGC\\,2244 and Bochum\\,1 (of comparable mass), and Pismis\\,5 and NGC\\,1931 (of significantly lower mass), they are heading towards dissolution. This interpretation is also consistent with the super-virial state of both clusters (as well as the dissolving OC NGC\\,2244), which have velocity dispersions much higher than the $\\sigma_v\\approx1\\,\\kms$ expected for nearly virialised clusters of similar mass and size as Cr\\,197 and vdB\\,92. In this context, Cr\\,197 and vdB\\,92 may be taken as a link between embedded clusters and young stellar associations. We provide evidence that early star cluster dissolution may be detected, for instance, by means of important and systematic irregularities in the stellar radial density profile of clusters as massive as several $10^2\\,\\ms$. Studies like the present one are important for a better understanding of the crucial early evolution of embedded star clusters - and the dependence on mass and environment - that rarely result in a {\\em classical} bound OC or, more frequently, lead to their dissolution into the field." }, "1003/1003.1564_arXiv.txt": { "abstract": "{ We previously attempted to ascertain why the Li~{\\sc i}~6708 line-strengths of Sun-like stars differ so significantly despite the superficial similarities of stellar parameters. We carried out a comprehensive analysis of 118 solar analogs and reported that a close connection exists between the Li abundance ($A_{\\rm Li}$) and the line-broadening width ($v_{\\rm r+m}$; mainly contributed by rotational effect), which led us to conclude that stellar rotation may be the primary control of the surface Li content.} { To examine our claim in more detail, we study whether the degree of stellar activity exhibits a similar correlation with the Li abundance, which is expected because of the widely believed close connection between rotation and activity. } { We measured the residual flux at the line center of the strong Ca~{\\sc ii} 8542 line, $r_{0}(8542)$, known to be a useful index of stellar activity, for all sample stars using newly acquired spectra in this near-IR region. The projected rotational velocity ($v_{\\rm e}\\sin i$) was estimated by subtracting the macroturbulence contribution from $v_{\\rm r+m}$ that we had already established. } { A remarkable (positive) correlation was found in the $A_{\\rm Li}$ versus (vs.) $r_{0}(8542)$ diagram as well as in both the $r_{0}(8542)$ vs. $v_{\\rm e}\\sin i$ and $A_{\\rm Li}$ vs. $v_{\\rm e}\\sin i$ diagrams, as had been expected. With the confirmation of rotation-dependent stellar activity, this clearly shows that the surface Li abundances of these solar analogs progressively decrease as the rotation rate decreases. } { Given this observational evidence, we conclude that the depletion of surface Li in solar-type stars, probably caused by effective envelope mixing, operates more efficiently as stellar rotation decelerates. It may be promising to attribute the low-Li tendency of planet-host G dwarfs to their different nature in the stellar angular momentum. } ", "introduction": "Since Li nuclei are burned and destroyed on their arrival at the hot stellar interior ($T \\ga 2.5 \\times 10^{6}$~K), we can gain valuable information from the surface Li composition of a star about the past history and the physical mechanism of stellar envelope mixing. It has been known, however, that Li abundances ($A_{\\rm Li}$) in Sun-like stars exhibits puzzling behaviors: \\begin{itemize} \\item A markedly large diversity (by more than $\\sim 2$~dex) of $A_{\\rm Li}$ is seen despite the similarity of stellar parameters. \\item Planet-host stars tend to show appreciably lower $A_{\\rm Li}$ than non-planet-host stars (cf. Israelian et al. 2004, 2009). \\end{itemize} As these characteristic trends cannot be explained by the naive classical picture of surface Li being determined by age (which relates to the duration time of gradual Li depletion by way of convective mixing) and $T_{\\rm eff}$ (which affects the depth of the convection zone), it has been important to find the hidden parameter(s) responsible for these observed findings. To elucidate this problem, Takeda et al. (2007, hereinafter referred to as Paper I) conducted an extensive high-precision study of stellar parameters as well as of $A_{\\rm Li}$ for 118 solar analogs and found that $A_{\\rm Li}$ values, exhibiting a large dispersion themselves, are closely correlated with the line-width, which is characterized by the macroscopic velocity dispersion ($v_{\\rm r+m}$) including the rotational as well as the macroturbulent broadening effect. We then speculated that $v_{\\rm e}$ (equatorial rotation velocity) would be the most important factor affecting $A_{\\rm Li}$, since the star-to-star variation in $v_{\\rm e} \\sin i$ may be responsible for the spread in $v_{\\rm r+m}$, any considerable fluctuation in the macroturbulent velocity field among similar solar-type stars being difficult to imagine. The motivation of the present paper, the second in a series, is to check (or substantiate) the hypothesis that stellar rotation is the decisive factor which determines the surface Li content of solar-analogs. One useful way to accomplish this would be to examine the stellar activity, which is considered to be of dynamo origin and thus deeply related to the intrinsic rotational rate. That is, if we could confirm that $A_{\\rm Li}$ is closely correlated with the degree of activity, our speculation would be reasonably justified. As an indicator of stellar activity, we adopt $r_{0}$(8542) ($\\equiv f_{0}/f_{\\rm cont}$), which is the residual flux (normalized by the continuum) at the line center of Ca~{\\sc ii} 8542.09, the strongest line of the near-IR 8498/8542/8662 triplet of mulptiplet 2 for the $^{2}$D--$^{2}$P$^{\\rm o}$ transition. This is known to reflect the chromospheric activity of a star; i.e., as the activity is enhanced, the core flux increases because of the greater amount of filled-in emission from the chromosphere (see, e.g., Linsky et al. 1979). This quantity is known to be well correlated with the more traditional Ca~{\\sc ii} H+K emission index ($\\log R'_{\\rm HK}$) and thus serves as a useful tool for diagnosing the activity level of late-type stars (e.g., Foing et al. 1989; Chmielewski 2000; Bus\\`{a} et al. 2007). In this study we aim to determine $r_{0}$(8542) (a measure of stellar activity\\footnote{Of course, this $r_{0}$(8542) index depends not only on the chromospheric activity but also on atmospheric parameters such as $T_{\\rm eff}$ (effective temperature), $\\log g$ (surface gravity), and [Fe/H] (metallicity) (e.g., Mallik 1997; Chmielewski 2000). However, in the present sample of solar analogs similar to each other, the mutual differences of these stellar parameters are of secondary importance and can be neglected to a first approximation.}) for each of the 118 stars studied in Paper I (a bona-fide sample of solar analogs), based on our new spectroscopic data obtained at Okayama Astrophysical Observatory, and examine whether or not they show any correlation with $A_{\\rm Li}$, to test our conclusion in Paper I. The remainder of this paper is organized as follows. In Sect. 2, we describe the observational material and the measurement of $r_{0}$(8542). Before discussing the results of stellar activity, the projected rotational velocity ($v_{\\rm e} \\sin i$) for each star is derived in Sect. 3 by appropriately subtracting the contribution of macroturbulence from the macrobroadening parameter ($v_{\\rm r+m}$) discussed in Paper I. The discussion about the resulting relationship between $r_{0}$(8542), $A_{\\rm Li}$, and $v_{\\rm e}\\sin i$ is presented in Sect. 4, where we show that the arguments in Paper I have been confirmed, and our conclusions are summarized in Sect. 5. Two additional appendices are included. Appendix A describes the results of our reanalysis of stellar parameters (including Li abundance) for HIP~41484, since another star (actually HIP~41184) was erroneously observed and analyzed as if it were HIP~41484 in Paper I. Appendix B is devoted to discussing the sensitivity difference between two representative activity indicators, Ca~{\\sc ii} triplet in near-IR (multiplet 2) and Ca~{\\sc ii} H+K lines in violet region (multiplet 1), based on some test results of non-LTE line profiles simulated with trial models. ", "conclusions": "In our previous study of Paper I, we carried out a comprehensive spectroscopic analysis of 118 solar analogs to clarify why the strengths of Li~{\\sc i}~6708 line in these Sun-like stars are considerably diversified despite that they have stellar parameters quite similar to each other, and interestingly found a close relationship between the Li abundance and its line-width. We then proposed that stellar rotation may be the most important parameter in determining the surface Li content. In this paper, we have tried to test this hypothesis by examining whether any correlation exists between the stellar activity and the Li abundance, as expected because of the widely believed rotation--activity connection. As an indicator of stellar activity, we used the residual line-center flux of the strong Ca~{\\sc ii} 8542 line ($r_{0}$), which was measured from the high-dispersion near-IR spectra obtained with the 188~cm reflector and the HIDES spectrograph at Okayama Astrophysical Observatory. The projected rotational velocity ($v_{\\rm e}\\sin i$) was reasonably accurately estimated by subtracting the contribution of the macroturbulence effect from the line-broadening width ($v_{\\rm r+m}$) as we already established in Paper~I. Clear correlations have been confirmed in the diagrams $A_{\\rm Li}$ vs. $r_{0}(8542)$, $r_{0}(8542)$ vs. $v_{\\rm e}\\sin i$, and $A_{\\rm Li}$ vs. $v_{\\rm e}\\sin i$), which support the arguments that (1) the stellar activity surely depends upon the rotational rate, and that (2) the atmospheric Li abundance of solar-analog stars declines progressively as the rotational velocity decreases. We thus concluded that a Li-depletion mechanism in these Sun-like stars, most probably caused by effective envelope mixing, operates more efficiently as the stellar rotation slows down. In this context, it may be interesting/enlightening to interpret the observational finding of a low-Li tendency of planet-host G dwarfs within the framework of the rotational properties (i.e., difference in the angular momentum), as stated in the theoretical prediction by Bouvier (2008). Additional detailed investigations along those lines would be worthwhile. However, the cause of this interconnection, which is found for comparatively high-rotation/activity/Li stars, remains unclear for the group of stars with low-rotation/activity/Li, where $r_{0}(8542)$ tends to converge and stabilize at $\\sim 0.2$ and can no longer be a useful activity indicator. Since we found from our non-LTE calculation that Ca~{\\sc ii} H+K violet lines at 3968/3934~$\\rm\\AA$ are more sensitive and useful (than Ca~{\\sc ii} IR triplet lines) for investigating the mild stellar activity of comparatively slow rotators, it would be beneficial to revisit this problem by studying these Ca~{\\sc ii} H+K lines in greater detail." }, "1003/1003.3561_arXiv.txt": { "abstract": "The early evolution of Earth's atmosphere and the origin of life took place at a time when physical conditions at the Earth where radically different from its present state. The radiative input from the Sun was much enhanced in the high-energy spectral domain, and in order to model early planetary atmospheres in detail, a knowledge of the solar radiative input is needed. We present an investigation of the atmospheric parameters, state of evolution and high-energy fluxes of the nearby star $\\kappa^1$~Cet, previously thought to have properties resembling those of the early Sun. Atmospheric parameters were derived from the excitation/ionization equilibrium of \\ion{Fe}{1} and \\ion{Fe}{2}, profile fitting of H$\\alpha$ and the spectral energy distribution. The UV irradiance was derived from FUSE and HST data, and the absolute chromospheric flux from the H$\\alpha$ line core. From careful spectral analysis and the comparison of different methods we propose for \\kap the following atmospheric parameters: \\Teff = 5665 $\\pm$ 30 K (H$\\alpha$ profile and energy distribution), $\\log g = 4.49 \\pm 0.05$ dex (evolutionary and spectroscopic) and $[Fe/H]=+0.10 \\pm0.05$ (\\ion{Fe}{2} lines). The UV radiative properties of \\kap indicate that its flux is some 35\\% lower than the current Sun's between 210 and 300 nm, it matches the Sun's at 170 nm and increases to at least 2--7 times higher than the Sun's between 110 and 140 nm. The use of several indicators ascribes an age to \\kap in the interval $\\sim$ 0.4--0.8~Gyr and the analysis of the theoretical HR diagram suggests a mass $\\sim$ 1.04~M$_\\odot$. This star is thus a very close analog of the Sun when life arose on Earth and Mars is thought to have lost its surface bodies of liquid water. Photochemical models indicate that the enhanced UV emission leads to a significant increase in photodissociation rates compared with those commonly assumed of the early Earth. Our results show that reliable calculations of the chemical composition of early planetary atmospheres need to account for the stronger solar photodissociating UV irradiation. ", "introduction": "The irradiation from the parent star is, by far, the most important source of energy in planetary atmospheres. Most of the physical and chemical properties of the atmosphere of a planet are largely driven by the stellar input, which ultimately determines its structure, composition and, even, its mere existence. Most of the radiation emitted by the Sun comes from the photosphere and, at the solar effective temperature today (${T_{\\rm eff}}_{\\odot} = 5780$ K), it is the dominant source at wavelengths above 170 nm. The solar photospheric flux is quite stable over short timescales (decades) and it only suffers variations driven by sunspots and faculae with an amplitude below 0.2--0.3\\% peak to peak \\citep{froehlichlean2004}. Over long timescales, the variations have been larger and are related to the nuclear evolution of the Sun. Model predictions indicate that the young ZAMS Sun could have had a luminosity some 35\\% lower than today \\citep[unless a scenario of heavy mass loss is considered;][]{sackmannboothroyd2003}. The energetic end of the solar spectrum (i.e., below 170 nm) is dominated by the emissions from high-temperature plasma in the chromosphere, transition region and corona. Such high-energy fluxes are strongly variable over short- and mid-term timescales because of flare events, rotational modulation, activity cycles, etc. \\citep[e.g.,][]{leanetal1997}. Over longer timescales, several past studies \\citep{ZW82,A97} established that the young Sun's high-energy emissions were possibly up to several orders of magnitude stronger than currently. The Sun in Time programme \\citep[hereafter Paper I]{dorrenguinan1994,ribasetal2005} focused on a small sample of carefully-selected and well-studied stellar proxies that represent key stages in the evolution of the Sun. This approach allowed the study in the X-ray, EUV, and FUV domains, where the variations are of one order of magnitude or more. However, the stellar proxy technique becomes increasingly uncertain in the UV since the stars are not perfect matches to the Sun (their masses are within 10\\% of 1~M$_\\odot$) and the expected flux variations are in the range of tens of a percent. For this reason, most of the Sun in Time stars employed in Paper I cannot be reliably used to infer the solar flux evolution at wavelengths longer than about 130 nm. But at the same time, the wavelength interval between 130 and 200 nm is an essential energy input to planetary atmospheres since it drives most of the photochemical reactions \\citep[e.g.,][]{CLA82,CLA83}. In addition, its impact on living organisms must also be considered, although it is unlikely that photons in this range could penetrate a dense atmosphere \\citep{cockell2000}. Among the solar proxies studied in the Sun in Time, \\kap (HD 20630, HIP 15457, $d=9.16$~pc, $V=4.84$) stands out as potentially having a mass very close to solar and a young age (see Paper I). This could be a very good analog of the Sun at the critical time when life is thought to have originated on Earth 3.8 Gyr ago, at the start of the Archean epoch \\citep{mojzsisetal1996}. It is also about the time when Mars lost its liquid water inventory at the end of the Noachian epoch some 3.7 Gyr ago \\citep{jakoskyphillips2001}. For these reasons, \\kap deserves attention as a possible precise match to the young Sun, thus providing information on the radiation environment that determined the properties and chemical composition of the planetary atmospheres. \\citet{cnossenetal2007} presented a study based on \\kap with a goal set on assessing the biological implications of the high-energy radiations. The authors did not use a real high-energy spectrum but generated synthetic data using plasma models from an inferred emission measure distribution. In this paper we carry out an in-depth analysis of $\\kappa^1$~Cet, including its radiative properties, chemical abundances, atmospheric parameters and state of evolution, with the ultimate goal of understanding the Sun's UV emissions at a critical time in the past. The paper is organized as follows. In \\S 2 we describe the spectroscopic analysis and derivation of atmospheric parameters, along with its Li abundance. The high-energy radiative processes and magnetic activity of \\kap are discussed in the context of active stars in \\S 3. In \\S 4 we determine its mass and state of evolution, and in \\S 5 we employ its properties to discuss the young Sun in the frame of photodissociation calculations. Our conclusions are drawn in \\S 6. ", "conclusions": "In this paper we have carried out an in-depth study of the bright, nearby solar analog $\\kappa^1$~Cet. Several methods have been used to estimate its effective temperature and chemical composition, yielding preferred values of \\Teff = 5665 $\\pm$ 30~K, and $[Fe/H] = +0.10 \\pm 0.05$. The systematic offset between \\Teff values obtained from photometry/line profiles and the excitation/ionization \\ion{Fe}{1} and \\ion{Fe}{2} equilibria is evidence of non-LTE effects, probably related to UV overionization due to the strong magnetic activity of this star. Adopting as the best atmospheric parameters the photometry and H$\\alpha$ profile \\Teff and the \\ion{Fe}{2} metallicity, we have been able to set constraints to the stellar age, which should be between 0.4 and 0.8~Gyr. All the information gathered indicates that \\kap is a star with nearly 1 solar mass in a relatively unevolved evolutionary stage. As such, it is an excellent match to the Sun as it was some 3.7--4.1 Gyr ago. The radiation from the young Sun must have played an essential role in shaping the atmospheres of the Solar System planets. In particular, the UV flux is responsible for the photochemical processes in the atmosphere. We have been able to compile data, taken both with FUSE and HST, covering the entire UV and it shows that $\\kappa^1$~Cet's flux is some 35\\% lower than the current Sun's between 210 and 300 nm, it matches the Sun's at 170 nm and increases to at least 2--7 times higher than the Sun's between 110 and 140~nm. We have compared these fluxes with a ``theoretical'' young Sun estimated by adding the current chromospheric flux to a photospheric model with the correct radiative properties. We have used a photochemical model to calculate the photodissociation rates of the most relevant molecules in the assumed composition of early Earth's atmosphere. The results indicate that such rates should have been several times higher than those resulting from a simplistic ``theoretical'' solar spectrum. Our calculations demonstrate that self-consistent planetary atmosphere calculations must account for the much stronger photodissociating radiation of the young Sun. The resulting chemistry could be significantly different from that commonly assumed. This is obviously very relevant at a significant point in the Solar System evolution, when life was gaining a secure foothold on Earth and Mars lost its liquid water inventory." }, "1003/1003.0754_arXiv.txt": { "abstract": " ", "introduction": "During the last thirty years the spatial distribution of galaxies has been investigated from the point of view of geometrical and physical theories. One first target was to reproduce the two-point correlation function $\\xi (r)$ for galaxies which on average scales as $\\approx (\\frac {r} {5.7Mpc})^{-1.8}$, see \\citet{Jones2005} and \\citet{Gallagher2000}. The statistical theories of spatial galaxy distribution can be classified as \\begin{itemize} \\item {\\bf Levy flights}: the random walk with a variable step length can lead to a correlation function in agreement with the observed data, see \\citet{Mandelbrot1975b}, \\citet{Soneira1977}, \\citet{Soneira1978} and \\citet{Peebles1980}. \\item {\\bf Percolation}: the theory of primordial explosions can lead to the formation of structures, see \\citet{Charlton1986} and \\citet{ferraro}. Percolation is also used as a tool to organize : (i) the mass and galaxy distributions obtained in 3D simulations of cold dark matter (CDM) and hot dark matter (HDM), see \\citet{Klypin1993}, (ii) the galaxy groups and clusters in volume-limited samples of the Sloan Digital Sky Survey (SDSS), see \\citet{Berlind2006} \\end{itemize} The geometrical models are well represented by the {\\bf Voronoi Diagrams}. The applications to galaxies started with \\citet{icke1987}, where a sequential clustering process was adopted in order to insert the initial seeds, and they continued with \\citet{Weygaert1989}, \\citet{pierre1990}, \\citet{barrow1990}, \\citet{coles1991}, \\citet{Weygaert1991a}, \\citet{Weygaert1991b}, \\citet{Subba1992}, \\citet{Ikeuchi1991} and \\citet{Goldwirth1995}. An updated review of the 3D Voronoi Diagrams applied to cosmology can be found in \\citet{Weygaert2002} or \\citet{Weygaert2003}. The 3D Voronoi tessellation was first applied to identify groups of galaxies in the structure of a super-cluster, see \\citet{Ebeling1993}, \\citet{Bernardeau1996}, \\citet{Schaap2000}, \\citet{Marinoni2002}, \\citet{Melnyk2006}, \\citet{Schaap2009} and \\citet{Elyiv2009}. The physical models that produce the observed properties of galaxies are intimately related, for example through the Lagrangian approximation, and can be approximately classified as \\begin{itemize} \\item {\\bf Cosmological N-body}: Through N-body experiments by \\citet{Aarseth1978} it is possible to simulate groups which are analogous to the studies of groups among bright Zwicky-catalog galaxies, see \\citet{Gott1979a} or covariance functions in simulations of galaxy clustering in an expanding universe which are found to be power laws in the nonlinear regime with slopes centered on 1.9 \\citet{Gott1979b}. Using gigaparticle N-body simulations to study galaxy cluster populations in Hubble volumes, \\citet{Evrard2002} created mock sky surveys of dark matter structure to z~=1.4 over $10000^{\\circ}~sq.~deg$ and to z~=0.5 over two full spheres. In short, N-body calculations seek to model the full nonlinear system by making discrete the matter distribution and following its evolution in a Lagrangian fashion, while N-body simulations are usually understood to concern gravity only. \\item {\\bf Dynamical Models}: Starting from a power law of primordial inhomogeneities it is possible to obtain a two-point correlation function for galaxies with an exponent similar to that observed, see \\citet{Peebles1974a,Peebles1974b,Gott1975}. Another line of work is to assume that the velocity field is of a potential type; this assumption is called the Zel'dovich approximation, see \\citet{ZelDovich1970,ZelDovich1989,Coles1995}. The Zel'dovich formalism is a Lagrangian approximation of the fully nonlinear set of equations. In this sense it is ``gravity'' only and does not include a pressure term. \\item {\\bf The halo models}: The halo model describes nonlinear structures as virialized dark-matter halos of different mass, placing them in space according to the linear large-scale density field which is completely described by the initial power spectrum, see \\citet{Neyman1952,Scherrer1991,Cooray2002}. Figure~19 in \\citet{Jones2005}, for example, reports the exact nonlinear model matter distribution compared with its halo-model representation. \\end{itemize} The absence of clear information on the 3D displacement of the physical results as a function of the redshift and the selected magnitude characterize the cosmological N-body, the dynamical and the hydrodynamical models. This absence of detailed information leads to the analysis of the following questions: \\begin {itemize} \\item Is it possible to compare the theoretical and observational number of galaxies as a function of the redshift for a fixed flux/magnitude ? \\item What is the role of the Malmquist bias when theoretical and observed numbers of galaxies versus the redshift are compared? \\item Is it possible to find an algorithm which describes the intersection between a slice that starts from the center of the box and the faces of irregular Poissonian Voronoi Polyhedrons? \\item Is it possible to model the intersection between a sphere of a given redshift and the faces of irregular Poissonian Voronoi Polyhedrons? \\item Does the developed theory match the observed slices of galaxies as given, for example, by the 2dF Galaxy Redshift Survey? \\item Does the developed algorithm explain the voids appearance in all sky surveys such as the RC3? \\item Can voids between galaxies be modeled trough the Voronoi normalized volume distribution? \\item Is it possible to evaluate the probability of having a supervoid once the averaged void's diameter is fixed? \\item Is it possible to compute the correlation function for galaxies by introducing the concept of thick faces of irregular Voronoi polyhedrons? \\item Is it possible to find the acoustic oscillations of the correlation function at $\\approx~100Mpc$ in simulated slices of the Voronoi diagrams. \\end{itemize} In order to answer these questions, Section~\\ref{formulary} briefly reviews the standard luminosity function for galaxies. An accurate test of the number of galaxies as a function of the redshift is performed on the 2dF Galaxy Redshift Survey (2dFGRS), see Section~\\ref{test}. Section~\\ref{Voronoisec} reports the technique which allows us to extract the galaxies belonging to the Voronoi polyhedron and Section~\\ref{cellular} simulates the redshift dependence of the 2dFGRS as well as the overall Third Reference Catalog of Bright Galaxies (RC3). Section \\ref{sec_corr} reports the simulation of the correlation function computed on the thick faces of the Voronoi polyhedron. ", "conclusions": "{\\bf Photometric maximum} The observed number of galaxies in a given solid angle with a chosen flux/magnitude versus the redshift presents a maximum that is a function of the flux/magnitude. From a theoretical point of view, the photometric properties of the galaxies depend on the chosen law for the luminosity function. The luminosity function here adopted (the Schechter function) predicts a maximum in the theoretical number of galaxies as a function of the redshift once the apparent flux/magnitude is fixed. The theoretical fit representing the number of galaxies as a function of the redshift can be compared with the real number of galaxies of the 2dFGRS which is theory-independent. The superposition of theoretical and observed fit is satisfactory and the $\\chi^2$ has been computed, see Figure~\\ref{maximum_flux}. The position of the maximum in the number of galaxies for different magnitudes is a function of the redshift and in the interval $15 < bJmag< 18.5 $ the comparison between observed and theoretical data is acceptable, see Figure~\\ref{zeta_max_flux}. Particular attention should be paid to the Malmquist bias and to equation~(\\ref{range}) that regulate the upper value of the redshift that defines the complete sample. \\noindent {\\bf 3D Voronoi Diagrams} The intersection between a plane and the 3D Voronoi faces is well known as $V_p(2,3)$. The intersection between a slice of a given opening angle, for example $3^{\\circ}$, and the 3D Voronoi faces is less known and has been developed in Section~\\ref{faces}. This intersection can be calibrated on the astronomical data once the number of Poissonian seeds is such that the largest observed void matches the largest Voronoi volume. Here the largest observed void is 2700 $Km/sec$ and in order to simulate, for example, the 2dFGRS, 137998 Poissonian seeds were inserted in a volume of $(131908~Km/sec)^3$. The intersection between a sphere and the 3D Voronoi faces represents a new way to visualize the voids in the distribution of galaxies, see Section~\\ref{faces}. In this spherical cut the intersection between a sphere and the 3D Voronoi faces is no longer represented by straight lines but by curved lines presenting in some cases a cusp behavior at the intersection, see Figure~\\ref{aitof_sphere}. In line of principle the spatial distribution of galaxies at a given redshift should follow such curved lines. \\noindent {\\bf Statistics of the voids} The statistical properties of the voids between galaxies can be well described by the volume distribution of the Voronoi Polyhedra. Here two distributions of probability were carefully compared: the old Kiang function here parametrized as a function of the dimension $d$ , see formula~(\\ref{kiang}), and the new distribution of Ferenc ~\\&~ Neda~ , see formula~(\\ref{rumeni}), which is a function of the selected dimension $d$. The probability of having voids as large as the Eridanus super-void was computed, see Section~\\ref{eridanus}. \\noindent {\\bf Simulations of the catalogs of galaxies} By combining the photometric dependence in the number of galaxies as a function of the redshift with the intersection between a slice and the Voronoi faces, it is possible to simulate the astronomical catalogs such as the 2dFGRS, see Section~\\ref{cat2dFGRS}. Other catalogs such as the RC3 which covers all the sky ( except the Zone of Avoidance ) can be simulated through a given number of spherical cuts, for example 25, with progressive increasing redshift. This simulation is visible in Figure~\\ref{mix_rc3} in which the theoretical influence of the Zone of Avoidance has been inserted, and in Figure~\\ref{noavoid_rc3} in which the theoretical RC3 without the Zone of Avoidance has been modeled. Figure~\\ref{rc3_radio} reports the subset of the galaxies which are radiogalaxies. \\noindent {\\bf Correlation function} The standard behavior of the correlation function for galaxies in the short range $[0-10~Mpc/h]$ can be simulated once 12 Poissonian seeds are inserted in a box of volume $( 96.24~Mpc/h )^3$ . In this case the model can be refined by introducing the concept of galaxies generated in a thick face belonging to the Voronoi Polyhedron. The behavior of the correlation function in the large range $[40-200~Mpc/h]$ of the Voronoi simulations of the 2dFGRS presents minimum variations from the processed astronomical data, see Figure ~\\ref{correlation_due}. We now extract a question from the conclusions of \\citet{Martinez2009} ``Third, the minimum in the large-distance correlation functions of some samples demands explanation: is it really the signature of voids?'' Our answer is ``yes''. The minimum in the large scale correlation function is due to the combined effect of the large empty space between galaxies ( the voids ) and to the photometric behavior of the number of galaxies as a function of the red-shift. \\appendix \\setcounter{equation}{0} \\renewcommand{\\theequation}{\\thesection.\\arabic{equation}}" }, "1003/1003.2567_arXiv.txt": { "abstract": "We present results of long-slit spectroscopy in several slit positions that cover different morphological structures of the central parts of three bright Galactic \\hii\\ regions: M8, M17 and \\ngc. We study the spatial distributions of a large number of nebular parameters such as the extinction coefficient, line fluxes, physical conditions and ionic abundances at the maximum spatial resolution attainable with our instrumentation. Particularly, our goal is to study the behaviour of the abundance discrepancy factor of \\ioni{O}{2+}, \\adfo, defined as the logarithmic difference of the \\ioni{O}{2+} abundances derived from collisionally excited and recombination lines. We find that the \\adfo\\ remains fairly constant along the slit positions of M8 and M17. In the case of \\ngc, we only detect the \\ion{O}{ii} recombination lines in the integrated spectrum along the whole slit, where the \\adfo\\ reaches a remarkably high value of about 0.59 dex. We compare our results with previous ones obtained for the Orion Nebula. We find several evidences that suggest the presence of a candidate to Herbig-Haro object in M8. ", "introduction": "\\label{intro} The study of the elemental abundances in \\hii\\ regions is an essential tool for our knowledge of the chemical evolution of the universe. Traditionally, ionic abundances relative of the elements heavier than He have been determined from the strong collisionally excited lines (CELs). More than 20 years ago, \\citet{french83} obtained the first determination of the \\ioni{C}{2+}/\\ioni{H}{+} ratio derived from the faint recombination line (RL) \\ion{C}{ii} 4267 \\AA\\ for a planetary nebula (PN), finding that it was several orders of magnitude larger than the abundance obtained from the CELs of this ion. Later, this result was confirmed in other PNe \\citep[$e.g.$][]{rolastasinska94,mathisliu99}. A similar qualitative result was also found by \\cite{peimbertetal93} for the \\ioni{O}{2+}/\\ioni{H}{+} ratio in the Orion Nebula: the abundances obtained from the flux of the faint RLs were higher than those derived using the standard method based on CELs. Currently, this observational fact is a classical problem in the understanding of the physics of photoionized nebulae known as ``Abundance Discrepancy\" (AD) problem. This disagreement is quantified by means of the Abundance Discrepancy Factor (ADF) that can be defined as the ratio, or the logarithmic difference, between abundances of a same ion derived from RLs and CELs. In the case of the \\ioni{O}{2+}/\\ioni{H}{+} ratio, the ADF has rather similar values between 0.1 and 0.3 dex for extragalactic and Galactic \\hii\\ regions \\citep[$e.g.$][]{garciarojasetal05,garciarojasesteban07,estebanetal09}, while for PNe the ADF shows a much wider range of values, becoming substantially larger in some object \\citep[$e.g.$][]{liuetal00,liuetal06,tsamisetal04,tsamisetal08}. What causes the AD problem is nowadays debated. On the one hand, the predictions of the temperature fluctuations paradigm proposed by \\cite{peimbert67}, and characterized by the mean square of the spatial distribution of temperature --the so-called temperature fluctuations parameter, \\tf-- seems to explain the ADF observed in \\hii\\ regions, as it is argued by \\cite{garciarojasesteban07}. Under this scheme, the AD problem is a direct consequence of the different temperature dependence of the emissivities of the lines used: in the case of CELs, it depends exponentially on the electron temperature, \\te, of the ionized gas, while the emissivity of RLs has a power law temperature dependence, similar to those of the Balmer lines used as reference to determine the ionic abundance ratio relative to \\ioni{H}{+}. On the other hand, the hypothesis suggested by \\cite{liuetal00}, where most of the emission of RLs come from a cold hydrogen-poor component immersed in the ambient gas which emits the bulk of CELs, seems to solve the AD problem in PNe with high ADF values. This hypothesis is based on the observed fact that certain PNe contain well resolved H-deficient knots which are strong metallic RL emitters \\citep[$e.g.$ Abell 30,][]{harringtonfeibelman84}. The existence and origin of the temperature fluctuations is controversial because high values of the \\tf\\ parameter are not reproduced by standard photoionization models \\citep{kingdonferland95,rodriguezgarciarojas10} and additional mechanisms are proposed in order to explain the presence of temperature fluctuations \\citep[see revisions of][]% {esteban02,peimbertapeimbert06}. In the same way, new scenarios try to find a solution to the AD problem in \\hii\\ regions putting forward new physical mechanisms. This is the case of the hypothesis presented by \\cite{tsamispequignot05} and \\cite{stasinskaetal07}. Based on the chemical model for the heavy-elements mixing of \\cite{tenoriotagle96}, these authors proposed the presence of two components of different chemical composition and physical conditions in \\hii\\ regions. The component responsible of most of the emission of RLs consists of cold metal-rich droplets from supernova ejecta still not mixed with the ambient gas of the \\hii\\ region where most of the CEL emission would be produced. Then, note that assuming a chemically inhomogeneous model in \\hii\\ regions in order to explain the AD problem, the abundance derived from RLs and CELs would be upper and lower limits, respectively, of the real abundance of the ionized gas \\citep{stasinskaetal07}. Recently, a new scenario have been proposed by \\cite{ercolano09} based on the existence of high-density quasi-neutral clumps --embedded in the nebular gas ionized by the extreme-ultraviolet (EUV) radiation-- which are ionized mainly by the X-ray emission from the central star. Under this scheme, the CEL emission mainly comes from the region ionized by the EUV radiation (E region), while the RLs are emitted in different proportions from the clumps (X region) and the E region. In this sense, the abundances of the E region would be representative of the nebula and those from CELs would be easier to correct than RL ones. Contrary to the model proposed by \\cite{tsamispequignot05} and \\cite{stasinskaetal07}, the nebular model of \\cite{ercolano09} has homogeneous abundances. In a previous paper \\citep{mesadelgadoetal08} we explored the behaviour of the AD at small spatial scales and its dependence on different nebular parameters and physical conditions in the Orion Nebula in order to shed light on the origin of the AD problem. In that study, we used long-slit spectroscopy at spatial scales of 1\\farcs2, finding high \\adfo\\ values related to the presence of Herbig-Haro (HH) objects and temperature spikes at the position of the protoplanetary discs (proplyds). A subsequent detailed analysis of HH~202 was carried out by \\cite{mesadelgadoetal09a} using integral field spectroscopy confirming a high \\adfo\\ value at the main knot of HH~202 in agreement with the results of \\cite{mesadelgadoetal08}. Another important result of \\cite{mesadelgadoetal09a} was the obtention --for the first time in an \\hii\\ region-- of a map of the Balmer temperature and of the temperature fluctuations in the observed field, finding no correlation between the \\adfo\\ and the \\tf\\ parameter. Following the same goals and methodology of \\cite{mesadelgadoetal08}, in this paper we have used long-slit spectroscopy at intermediate spectral resolution in order to study the spatial distribution of the \\adfo\\ and other main nebular properties as well as their relation with the local morphological structures ($e.g.$ density condensations, ionization fronts or HH objects) in other bright Galactic \\hii\\ regions, namely, M8, M17 and \\ngc. After the Orion Nebula, M8 and M17 are probably the most studied Galactic \\hii\\ regions. M8 forms a blister of photoionized material on the surface of a giant molecular cloud. Near the optical center the region with the highest surface brightness of the nebula is found, the Hourglass (HG) region. This region is mainly ionized by the O star Herschel~36 (Her~36), while the stars HD~165052 and 9~Sgr ionize the rest of the nebula \\citep{woodwardetal86}. M17 is a cavity with a V shape and the V opening in the line of sight. The main ionization source of M17 is a group of O3-O4 stars, which belong to the open cluster NGC~6618, located in the dark bay of the nebula \\citep{hansonconti95}. A singular characteristic of M17 is its rather high ionization degree --\\ioni{O}{2+}/\\ioni{O}{+} ratio-- in comparasion with other Galactic \\hii\\ regions. The chemical composition of M8 and M17 have been widely studied by several authors in all spectral ranges from low to high spectral resolution \\citep[$e.g.$][]{rubin69,peimbertcostero69,sanchezpeimbert91,% rodriguez99b,estebanetal98,tsamisetal03}. Based on high resolution and deep echelle spectrophotometry, \\cite{garciarojasetal07} have provided a complete revision of the chemical abundances of these \\hii\\ regions using CELs and RLs. Nevertheless, we have not found in the literature detailed studies about the spatial behaviour of the nebular properties of these regions, excluding that of \\cite{peimbertetal92} along of 17 areas of M17 at low spectral resolution. Our third region, \\ngc, is not a classical \\hii\\ region. This nebula is an interstellar bubble formed by the interaction of the stellar wind of the O6.5~IIIf star BD$+$60~2522 with the surrounding interstellar medium. Recently works assume that the ram pressure of the stellar wind is balanced by the surrounding gas pressure due to the similarity between the velocities of the the molecular cloud and the bright nebulosity places of the nebula \\citep{christopoulouetal95,mooreetal02b}. The study of physical conditions and chemical abundances in \\ngc\\ have been restricted to some selected zones \\citep{talentdufour79,% rodriguez99b,mooreetal02b}. \\cite{mooreetal02b} have been the first in exploring the spatial distributions of several bright emission lines --[\\ion{O}{iii}] 5007 \\AA, H$\\alpha$ and [\\ion{N}{ii}] 6584 \\AA-- along the set of knots located at northwest of the central star and the rim of the bubble. These authors also obtained the first density spatial profile along the slit position that covered the knots. \\begin{figure*} \\centering \\includegraphics[scale=0.7]{f1.ps} \\caption[f1]{Observed slit positions over the central parts of the three Galactic \\hii\\ regions. The slits only show the total extraction area, while the total slit length is 3\\farcm8. The positions of the candidates to proplyd \\citep{stecklumetal98} and HH object (see \\S\\ref{newhh}) are indicated as UC and HH?, respectively. The images are combinations of exposures taken with different narrow-band filters. For all regions, emission from [\\ion{O}{iii}] is shown in blue, emission from H$\\alpha$ is shown in green and emission from [\\ion{S}{ii}] is shown in red. M8 and \\ngc\\ images are sections of combinations of WFPC2 images --\\cite{caulet97} and \\cite{mooreetal02b}, respectively. The M17 image is a section of the original mosaic obtained by the amateur astronomer I. de la Cueva Torregrosa.} \\label{pos} \\end{figure*} In Section~\\ref{obsred} we describe the observations of the Galactic \\hii\\ regions, the reduction procedure and the extraction of the one-dimensional spectra. In Section~\\ref{lir} we enumerate the selected emission lines and describe the procedure used to measure the fluxes and the extinction correction applied for each nebula. In Section~\\ref{phyab} we describe the method used to determine the physical conditions and the ionic abundances from both kinds of lines, CELs and RLs. In Section~\\ref{spatial} we present and discuss the spatial distributions along the slit positions of several nebular quantities for each nebula. In Section~\\ref{rlngc} we show the physical conditions and the ionic abundances for the individual extractions of \\ngc\\ as well as discuss the puzzling abundance pattern found in this object. In Section~\\ref{discu} we compare the results on the light of the different lineal spatial resolution attained in this paper with that of the observations of \\cite{mesadelgadoetal08} in the case of the Orion Nebula. We also present several arguments for the presence of a new HH object in M8 and discuss the possible causes of the high \\adfo\\ found in \\ngc. Finally, in Section~\\ref{conclu} we summarize the main conclusions of the paper. ", "conclusions": "\\label{conclu} In this article, we have carried out long-slit spectrophotometry at intermediate spectral resolution of the Galactic \\hii\\ regions M8, M17 and \\ngc. The one-dimensional spectra were extracted with a resolution of 1\\farcs2 for M8 and M17, and 3\\arcsec\\ in the case of \\ngc. Additional extractions with a spatial size of 4\\farcs8 were necessary in order to measure the faint \\ion{C}{ii} and \\ion{O}{ii} RLs in M8. We have studied the spatial distributions of a large number of nebular quantities along several slit positions covering different morphological structures such as HH objects, ionization fronts or bright knots. The studied quantities were \\chb, \\nel, \\te([\\ion{O}{iii}]), \\te([\\ion{N}{ii}]), the observed and dereddened flux of several emission lines ([\\ion{Fe}{iii}] 4658 \\AA, \\ion{C}{ii} 4267 \\AA, \\ion{O}{ii} 4649 \\AA, [\\ion{O}{iii}] 4959,5007 \\AA\\ and [\\ion{O}{ii}] 7320 \\AA) and the \\ioni{O}{2+} abundances derived from CELs and RLs, as well as the difference of both determinations, \\adfo. The density spatial distributions show a large range of variation across the different slit positions. We have found local maxima associated with the HG region, HH~870, knots and regions with a high surface brightness. On the other hand, the temperature spatial profiles do not show important variations related to the cited structures. The temperatures obtained from the different indicators present the classical behaviour for M8 and M17: those derived from [\\ion{N}{ii}] lines are higher than those derived from [\\ion{O}{iii}] lines, as expected for ionization-bounded nebulae. In the case of \\ngc, both temperatures seem to be very similar considering the error due to the own structure of the Bubble nebula which is matter-bounded. We have also explored the spatial behaviour of the \\adfo\\ along the slit positions of M8 and M17, which remains rather constant, finding values between 0.3 and 0.5 dex with an average error of about 0.1 dex. We have analysed the physical conditions and chemical composition of four additional extractions of \\ngc: three of them centered on the knots --K1 and K2-- and the rim of the bubble; and the last one corresponding to the ``whole slit\" spectrum. On the one hand, the comparison of the \\ioni{O}{2+}/\\ioni{H}{+} ratio determined from CELs and RLs in the ``whole slit\" spectrum produces an \\adfo\\ of about 0.59 dex. Assuming that the AD problem is related to temperature fluctuations, we have obtained a \\tf\\ parameter of 0.071. We have found a puzzling pattern in the total abundances derived for \\ngc. The total abundances obtained for the knots are slightly higher than those of the rim. The total abundances were compared with those expected by the Galactic abundance gradients, finding that there are discrepancies, specially in the case of C. We suspect that \\ion{C}{ii} 4267 \\AA\\ RL may be abnormally enhanced in \\ngc\\ due to an unknown physical process, whose investigation is outside the scope of this paper but deserves further more detailed observations. Comparing our observations with those of \\cite{mesadelgadoetal08}, we conclude that proplyds --and the associated variations of the local properties of the gas-- with linear sizes similar to those found in the Orion Nebula cannot be resolved with the observations reported in this paper. Angular resolutions of the order or smaller than the minimun seeing reached from the ground-based telescope would be needed to distinguish the presence of proplyds. Finally, we have found several evidences that point out to a possible new candidate to HH object in M8. This new object is located 16\\arcsec\\ east and 44\\arcsec\\ north from Her~36 where we have found enhancements in the spatial profile of [\\ion{Fe}{iii}] 4658 \\AA\\ and in [\\ion{S}{ii}]/H$\\alpha$ map presented by \\cite{ariasetal06}." }, "1003/1003.6108_arXiv.txt": { "abstract": "The {\\it Planck} satellite, successfully launched on May 14th 2009 to measure with unprecedented accuracy the primary Cosmic Microwave Background (CMB) anisotropies, is operating as expected. The Standard Model of the Universe (``concordance'' model) provides the current realistic context to analyze the CMB and other cosmological/astrophysical data, inflation in the early Universe being part of it. The {\\it Planck} performance for the crucial primordial parameter $r$, the tensor--to--scalar ratio related to primordial $B$ mode polarization, will depend on the quality of data analysis and interpretation. The Ginzburg-Landau approach to inflation allows to take high benefit of the CMB data. The fourth degree double well inflaton potential gives an excellent fit to the current CMB+LSS data. We evaluate the {\\it Planck} precision to the recovery of cosmological parameters, taking into account a reasonable toy model for residuals of systematic effects of instrumental and astrophysical origin based on publicly available information. We use and test two relevant models: the $\\Lambda$CDM$r$ model, i.e. the standard $\\Lambda$CDM model augmented by $ r $, and the $\\Lambda$CDM$r$T model, where the scalar spectral index, $n_s$, and $r$ are related through the theoretical ``banana-shaped'' curve $ r = r(n_s)$ coming from the Ginzburg-Landau theory with double--well inflaton potential. In the latter case, the analytical expressions for $ n_s $ and $ r $ are imposed as a hard constraint in the Monte Carlo Markov Chain (MCMC) data analysis. We consider two $C_\\ell-$likelihoods (with and without $B$ modes) and take into account the white noise sensitivity of {\\it Planck} (LFI and HFI) in the 70, 100 and 143 GHz channels as well as the residuals from systematics errors and foregrounds. We also consider a cumulative channel of the three ones. We produce the sky (mock data) for the CMB multipoles $ C_l^{TT}, \\; C_l^{TE}, \\; C_l^{EE} $ and $ C_l^{BB} $ from the $\\Lambda$CDM$r$ and $\\Lambda$CDM$r$T models and obtain the cosmological parameter marginalized likelihood distributions for the two models. Foreground residuals turn to affect only the cosmological parameters sensitive to the $ B $ modes. As expected, the likelihood $ r $ distribution is much clearly peaked near the fiducial value ($ r = 0.0427 $) in the $\\Lambda$CDM$r$T model than in the $\\Lambda$CDM$r$ model. The best value for $ r $ in the presence of residuals turns to be about $ r \\simeq 0.04 $ for both the $\\Lambda$CDM$r$ and the $\\Lambda$CDM$r$T models. The $\\Lambda$CDM$r$T model turns to be very stable, its distributions do not change by including residuals and the $ B $ modes. For $ r $ we find $ 0.028 < r < 0.116 $ at 95 \\% CL with the best value $ r = 0.04 $. We also compute the $B$ mode detection probability by the most sensitive HFI-143 channel. At the level of foreground residual equal to 30\\% of our toy model only a 68\\% CL (one sigma) detection is very likely. For a 95\\% CL detection (two sigmas) the level of foreground residual should be reduced to 10\\% or lower of the adopted toy model. The lower bounds (and most probable value) we infer for $ r $ support the searching of CMB $B$ mode polarization in the current data as well as the planned CMB missions oriented to $B$ polarization. ", "introduction": "The {\\it Planck} satellite\\footnote{http://www.rssd.esa.int/planck} was successfully launched on May 14th 2009 to measure the primary Cosmic Microwave Background (CMB) temperature and polarization anisotropies on the whole sky with unprecedented accuracy. It is now in normal operation, with the expected performances \\citep{PlanckBlueBook,prelauB,prelauM,lamarre2010,Maffei_prelaunch}. {\\it Planck} will improve the measurement of most cosmological parameters by several factors with respect to current experiments, in particular the Wilkinson Microwave Anisotropy Probe (WMAP) satellite\\footnote{http://lambda.gsfc.nasa.gov/}. The expected CMB polarization measurements from {\\it Planck} will allow to push both ($E$ and $B$) polarization results well beyond the present knowledge and considerably constrain the tensor ($B$ modes) to scalar ratio parameter $ r $, if not to obtain a detection on it. In this respect, the way of extracting and physically interpreting cosmological parameters (once the CMB data cleaned from the different astrophysical foregrounds) will be important. In other words, the {\\it Planck} actual performance for the crucial primordial parameter $ r $ will depend on the adopted physical modeling and on the quality of data analysis and interpretation. It is then important and timely to make forecasts for the {\\it Planck} determination of $ r $ and other cosmological parameters taking into account the theoretical progress in the field and WMAP results. \\medskip The Standard Model of the Universe (or ``concordance'' model) provides the current realistic context to analyze the CMB and other cosmological/astrophysical data. Inflation (quasi-exponential accelerated expansion) of the early Universe is a part of this model and one important goal of CMB experiments is probing the physics of it. Inflation solves the shortcomings of the decelerated expanding cosmology (horizon problem, flatness, entropy of the Universe), and explains the observed CMB anisotropies providing the mechanism for the generation of scalar and tensor perturbations seeding the large scale structures (LSS) and primordial (still undetected) gravitational waves ($B$ mode polarization). \\medskip The current CMB $+$ LSS data support the standard inflationary predictions of a nearly spatially flat Universe with adiabatic and nearly scale invariant initial density perturbations. These data are validating the single field slow-roll inflationary scenario \\citep{WMAP5}. Single field slow-roll models provide an appealing, simple and fairly generic description of inflation \\citep{libros,reviu}. The inflationary scenario is implemented using a scalar field, the \\emph{inflaton} with a potential $ V(\\varphi) $, self-consistently coupled to the space-time metric. In the effective theory based on the Ginzburg-Landau (G-L) approach to inflation \\citep{reviu}, the potential is a polynomial in the field starting by a constant term. Linear terms can always be eliminated by a constant shift of the inflaton field. The mass (quadratic) term can have a positive or a negative sign associated to unbroken symmetry (chaotic inflation) or to broken symmetry (new inflation), respectively. The fourth degree double--well inflaton potential gives an excellent fit of the present CMB $+$ LSS data \\citep{reviu}. A cubic term does not improve the fit and can be omitted \\citep{mcmc1}. Adding higher order terms with additional parameters does not improve significantly the fits \\citep{high}. The G-L framework is not just a class of physically well motivated inflaton potentials, among them the double and single well potentials. This approach provides the effective theory for inflation, with powerful gain in the physical insight and analysis of the data. The present set of data with the effective theory of inflation favor the double well potential \\citep{reviu,mcmc1}. Analyzing the present data without the relation between $r$ and $n_s$ does not allow to discriminate among different classes of models for the inflaton potential in the considered framework. Although the G-L effective theory approach to inflation is quite general, it predicts precise order of magnitude estimates for $ n_s $, $ r $ and the running of the spectral index $ dn_s/d \\ln k $ \\citep{reviu} $$ n_s - 1 = {\\cal O}\\left(\\frac1{N}\\right) \\; , \\quad r = {\\cal O}\\left(\\frac1{N}\\right) \\; , \\quad \\frac{d n_s}{d \\ln k} = {\\cal O}\\left(\\frac1{N^2}\\right) \\; ; $$ here $ N \\sim 60 $ is the number of efolds since the cosmologically relevant modes exit the horizon till inflation ends. The WMAP values for $ n_s $ and the upper bounds for $ r $ and $ dn_s/d \\ln k $ agree with these estimates. Since in this framework the estimated running, $ dn_s/d \\ln k \\sim 3 \\times 10^{-4} $, is very small, in this paper we will concentrate on $ n_s $ and $ r $. \\medskip In this work, we evaluate the accuracy in the recovery of the cosmological parameters expected from the {\\it Planck} data. First, we do this forecast without including the systematic effects of instrumental and/or astrophysical origin or their coupling, affecting the {\\it Planck} measurements, and then by including the systematic effects. In this study we exploit the {\\it Planck} sensitivity and resolution at its three favorite cosmological channels, i.e. at the frequencies of 70, 100, and 143 GHz. Table \\ref{table:sens} reports the {\\it Planck} performance at these frequencies, based on \\citet{PlanckBlueBook} and, for the LFI channel at 70 GHz, as updated in \\citet{prelauM}, \\citet{prelauB}, \\citet{sandri_etal_2010}. These sensitivities do not include the degradation in accuracy that could come from various sources of systematic effects, of both instrumental and/or astrophysical origin, or their coupling. In Sect. \\ref{toym} we discuss the current published estimates for the residuals of systematic effects and foregrounds affecting the {\\it Planck} CMB measurements: straylight, main beam asymmetry, leakage, time constants, glitches, and foregrounds. In general, we do not use in this work a precise (still not completely available) description of the considered systematic effects, but only suitable representations of them, as described in Sect. \\ref{toym}. This is done in a parametric approach, identifying the corresponding levels at which the control of the systematic effects is necessary not to spoil the {\\it Planck} data scientific accuracy. We technically implement this rescaling with a multiplicative constant on the residuals of the systematic effects on the CMB multipoles $ C_\\ell $. Obviously, the real analysis of {\\it Planck} data will have to properly consider all possible systematic effects of optical, thermal, and instrumental (radiometric and bolometric) origin, with an even better accuracy than those achieved in past projects. In parallel, a significantly improved separation of CMB from astrophysical components will be needed, a task in principle possible for {\\it Planck} thanks to its wide frequency coverage. \\medskip Instrumental systematics on CMB tensors-to-scalar have been studied by \\citet{whu,shimon,yadav}. \\medskip We use and test two relevant models: the $\\Lambda$CDM$r$ model, that is the standard $\\Lambda$CDM model augmented by the tensor--to--scalar ratio $ r $, and the $\\Lambda$CDM$r$T model, that is the $\\Lambda$CDM$r$ model in which the double--well inflaton potential (see Eq. (\\ref{binon}) in the next section) is imposed. Namely, $ n_s $ and $ r $ are constrained by the analytic relation $ r = r(n_s) $ to lay on the theoretical banana-shaped curve (the upper border of the banana-shaped region Fig. \\ref{banana}). The novelty in the MCMC analysis of the CMB data with the $\\Lambda$CDM$r$T model is in the fact that we impose the analytical expressions for $ n_s $ and $ r $ derived from the inflaton potential as a hard constraint \\citep{mcmc1}. We take both models, $\\Lambda$CDM$r$ and $\\Lambda$CDM$r$T, as fiducial models in our Monte Carlo Markov Chains (MCMC) simulations to produce the corresponding skies (mock data). In the $\\Lambda$CDM$r$ model the independent cosmological parameters are $\\Omega_b \\, h^2, \\; \\Omega_c \\, h^2, \\; \\theta , \\; \\tau, \\; A_s , \\; n_s $ and $ r $, while all other independent parameters are assumed to vanish, {\\em e.g.} $ \\Omega_\\nu=0 $, or have the standard values, {\\em e.g.} $ w=-1 $. The aforementioned $\\Lambda$CDM$r$T model includes the same parameters but with $ n_s $ and $ r $ not being independent, but related by the curve $r=r(n_s)$ as widely discussed in Sect. \\ref{LGtheory}. We produce one sky (mock data) for the anisotropy CMB multipoles $ C_l^{TT}, \\; C_l^{TE}, \\; C_l^{EE} $ and $ C_l^{BB} $ from the $\\Lambda$CDM$r$ model and from the $\\Lambda$CDM$r$T model, with the parameters in Table \\ref{tab2}. We describe the detailed procedure in Sect. \\ref{mock}. We run Monte Carlo Markov Chains from this sky and obtain the marginalized likelihood distributions for the cosmological parameters ($ \\Omega_b \\, h^2 , \\; \\Omega_c \\, h^2 ,\\; \\theta \\; \\tau, \\; \\Omega_{\\Lambda} $, Age of the Universe, $ z_{re}, \\; H_0, \\; A_s, \\; n_s $ and $ r $) in the two test models $\\Lambda$CDM$r$ and $\\Lambda$CDM$r$T . We study the independent $\\Lambda$CDM$r$ parameters with the mock data produced from $\\Lambda$CDM (first row of Table \\ref{tab2}) and the independent parameters of both $\\Lambda$CDM$r$ and $\\Lambda$CDM$r$T with the mock data produced from $\\Lambda$CDM$r$T (second row in Table \\ref{tab2}). The fiducial values, $ r = 0.0427 $ and $ n_s = 0.9614 $ correspond to the best fit to the CMB-LSS data with the $\\Lambda$CDM$r$T model using the double--well inflaton potential expressed by Eq. (\\ref{binon}). Namely, these are the best fit values to $ r $ and $ n_s $ within the Ginsburg-Landau effective theory approach. Not using the Ginsburg-Landau approach, lower bounds for $ r $ are not obtained and the best fit value for $ r $ can be much smaller than $ r = 0.04 $ \\citep{kinney08,PeirisEasther}. We consider two choices for the $C_\\ell-$likelihood, one without the $ B $ modes and one with the $B$ modes and take into account the white noise sensitivity of {\\it Planck} (LFI and HFI) in the 70, 100 and 143 GHz channels \\citep{PlanckBlueBook}. We also consider a cumulative channel whose $ \\chi^2 $ is the sum of the $ \\chi^2 $'s of the three channels above. When using different channels in the MCMC analysis, we use different noise realizations while keeping the same sky, that is the same realization of the Gaussian process that generated the primordial fluctuations. In our MCMC analysis we always take standard flat priors for the cosmological parameters. In particular we assume the flat priors $ 0 \\leq r < 0.2 $ in the $\\Lambda$CDM$r$ model and $ 0 \\leq r< 8/60$, where $8/60 \\simeq 0.133$ is the theoretical upper limit for $ r $ in the $\\Lambda$CDM$r$T model. We performed the MCMC simulations using the publicly available CosmoMC code\\footnote{http://cosmologist.info/cosmomc/} \\citep{mcmc} interfaced to the Boltzmann code CAMB\\footnote{http://camb.info/} (see \\citet{2000ApJ...538..473L} and references therein). \\medskip Our findings without including the systematic effects are summarized in Figs. \\ref{r0r04}-\\ref{ban_y} where the marginalized likelihood distributions of the cosmological parameters are plotted for several different setups. In Tables \\ref{tab3} to \\ref{tab5} we list the corresponding relevant numerical values. Clearly, in the case of the ratio $ r $, due to the specific form of its likelihood distribution, it is more interesting to exhibit upper and lower bounds rather than mean values and standard deviations as in Tables \\ref{tab3} and \\ref{tab4}. We report the upper bounds and, when present, the lower bounds in Tables \\ref{tab5} and \\ref{tab6}. Our conclusions without including the systematic effects are: \\begin{itemize} \\item{The upper bound on $ r $ and the best value of $ n_s $ do not require to include the $ B $ modes in the likelihood, and can be obtained with the $\\Lambda$CDM$r$ model alone, (i.e. $ r < 0.068 $ and $ n_s = 0.9549 $ at 95 \\% CL ). See Tables \\ref{tab3}, \\ref{tab5} and Fig. \\ref{r0r04}. The inclusion of $ B $ modes, for a non vanishing fiducial value, ($ r = 0.0427 $), allows peaked marginalized distributions for $ r $ and a lower bound for $ r $. See Table \\ref{tab6} and Figs. \\ref{r0r04}, \\ref{br0br04} and \\ref{r04br04T}. We obtain $ 0.013 < r < 0.045 $ at 95 \\% CL in the $\\Lambda$CDM$r$ model, with the best values $r = 0.0240, n_s = 0.9597 $. This shows a substantial progress in the forecasted bounds for $ r $ with respect to the WMAP+LSS data set for which $ r < 0.20 $ in the pure $\\Lambda$CDM$r$ model \\citep{WMAP5,WMAP7}.} \\item{Lower bounds on $ r $ and most probable $ r $ values are always obtained (with or without the $ B $ modes) with the $\\Lambda$CDM$r$T model. See Tables \\ref{tab3}, \\ref{tab4}, \\ref{tab6} and Fig. \\ref{r04br04T}. The $\\Lambda$CDM$r$T model provide well peaked distributions for $ r $ on nonzero values $ r \\simeq 0.04 $. We obtain $ r > 0.039 $ at 68 \\% CL and $ r > 0.030 $ at 95 \\% CL in the $\\Lambda$CDM$r$T model.} \\end{itemize} In Sect. \\ref{fortoy} we include in the forecasts the systematic effects discussed in Sects. \\ref{sensit} and \\ref{toym}. Our conclusions including the systematic effects and foreground residuals are: \\begin{itemize} \\item{The likelihood distributions with and without $ B $ modes result almost the same when including the residuals. Only the cosmological parameters sensitive to the $B$ modes appear to be affected by the residuals, namely, $ \\tau, \\; z_{re} $ and $ r $. The main numbers are displayed in Tables \\ref{tab5} and \\ref{tab6}.} \\item{The marginalized likelihood $ r $ distribution for fiducial ratio $ r = 0.0427 $ is much clearly peaked on a value of $ r $ near the fiducial one in the $\\Lambda$CDM$r$T model than in the $\\Lambda$CDM$r$ model (compare Figs. \\ref{resbr0br04} and \\ref{resbr04T}). In any case, the best value for $ r $ in the presence of residuals is about $ r \\simeq 0.04 $ (near the fiducial value) both for the $\\Lambda$CDM$r$ and the $\\Lambda$CDM$r$T models. The $\\Lambda$CDM$r$T model turns to be robust, it is very stable (its distributions do not change) with respect to the inclusion of residuals (and they do not change neither with respect to the inclusion of $B$ modes). The main numbers are included in Tables \\ref{tab5} and \\ref{tab6}. With the $\\Lambda$CDM$r$T model we have for $ r $ at 95 \\% CL: $$ 0.028 < r < 0.116 \\quad {\\rm with ~the ~best ~values} \\quad r = 0.04 \\quad n_s = 0.9608 \\; . $$} \\end{itemize} \\medskip It must be stressed that, in the $\\Lambda$CDM$r$T model, future improvements in the precision $ \\delta $ on the measured value of $ n_s $ alone will immediately give an improvement $ dr/dn_s \\; \\delta $ on the prediction for $ r $ as well as for its lower bound. Better measurements for $ n_s $ will thus improve the prediction on $ r $ from the $T$, $TE$ and $E$ modes even if a secure detection of $B$ modes will be still lacking. \\medskip In order to assess the probability for {\\it Planck} to detect $ r $ we also compute the $B$ mode detection probability by the most sensitive HFI-143 channel; this is done in Sect. \\ref{143}. We extract $10^5$ skies obtaining the corresponding multipoles $ A_{lm} $ from the $\\Lambda$CDM$r$T model according to the procedure described in Sect. \\ref{mock}, adopting $ r=0.0427 $ as fiducial value. We compute all the corresponding likelihood profiles only for $ r $ and their interesting properties, like the most likely value $ r_{max} $, the mean value $ r_{mean} $, the standard deviation $ \\Delta r_{\\max} $ of the $ r_{max} $ distributions, the skewness and the kurtosis, (which measures the departure from a Gaussian likelihood), Fig. \\ref{probd}. We finally compute the 99\\% CL, 95\\% CL, and 68\\% CL lower bounds for $ r $. The probabilities of detection of $ r $ are displayed in Fig. \\ref{probd2}. At the level of foreground residual equal to 30\\% of the considered toy model, only a 68\\% CL (one sigma) detection is very likely. For a 95\\% CL detection (two sigmas) the level of foreground residual should be reduced to 10\\% of the considered toy model, or lower. Lensing acts on the B-modes as a contamination by transforming E-modes into B-modes. It is a frequency independent effect while residuals are frequency dependent. Lensing weakens the signal around $ \\ell \\sim 90 $ where the primordial B-modes peak but not in the small $ \\ell $ modes range where the reionization bump dominates. On the other hand foreground residuals are larger at small $ \\ell $ than at $ \\ell \\sim 90 $. Namely, residuals and lensing affect the detection of B-modes in complementary ways, with the effect of residuals stronger than that of lensing. As a consequence, lensing plus residuals can spoil the detection of $ r $ even when residuals are assumed at the 30\\% level of the considered toy model. On the contrary, lensing in the absence of residuals still allows a detection of $ r $. For example, several MCMC simulations show that our lower bounds on $ r $ are not significantly affected by lensing in the absence of foreground residuals. Let us make clear, at any rate, that lensing was not considered in the analysis of the $r-$detection probability in Sect. \\ref{143}. Finally, it should be clear that if the theoretical constraint $ r = r(n_s) $ of the $\\Lambda$CDM$r$T model is imposed on the MCMC analysis, $ r $ has always well defined lower bounds regardless of lensing and/or residuals. The forecasted probability of detecting $ r $ is based on the statistics of the shape of the $r$ -likelihood. This shape determines whether a detection of $ r $ can be claimed with a given confidence level. But real CMB experiments can observe only one sample: the observed sky. So, the possibility of inferring $ r $ from one single (albeit very large) sample depends on the sample itself, and therefore, whether $ r $ will be or will be not detected depends also of a question on luck. In addition, the results for many skies presented in Sect. \\ref{143} show the consistency of our whole approach to determine $ r $. \\medskip Finally, in Sect. \\ref{efbias} we consider the bias effect in the foreground residuals implemented as a linear perturbation affecting the $ C_l'$s and explore how the cosmological parameter distributions are affected by the bias. We implement two extreme cases: in case (i) the bias fluctuates randomly around zero and in case (ii) the bias fluctuates around a non-zero value, staying significantly non-zero. In case (i) the cosmological parameters are practically unaffected while in case (ii) the peaks of the cosmological parameter distributions are shifted within one or two sigmas of the WMAP values. In particular, $ r $ is not anymore detected in case (ii). \\medskip The best and mean values reported here for $ r $ and the other cosmological parameters do not correspond to the true sky data but to mock skies generated from the MCMC simulations as explained above. Nevertheless, the deviations between the best and the fiducial values are relevant indicators for $ r $ as well as the lower and upper bounds and the standard deviation. The fact that the fiducial and mean values of $ r $ are very close and that $ \\Delta r_{\\max} $ coincides with the mean value of the standard deviation of $ r $ indicate that {\\it Planck} can provide detections of high quality. More in general, our results support the quest for $B$ mode polarization in the current CMB data and future $B$ oriented polarization missions under study by both ESA\\footnote{http://www.b-pol.org/index.php} and NASA\\footnote{http://cmbpol.uchicago.edu/} \\citep{2009ExA....23....5D,2006AAS...209.4907B}. ", "conclusions": "In this paper we provide a precise forecast for the {\\it Planck} results on cosmological parameters, in particular for the tensor--to--scalar ratio $ r $. These new forecasts go far beyond the published ones (see e.g. \\citet{PlanckBlueBook}, \\citet{otros}) and pave the road for a promising scientific exploitation and interpretation of the {\\it Planck} data (once cleaned from the different astrophysical foregrounds). We appropriately combined the following, as main ingredients: the current public available knowledge of {\\it Planck} instrument sensitivity and a reasonable toy model estimation of the residuals from systematic errors and foregrounds; the highly predictive theory setup \\citep{reviu,mcmc1,mcmc2} provided by the Ginzburg-Landau approach to inflation to produce and analyze the skies (mock data) which allows a decisive gain in the physical insight and data analysis; precise MCMC methods to produce the skies (mock data) and to analyze them. This turns into an improvement in the physical analysis, in particular for the ratio $ r $. It must be also stressed that, in the considered framework, better measurements for $ n_s $ will improve the predictions on $ r $ from the $T$, $TE$ and $E$ modes even if a secure detection of $B$ modes will be still lacking. We remark also that the model is falsifiable in the case of constraints on $ n_s $ and $ r $ not compatible with the banana shape of the considered framework. The lower bounds and most probable value inferred from WMAP for $ r $ ($ r \\simeq 0.04 $) in the considered framework support the search for $B$ mode polarization in {\\it Planck} data and the future CMB $B$ oriented polarization missions." }, "1003/1003.4427_arXiv.txt": { "abstract": "{ The primary target in the seismo-field of \\emph{CoRoT}, HD 49434, known to be a bona fide hybrid $\\gamma$ Doradus/$\\delta$ Scuti, shows excited intermediate-order g modes. Time-Dependent Convection models, however, predict a range in frequency that is stable to pulsations, between the simultaneously excited high-order g modes ($\\gamma$ Dor) and low-order p and g modes ($\\delta$ Sct). Furthermore, theoretical studies based on model computations of $\\delta$ Sct stars suggest that stochastically excited modes are likely to be observed. A pertinent question would then be to ask: Might those observed intermediate-order g modes be stochastically excited? By employing a statistical method which searches for the signature of stochastic excitation in stellar pulsations, we investigate the nature of those modes with possible implications on the identification of their excitation mechanism. Preliminary results are rather inconclusive about the presence of stochastic excitation. A new analysis that thoroughly takes into account sampling effects is necessary in order to get more reliable results. } ", "introduction": "Hybrid $\\gamma$ Dor/$\\delta$ Sct stars are of great interest because they offer additional constraints on stellar structure and may be used to test theoretical models. $\\gamma$ Dor stars pulsate in high-order g modes with periods of the order of 1 day, driven by convective flux blocking at the base of their convective envelopes. In turn, $\\delta$ Sct stars pulsate in low-order p and g modes with periods of the order of 2 hours, driven by the $\\kappa$-mechanism operating in the \\ion{He}{ii} ionization zone. Most of the already confirmed hybrid stars clearly display separate frequency domains. However, some of these objects are known to exhibit a range of excited intermediate modes. Such a feature has been detected for the first time in the frequency power spectrum of the Am star HD 8801 by \\citet{Handler}. That work also shows that the amplitudes of these intermediate oscillation modes are of the same order of magnitude as of those detected in the classical $\\delta$ Sct domain. The presence of intermediate-order modes opens new questions concerning the possible mechanism responsible for their excitation. In this regard, the present work aims at providing insight into the possible excitation mechanism of such intermediate-order modes, through the analysis of the frequency power spectrum of the hybrid star HD 49434, which has been selected for the asteroseismic core programme of the \\emph{CoRoT} satellite \\citep{CoRoT}. HD 49434 (spectral type F1V) is a bright ($V\\!=\\!5.75$) and multiperiodic pulsator located near the blue edge of the $\\gamma$ Dor instability strip (IS) and inside the $\\delta$ Sct IS. Photometric data recently made available by \\emph{CoRoT} \\citep{Chapellier} allowed confirming the hybrid nature of HD 49434. This star had been referenced as a candidate hybrid $\\gamma$ Dor/$\\delta$ Sct by \\citet{Uytterhoeven}, following an extensive photometric and spectroscopic ground-based campaign. A compelling feature of its frequency power spectrum is the presence of excited intermediate-order g modes. From a theoretical point of view, different approaches are currently being adopted in order to explain such spectral features. Theoretical studies based on model computations of cool $\\delta$ Sct stars \\citep{Houdek,Samadi} predict the presence of solar-like oscillations with amplitudes large enough to be detectable with ground-based instruments, although so far not confirmed. It should be stressed here that these models lie exactly on or near the $\\gamma$ Dor IS determined afterwards by \\citet{Dupret}. Therefore, it is plausible considering that solar-like, $\\gamma$ Dor-like and $\\delta$ Sct-like oscillations might be simultaneously excited in such intermediate-mass, main-sequence stars. Rotational splitting has already been invoked as a possible explanation for these observed frequencies \\citep{Uytterhoeven,Bouabid}. More recently, \\citet{Kallinger} suggested that most of the peaks in the rich frequency spectra of the $\\delta$ Sct stars already analysed using data from \\emph{CoRoT} \\citep[see e.g.,][]{GH} could be the signature of non-white granulation background noise, meaning that only a few tens of those frequencies should be interpreted as stellar p modes. On the other hand, based on the Time-Dependent Convection (TDC) of \\citet{Grig05}, \\citet{Dupret} showed that TDC models predict the likely existence of hybrid stars with both $\\delta$ Sct p-mode and $\\gamma$ Dor g-mode oscillations. Furthermore, TDC models predict the existence of a frequency \\emph{gap} that is stable to pulsations in the range 5--15$\\:\\rm{d^{-1}}$ for low-degree modes. Therefore, according to these models, neither the classical $\\kappa$-mechanism nor the convective flux blocking at the bottom of the convective envelope can explain the excitation of observed intermediate-order g modes. In the present work, we address the possibility that such intermediate-order g modes are excited by a stochastic mechanism. In order to do so, a first approach to the problem can be easily implemented by considering the characteristics of this type of excitation, in particular those concerning the statistical behaviour of the mode amplitudes. A simple diagnostic method has been established by \\citet{Pereira05} that probes the stellar pulsations' excitation mechanism by analysing the temporal variation in the amplitude of oscillation modes. Numerical simulations and the application to the $\\gamma$ Dor star HD 22702 \\citep{Pereira07} serve as a test of the method. In this work we employ this statistical method to investigate the nature of the intermediate-order g modes visible in the spectrum of HD 49434. ", "conclusions": "Two remarks should be made after inspection of the excitation diagram (Fig.~\\ref{excitation}): \\begin{enumerate} \\item The way the sampling has been performed, i.e., the number of amplitude measurements and the type of sampling (linear with separation 1, in other words, sampling from contiguous subseries), does not allow for a good estimate of $\\sigma(A)/\\langle A \\rangle$ if we assume stochastic excitation. \\item The fact that the observational results lie outside the confidence interval (for the two strongest modes) or just on the lower $1\\sigma$ bound (for the two faintest modes) might tempt us to conclude that these modes are not stochastically excited. However, we need to be cautious and a new analysis should definitely be carried out where we increase the number of amplitude measurements (with the corresponding loss in frequency resolution of each of the subseries), thus increasing the significance of the statistic $\\sigma(A)/\\langle A \\rangle$. \\end{enumerate}" }, "1003/1003.3031_arXiv.txt": { "abstract": "We use the supernova remnants (SNRs) in the two Magellanic Clouds (MCs) as a supernova (SN) survey, ``conducted'' over tens of kyr, from which we derive the current SN rate, and the SN delay time distribution (DTD), i.e., the SN rate vs. time that would follow a hypothetical brief burst of a star formation. In a companion paper (Badenes, Maoz, \\& Draine 2010) we have compiled a list of 77 SNRs in the MCs, and argued that it is a fairly complete record of the SNRs that are now in the Sedov phase of their expansions. We recover the SN DTD by comparing the numbers of SNRs observed in small individual ``cells'' in these galaxies to the star-formation histories of each cell, as calculated from resolved stellar populations by Harris \\& Zaritsky. We identify the visibility times of SNRs in each cell with the Sedov-phase lifetimes, which depend on the local ambient densities. The local densities are estimated from 21-cm emission, from an inverse Schmidt-Kennicutt law based on either H$\\alpha$ emission or on the star-formation rate from the resolved stellar populations, and from combinations of these tracers. This is the first SN DTD that is based on resolved stellar populations. We detect a population of ``prompt'' type-Ia SNe (that explode within 330~Myr of star formation) at $>99\\%$ confidence level (c.l.). The best fit for the number of prompt type-Ia SNe per stellar mass formed is $2.7-11.0\\times 10^{-3}M_\\odot^{-1}$, depending on the density tracer used. The $95\\%$ c.l. range for a ``delayed'' (from $330$~Myr to a Hubble time) type-Ia component is $<1.6\\times 10^{-13}$~SN~yr$^{-1}M_\\odot^{-1}$, consistent with rate measurements in old populations. The current total (core-collapse+Ia) SN rate in the MCs is 2.5-4.6 SNe per millenium (68\\% c.l.+systematics), or 1.7-3.1 SNuM [SNe $(100~{\\rm yr}~10^{10}M_\\odot)^{-1}$], in agreement with the historical record and with rates measured in other dwarf irregulars. Conversely, assuming the SNRs are in free-expanion, rather than in their Sedov phase, would impose on the SNRs a maximum age of 6~kyr, and would imply a MC SN rate per unit mass that is 5 times higher than in any type of galaxy, and a low-mass limit for core-collapse progenitors in conflict with stellar evolution theory. ", "introduction": "Supernova (SN) explosions and their remnants touch upon multiple aspects of astrophysics and cosmology, whether as endpoints in stellar evolution, as sources of energy and enriched material to the interstellar and intergalactic media, as sites of cosmic-ray acceleration, or as standard candles for cosmography. However, much remains to be understood regarding these events, both the core-collapse (CC) SN explosions that are thought to end the lives of some or all massive ($\\ga 8M_\\odot$) stars; and the type-Ia SNe (SNe~Ia), that are believed to be the thermonuclear combustions of CO white dwarfs (WDs) that have approached the Chandrasekhar mass by accreting material from, or merging with, a companion star. The lower limit on the initial mass of stars that eventually undergo CC is poorly known, both theoretically \\citep{poelarends08:SuperAGB_SNe} and observationally \\citep[from identification of progenitors in pre-explosion images; e.g.][]{smartt09:CCSN_progenitors}, and could be as low as 7 $M_\\odot$ or as high as 12 $M_\\odot$. The upper mass limit on the intial mass that leads to a SN explosion, rather than (perhaps) to an explosion-less collapse into a black hole, is also uncertain theoretically \\citep{heger03:CCSNe} and observationally \\citep{kochanek08:survey_nothing}. Among the CC-SN progenitors, it is still not clear which mass range leads to which SN subtypes, such as IIP, IIn, Ib, and Ic \\citep{smartt09:CCSN_progenitors}. Parameters other than mass -- binarity, rotation, metallicity -- likely also play an important role in determining CC-SN type \\citep[e.g.][]{eldridge08:CCSN_Progenitors_Binaries}. For SNe~Ia, our ignorance is even greater. The progenitor systems of these events are still unknown, with two distinct models generally considered: the single-degenerate (SD) scenario in which the WD accretes material from a normal star \\citep{whelan73:SNI_SD,nomoto82:accretingWD}, and the double-degenerate (DD) model, where the WD merges with another WD \\citep{iben84:typeIsn,webbink84:DDWD_Ia_progenitors}. Both models suffer from theoretical and observational problems. For SD scenarios, it is unclear how the material accreted from the donor can burn quietly on the WD surface until the required mass (close to the Chandrasekhar limit) is reached \\citep{cassisi98:H-accreting-CO-WD}. The narrow H or He lines that are expected in nebular SN spectra for SD progenitors have never been found in normal Ia events \\citep{leonard07:H_nebular_Ia_spectra}, nor have the integrated X-ray emission from a population of accreting WDs undergoing slow nuclear burning on their surfaces \\citep{distefano10:sss_Ia_progenitors,gilfanov10:Ia_progenitors}. The claimed identification of the surviving donor star in the nearby Type Ia supernova remnant (SNR) Tycho \\citep{ruiz-lapuente04:Tycho_Binary,hernandez09:Tycho_G} has been recently questioned by \\citet{kerzendorf09:Tycho_G}. In the case of DD systems, there are concerns regarding the fate of merging WDs, which might lead to accretion induced collapse instead of a SN Ia explosion \\citep{saio85:DD_Mergers}. It is also unknown whether there are enough DD progenitor candidates in the Milky Way to produce the observed Ia SN rate \\citep{napiwotzki04:SPY_04,nelemans05:SPY_IV}, although current searches might clarify this point in the near future \\citep{badenes09:SWARMS_I}. Major progress in resolving these questions could come from knowledge of the elapsed times between the formation of a stellar population and the explosion of some of its members as different types of SNe. Indeed, a major objective of SN studies has been the recovery of the so-called delay-time distribution (DTD). The DTD is the SN rate as a function of time that would be observed following a $\\delta$-function burst of star formation. (In other contexts, the DTD would be called the delay function, the transfer function, or Green's function). Knowledge of the DTD would be useful for understanding the route along which cosmic metal enrichment and energy input by SNe proceed, but no less important, for obtaining clues about the SN progenitor systems. Different progenitor stars, binary systems, and binary evolution scenarios predict different DTDs. The ``border'' in progenitor masses between SNe~Ia and CC-SNe will lead to a time border in the DTD, dictated by stellar evolution, between the two types. The age corresponding to a particular mass border will depend on metallicity. For example, an $8M_\\odot$, $0.4 Z_\\odot$ (where $Z_\\odot$ is the solar metallicity) star lives for 41 Myr. For reference here and in the rest of this paper, Table~\\ref{lifetimes} lists the \\citep{girardi00:Low_Intermediate_Mass_Stars} lifetimes of stars of various zero-age main sequence masses, for various metallicities. The mean metallicities of the Small and Large Magellanic Clouds, the two galaxies that will be at the focus of this paper, are $0.1 Z_\\odot$ and $0.4 Z_\\odot$, respectively (Russell \\& Dopita 1992). \\begin{table} \\label{lifetimes} \\centering \\begin{minipage}{100mm} \\caption{Stellar Lifetimes} \\begin{tabular*}{80mm}{@{}crrr@{}} \\hline Mass & \\multicolumn{3}{r}{Lifetime\\footnote{Girardi et al. 2000} [Myr]} \\\\ $[M_\\odot]$ & $0.05Z_\\odot$ & $0.4Z_\\odot$ & $Z_\\odot$\\\\ \\hline 3&318 &363 & 477 \\\\ 4&167 &182 & 214 \\\\ 5&105 &112& 121 \\\\ 6&74 &76 &78 \\\\ 7&55 &56& 55 \\\\ 8&40 &41&40 \\\\ 9&31 &33&31 \\\\ \\hline \\end{tabular*} \\end{minipage} \\end{table} The shape of the DTD is especially important to constrain the hotly debated nature of SN~Ia progenitors. For both of the currently popular progenitor scenarios, SD and DD, calculations of the DTD depend on a series of assumptions regarding initial conditions (initial mass function -- IMF , binarity fraction, mass ratio distribution, separation distribution), and complex physics (mass loss, mass transfer, common envelope evolution, accretion) that are often computationally intractable except in the most rudimentary, parametrized forms \\citep[e.g.,][Bogomazov \\& Tutukov 2009; Mennekens et al. 2010]{yungelson98:SN_Ia_progenitors,hurley02:binaries_tides,han04:SDchannel_for_SNIa,nelemans05:common_envelope_in_WD_binaries,beer07:common_envelope,ruiter09:SNIa_rates_delay_times,bear10:common_envelope}. Observational estimates of the DTD could rule out particular DTD predictions, or at least could provide some input and generic features that successful models will need to reproduce. However, to date, observational SN DTD estimates have been few and controversial. One approach has been to compare the SN rate in field galaxies, as a function of redshift, to the cosmic star formation history (SFH). Given that the DTD is the SN ``response'' to a short burst of star formation, the SN rate versus cosmic time will be the convolution of the full SFH with the DTD. \\citet{gal-yam04:redshift_Ias} carried out the first such comparison, using a small sample of SNe Ia out to $z=0.8$, and concluded that the results were strongly dependent on the poorly known cosmic SFH, a conclusion echoed by \\citet{forster06:SNIa_progenitor_delays}. With the availability of SN rate measurements to higher redshifts, \\citet{barris04:Ia_distances} found a SN~Ia rate that closely tracks the SFH out to $z\\sim 1$, and concluded that the DTD must be concentrated at short delays, $\\lesssim 1$~Gyr. Similar conclusions have been reached, at least out to $z\\sim 0.7$, by \\citet{sullivan06:SNIa_rate_host_SFR}. In contrast, \\citet{dahlen04:high_z_SN_rates,dahlen08:extended_HST_survey} and \\citet{strolger04:higher_z} have argued for a DTD that is peaked at a delay of $\\sim 3$~Gyr, with little power at short delays, based on a decrease in the SN~Ia rate at $z>1$. However, \\citet{kuznetsova08:reanalysis_high_z_SN_rates} have re-analyzed some of these datasets and concluded that the small numbers of SNe and their potential classification errors preclude reaching a conclusion. Similarly, \\citet{poznanski07:SNIa_rate_subaru_df} performed new measurements of the $z>1$ SN~Ia rate, and found that, within uncertainties, the SN rate could be tracking the SFH. This, again, would imply a short delay time. \\citet{greggio08:SN_Ia_rates} pointed out that these results could also be affected by an underestimated extinction for the highest-$z$ SNe, which are observed in their rest-frame ultraviolet emission. A second approach to recovering the DTD has been to compare the SN rates in stellar populations of different characteristic ages. Using this approach, \\citet{mannucci05:SN_rate,mannucci06:two_progenitor_populations_SNIa}, \\citet{scannapieco05:A+B_models}, \\citet{sullivan06:SNIa_rate_host_SFR}, and \\citet{raskin09:prompt_Ia_SNe} have all found evidence for the co-existence of two SN~Ia populations, a ``prompt'' population\\footnote{We note that different authors have used the term ``prompt'' to describe different delays, from $<100$~Myr to as long as $< 1$~Gyr. In this work, we will refer as prompt SNe Ia to those with delays $<330$~Myr.} that explodes within of order $10^8$~yr, and a delayed channel that produces SNe~Ia on timescales of order 5~Gyr. Naturally, these two ``channels'' may in reality be just integrals over a continuous DTD on two sides of some time border \\citep{greggio08:SN_Ia_rates}. \\citet{totani08:SNIa_DTD} have used a similar approach to recover the DTD, by comparing SN~Ia rates in early-type galaxies of different characteristic ages. They find a DTD consistent with a $t^{-1}$ form. Additional recent studies using this approach, but which may be influenced by selection effects and by \\textit{a posteriori} statistics \\citep[because they focus on the properties of SN host galaxies; see][]{maoz08:fraction_intermediate_stars_Ia_progenitors} can be found in \\citet{aubourg07:massive_stars_SNIa}, \\citet{cooper09:metallicity_bias_SNIa}, Schawinski (2009), and \\citet{yasuda10:SNIa_luminosity}. A third approach for recovering the DTD is to measure the SN-rate vs. redshift in massive galaxy clusters. Optical spectroscopy and muliwavelength photometry of cluster galaxies has shown that the bulk of their stars were formed within short episodes at $z\\sim 2-3$ \\citep[e.g.][]{eisenhardt08:clusters}. Thus, the observed SN rate vs. cosmic time (since the stellar formation epoch) essentially provides a direct measurement of the form of the DTD. Furthermore, the record of metals stored in the intracluster medium (ICM) constrains the number of SNe that have exploded, and hence the normalization of the DTD. \\citet{maoz04:Ia_rate_clusters} analysed the then-available cluster SN~Ia rates \\citep{gal-yam02:HST_SNIa_cluster_rates}. The low observed SN~Ia rates out to $z\\sim 1$ implied that the large number of events needed to produce the bulk of the iron occurred at even higher redshifts, beyond the range of the then-existing observations. They concluded that most of the cluster SNe~Ia exploded during the relatively brief time interval between star formation in massive clusters (at $z\\sim 2-3$) and the highest-redshift cluster SN rate measurements (at $z\\sim 1$). Cluster SN rate measurements have by now greatly improved in the redshift range from zero to 1, and beyond \\citep[][]{sharon07:SNIa_cluster_rate,gal-yam08:SN_rates_low_z_clusters,mannucci08:SN_rate_nearby_clusters,graham10:SNIa_rate,sharon10:hst_cluster_rates,dilday10:sdss_clusters}. Analysis of these rates (Maoz et al. 2010b) reinforces the conclusion that, although the SN~Ia DTD may have a low-level tail at delays of a few Gyr, the bulk of the events must occur within $\\sim 1$~Gyr of star formation. Finally, \\cite{maoz10:SN_DTD_LOSS} have presented a method for recovering the DTD from a SN survey, where the SFHs of the individual galaxies surveyed are taken into account. They applied the method to a subset of the Lick Observatory SN Search (LOSS -- Filippenko in prep.; Leaman et al. in prep.; Li et al. in prep.) that overlaps with the Sloan Digital Sky Survey \\citep[SDSS;][]{york00:SDSS_Technical}, for which SFH reconstruction is availale based on the SDSS spectroscopy \\citep{tojeiro09:SDSS_VESPA}. They found that a ``prompt'' ($<420$~Myr) SN~Ia component is required by these data at the $>99\\%$ confidence level. In addition, a delayed SN~Ia population, with delays of $>2.5$~Gyr, is detected at the $4\\sigma$ level. A related approach has been used by \\cite{brandt10:Ia_progenitors}, with similar conclusions. In this paper, we seek to recover the SN DTD in yet another way, by analyzing the SNRs in the Magellanic Clouds (MCs). The MC SNRs present several advantages. First, they are all at the known distances of their two host galaxies. Second, both MCs have been surveyed to large depths in the radio, with individual sources followed up at multiple wavelengths and classified. In a companion paper (Badenes, Maoz, \\& Draine 2010, hereafter Paper I), we have compiled a sample of MC SNRs, and argued that it is largely complete. Third, the MCs are close enough to permit a region-by-region fitting of their resolved stellar populations with detailed stellar evolution isochrones. The SFHs of the Clouds can therefore be reconstructed with better spatial and temporal detail than in any other galaxies. \\citet{harris04:SMC_SFH} and \\citet{harris09:LMC_SFH} have recently carried out such a program of SFH reconstruction of the MCs. \\citet{badenes09:SNRs_LMC} have compared the SFHs of the MC regions hosting specific SNRs to the properties of the remnants, in order to deduce constraints on the nature of some of the explosions and their delay times. Here, we go a step further, and treat the MCs and their SNRs as an effective SN survey conducted in a sample of galaxy subunits, where the detailed SFH of each subunit is known. In \\S2, below, we briefly review the reconstruction of MC SFHs by \\citet{harris04:SMC_SFH} and \\citet{harris09:LMC_SFH}. In \\S3 we review our MC SNR sample from Paper I, and the physical model that reproduces the observed size distribution of this SNR sample. We show how this same model permits estimating the relative visibility times of SNRs at different locations in the MCs, given the local ambient densities. As in Paper I, we use three different tracers of MC density to estimate the SNR visibility times as a function of location. In \\S4, we review the method of \\cite{maoz10:SN_DTD_LOSS} for recovering the most likely SN DTD and its uncertainty, given the spatially resolved SFHs, the SNR visibility times, and the observed number of SNRs. In \\S5 we apply the method to the MCs with their sample of SNRs, and derive the DTD. In \\S6 we use the visibility times to obtain also the current SN rate in the MCs, and discuss the emerging picture. For the purpose of this paper, we assume that the DTD is a universal function: it is the same in all galaxies, independent of environment, metallicity, and cosmic time --- a simplifying assumption that may be invalid at some level. For example, a dependence of SN delay time on metallicity is expected in some models (e.g., Kobayashi et al. 2000). Similarly, variations in the initial mass function (IMF) with cosmic times or environment would also lead to a variable DTD, but we will again ignore this possibility in the present context. ", "conclusions": "In Paper~I, we assembled a multi-wavelength compilation of the 77 known SNRs in the MCs, collected from the existing literature. We verified that this compilation is fairly complete, and that the size distribution of SNRs is approximately flat, within the allowed uncertainties, up to a cutoff at $r\\sim30$ pc, as noted by other authors before. We then proposed a physical model to explain this size distribution. According to our model, most of the SNRs are in their Sedov stage, quickly fading below detection as soon as they reach the radiative stage. Under these circumstances, a flat distribution of SNR sizes can be obtained, provided that the distribution of densities in the MCs follows a power law with index $-1$. Finally, in Paper~I, we used three different density tracers (HI column density, H$\\alpha$ flux, and SFR) to demonstrate that the distribution of densities in the MCs indeed follows a power law of index $-1$. In this paper, we have used the Paper~I sample of SNRs as an effective supernova survey, conducted over tens of kyr. We have applied a novel technique to this SNR sample to derive, in these galaxies, the SN rates and the SN DTD. In order to accomplish this, we have used the three tracers of the density to scale the visibility time of SNRs in the MCs. From the scaled visibility times and the SFH maps from the resolved stellar populations in the MCs published by \\cite{harris04:SMC_SFH} and \\cite{harris09:LMC_SFH}, we have calculated the DTD. The DTDs we have derived are the first obtained using SFHs from resolved stellar populations. Our main findings are the following:\\\\ 1. We detect at the $>99\\%$ confidence level a ``prompt'' SN~Ia population, defined here as one that explodes within 330~Myr of star formation. This finding joins a growing number of measurements of such a component (see \\S1). However, our measurement is based on our statistically robust DTD recovery method, with its avoidance of the averaging inherent to many previous measurements, and using the most reliable SFHs, based on resolved stellar populations. The yield of prompt SN~Ia, in terms of SNe per stellar mass formed, is also consistent with other measurements. \\\\ 2. We obtain upper limits on the delayed SN~Ia population, which are consistent with other measurements. Using these upper limits, we find that roughly half, but possibly more, of SNe~Ia are ``prompt'', as defined above. This again joins previous results on the large relative fraction of the prompt population.\\\\ 3. We use our SNR sample and our SNR visibility times to derive current MC SN rates. The SN rate in the MCs agrees well with the historical record for these galaxies. The mass-normalised SN rate in the MCs is in excellent agreement with the rates measured by SN surveys in galaxies of this type. This lends support to the physical model we have presented to explain the SNR size distribution and the clouds, and which we have used to derive visibility times for our DTD and rate calculations.\\\\ 4. Conversely, a ``free expansion'' model for the SNRs, as has been invoked for the MCs and other galaxies, would imply SN rates in conflict with the historical record and with the rates in other star-forming dwarf galaxies. Furthermore, this model would indicate an unreasonably high yield of CC-SNe per stellar mass formed. The main limitation of our study has been the relatively small number of SNRs in the MC sample, which has forced us to use coarse time bins and has led to large Poisson errors. Construction of significantly larger samples of SNRs is, however, possible by means of deep radio surveys of additional galaxies that are near enough for deriving SFHs of their individual regions via resolved stellar populations, namely M31 and M33 (see Paper I). Such data would permit a similar analysis to the one we have done, but with larger SNR numbers, permitting a more accurate determination of the SN DTD, and bringing into better focus the properties of SNe, their remnants, and the connections between them." }, "1003/1003.4757_arXiv.txt": { "abstract": "After recent sensitivity upgrades at the Keck Interferometer (KI), systematic interferometric 2~$\\mu$m studies of the innermost dust in nearby Seyfert nuclei are within observational reach. Here, we present the analysis of new interferometric data of NGC~4151, discussed in context of the results from recent dust reverberation, spectro-photometric and interferometric campaigns. { The complete data set gives a complex picture, in particular the measured visibilities from now three different nights appear to be rather insensitive to the variation of the nuclear luminosity.} KI data alone indicate two scenarios: the $K$-band emission is either dominated to $\\sim 90\\%$ by size scales smaller than 30~mpc, which falls short of any dust reverberation measurement in NGC~4151 and of theoretical models of circum-nuclear dust distributions. Or contrary, and more likely, the $K$-band continuum emission is dominated by hot dust ($\\gtrsim 1300~K$) at linear scales of about 50~mpc. The linear size estimate varies by a few tens of percent depending on the exact morphology observed. { Our interferometric, deprojected centro-nuclear dust radius estimate of $55\\,\\pm\\,5$~mpc} is roughly consistent with the { earlier published} expectations from circum-nuclear, dusty radiative transfer models, and spectro-photometric modeling. { However, our data do not support the notion that the dust emission size scale follows the nuclear variability of NGC~4151 as a $R_{\\rm dust}\\,\\propto\\,L_{\\rm nuc}^{0.5}$ scaling relation.} Instead variable nuclear activity, lagging, and variable dust response to illumination changes need to be combined to explain the observations. ", "introduction": "\\label{sec:1} Nearby Active Galactic Nuclei (AGN) are rosetta stones to understand the astrophysics close to an actively accreting super-massive black hole. Sub-parsec resolution is required to identify answers to the key questions of AGN physics: (i) How does galactic material flow down to the accretion disk to feed luminosities close to the Eddington limit? (ii) Which role do play outflows and jets in the energetic interconnection between the AGN and its host? and (iii) Are AGN of different luminosity intrinsically similar? Studying the spectral energy distribution of quasars and type 1 AGN from the optical to the far-infrared reveals remarkably similar SED shapes, suggesting similar physics \\citep{1986ApJ...308...59E,1989ApJ...347...29S,1993ApJ...404...94K,2006A&A...457...61R}. Therefore, resolving astrophysical phenomena in nearby AGN enables to better understand farther, unresolved nuclei, and to gauge correlations between nuclear luminosity and properties of the surrounding circum-nuclear environment. Being one of the brightest AGN on the sky, \\object{NGC~4151} reveals this unique potential of nuclei closer than a few tens of Mpc's in numerous publications. In this article, we concentrate on the origin of the near-infrared (NIR) emission of \\object{NGC 4151}, and its relation to the nuclear luminosity. NIR-AGN emission is characterized by a steep decline from optical wavelengths down to about 1~$\\mu$m, and on the long wavelength side by the quick rise of a separate IR emission bump peaking at about 3~$\\mu$m \\citep{1986ApJ...308...59E,1993ApJ...404...94K}. The origin of this NIR excess has been discussed throughout the past four decades \\citep[probably one of the first were][]{1968AJ.....73..870P, 1969Natur.223..788R}. While it is not in question that the original power originates in the accretion disk, it is uncertain if the NIR emission excess derives directly from nuclear emission processes, or if mostly dust outside the broad line region (BLR) re-radiates the nuclear emission \\citep{1986ApJ...308...59E,1988Natur.336..749E,1987ApJ...320..537B,1992ApJ...400..502B}. The proper measurement of the nuclear near-infrared emission (and its discrimination from dust-re-processed light) is crucial to derive the intrinsic SED of accreting supermassive black holes (SMBH) in the centers of galaxies and to understand the radiative processes involved. Even at the relatively detailed linear scale of NGC~4151 ($\\sim\\,82\\,{\\rm mpc\\,mas^{-1}}$, footnote~$^{(a)}$ in Table~\\ref{tab:33}), the accretion disk itself and the surrounding BLR are too small to be spatially resolved by the current generation of optical or infrared telescopes or interferometers \\citep{2006ApJ...644..133B}. However, the circum-nuclear dust distribution is now within reach of observations with infrared telescope arrays at sub-50~mas angular resolution \\citep{2003ApJ...596L.163S,2004A&A...418L..39W,2004Natur.429...47J}. Mid-infrared data of the VLT interferometer typically find in Seyfert nuclei 10~$\\mu$m emission sizes of a few parsecs \\citep{2009A&A...502...67T, 2009ApJ...705L..53B, 2009MNRAS.394.1325R}. Dust plays an exceptional role among the circum-nuclear components. It is not only assumed to significantly contribute to the near-to-mid-infrared continuum radiation of an AGN. Also, a non-spherically symmetric dust distribution is widely assumed to explain the type~1~/~2 dichotomy of AGN emission line spectra. Radiative transfer models of clumpy distributions of dust clouds extending out of the equatorial plane are currently the favored explanation of the steadily increasing amount of observational data \\citep{1987ApJ...320..537B,1992ApJ...400..502B,2005A&A...437..861S,2006A&A...452..459H,2008ApJ...685..147N,2008ApJ...685..160N,2009arXiv0909.4539H}. In the following we refer to the dust distribution as {\\it torus}, keeping in mind that its detailed morphology is rather uncertain, might not resemble closely a smooth torus. It is currently an open debate, how the findings and models at MIR wavelengths connect to the innermost hot dust, which is expected to contribute to the continuum emission at 2~$\\mu$m. \\citet{2009ApJ...705..298M} argue for a spatially and chemically distinct inner component, based on $2\\,-\\,35\\,\\mu$m spectra. Also, the interferometric MIR data of NGC~4151 appear to reject that the nuclear flux at 10~$\\mu$m is dominated by the outer, cooler part of the same structure, which dominates the 2~$\\mu$m continuum \\citep{2009ApJ...705L..53B}. Such results discourage simple attempts to unify the dust emission structures around AGN, and suggest multi-component models, or at least a significant change of dust composition and grain size distribution with the radial distance to the central engine \\citep{2005A&A...437..861S, 2009A&A...493L..57K}. We report in this article on first results of a new campaign using the Keck Interferometer (KI) to add observational constraints to the origin of the NIR continuum of NGC~4151. The 85~m baseline and the sensitivity of the KI is adequate to resolve the linear distances of order 30-200 mpc around NGC~4151 which matches the theoretically calculated dust sublimation radii for this source. Recent sensitivity improvements of the KI \\citep[][\\footnote{Current KI performance numbers are given at the website of NASA Exoplanet Science Institute ({\\sc Nexsci}): {\\tt http://nexsci.caltech.edu/software/KISupport/v2/v2sensitivity.html}}]{2006SPIE.6268E..21W, 2008SPIE.7013E..10R} enabled us to repeat the early Swain et al. KI measurement, although the variable nucleus of NGC~4151 has been at a significantly fainter state during the time of observation. After a description of the observations (Section~\\ref{sec:2}), we discuss our findings and the implications on the interpretation of the combined visibility data set { \\citep[from][and the here presented observations]{2003ApJ...596L.163S, 2009A&A...507L..57K} } in Section~\\ref{sec:5}. { We concentrate our discussion on issues left open by the previous high resolution NIR observations of the AGN: (i) is the $K$-band emission rather dominated by unresolved accretion disk emission \\citep[as favored by ][]{2003ApJ...596L.163S}, or by resolved circum-nuclear dust emission \\citep{2009A&A...507L..57K}? (ii) How closely related are accretion disk luminosity and circum-nuclear dust location in NGC~4151 \\citep{2009ApJ...700L.109K}? (iii) Is the currently described scatter between direct interferometric and indirect reverberation measurements of the torus size a systematic offset reflecting that the respective method probes slightly different dust reservoirs, or does the scatter rather reflect the currently achievable accuracy of either method in deriving the actual dust location \\citep{2009ApJ...700L.109K,2009A&A...507L..57K}? } The concluding remarks in Section~\\ref{sec:6} summarize our current results and give a brief outlook on the project. ", "conclusions": "\\label{sec:6} The results of a new interferometric NIR observing campaign of NGC~4151 are presented. Although this time in a 2 times fainter state, we managed to re-observe the AGN with the Keck Interferometer thanks to recent sensitivity improvements of the instrument. The interferometric visibilities and size estimates were compared to previously published interferometric and single telescope data. The current KI datasets suggest that the major part of the radiation originates in a (possibly inclined) toroidal structure of an intrinsic, deprojected radius of about $55\\,\\pm\\,5$~mpc, comparable to $V$-to-$K$-continuum reverberation measurements, without constraining the morphology in further detail. Our dataset { and its comparison to published data enables for the first time to study the sensitivity and response of near-infrared interferometric visibilities of an AGN to intrinsic flux variations}. The observations show that NIR-interferometry on bright AGN ($K\\,\\sim\\,$10~mag) are now feasible at a high precision of a few percent. We did not detect a significant visibility dependence on doubling the $K$-band emission of NGC~4151. This supports the notion, that, for an individual AGN, the size of its circum-nuclear NIR emission structure does not strictly depend on the respective momentary nuclear luminosity. The direct interferometric size estimates appear to be more robust estimates of the average location of the dust than NIR continuum reverberation estimates. Being flux-independent, this average location of the hot dust probably relates to the sublimation radii of the AGN at its high activity state. Future dust reverberation campaigns, contemporaneous to direct interferometric measurements are highly desirable to study the circum-nuclear NIR emission region, to correctly interpret the reverberation measurements, and to investigate apparent changes in the dust illumination, heating efficiency, or covering factors." }, "1003/1003.5751_arXiv.txt": { "abstract": "We introduce a novel class of field theories where energy always flows along timelike geodesics, mimicking in that respect dust, yet which possess non-zero pressure. This theory comprises two scalar fields, one of which is a Lagrange multiplier enforcing a constraint between the other's field value and derivative. We show that this system possesses no wave-like modes but retains a single dynamical degree of freedom. Thus, the sound speed is always identically zero on all backgrounds. In particular, cosmological perturbations reproduce the standard behaviour for hydrodynamics in the limit of vanishing sound speed. Using all these properties we propose a model unifying Dark Matter and Dark Energy in a single degree of freedom. In a certain limit this model exactly reproduces the evolution history of $\\Lambda$CDM, while deviations away from the standard expansion history produce a potentially measurable difference in the evolution of structure. ", "introduction": "How can one obtain dust from a scalar field? One can imagine a canonical scalar-field where the kinetic term is constrained to be equal to the potential. We can implement this property by introducing a Lagrange multiplier, $\\lambda$, in the Lagrangian, \\[ \\mathcal{L}=\\lambda\\left(\\frac{1}{2}(\\partial\\varphi)^{2}-V(\\varphi)\\right)\\,.\\] We then find that the pressure is identically vanishing on all solutions and energy follows geodesics. This model describes the usual dust without vorticity. How can we obtain ``dust with pressure''? We can generalise the above by adding some function of the scalar field and its derivatives to the Lagrangian,\\[ \\mathcal{L}=K\\left(\\varphi,\\partial\\varphi\\right)+\\lambda\\left(\\frac{1}{2}(\\partial\\varphi)^{2}-V(\\varphi)\\right)\\,.\\] The constraint remains in effect and standard scalar-field dynamics are not restored. In fact, we will show that fluid elements in all such theories \\emph{also }always\\emph{ }flow along geodesics, mimicking in that respect standard dust, yet the fluid has non-vanishing pressure. With this simple idea we have separated the notion that the pressure of the fluid is tied to the motion of a fluid element as is the situation in the usual case, e.g. radiation or cold dark matter. A parcel of such fluid will flow along geodesics, yet a manometer will record a pressure changing with time. In this paper, we introduce this new class of scalar-field models, which we will call $\\lambda\\varphi$-\\emph{fluids}. These theories are described by an action containing two scalar fields, $\\varphi$ and $\\lambda$, where the latter plays the role of a Lagrange multiplier and enforces a constraint relating the value of the scalar field $\\varphi$ to the norm of its derivative. This constraint forces the dynamics of the $\\lambda\\varphi$-\\emph{fluid} to be driven by a system of two first-order \\emph{ordinary }differential equations, one for the field $\\varphi$, the other for the Lagrange multiplier. As a consequence, there are no propagating wave-like degrees of freedom and the sound speed for perturbations is exactly zero irrespective of the background solution. However, the initial-value problem still requires the specification of two functions on the initial time slice. Thus, effectively, a single dynamical degree of freedom remains. Provided that the derivatives of the scalar field $\\varphi$ are time-like, the system can be interpreted as a perfect fluid. However, for a $\\lambda\\varphi$-\\emph{fluid} given by a particular action, the relation between the pressure $p$ and the energy density $\\varepsilon$ is solution dependent. We show that an arbitrary effective equation of state, including phantom ones, can be obtained by choosing the form of the Lagrangian appropriately. In addition, for all $\\lambda\\varphi$-\\emph{fluids}, there always exists a region in their phase spaces in which the $\\lambda\\varphi$\\emph{-fluid} is effectively pressureless. We will exploit this feature to model the evolution of the cosmological background from matter domination through to the acceleration era as being driven by the dynamics of a single degree of freedom provided by the $\\lambda\\varphi$\\emph{-fluid}. The key novel aspect of this class of theories is that the four-acceleration is always zero, even when the pressure does not vanish. This is a result of the constraint's eliminating those fluid configurations where the pressure has a gradient orthogonal to the fluid velocity. Motivated by these properties, we will use the $\\lambda\\varphi$\\emph{-fluid} to frame the problem of the dark sector in cosmology in a unified manner. The existence of the dark sector in the Universe's energy budget is now established beyond reasonable doubt. The standard model of cosmology, $\\Lambda$CDM, splits it into two constituents: cold dark matter (``CDM'')---a pressureless fluid (``dust'') which clusters allowing baryonic structures to form in its potential wells and detectable to this day in the form of halos around galaxies---and dark energy (``DE'')---a form of energy that appears to be smooth and to have an equation of state close to a cosmological constant. This dichotomy of phenomenology has made it difficult to build a compelling model which would treat the two dark components in a unified fashion. Nonetheless, some models exist in the literature: \\cite{Padmanabhan:2002sh, Scherrer:2004au,Bertacca:2007ux, Arbey:2006it, Kamenshchik:2001cp, Bento:2004uh, Capozziello:2005tf, Nojiri:2005pu, Quercellini:2007ht, Linder:2008ya, Piattella:2009kt, Gao:2009me,Creminelli:2009mu}. It is not hard to see that, given a fluid which clusters like dust yet has arbitrary pressure, we can construct such a unified model---which we call Dusty Dark Energy (``DDE''). We present the main results of our application of $\\lambda\\varphi$\\emph{-fluid} to such a cosmology beneath. We refer the reader to the full analysis in the main body of the paper, section \\ref{s:DDE}. \\subsection{Summary of Cosmological Results} We have studied cosmological perturbations in the case when an arbitrary $\\lambda\\varphi$\\emph{-fluid} dominates the\\emph{ }Universe. We have derived the closed-form equation for the evolution of the Newtonian potential $\\Phi$ which turns out to recover the standard result for general hydrodynamics in the limit of vanishing sound speed. This evolution is determined by background expansion history only and in the limit of the $\\Lambda$CDM expansion history the evolution of perturbations is exactly as in the standard case. We have also written down an action for perturbations which explicitly shows that there are no ghosts in this theory when the equation of state for the $\\lambda\\varphi$-\\emph{fluid} is non-phantom. We demonstrate that on a classical level our model can cross the phantom divide without singularities while linear perturbations continue to evolve stably; however, the perturbations do become ghosts at this point, hence making the system unstable (in particular quantum-mechanically) when interactions are taken into account. In this paper we put the investigation of that instability and issues related to the possible strong coupling scales aside. We consider a universe comprising only radiation and the $\\lambda\\varphi$-\\emph{fluid} which will describe both CDM and DE. To illustrate more concretely the phenomenology we have focused on a specific family of models which is parameterised by $w_{\\text{fin}}$---the equation of state of the $\\lambda\\varphi$-\\emph{fluid} in the asymptotic future. Given that the initial values for the $\\lambda\\varphi$-\\emph{fluid} are chosen appropriately, the radiation becomes subdominant while the $\\lambda\\varphi$-\\emph{fluid} is still far off its final attractor (given by $w_{\\text{fin}}$) and evolves approximately like dust, giving an epoch of matter domination. The duration of this epoch is also determined by the initial values. In the limit $w_{\\text{fin}}\\rightarrow-1$ this family of models recovers exactly the background evolution and growth of structure of $\\Lambda$CDM. However, if $w_{\\text{fin}}\\neq-1,$ the evolution of the background and perturbations differs from a ``$w$CDM'' model comprising cold dark matter and dark energy with a constant equation of state. We illustrate the evolution of the effective equation of state for the dark sector in a selection of different $w_{\\text{fin}}$-cosmologies in Fig.\\ \\ref{fig:wevol}. We find that the transition from matter-domination to dark energy domination differs from that of $w$CDM (Fig.\\ \\ref{fig:wprime}). We compare the growth of linear perturbations with that of $\\Lambda$CDM in Fig.\\ \\ref{fig:Ratio}. We find that models with $1+w_{\\text{fin}}>0$ exhibit a growth factor suppressed by a few tens of percent. Also the Newtonian potential here is lower by a few percent than in $\\Lambda$CDM, which would decrease the integrated Sachs-Wolfe effect. \\begin{figure} \\includegraphics[width=1\\columnwidth]{w_evol} \\caption{Time evolution of the total effective equation of state for the dark sector. The black solid line represents the evolution in $\\Lambda$CDM which is identical to that of the $w_{\\text{fin}}=-1$ model. Models with final equations of state $1+w_{\\text{fin}}>0$ begin to deviate from matter domination earlier and the transition is slower than $\\Lambda$CDM. The opposite is true in phantom models. The evolution is normalised such that the equation of state at $a=1$ matches the best-fit result for the $\\Lambda$CDM cosmology as determined by WMAP7 results, $w_{0}=-0.74$ \\cite{Komatsu:2010fb}. \\label{fig:wevol} } \\end{figure} \\begin{figure} \\includegraphics[width=1\\columnwidth]{wXp}\\caption{Comparison of the derivative of the effective equation of state for Dusty Dark Energy (Eq.\\ \\eqref{LambdaEvolutionWeqConst}) with that for a dark matter plus dark energy with a constant equation of state, $w$CDM, (Eq.\\ \\eqref{WevolutionMixture}). The magnitude of the derivative determines the duration of the transition between matter domination and the acceleration era. The left panel shows that for $w_{\\text{fin}}>-1$, the transition in the DDE model is more rapid than the corresponding $w$CDM model. On the other hand, for phantom $w_{\\text{fin}}$ this transition is slower than for the corresponding $w$CDM model, as shown in the panel on the right. \\label{fig:wprime}} \\end{figure} \\begin{figure} \\includegraphics[width=1\\columnwidth]{Phi_plot}\\caption{The comparison of the total growth of perturbation amplitude between Dusty Dark Energy and $\\Lambda$CDM. The evolution of the Newtonian potential, $\\Phi$, is determined by Eq.\\ \\eqref{PhiN} and deviates by a few percent from its $\\Lambda$CDM values by a few percent. This evolution will affect the strength of the ISW signal in the CMB. On the other hand, the evolution of the density perturbation on subhorizon scales is determined by Eq.\\ \\eqref{dotRelativeDeltaE} and is affected much more strongly.\\textbf{\\label{fig:Ratio}}} \\end{figure} \\subsection{Open Questions and Future Directions} Finally, we mention open issues in this setup which remain to be addressed \\begin{itemize} \\item The nature of caustics: It is well known that in non-canonical field theories caustic can develop, e.$\\,$g. see Ref. \\cite{Felder:2002sv} on caustics in Sen's string-theoretical tachyon matter \\cite{Sen:2002in,Sen:2002nu} and Ref. \\cite{ArkaniHamed:2005gu} on caustics in the ghost condensate \\cite{ArkaniHamed:2003uy} and the discussion in Refs \\cite{Mukohyama:2009tp,Blas:2009yd} on caustics in Ho\\v{r}ava gravity \\cite{Horava:2009uw}. We expect that the $\\lambda\\varphi$\\emph{-fluid }will develop caustics. The question is how to interpret the multivalued regions once this occurs. \\item Can $\\lambda\\varphi$\\emph{-fluids} virialise? If the $\\lambda\\varphi$\\emph{-fluid} is to really model non-linear structure, it must be able to form static and stable configurations (e.g. halos). \\item What is the origin for the initial conditions for Dusty Dark Energy introduced in section \\ref{s:DDE}? Could a more generalised setup provide a solution for the coincidence problem? In our model, we require that at some point during the radiation epoch the energy density of the $\\lambda\\varphi$\\emph{-fluid} be equal to the energy density of CDM: we have no alternative to the standard dark-matter freeze out scenario which would produce this in a natural manner. \\item What is the Hamiltonian structure of this theory? How to quantise it? See, for example \\cite{Brown:1994py}. Is the structure of this theory stable as a result of radiative corrections: would a kinetic term for $\\lambda$ be generated? \\item What is the strong-coupling scale for perturbations? We should note that our model is rather similar to a potentially singular limit of Ho\\v{r}ava gravity \\cite{Horava:2009uw}. There it was found \\cite{Blas:2009yd} that, contrary to our case, the sound speed for cosmological perturbations becomes imaginary, and that the strong coupling scale may be extremely low. \\item What is the rate of instability for the phantom case? Is this instability catastrophic as it is usually for ghost degrees of freedom? It is possible that the absence of propagating wave-like degrees of freedom may change the standard picture \\cite{Carroll:2003st,Dubovsky:2005xd,Woodard:2006nt,Cline:2003gs,Emparan:2005gg}. \\item The action formulation of this theory allows us to consider couplings to standard model fields. It is natural, for example, to consider Lorentz violation in this framework in analogy to Einstein aether theories \\cite{Jacobson:2004ts,Carroll:2004ai,Lim:2004js}. \\item Can this theory be a low-energy limit of a more fundamental theory? \\item Can our conceit with the constraint be usefully extended beyond scalar fields to theories with fermions, vector fields or many degrees of freedom? \\end{itemize} ", "conclusions": "In this paper, we introduced a novel class of field theories with a single dynamical degree of freedom which have a perfect-fluid interpretation. The key feature of this theory is that its fluid velocity flows along geodesics---hence mimicking ``dust'' in this respect. On the other hand, unlike a standard cold-dark-matter fluid, it carries pressure parallel to its fluid velocity. In cosmology, this pressure affects the expansion history. This sleight-of-hand is achieved by means of a ``Lagrange multiplier'' field, employed to constrain by equation of motion the above-mentioned behaviour of the fluid velocity vector. Our system then consists of \\emph{two first order equations of motion}, and hence effectively a single degree of freedom. This dynamic cannot be reproduced by usual scalar field theories such as k-\\emph{essence} or higher derivative theories. As an application, we consider the evolution and effects of this fluid in cosmology. We show that there exists a class of scaling solutions which have an attractor solution with fixed equation of state $w_{\\mathrm{fin}}$. Off the attractor, this model possess an interesting dynamic where part of the energy density redshifts as dust while part of the energy density tracks any dominant background energy density. We use this curious property to construct a unified dark energy/dark matter model where the limit $w_{\\mathrm{fin}}=-1$ corresponds to standard $\\Lambda$CDM. We also show that in this class of models, we can construct phantom models with $w_{\\mathrm{fin}}<-1$ where there is no pathology when crossing the ``phantom divide'' at $w=-1$, at least classically. If we insist that the system satisfy the Null Energy Condition (i.e. $w\\geq-1$), then we show that the kinetic term for perturbations is positive definite. Finally, we would like to conclude by emphasising that this class of theories provides a novel framework for cosmological model building and exploring exotic states of matter. \\medskip{}" }, "1003/1003.0945_arXiv.txt": { "abstract": "WZ~Sge-type dwarf novae are characterized by long recurrence times of outbursts ($\\sim 10\\;{\\rm yr}$) and short orbital periods ($\\lesssim 85\\;{\\rm min}$). A significant part of WZ~Sge stars may remain undiscovered because of low outburst activity. Recently, the observed orbital period distribution of cataclysmic variables (CVs) has changed partly because outbursts of new WZ~Sge stars have been discovered routinely. Hence, the estimation of the intrinsic population of WZ~Sge stars is important for the study of the population and evolution of CVs. In this paper, we present a Bayesian approach to estimate the intrinsic period distribution of dwarf novae from observed samples. In this Bayesian model, we assumed a simple relationship between the recurrence time and the orbital period which is consistent with observations of WZ~Sge stars and other dwarf novae. As a result, the minimum orbital period was estimated to be $\\sim 70\\;{\\rm min}$. The population of WZ~Sge stars exhibited a spike-like feature at the shortest period regime in the orbital period distribution. These features are consistent with the orbital period distribution previously predicted by population synthesis studies. We propose that WZ~Sge stars and CVs with a low mass-transfer rate are excellent candidates for the missing population predicted by the evolution theory of CVs. ", "introduction": "Cataclysmic variables (CVs) are close binary systems containing a white dwarf and a Roche-lobe filling normal star. They have several subclasses, such as, dwarf novae (DNe), nova-like variables, novae, and magnetic CVs (\\cite{war95book}). CVs are considered to be one of the final stages of the evolution of low mass binaries. The formation and evolution of low mass binaries can be explored through the orbital period ($P_{\\rm orb}$) distribution of CVs (for a review, see \\cite{kin88binaryevolution}). Stable mass-transfer processes in CVs are driven by angular momentum removal from the binaries. In short period CVs having $P_{\\rm orb}\\lesssim 3$~hr, the driving mechanism is considered to be angular momentum removal associated with gravitational radiation (\\cite{pac81cvevolution}; \\cite{rap82cvevolution}). A CV evolves toward a short-$P_{\\rm orb}$ region by losing angular momentum. This phase is terminated when the core of the secondary star becomes degenerate. Then, the degeneracy pressure changes the mass--radius relationship of the secondary star. As a result, the CV evolves toward a long-$P_{\\rm orb}$ regime after passing through the minimum $P_{\\rm orb}$ ($P_{\\rm min}$). The $P_{\\rm orb}$ distribution near $P_{\\rm min}$ has received attention, since the observed $P_{\\rm orb}$ distribution has different characteristics from that expected from theoretical studies. It is well known that $P_{\\rm min}$ is $\\sim 80$~min in the observed CV population (\\cite{pac81cvevolution}; \\cite{rap82cvevolution}). However, the theoretically predicted $P_{\\rm min}$ is 60--70~min, significantly shorter than the observed value (e.g. \\cite{kol93CVpopulation}). This is the so-called ``period minimum problem''. According to population synthesis studies, most CVs have already passed through $P_{\\rm min}$. It has been proposed that the evolution time-scale is long for an evolutionary-advanced system with a low mass secondary star, in other words, the mass-transfer rate ($\\dot{M}$) from the secondary star is low (e.g. \\cite{kol93CVpopulation}; \\cite{how97periodminimum}). As a result, a spike-like feature is expected to appear near $P_{\\rm min}$ in the $P_{\\rm orb}$ distribution due to the accumulation of systems. However, the observed distribution is rather flat. This is the so-called ``period spike problem'' (\\cite{kol99CVperiodminimum}; \\cite{ren02CVminimum}). Thus, it is known that there is a ``missing'' population in the observed CVs with a $P_{\\rm orb}$ range of 70--80~min when compared with the predicted population. \\citet{lit08postPmin} have recently reported three CVs having a brown dwarf secondary, which are candidates for the missing population (\\cite{lit06sdss1035}; \\cite{lit07sdss1507}). This discovery indicates that CVs can actually reach the post-$P_{\\rm min}$ regime as predicted by the theories, while the number of such systems is still not properly known. It has been suggested that the problems can be reconciled by another mechanism of angular momentum removal in addition to gravitational radiation (\\cite{rez01gr}), such as the effects of magnetic stellar wind braking (\\cite{kin02CVperiodminimum}), the circumbinary disk (\\cite{wil05circumdisk}), and a combination of magnetic propeller and accretion disk resonances (\\cite{mat06mp}). On the other hand, it is possible that the missing population exists, but remains undiscovered. According to the disk instability model for DNe, the recurrence time of superoutbursts (supercycle; $T_s$) is longer in a system with a lower $\\dot{M}$ (\\cite{osa95wzsge}). WZ~Sge-type DNe are known to have long $T_s$ ($\\gtrsim 10\\;{\\rm yr}$), which indicate a low $\\dot{M}$ (\\cite{how95TOAD}; \\cite{kat01hvvir}). It has been proposed that several WZ~Sge stars, or CVs with a low $\\dot{M}$ have remained undiscovered because of their low activity (\\cite{mey98wzsge}). Since WZ~Sge stars are found concentrating in a short $P_{\\rm orb}$ regime, these objects are a candidate for the missing population (\\cite{pat98evolution}). However, their intrinsic contribution to the whole CV population has been poorly understood because their $T_s$ were uncertain and their quiescent luminosity is low. \\citet{gan09sdssCV} have recently reported that the CV sample from the Sloan Digital Sky Survey (SDSS) database shows a period spike feature in the $P_{\\rm orb}$ range of 80--86~min. The spike feature is a consequence of deep images and a homogeneously selected CV sample obtained by SDSS, which could detect intrinsically faint CVs with a low $\\dot{M}$. Recently, in addition, $T_s$ of several WZ~Sge stars have been well determined owing to long monitoring of DNe by amateur observers (e.g. \\cite{waa07gwlib}) In this paper, we investigate the contribution of long-$T_s$ DNe to the whole population of CVs by estimating their intrinsic population. In \\textsection~2, we demonstrate that the whole population of CVs has changed significantly due to a recent increase of the number of short $P_{\\rm orb}$ CVs. In \\textsection~3, we estimate the population of DNe by considering the detection probability of superoutbursts using Bayesian analysis. We discuss the implications of our results to the study of CV evolution in \\textsection~5. In the final section, we summarize our findings. ", "conclusions": "\\subsection{Dependence of $\\alpha$ and $P_{\\rm min}$ on the assumed model parameters} In the previous section, we assumed several parameters in our model to estimate $\\alpha$ and $P_{\\rm min}$, such as $n$, $P_{\\rm crit}$ and $M_V$ at supermaximum. Here, we discuss the dependence of the result on these parameters. We present $\\alpha$ and $P_{\\rm min}$ estimated with several different sets of these parameters in table~\\ref{tab:post}. As can be seen in table~\\ref{tab:post}, $\\alpha$ is sensitive to both $n$ and $P_{\\rm crit}$. Smaller $n$ or $P_{\\rm crit}$ leads to a smaller $\\alpha$. $\\alpha$ also depends on the model of $M_V$. It is possible that $M_V$ takes a form different from equation~(7). According to \\citet{har04MVPorb}, $M_{V,{\\rm SU}}$ for superoutbursts has a large dispersion and does not strictly follow the $M_V$--$P_{\\rm orb}$ relationship for normal outbursts. This implies that the $P_{\\rm orb}$ dependence of $M_{V,{\\rm SU}}$ is weak. Assuming $M_{V,{\\rm SU}}=M_{V,{\\rm WZ}}={\\rm const.}$, we obtained a smaller $\\alpha$ as shown in table~\\ref{tab:post}. In all the above cases, $\\alpha$ takes high values of $\\gtrsim 0.6$ and a spike near $P_{\\rm min}$ appears in the $P_{\\rm orb}$ distribution. The spike feature disappears only when $n$ is extremely small. For example, we show the case of $n=0.2$ in table~\\ref{tab:post}. Such a small $n$ is, however, unfavorable for the observed $D_{\\rm outb}$--$P_{\\rm orb}$ relationship, as shown in figure~\\ref{fig:doutb}. Table~\\ref{tab:post} indicates that $P_{\\rm min}$ is sensitive to $n$; a small $n$ yields a large $P_{\\rm min}$. On the other hand, the estimated $P_{\\rm min}$ in table~2 are significantly smaller than the observed value, even in the extreme case of $n=0.2$. Thus, the result is robust within the reasonable ranges of the model parameters. As mentioned in \\textsection~3.3, we used $n=2.0$, which is a lower limit of $n$ in the case of $P_{\\rm min}=72.6\\;{\\rm min}$. $n$ can be smaller than 2.0 only when $P_{\\rm min}> 72.6\\;{\\rm min}$. As shown in table~2, all $P_{\\rm min}$ satisfies $P_{\\rm min}<72.6\\;{\\rm min}$, except for the extreme case of $n=0.2$. Therefore, it is less likely that $n$ is smaller than 2.0 in our model. Our result highly depends on the assumption for the $D_{\\rm outb}$--$P_{\\rm orb}$ relationship. The assumption might be incorrect if a number of short-$T_s$ systems remain undiscovered in the shortest $P_{\\rm orb}$ regime, such as {\\it Group X} identified in Paper~I. The assumption for the $D_{\\rm outb}$--$P_{\\rm orb}$ relationship should thus be reconsidered if the intrinsic population of those objects is found to be much larger than the currently observed population. However, it is unlikely that there are significant undetected systems because the outburst detection probability must be high in such short-$T_s$ systems and hence their outbursts should have already been discovered. \\subsection{Implications for CV evolution} In this section, we discuss the implications of our result to CV evolution, considering the presence of post-$P_{\\rm min}$ CVs, magnetic CVs, and past theoretical and observational works in this field. In the previous section, we regarded all systems as pre-$P_{\\rm min}$ CVs. However, it has been proposed that about $70\\%$ of CVs have already passed through $P_{\\rm min}$ (\\cite{kol93CVpopulation}; \\cite{how97periodminimum}). It is posible that all known DNe are pre-$P_{\\rm min}$ CVs, while most CVs have already passed through $P_{\\rm min}$ and remained undiscovered. In this case, the intrinsic $P_{\\rm orb}$ distribution would have a taller spike near $P_{\\rm min}$ than that in figure~\\ref{fig:model} because of the contribution of the undiscovered post-$P_{\\rm min}$ CVs. On the other hand, a fraction of the long-$T_s$ DNe in our sample might actually be post-$P_{\\rm min}$ systems because post-$P_{\\rm min}$ systems presumably have quite low $\\dot{M}$ and long $T_s$. Even if most known WZ~Sge stars are post-$P_{\\rm min}$ CVs, the spike structure in figure~\\ref{fig:model} is still expected to appear because the spike feature depends, in principle, on the number of long-$T_s$ DNe near $P_{\\rm min}$ and is independent of the nature of the systems, namely, whether they are pre- or post-$P_{\\rm min}$ systems. In contrast to the spike feature, the real $P_{\\rm min}$ is possibly sensitive to the number of pre-$P_{\\rm min}$ systems in the shortest $P_{\\rm orb}$ regime. If the true fraction of post-$P_{\\rm min}$ systems is large in known WZ~Sge stars, then it means that superoutbursts of the post-$P_{\\rm min}$ systems are detectable for us. If this is the case, we should be able to detect outbursts of DNe below the observed $P_{\\rm min}$. The estimated $P_{\\rm min}$ could then be close to the observed value. The true intrinsic fraction of post-$P_{\\rm min}$ systems in the shortest $P_{\\rm orb}$ regime is poorly known, while recent observational studies suggest the presence of post-$P_{\\rm min}$ systems. \\citet{two09mdot} investigated the temperature of the white dwarf in CVs as a probe of $\\dot{M}$. Their result shows that there are anomalies having quite low temperatures of the white dwarf compared with most CVs in the shortest $P_{\\rm orb}$ regime. The $\\dot{M}$ of those anomalies are possibly lower than the ordinary ones, and hence, suggesting that they are post-$P_{\\rm min}$ systems. \\citet{lit08postPmin} reported three CVs in which the post-$P_{\\rm min}$ nature was dynamically confirmed (also see, \\cite{lit06sdss1035}; \\cite{lit07sdss1507}). Among known DNe, WZ~Sge has the longest $T_s$ and hence is a good candidate for post-$P_{\\rm min}$ systems. \\citet{ste07wzsge}, however, reported that its secondary mass is so high that WZ~Sge is in fact a pre-$P_{\\rm min}$ system. The nature of WZ~Sge implies that most DNe whose $T_s$ are known are still in the pre-$P_{\\rm min}$ regime. As mentioned in \\textsection~3.3, our model of $D_{\\rm outb}$ assumed a decrease of $\\dot{M}$ toward $P_{\\rm min}$ in the region $P_{\\rm orb}<86\\;{\\rm min}$. This was required from the observed $T_s$--$P_{\\rm orb}$ relationship, as can be seen in figure~\\ref{fig:doutb}. However, this is apparently inconsistent with theoretical calculations for the CV evolution, which predict a rather constant $\\dot{M}$ for pre-$P_{\\rm min}$ systems even near $P_{\\rm min}$ (e.g. \\cite{sha92DNperiod}; \\cite{kol93CVpopulation}; \\cite{bar03CVevolv}). If the theoretical prediction is true, we would expect many short-$T_s$ systems in the shortest $P_{\\rm orb}$ regime. The prediction is, however, unfavorable for the observed population, because there are only a few short-$T_s$ systems in $P_{\\rm orb}\\lesssim 86\\;{\\rm min}$, such as {\\it Group X} in Paper~I. This inconsistency implies that the current understanding of the secondary star is imperfect, since the theoretical $\\dot{M}$--$P_{\\rm orb}$ relationship depends on the structure of the secondary star. \\citet{mey99diskviscosity} proposed that the magnetic activity of the secondary star reduces rapidly as it becomes degenerate, and therby, the viscosity in the disk reduces and $T_s$ increases. The observed $T_s$--$P_{\\rm orb}$ relationship may be reprocuded if $T_s$ increases by this mechanism in $P_{\\rm orb}<86\\;{\\rm min}$. It has been reported that the $P_{\\rm orb}$ distribution of magnetic CVs is similar to that of non-magnetic CVs (e.g. \\cite{wil05circumdisk}). However, this discussion was based on samples obtained from various observation processes, such as samples in RKcat. As mentioned in section~2, such a sample is not suitable to be compared with the theoretical $P_{\\rm orb}$ distribution obtained from population systhesis studies. Our result suggests that, in the case of DNe, the period minimum and spike problems are reconciled with a number of undiscovered WZ~Sge stars. Since AM~Her stars have no accretion disk, and thereby, experience no DN-type outbursts, our result also implies that another mechanism to reduce short-$P_{\\rm orb}$ systems should work for magnetic CVs (e.g. \\cite{kol95magCV}; \\cite{mey99amherevolution}). The reduce of magnetic braking in magnetic CVs may lead to a different $P_{\\rm orb}$ distribution from non-magnetic CVs (\\cite{web02polarperiodgap}; \\cite{ara05mCV}). \\citet{gan09sdssCV} have recently reported the discovery of the period spike feature in the CV sample selected from the SDSS database. The appearance of the spike is in agreement with our result; we reproduced a spike based on observed outburst activity, while \\citet{gan09sdssCV} obtained it directly from the deep survey. They discovered that CVs having a white-dwarf dominated spectrum form a major population in the CVs making the period spike feature. Their spectrum indicates that $\\dot{M}$ of those CVs are quite low. Hence, they may be WZ~Sge stars, and their outbursts could be observed with a long $T_s$. Actually, SDSS~J080434.20+510349.2, a system having a white-dwarf dominated spectrum, experienced a WZ~Sge-type superoutburst in 2005 (\\cite{pav07j0804}). The SDSS CVs having a white-dwarf dominated spectrum probably have the same nature as the objects that our analysis predicts to be present in the shortest $P_{\\rm orb}$ regime. On the other hand, the implications for the period minimum problem are different between our result and \\citet{gan09sdssCV}. In our result, the observed $P_{\\rm min}$ is significantly longer than the real $P_{\\rm min}$ because $T_s$ is so long that the detection frequency is low near the real $P_{\\rm min}$. The SDSS images are so deep that we can expect to detect CVs with a low $\\dot{M}$ near the real $P_{\\rm min}$. However, a sharp cut-off exists at $P_{\\rm orb}\\sim 80\\;{\\rm min}$ in the $P_{\\rm orb}$ distribution of the SDSS sample and only a few sources were found in the $P_{\\rm orb}$ range of 60--78~min. The period minimum problem, hence, remains unresolved in \\citet{gan09sdssCV}. The lack of the sources in the $P_{\\rm orb}$ range 60--78~min may indicate that a fraction of known WZ~Sge stars are post-$P_{\\rm min}$ CVs, as discussed above. Alternatively, some CVs with a very low $\\dot{M}$ may remain undiscovered even in the SDSS survey. \\citet{gan09sdssCV} identified CVs based on spectral features of objects that were selected by several color criteria. A part of the color criteria plays a role in excluding single white dwarfs from the CV sample. The criteria, however, possibly excludes evolved CVs which have a quite low $\\dot{M}$, and thereby, have optical spectra dominated by the emission from a white dwarf." }, "1003/1003.5567_arXiv.txt": { "abstract": "Strong lensing is a powerful tool to address three major astrophysical issues: understanding the spatial distribution of mass at kpc and sub-kpc scale, where baryons and dark matter interact to shape galaxies as we see them; determining the overall geometry, content, and kinematics of the universe; studying distant galaxies, black holes, and active nuclei that are too small or too faint to be resolved or detected with current instrumentation. After summarizing strong gravitational lensing fundamentals, I present a selection of recent important results. I conclude by discussing the exciting prospects of strong gravitational lensing in the next decade. ", "introduction": "As photons from distant sources travel across the universe to reach our telescopes and detectors, their trajectories are perturbed by the inhomogeneous distribution of matter. Most sources appear to us slightly displaced and distorted in comparison with the way they would appear in a perfectly homogeneous and isotropic universe. This phenomenon is called weak gravitational lensing \\citep[e.g.][and references therein]{Ref03}. Under rare circumstances, the deflection caused by foreground mass overdensities such as galaxies, groups, and clusters is sufficiently large to create multiple images of the distant light source. This phenomenon is called strong gravitational lensing. Due to space limitations, this article will focus on cases where gravitational lensing is caused primarily by a galaxy-sized deflector (or lens). The first strong gravitational lens was discovered more than thirty years ago, decades after the phenomenon was predicted theoretically \\citep[see][and references therein]{B+N92}. However, in the past decade there has been a dramatic increase in the number of known lenses and in the quality of the data. At the time of the review by \\citet{B+N92}, the 11 ``secure'' known galaxy-scale lenses could all be listed in a page and discussed individually. At the time of this writing, the number of known galaxy-scale lens systems is approximately 200, most of which have been discovered as part of large dedicated surveys with well defined selection functions. This breakthrough has completed the transformation of gravitational lensing from an interesting and elegant curiosity to a powerful tool of general interest and statistical power. Three properties make strong gravitational lensing a most useful tool to measure and understand the universe. Firstly, strong lensing observables - such as relative positions, flux ratios, and time delays between multiple images - depend on the gravitational potential of the foreground galaxy (lens or deflector) and its derivatives. Secondly, the lensing observables also depend on the overall geometry of the universe via angular diameter distances between observer, deflector, and source. Thirdly, the background source often appears magnified to the observer, sometimes by more than an order of magnitude. As a result, gravitational lensing can be used to address three major astrophysical issues: i) understanding the spatial distribution of mass at kpc and sub-kpc scale where baryons and DM interact to shape galaxies as we see them; ii) determining the overall geometry, content, and kinematics of the universe; iii) studying galaxies, black holes, and active nuclei that are too small or too faint to be resolved or detected with current instrumentation. The topic of strong lensing by galaxies is too vast to be reviewed entirely in a single Annual Review Article. This article is meant to provide an overview of a selection of the most compelling and promising astrophysical applications of strong gravitational lensing at the time of this writing. The main focus is on recent results (after $\\sim$2005). For each application, I discuss the context, recent achievements, and future prospects. Of course, lensing is only one of the tools of the astronomers' trade. When needed, I discuss scientific results that rely on strong lensing in combination with other techniques. For every astrophysical problem, I also present a critical discussion of whether strong gravitational lensing is competitive with alternative tools. Excellent reviews and monographs are available to the interested reader for more details, different points of view, history of strong lensing, and a complete list of pre-2005 references. The Saas Fee Lectures by Schneider, Kochaneck and Wambsganss (2006) provide a comprehensive and pedagogical treatment of lensing fundamentals, theory and observations until 2006. Additional information can be found in the review by \\citet{Fal05} and that by \\citet{CSS02}. The classic monograph by \\citet{SEF92} and that by \\citet{PLW01} are essential references for strong gravitational lensing theory. This review is organized as follows. First, for convenience of the reader and to fix the notation and terminology, in \\S~\\ref{sec:theory} I give a very brief summary of strong lensing theory. Then, in \\S~\\ref{sec:over}, I present an overview of the current observational landscape. The following four sections cover the main astrophysical applications of gravitational lensing: ``The mass structure of galaxies'' (\\S~\\ref{sec:mass}), ``Substructure in galaxies'' (\\S~\\ref{sec:sub}), ``Cosmography'' (\\S~\\ref{sec:cosmo}), and ``Lenses as cosmic telescopes'' (\\S~\\ref{sec:telescopes}). After the four main sections, the readers left with an appetite for more results from strong gravitational lensing will be happy to learn about the many promising ongoing and future searches for more gravitational lenses described in \\S~\\ref{sec:searches}. Some considerations on the future of strong gravitational lensing - when the number of known systems should be well into the thousands -- are given in~\\S~\\ref{sec:future}. ", "conclusions": "" }, "1003/1003.0280.txt": { "abstract": "Studying the radial variation of the stellar mass function in globular clusters (GCs) has proved a valuable tool to explore the collisional dynamics leading to mass segregation and core collapse. Recently, Pasquato et al. (2009) used the mass segregation profile to investigate the presence of an intermediate-mass black hole (IMBH) in NGC\\,2298. As a relaxed cluster with a large core, M\\,10 (NGC\\,6254) is suitable for a similar investigation. In order to study the radial dependence of the luminosity and mass function of M\\,10, we used deep high resolution archival images obtained with the Advanced Camera for Survey (ACS) on board the Hubble Space Telescope (HST), reaching out to approximately the cluster's half-mass radius ($r_{hm}$), combined with deep Wide Field and Planetary Camera 2 (WFPC2) images that extend our radial coverage to more than $2 \\, r_{hm}$. From our photometry, we derived a radial mass segregation profile and a global mass function that we compared with those of simulated clusters containing different energy sources (namely hard binaries and/or an IMBH) able to halt core collapse and to quench mass segregation. A set of direct N-body simulations of GCs, with and without an IMBH of mass $1\\%$ of the total cluster mass, comprising different initial mass functions (IMFs) and primordial binary fractions, was used to predict the observed mass segregation profile and mass function. The mass segregation profile of M\\,10 is not compatible with cluster models without either an IMBH or primordial binaries, as a source of energy appears to be moderately quenching mass segregation in the cluster. Unfortunately, the present observational uncertainty on the binary fraction in M\\,10 does not allow us to confirm the presence of an IMBH in the cluster, since an IMBH, a dynamically non-negligible binary fraction ($\\sim 5\\%$), or both can equally well explain the radial dependence of the cluster mass function. ", "introduction": "Globular Clusters (GCs) are important astrophysical laboratories for the study of both stellar evolution and stellar dynamics. In recent years it has become clear that these two astrophysical processes cannot be studied independently: physical interactions between single stars as well as the formation, evolution, survival and interactions of binary systems have a significant role in the evolution of the clusters and of their stellar populations~\\citep{spi87}. Such interactions change the energy budget of the cluster and therefore influence the time scales on which mass segregation, core collapse and other dynamical processes occur. On the evolutionary side, they can generate objects that cannot be explained by standard stellar evolution \\citep[like blue stragglers, X-ray binaries, millisecond pulsars, etc.; see][and references therein]{fe06}. Recently, various clues have emerged that point to the presence of intermediate-mass black holes (IMBHs) in GCs. On the theoretical side at least three different formation mechanisms have been proposed, namely merging of stellar-mass black holes via four-body interactions \\citep{mil02}, runaway merging of massive stars \\citep{por04}, and the death of Pop III stars in the early universe, which would leave IMBHs behind as remnants, as shown e.g. by \\citet{mad01} \\citep[see also][]{tre07a}. Detailed collisional N-body simulations~\\citep{ba05,tre07b} and theoretical arguments~\\citep{he07} have shown that an IMBH can be at the origin of a shallow central cusp in the surface brightness profile (SBP) of GCs that have been observed in the projected density profile of a few GCs by~\\citet{noy06, noy07},~\\citet{la07}, and \\citet{iba09}. Line-of-sight velocity studies were undertaken by \\citet{bau03a}, \\citet{bosch06}, and \\citet{chak06} on M\\,15, by \\citet{geb02},~\\citet{bau03b},~\\citet{geb05} on G\\,1, by \\cite{noy08},~\\citet{and09},~\\citet{sol09} on $\\omega$ Cen, resulting for now in no undisputed definitive detection. \\citet{iba09} recently found a stellar density cusp and a velocity dispersion increase in the center of the globular cluster M\\,54. This could be explained by the presence of a $\\sim 9400 M_{\\odot}$ black hole if the cusp stars possess moderate radial anisotropy. Proper motion studies from Hubble Space Telescope (HST) data are likely the best way to attack the problem, but they are technologically challenging and inherently time-consuming, requiring multi-epoch HST-quality observations. Numerical simulations by \\citet{bau04} have shown that IMBHs produce a clear photometric signature in GCs, in terms of a large core to half-mass-radius ratio. Moreover, \\cite{gill08} show that a quenching of mass segregation is predicted by N-body simulations of clusters harboring an IMBH, resulting in an observable signature on the radial dependence of the stellar mass function in GCs, which can be used to either suggest or rule out the presence of an IMBH. \\cite{pas09} provide the first application of this method, using the observed mass segregation profile to argue against the presence of an IMBH in NGC\\,2298. More generally, studying the radial dependence of the stellar mass function in collisionally relaxed GCs can lead to valuable insights on the underlying dynamics of the relaxation processes, such as stellar evaporation, mass segregation and core collapse. Mass segregation occurs in star clusters as two-body relaxation processes drive the system towards energy equipartition. Heavier stars sink towards the center of the cluster, while the lighter stars preferentially live in the outer regions and are more likely to attain escape velocity from the system. If an additional energy source is present, such as a dynamically significant population of binaries or an IMBH, the mass segregation and the core collapse processes are slowed down and the arising gradient of the IMF is less steep. Also a significant population of stellar mass black holes can in principle segregate in the center of the cluster and sustain a large core \\citep[see][]{mac08}, thereby acting as an energy source. This effect is observed in our N-BODY simulations with a \\cite{sal55} IMF (see Section~\\ref{MassSeg}), which produces comparatively more remnants than a \\cite{ms79} IMF. However, in such simulations this effect typically lasts a few Gyr, because stellar mass black holes kick each other out of the system via dynamical interactions, and is therefore unlikely to be at work in M\\,10, which is approximately $11.8 \\pm 1.1$ Gyr old \\citep[][]{sala02}. In any case, a \\cite{sal55} IMF likely overestimate the number of large stellar-mass black holes produced in a real GC, so any effect of such a \\emph{dark core} on the long-term dynamics of the cluster is effectively ruled out. IMBHs and binaries both work towards reducing the degree of mass segregation attained by a GC at a given (dynamical) age. A realistic binary fraction can partly mimick the effects of an IMBH, so a detailed knowledge of the binary fraction allows more robust conclusions either in favor or against the presence of an IMBH. So far, however, only for a few GCs has a reliable binary fraction been estimated with robust photometric methods based on main sequence (MS) broadening \\citep[see][and references therein]{sol07,mil08}, so the type of study that we discuss here is only possible for a handful of clusters. M\\,10 is a reasonable candidate for harboring an IMBH since it displays a ratio of almost $0.4$ \\citep{mc05} between the core radius and the half-light radius, which in dynamically old clusters is usually interpreted as the signature of an extra energy source~\\citep[see][]{ves94}. In addition, M10 is an ideal candidate for the mass segregation method described and employed for NGC2298 by Pasquato et al. (2009) because, with a mass of about $1.5 \\times {10}^5$ $M_{\\odot}$ \\citep[][]{mc05}, it is dynamically relaxed ($\\log(t_{h})=8.86 yrs$; Harris 1996) and has not been overly influenced by the Galactic tidal field ($r_t/r_{hm}$ = 11.86 and $R_{gc}$ = 4.6 kpc; Harris 1996). These conditions ensure that it has experienced many initial half-mass relaxation times and thus has achieved equilibrium with respect to mass segregation. The structure of the paper is as follows. We present the data analysis in Sect.~\\ref{Observations} and the resulting color-magnitude diagram of M\\,10 Sect.~\\ref{sec_cmd}. The radial dependence of the mass and luminosity function is showed in Sect.~\\ref{lumin}, while in Sect.~\\ref{sec_gmf} the global mass function (GMF) of the cluster is derived. In Sect.~\\ref{MassSeg} we compare the observed mass segregation profile with predictions from our N-body simulations. A summary and conclusions follow in Sect.~\\ref{Conc}. ", "conclusions": "\\label{Conc} In this work we present a study of the dynamical state of the globular cluster M\\,10 through the characterization of the radial properties of its stellar LF and MF. We used a combination of archival deep ACS and WFPC2 images to sample cluster stars from the core regions through to $\\sim 2.5\\,r_h$ in the mass range $0.25 - 0.8$\\,M$_{\\odot}$. The full data set was divided in six regions at different distances from the cluster center and characterized by the same photometric completeness. The LF was calculated for each region and converted to a MF using the mass--luminosity relationship of \\citet{ba97} for the metallicity of the cluster. Each of the local MFs was fitted using a simple power law of the type $dN/dm \\propto m^\\alpha$ and the best fitting indices were found to be $\\alpha=0.7, 0.4, 0.1, -0.3, -0.6, -0.9$ from the center outwards. A positive value of $\\alpha$ means that the number of stars decreases with decreasing mass. This radial change in the MF index is a clear sign of mass segregation in M\\,10. To better characterise this effect, we built a multi-mass Michie--King model to reproduce the observed radial variation of the MF and the cluster's SBP and veolcity dispersion and found that the distribution of the stars in the cluster is compatible with a condition of equipartition of energy. The GMF as obtained by the Michie--King model is a power-law with index $\\alpha=-0.7$, in agreement with the value $\\alpha\\sim-0.6$ near the half mass radius ($r_{hm} = {124}^{\\prime\\prime}$). In order to investigate in more detail the dynamical history of M\\,10, we used N-body simulations that include the presence and evolution of binaries, with the goal to understand whether the current mass segregation profile requires the presence of an IMBH. Within the current uncertainties, we cannot exclude an IMBH of mass $\\sim 1\\,\\%$ of the total cluster mass, although a plausible alternative explanation of the observed mass segregation profile is the presence of a substantial amount of binaries, likely primordial in origin, capable of significant dynamical effects. We show that N-body simulations initialized with a primordial binary fraction in the range of $3-5\\,\\%$ are effectively able to correctly predict the observed mass segregation profile of M\\,10. M\\,10 is a better candidate than NGC\\,2298 for an observational follow-up aimed at detecting an IMBH. While the latter cluster does not contain an IMBH according to \\cite{pas09}, M\\,10 shows a shallow mass segregation profile which could be interpreted as the fingerprint of an IMBH. The comparison of the observed mass segregation with the simulations clearly demonstrates that the cluster binary stars play a key role in defining its dynamical state and therefore cannot be ignored. The binary fraction of the cluster needs to be better constrained." }, "1003/1003.1635_arXiv.txt": { "abstract": "Until now, there is no experimental evidence on the gravitational behaviour of antimatter. While we may be confident that antimatter attracts antimatter, we do not know anything on the interaction between matter and antimatter. We investigate this issue on theoretical grounds. Starting from the CPT invariance of physical laws, we transform matter into antimatter in the equations of both electrodynamics and gravitation. In the former case, the result is the well-known change of sign of the electric charge. In the latter, we find that the gravitational interaction between matter and antimatter is a mutual repulsion. This result supports cosmological models attempting to explain the Universe accelerated expansion in terms of a matter-antimatter symmetry. ", "introduction": "\\label{intro} The discovery of antimatter (in 1932) raised the question about its reaction to a gravitational field. Until now no clear experimental answer could be obtained, due to the weakness of gravitation compared to the electromagnetic forces governing antiparticle motion in accelerators, and, even when dealing with electrically neutral antihydrogen, due to its fast annihilation with matter. While CPT invariance of physical laws assures that antimatter is gravitationally attracted by antimatter exactly as matter by matter, no definitely convincing theoretical argument has been so far proposed to discriminate whether matter and antimatter mutually attract or repel. Here we show that the answer can be found in the equations of general relativity, when the above mentioned CPT operation is properly applied. The result is that matter and antimatter repulse each other. Besides the obvious importance of this result in terms of pure knowledge, it also supports the recent attempts to explain the observed accelerated expansion of the Universe with matter-antimatter models, which look much simpler than the $\\Lambda$-CDM model, that is based on the dominant existence of the so-called `dark energy', of unknown origin. ", "conclusions": "\\label{sec:4} This theoretical derivation of the gravitational repulsion between matter and antimatter supports cosmological models attempting to explain the observed accelerated expansion of the Universe through such a repulsion between equal amounts of the two components. Very promising appears to be the model by Benoit-L\\'evy \\& Chardin \\cite{Ref4,Ref5}. The gravitational repulsion would prevent the mutual annihilation of isolated and alternated systems of matter and antimatter. The location of antimatter could be identified with the well-known large-scale (tens of Mpc) voids observed in the distribution of galaxy clusters and superclusters. Indeed, Piran \\cite{Ref6} showed that these voids can originate from small negative fluctuations in the primordial density field, which ({\\it acting as if they have an effective negative gravitational mass}) repel surrounding matter, and grow as the largest structures in the Universe. These new cosmological scenarios could eliminate the uncomfortable presence of an unidentified dark energy, and maybe also of cosmological dark matter, which, according to the $\\Lambda$-CDM concordance model, would together represent more than the 95\\% of the Universe content. If large-scale voids are the location of antimatter, why should we not observe anything there? There is more than one possible answer, which will be investigated elsewhere." }, "1003/1003.4821_arXiv.txt": { "abstract": "In this paper we continue our study on the accretion process onto super-spinning Kerr objects with no event horizon (super-spinars). We discuss the counterpart of the Bondi accretion onto black holes. We first report the results of our numerical simulations. We found a quasi steady-state configuration for any choice of the parameters of our model. The most interesting feature is the presence of hot outflows. Unlike jets and outflows produced around black holes, which are thought to be powered by magnetic fields and emitted from the poles, here the outflows are produced by the repulsive gravitational force at a small distance from the super-spinar and are ejected around the equatorial plane. In some circumstances, the amount of matter in the outflow is considerable, which can indeed significantly reduce the gas mass accretion rate. Finally, we discuss a possible scenario of the accretion process in more realistic situations, which cannot be simulated by our code. ", "introduction": "The actual nature of the final product of the gravitational collapse of matter is an outstanding and longstanding problem in general relativity (GR). Under apparently reasonable assumptions, collapsing matter leads to the formation of space-time singularities~\\cite{hawking}. Here there are two possibilities: $i)$ the singularity is hidden behind an event horizon and the final product is a black hole (BH), or $ii)$ the singularity is naked. Since space-times with naked singularities present pathologies, some form of the Cosmic Censorship Conjecture is usually assumed and naked singularities are forbidden~\\cite{penrose}. Neglecting the electric charge, it turns out that, in four dimensions, the only stationary and asymptotically flat solution of the vacuum Einstein equation with an event horizon is the Kerr BH~\\cite{carter, robinson}. The Kerr metric is completely characterized by two parameters; that is, the mass $M$ and the spin $J$. The latter is often replaced by the Kerr parameter $a$, defined as $a = J/M$, or by the dimensionless Kerr parameter $a_* = a/M$. Using Boyer-Lindquist coordinates, the position of the horizon of a Kerr BH is given by \\be r_H = M \\left(1 + \\sqrt{1 - a^2_*}\\right) \\, , \\ee which demands the well known constraint $|a_*| \\le 1$. For $|a_*| > 1$, there is no horizon and the space-time has a naked singularity. Since causality can be violated in Kerr space-times with no event horizon~\\cite{carter2, chandra}, the assumption of the Cosmic Censorship Conjecture appears well motivated. Interestingly, several solutions of the Einstein equations that lead to the formation of naked singularities are known (see e.g. Refs.~\\cite{chr84, piran, joshi93, chr94, joshi94, joshi98}). Moreover, the reliability of GR at the singularity is surely questionable. GR is a classical theory and it is widely believed that the Planck scale, $E_{Pl} \\sim 10^{19}$~GeV, is its natural UV cut-off. New physics needs to replace the singularity with something else~\\cite{nakao}; then the space-time in the full theory may present no pathologies. If this is the case, the Cosmic Censorship Conjecture may not be necessary and super-spinning Kerr objects with no event horizon, or ``super-spinars'', may exist in the Universe~\\citep{horava}. A subtle point is the stability of these objects, but the issue is not so easy to address, especially because we do not know the actual space-time emerging from the full theory. The simplest kind of instability is represented by the process of accretion itself, which may reduce $|a_*|$ and convert the super-spinar into a BH~\\cite{defelice}. However, the evolution of the Kerr parameter due to accretion is usually negligible in stellar mass objects in binary systems and, as discussed in this paper, the accretion onto super-spinars is strongly suppressed, because of the presence of powerful outflows. A different issue is the intrinsic stability of the space-time. A ``true'' Kerr naked singularity is {\\it probably} unstable~\\cite{dotti06, dotti08}, but such a conclusion is based on an analysis of the whole Kerr space-time with $- \\infty < r < + \\infty$, which includes ``another Universe'', connected to ``our Universe'' through the singularity and admitting closed time-like curves. Another danger is represented by the so-called ergo-region instability~\\cite{cardoso1, cardoso2}. The stability of the space-time is, however, mainly determined by the boundary conditions, which are unknown in our case. The motivation to consider super-spinars is that high energy corrections to classical GR may replace the singularity with something else and that the space-time has no pathologies in the full theory. Since we do not know the full theory, we cannot predict the actual structure of the space-time at very small radii. Here we do not further discuss this point: we simply assume that super-spinars are stable or quasi-stable objects. We study their astrophysical implications. Let us also notice that, even if super-spinars were unstable, they may still be a possible intermediate state of the gravitational collapse, before decaying into spinning BHs after some characteristic time scale. Their astrophysical implications may still be very intriguing. It is thus interesting to look for observational signatures capable of distinguishing super-spinars from ordinary BHs, with present and future astrophysical experiments. Refs.~\\cite{bf09, bft09} discussed the implications on the apparent shape. There it was found that, even if the bound $|a_*| \\le 1$ is violated by a small amount, the shadow cast by the super-spinar (i.e. how it blocks light coming to us from an object behind it) changes significantly from the BH case: the shadow for the super-spinar is about an order of magnitude smaller as well as distorted. This distinction can be used as an observational signature in the search for these objects. Based on recent observations at mm wavelength of the super-massive BH candidate at the center of the Galaxy~\\cite{doeleman}, in~\\cite{bf09} one of us speculated on the possibility that it might violate the Kerr bound. The X-ray thermal spectrum emitted by an optically thick and geometrically thin accretion disk around a super-spinar was instead investigated in~\\cite{th10}. Surprisingly, for any BH with Kerr parameter $a_*^{BH}$, there is a super-spinar with Kerr parameter $a_*^{SS}$ in the range $[5/3 \\, ; \\, 8\\sqrt{6}/3]$ whose spectrum is very similar (and practically indistinguishable) from the one of the BH. The result is that the X-ray thermal spectrum observed from current BH candidates cannot be used to confirm the Kerr bound $|a_*| \\le 1$. However, the contrary is {\\it not} true: the X-ray thermal spectrum of a super-spinar with $1 < |a_*| < 5/3$ can be significantly different from the one of BH. The disk temperature is higher and the radiation flux larger, even by a few orders of magnitude for $|a_*|$ slightly larger than 1. In principle, this could be used to distinguish super-spinars from BHs. In this paper, we extend the study on the accretion process onto super-spinars started in~\\cite{bfhty09}. Here we consider the counterpart of the Bondi accretion onto BHs. In the case of BHs, this is a quite inefficient mechanism to convert the gravitational energy into radiation, because a large fraction of the thermal energy is lost behind the horizon. The efficiency is $\\eta \\sim 10^{-4}$ and thus the Bondi accretion cannot explain many of the observed BH candidates. In the case of super-spinars, the efficiency can be much higher. As shown in ~\\cite{bfhty09}, near super-spinars there are space regions where the gravitational force is repulsive. In the Bondi-like case, the gas approaches the compact object with high velocity, enters the region with repulsive gravitational force, and slows down. Some amount of gas, which depends on the spin and the actual size of the super-spinar, is eventually ejected away, preferably on the equatorial plane. Unlike the case of an accretion flow of low temperature and low velocity discussed in~\\cite{bfhty09}, here the gas reaches the center preferably from the axis of symmetry (rather than from the equatorial plane) and we find the production of hot outflows, which can have quite interesting implications. Since outflows produced in the accretion process onto BHs are expected to be emitted from their poles, the possible observation of powerful equatorial outflows from a BH candidate could be evidence that such an astrophysical object is actually a super-spinar or, more in general, a spinning super-compact object with no event horizon. The paper is organized as follows. In Sec.~\\ref{s-pworks}, we briefly summarize the results obtained in~\\cite{bfhty09}. In Sec.~\\ref{s-bondi}, we present the results of our simulations of the Bondi-like accretion onto super-spinars. In Sec.~\\ref{s-disc}, we discuss the role of the free parameters of our model and how their variation changes the accretion process. On the basis of such numerical study, in Sec.~\\ref{s-astro} we discuss a possible scenario in more realistic cases, mentioning physical and astrophysical implications. Summary and conclusions are reported in Sec.~\\ref{s-concl}. Throughout the paper we use Boyer-Lindquist coordinates to describe the Kerr background and natural units $G_N = c = k_B = 1$. ", "conclusions": "} In this paper, we have investigated the counterpart of the Bondi accretion in the case of super-spinars. Our numerical simulations show that the repulsive gravitational force in the neighborhood of the center can produce powerful outflows of gas, suppressing the accretion process. It is important to notice that the gravitational force becomes repulsive already at radii $r \\sim M$ where, for $M \\gg M_{Pl}$, classical GR is presumably reliable. The absence of the horizon can thus lead to the interesting astrophysical implications discussed in this work. In our simulations, the accretion depends on three parameters: the maximum temperature of the gas, $T_{max}$, the radius of the inner boundary, $r_{in}$, and the dimensionless Kerr parameter, $a_*$. We have studied numerically how the variation of $T_{max}$, $r_{in}$, and $a_*$ changes our results. Then, we have tried to discuss what can happen in more realistic circumstances that cannot be simulated by the code. Assuming that the super-spinar can be considered as a compact object with a finite radius and an absorbing surface, and that even in its neighborhood deviations from the classical Kerr metric are not significant, we suggest the following accretion scenario. Only a small fraction of the accreting gas can really reach the surface of the super-spinar and be absorbed, since at smaller and smaller radii the gas velocity decreases, due to two effects: the gravitational force is repulsive and the gas pressure becomes higher and higher. Most of the accreting gas is pushed back to larger radii. Some blobs of hot gas are ejected away with high energy and can escape from the gravitational well of the super-spinar. The rest of the gas is instead cooled and gravitationally bound to the massive object. In this way, the cloud of the accreting gas surrounding the super-spinar cools, less and less gas can enter the region with repulsive gravitational force, and eventually the formation of hot outflows stops. We argue that in this case we recover the accretion process discussed in~\\cite{bfhty09}: the absence of hot outflows let the gas reach the center from the equatorial plane (for $|a_*| \\ge 1.4$, since in this case the gravitational force is everywhere attractive on the equatorial plane) and then be absorbed by the surface of the super-spinar. For astrophysical black holes, it is widely believed that jets and outflows are powered by magnetic fields (or maybe even by radiation for high mass accretion rates) and are expected to be ejected from the poles. It is thus likely that the possible observation of powerful equatorial outflows from a black hole candidate is a strong indication that the object is not a black hole, but a super-spinar or, more in general, a spinning super-compact object with no event horizon. Regions with repulsive gravitational force with features similar to the ones shown in Fig.~\\ref{f-geod} seem indeed common in space-times describing the gravitational field of spinning masses and containing naked singularities. At least during the Bondi-like accretion phase, super-spinars should be objects much brighter than usual black holes, because the gas is not quickly lost behind the horizon, but orbits around it. In this phase, there might also be the possibility of producing exotic matter: scattering of the particles near the center, where the density and the temperature can be very high, may produce heavy particles. Another intriguing possibility deserving some attention is that super-spinars might work as the central engine for long GRBs. Indeed, they seem to be able to produce collimated jets with a high Lorentz factor." }, "1003/1003.1259_arXiv.txt": { "abstract": "\\small{ In this work, we study a class of early \\de~(EDE) models, in which, unlike in standard \\de~models, a substantial amount of \\de~exists in the matter-dominated era. We self-consistently include \\de~\\perts, and constrain these models using current observations. We consider EDE models in which the \\de~equation of state is at least $\\wm \\ggeq -0.1$ at early times, which could lead to a early \\de~density of up to $\\Omega_{DE} (z_{CMB})= 0.03 \\om (z_{CMB})$. Our analysis shows that, marginalizing over the non-\\de~parameters such as $\\om, H_0, n_s$, current CMB observations alone can constrain the scale factor of transition from early \\de~to late time \\de~to $\\at \\ggeq 0.44$ and width of transition to $\\dt \\lleq 0.37$. The equation of state at present is somewhat weakly constrained to $\\w \\lleq -0.6$, if we allow $H_0 < 60$ km/s/Mpc. Taken together with other observations, such as supernovae, HST, and SDSS LRGs, $\\w$ is constrained much more tightly to $\\w \\lleq -0.9$, while redshift of transition and width of transition are also tightly constrained to $\\at \\lleq 0.19, \\dt \\lleq 0.21$. The evolution of the equation of state for EDE models is thus tightly constrained to $\\ld$CDM-like behaviour at low redshifts. Incorrectly assuming \\de~\\perts~to be negligible leads to different constraints on the equation of state parameters-- $\\w \\lleq -0.8, \\at \\lleq 0.33, \\dt \\lleq 0.31$, thus highlighting the necessity of self-consistently including \\de~\\perts~in the analysis. If we allow the spatial curvature to be a free parameter, then the constraints are relaxed to $\\w \\lleq -0.77, \\at \\lleq 0.35, \\dt \\lleq 0.35$ with $-0.014 < \\omk < 0.031$ for CMB+other observations. For perturbed EDE models, the $2\\sigma$ lower limit on $\\sig$ ($\\sig \\geq 0.59$) is much lower than that in $\\ld$CDM ($\\sig \\geq 0.72$), thus raising the interesting possibility of discriminating EDE from $\\ld$CDM using future observations such as halo mass functions or the Sunyaev-Zeldovich power spectrum. } ", "introduction": "Over the last decade, the unexpected faintness of distant Type Ia supernovae have shown that the expansion of the universe is accelerating at present \\citep{accl1, accl2, accl3, accl4, accl5, accl6, accl7, accl8, const, sdss_sne}. This remarkable discovery points to the existence of \\de~(DE), a negative pressure energy component which dominates the energy content of the universe at present. Other, complementary, probes such as the Cosmic Microwave background (CMB) and various large scale structure surveys have also confirmed the existence of this mysterious component of energy \\cite{bao, wmap5, lrg}. Several theories have been propounded to explain this phenomenon, the simplest of which is the cosmological constant $\\ld$, with a constant energy density and a constant equation of state $w=-1$. The cosmological constant is fit well by the current data \\citep{const}, however, there are no strong constraints on the time evolution of \\de~at present. Thus, evolving models of \\de~remain viable as alternative candidates for dark energy. Many non-cosmological constant phenomenological explanations for cosmic acceleration have been suggested \\citep[see reviews][and references therein]{rev1, rev2, rev3, rev4, rev5, rev6, rev7}. These are based either on the introduction of new physical fields (quintessence models, Chaplygin gas, etc.), or on modifying the laws of gravity and therefore the geometry of the universe (scalar-tensor gravity, $f(R)$ gravity, higher dimensional `Braneworld' models \\etc). As of now, there is no consensus on the true nature of \\de. An interesting class of models which have been suggested in the literature are early \\de~models, a class of \\de~in which the early universe contained a substantial amount of \\de. These models were studied theoretically in \\citep{ede1,ede2,ede3} and references therein, and have been analyzed with respect to observations extensively in recent times in \\citep{ede4,ede5,ede6,ede7,ede8}. For now, there are no strong observational constraints on the EDE models, and it is especially difficult to discriminate EDE models which have $w = -1$ at present from the $\\ld$CDM model of \\de. In this work we use a parameterization of the equation of state of \\de~to study and constrain EDE models using the currently available data. We attempt to see if bounds can be put on the transition from early \\de~to the present day \\de~content of the universe. Section~\\ref{meth} explains the methods and data used for this analysis, section~\\ref{res} shows the results, and in section~\\ref{concl} we conclude. ", "conclusions": "In this work, we have studied early \\de~models using current observations. We find that, if a sizeable amount of \\de~exists in early times ($\\omde (z_{CMB}) \\simeq 0.03 \\om (z_{CMB})$), we may put tight constraints on the transition of this \\de~to its present day value, and that the present day value of the \\de~equation of state must be close to the $\\ld$CDM value. If the \\de~\\perts~are correctly accounted for, then the current \\de~equation of state is constrained to $\\w < -0.89$, while the transition from early \\de~must occur at redshifts of $z_t > 4.2$, with a narrow transition width of $\\dt < 0.21$. Incorrectly assuming that \\de~\\perts~are negligible leads to a different result-- $\\w < -0.8, z_t > 2, \\dt > 0.31$, thus showing that it is vital to include the \\de~\\perts~self-consistently in any analysis that uses perturbative data such as CMB or the matter power spectrum. Leaving $\\omk$ to be a free parameters leads to a weakening of the constraints on the \\de~parameters, with $\\w < -0.77, \\at < 0.35$ (\\ie~$z_t \\ggeq 2$), $\\dt < 0.35$ for $-0.014 < \\omk < 0.031 $. We note that, for the flat universe in which \\de~\\perts~are considered, the value of $\\sig$ is much lower than that in corresponding $\\ld$CDM models. As will be shown in a companion paper \\cite{ede_lss}, this may lead to interesting constraints from future large scale structure data such as halo mass functions, as also from the Sunyaev-Zeldovich power spectrum." }, "1003/1003.3540_arXiv.txt": { "abstract": "An empirical model has been developed to reproduce the drift of the spectrum recorded by the EIS on {\\it Hinode} using instrumental temperatures and relative motion of the spacecraft. The EIS spectrum shows an artificial drift in wavelength dimension in sync with the revolution of the spacecraft, which is caused by temperature variations inside the spectrometer. The drift amounts to $70$~km~s$^{-1}$ in Doppler velocity and introduces difficulties in velocity measurements. An artificial neural network is incorporated to establish a relationship between the instrumental temperatures and the spectral drift. This empirical model reproduces observed spectrum shift with an rms error of 4.4~km~s$^{-1}$. This procedure is robust and applicable to any spectrum obtained with EIS, regardless of of the observing field. In addition, spectral curvatures and spatial offset in the north -- south direction are determined to compensate for instrumental effects. ", "introduction": "The EUV Imaging Spectrometer (EIS; \\opencite{culhane2007}, \\opencite{korendyke2006}, \\opencite{lang2006}) on {\\it Hinode} \\cite{kosugi2007} is a powerful instrument for plasma diagnostics in the solar corona. The spectrometer can obtain a lot of EUV spectral lines emitted from a wide temperature range of plasma and allows us to study the dynamics of the solar corona. The spectrometer is capable of measuring spectrum shift in prominent emission lines with an accuracy better than 3~km~s$^{-1}$. However, one of the challenges in EIS data analysis is that the spectrum shows a quasi-periodic shift in the wavelength dimension. Since the shift is synchronized with the spacecraft revolution around the Earth, it must be an instrumental effect. The artificial spectral drift amounts to $70$~km~s$^{-1}$ in Doppler velocity scale and introduces a significant effect on velocity measurements. A commonly used method so far is to assume that a net Doppler shift vanishes in a quiet region. As a quiet region is not always included in the EIS field of view, the assumption of zero net velocity is not applicable in some cases. It has been reported that the spectrum drift is connected with temperature variations in the spectrometer. \\inlinecite{brown2007} showed that the spectral drift is correlated with grating temperature variation. But their relationship is not simple because a slight phase difference is found between them. A qualitative comparison between the spectrum drift and instrumental temperatures shows that the spectral drift reflects all aspects of the temperature variations. \\inlinecite{rybak1999} demonstrated that periodic spectral shifts of the SUMER spectrometer can be corrected by using instrumental temperatures. However, a more complex method is needed to reproduce the spectral drifts of the EIS. Our goal is to develop an empirical model to reproduce spectral drift from instrumental temperatures so that the Doppler velocity can be determined from any spectra obtained with EIS. The organization of the paper is as follows: The scheme of the model and definition of their input and output are described in Section 2. A comparison between the model and the measured spectral position is presented in Section 3. In Section 4, the performance of the model and its application to observations are discussed. The curvatures and spatial offset of the spectrum, which are essential for data analysis with EIS, are described in the Appendices. ", "conclusions": "An empirical method for velocity calibration is constructed using data from three years of observations with the EIS. It estimates instrumental spectral drift from temperatures inside the spectrometer and relative motion of the spacecraft. The model reproduces observed spectral position on the detector with a moderate accuracy of 0.13 pixel or 4.4~km~s$^{-1}$ at the Fe~{\\sc xii} $\\lambda$195.12~\\AA~ emission line. It proves that the spectral drift, which introduces difficulties in velocity measurement with EIS, is primarily caused by temperature variations inside the spectrometer. This model works in the eclipse period, when the spacecraft experiences severe temperature variation in the day -- night cycle, as well as in the day period. Since our new method requires only spacecraft status data, it is applicable to any spectrum recorded by the EIS. It can be applied to both the short and the long wavelength bands, as the instrumental effect produces the same amount of spectral drift in the entire wavelength range of the EIS. The correction method would be particularly useful for active region and flare observations since no velocity reference is needed. This procedure is provided as a part of the SSW. \\appendix" }, "1003/1003.1897_arXiv.txt": { "abstract": "We study the velocity field of umbral dots (UDs) at a resolution of 0\\farcs14. Our analysis is based on full Stokes measurements of a pore taken with the Crisp Imaging Spectro-Polarimeter at the Swedish 1 m Solar Telescope. We determine the flow velocity at different heights in the photosphere from a bisector analysis of the \\ion{Fe}{1} 630~nm lines. In addition, we use the observed Stokes $Q$, $U$, and $V$ profiles to characterize the magnetic properties of these structures. We find that most UDs are associated with strong upflows in deep photospheric layers. Some of them also show concentrated patches of downflows at their edges, with sizes of about 0\\farcs25, velocities of up to 1000~m~s$^{-1}$, and enhanced net circular polarization signals. The downflows evolve rapidly and have lifetimes of only a few minutes. These results appear to validate numerical models of magnetoconvection in the presence of strong magnetic fields. ", "introduction": "At high angular resolution, sunspot umbrae exhibit small bright features embedded in a darker, smoothly varying background. These features are called umbral dots (UDs) and occur in essentially all sunspots but also in pores \\citep[for a review, see][]{sobotka97}. Frequently a distinction is made between central and peripheral UDs based on their location within the umbra \\citep[e.g.][]{gd86,rieth08a,sob09}. Central UDs do not show measurable vertical flows and have similar magnetic field inclinations as the surrounding umbra. In contrast, upflows and more horizontal magnetic fields are detected in peripheral UDs. The importance of UDs lies in the fact that they may be the signature of convection in sunspots. We know that convective motions must exist because radiation alone cannot explain the brightness of the umbra \\citep{deinzer65}. The question is whether these motions are field free or magnetized. Recent simulations of magnetoconvection in a strong background field \\citep{simu06} suggest that the convective energy transport inside the umbra is achieved in the form of narrow plumes of rising hot plasma and strongly reduced magnetic fields. The plumes are associated with bright structures that share many similarities with real UDs, including their sizes (about 300 km) and lifetimes (30 minutes). The simulated UDs tend to be elliptical in shape, and many have a central dark lane. According to \\cite{simu06}, the pile-up of material at the top of the plumes increases the gas density and moves the $\\tau=1$ level toward higher layers, where the temperature is lower. This produces the dark lanes of the simulated UDs. The upflows attain maximum velocities of 3-4 km~s$^{-1}$ just below the solar surface. Near $\\tau=1$ they turn horizontal and move in the direction of the dark lane, until the gas returns to deeper layers along narrow channels. The downflows occur mainly at the endpoints of the dark lanes, but also in the immediate surroundings of the UDs. They usually appear as tiny patches of concentrated flows with velocities of up to 1200 m\\,s$^{-1}$ at $z = 0$ km. The velocities at $\\tau = 1$ are smaller because the large opacity of the gas hides the layers where the flows are stronger \\citep{bhar10}. The simulations of \\cite{simu06} suggest that UDs are a natural consequence of convection in a strong magnetic field, which seems to favor the monolithic sunspot model over the cluster model. However, the basic physical process at work in the simulations, namely overturning convection, has not yet been confirmed observationally. Upflows are known to exist in UDs; what is missing is a clear, unambiguous detection of return downflows. The velocity field of UDs has been studied by a number of authors. Most of them report negligible velocities in central UDs \\citep[e.g.,][]{schmidt94, rim08, rieth08a,sob09} and upflows ranging from 100 m s$^{-1}$ to 1000 m\\,s$^{-1}$ in peripheral UDs \\citep[e.g.,][]{hector04,rim04,rieth08a, sob09}. \\citet{bar07} claim to have measured upflows of 400 m\\,s$^{-1}$ and downflows of 300 m\\,s$^{-1}$ in UDs, using the Universal Birrefringent Filter at the Dunn Solar Telescope. This is a remarkable result which might have been compromised by the fact that no adaptive optics, phase diversity, or any postprocessing technique such as speckle reconstruction were available at the time of the observations. In any case, the downflows described by \\citet{bar07} seem to be different from those reported by \\citet{simu06} in that they extend over much larger areas. The same happens to the downflow patches detected by \\cite{bar09} using {\\it Hinode} observations. Interestingly, the spectropolarimetric measurements of \\cite{rieth08a} and \\cite{sob09}, also from {\\it Hinode}, do not show downward motions despite their better angular resolution (0\\farcs3 versus 0\\farcs6). To detect the downflows predicted by the simulations, one would like to have extremely high spatial resolution --higher than typically achieved by current instruments-- and line spectra to probe the velocity field at different heights in the atmosphere. Until now it has been difficult to fulfill the two requirements simultaneously. Here, we present the first spectropolarimetric observations of UDs approaching a resolution of 0\\farcs1. This unique data set, acquired at the Swedish 1 m Solar Telescope (SST), allows us to investigate the flow field and magnetic properties of UDs with unprecedented detail. The measurements show localized patches of downflows at the edges of UDs, providing firm evidence for magnetoconvection in sunspot umbrae. The paper is organized as follows. The observations and the data reduction are described in Section~\\ref{data}. Section \\ref{results} deals with the morphology, flow field, temporal evolution, and magnetic properties of UDs as derived from the observed Stokes spectra. Finally, in Section~\\ref{disc} we discuss our results and compare them with previous works. ", "conclusions": "\\label{disc} We have presented the first spectropolarimetric measurements of UDs at a resolution of 0\\farcs14. Our observations reveal the existence of substructures within UDs in the form of dark lanes. Only \\citet{bhar07b}, \\citet{rim08}, and \\citet{sobpus09} have detected these structures before. \\citet{rim08} estimated their size to be 0\\farcs12 in G-band images, right at the diffraction limit of the Dunn Solar Telescope, while \\citet{sobpus09} reported widths of less than 0\\farcs14 from broadband 602~nm filtergrams taken at the SST. Our observations also indicate sizes of 0\\farcs14. At the resolution of CRISP, however, not all UDs possess a dark lane. In this paper, we have focused on the velocity field of UDs. One can find in the literature a variety of works describing the morphological properties of UDs, but very few attempts have been made to determine their velocities. However, very few attempts have been made to determine their velocities. A good knowledge of the flow field of UDs is important both to understand the energy transport in strongly magnetized media and to discern between models of sunspot structure. Measuring UD velocities is a difficult task, as recognized by several authors \\citep[e.g.,][]{rim04,hector04,simu06}. Our results, and those from previous works, point to the existence of large LOS velocities and reduced magnetic field strengths, but only in deep layers of the photosphere --near the continuum forming region-- definitely deeper than the layers traced by the core of many photospheric lines. For example, \\citet{rim04} finds upflows of more than 1 km s$^{-1}$ in the \\ion{C}{1} 538.0~nm line (whose core forms at approximately 40 km), while he observes velocities of less than 300 ms$^{-1}$ in \\ion{Fe}{1} 557.6~nm (formation height of around 320~km). \\citet{simu06} also mention the difficulty of observing flows and reduced field strengths in UDs because the surfaces of equal optical depth are locally elevated. Other factors such as the small sizes and the rapid evolution of these structures complicate the determination of their velocities and magnetic fields. Observational studies have found upflows of 100 m\\,s$^{-1}$ \\citep{hector04}, 600 m\\,s$^{-1}$ \\citep{sob09}, 800 m\\,s$^{-1}$ \\citep{rieth08a}, and even 1000 m\\,s$^{-1}$ \\citep{rim04} in peripheral UDs. Our data confirm the existence of relatively strong upflows of around 1000 m\\,s$^{-1}$, with peaks up to 1500 m\\,s$^{-1}$. The upward velocities observed in MHD simulations do not exceed 2000 m\\,s$^{-1}$ at photospheric levels \\citep{bhar10}. The most relevant and novel result of this paper is the finding of {\\em downflows} at the edges of some UDs. If the UD has a dark lane, then the downflows are observed at its endpoints. The downflow velocities range from 400 m\\,s$^{-1}$ to almost 1000 m\\,s$^{-1}$. The existence of downflows seems to fit the scenario proposed by \\citet{simu06}, in which UDs are the result of magnetoconvection in regions of reduced field strengths with both upflows and return downflows. We believe this is the first time that the predicted downflows are reliably measured within UDs, mostly as a consequence of the high resolution provided by CRISP and the SST. We have also examined the LOS velocities at different intensity levels. Since the 80$\\%$ intensity level shows the strongest velocities and motions at other intensity levels are gradually of lesser magnitude, a gradient of velocity with height appears to exist. This suggests that the flows occur preferentially in the deep layers of the photosphere. A comparison between our Figure~\\ref{fig2}(f) and Figure 1(b) of \\citet{simu06} reveals some similarities, like the fact that the upflows are co-located with the centers of the UDs and their dark lanes. Also in good agreement is the fact that the downflows are observed around the UD structure, but most prominently at the endpoints of the dark lanes. However, our measured velocities appear to be slightly larger than those predicted by the simulations \\citep[see, e.g., Figure~16 of][]{bhar10}. In the same manner, the temporal evolution shown in our Figure~\\ref{fig3} may be compared with Figure 5 of \\citet{simu06}, where the highly dynamic and transient nature of the UDs can be seen. We have also determined the magnetic properties of UDs. New in this work is the description of the properties of their substructure, the dark lanes. UDs exhibit weaker magnetic fields and more inclined field lines than the surrounding umbra, confirming previous analyses \\citep[e.g.][]{hector04,rieth08a,sob09}. We report a magnetic field weakening of up to 500~G between UDs and their adjacent umbral background, similar to the values given by \\citet{rieth08a}. According to our data, the magnetic field in the UDs is more horizontal than in the immediate surroundings by as much as $\\sim$20$^{\\circ}$. Before this work, the magnetic properties of the dark lanes were largely unknown. We report that they exhibit an even weaker field than their associated UDs, up to 500~G less. The dark lanes also show smaller linear and circular polarization signals. Their inclinations, however, do not differ from those of the associated UDs. Another new result is the existence of enhanced NCP in UDs. We observe tiny patches of NCP that tend to coincide with the downflowing areas. In those regions, the NCP sign indicates a reduction of the field strength towards deeper photospheric layers. This again appears to confirm the results of numerical simulations \\citep{simu06}." }, "1003/1003.0858_arXiv.txt": { "abstract": "We present the results of a 5--8~$\\mu$m spectral analysis performed on the largest sample of local ultraluminous infrared galaxies (ULIRGs) selected so far, consisting of 164 objects up to a redshift of $\\sim$0.35. The unprecedented sensitivity of the Infrared Spectrograph onboard \\textit{Spitzer} allowed us to develop an effective diagnostic method to quantify the active galactic nucleus (AGN) and starburst (SB) contribution to this class of objects. The large AGN over SB brightness ratio at 5--8~$\\mu$m and the sharp difference between the spectral properties of AGN and SB galaxies in this wavelength range make it possible to detect even faint or obscured nuclear activity, and disentangle its emission from that of star formation. By defining a simple model we are also able to estimate the intrinsic bolometric corrections for both the AGN and SB components, and obtain the relative AGN/SB contribution to the total luminosity of each source. Our main results are the following: \\\\ 1) The AGN detection rate among local ULIRGs amounts up to 70 per cent, with 113/164 convincing detections within our sample, while the global AGN/SB power balance is $\\sim$1/3. \\\\ 2) A general agreement is found with optical classification; however, among the objects with no spectral signatures of nuclear activity, our IR diagnostics find a subclass of \\textit{elusive}, highly obscured AGN. \\\\ 3) We analyse the correlation between nuclear activity and IR luminosity, recovering the well-known trend of growing AGN significance as a function of the overall energy output of the system: the sources exclusively powered by star formation are mainly found at $L_\\mathit{IR}<10^{12.3} L_\\odot$, while the average AGN contribution rises from $\\sim$10 to $\\sim$60 per cent across the ULIRG luminosity range. \\\\ 4) From a morphological point of view, we confirm that the AGN content is larger in compact systems, but the link between activity and evolutionary stage is rather loose. \\\\ 5) By analysing a control sample of IR-luminous galaxies around $z \\sim 1$, we find evidence for only minor changes with redshift of the large-scale spectral properties of the AGN and SB components. This underlines the potential of our method as a straightforward and quantitative AGN/SB diagnostic tool for ULIRG-like systems at high redshift as well, and hints to possible photometric variants for fainter sources. ", "introduction": "Ultraluminous infrared galaxies (ULIRGs, $L_\\mathit{IR} \\sim L_\\mathit{bol} > 10^{12} L_\\odot$) emit the bulk of their energy at 8--1000~$\\mu$m, and are the most luminous among the local sources (Sanders \\& Mirabel 1996; Lonsdale, Farrah \\& Smith 2006). Moreover, their high-redshift counterparts represent a key component of the early Universe (Blain et al. 2002; Caputi et al. 2007). Providing a comprehensive picture of the neighbouring population is therefore a task whose implications are manifold and far-reaching. In the last years a consensus view has rapidly grown about ULIRGs, thanks to the multiwavelength approach to their study. Most of these systems are the result of a major merger (e.g. Kim, Veilleux \\& Sanders 2002; Dasyra et al. 2006a), that may eventually lead to the formation of elliptical galaxies with moderate mass (Dasyra et al. 2006b). According to numerical simulations, the tidal interaction between the progenitors drives an inflow of gaseous material that triggers and feeds both intense star formation and black hole accretion (Mihos \\& Hernquist 1996; Springel, Di Matteo \\& Hernquist 2005). The primary radiation field from the starburst (SB) and active galactic nucleus (AGN) components is reprocessed by the surrounding dust, giving rise to the huge infrared (IR) emission of ULIRGs. Anyway, the great opacity of the nuclear environment hinders a clear identification of the underlying power source, and the detection of faint or highly obscured AGN components inside ULIRGs has always been a major challenge. The most effective way to address this matter is to search for relic signatures of the buried engine within the dust emission itself. The advent of the Infrared Spectrograph (IRS; Houck et al. 2004) onboard the \\textit{Spitzer Space Telescope} (Werner et al. 2004) has opened a new era in the study of ULIRGs, providing access to a wealth of diagnostic tools at $\\sim$5--35~$\\mu$m, such as the [Ne~\\textsc{v}] $\\lambda$14.32 high-ionization line, the 9.7~$\\mu$m silicate absorption, the complex of polycyclic aromatic hydrocarbon (PAH) features at 6.2--11.3~$\\mu$m, and the 6, 15 and 30~$\\mu$m continuum colours (e.g. Farrah et al. 2007; Spoon et al. 2007; Veilleux et al. 2009a). In the end, the existence of a sizable population of AGN which are elusive not only in the optical (Maiolino et al. 2003) but also at mid-IR wavelengths has been safely ruled out. \\\\ Thanks to the high accuracy of the \\textit{Spitzer}-IRS observations, we have developed an AGN/SB decomposition method based on 5--8~$\\mu$m rest-frame spectroscopy, that has been successfully tested on a representative sample of local ULIRGs (Nardini et al. 2008, 2009; hereafter Paper~I and Paper~II, respectively). At 5--8~$\\mu$m an accurate determination of the AGN and SB components can be obtained by means of spectral templates: the large difference between the average spectral properties of SB galaxies and AGN, along with the little dispersion within the separate classes (e.g. Brandl et al. 2006; Netzer et al. 2007), makes this wavelength range very favourable to solid ULIRG diagnostics. We are therefore able to unveil even faint or obscured nuclear activity, and to assess its contribution to the luminosity of each source. In the present work we investigate the properties of the AGN population within the largest sample of local ULIRGs studied so far, consisting of 164 objects and extending to redshifts up to $z \\simeq 0.35$. This allows us to better constrain the correlation between nuclear activity and extreme IR emission. In fact, the incidence of black hole accretion on the luminosity of ULIRGs is known to increase with the total energy output of the system, from both optical classification (Veilleux, Kim \\& Sanders 1999) and early mid-IR spectroscopy (e.g. Tran et al. 2001). Such a trend is indeed suggested to involve also the IR systems at lower luminosities, as recently confirmed by many \\textit{Spitzer}-based studies (e.g. Imanishi 2009; Valiante et al. 2009). We also take into account a compilation of IR-luminous galaxies at $z \\sim 1$, extracted from different samples in the literature, in order to test our method against the possible evolution with redshift of the AGN/SB templates and bolometric corrections (i.e. the ratios between the 5--8~$\\mu$m and the 8--1000~$\\mu$m luminosities). This paper is arranged as follows: in the first part (Sections~2--3) we describe the selection of our local sample and the data reduction. In Section~4 we briefly summarize the main steps of our diagnostic method, introducing the model and computing the AGN and SB bolometric corrections from which the relative AGN/SB contribution can be derived. We then discuss how our results improve the optical classification and fit into the question of the growing AGN contribution with IR luminosity (Section~5). The application of our method to high-redshift sources and possible, alternate variants are dealt with in Section~6, while the conclusions are drawn in Section~7. Throughout this paper we have made use of the concordance cosmology from the \\textit{Wilkinson Microwave Anisotropy Probe} (\\textit{WMAP}) sky survey, with $H_0=70.5$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_m=0.27$ and $\\Omega_\\Lambda=0.73$ (Hinshaw et al. 2009). ", "conclusions": "The \\textit{Spitzer}-IRS unprecedented sensitivity allowed a deeper investigation of the role of supermassive black hole accretion and intense star formation as the engine underlying extreme IR activity. In particular, the 5--8~$\\mu$m rest-frame wavelength range has proven to be very suitable for solid diagnostics of ULIRGs: the large difference in this band between the typical spectral signatures of SB galaxies and AGN makes it possible to fully characterize both components and disentangle their contribution. In this paper we have presented the application of our spectral decomposition method to the largest sample of local ULIRGs studied so far, consisting of 164 sources up to $z \\simeq 0.35$. The size of our sample, which is not affected by any significant bias with respect to the nature of the energy source, allowed us to obtain the following results: \\\\ 1) Our method has proven to be very effective in the discovery of faint/obscured AGN components, with an AGN detection rate among ULIRGs of $\\sim$70 per cent (113/164 secure detections, plus a dozen of ambiguous cases among the fainter objects). This rate is comparable to that achieved by collecting together all the other multiwavelength diagnostics. \\\\ 2) In terms of global contribution, star formation is confirmed as the dominant power source for extreme IR activity: only $\\sim$27 per cent of the total luminosity of our sample turns out to have a gravitational origin. Nevertheless, even if they are usually minor contributors with respect to the concurrent SB events, AGN play a key role in the ULIRG phenomenon, since obscured nuclear activity is actually in place in many sources that are optically classified as H~\\textsc{ii} regions (or LINERs) as well. Elusive AGN components are important in $\\sim$10 per cent of ULIRGs. \\\\ 3) The coverage of the entire ULIRG luminosity range allows a comprehensive and \\textit{quantitative} re-analysis of the well-known correlation between the relative contribution of nuclear activity and the overall energy output. The increasing trend as a function of the total IR luminosity is clearly recovered: the average AGN contribution is almost negligible at the lower luminosities, but it gradually grows and eventually outshines the SB counterpart. A physical turning point is suggested to exist around $L_\\mathit{IR} \\simeq 3 \\times 10^{12} L_\\odot$, possibly related to the complex interplay between the two processes at work. \\\\ 4) As for the morphological properties, the AGN content turns out to be larger in correspondence of a more advanced merger stage; this supports the possibility of a time dependence in the relative AGN/SB contribution to the luminosity of interacting systems, consistent with the evolutive scenario according to which some ULIRGs may represent an intermediate phase in the formation of optically bright quasars. \\\\ 5) The analysis of a control sample of 52 IR-luminous sources at $z \\sim 1$ shows that to first approximation the spectral properties and the large-scale SED shapes of the AGN and SB components do not suffer substantial evolution with redshift. Our method then provides also an initial measure of the role of black hole accretion and star formation among ULIRG-like systems at earlier cosmic epochs. Moreover, the large spectral coverage that will be achieved with the upcoming IR facilities hints at the photometric extension of our diagnostic technique, based on accessible indicators such as the $\\sim$5~$\\mu$m rest-frame slope and the bolometric correction." }, "1003/1003.2894_arXiv.txt": { "abstract": "{ We combine published optical and near-infrared photometry to identify new low-mass candidate members in an area of about 0.64~deg$^2$ in Corona Australis with the $S$-parameter method. Five new candidate members of the region are selected. They have estimated ages between 3 and 15~Myr and masses between 0.05 and 0.15~$M_{\\odot}$. With {\\em Spitzer} photometry we confirm that these objects are not surrounded by optically thick disks. However, one of them is found to display excess at 24~$\\mu$m, thus suggesting it harbors a disk with an inner hole. With an estimated mass of 0.07$M_{\\odot}$ according to the SED fitting, this is one of the lowest-mass objects reported to possess a transitional disk. Including these new members, the fraction of disks is about 50\\% among the total Corona Australis population selected by the same criteria, lower than the 70\\% fraction reported earlier for this region. Even so, we find a ratio of transitional to primordial disks (45\\%) very similar to the value derived by other authors. This ratio is higher than for solar-type stars (5-10\\%), suggesting that disk evolution is faster in the latter, and/or that the ``transitional disk'' stage is not such a short-lived step for very low-mass objects. However, this impression needs to be confirmed with better statistics. }{}{}{}{} ", "introduction": "\\label{sec:intro} In the last two decades, our understanding of the formation and early evolution of low-mass stars and substellar objects has been remarkably improved, but there are still many unknowns regarding the exact mechanisms at play in these processes (see e.g. Whitworth et al. \\cite{whitworth07} for a review). The census of the low-mass population for an increasing number of star-forming regions is important to address the problem of the universality of the initial mass function (IMF). The comparison of the disk fractions (that is, the number of sources with disks with respect to the total number of members) in clusters of similar and different ages also helps to constrain the timescales of disk evolution (e.g. Haisch et al. \\cite{haisch01}). Moreover, the disk fraction and the spatial distribution of stars with and without disks provide important clues to understand the star-forming history and the disk evolution of a particular region. Most of our current knowledge about these issues comes from the study of star-forming regions and young clusters at distances within 100-300~pc from the Sun (see the reviews collected in Reipurth \\cite{reipurth08}). However, in the vast majority of these studies the identification of cluster members is based on the detection of features related to accretion and circum(sub)stellar disks (e.g. H$\\alpha$ emission, infrared excess), or indicative of strong activity (e.g. hard X-ray emission). Thus, these surveys are likely to be missing a fraction of the young low-mass population, especially objects that have dissipated all or most of their disks, that are in quiescence at the moment of the observations, and/ or whose activity level lies below the detection threshold of current X-ray surveys. In a recent paper Comer\\'on, Spezzi \\& L\\'opez Mart\\'{\\i} (2009, hereafter \\cite{comeron09}) reported on a large-scale optical survey of the \\object{Lupus} clouds, selected to include most of the areas surveyed in the {\\em Spitzer} Legacy program {\\em ``From molecular cores to planet-forming disks''} ({\\em c2d}). Combining their data with 2MASS near-infrared photometry, these authors developed an analysis procedure based on a dimensionless parameter $S$ that allowed them to easily identify very low-mass members of the dark clouds as an excess in the expected S-parameter distribution of field cool dwarfs and giants (see Sect.~\\ref{sec:analysis} for details). In this way they were able to unveil a large population of objects belonging to Lupus I and III, which seems to be composed of very low-mass stars and brown dwarfs that have lost their inner disks on a timescale of a few Myr. The discovery of a substantial and even dominant population of thus far unnoticed members of one of our nearest star-forming regions stresses the important unknowns that still subsist in the observational characterization of young very low-mass objects. It also poses the important question of whether similar populations of yet unknown very low-mass members exist in other regions. In this paper, we apply the $S$-parameter analysis to a set of optical and near-infrared observations of \\object{Corona Australis}, another of the nearest regions with ongoing or recent intermediate- and low-mass star formation. Our aim is the identification of new very low-mass candidate members of this dark cloud. We focus on the core in the direction of the intermediate-mass star \\object{R~CrA}, which contains a compact embedded stellar cluster known as the ``\\object{R~CrA cluster}'' or ``the {\\em\\object{Coronet}} cluster'' (Taylor \\& Storey \\cite{taylor84}). With an age of about 3~Myr, its distance has been estimated to be within 50-170~pc (see Neuh\\\"auser \\& Forbrich \\cite{neuhauser08} for a review). The young stellar population in this region has been studied with different techniques, including H$\\alpha$ surveys (e.g. Marraco \\& Ryndgren \\cite{marraco81}), optical spectroscopy (e.g. Walter et al. \\cite{walter97}), near-infrared mapping (e.g. Wilking et al. \\cite{wilking92}), mid-infrared surveys (e.g. Olofsson et al. \\cite{olofsson99}) and X-ray observations (e.g. Neuh\\\"auser et al. \\cite{neuhauser00}; Forbrich \\& Preibisch \\cite{forbrich07}). Some brown dwarfs and brown dwarf candidates have been reported (Wilking et al. \\cite{wilking97}; Fern\\'andez \\& Comer\\'on \\cite{fernandez01}; Bouy et al. \\cite{bouy04}; L\\'opez Mart\\'{\\i} et al. \\cite{lm05}). Very recently, this cloud has also been the target of {\\em Spitzer} observations (Sicilia-Aguilar et al. \\cite{sicilia08}; L\\'opez Mart\\'{\\i} et al. \\cite{lm09}). The structure of the paper is as follows: In Section~\\ref{sec:data} we present the data used in our analysis. Section~\\ref{sec:analysis} summarizes the basis of the $S$-parameter formalism, discusses the contamination expected in our sample and describes the object selection. Section~\\ref{sec:disc} is a discussion of the membership status and the properties of our new candidate members. Finally, in Section~\\ref{sec:concl} we draw our conclusions. ", "conclusions": "\\label{sec:concl} We performed an analysis of optical and near-infrared data of the Corona Australis star-forming region based on the $S$-parameter formalism. In our surveyed area of $\\sim0.64$~deg$^2$, we identified fourteen previous candidate members and five new candidate members of this dark cloud. The new candidates have estimated effective temperatures between 2900 and 3400~K, corresponding to masses between 0.05 and 0.13$M_{\\odot}$. They are thus very low-mass stars and massive brown dwarfs. Their ages span between 3 and 15~Myr, which is consistent with the reported age spread in Corona Australis ($\\sim$1-10~Myr), given the uncertainties in the models, in the distance to the cloud and in the SED fitting procedure. The membership of these objects to the star forming region is further supported by their spatial distribution and, when available, proper motion information. The source for which membership is more uncertain is CrA~J190151.7-371048.4, the oldest star in our sample according to the SED fitting results. The SEDs of four of our new candidates are nearly photospheric. The exception is CrA~J190111.6-364532.0: This source displays excess at 24~$\\mu$m, which is suggestive of a transition disk with an inner hole. With an estimated mass of 0.07$M_{\\odot}$, this is one of the lowest-mass objects reported to possess such a disk. We calculated the disk fraction of the Corona Australis population selected with our method, which is 50\\%. This value is lower than the one reported in a previous study by Sicilia-Aguilar et al. (\\cite{sicilia08}). The ratio of transitional to primordial disks (45\\%) though agrees well with the fraction of 50\\% reported by Sicilia-Aguilar et al. (\\cite{sicilia08}) and is remarkably higher than the value measured in other clusters of similar age. This suggests that transitional disks around brown dwarfs may have longer lifetimes than around low-mass stars. However, this impression should be confirmed with a larger, more statistically meaningful sample. The results from this work stress the need to properly characterize the diskless population of a region to derive meaningful disk fractions. This is especially important to understand the dependence of the disk fractions with the mass of the central objects, in particular for the very low-mass population. Even in a relatively small survey like the one reported here and in a well-studied and relatively low-density region like Corona Australis, 25\\% of the total number of cloud members selected with our method had not been detected in previous studies based on accretion, disk or activity signatures. An analogous study of an eventual larger-scale survey covering most of the dark cloud would probably identify a significant number of young very low-mass objects belonging to this star-forming region." }, "1003/1003.0405_arXiv.txt": { "abstract": "{Recent observational work by Israelian et al. has shown that sun-like planet host stars in the temperature range $5700$\\,K$ < T_{eff} <5850$\\,K have lithium abundances that are significantly lower than those observed for ``single'' field stars. In this letter we use stellar evolutionary models to show that differences in stellar mass and age are not responsible for the observed correlation. This result, along with the finding of Israelian et al., strongly suggest that the observed lithium difference is likely linked to some process related to the formation and evolution of planetary systems. ", "introduction": "The increasing number of discovered planets continues to give new insights into the planet formation processes \\citep[e.g.][]{Udry-2007b}. The study of planet host stars is providing crucial observational evidence for this. The first large uniform studies \\citep[][]{Santos-2001,Santos-2004b} have shown that stars with giant planets have chemical abundances that are distinctly different from those found in ``single'' stars, a result that was later confirmed by other authors \\citep[e.g.][]{Fischer_Valenti-2005}. This result provided important constraints for the models of planet formation and evolution \\citep[e.g. ][]{Pollack-1996, Mordasini-2009, Boss-2002}. Although most studies of chemical abundances in stars with planets have concentrated on the measurement of iron abundances as a metallicity proxy, several works have been made to study the abundances for a variety of refractory and volatile elements in planet host stars \\citep[e.g.][]{Gilli-2006, Ecuvillon-2006b, Takeda-2007, Neves-2009}. Besides the general chemical enrichment found for stars hosting giant planets \\citep[interestingly not found for stars hosting very low mass planets, see: ][]{Sousa-2008}, none of these found compelling evidence for other chemical peculiarities. The possibility that stars with planets present different abundances of the light elements lithium and beryllium has also been debated \\citep[][]{King-1997, GarciaLopez-1998, Deliyannis-2000, Cochran-1997, Ryan-2000, Gonzalez-2000, Gonzalez-2008, Israelian-2004, Santos-2004c, Takeda-2005, Takeda-2007, Chen-2006, Luck-2006}. Different abundances could indicate that planets or planetary material were engulfed by the star during its lifetime (and in which quantity) \\citep[e.g.][]{Israelian-2001}, or suggest that the rotational history of the stars depends on the existence of planets or indirectly on the process for planet formation \\citep[][]{Bouvier-2008, Castro-2009}. Recently, \\citet[][]{Israelian-2009} presented a large uniform study of lithium abundances in a sample of stars from the HARPS planet search programme \\citep[][]{Mayor-2003}. In their work they reported a significant difference regarding the depletion of lithium for planet host stars when compared with stars with no detected planet in the range $5700$\\,K$ < T_{eff} <5850$\\,K, confirming former suspicions \\citep[e.g.][]{Israelian-2004,Takeda-2007}. According to Israelian et al., stars with planets in the temperature range around the solar temperature (solar analogues) have significantly lower lithium abundances when compared with ``single'' stars (for which no planets were detected so far). The uniformity of the HARPS sample (composed mainly of old inactive stars) allowed Israelian et al. to exclude effects like stellar rotation, stellar activity, or chemical abundances as the cause for the observed difference. In this paper we use stellar evolution models to explore the possibility that stellar age and mass could be responsible for the observed difference. In Sect.\\,2 we present the procedure used for the determination of precise and uniform masses and ages for our sample. In Sect.\\,3 we explore possible correlations between these two parameters and the lithium abundances. We conclude in Sect.\\,4. ", "conclusions": "We derived uniform values for the stellar ages and masses in a sample of solar analogue stars, for which \\citet[][]{Israelian-2009} have found a clear difference in the lithium abundances correlated to the existence of planets. Our analysis has shown that age and mass cannot explain the lithium abundance differences. This strongly suggests that the observed differences are not related to stellar intrinsic properties. This result confirms the uniformity of the studied sample. Our results are directly linked with Li observations in open clusters. A large high-quality database for Li in open clusters acquired in the last years has permitted us to draw a secure picture of the empirical behaviour and evolution of Li in solar analogue stars \\citep[][]{Randich-2008, Randich-2009, Sestito-2005}. Numerous observations show that stars with the same age, temperature, and metallicity can be affected by different amounts of Li depletion. A dispersion in Li of at least a factor of 10 has been observed in the solar-age, solar metallicity clusters M 67 \\citep[][]{Pasquini-2008} and several other old, solar metallicity or metal rich clusters \\citep[][]{Randich-2008, Randich-2009}. This spread strongly suggests that Li depletion must be affected by an additional parameter besides mass, age, and chemical composition. This parameter could be an initial angular momentum of the star affected by the formation and evolution of planets. Similar Li dispersion (or often cited as bi-modality) is also observed in solar type field stars \\citep[][]{Favata-1996, Galeev-2004, Chen-2001, Lambert-2004}. We think that the same mechanism is responsible for the large Li dispersion in the field and open cluster solar type stars. Figure \\ref{fig_calibration_proc} clearly supports this suggestion. No clear explanation has been found for the lithium abundance difference observed between stars with and without planets, though a few possibilities have been suggested. These include the star-planet interaction \\citep[][]{Castro-2009}, the infall of planets into the star (leading to higher mixing - Theado et al., priv. comm.), or a difference in the rotational history of the star due to star-disc interaction. Massive proto-planetary discs capable of forming planets will likely help to break stellar rotation, thus changing the depletion rate of lithium \\citep[e.g.][]{Charbonnel-2005, Cochran-1997, Bouvier-2008, Pinsonneault-2010}. Interestingly, the same mechanism could have been responsible for the low lithium abundance observed in our own Sun, itself a planet host stars, and can be used to explain the dispersion of lithium abundances observed in clusters of different ages \\citep[][]{Sestito-2005}." }, "1003/1003.0633_arXiv.txt": { "abstract": "The long-term habitability of Earth-like planets requires low orbital eccentricities. A secular perturbation from a distant stellar companion is a very important mechanism in exciting planetary eccentricities, as many of the extrasolar planetary systems are associated with stellar companions. Although the orbital evolution of an Earth-like planet in a stellar binary is well understood, the effect of a binary perturbation to a more realistic system containing additional gas giant planets has been very little studied. Here we provide analytic criteria confirmed by a large ensemble of numerical integrations that identify the initial orbital parameters leading to eccentric orbits. We show that an extra-solar earth is likely to experience a broad range of orbital evolution dictated by the location of a gas-giant planet, necessitating more focused studies on the effect of eccentricity on the potential for life. ", "introduction": "\\subsection{Habitability} Habitability of a planet is conventionally defined as the capability for liquid water to be sustained on a planetary surface. A high (or low) level of stellar flux incident upon a planetary surface results in the loss of liquid water through evaporation by the runaway-greenhouse effect (or through freezing by global refrigeration) (Kasting \\textit{et al.}, 1993). Planets with eccentric orbits experience varying levels of stellar flux throughout a year as the distance from the star fluctuates during an orbit. The time-averaged stellar flux $$ over the orbital period is \\begin{equation} ={L\\over 4\\pi a^2 (1-e^2)^{1/2} }, \\end{equation} where $L$ is the luminosity of the host star, and $a$ and $e$ are the semi-major axis and eccentricity of the planet, respectively (Williams and Pollard, 2002). As seen in Eq. 1, the eccentricity growth of a planet results in an increase in the time-averaged stellar flux and thus eventually leads to unhabitable conditions. Climate models show that the Earth would start to lose its surface water with a hypothetical eccentricity of 0.4 around the Sun; if the eccentricity exceeds 0.7, the Earth would lose all of its liquid water through the runaway-greenhouse effect (Williams and Pollard, 2002). Thus, a perturbation that causes a planet to deviate from a circular orbit can greatly impact its long-term climatic stability. This in turn could disturb the possible origin, evolution, and prevalence of life on the planet. \\subsection{Kozai Mechanism} A secular perturbation from a stellar companion is one of the most efficient mechanisms in exciting planetary eccentricities to very large values. Unlike our own Solar System, at least 20\\,\\% of the $\\sim240$ extrasolar planetary systems detected as of 2007 are members of multiple-star systems (Raghavan \\textit{et al.}, 2006; Desidera and Barbieri, 2007; Eggenberger \\textit{et al.}, 2007). Moreover, the multiplicity among the current sample of planet-hosting stars should be higher than 20\\,\\% as the photometric searches for stellar companions around known planetary systems are still ongoing. These stellar companions, despite their large distances from the planetary systems (typical binary separations range from $\\sim 10^2 $ -- $10^4\\,$AU), can still secularly perturb the planetary orbits around the primaries through a unique three-body interaction called the ``Kozai mechanism'', in which the cyclic angular momentum exchange between a binary companion and a planet results in a large-amplitude eccentricity oscillation of the planet (Kozai, 1962; Holman \\textit{et al.}, 1997; Innanen \\textit{et al.}, 1997). The Kozai mechanism takes place when the initial relative inclination $\\ipcomp$ between the planetary and binary orbits exceeds the critical Kozai angle $\\ikoz = 39.23^{\\circ}$. Under such an initial configuration, the planet's eccentricity grows from $\\sim0$ and oscillates with an amplitude that is constant to the lowest order as \\begin{equation} e_{\\rm max} \\simeq \\sqrt{1-{5 \\over 3} \\cos^2{\\ipcomp}}. \\end{equation} (Holman \\textit{et al.}, 1997). For example, an Earth-mass planet at 1\\,AU would reach the habitability limits of $e =$0.4 and 0.7 if there is a stellar companion initially inclined by $\\ipcomp = 45^\\circ$ and $56^\\circ$, respectively. Note that the distribution of $ \\ipcomp$ in space is most likely isotropic because the orbital orientation of binaries with separations greater than $\\sim 100\\,$AU are not expected to be correlated with the invariable plane of the planetary systems around the primaries (Hale, 1994; Takeda \\textit{et al.}, 2008). During Kozai cycles, a stellar companion repeatedly applies a small torque on the planetary orbit. This torque accumulated over many binary orbits results in the precession of the planet's pericenter argument and hence the oscillation of the orbital eccentricity. The timescale of the pericenter and eccentricity evolution can be analytically estimated as \\begin{equation} \\tkoz \\approx {2\\over3\\pi}{P_{\\rm bin}^{2}\\over P_1}{m_0+m_1+m_{\\rm bin} \\over m_{\\rm bin}}(1-e_{\\rm bin}^2)^{3/2} \\label{tkoz} \\end{equation} (Kiseleva \\textit{et al.}, 1998; Ford \\textit{et al.}, 2000), where the subscripts $bin$, $0$, and $1$ refer to the binary companion, the primary star, and the planet, respectively. If there is another source of perturbation that precesses the planetary orbit on the timescale shorter than $\\tkoz$ given in Eq.~\\ref{tkoz}, then the torque applied from the stellar companion to the planetary orbit secularly averages out to zero, and no eccentricity growth by the companion would be observed (Wu and Murray, 2003). The Kozai timescale may largely vary depending on the system parameters; for example, an Earth-size planet with a solar-mass stellar companion at 750\\,AU would experience Kozai oscillations with a period $\\tkoz \\approx 320\\,$Myr. If a stellar companion is closer, then the oscillation timescale would be largely reduced; for example at 250\\,AU, $ \\tkoz \\approx 15\\,$Myr. Note that this oscillation period is also relevant to the evolution of life on a planet. For example, if the Earth had an eccentricity of 0.4, even though some liquid water would remain on the surface, the surface temperatures would exceed $70^{\\circ}\\,C$ near the pericenter of the orbit (Williams and Pollard, 2002). At this temperature, thermophiles and other simple organisms may be able to comfortably survive, but complex life could be eliminated (Levy and Miller, 1998; Daniel and Cowan, 2000). However, if the eccentricity oscillation period is sufficiently long (on the order of hundred million years), the low eccentricity duration of the cycle may provide an intriguing avenue for the evolution of more complex organisms. Although the eccentricity variation of a planet in a three-body system is understood well, in reality, a single Earth-like planet in a binary is probably not a common configuration. Both observations and numerical planet-formation simulations rather suggest that planetary systems naturally form in multiple configurations, containing one or a few gas giant planets (Ida and Lin, 2005; Thommes \\textit{et al.}, 2008). Thus, it is certainly relevant to discuss the dynamical evolution of a habitable Earth-like planet, secularly perturbed by other planets and a distant stellar companion. However, the evolution of gravitationally-coupled multiple planets under the influence of a secular binary perturbation is significantly more complex than the simple three-body Kozai mechanism and has been very little understood in previous studies. ", "conclusions": "\\subsection{Using the Timescale Analysis} The timescale analysis provides a method to quantify the strength of the dynamical interactions, thus creating the ability to predict the behaviors that should arise from a wide range of planetary systems hosted within a wide-binary. The consistency between the behaviors seen in the numerical simulations with the behaviors predicted by the timescale analysis demonstrates the effectiveness of this approach, as shown in Figure 8. Using the timescale analysis, we can show, without exploring the parameter space with numerical simulations, that a different range of parameters would result in a similar set of orbital behaviors. Figure 10 shows the parameters at which an earth-like planet will experience the high eccentricities observed in Region D. The parameters were determined by obtaining the specific conditions determined by the timescale analysis. \\subsection{Distribution of Eccentricities} Most notably, our simulations indicate that an earth is likely to exhibit high eccentricities across a wide range of parameters. Figure 11 illustrates the probability of the earth reaching the habitable eccentricity limits of 0.4 and 0.7 as a function of the jupiter's semimajor axis. The similarity between the dotted and solid lines at the dynamically rigid and nodal libration regions ($a_{\\rm 2} < 20$) indicate that the Earth is likely to experience extremely high eccentricity oscillations rather than a distribution of high and moderate eccentricities like the single-planet case. As shown with the timescale analysis and additional simulation sets, a similar distribution should be seen across a range of parameters. Thus, an earth-like planet within a binary system will most likely exhibit convincingly non-habitable eccentricities. However, it is important to note the length of time at which the planet experiences the high orbital eccentricities. As described above, the earth-like planet may experience long periods of low eccentricities between its periods of high eccentricities. Such secular behavior must be carefully studied in the context of the evolution of life before ruling out the possibility of a habitable earth. \\subsection{Conclusion} Our results demonstrate the variety of orbits that an earth can exhibit within a binary system when accompanied by a second planet. With such a high prevalence of binaries, in addition to the strong likelihood for multiple planets within a system, we should expect eccentric earths even when the companion is distant. As the ability to search for extrasolar earths increases with observational advances, the possibility of other bodies within the system must also be taken into account when discussing habitability. The diverse orbital behavior that result from the presence of such bodies highlights the necessity to study in detail extrasolar earths without stable circular orbits. Thus, we call on climatologists, biologists, and astrobiologists to carefully consider the broad range of possible eccentricities when studying the potential for the origin, evolution, and prevalence of life on a planet." }, "1003/1003.2636_arXiv.txt": { "abstract": " ", "introduction": "\\label{sec:0} White dwarf binaries are thought to be the most common binaries in the Universe, and in our Galaxy their number is estimated to be as high as 10$^8$. In addition most stars are known to be part of binary systems, roughly half of which have orbital periods short enough that the evolution of the two stars is strongly influenced by the presence of a companion. Furthermore, it has become clear from observed close binaries, that a large fraction of binaries that interacted in the past must have lost considerable amounts of angular momentum, thus forming compact binaries, with compact stellar components. The details of the evolution leading to the loss of angular momentum are uncertain, but generally this is interpreted in the framework of the so called ``common-envelope evolution'': the picture that in a mass-transfer phase between a giant and a more compact companion the companion quickly ends up inside the giant's envelope, after which frictional processes slow down the companion and the core of the giant, causing the ``common envelope'' to be expelled, as well as the orbital separation to shrink dramatically \\cite{Taam and Sandquist (2000)}. Among the most compact binaries know, often called ultra-compact or ultra-short binaries, are those hosting two white dwarfs and classified into two types: \\emph{detached} binaries, in which the two components are relatively widely separated and \\emph{interacting} binaries, in which mass is transferred from one component to the other. In the latter class a white dwarf is accreting from a white dwarf like object (we often refer to them as AM CVn systems, after the prototype of the class, the variable star AM CVn; \\cite{warn95,Nelemans (2005)}). \\begin{figure} \\includegraphics[height=7.5cm,angle=0]{P_Mtot_SPYnew1} \\caption{Period versus total mass of double white dwarfs. The points and arrows are observed systems \\cite{Nelemans et al. (2005)}, the grey shade a model for the Galactic population. Systems to the left of the dashed line will merge within a Hubble time, systems above the dotted line have a combined mass above the Chandrasekhar mass. The top left corner shows the region of possible type Ia supernova progenitors, where the grey shade has been darkened for better visibility (adapted from \\cite{Nelemans (2007)}). } \\label{fig:P_Mtot} \\end{figure} In the past many authors have emphasised the importance of studying white dwarfs in DDBs. In fact, the study of ultra-short white dwarf binaries is relevant to some important astrophysical questions which have been outlined by several author. Recently, \\cite{Nelemans (2007)} listed the following ones: \\begin{itemize} \\item {\\em Binary evolution} Double white dwarfs are excellent tests of binary evolution. In particular the orbital shrinkage during the common-envelope phase can be tested using double white dwarfs. The reason is that for giants there is a direct relation between the mass of the core (which becomes a white dwarf and so its mass is still measurable today) and the radius of the giant. The latter carries information about the (minimal) separation between the two components in the binary before the common envelope, while the separation after the common envelope can be estimated from the current orbital period. This enables a detailed reconstruction of the evolution leading from a binary consisting of two main sequence stars to a close double white dwarf \\cite{Nelemans et al.(2000)}. The interesting conclusion of this exercise is that the standard schematic description of the common envelope -- in which the envelope is expelled at the expense of the orbital energy -- cannot be correct. An alternative scheme, based on the angular momentum, for the moment seems to be able to explain all the observations \\cite{Nelemans and Tout (2005)}. \\item {\\em Type Ia supernovae} Type Ia supernovae have peak brightnesses that are well correlated with the shape of their light curve \\cite{Phillips (1993)}, making them ideal standard candles to determine distances. The measurement of the apparent brightness of far away supernovae as a function of redshift has led to the conclusion that the expansion of the universe is accelerating \\cite{Perlmutter et al. (1998),Riess et al.(2004)}. This depends on the assumption that these far-away (and thus old) supernovae behave the same as their local cousins, which is a quite reasonable assumption. However, one of the problems is that we do not know what exactly explodes and why, so the likelihood of this assumption is difficult to assess \\cite{Podsiadlowski et al. (2006)}. One of the proposed models for the progenitors of type Ia supernovae are massive close double white dwarfs that will explode when the two stars merge \\cite{Iben and Tutukov (1984)}. In Fig.~\\ref{fig:P_Mtot} the observed double white dwarfs are compared to a model for the Galactic population of double white dwarfs \\cite{Nelemans et al. (2001)}, in which the merger rate of massive double white dwarfs is similar to the type Ia supernova rate. The grey shade in the relevant corner of the diagram is enhanced for visibility. The discovery of at least one system in this box confirms the viability of this model (in terms of event rates). \\item {\\em Accretion physics} The fact that in AM CVn systems the mass losing star is an evolved, hydrogen deficient star, gives rise to a unique astrophysical laboratory, in which accretion discs made of almost pure helium \\cite{Marsh et al. (1991),Schulz et al.(2001),Groot et al. (2001),Roelofs et al. (2006),Werner et al. (2006)}. This opens the possibility to test the behaviour of accretion discs of different chemical composition. \\item {\\em Gravitational wave emission} Untill recently the DDBs with two NSs were considered among the best sources to look for gravitational wave emission, mainly due to the relatively high chirp mass expected for these sources, In fact, simply inferring the strength of the gravitational wave amplitude expected for from \\cite{Evans et al. (1987)} \\begin{equation} h = \\left[ \\frac{16 \\pi G L_{GW}} {c^3 \\omega^2_g 4 \\pi d^2} \\right] ^{1/2} = 10^{-21} \\left( \\frac{{\\cal{M}}}{\\msun} \\right)^{5/3} \\left ( \\frac{P_{orb}}{\\rm 1 hr} \\right)^{-2/3} \\left ( \\frac{d}{\\rm 1 kpc} \\right)^{-1} \\end{equation} where \\begin{equation} L_{GW} = \\frac{32}{5}\\frac{G^4}{c^5}\\frac{M^2 m^2 (m+M)}{a^5} ; \\end{equation} \\begin{equation} {\\cal{M}}=\\frac{(Mm)^{3/5}}{(M+m)^{1/5}} \\end{equation} where the frequency of the wave is given by $f = 2/P_{orb}$. It is evident that the strain signal $h$ from DDBs hosting neutron stars is a factor 5-20 higher than in the case of DDBs with white dwarfs as far as the orbital period is larger than approximatively 10-20 minutes. In recent years, AM CVns have received great attention as they represent a large population of guaranteed sources for the forthcoming \\textit{Laser Interferometer Space Antenna} \\cite{2006astro.ph..5722N,2005ApJ...633L..33S}. Double WD binaries enter the \\textit{LISA} observational window (0.1 $\\div$ 100 mHz) at an orbital period $\\sim$ 5 hrs and, as they evolve secularly through GW emission, they cross the whole \\textit{LISA} band. They are expected to be so numerous ($\\sim 10^3 \\div 10^4$ expected), close on average, and luminous in GWs as to create a stochastic foreground that dominates the \\textit{LISA} observational window up to $\\approx$ 3 mHz \\cite{2005ApJ...633L..33S}. Detailed knowledge of the characteristics of their background signal would thus be needed to model it and study weaker background GW signals of cosmological origin. \\end{itemize} \\begin{figure} \\includegraphics[height=12cm,angle=-90]{Galactic_GWR1} \\caption{Expected signals of ultra-compact binaries, the ones with error bars from (adapted from \\cite{Roelofs et al. (2006),Nelemans (2007)}.} \\label{fig:fh_HST} \\end{figure} A relatively small number of ultracompact DDBs systems is presently known. According to \\cite{2006MNRAS.367L..62R} there exist 17 confirmed objects with orbital periods in the $10 \\div 70$ min in which a hydrogen-deficient mass donor, either a semi-degenerate star or a WD itself, is present. These are called AM CVn systems and are roughly characterized by optical emission modulated at the orbital period, X-ray emission showing no evidence for a significant modulation (from which a moderately magnetic primary is suggested, \\cite{2006astro.ph.10357R}) and, in the few cases where timing analyses could be carried out, orbital spin-down consistent with GW emission-driven mass transfer. In addition there exist two peculiar objects, sharing a number of observational properties that partially match those of the ``standard'' AM CVn's. They are \\src\\ and \\srcm, whose X-ray emission is $\\sim$ 100\\% pulsed, with on-phase and off-phase of approximately equal duration. The single modulations found in their lightcurves, both in the optical and in X-rays, correspond to periods of, respectively, 321.5 and 569 s (\\cite{2004MSAIS...5..148I,2002ApJ...581..577S}) and were first interpreted as orbital periods. If so, these two objects are the binary systems with the shortest orbital period known and could belong to the AM CVn class. However, in addition to peculiar emission properties with respect to other AM CVn's, timing analyses carried out by the above cited authors demonstrate that, in this interpretation, these two objects have shrinking orbits. This is contrary to what expected in mass transferring double white dwarf systems (including AM CVn's systems) and suggests the possibility that the binary is detached, with the orbit shrinking because of GW emission. The electromagnetic emission would have in turn to be caused by some other kind of interaction. Nonetheless, there are a number of alternative models to account for the observed properties, all of them based upon binary systems. The intermediate polar (IP) model (\\cite{Motch et al.(1996),io99,Norton et al. (2004)}) is the only one in which the pulsation periods are not assumed to be orbital. In this model, the pulsations are likely due to the spin of a white dwarf accreting from non-degenerate secondary star. Moreover, due to geometrical constraints the orbital period is not expected to be detectable. The other two models assume a double white dwarf binaries in which the pulsation period is the orbital period. Each of them invoke a semi-detached, accreting double white dwarfs: one is magnetic, the double degenerate polar model (\\cite{crop98,ram02a,ram02b,io02a}), while the other is non-magnetic, the direct impact model (\\cite{Nelemans et al. (2001),Marsh and Steeghs(2002),ram02a}), in which, due to the compact dimensions of these systems, the mass transfer streams is forced to hit directly onto the accreting white dwarfs rather than to form an accretion disk . \\begin{table}[!ht] \\caption{Overview of observational properties of AM CVn stars (adapted from \\cite{Nelemans (2005)})} \\label{tab:overview} \\smallskip \\begin{center} \\hspace*{-0.5cm} {\\small \\begin{tabular}{lllllllcc}\\hline Name & $P_{\\rm orb}^a$ & & $P_{\\rm sh}^a$ & Spectrum & Phot. var$^b$ & dist & X-ray$^c$ & UV$^d$ \\\\ & (s) & & (s) & & & (pc) & & \\\\ \\hline \\hline ES Cet & 621 &(p/s) & & Em & orb & 350& C$^3$X & GI \\\\ AM CVn & 1029 &(s/p) & 1051 & Abs & orb & 606$^{+135}_{-95}$ & RX & HI \\\\ HP Lib & 1103 &(p) & 1119 & Abs & orb & 197$^{+13}_{-12}$ & X & HI \\\\ CR Boo & 1471 &(p) & 1487 & Abs/Em? & OB/orb & 337$^{+43}_{-35}$& ARX & I \\\\ KL Dra & 1500 &(p) & 1530 & Abs/Em? & OB/orb & & & \\\\ V803 Cen & 1612 &(p) & 1618 & Abs/Em? & OB/orb & & Rx & FHI \\\\ SDSSJ0926+36 & 1698.6& (p) & & & orb & & & \\\\ CP Eri & 1701 &(p) & 1716 & Abs/Em & OB/orb & & & H \\\\ 2003aw & ? & & 2042 & Em/Abs? & OB/orb & & & \\\\ SDSSJ1240-01 & 2242 &(s) & & Em & n & & & \\\\ GP Com & 2794 &(s) & & Em & n & 75$\\pm2$ & ARX & HI \\\\ CE315 & 3906 &(s) & & Em & n & 77? & R(?)X & H \\\\ & & & & & & & & \\\\ Candidates & & & & & & & & \\\\\\hline\\hline RXJ0806+15 & 321 &(X/p) & & He/H?$^{11}$ & ``orb'' & & CRX & \\\\ V407 Vul & 569 &(X/p) & & K-star$^{16}$ & ``orb'' & & ARCRxX & \\\\ \\hline \\end{tabular} } \\end{center} {\\small $a$ orb = orbital, sh = superhump, periods from {ww03}, see references therein, (p)/(s)/(X) for photometric, spectroscopic, X-ray period.\\\\ $b$ orb = orbital, OB = outburst\\\\ $c$ A = ASCA, C = Chandra, R = ROSAT, Rx = RXTE, X = XMM-Newton {kns+04}\\\\ $d$ F = FUSE, G = GALEX, H = HST, I = IUE } \\end{table} After a brief presentation of the two X--ray selected double degenerate binary systems, we discuss the main scenario of this type, the Unipolar Inductor Model (UIM) introduced by \\cite{2002MNRAS.331..221W} and further developed by \\cite{2006A&A...447..785D,2006astro.ph..3795D}, and compare its predictions with the salient observed properties of these two sources. \\subsection{\\src} \\label{0:j0806} \\src\\ was discovered in 1990 with the \\R\\ satellite during the All-Sky Survey (RASS; \\cite{beu99}). However, it was only in 1999 that a periodic signal at 321\\,s was detected in its soft X-ray flux with the \\R\\ HRI (\\cite{io99,bur01}). Subsequent deeper optical studies allowed to unambiguously identify the optical counterpart of \\src, a blue $V=21.1$ ($B=20.7$) star (\\cite{io02a,io02b}). $B$, $V$ and $R$ time-resolved photometry revealed the presence of a $\\sim 15$\\% modulation at the $\\sim 321$\\,s X-ray period (\\cite{io02b,ram02a}. \\begin{figure*}[htb] \\resizebox{16pc}{!}{\\rotatebox[]{-90}{\\includegraphics{new_spec_norm.ps}}} \\caption{VLT FORS1 medium (6\\AA; 3900--6000\\AA) and low (30\\AA; above 6000\\AA) resolution spectra obtained for the optical counterpart of \\src. Numerous faint emission lines of HeI and HeII (blended with H) are labeled (adapted form \\cite{io02b}).} \\label{spec} \\end{figure*} \\begin{figure*}[hbt] \\centering \\resizebox{20pc}{!}{\\includegraphics{israel_f1.eps} } \\caption{Left panel: Results of the phase fitting technique used to infer the P-\\.P coherent solution for \\src: the linear term (P component) has been corrected, while the quadratic term (the \\.P component) has been kept for clarity. The best \\.P solution inferred for the optical band is marked by the solid fit line. Right panel: 2001-2004 optical flux measurements at fdifferent wavelengths.} \\label{timing} \\end{figure*} The VLT spectral study revealed a blue continuum with no intrinsic absorption lines \\cite{io02b} . Broad ($\\rm FWHM\\sim 1500~\\rm km~s^{-1}$), low equivalent width ($EW\\sim -2\\div-6$ \\AA) emission lines from the He~II Pickering series (plus additional emission lines likely associated with He~I, C~III, N~III, etc.; for a different interpretation see \\cite{rei04}) were instead detected \\cite{io02b}. These findings, together with the period stability and absence of any additional modulation in the 1\\,min--5\\,hr period range, were interpreted in terms of a double degenerate He-rich binary (a subset of the AM CVn class; see \\cite{warn95}) with an orbital period of 321\\,s, the shortest ever recorded. Moreover, \\src\\ was noticed to have optical/X-ray properties similar to those of \\srcm, a 569\\,s modulated soft X-ray source proposed as a double degenerate system (\\cite{crop98,ram00,ram02b}). In the past years the detection of spin--up was reported, at a rate of $\\sim$6.2$\\times$10$^{-11}$\\, s~s$^{-1}$, for the 321\\,s orbital modulation, based on optical data taken from the Nordic Optical Telescope (NOT) and the VLT archive, and by using incoherent timing techniques \\cite{hak03,hak04}. Similar results were reported also for the X-ray data (ROSAT and Chandra; \\cite{stro03}) of \\src\\ spanning over 10 years of uncoherent observations and based on the NOT results \\cite{hak03}. A Telescopio Nazionale Galileo (TNG) long-term project (started on 2000) devoted to the study of the long-term timing properties of \\src\\ found a slightly energy--dependent pulse shape with the pulsed fraction increasing toward longer wavelengths, from $\\sim$12\\% in the B-band to nearly 14\\% in the I-band (see lower right panel of Figure~\\ref{QU}; \\cite{2004MSAIS...5..148I}). An additional variability, at a level of 4\\% of the optical pulse shape as a function of time (see upper right panel of Figure~\\ref{QU} right) was detected. The first coherent timing solution was also inferred for this source, firmly assessing that the source was spinning-up: P=321.53033(2)\\,s, and \\.P=-3.67(1)$\\times$10$^{-11}$\\,s~s$^{-1}$ (90\\% uncertainties are reported; \\cite{2004MSAIS...5..148I}). Reference \\cite{2005ApJ...627..920S} obtained independently a phase-coherent timing solutions for the orbital period of this source over a similar baseline, that is fully consistent with that of \\cite{2004MSAIS...5..148I}. See \\cite{2007MNRAS.374.1334B} for a similar coherent timing solution also including the covariance terms of the fitted parameters. \\begin{figure}[hbt] \\centering \\resizebox{33pc}{!}{\\rotatebox[]{0}{\\includegraphics{israel_f2.eps}}} \\caption{Left Panel: The 1994--2002 phase coherently connected X--ray folded light curves (filled squares; 100\\% pulsed fraction) of \\src, together with the VLT-TNG 2001-2004 phase connected folded optical light curves (filled circles). Two orbital cycles are reported for clarity. A nearly anti-correlation was found. Right panels: Analysis of the phase variations induced by pulse shape changes in the optical band (upper panel), and the pulsed fraction as a function of optical wavelengths (lower panel). } \\label{QU}% \\end{figure} The relatively high accuracy obtained for the optical phase coherent P-\\.P solution (in the January 2001 - May 2004 interval) was used to extend its validity backward to the ROSAT observations without loosing the phase coherency, i.e. only one possible period cycle consistent with our P-\\.P solution. The best X--ray phase coherent solution is P=321.53038(2)\\,s, \\.P=-3.661(5)$\\times$10$^{-11}$\\,s~s$^{-1}$ (for more details see \\cite{2004MSAIS...5..148I}). Figure~\\ref{QU} (left panel) shows the optical (2001-2004) and X--ray (1994-2002) light curves folded by using the above reported P-\\.P coherent solution, confirming the amazing stability of the X--ray/optical anti-correlation first noted by (\\cite{2003ApJ...598..492I}; see inset of left panel of Figure\\,\\ref{QU}). \\begin{figure} \\centering \\resizebox{20pc}{!}{\\rotatebox[]{-90}{\\includegraphics{israel_f3.eps}}} \\caption{The results of the \\xmm\\ phase-resolved spectroscopy (PRS) analysis for the absorbed blackbody spectral parameters: absorption, blackbody temperature, blackbody radius (assuming a distance of 500\\,pc), and absorbed (triangles) and unabsorbed (asterisks) flux. Superposed is the folded X-ray light curve. } \\label{xmm}% \\end{figure} \\begin{figure} \\centering \\resizebox{16pc}{!}{\\rotatebox[]{-90}{\\includegraphics{israel_f4.eps}}} \\caption{ Broad-band energy spectrum of \\src\\ as inferred from the \\AXAF, \\xmm, VLT and TNG measurements and {\\it EUVE\\/} upper limits. The dotted line represents one of the possible fitting blackbody models for the IR/optical/UV bands.} \\label{xmm}% \\end{figure} On 2001, a Chandra observation of \\src\\ carried out in simultaneity with time resolved optical observation at the VLT, allowed for the first time to study the details of the X-ray emission and the phase-shift between X-rays and optical band. The X-ray spectrum is consistent with a occulting, as a function of modulation phase, black body with a temperature of $\\sim$60\\,eV \\cite{2003ApJ...598..492I}. A 0.5 phase-shift was reported for the X-rays and the optical band \\cite{2003ApJ...598..492I}. More recently, a 0.2 phase-shift was reported by analysing the whole historical X-ray and optical dataset: this latter result is considered the correct one \\cite{2007MNRAS.374.1334B}. On 2002 November 1$^{\\rm st}$ a second deep X-ray observation was obtained with the \\xmm\\ instrumentations for about 26000\\,s, providing an increased spectral accuracy (see eft panel of Figure~\\ref{xmm}). The \\xmm\\ data show a lower value of the absorption column, a relatively constant black body temperature, a smaller black body size, and, correspondingly, a slightly lower flux. All these differences may be ascribed to the pile--up effect in the Chandra data, even though we can not completely rule out the presence of real spectral variations as a function of time. In any case we note that this result is in agreement with the idea of a self-eclipsing (due only to a geometrical effect) small, hot and X--ray emitting region on the primary star. Timing analysis did not show any additional significant signal at periods longer or shorter than 321.5\\,s, (in the 5hr-200ms interval). By using the \\xmm\\ OM a first look at the source in the UV band (see right panel of Figure~\\ref{xmm}) was obtained confirming the presence of the blackbody component inferred from IR/optical bands. Reference \\cite{2003ApJ...598..492I} measured an on-phase X-ray luminosity (in the range 0.1-2.5 keV) $L_X = 8 \\times 10^{31} (d/200~\\mbox{pc})^2$ erg s$^{-1}$ for this source. These authors suggested that the bolometric luminosity might even be dominated by the (unseen) value of the UV flux, and reach values up to 5-6 times higher. The optical flux is only $\\sim$ 15\\% pulsed, indicating that most of it might not be associated to the same mechanism producing the pulsed X-ray emission (possibly the cooling luminosity of the WD plays a role). Given these uncertainties and, mainly, the uncertainty in the distance to the source, a luminosity $W\\simeq 10^{32} (d/200~\\mbox{pc})^2$ erg s$^{-1}$ will be assumed as a reference value.\\\\ \\subsection{\\srcm\\ } \\label{0:j1914} The luminosity and distance of this source have been subject to much debate over the last years. Reference \\cite{2002MNRAS.331..221W} refer to earlier ASCA measurements that, for a distance of 200-500 pc, corresponded to a luminosity in the range ($4\\times 10^{33} \\div 2.5 \\times 10^{34}$) erg s$^{-1}$. Reference \\cite{2005MNRAS.357...49R}, based on more recent XMM-Newton observations and a standard blackbody fit to the X-ray spectrum, derived an X-ray luminosity of $\\simeq 10^{35} d^2_{\\mbox{\\tiny{kpc}}}$ erg s$^{-1}$, where $d_{\\mbox{\\tiny{kpc}}}$ is the distance in kpc. The larger distance of $\\sim$ 1 kpc was based on a work by \\cite{2006ApJ...649..382S}. Still more recently, \\cite{2006MNRAS.367L..62R} find that an optically thin thermal emission spectrum, with an edge at 0.83 keV attributed to O VIII, gives a significantly better fit to the data than a blackbody model. The optically thin thermal plasma model implies a much lower bolometric luminosity of L$_{\\mbox{\\tiny{bol}}} \\simeq 10^{33}$ d$^2_{\\mbox{\\tiny{kpc}}}$ erg s$^{-1}$. \\\\ Reference \\cite{2006MNRAS.367L..62R} also note that the determination of a 1 kpc distance is not free of uncertainties and that a minimum distance of $\\sim 200$ pc might still be possible: the latter leads to a minimum luminosity of $\\sim 3 \\times 10^{31}$ erg s$^{-1}$. \\\\ Given these large discrepancies, interpretation of this source's properties remains ambiguous and dependent on assumptions. In the following, we refer to the more recent assessment by \\cite{2006MNRAS.367L..62R} of a luminosity $L = 10^{33}$ erg s$^{-1}$ for a 1 kpc distance.\\\\ Reference \\cite{2006MNRAS.367L..62R} also find possible evidence, at least in a few observations, of two secondary peaks in power spectra. These are very close to ($\\Delta \\nu \\simeq 5\\times 10^{-5}$ Hz) and symmetrically distributed around the strongest peak at $\\sim 1.76 \\times 10^{-3}$ Hz. References \\cite{2006MNRAS.367L..62R} and \\cite{2006ApJ...649L..99D} discuss the implications of this possible finding. ", "conclusions": "\\label{conclusions} The observational properties of the two DDBs with the shortest orbital period known to date have been discussed in relation with their physical nature. \\\\ The Unipolar Inductor Model and its coupling to GW emission have been introduced to explain a number of puzzling features that these two sources have in common and that are difficult to reconcile with most, if not all, models of mass transfer in such systems.\\\\ Emphasis was put on the relevant new physical features that characterize the model. In particular, the role of spin-orbit coupling through the Lorentz torque and the role of GW emission in keeping the electric interaction active at all times has been thoroughly discussed in all their implications. It has been shown that the model does work over arbitrarily long timescales.\\\\ Application of the model to both \\src\\ and \\srcm\\ accounts in a natural way for their main observational properties. Constraints on physical parameters are derived in order for the model to work, and can be verified by future observations.\\\\ It is concluded that the components in these two binaries may be much more similar than it may appear from their timing properties and luminosities. The significant observational differences could essentially be due to the two systems being caught in different evolutionary stages. \\srcm\\ would be in a luminous, transient phase that preceeds its settling into the dimmer steady-state, a regime already reached by the shorter period \\src\\ . Although the more luminous phase is transient, its lifetime can be as long as $ \\sim 10^5$ yrs, one or two orders of magnitude longer than previously estimated.\\\\ The GW luminosity of \\srcm\\ could be much larger than previously expected since its orbital evolution could be largely slowed down by an additional torque, apart from GW emission.\\\\ Finally, we stress that further developements and refinements of the model are required to address more specific observational issues and to assess the consequences that this new scenario might have on evolutionary scenarios and population synthesis models." }, "1003/1003.6129_arXiv.txt": { "abstract": "We model the evolution of the mean galaxy occupation of dark-matter halos over the range $0.1 0.1 M_0$ merger event occuring between redshifts of 0.5 and 1.0. Futhermore, we find that more luminous galaxies are found to occupy more massive halos irrespectively of the redshift. Finally, the average number of galaxies per halo shows little increase from redshift z$\\sim$ 1.0 to z$\\sim$ 0.5, with a sharp increase by a factor $\\sim$3 from z$\\sim$ 0.5 to z$\\sim$ 0.1, likely due to the dynamical friction of subhalos within their host halos. ", "introduction": "\\label{intro} The correlation function of galaxies is a simple yet powerful tool that allows one to constrain cosmological parameters and models of galaxy formation. Furthermore, with the help of high redshift surveys, the evolution in the clustering of galaxies allows for a better discrimination between theoretical models degenerate at the present epoch \\citep{pea97}. Until recently, the galaxy correlation function had been thought to follow a power law % \\citep{tot69,pee74,got79}. Subsequently, a departure from the power law on small scales (of the order 1 to a few Mpc/h) in the galaxy correlation function was noticed in pioneering surveys of the eighties, e.g. \\citep{guz91}, % and has now been fully confirmed by many large and deep galaxy surveys. These consist of various surveys at low and intermediate redshifts (z $\\le$ 1.5) such as the SDSS \\citep{con02,zeh04}, 2dFGRS \\citep{mag03}, % VVDS \\citep{lef05b, pol06}, COMBO-17 \\citep{phl06}, DEEP2 \\citep{coi06} and for lyman break galaxies (LBGs) at high redshifts in the SXDS and GOODS surveys \\citep{ouc05, lee06}. % Earliest measures showing a non-power law for the galaxy correlation function were difficult to interpret as they relied heavily upon the angular correlation function, $\\omega(\\theta)$, at low redshifts \\citep{con02}. This meant an integration over a wide range of galaxy luminosities and redshifts as well as complicated correlations between the statistical errors \\citep{zeh04}. This was followed by measurements of the projected correlation function, $w_p(r_p)$, and $\\omega(\\theta)$ in larger and deeper surveys. It has been seen that the deviation from a power law becomes more pronounced for bright galaxy samples with $L$ $>$ $L_\\star$ \\citep{coi06,pol06} and for LBGs at high redshifts \\citep{lee06}. Likewise, in SPH simulations % the luminous and more strongly clustered galaxies show a similar behaviour where the 'kink' is clearer than in the case of the full galaxy sample \\citep{wei04}. Interestingly enough, the correlation function for dark matter in N-body simulations is well known to be non-adherent to a power law (Jenkins et al. 1998, Kauffmann et al. 1999, Cooray \\& Sheth 2002 and references therein). The natural question that arises is how biased is the galaxy distribution with respect to the underlying matter distribution? The overall shape of the dark matter correlation function is mostly unaffected as one goes to higher redshifts as seen in SPH simulations \\citep{wei04} and N-body simulations \\citep{jen98}. This is different to what is seen for high-z galaxies \\citep{ouc05, lee06, zhe07}, where the so-called 'break' is more prominent implying that the biasing of the galaxy distribution on large scales, spatial exclusion of dark matter halos on small scales, alongwith a host of other complex physical processes, such as dynamical friction, feedback from supernovae, ram-pressure stripping etc. conspire in a non-trivial way to produce differences in the galaxy correlation function at different redshifts. The break in the power law can be physically interpreted in the language of the halo model (see Cooray \\& Sheth 2002 for a detailed review) as the transition between two scales - small scales lying within the halo to those larger than the halo. It is only natural to use a halo-based prescription where galaxies form by the cooling of gas within dark matter halos \\citep{whi78}, which are bound, virialized clumps of dark matter that are roughly 200 times the background density at that time \\citep{gun72}. The galaxies occupy dark matter halos following a HOD (halo occupation distribution) model. In turn the HOD fully describes the bias in the distribution of galaxies with respect to the underlying dark matter distribution \\citep{ber02}. The motivation for HOD based models arose when it was noticed that the clustering of galaxies could be reproduced by populating halos in semi-analytic models with galaxies following a particular probability distribution cite{kau99,ben00,ber02}. Furthermore, it has been seen that without the help of a proper halo based description the strong clustering of red galaxies (at z $\\simeq$ 3) can be explained by high and unrealistic (anywhere in the range of 70-200 galaxies per halo) occupation numbers to match the observed number density and strong clustering of a small number of high mass halos \\citep{zhe04}. % Recently, several groups have studied the galaxy correlation in light of the HOD models \\citep{mag03, van03, ham04, zeh05, ouc05, phl06, con06, zhe07}. Most of these works have mainly concentrated on obtaining best-fit HOD parameters and consequently the evolution in the HOD and information on the underlying dark matter distribution. In some cases, data from different surveys having different selections were used to study the evolution \\citep{con06, zhe07} The work in this paper complements and extends these analyses by studying the evolution in the HOD over a redshift range z $\\approx$ 0.1 - 1.3 for a variety of luminosity-limited samples but always for data from the same survey, the VIMOS VLT Deep Survey (VVDS). We also study two different HOD models in order to obtain a better understanding on the degeneracies between the various parameters. The paper is organized as follows. In section~\\ref{data} we briefly describe the data used from the VVDS survey. This section is followed by Section~\\ref{model} giving an outline of the theoretical framework. Section~\\ref{results} describes the fitting procedure alongwith the best-fit parameters obtained for the different models. The estimates for the average halo mass and number of galaxies per halo (galaxy weighted) for the different luminosity-threshold samples are also presented. Finally, Section~\\ref{disc} wraps up with a discussion and conclusion of the results. Throughout we will assume a flat $\\Lambda$CDM model for which ($\\Omega_0,h,\\sigma_8$) = (0.3,0.7,0.9) at $z=0$. Here $\\Omega_0$ is the density in units of critical density today, $h$ is the Hubble constant today in units of 100 km s$^{-1}$ Mpc$^{-1}$, and $\\sigma_8$ describes the rms fluctuations of the initial field evolved to the present time using linear theory and smoothed with a top-hat filter of radius 8 Mpc/$h$. All absolute magnitudes are in the AB system. ", "conclusions": "\\label{disc} The comparison of analytical models and data provides useful information of how the distribution of galaxies depends on the underlying dark matter. Subsequently, the best-fitting parameters obtained as a result of this comparison provide physical information regarding the dark matter halos and galaxies. The size of the VVDS dataset allows one to study, with a unique sample, the global change in the underlying halo properties of an average galaxy down to $z \\sim 1$. We attempt to follow the evolution in some properties of a magnitude selected sample, evolving the magnitude cut-off based on accurate measurements of galaxy evolution. We have presented results of the fitting of analytical halo models, incorporating simple HOD models with minimal number of free parameters, to data (in this case the projected 2-point correlation function) from the VVDS survey. This allowed us to study the evolution of the average number weighted halo mass and satellite fraction. On different scales there are contributions from central-satellite, satellite-satellite, and central-central pairs of galaxies to the correlation function thereby providing constraints on the evolution of the galaxy satellite fraction. The evolution was obtained from data observed in the same restframe band, and provides for simpler interpretations as compared to previous studies using data from different restframe bands. Various luminosity threshold samples at different redshifts were selected and the corresponding best-fit HOD parameters for two similar HOD models obtained. This is done in order to single out possible degeneracies and inconsistencies with the fitting procedure at high redshifts. On the whole, both models are in agreement with each other and show similar trends in evolution. The impact of our selection on the average halo mass is addressed using the Millennium simulation. We find that a growth in halo mass as seen in the data could rather be an underestimation of $\\sim 10\\%$ to what is seen in an 'ideal' sample containing all the descendants. Therefore a measure in the growth of mass of a halo can be mainly attributed to the hierarchical formation of structure and not due to the typology of the selection. We find that the number-weighted average halo mass grows by $\\sim 90 \\%$ from redshift 1.0 to 0.5. This is the first time a growth in the underlying halo mass has been measured {\\it at high redshifts} within a single data survey, and provides evidence for the rapid accretion phase of massive halos. The mass accretion history follows the form given in Wechsler et al. (2002) with $M(z) = M_0 e^{-\\beta z}$, where $\\beta \\sim 1.07 \\pm 0.57 (1.54 \\pm 0.13)$ when only the VVDS points were used and $\\beta \\sim 1.94 \\pm 0.10 (2.09 \\pm 0.04)$ after including the SDSS data as reference points at low redshift and depending on the model used to obtain the best fits. The addition of the low redshift SDSS points adds complications due to the addition of possible systematics by comparing data from two different rest-frame bands, even after conversion to a common fiducial band. We adopt the average value of $\\beta \\sim 1.3 \\pm 0.30$ from the VVDS points when discussing a growth in halo mass, and found to be slightly higher than the results from N-body simulations. If we express this result in terms of the expected halo mass at present times, $M_0 \\simeq$ $10^{13.5} h^{-1}M_\\odot$, such halos appear to accrete $m \\sim$ 0.25 $M_0$ between redshifts of 0.5 and 1.0. Stewart et al. (2007) have shown that $\\sim$ 25\\% (80\\%) of $M_0=10^{13} h^{-1}M_\\odot$ halos experienced an $m > 0.3 M_0$ ($m > 0.1 M_0$) merger event in the last 10 Gyr, this would translate into a $m > 0.1 M_0$ merger event over the redshift range z=[0.5-1.0] for the high mass halos here. From merger rate studies one finds that 30\\% of the stellar mass of massive galaxies with $10^{10} M_\\odot < M < 10^{11} M_\\odot$ has been assembled through mergers since z=1 (de Ravel et al. 2009, and references therein). The integrated stellar mass growth obtained can then be compared to the halo mass growth obtained here. For samples at similar redshifts we see that the average halo mass, $$, generally increases with the luminosity threshold of the sample, with a very mild hint of a decreasing galaxy satellite fraction. This implies that galaxies in the faint sample show a stronger probability of being satellites in low mass halos as compared to bright galaxies in massive halos. We also find that the satellite fraction or average number of satellite galaxies appears to slowly increase over the redshift interval [0.5,1.0], but a stronger increase by a factor of $\\sim$3 over z=[0.1,0.5] is seen. This can be understood in terms of the dynamical friction that subhalos hosting satellite galaxies encounter within their host halos. The efficiency of dynamical friction depends on the relative subhalo to halo mass. Subhalos experience more efficient dynamical friction in low mass halos, which can be thought of as progenitor halos at high redshift. The subhalos are continuously subjected to tidal stripping and gravitational heating within the dense environments and get eroded if not completely. As time evolves the halo accretes mass and undergoes mergers with other halos. The subhalos that form as remnants of halo mergers are now more likely to remain intact within the higher mass halo, in turn leading to a larger number of satellite galaxies in present-day halos. A comparison with the SDSS results shows a few interesting features. The value for $M_{min}$, which is the mass of a halo hosting at least one central galaxy on average, in $40 \\%$ of the luminosity threshold VVDS samples is similar to values for local SDSS galaxies. Whereas, $M_1$ is generally higher for VVDS galaxies as compared to what is seen locally. The ratio of $M_1/M_{min}$ is found to be considerably higher (almost a factor of 2) in the VVDS as compared to the SDSS results. This shows that in order to begin hosting satellite galaxies, halos at high redshift need to accrete a larger amount of mass than is seen locally. Hence one would observe roughly twice as many local satellite galaxies than high redshift ones within the same evolved halo mass. This is another line of evidence in favor of the lower observed satellite fraction at high redshift and high local satellite fraction. This interpretation is highly simplified in light of the fact that the results have been obtained with data taken in different restframe bands. In order to investigate further and better constrain the mass growth and evolution in the number of satellite galaxies per halo over a larger redshift range, one needs to have samples from the same survey at low redshifts. This can be done with samples from deeper and wider redshift surveys. Here we have concentrated on luminosity-threshold samples leading to a link between the luminosity of galaxies and the underlying dark matter distribution. The present paper can be seen as a precursor to many studies that can be carried out with larger samples than the VVDS, including CLF - conditional luminosity function studies (e.g. van den Bosch et al. 2003, etc.), analyses with galaxy samples of different stellar masses (Zheng et al. 2007), etc.. They will certainly add to the understanding of the vast pool of underlying dark matter properties and hopefully obtain tighter constraints on models of galaxy formation." }, "1003/1003.1365_arXiv.txt": { "abstract": "% We have observed the $J=3-2$ transition of $\\mathrm{N_2H^+}$ and $\\mathrm{N_2D^+}$ to investigate the trend of deuterium fractionation with evolutionary stage in three selected regions in the Infrared Dark Cloud (IRDC) G28.34+0.06 with the Submillimeter Telescope (SMT) and the Submillimeter Array (SMA). A comprehensible enhancement of roughly 3 orders of magnitude in deuterium fractionation over the local interstellar $\\mathrm{D/H}$ ratio is observed in all sources. In particular, our sample of massive star-forming cores in G28.34+0.06 shows a moderate decreasing trend over a factor of 3 in the $N(\\mathrm{N_2D^+})/N(\\mathrm{N_2H^+})$ ratio with evolutionary stage, a behavior resembling what previously found in low-mass protostellar cores. This suggests a possible extension for the use of the $N(\\mathrm{N_2D^+})/N(\\mathrm{N_2H^+})$ ratio as an evolutionary tracer to high-mass protostellar candidates. In the most evolved core, MM1, the $\\mathrm{N_2H^+ \\; (3-2)}$ emission appears to avoid the warm region traced by dust continuum emission and emission of $\\mathrm{^{13}CO}$ sublimated from grain mantles, indicating an instant release of gas-phase CO. The majority of the $\\mathrm{N_2H^+}$ and $\\mathrm{N_2D^+}$ emission is associated with extended structures larger than $8^{\\prime\\prime}$ ($\\sim 0.2 \\; \\mathrm{pc}$). ", "introduction": "% In the early evolutionary stages of star formation process, sequential depletion of molecular species on grain mantles nurtures a peculiar low-temperature chemistry due to the removal of important gas-phase reactants, starting with sulfur-bearing species and followed by even volatile molecules such as CO (Bergin \\& Tafalla 2007). Besides $\\mathrm{H_2}$, $\\mathrm{N_2}$ is thought to be least affected in this condensation process and results in an enrichment of its daughter products, $\\mathrm{NH_3}$ and $\\mathrm{N_2H^+}$ (Bergin {et~al.} 2002). The removal of the gas-phase $\\mathrm{CO}$ also promotes ion-molecular reactions and induces a sharp increase in the abundance of deuterated molecules in dense cores (Millar, Bennett, \\& Herbst 1989). Indeed, an enhancement of $2-3$ orders of magnitude in the $\\mathrm{D/H}$ ratio in star-forming cores (Crapsi {et~al.} 2005; Fontani {et~al.} 2006; Pillai {et~al.} 2007) over the local interstellar value of $1.51 \\times 10^{-5}$ (Oliveira {et~al.} 2003) has been observed. In particular, the deuterium fractionation of $\\mathrm{N_2H^+}$, $D_\\mathrm{frac} \\equiv N(\\mathrm{N_2D^+})/N(\\mathrm{N_2H^+})$, in low-mass star-forming cores shows an increasing trend with dynamical age in the prestellar phase (Crapsi {et~al.} 2005) but a decreasing trend in the protostellar phase (Emprechtinger {et~al.} 2009). Chemical models anticipate $D_\\mathrm{frac}$ to be affected by a few factors such as the kinetic temperature and, for ion species, the electron abundance (Roueff {et~al.} 2005) as well as the gas-phase CO abundance (Aikawa {et~al.} 2005). The abundance of gas-phase CO is expected to decline in the prestellar phase through molecular depletion onto grain surfaces but to rise up in the protostellar phase through sublimation of ice mantles as the envelope warms up. The correlation between deuterium fractionation and CO depletion factor has been recognized in a compiled sample of prestellar and protostellar cores (Crapsi {et~al.} 2005; Emprechtinger {et~al.} 2009). In a subsample of Taurus cores, a better correlation is found and leads to the speculation of external environment being influential to the evolution of a core (Crapsi {et~al.} 2005). On the other hand, there is no evidence of a consistent behavior of $D_\\mathrm{frac}$ in the case of high-mass protostellar candidates although a clear but less dramatic enhancement has been observed in a number of massive protostellar objects (Fontani {et~al.} 2006; Pillai {et~al.} 2007). The environs of high-mass protostars may not sustain low temperature long enough to build up deuterated species as abundant as those of low-mass objects. In this study, we investigated a possible trend in the deuterium fractionation of $\\mathrm{N_2H^+}$ with dynamical age using a sample of massive star-forming cores from one single IRDC to reduce the environmental fluctuations among the selected cores. Infrared dark clouds (IRDCs) were first discovered by Infrared Space Observatory (ISO) and the Midcourse Space Experiments (MSX) through silhouette against the bright, diffuse infrared background emission of the Galactic plane (Egan {et~al.} 1998; Carey {et~al.} 1998, 2000; Hennebelle {et~al.} 2001; Rathborne, Jackson, \\& Simon 2006). Because of their large mass ($M \\gtrsim 10^3 M_\\odot$), low temperature ($T < 20 \\; \\mathrm{K}$), and high density ($n_\\mathrm{H_2} \\gtrsim 10^5 \\; \\mathrm{cm^{-3}}$), IRDCs have been proposed to be in the earliest stage of massive star formation. At a distance of $4.8 \\; \\mathrm{kpc}$, IRDC G28.34+0.06 (hereafter G28) is associated with roughly $10^3 \\; M_\\odot$ in the infrared absorption region and contains several dense cores in different evolutionary stages (Carey {et~al.} 2000; Wang {et~al.} 2008). We have selected three dense cores in G28 to form an evolutionary sequence: starting with MM9 in an early stage of mass collection with one weak continuum source, followed by MM4 in a stage of mass fragmentation with at least five continuum sources, and MM1 in a later stage with embedded massive protostellar objects with a total luminosity of $10^3 \\; L_\\odot$. All the selected regions are associated with water masers, indicating star-forming activities (Wang {et~al.} 2006). Our study is enabled by the Arizona Radio Observatory (ARO) Submillimeter Telescope (SMT) and the Submillimeter Array\\footnotemark[4] (SMA). The SMT offers good sensitivity to detect weak line emission in extended structures while the SMA can preferentially image compact structures in the regions of interest. \\footnotetext[4]{The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics, and is funded by the Smithsonian Institution and the Academia Sinica.} ", "conclusions": "% \\subsection{Deuterium Fractionation as an Evolutionary Probe} % The $J=3-2$ transition of $\\mathrm{N_2H^+}$ and $\\mathrm{N_2D^+}$ were detected in all three sources with the SMT (Fig.~\\ref{fig_g28smt}). Since both transitions contain numerous hyperfine components, we fit each spectrum of every source with a model comprised of thirty-eight hyperfine components with updated line frequencies and spontaneous emission rates (Pagani, Daniel, \\& Dubernet 2009). For each individual source, all the hyperfine components of every $J$-level are assumed to be in thermal equilibrium at a single excitation temperature, $T_\\mathrm{ex}$, adopted from the ammonia observations with an angular resolution of 40\\arcsec\\ (Pillai {et~al.} 2006). The models are described by three more parameters: total column density, $N$, systemic velocity, $\\upsilon_\\mathrm{LSR}$, full-width at half maximum (FWHM) as line width, $\\Delta \\upsilon$. Model spectra are optimized with the minimization of the reduced $\\chi^2$ value, $\\overline{\\chi^2}$, and the results are listed in Table~\\ref{table_smtfits}. The $\\mathrm{N_2H^+}$ spectrum towards MM1 appears to be doubly peaked, possibly affected by the presence of multiple sources (Zhang {et~al.} 2009) and resulted in a significantly broader line width and a large $\\overline{\\chi^2}$. \\placefigure{f1} % \\placetable{table_smtfits} % In cold clouds, a correction of the cosmic background temperature, $T_\\mathrm{bg} = 2.7 \\; \\mathrm{K}$, is necessary when extracting the optical depth information of an observed spectrum with \\begin{equation} % \\tau(\\upsilon) = -\\ln \\left[ 1 - \\frac{T_\\mathrm{mb}(\\upsilon)}{J(T_\\mathrm{ex}) - J(T_\\mathrm{bg})} \\right], \\end{equation} % where $T_\\mathrm{mb}(\\upsilon)$ is the main beam temperature of the spectra, and $J(T) = (h\\nu/k)(e^{h\\nu/kT}-1)$. All the line emissions appear to have fairly small optical depths. Meanwhile, the optimized model also provides an estimate of optical depth by integrating optical depths of all the hyperfine components. The maximum optical depth, $\\tau_\\mathrm{max}$, is less than 0.21 (Table~\\ref{table_smtfits}). The emission of all the observed lines is optically thin. Over all, the fitted line widths, $\\Delta \\upsilon$, of all the transitions are much broader than their thermal line width, $\\Delta \\upsilon_\\mathrm{th} \\equiv \\sqrt{8 \\ln 2 \\, k T_\\mathrm{ex}/m_\\mathrm{N2H^+}} \\simeq 0.16 \\; \\mathrm{km \\, s^{-1}}$, suggesting a significant contribution from nonthermal motions. Our $\\mathrm{N_2H^+ \\; (3-2)}$ spectra also have larger line widths when compared to the $\\mathrm{NH_3 \\; (1,1)}$ spectra observed with the Very Large Array (VLA) and the Effelsberg 100m telescope (Pillai {et~al.} 2007; Wang {et~al.} 2008). Since the two transitions have similar upper state energies, $E_\\mathrm{up} = 26.8 \\; \\mathrm{K}$ for $\\mathrm{N_2H^+}$ and $23.4 \\; \\mathrm{K}$ for $\\mathrm{NH_3}$, but fairly different critical densities, $n_\\mathrm{crit} \\sim 10^6 \\; \\mathrm{cm^{-3}}$ for $\\mathrm{N_2H^+}$ (Daniel {et~al.} 2005; Pagani {et~al.} 2009) and $10^3 \\; \\mathrm{cm^{-3}}$ for $\\mathrm{NH_3}$ (Evans 1999), our $\\mathrm{N_2H^+}$ observations tend to trace denser clumps possibly embedded in the inner region that is affected by star-forming activities as suggested by $\\mathrm{H_2O}$ masers (Wang {et~al.} 2006). Observations of lower $J$ transitions of $\\mathrm{N_2H^+}$ such as the $J=1-0$ ($E_\\mathrm{up} = 4.5 \\; \\mathrm{K}$ and $n_\\mathrm{crit} \\sim 10^5 \\; \\mathrm{cm^{-3}}$) line are needed to verify this interpretation. The deuterium fractionation, $D_\\mathrm{frac}$, is calculated for every source (Table~\\ref{table_smtfits}) and shows a significant enhancement of roughly 3 orders of magnitude higher than the average $\\mathrm{D/H}$ ratio of $1.51 \\times 10^{-5}$ in the local interstellar medium (Oliveira {et~al.} 2003). Such an enhancement has been observed in a large sample of high-mass star-forming cores (Fontani {et~al.} 2006) as well as low-mass prestellar and protostellar cores (Crapsi {et~al.} 2005; Roberts \\& Millar 2007; Emprechtinger {et~al.} 2009). In particular, we find a decreasing trend with evolutionary stage, from $D_\\mathrm{frac} = 0.051$ in the younger MM9 core to $D_\\mathrm{frac} = 0.016$ in the more evolved MM1 core, a change over a factor of roughly $3$. Although a decreasing trend over a factor of roughly $8$ has been identified in low-mass protostellar cores (Emprechtinger {et~al.} 2009), there was no evidence for such trend in a sample of massive star-forming cores associated with different molecular clouds (Fontani {et~al.} 2006). Variations in thermal history, external environment, and initial chemical abundances across different molecular clouds may cause undesired confusion. Although our sample contains only three sources, the association in one single IRDC helps to minimize the environmental variations among the sources and to reveal the gentle drop in deuterium fractionation as the dusty envelope warms up. \\subsection{Emission Contained in Compact Structures \\label{g28comp}} Using the SMA, we further imaged the $\\mathrm{N_2H^+}$ and $\\mathrm{N_2D^+}$ emission to study more centrally concentrated structures. The $\\mathrm{N_2H^+}$ emission shows clear detections in MM1 and MM4 as well as a marginal detection in MM9 (Fig.~\\ref{fig_g28n2hp}). Since an interferometer is insensitive to structures larger than the scale corresponding to its shortest projected baseline, our SMA observations serve as a spatial filter to probe the fractional flux contained in clumps smaller than 8\\arcsec\\ (Wilner \\& Welch 1994). A comparison of the integrated intensity observed with the SMA, $W_\\mathrm{SMA}$, to that with the SMT, $W_\\mathrm{SMT}$, can be made by convolving the SMA maps with the SMT beam after masking out regions below $-3\\sigma$, which are believed to be artifacts induced by the lack of short baselines. For every source, we compare the $\\mathrm{N_2H^+}$ integrated intensity observed with the SMA to that with the SMT (Fig.~\\ref{fig_g28smt}) and estimate the fraction of the SMA integrated intensity, $W_\\mathrm{SMA}/W_\\mathrm{SMT}$, which indicates the state of a mass concentration process (Table~\\ref{table_mms}). The small values of $W_\\mathrm{SMA}/W_\\mathrm{SMT}$ suggest that most of the $\\mathrm{N_2H^+}$ emission is in structures larger than 8\\arcsec. An increasing trend from MM9 to MM1 is found and agrees with the presumed evolutionary stage of the central sources. On the other hand, $\\mathrm{N_2D^+}$ is not detected in all regions at a level of $4\\sigma \\simeq 2.3 \\; \\mathrm{K \\, km \\, s^{-1}}$, translating to an upper limit of $N(\\mathrm{N_2D^+}) = 2.0 \\times 10^{12} \\; \\mathrm{cm^{-2}}$ at $T_\\mathrm{ex} = 13 \\; \\mathrm{K}$. Compared to the SMT integrated intensity in MM9, this detection limit sets a maximum fraction of 19\\% for emission coming from compact structures. Given the high critical density of $n_\\mathrm{crit} \\sim 10^6 \\; \\mathrm{cm^{-3}}$, the large fraction of the $\\mathrm{N_2H^+} \\; (3-2)$ emission missed by the SMA observations strongly suggests the presence of cold and dense gas in scales larger than 8\\arcsec\\ ($\\sim 0.2 \\; \\mathrm{pc}$). Similar to $\\mathrm{N_2H^+}$, most of the $\\mathrm{N_2D^+}$ emission is in extended structures. \\placefigure{f2} % \\placetable{table_mms} % Continuum emissions at $284.2$ and $236.5 \\; \\mathrm{GHz}$ are detected in all sources with similar morphology. For the first time, a compact continuum source is detected in MM9 (Fig.~\\ref{fig_g28n2hp}). Since the visibility coverages of the two observation runs were quite different, we made continuum maps with visibilities of projected baselines within $13-57 \\; \\mathrm{k\\lambda}$, the range that all sources have in common at the two frequencies. If the dust emission is optically thin, the continuum flux density $F_\\nu \\propto \\kappa_\\nu \\, B_\\nu(T_d)$, where $\\kappa_\\nu = 0.006 \\, (\\nu/245 \\; \\mathrm{GHz})^{\\beta} \\; \\mathrm{cm^2 \\, g^{-1}}$ (Shepherd \\& Watson 2002) is the dust opacity at the observing frequency $\\nu$, and $B_\\nu(T_d)$ is the Planck function at a dust temperature $T_d$. The opacity spectral index, $\\beta \\equiv \\Delta \\log \\kappa_\\nu/\\Delta \\log \\nu$, can be measured by comparing flux densities at two observing frequencies. In cold clouds, the condition $h \\nu \\sim k T_d$ makes the Rayleigh-Jeans approximation inappropriate. We estimate the opacity spectral index in the selected regions with \\begin{equation} % \\beta = \\frac{\\log (F_{\\nu_2}/F_{\\nu_1}) + \\log [(e^{h\\nu_2/kT_d}-1)/(e^{h\\nu_1/kT_d}-1)]}{\\log (\\nu_2/\\nu_1)} - 3, \\end{equation} % and the core mass with $M_\\mathrm{core} = F_{\\nu} D^2/\\kappa_{\\nu} B_{\\nu}(T_d)$, where $D$ is the distance of the source. Assuming thermal equilibrium between gas and dust over similar spatial scales in the calculations of $\\beta$ and $M_\\mathrm{core}$, we adopt the gas temperature derived from the VLA ammonia observations with a comparable angular resolution of $5\\arcsec \\times 3\\arcsec$ (Zhang {et~al.} 2009) to be the dust temperature, $T_d$. The results are listed in Table~\\ref{table_mms}. Because of the limited frequency span of $48 \\; \\mathrm{GHz}$, we note that a systematic uncertainty of 15\\% in flux density measurements will produce a fairly large uncertainty of $\\Delta \\beta \\simeq 1.2$ in our calculations. The dust opacity spectral index shows an increasing trend from MM9 to MM1 as the central sources evolve although the multiplicity in MM4 may complicate the interpretation of its averaged $\\beta$. Nevertheless, a smaller value of $\\beta$ is observed in MM9 with respect to MM1. At millimeter wavelengths, the value of $\\beta$ is a good probe for the size distribution of dust grains (Miyake \\& Nakagawa 1993), and a smaller $\\beta$ suggests that MM9 is surrounded by larger dust grains as a result of grain growth in high-density environment. On the other hand, the larger $\\beta$ in the more evolved region MM1 may be attributed to possible changes in the size distribution and chemical composition of the surrounding dust grains, which have been exposed to strong radiation fields generated by the associated massive YSOs. \\subsection{Deficiency of N$_2$H$^+$ in a Warm Region} % In cold, dense environment, CO tends to freeze out onto dust grains while $\\mathrm{N_2H^+}$ and $\\mathrm{NH_3}$ suffer less from depletion (Bergin {et~al.} 2002; Tafalla {et~al.} 2004). Since CO is the major destroyer of molecular ions, its removal from the gas phase results in a subsequent enrichment of $\\mathrm{N_2H^+}$ (Aikawa {et~al.} 2005). In later evolutionary stages when the internal heating of YSOs becomes important, the temperature will rise up and lead to sublimation of ice mantle, which returns volatile species such as CO back to the gas phase. The warmer temperature together with the reappearance of gas-phase CO can alter the competition among chemical reaction routes and destroy $\\mathrm{N_2H^+}$ that has been produced during the cold early phase (Roueff {et~al.} 2005; Aikawa {et~al.} 2005). In MM1, a deficiency of $\\mathrm{N_2H^+}$ is observed in the location of the dust continuum emission as well as the $\\mathrm{^{13}CO \\; (2-1)}$ emission (Fig.~\\ref{fig_g28mm1}) in a previous study (Zhang {et~al.} 2009). Observationally, a centrally heated temperature structure from $16$ to $30 \\; \\mathrm{K}$ is derived in MM1 over spatial scales of roughly $1$ to $0.1 \\; \\mathrm{pc}$ (Pillai {et~al.} 2006; Zhang {et~al.} 2009). Theoretically, a significant drop in $D_\\mathrm{frac}$ is predicted across this temperature range (Roueff {et~al.} 2005), and an instant release of CO is expected due to a significant drop of the CO sublimation timescale from $10^8 \\; \\mathrm{yr}$ at $T_d \\simeq 12 \\; \\mathrm{K}$ to $0.1 \\; \\mathrm{yr}$ at $T_d \\simeq 20 \\; \\mathrm{K}$ (Collings {et~al.} 2003). This hypothesized release of CO in the warm region can be traced in MM1 with the $\\mathrm{^{13}CO}$ emission and dust continuum emission. A contrast in chemical composition can occur across the boundary between the cold, outer part and the warm, inner part that has been altered by the newly released gas-phase reactants. This chemical contrast has recently been observed in AFGL~5142 by comparing emissions of $\\mathrm{N_2H^+}$ and $\\mathrm{NH_3}$ (Busquet {et~al.} 2009). \\placefigure{f3} %" }, "1003/1003.1686_arXiv.txt": { "abstract": "We consider the potential dominated era of Friedmann-Lema\\^{\\i}tre-Robertson-Walker flat cosmological models in the framework of general Jordan frame scalar-tensor theories of gravity with arbitrary coupling functions, and focus upon the phase space of the scalar field. To study the regime suggested by the local weak field tests (i.e. close to the so-called limit of general relativity) we propose a nonlinear approximation scheme, solve for the phase trajectories, and provide a complete classification of possible phase portraits. We argue that the topology of trajectories in the nonlinear approximation is representative of those of the full system, and thus can tell for which scalar-tensor models general relativity functions as an attractor. ", "introduction": "The unknown source of observed present day acceleration % of the Universe, called dark energy, is inspiring thorough investigations of different extensions of general relativity (GR) and $\\Lambda$CDM cosmology (for recent reviews see Refs. \\cite{de:theory}). The scalar-tensor theory of gravity (STG) \\cite{books} offers one such consistent possibility. Besides the usual spacetime metric tensor $g_{\\mu\\nu}$ it employs a scalar field $\\Psi$, playing the role of a variable gravitational ``constant\", to describe the gravitational interaction. In the Jordan frame STG is specified by two functions \\cite{flanagan}, e.g. a coupling $\\omega(\\Psi)$ and a scalar potential $V(\\Psi)$. In fact, a wide class of theories of gravitation, including higher order theories \\cite{sotiriou}, theories of variable speed of light \\cite{magueijo}, as well as low energy approximations of brane world models and string theories \\cite{gasperini} can be cast into the general form of STG. The weak field tests \\cite{ppn} pose a restriction to all alternative models of gravity including STG, since the Universe around us tends to be described by the Einstein tensorial gravity very precisely \\cite{constraints}. This means that only those STG models are physically viable which in their late time cosmological evolution imply local consequences very close to those of GR. Several authors have studied how general relativity acts as an attractor for a wide class of STGs, such that the Solar System weak field (PPN) constraints spontaneously come to be satisfied at late times \\cite{dn,bp,asymtotics}. In our recent papers \\cite{meie4, meie5} we have proposed a limiting process for the scalar field which describes scalar-tensor cosmological models relaxing to satisfy the Solar System constraints, with an indication for which classes of STGs the attractor behavior is realized. The methods of dynamical systems have proved to be a useful tool when explicit analytic solutions are hard to find. In STG cosmology several authors have performed the analysis for different specific choices of the coupling and potential \\cite{stg:dynsys}, while Refs. \\cite{faraoni1,burd_coley,faraoni3,meie4} study the phase space and dynamics in the general case. The STG phase space point corresponding to the limit of GR is peculiar in the sense that the standard linearization process there is hampered by ratios which turn out to be indeterminate. In our previous studies \\cite{meie4, meie5} we assumed that these indeterminate terms vanish, which allows to treat this singular point as a standard fixed point. In the present paper we focus upon the potential dominated era of Friedmann-Lema\\^{\\i}tre-Robertson-Walker (FLRW) flat cosmological models in the phase space of the (decoupled) scalar field $\\Psi$ and its (cosmological) time derivative ${\\dot \\Psi} \\equiv \\Pi$, and propose an approximation which takes into account all possible finite values of these indeterminate ratios, thus preserving the leading nonlinear term in the field equations. We give a comprehensive description of the phase space trajectories of the approximate nonlinear system near the GR point and classify all trajectories allowed by the parameters of the theory. While the topology of trajectories differs in linear and nonlinear approximation, there is a correspondence in the final asymptotics, i.e. whether the trajectories end up at the GR point or are repelled from it. We argue that the nonlinear system accurately captures the key properties of the full system of STG equations near this singular point, and the topology of trajectories of the nonlinear approximation is representative of those of the full system. Therefore for any given STG model with a reasonable coupling $\\omega(\\Psi)$ and scalar potential $V(\\Psi)$ our results predict whether GR is an attractor. The paper is organized as follows. In the next section we review very briefly the scalar-tensor theory of gravity in the potential dominated era, write down the field equations of FLRW cosmology in the form of a dynamical system, and make some general remarks about the phase space including the singular GR point. In Sect. \\ref{ae} we introduce the approximation method and present linear and nonlinear systems of equations in the neighbourhood of the singular point. In Section \\ref{pt} we present solutions (phase trajectories) of the nonlinear system and a classification of trajectories, summarized in Table 1 and illustrated on Figure 1. In Section \\ref{trivial} our claims are backed up by a simple example. Finally, Section \\ref{disc} provides a summary and some remarks for future work. ", "conclusions": "\\label{disc} This paper considers general scalar-tensor gravity (STG) in the Jordan frame with arbitrary coupling $\\omega(\\Psi)$ and potential $V(\\Psi)$. We have presented and justified an approximate theory for the behavior of the scalar field in flat FLRW cosmological STG models in a potential dominated era and in the regime where the local weak field experiments are satisfied. In terms of the phase space ($\\Psi$, $\\dot{\\Psi} \\equiv \\Pi$) the latter is understood as the neighbourhood of the `GR point' ($\\Psi_{\\star}$, $\\Pi_{\\star}$), defined by (a) $\\frac{1}{2 \\omega(\\Psi_{\\star}) +3} = 0$, (b) $\\Pi_{\\star}=0$. We propose that if (c) $\\frac{d}{d \\Psi}(\\frac{1}{2 \\omega(\\Psi_{\\star})+3}) \\neq 0$ and (d) the higher derivatives of $\\frac{1}{2 \\omega(\\Psi_{\\star})+3}$ do not diverge, then in the neighbourhood of the GR point the nonlinear system (\\ref{mlin_x}), (\\ref{mlin_y}) can be considered as an adequate approximate description of the full dynamical system (\\ref{dynsys_x}), (\\ref{dynsys_y}), since both are endowed with the same characteristic features. The phase portraits, summarized in Table 1 and depicted on Figure 1, typically show many trajectories passing through the GR point either once on multiple times. In the end, only if \\begin{displaymath} V(\\Psi_{\\star})>0 \\,, \\qquad \\quad \\frac{d}{d \\Psi}\\left( \\frac{1}{2 \\omega(\\Psi)+3} \\right)\\Bigg|_{\\Psi_{\\star}} \\left(2V(\\Psi)-\\frac{dV(\\Psi)}{d\\Psi}\\Psi \\right) \\Bigg|_{\\Psi_{\\star}} <0 \\, \\end{displaymath} does the GR point function as an asymptotic attractor for the flow of all trajectories in the vicinity. These analytic results could not have been predicted by numerical simulations, as the numerical calculations become rather problematic near the GR point due to the indeterminacy present in the equations. It would be very intersting to study how the different looping behaviors through the GR point manifest themselves in terms of observational predictions. It would also be of obvious physical relevance to extend the analysis to the matter dominated case, although the treatment of the problem would face a difficulty of having an additional phase space dimension to deal with. \\bigskip {\\bf Acknowledgments} \\smallskip This work was supported by the Estonian Science Foundation Grant No. 7185 and by Estonian Ministry for Education and Science Support Grant No. SF0180013s07. M.S. also acknowledges the Estonian Science Foundation Postdoctorial research Grant No. JD131. \\medskip" }, "1003/1003.4892_arXiv.txt": { "abstract": "In this article we give a full description of the dynamics of the flat anisotropic (4+1)-dimensional cosmological model in the presence of both Gauss-Bonnet and Einstein contributions. This is the first complete description of this model with both terms taken into account. Our data is obtained using the numerical analysis, though, we use analytics to explain some features of the results obtained, and the same analytics could be applied to higher-dimensional models in higher-order Lovelock corrections. Firstly, we investigate the vacuum model and give a description of all regimes; then, we add a matter source in the form of a perfect fluid and study the influence the matter exerts upon the dynamics. Thus, we give a description of matter regimes as well. Additionally, we demonstrate that the presence of matter not only ``improves'' the situation with a smooth transition between the standard singularity and the Kasner regime, but also brings additional regimes and even partially ``erases'' the boundaries between different regimes inside the same triplet. Finally, we discuss the numerical and analytical results obtained and their generalization to the higher-order models. ", "introduction": "The idea of extra dimensions traces back to the beginning of last century, to the papers by Nordstr\\\"om~\\cite{Nord1914}, Kaluza~\\cite{Kaluza} and Klein~\\cite{Klein}. At that time such ideas were considered as mere mathematical speculations, but with time they have come to hold a firm place in the minds of scientists. The 70s gave rise to an interest in the extradimensional theories due to the development of superstring and supergravity theories (see, e.g., \\cite{super}). Later in the 90s the interest has only increased owing to the possibility to solve the hierarchy problem~\\cite{hierarchy}, lower the grand unification scale in the M-theory~\\cite{M-GUT} and in many others (see, e.g.,\\cite{rec-rev} for recent reviews). Dealing with extra dimensions in the cosmological context, one usually uses a modified gravity (see the forementioned reviews). One of modifications of this kind is the Lovelock gravity~\\cite{Lovelock}. This theory is equivalent to Einstein's theory in (3+1), but starting from (4+1) it gives rise to higher-order (in powers of curvature) corrections. The first of these corrections is known as the Gauss-Bonnet term, first found by Lanczos~\\cite{Lanczos} (therefore it is sometimes referred to as the Lanczos term). In the cosmological context the Lovelock (and Gauss-Bonnet in particular) gravity was intensively studied over the past two decades~\\cite{studied}. In this paper we investigate the dynamics of the flat anisotropic Universe in Lovelock gravity. Since we work in a (4+1)-dimensional case, Lovelock contribution reduces to the Gauss-Bonnet term. We took both the General Relativity (GR) and Gauss-Bonnet (GB) contributions into account for a reason. Namely, with only one of them taken into account, some exact solutions (see, e.g., \\cite{il1, we2, we3, iv1, iv2}) could be obtained (and that is true not just for GB, but for any order of Lovelock corrections and in any number of dimensions~\\cite{prd}). Though when we take into account different contributions, the equations become much more complicated making it almost impossible to obtain any exact solution in general case (for the first time solutions with both terms taken into account were found in~\\cite{deser}; in~\\cite{we3} a cosmological solution in exponential form was obtained). Yet, the influence of the other (than leading) terms could be significant -- for instance, in~\\cite{we2}, while investigating a (4+1)-dimensional model with matter in the form of a perfect fluid in a ``pure'' Gauss-Bonnet case, we discovered an unusual behavior in $w < 1/3$ case, and we believe it was caused by neglecting the GR term in the equations of motion. The (4+1)-dimensional case with both GR and GB contributions was studied in~\\cite{ind}, but a full analysis of all regimes was not performed there. In this paper we are going to do this -- we will give a description of all regimes, as well as try to explain some of them analytically. The second aim of this paper is to investigate the regimes in presence of matter in the form of a perfect fluid. We reported some features of the influence of matter in~\\cite{we3}, but no actual description was given. Namely, we reported that the presence of matter leads to an increase of the ``probability'' of smooth transition between the low- and high-energy Kasner regimes. In the current paper we give a full description of all types of transitions as well as describe the influence of the matter in large. Finally, we generalize the results obtained -- both for the vacuum and matter cases -- to higher even-dimensional models with all possible Lovelock corrections taken into account. This article is the logical continuation of~\\cite{prd}. Indeed, in~\\cite{prd} we considered the flat anisotropic cosmological models in the Lovelock gravity, derived the equations of motion and investigated the dynamics with only the highest Lovelock correction taken into account. In this paper we investigate the lowest model ($D=4$, $n=2$) with all corrections taken into account and based on the results predict the behavior of even-dimensional models with all possible Lovelock corrections considered. ", "conclusions": "In this paper we gave the first complete description of all the regimes in the (4+1)-dimensional flat anisotropic cosmological model in the Einstein-Gauss-Bonnet gravity. All through the paper we hold the condition on the initial values of Hubble parameters: $\\sum H_i^{(0)} > 0$. With the $\\sum H_i^{(0)} < 0$ case taken into account, our description will be complete. And the $\\sum H_i^{(0)} < 0$ case could be easily obtained from the $\\sum H_i^{(0)} > 0$ case: indeed, assuming $\\sum H_i^{(0)} < 0$ we effectively only make a $t\\to -t$ transform, so we will have the same, but time-reversed results, such as, ``standard singularity $\\to$ Kasner'' turns into ``Kasner $\\to$ recollapse'' and so on. With this remark in mind, our description of the regimes is complete for all the possible values of the initial Hubble parameters. We have found that the type-II transition (``standard singularity $\\to$ recollapse'') dominates over all the other types of trajectories in the vacuum case, and describe it both qualitatively and quantitatively. We as well found that the type-I transition (``standard singularity $\\to$ Kasner regime'') is spread more widely than it was assumed~\\cite{ind}, but still not enough to compete even with type IV. If we add matter in the form of a perfect fluid, the situation changes drastically. Not only the abundance of the type-I transition becomes comparable with abundances of other types~\\cite{we3}, but we also obtain new regimes that are unknown for the vacuum case\\footnote{Let us note -- considering the $\\sum H_i^{(0)} < 0$ case would not bring these two cases for they are time-reversal to each other.} (type-III and V). With these two, in the matter case we have all the possible transitions between two possibilities for the past (standard and nonstandard singularities) and three for the future (the Kasner regime, recollapse and nonstandard singularity) evolutions. We gave the description of the influence of matter on all the initially possible vacuum regimes. Apart from the description of the matter regimes, we found that at intermediate--high densities all the regimes (inside their triplet) look alike: positive triplets have III$\\to$I$\\to$IV ``regime stratification'' (with a decrease of $w$) while negative ones have VI$\\to$II ``regime sequence'' (with an increase of $\\rho$). Thus, the influence of matter is stronger than we anticipated: it does not just significantly increase the abundance of the type-I transition and adds the two remaining possible trajectories, but also almost ``erases the bounds'' between different trajectories inside the triplet. But there are exceptions like II$-$ with $0 < w_{st} < 1/3$ that we discussed in the relevant section. Finally we generalized the obtained results on the higher-dimensional case. We expect that the vacuum even-dimensional flat anisotropic models in the Lovelock gravity of the $n$ order are dominated by the type-II transition. The presence of the other types of transitions is in general more and more suppressed with the growth of $n$. The presence of matter acts the same (as described for (4+1)) way in the case of negative multiplets and the same but with decreases of the presence of type-I trajectories with the growth of $n$ in the case of positive multiplets. Therefore, the even-dimensional vacuum models are degenerative from the point of view of the presence of a smooth transition between the standard singularity and the cosmological expansion; the presence of matter in the form of a perfect fluid lifts this degeneration but with the growth of $n$ this improvement decreases. Let us note, though, that ($2+1$) in $L_1$ differs from the described above picture~\\cite{2plus1einst}; hence, the analysis is valid only a t $n\\geqslant 2$." }, "1003/1003.1962_arXiv.txt": { "abstract": "We report on the results of two epochs of Very Long Baseline Array (VLBA) observations of the 22 GHz water masers toward IRAS 19190+1102. The water maser emission from this object shows two main arc-shaped formations perpendicular to their NE-SW separation axis. The arcs are separated by $\\sim$280 mas in position, and are expanding outwards at an angular rate of 2.35 mas yr$^{-1}$. We detect maser emission at velocities between -53.3 km s$^{-1}$ to +78.0 km s$^{-1}$ and there is a distinct velocity pattern where the NE masers are blueshifted and the SW masers are redshifted. The outflow has a three-dimensional outflow velocity of 99.8 km s$^{-1}$ and a dynamical age of about 59 yr. A group of blueshifted masers not located along the arcs shows a change in velocity of more than 25 km s$^{-1}$ between epochs, and may be indicative of the formation of a new lobe. These observations show that IRAS 19190+1102 is a member of the class of ``water fountain'' pre-planetary nebulae displaying bipolar structure. ", "introduction": "Intermediate mass stars (1-8 $M_{\\odot}$) evolve from being asymptotic giant branch (AGB) stars into planetary nebulae (PNs) via a short transition phase during which the stars are classified as pre-planetary nebulae (PPNs). AGB stars lose mass from a slow, dense wind with expansion velocities of 5--30 km s$^{-1}$ and exhibit roughly spherically symmetric circumstellar shells; however, PNs are often observed to have aspherical morphologies, including multipolar and elliptical structures (e.g., \\citealt{sah98}; \\citealt{sah00}; \\citealt{sah07}). Presumably it is during the PPN stage that some mechanism, responsible for the shaping of the wide variety of PN morphologies, becomes operational. Due to the short lifetime of PPNs, few objects have been studied during this specific stage and the details of the evolution of an AGB star into a PN still remain unclear. Understanding the short PPN stage is fundamental for understanding the final stages of intermediate mass stellar evolution. As PPNs exhibit bipolar outflows (e.g., \\citealt{sah07}), it is now believed that jets are important in shaping multipolar PNs \\citep{sah98}. Detailed observations of individual objects can provide information about the prevailing physical conditions under which the PN morphologies are formed, and about the progenitor star. To study the effects of jets in these objects, kinematical information concerning the outflows is crucial. Radio H$_{2}$O, OH, and SiO maser line emission can be used to aid classification of these objects in their late stages of evolution. In ``water fountain'' nebulae, high-velocity H$_{2}$O masers (velocity spreads $>$ 50 km s$^{-1}$) are believed to trace high-velocity outflows. Such high-velocity water masers were first discovered in \\objectname{IRAS 16342-3814} \\citep{lik88}, and several more water fountain PPN candidates have been discovered via their single dish spectra. High angular resolution observations of water masers in IRAS 16342-3814, OH 12.8-0.9, IRAS 19134+2131, W43A, and IRAS 16552-3040 \\citep{cla09,bob07,ima07,sua08} have confirmed their classification as water fountains, by showing a spatial and kinematical structure consistent with bipolar lobes. Individual maser features persist on 1--3 year periods \\citep{eng02} and can thus be used to trace dynamics of the gas. Accurate distance measures to PPNs as well as PNs in our galaxy are sparse, thus limiting the accuracy of derived stellar properties. Observations using the VLBA allow high precision ($<$ 1 mas at 22 GHz) astrometric studies of the maser features in PPNs, therefore enabling the possibility of performing trigonometric parallax measurements. The measured proper motions of the masers is a combination of three components: motion of the gas relative to the central object, peculiar motion of the source within the galaxy, and parallax. By measuring these motions we can determine distances to and dynamical ages of PPNs, thereby affording further understanding of the transition from the AGB to PN phase, including luminosity, mass, and mass loss rate. \\objectname{IRAS 19190+1102} is an OH and H$_{2}$O source with {\\it Infrared Astronomical Satellite} ({\\it IRAS}) colors characteristic of evolved stars ([25-60]$\\lesssim$1.5 and [12-25]$\\gtrsim$1.4, where [a-b]=2.5log($S_{b}/S_{a}$)). \\citet{lik89} first observed the OH \\& H$_{2}$O emission in this source, finding the velocity range of the H$_{2}$O emission spanned more than 70 km s$^{-1}$, atypical of OH/IR stars. In this paper, we present high spatial resolution data identifying IRAS 19190+1102 as a water fountain PPN. We give initial results (two epochs) of observations of H$_{2}$O maser emission from IRAS 19190+1102, as well as a brief discussion of its far-infrared characteristics, and discuss its classification as a water fountain PPN. This source is our first target to which we will perform parallax measurements on (to be presented in a subsequent paper). \\begin{figure} \\centering \\includegraphics[angle=270,scale=0.8]{f1.eps} \\caption{2004 March 19 (Epoch 1) Top panel: log plot of the spectral distribution of maser emission in IRAS 19190+1102. Bottom panel: spatial distribution of masers in IRAS 19190+1102, relative to the feature at (0,0).\\label{fige1}} \\end{figure} ", "conclusions": "Using the VLBA, we have observed the H$_{2}$O masers toward IRAS 19190+1102 with very high angular resolution; these observations indicate that the high-velocity masers span a velocity range $\\Delta$V$>$130 km s$^{-1}$ and lie in a bipolar structure. The afore mentioned properties indicate that IRAS 19190+1102 is one of six known ``water fountain'' sources, i.e. very young pre-planetary nebulae in which the presence of collimated, high-velocity outflows is manifest. The arclike structures appear to lie perpendicular to a NE-SW axis, with the blueshifted features to the north and redshifted features to the south. The arclike distributions are suggestive of bow shocks produced by a collimated jet colliding with surrounding material. The proper motions of the H$_{2}$O masers suggest that the dynamical age of the jet is $\\sim$59 yr. Additionally, a group of blueshifted masers not located along the arcs was detected in both epochs. These ``central\" masers showed a dramatic change ($>$25 km s$^{-1}$) in velocity, which we suggest may signal the beginning of a new lobe due to an episodic or precessing jet. We have obtained five additional epochs of IRAS 19190+1102 using the VLBA, using the 22 GHz line of H$_{2}$O and the 1612, 1665, 1667, 1720 MHz OH lines. With these observations we anticipate establishing an accurate estimate to the distance of IRAS 19190+1102 via parallax measurements, and thus more precise estimates of its stellar properties." }, "1003/1003.2961_arXiv.txt": { "abstract": "The Antarctic Impulsive Transient Antenna (ANITA) completed its second Long Duration Balloon flight in January 2009, with 31 days aloft (28.5 live days) over Antarctica. ANITA searches for impulsive coherent radio Cherenkov emission from 200 to 1200~MHz, arising from the Askaryan charge excess in ultra-high energy neutrino-induced cascades within Antarctic ice. This flight included significant improvements over the first flight in payload sensitivity, efficiency, and flight trajectory. Analysis of in-flight calibration pulses from surface and sub-surface locations verifies the expected sensitivity. In a blind analysis, we find 2 surviving events on a background, mostly anthropogenic, of $0.97\\pm0.42$ events. We set the strongest limit to date for $10^{18}-10^{21}$~eV cosmic neutrinos, excluding several current cosmogenic neutrino models. ", "introduction": "The existence of cosmic-ray particles of energies above $10^{19}$~eV, first established in the early 1960's, has become a problem of the first rank in particle astrophysics. Models for their production must generate particle energies many orders of magnitude higher than achievable on Earth, and these models in turn require extreme source physics that has not yet been formulated in a self-consistent manner. Even more problematic, the propagation of such particles is limited by strong energy loss due to the Greisen-Zatsepin-Kuzmin (GZK) process~\\cite{Greisen}. Hadrons are produced via the Delta-photoproduction resonance by interactions with cosmic microwave background (CMB) photons. This GZK cutoff in energy limits the propagation distance of the ultra-high energy cosmic rays (UHECRs) to within 100-200~Mpc in the current epoch, and severely distorts the observed energy spectrum. Astronomy using charged-particle UHECRs is thus limited to the local universe, suffering from both the loss of source spectral information, and the difficulty in back-tracing UHECRs through intergalactic magnetic fields. At distances of several Gpc, corresponding to the star-formation maximum at redshift $z \\sim 1$, the higher energy and density of the CMB photons leads to even greater restrictions on cosmic-ray propagation, and a more rapid energy loss to Delta photoproduction. However, information about the source particles does survive in the form of secondary neutrinos in the decay chain, known as the ultra-high energy (UHE) cosmogenic neutrinos, first described by Berezinsky \\& Zatsepin (BZ)~\\cite{BZ}. Their momenta are unaffected by magnetic fields, and they propagate without energy loss directly to Earth, retaining information about the cosmic distribution of UHECRs and their sources. The ANITA Long Duration Balloon experiment was designed to search for cosmogenic neutrinos via electromagnetic cascades initiated by the neutrinos in Antarctic ice, the most massive body of accessible, solid, radio-transparent dielectric material on Earth. We previously placed limits on the UHE cosmic neutrino flux from the first flight of ANITA~\\cite{ANITA-1}, and provided a separate detailed description~\\cite{ANITA-inst} of the ANITA instrument, flight system, data acquisition, and analysis methods. In this article we detail upgrades and augmentations beyond the instrument and methodology previously reported. ", "conclusions": "Our model-independent~\\cite{Anch02, FORTE04} 90\\%~CL limit on neutrino fluxes is based on the 28.5 day livetime, energy-dependent analysis efficiency (68\\%-42\\% from $10^{18}$-$10^{23}$~eV), the average acceptance from the two independent simulations,~\\cite{ANITA-inst}\\footnote{The acceptance used from $10^{18}$-$10^{23}$~eV in half-decade energy steps is: $4.3\\times10^{-4}$, $5.0\\times10^{-2}$, 0.92, 6.6, 36., 108., 259., 602., 1140, 1950, and 3110~km$^2$-sr} and $0.97\\pm0.42$ expected background events including the systematic effects described above. Relative to the revised ANITA-I limit~\\cite{hooverThesis,ANITA-1} shown in Fig.~\\ref{lim10}, the {\\it expected} limit from this data, in the absence of signal, is a factor of four more sensitive. We set the {\\it actual} limit, shown in Fig.~\\ref{lim10}, using our 2 observed candidates. Because ANITA-II saw more than the expected background, % the {\\it actual} limit is only a factor of two better than ANITA-I even though the {\\it a priori} sensitivity is four times higher for ANITA-II. The ANITA-II limit supercedes the ANITA-I limit and would not significantly be improved by combining the results. \\begin{figure}[ht!] \\begin{center} \\includegraphics[width=3.2in]{submission-fig-5.eps} \\\\ \\caption{ANITA-II limit for 28.5 days livetime. The red curve is the expected limit before unblinding, based on seeing a number of candidates equal to the background estimate. The blue curve is the actual limit, based on the two surviving candidates. Other limits are from AMANDA~\\cite{AMANDA08}, RICE~\\cite{RICE06}, Auger~\\cite{Auger07}, HiRes~\\cite{Hires08}, and a revised limit from ANITA-I~\\cite{hooverThesis}. The BZ (GZK) neutrino model range is determined by a variety of models~\\cite{PJ96,Engel01,Kal02,Kal02a,Aramo05,Ave05,Barger06}.} \\label{lim10} \\end{center} \\end{figure} Table~\\ref{table2} gives integrated event totals for a range of cosmogenic neutrino models with widely varying assumptions. We also include for reference the expected number of events for a pure power-law neutrino spectrum that matches the Waxman-Bahcall flux bounds for both evolved and standard UHECR sources~\\cite{WB}. ANITA-II's constraint on cosmogenic neutrino models strongly excludes models with maximally energetic UHECR source spectra which saturate other available bounds~\\cite{Yosh97,Kal02,Aramo05}. These models generally assume very flat source energy spectra which may extend up to $10^{23}$~eV; our results are incompatible with a combination of both of these features. ANITA-II is now probing several models with strong source evolution spectra that are plausible within current GZK source expectations~\\cite{Kal02,Aramo05,VB05,Barger06,Yuksel07}, some at $>90$\\% confidence level. The ANITA-II 90\\% CL integral flux limit on a pure $E^{-2}$ spectrum for $10^{18}$~eV $\\leq E_{\\nu} \\leq 10^{23.5}$~eV is $E_{\\nu}^2 F_{\\nu} \\leq 2 \\times 10^{-7}$~GeV~cm$^{-2}$~s$^{-1}$~sr$^{-1}$. These differential and integral limits, as well as the individual model limits above, are the strongest constraints to date on the cosmogenic UHE neutrino flux." }, "1003/1003.5546_arXiv.txt": { "abstract": "Observations of the cores of nearby galaxy clusters show H$\\alpha$ and molecular emission line filaments. We argue that these are the result of {\\em local} thermal instability in a {\\em globally} stable galaxy cluster core. We present local, high resolution, two-dimensional magnetohydrodynamic simulations of thermal instability for conditions appropriate to the intracluster medium (ICM); the simulations include anisotropic thermal conduction along magnetic field lines and adiabatic cosmic rays. Thermal conduction suppresses thermal instability along magnetic field lines on scales smaller than the Field length ($\\gtrsim$10 kpc for the hot, diffuse ICM). We show that the Field length in the cold medium must be resolved both along and perpendicular to the magnetic field in order to obtain numerically converged results. Because of negligible conduction perpendicular to the magnetic field, thermal instability leads to fine scale structure in the perpendicular direction. Filaments of cold gas along magnetic field lines are thus a natural consequence of thermal instability with anisotropic thermal conduction. This is true even in the fully nonlinear regime and even for dynamically weak magnetic fields. The filamentary structure in the cold gas is also imprinted on the diffuse X-ray emitting plasma in the neighboring hot ICM. Nonlinearly, filaments of cold ($\\sim 10^4$ K) gas should have lengths (along the magnetic field) comparable to the Field length in the cold medium $\\sim 10^{-4}$ pc! Observations show, however, that the atomic filaments in clusters are far more extended, $\\sim 10$ kpc. Cosmic ray pressure support (or a small scale turbulent magnetic pressure) may resolve this discrepancy: even a small cosmic ray pressure in the diffuse ICM, $\\sim 10^{-4}$ of the thermal pressure, can be adiabatically compressed to provide significant pressure support in cold filaments. This is qualitatively consistent with the large population of cosmic rays invoked to explain the atomic and molecular line ratios observed in filaments. ", "introduction": "The thermal instability has been studied extensively in the context of the interstellar medium (ISM; \\citealt{fie65,koy00,san02,kri02,pio04,aud05}) and the formation of solar prominences \\citep[e.g.,][]{kar88}, but its role in galaxy clusters has not received as much attention. The cooling time in the intracluster medium (ICM) near the centers of galaxy clusters may be as short as 10-100 Myr. Observations show that there is a dramatic {\\em lack} of plasma below $\\sim 1$ keV \\citep[e.g.,][]{pet03}, inconsistent with the prediction of the original cooling flow models \\citep[e.g.,][]{fab94}. In addition, the star formation rate in the central galaxy is 10--100 times smaller than if the gas cooled at the predicted rate \\citep[e.g.,][]{ode08}. This implies that cooling in the intracluster medium (ICM) is balanced by some form of heating that maintains an approximate global thermal equilibrium. Feedback from a central Active Galactic Nucleus (AGN) is an energetically plausible source of the required heating (e.g., \\citealt{guo08}). However, precisely how the AGN provides this heating is not understood in detail; nor are other heating mechanisms ruled out. Many clusters with a short central cooling time ($\\lesssim$ 1 Gyr; or equivalently a low central entropy) show both star formation and H$\\alpha$ emission, indicative of cool plasma at $\\lesssim 10^4$ K \\citep[e.g.,][]{cav08}. Even if heating balances cooling in a global sense, the ICM plasma is expected to be {\\em locally} thermally unstable because of the form of the cooling function \\citep[e.g.,][]{fie65}. This local thermal instability is an attractive mechanism for producing the H$\\alpha$ and molecular filaments seen in clusters with short cooling times. Recently, the atomic and molecular filaments in the core of the Perseus cluster have been spatially resolved \\citep[e.g.,][]{con01,sal06}. Based on the narrowness and coherence of these filaments, \\citet{fab08} suggested that magnetic fields play a critical role in the dynamics of the filaments. In addition to the possible role of magnetic pressure and tension, the magnetic field also modifies the microscopic transport processes in the ICM because the mean free path along magnetic field lines is orders of magnitude larger than the gyroradius. As a result, thermal conduction is primarily along magnetic field lines \\citep[][]{bra65}. This paper centers on carrying out magnetohydrodynamic (MHD) simulations of thermal instability with thermal conduction along magnetic field lines. We focus on understanding the physics of the thermal instability in the ICM, rather than on making detailed comparisons with observations. Throughout this paper we ignore the possible presence of a background gravitational field. This allows us to study the physics of the thermal instability without the added complication of buoyant motions, inflow, etc. In addition to including the effects of anisotropic heat transport, we also include cosmic rays as a second fluid. Even an initially small cosmic ray pressure can become energetically important in cold filaments. Part of the motivation for including cosmic rays is that a significant population of energetic ions ($\\gg$ eV, the temperature of optical filaments) appear required to explain the observed line ratios in filaments in galaxy clusters \\citep[][]{fer09}. The multiphase nature of the ICM is physically analogous to the well-studied multiphase ISM. The ISM has three dominant phases: a molecular phase at $\\sim$ 100 K, an atomic phase at $\\sim 10^4$ K, and the hot phase at $10^6$ K. The cooling function in the ISM is thermally bistable with thermally stable phases at $\\sim 100$ K and $\\sim 10^4$ K. The hot phase is thermally unstable but is probably maintained at its temperature by supernova heating \\citep[][]{mck77}. The same physical considerations apply for the ICM, except that the hot phase is maintained by a still poorly understood heating process (e.g., AGN feedback). This paper is organized as follows. Section 2 summarizes our model equations and the results of a linear stability analysis including conduction along field lines, cosmic rays, and magnetic fields (see appendix A). Section 3 presents the numerical set-up and the results of our numerical simulations. Section 4 discusses the astrophysical implications of our results. ", "conclusions": "" }, "1003/1003.0680_arXiv.txt": { "abstract": "We examine the nuclear morphology, kinematics, and stellar populations in nearby S0 galaxy NGC~404 using a combination of adaptive optics assisted near-IR integral-field spectroscopy, optical spectroscopy, and HST imaging. These observations enable study of the NGC~404 nucleus at a level of detail possible only in the nearest galaxies. The surface brightness profile suggests the presence of three components, a bulge, a nuclear star cluster, and a central light excess within the cluster at radii $<$3~pc. These components have distinct kinematics with modest rotation seen in the nuclear star cluster and counter-rotation seen in the central excess. Molecular hydrogen emission traces a disk with rotation nearly orthogonal to that of the stars. The stellar populations of the three components are also distinct, with half of the mass of the nuclear star cluster having ages of $\\sim$1~Gyr (perhaps resulting from a galaxy merger), while the bulge is dominated by much older stars. Dynamical modeling of the stellar kinematics gives a total nuclear star cluster mass of $1.1 \\times 10^7$~M$_\\odot$. Dynamical detection of a possible intermediate mass black hole is hindered by uncertainties in the central stellar mass profile. Assuming a constant mass-to-light ratio, the stellar dynamical modeling suggests a black hole mass of $< 1 \\times 10^5$~M$_\\odot$, while the molecular hydrogen gas kinematics are best fit by a black hole with mass of $4.5 ^{+3.5}_{-2.0}\\times 10^5$~M$_\\odot$. Unresolved and possibly variable dust emission in the near-infrared and AGN-like molecular hydrogen emission line ratios do suggest the presence of an accreting black hole in this nearby LINER galaxy. ", "introduction": "The centers of galaxies contain both massive black holes (MBHs) and nuclear star clusters (NSCs). The presence of MBHs has been dynamically measured in about 50 galaxies and it appears that most massive galaxies probably have a MBH \\citep[e.g.][]{richstone98,graham08b}. Nuclear star clusters are compact ($r_{eff} \\sim 5$~pc), massive ($\\sim10^7$~M$_\\odot$) star clusters found at the center of a majority of spirals and lower mass ellipticals \\citep{boker02,carollo02,cote06}. Unlike normal star clusters, they have multiple stellar populations with a wide range of ages \\citep{long02,walcher06,rossa06,seth06,siegel07}. Nuclear star clusters coexist with MBHs in some galaxies \\citep{filippenko03, seth08, shields08, graham09}. However, the nearby galaxies M33 and NGC~205 have NSCs but no apparent central black hole \\citep{gebhardt01,valluri05}, while some high mass core elliptical galaxies have MBHs but lack NSCs \\citep{cote06}. Occupation fractions and masses for MBHs in lower mass galaxies retain the imprint of the seed black holes (BHs) from which they form, information which has been lost due to subsequent accretion in higher mass galaxies \\citep[e.g.][]{volonteri08}. However, the presence and mass of MBHs in lower mass galaxies is very poorly constrained. Most galaxies are too far away to measure the dynamical effect of a BH with mass $\\lesssim$10$^6$~M$_\\odot$ (often referred to as intermediate mass black holes; IMBHs) with current instrumentation. Thus the presence of IMBHs in galaxy centers has only been inferred when AGN activity is observed \\citep[e.g.][]{filippenko89,greene04,satyapal07}. These AGN provide only a lower limit on the number of MBHs in lower mass galaxies, and the BH mass estimates from the AGN are quite uncertain. Dynamical measurements of MBHs in nearby massive galaxies have revealed that the mass of a galaxy's central MBH is correlated with its bulge mass \\citep{kormendy95,magorrian98,haring04}. The scaling of BH mass with the large-scale properties of galaxies extends to many measurable quantities including the bulge velocity dispersion \\citep{ferrarese00,gebhardt00,graham08b,gultekin09}. More recently, \\citet{ferrarese06}, \\citet{wehner06}, and \\citet{rossa06} have presented evidence that NSCs scale with bulge mass and dispersion in elliptical and early-type spiral galaxies in almost exactly the same way as MBHs. The similarity between the NSC and MBH scaling relationships led \\citet{ferrarese06} and \\citet{wehner06} to suggest that MBHs and NSCs are two different types of central massive objects (CMO) both of which contain a small fraction ($\\sim$0.2\\%) of the total galaxy mass. This correlation of the CMO mass with the large scale properties of galaxies suggests a link between the formation of the two. However, the nature of this connection is unknown, as is the relation between NSCs and MBHs. Studies of NSCs can help address these issues. Their morphology, kinematics, and stellar populations contain important clues about their formation and the accretion of material into the center of galaxies \\citep[e.g.][]{hopkins10a}. The scaling relationships of large-scale galaxy properties with the mass of the CMO might indicate that galaxy centers should be simple systems. The Milky Way center clearly shows this not to be the case. It is an incredibly complicated environment with a black hole \\citep[$M_{BH} = 4 \\times 10^6$~M$_\\odot$;][]{ghez08} and a nuclear star cluster \\citep[$M_{NSC} \\sim 3 \\times 10^7$~M$_\\odot$][]{genzel96,schodel07,trippe08} that contains stars of many ages and in different substructures including disks of young stars in the immediate vicinity of the black hole \\citep[e.g.][]{lu09,bartko09}. Only through understanding these complex structures in the Milky Way and finding other examples in nearby galaxies can we hope to fully understand MBH and NSC formation, and the links between these objects and their host galaxies. In our current survey, we are looking at a sample of the nearest galaxies ($D < 5$~Mpc) that host NSCs. We are resolving their properties using a wide range of observational data to understand how NSCs form and their relation to MBHs. In our first paper on nearby edge-on spiral NGC~4244 \\citep{seth08b}, we showed evidence that the NSC kinematics are dominated by rotation, suggesting that it was formed by episodic accretion of material from the galaxy disk. This paper focuses on the NSC and possible MBH in NGC~404, the nearest S0 galaxy. Table~\\ref{proptab} summarizes its properties. The stellar populations of the galaxy are predominantly old and can be traced out to 600\\asec~(9~kpc) \\citep{tikhonov03,williams10}. However, HI observations show a prominent nearly face-on HI disk at radii between 100-400\\asec, with detectable HI out to 800\\asec~\\citep{delrio04}. CO observations and optical color maps show the presence of molecular gas and dust within the central $\\sim$20\\asec~(300~pc) of the galaxy \\citep{wiklind90,tikhonov03}, lying primarily to the NE of the nucleus. A nuclear star cluster in the central arcsecond of NGC~404 was noted by \\citet{ravindranath01} from an analysis of NICMOS data. \\begin{deluxetable}{lr} \\tablewidth{\\columnwidth} \\tablecaption{NGC~404 Properties \\label{proptab}} \\startdata \\tableline Distance$^a$ & 3.06 Mpc\\\\ \\hspace{0.2in}$m-M$ & 27.43\\\\ \\hspace{0.2in}pc/\\asec & 14.8\\\\ Galaxy $M_{V,0}$, $M_{I,0}$$^b$ & -17.35, -18.36\\\\ Bulge/Total Luminosity$^c$ & 0.76\\\\ Bulge $M_{V,0}$, $M_{I,0}$$^d$ & -17.05, -18.06\\\\ Bulge Mass$^e$ & 9.2$\\times$10$^8$~M$_\\odot$\\\\ HI Gas Mass$^f$ & 1.5$\\times$10$^8$~M$_\\odot$\\\\ Molecular Gas Mass$^g$ & 6$\\times$10$^6$~M$_\\odot$\\\\ Central Velocity$^h$ & -58.9 km$\\,$s$^{-1}$ \\enddata \\tablecomments{$(a)$ TRGB measurement from \\citet{karachentsev04}, $(b)$ at $r < 200$\\asec~from \\citet{tikhonov03} corrected for foreground reddening, $(c)$ for $r < 200$\\asec~from \\citet{baggett98}, $(d)$ combining above values; we derive a slightly brighter bulge $M_{I,0} = -18.20$ in \\S3, $(e)$ using $M/L_I = 1.28$ from \\S5.4, $(f)$ \\citet{delrio04}, $(g)$ \\citet{wiklind90}, corrected to $D=3.06$~Mpc, $(h)$ heliocentric velocity, see \\S4.1.} \\end{deluxetable} The presence of an AGN in NGC~404 is a controversial topic. The optical spectrum of the nucleus has line ratios with a LINER classification \\citep{ho97a}. Assuming a distance of $\\sim$3~Mpc, NGC~404 is the nearest LINER galaxy; other nearby examples include M81 and NGC~4736 (M94). Recent studies have shown that most LINERS do in fact appear to be AGN, with a majority of them having detected X-ray cores \\citep{dudik05,gonzalezmartin06,zhang09}, radio cores \\citep{nagar05}, and many of them possesing mid-IR coronal lines \\citep{satyapal04} and UV variable cores \\citep{maoz05}. However, NGC~404 is quite unusual in its properties. No radio core is observed down to a limiting flux of 1.3~mJy at 15~GHz with the VLA in A array \\citep{nagar05}, however, \\citet{delrio04} do detect an unresolved 3~mJy continuum source at 1.4~GHz using the C array. A compact X-ray source is detected \\citep{lira00,eracleous02}, but its low luminosity and soft thermal spectrum indicates that it could be the result of a starburst. Signatures of O stars are seen in the UV spectrum of the nucleus, however dilution of these lines suggests that $\\sim$60\\% of the UV flux could result from a non-thermal source \\citep{maoz98}. This suggestion is supported by more recent UV observations which show that the UV emission is variable, declining by a factor of $\\sim$3 between 1993 and 2002 \\citep{maoz05}. HST observations of H$\\alpha$ show that the emission occurs primarily in a compact source 0.16\\asec~north of the nucleus and in wispy structures suggestive of supernova remnants \\citep{pogge00}. The [\\ion{O}{3}] emission has a double lobed structure along the galaxy major-axis \\citep{plana98}, and has a higher velocity dispersion than H$\\alpha$ near the galaxy center \\citep{bouchard10}. The mid-IR spectrum of NGC~404 shows evidence for high excitation consistent with other AGN \\citep{satyapal04}. In particular the ratio of the [\\ion{O}{4}] flux relative to other emission lines ([\\ion{Ne}{2}], [\\ion{Si}{2}]) is higher than any other LINERs in the \\citet{satyapal04} sample and is similar to other known AGN. However, [NeV] lines, which are a more reliable indicator of AGN activity \\citep{abel08}, are not detected. In summary, the case for an accreting MBH in NGC~404 remains ambiguous, with the variable UV emission providing the strongest evidence in favor of its existence. In this paper we take a detailed look at the central arcseconds of NGC~404 and find evidence that it contains both a massive nuclear star cluster and a black hole. In \\S2 we describe the data used in this paper including adaptive optics Gemini/NIFS observations. We then use this data to determine the morphology in \\S3, the stellar and gas kinematics in \\S4, and the stellar populations in \\S5. In \\S6 we present dynamical modeling from which we derive a NSC mass of $1.1 \\times 10^7$~M$_\\odot$ and find mixed results on detecting a possible MBH with mass $<10^6$~M$_\\odot$. In \\S7 we discuss these results, concluding in \\S8. \\newpage ", "conclusions": "This paper is the second resulting from our survey of nearby nuclear star clusters, and demonstrates the rich detail we can obtain for these objects. The NGC~404 nucleus is a complicated environment, with both a nuclear star cluster and a possible black hole. We have found that the surface brightness profile of the inner part of NGC~404 suggests the galaxy can be broken into three components: (1) a bulge that dominates the light beyond 1\\asec, with $M_{bulge} \\sim 9 \\times 10^8$~M$_\\odot$, $r_{eff} = 640$~pc, and a S\\'ersic index of $\\sim$2.5, (2) a NSC that dominates the light in the central arcsecond with $r_{eff}=10$~pc and a dynamical mass of $1.1 \\pm 0.2 \\times 10^7$~M$_\\odot$, and (3) a central excess at $r < 0\\farcs2$, composed of younger stars, dust emission, and perhaps AGN continuum. NIFS IFU spectroscopy shows that the NSC has modest rotation along roughly the same axis as the HI gas at larger radii, while the central excess counter-rotates relative to the NSC. Furthermore, molecular gas traced by H$_2$ emission shows rotation perpendicular to the stellar rotation. A stellar population analysis of optical spectra indicate that half of the stars in the NSC formed $\\sim$1~Gyr ago. Some ancient and very young ($<$10~Myr) stars are also present. This star formation history is dramatically different from the rest of NGC~404, which is dominated by stellar populations $>$5~Gyr in age. We suggest a possible scenario where the burst of star formation in the NSC $\\sim$1~Gyr ago resulted from the accretion of gas into the galaxy center during a merger. This formation scenario is quite different from the episodic disk accretion suggested by observations NSCs in late-type galaxies. Our dynamical modeling of the stellar and gas kinematics provide mixed evidence for the presence of a black hole in NGC~404. Assuming a constant $M/L$ within the nucleus, the stellar dynamical model suggests an upper limit of $1 \\times 10^5$~M$_\\odot$, as well as measuring a $M/L_I = 0.70 \\pm 0.04$ for the NSC. The gas kinematics are best fit by models including the presence of a black hole with $M_{\\rm BH}= 4.5^{+3.5}_{-2.0} \\times 10^5$ M$_\\odot$. Both dynamical BH mass estimates rely on a model for the stellar mass that we construct from HST $F814W$-band imaging. Uncertainties in the light profile (due to variability) and $M/L$ (due to stellar population changes) within the central excess are of the same order as the difference between the two black hole mass estimates. We have proposed to measure the mass model by using HST/STIS spectroscopy to resolve the stellar populations within the nucleus and additional multi-band imaging to extend this model to two dimensions. If successful we will combine this mass model with the kinematic observations presented here as well as larger-scale kinematics obtained from the MMT to more robustly determine the BH mass. In addition to the direct evidence of the BH from the dynamical models, we find two other properties of the nucleus which suggest the presence of an AGN. First, we find unresolved hot dust emission at the center of the NSC with a luminosity of $\\sim1.4 \\times 10^{38}$~ergs$\\,$s$^{-1}$. Comparison of the NIFS light profile to previous NICMOS observations suggests that this dust emission may be variable as is seen in other AGN. Second, the H$_2$ line ratios within the central arcsecond indicate thermal excitation in dense gas similar to what is seen in other AGN. Our proposed HST observations request multi-epoch UV through NIR imaging to search for definitive evidence of BH accretion in NGC~404. Our nearby NSC survey includes 13 galaxies within 5~Mpc with $M_B$ between -15.9 and -18.8 for which we will have comparable data to that presented here obtained using MMT, VLT, and Gemini. These galaxies span a wide-range of Hubble types in which we can examine the process of NSC formation; there are four early type E/S0 galaxies and 9 spirals of type Sc and later. NGC~404 represents one of the stronger cases for finding a black hole, given its LINER emission and proximity. However, we expect to be able detect or place upper limits of $\\lesssim10^5$~M$_\\odot$ on black holes in each of the sample galaxies. \\vspace{0.1in} \\noindent {\\em Acknowledgments:} We thank the referee, Jenny Greene, for helpful comments, St\\'ephane Charlot for sharing his models, Christy Tremonti for sharing her code, the staff at Gemini and MMT, NED, and ADS. AS acknowledges support from the Harvard-Smithsonian CfA as a CfA and OIR fellow, and helpful conversations with Margaret Geller, Pat C\\^ot\\'e, and Davor Krajnovi\\'c. MC and NB acknowledge support from STFC Advanced Fellowships. NN acknowledges support from the Cluster of Excellence ``Origin and Evolution of the Universe\". Partially based on observations obtained at the Gemini Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc.. Gemini data was taken as part of program GN-2008B-Q-74. {\\it Facilities:} \\facility{Gemini:Gillett (NIFS/ALTAIR)}, \\facility{HST (ACS/WFC)}, \\facility{MMT}" }, "1003/1003.0013_arXiv.txt": { "abstract": "Previous surveys in a few metal-poor globular clusters (GCs) showed that the determination of abundances for Li and proton-capture elements offers a key tool to address the intracluster pollution scenario. In this Letter, we present Na, O, and Li abundances in a large sample of dwarf stars in the metal-rich GC 47 Tucanae. We found a clear Na$-$O anticorrelation, in good agreement with what obtained for giant members by Carretta et al. While lithium and oxygen abundances appear to be positively correlated with each other, there is a large scatter, well exceeding observational errors, and no anticorrelation with sodium. These findings suggest that Li depletion, due to mechanisms internal to the stars (which are cooler and more metal-rich than those on the Spite plateau) combines with the usual pollution scenario, responsible for the Na$-$O anticorrelation. ", "introduction": "\\label{sec:intro} The traditional paradigm of globular clusters (GCs) as the textbook example of a simple stellar population is now outdated. In the last several years, a wealth of observational studies (both photometric and spectroscopic) have been carried out, revealing the quite complex nature of this old Galactic population. With the famous exception of $\\omega$~Cen (Freeman \\& Rodgers 1975), and of the more recently scrutinized M~54 (Carretta et al. 2010), M22 (Marino et al. 2009), and Terzan 5 (Ferraro et al. 2009), GCs show homogeneous compositions in the iron peak and heavier $\\alpha$ elements (e.g., Ca and Ti). Abundance variations in the lighter elements, namely Li, C, N, O, Na, Mg and Al, have been instead recognized since the 1970's (e.g. Cohen 1978; see Gratton, Sneden \\& Carretta 2004 for a recent review). These peculiarities of GC stars require that some material must have been processed through the complete CNO cycle in hot H burning, through proton-capture reactions (Denisenkov \\& Denisenkova 1990): element pairs C and N, O and Na, and Mg and Al are anticorrelated, the abundances of C, O and Mg being depleted while those of N, Na and Al being enhanced. A previous generation of stars, which synthesized proton-capture elements in their interiors, is now commonly accepted as responsible for the above-mentioned chemical signatures. Regardless of the debated nature of these element polluters (asymptotic giant branch (AGB) stars experiencing hot bottom burning, e.g., Ventura \\& D'Antona 2009, or fast rotating massive stars (FRMS), e.g., Decressin et al. 2007), two fundamental observational facts deserve to be emphasized. (1) All the GCs surveyed so far show the Na$-$O anticorrelation (Carretta et al. 2009b): this indicates the presence of at least two populations (neither coeval nor chemically homogeneous) within each cluster. (2) The primordial nature of such a phenomenon is uniquely indicated by the presence of these chemical signatures also among not evolved turnoff (TO) or scarcely evolved (subgiant, SGB) members (see, e.g., Gratton et al. 2001). In this context, Li abundances play a fundamental role. In fact, this element can be easily destroyed in stellar interiors (starting at T$_{\\rm burn}$$\\approx$2.5$\\times$10$^6$~K); and since the CNO and NeNa cycles require much higher temperatures, it is expected that in the site(s) where these cycles occur no Li is left. In particular, the Na-poor, O and Li rich stars, that are the first population born in the cluster, share the chemical composition of field stars of the same metallicity, while the Na-rich (Li- and O-poor) stars form from gas progressively enriched by ejecta (which are rich in Na, depleted in O and Li) of the first generation. As a consequence, if primordial and processed material are mixed in different proportions, then Li and Na (Li and O) are expected to be anticorrelated (correlated) with each other. Measuring the Li abundances in unevolved stars provides a direct indication of the amount of pristine, and by difference of polluted (CNO-cycle processed), material present in each star. This makes Li a unique tracer of the dilution process which took place in the GC, supplying also fundamental insights on its early evolutionary stages. Another important point is that Li abundances can provide strong observational constraints on the origin of the polluters: since AGB stars might also produce Li (via the Cameron$-$Fowler mechanism, Cameron \\& Fowler 1971), while massive stars can only destroy it, if Li-rich stars are present in GCs, AGB stars would be definitely favored with respect to FRMS. To date, only three GCs have been surveyed for correlations between Li and proton-capture elements. Pasquini et al. (2005) obtained Li abundances in nine TO members of NGC 6752: they found a depletion reaching down to $\\sim$1 dex below the Spite plateau values, clearly anticorrelated with Na abundances. A similar result was obtained by Bonifacio et al. (2007) who found a scatter in Li abundances much larger than observational errors and anticorrelated with Na among (only) four stars in 47 Tuc. Very recently, Lind et al. (2009), targeting about 100 main-sequence (MS) and early subgiant branch stars in NGC~6397, detected for the first time a significant anticorrelation between Li and Na in this GC. This last investigation supersedes an older research by Bonifacio et al. (2002) based on only a few stars, and highlighting the importance of large samples of stars for similar studies. In this Letter we present Na, O, and Li abundance determinations in $\\sim$90 unevolved TO stars of the GC 47 Tucanae, providing the {\\it largest} database of this kind available in the literature so far. ", "conclusions": "\\label{sec:sum} We present in this Letter Li, Na, and O abundances for a large sample ($\\approx$90) of TO stars in the metal-rich GC 47 Tuc, providing the largest database of this kind available so far. Our main results can be briefly summarized as follows: \\begin{itemize} \\item[1.] We obtained a very clear Na$-$O anticorrelation, confirming the previous findings from giants derived by Carretta et al. (2009b). {\\it This is the first time that the Na$-$O distributions in dwarfs and giants can be directly compared for a cluster using large samples of comparable sizes for both evolutionary stages.} At least for 47 Tuc, evolutionary effects due to the RGB phase can be ruled out as contributors to the extent of Na$-$O anticorrelation, which in both cases span the same range (within the observational uncertainties). \\item[2.] As expected from stellar nucleosynthesis models in conjunction with multiple population scenarios, Li abundances should be positively correlated with O and anticorrelated with Na. At variance of the metal-poor GC NGC~6397, in the case of 47 Tuc the Li content does not show an anticorrelation with Na, and only a weak correlation appears with O, with a quite scatter distribution from both diagrams. Our result disagrees with the previous study by Bonifacio et al. (2007), who found Li$-$Na anticorrelation from a small sample of only four stars, and once again emphasizes the crucial role of statistics in this kind of analysis. A simple dilution model fails to reproduce the Li$-$Na$-$O distributions for this cluster and advocates the presence of some different mechanisms responsible for the observed Li pattern. \\item[3.] The scatter we find in Li abundances reminds of what has been detected, and reported in a large body of the literature, in Population~I stars of similar parameters (\\teff, $\\log{g}$), the most famous case being the old open cluster M~67 (see, e.g., Randich et al. 2000). \\end{itemize} We are not presently able to conclude if the trend we discovered in 47 Tuc is peculiar or, on the other hand, other GCs share a similar behavior. In fact, to date, only two GCs have been investigated from this point of view. In this context, we mention that we cannot explain the unlikeness in the Li$-$Na distributions between NGC~6397 and 47 Tuc; in particular, we cannot discriminate if the differences in \\teff's or metallicity can account for such a discrepancy. To probe this issue, it is crucial to enlarge the sample of simultaneous determinations of Li, Na, O in GC dwarf stars, by including other (nearby) clusters with different structural parameters (e.g., HB morphology, age, metallicity). \\begin{table*} \\begin{small} \\begin{center} \\caption{Properties of target stars.}\\label{t:tab1} \\setlength{\\tabcolsep}{1.3mm} \\begin{tabular}{lcccccccccccc} \\hline\\hline Star$_{\\rm ID}$ & R.A. & Decl. & $V$ & $B$ & T$_{\\rm eff}$ & $\\log~g$ & $\\log~n{\\rm (O)}$ & Err$_{\\rm O}$ & $\\log~n{\\rm (Na)}$ & Err$_{\\rm {Na}}$ & $\\log~n{\\rm (Li)}$ & Err$_{\\rm {Li}}$\\\\ & (deg) & (deg) & (mag) & (mag) & (K) & (dex) & (dex) &(dex) &(dex) &(dex) &(dex) & \\\\ \\hline & & & & & & & & & & & & \\\\ 001 & 6.15846 & $-$71.9632 & 17.347 & 17.913 & 5685 & 4.05 & 8.266 &\t0.050 & 5.715 & 0.103 & 2.105 & 0.053\\\\ 002 & 5.96746 & $-$71.9607 & 17.301 & 17.853 & 5676 & 4.02 & 8.061 &\t0.083 & 5.741 & 0.080 & 1.538 & 0.057\\\\ 003 & 6.33533 & $-$71.9429 & 17.347 & 17.913 & 5685 & 4.05 & 8.445 &\t0.050 & 9.999 & 9.999 & 1.998 & 0.130\\\\ 005 & 6.24487 & $-$71.9313 & 17.330 & 17.894 & 5681 & 4.04 & 8.112 &\t0.111 & 5.752 & 0.080 & 2.261 & 0.096\\\\ 006 & 6.16508 & $-$71.9222 & 17.372 & 17.941 & 5689 & 4.06 & 8.374 &\t0.050 & 5.689 & 0.080 & 2.453 & 0.108\\\\ 007 & 6.21133 & $-$71.9159 & 17.365 & 17.931 & 5688 & 4.05 & 8.493 &\t0.050 & 9.999 & 9.999 & 2.381 & 0.087\\\\ 008 & 6.27892 & $-$71.9146 & 17.377 & 17.940 & 5690 & 4.06 & 8.705 &\t0.050 & 9.999 & 9.999 & 2.240 & 0.049\\\\ 009 & 6.24204 & $-$71.9106 & 17.360 & 17.926 & 5687 & 4.05 & 8.327 &\t0.050 & 9.999 & 9.999 & 2.299 & 0.081\\\\ 010 & 6.11050 & $-$71.9085 & 17.335 & 17.899 & 5682 & 4.04 & 8.514 &\t0.050 & 9.999 & 9.999 & 2.294 & 0.112\\\\ 011 & 5.80258 & $-$71.9629 & 17.336 & 17.909 & 5682 & 4.04 & 9.999 &\t9.999 & 5.910 & 0.083 & 1.987 & 0.129\\\\ \\hline\\hline \\end{tabular} \\end{center} \\end{small} \\tablecomments{We report our ID in Column 1, R.A. and Decl. in Columns 2 and 3, while magnitudes $V$ and $B$ are given in Columns 4 and 5. Stellar parameters (T$_{\\rm eff}$ and $\\log g$) along with abundances and their errors are listed from Columns 6 to 13, respectively. (This table is available in its entirety in a machine-readable form in the online journal. A portion is shown here for guidance regarding its form and content.)} \\end{table*}" }, "1003/1003.5308_arXiv.txt": { "abstract": "{So far, GJ 710 is the only known star supposed to pass through outskirts of the solar system within 1 ly. We have reexamined the SIMBAD database for additional stellar candidates (from highest ratios of squared parallax to total proper motion) and compared them with new HIP2 parallaxes and known radial velocities. At the moment, the best nominee is double star HD 107914 in the constellation Centaurus at $\\approx 78.3$ pc from the Sun whose principal component is a white (A-type) giant. It does not seem to appear neither in general catalogues of radial velocities available at SIMBAD nor in authoritative Garc\\'{\\i}a-S\\'anchez et al. papers on stellar encounters with the solar system. Awaiting for the value $v_r$ of its radial velocity, unknown to the author, we have calculated limits of $|v_r|$ necessary to this star to pass within 1 ly and 1 pc from the Sun in linear approximation. A very accurate value of its total proper motion is also extremely important. In the case of $v_r=-100$ km/s and most ``advantageous'' HIP2 data, HD 107914 could pass as near as 8380 AU from the Sun in an almost direct collision course with the inner part of the solar system! Inversely, if $v_r$ had a great positive value, then HIP 60503 could be the creator of peculiar trajectories of detached trans-Neptunian objects like Sedna.} ", "introduction": "Many authors (\\emph{cf.}~\\cite{B,D,Ga1,Ga2}) have searched past and future stellar perturbers of the Oort cloud. Indeed, it turns out that a massive star passing within 1 ly would have a significant influence on long-period comets. More closely, at 10000 AU, such a star would have a serious direct influence on trans-Neptunian objects. \\smallskip Recently, Bobylev \\cite{B} has updated the list of stars supposed to transit within 2 pc from the Sun using new HIP2 parallaxes. HD 107914 is beyond the current scope of 30 pc in Bobylev's studies. This star has two currently known components CCDM J12242-3855AB of visual magnitudes 7.0 and 12.8 resp. CCDM attributes the spectral type A5 to the A-component while SIMBAD considers the double star as an A7/8 giant. [For the sake of comparison, A5-giant NSV 8327 at $\\approx$ 94.9 pc has the A-component of spectral type A2 with $V_{\\textrm{mag}}=6.0$ and B-component with $V_{\\textrm{mag}}=14.5$ according to CCDM.] Two components are separated by 4.6 arcsec which gives, at least, 350 AU at their current distance. The total mass of the system should be greater than $2M_{\\odot}$. ", "conclusions": "Nearby stars with very small proper motions are the best targets in the search of potential Nemeses. Our example shows that very accurate measurements of proper motions are indispensable for such stars. Then one can easily create a more or less full list of Nemesis candidates only from parallaxes and proper motions and calculate limit radial velocities. Finally, eliminating stars with small radial velocities, one can use elaborated models of the galactic potential (\\emph{cf.}~{\\cite{Ga2}}) to calculate stellar trajectories and minimal distances from target stars to the Sun. \\nocite{*}" }, "1003/1003.0875.txt": { "abstract": "{}{}{}{}{} % 5 {} token are mandatory \\abstract {Very few examples of luminous blue variable (LBV) stars or LBV candidates (LBVc) are known, particularly at metallicities below the SMC. The LBV phase is crucial for the evolution of massive stars, and its behavior with metallicity is poorly known. V39 in IC 1613 is a well-known photometric variable, with B-band changes larger than 1mag. over its period. The star, previously proposed to be a projection of a Galactic W Virginis and an IC 1613 red supergiant, shows features that render it a possible LBVc.} {We aim to explore the nature of V39 and estimate its physical parameters.} {We investigate mid-resolution blue and red VLT-VIMOS spectra of V39, covering a time span of 40 days, and perform a quantitative analysis of the combined spectrum using the model atmosphere code CMFGEN.} {We identify strong Balmer and \\ion{Fe}{ii} P-Cygni profiles, and a hybrid spectrum resembling a B-A supergiant in the blue and a G-star in the red. No significant Vrad variations are detected, and the spectral changes are small over the photometric period. Our analysis places V39 in the low-luminosity part of the LBV and LBVc region, but it is also consistent with a sgB[e] star. From this analysis and the data in the literature we find evidence that the [$\\alpha$/Fe] ratio in IC 1613 is slightly lower than solar.} {The radial velocity indicates that V39 belongs to IC 1613. The lack of Vrad changes and spectroscopic variations excludes binary scenarios. The features observed are not consistent with a W Virginis star, and this possibility is also discarded. We propose that the star is a B-A LBVc or sgB[e] star surrounded by a thick disk precessing around it. If confirmed, V39 would be the lowest metallicity resolved LBV candidate known to date. Alternatively, it could represent a new transient phase of massive star evolution, an LBV impostor.} ", "introduction": "Luminous blue variables (LBVs) constitute a short and rare phase in the life of massive stars, in which the stars suffer great mass loss with unpredictable outbursts and mass eruptions. Apart from the stages in which a nuclear runaway is produced in the stellar core, LBVs seem to be the shortest phase in the massive star's evolution, even shorter than Wolf-Rayet phases, with estimated lifetimes of 10$^4$-10$^5$ years. In spite of its brevity, this phase is extremely important for the evolution of the star and its surroundings, because it returns important amounts of processed material to the ISM, often forming a shell or a nebula around the star. % Danny recommends to delete it, why?: , and is later illuminated by the star. Moreover, the loss of matter modifies the star's subsequent evolution and it has been very recently suggested that LBVs might be {\\it direct} SNe progenitors (see \\cite{kotak06}, \\cite{vink09}), which may have an impact on our current ideas of massive star and galaxy evolution in the Universe. It is a tautology to state that LBVs are luminous (log(L/L$_\\odot$ $\\geq$ 5.4), blue ((B-V)$_0$ $\\leq$ 0.1) and variable stars. Nevertheless, their variations can be considerably different in frequency and intensity, going from intense and sudden eruptions in which the stellar magnitude increases by two magnitudes or more (like $\\eta$ Car) to cyclic variations of 1-2 magnitudes in time scales of years to decades (like the prototype S-Dor). These photometric changes are accompanied by the corresponding spectroscopic ones, in which the effective temperature varies and the bolometric luminosity may or may not remain constant \\citep{vink09}. However, outbursts and eruptions are not very common, and sometimes the star remains stable for centuries, as it is the case for P-Cygni. Thus, rather than from photometric variations, they are often searched for by means of their P-Cygni profiles in the Balmer lines (which are a signature of the presence of strong mass-loss processes), emission line profiles (at lower resolutions or NIR wavelengths) and of other lines which depend on the temperature of the star. When the star shows a characteristic spectrum with P-Cygni profiles, but no relatively large photometric and spectroscopic variations (or we do not have information about such variations), we speak of an LBV candidate, or LBVc. Sometimes, other signatures (like IR excess or nebulosity, see f.e. \\cite{clark05}) are considered further signs for the classification as LBVc. The mechanism producing LBVs is not known, neither is the exact duration of the phase (and hence, the exact impact on the star's life). It is assumed that the stellar mass and, very especially, the mass-loss rate play a role. As the winds of massive stars are driven by the momentum gained by the atoms from the radiation field through the absorption of photons by metal lines, LBVs should be rare at low metallicities. This is consistent with no LBVs or LBV candidates known in the Local Group at metallicities below that of the Small Magellanic Cloud (SMC). While this fact could be at least partly the result of the low number of massive stars in such galaxies, the sheer presence of an LBV at low Z should be considered significant, as their presence offers clues for mass-loss processes that could have been relevant in the early Universe. Beyond the Local Group few relevant examples of LBVs or possible LBVs are known. \\citet{drissen01} (see also \\citet{petit06}) have analyzed the star V1 in NGC 2363, a galaxy at about 3.3 Mpc (\\citet{tolstoy95}, \\citet{kara02}) with 12+log(O/H)= 7.9 \\citep{gonzalez94} or Z= 0.25 Z$_\\odot$ \\citep{luridiana99}. HST spectra were used for the analysis to avoid contamination because the star is in a dense stellar environment with strong nebulosity. NGC 2363-V1 shows strong photometric and spectroscopic variations \\citep{petit06}, thus fulfilling all requirements to be called an LBV. Using the CMFGEN code \\citep{hillier98}, \\citet{drissen01} obtained T$_{\\rm eff}$= 11000 -- 13000 K, log(L/L$_\\odot$)= 6.35 -- 6-40 and v$_\\infty$= 325 -- 290 km s$^{-1}$ for spectra collected in Nov. 1997 and July 1999, respectively, fully in agreement with the values expected for an LBV in transition. From the iron spectrum they get Z$\\approx$ 0.2 Z$_\\odot$, intermediate between \\citet{gonzalez94} and \\citet{luridiana99}. Very interesting, but less evident is the case of two very recent objects studied in very metal poor galaxies: knot 3 in DDO 68 (\\citet{pustilnik08}; \\citet{izotov09}) and PHL 293B \\citep{izotov09}. DDO 68 is a blue compact dwarf (BCD) galaxy at a distance of at least 5.9 Mpc (\\citet{makarova98} determined its distance from the average magnitude of the three brightest blue stars). %It has a very low metallicity: \\citet{pustilnik05} give 12+log(O/H)= 7.21 and \\citet{izotov07} give 7.14$\\pm$0.03. Knot 3 (DDO 68-3) is one of its strong HII regions, for which \\citet{pustilnik08} give 12+log(O/H)= 7.10$\\pm$0.06 and \\citet{izotov09} 7.15$\\pm$0.04. In spectra of this knot separated by about three years, \\citet{pustilnik08} detected spectral variations (broader emission in the Balmer lines and \\ion{He}{i} emission not previously present) at resolutions of 6 and 12 \\AA~ FWHM that they attributed to the outburst of an LBV inside the region. This explanation was adopted and further investigated by \\citet{izotov09}, who estimated the absolute magnitude, H$_\\alpha$ luminosity and wind terminal velocity from APO and MMT spectra at resolutions of 7 and 3 \\AA~ FWHM. However, note that the 1.5\" slit width used by \\citet{izotov09} corresponds to about 45 pc at the distance of DDO 68 and therefore the observed spectrum is a composite of the HII region and the underlying stellar population (including the possible LBV). A similar problem is found in the second object presented by \\citet{izotov09}, PHL 293B. This is another low-metallicity BCD (12+log(O/H)= 7.66$\\pm$0.04 according to \\citet{izotov07} and 7.72$\\pm$0.01 according to \\cite{izotov09}) but this time located at 21.4 Mpc (NED database). The spectra of \\cite{izotov09} were obtained with a slit width of 1.0\", which corresponds to more than 100 pc. The oxygen abundance quoted by \\cite{izotov09} has been obtained from a high resolution (0.2 \\AA~ FWHM) VLT-UVES spectrum that also displays broad emissions at the positions of H$_\\alpha$, H$_\\beta$ and H$_\\gamma$ and blue-shifted absorptions in all Balmer lines. No He broad emissions are seen. Neither the DDO 68-3 nor the PHL 293B spectra show metal lines that could be attributed to a star. In both cases, the wind terminal velocities quoted by \\cite{izotov09} ($\\approx$ 700-850 km s$^{-1}$) are significantly larger than the usual range of LBVs ( 100-250 km$^{-1}$; 500 km s$^{-1}$ for $\\eta$ Car), which is attributed by the authors to the very low metallicity. The mere presence of LBVs at these low metallicities represents a challenge for the theory, as there are few metals to drive the wind. Other processes have been suggested to cause the LBV outburst, like the metallicity-independent continuum-driven (instead of line-driven) wind proposed by \\citet{owocki04} (see also \\cite{smith06}). Note however that the star needs to be close to the Eddington limit for the mechanism to operate. There is a clear interest in investigating these processes, confront them with observations and apply the results to early Universe and Population III objects. Obviously, it would be very interesting to obtain spatially resolved spectra of the stars in DDO 68-3 and PHL 293B, both to confirm that the changes in the composite spectra are due to an LBV and to study the evolutionary processes at very low metallicity. However, because of the large distances involved, we have to look closer for resolved stars, such as NGC 2363-V1. IC\\,1613 is a dwarf irregular galaxy in the Local Group with a distance modulus of (m-M)$_0$= 24.27 \\citep{dolphin01}. It has a metallicity of log(O/H)+12= 7.80$\\pm$0.10 as determined from its B-supergiants \\citep{bresolin07}, whereas nebular studies vary between 7.60 and 7.90 \\citep{lee03}. It is therefore clearly below the SMC metallicity and intermediate between that of NGC 2363 and PHL 293B. Therefore, the observation of a resolved LBV in such a poor-metal galaxy would constitute a step forward, adding a second object to the list of LBVs at metallicities below that of the SMC and lowering the present metallicity limit for resolved LBVs. IC\\,1613 shows a recent and intense burst of massive star formation, particularly in its NE part. We have recently published a new catalog of OB associations in IC\\,1613 \\citep{garcia09} and their physical properties (Garcia et al., in prep.) as part of our effort to carry out an in-depth study of the young population of IC\\,1613. We also obtained spectra of some stars in this galaxy. In the field of IC\\,1613 we find the well-known variable star V39. In Table~\\ref{stardata} we offer an overview of its photometric data. \\begin{table*}[htdp] \\caption{Photometric data for V39. \\citet{udalski01}, \\citet{pietr06} and \\citet{garcia09} give mean values.} \\centering \\begin{tabular}{cccccccl} \\hline \\hline U & B & V & R & I & J & K & Ref \\\\ \\hline & 18.6 - 19.9 & & & & & & \\cite{sandage71}\\\\ & & 18.851 & & 17.617 & & & \\citet{udalski01}\\\\ & & 18.5 - 19.0 & & 17.45 - 17.7 & & & \\citet{mantegazza02}\\\\ & & & & & 16.815 & 15.818 & \\citet{pietr06}\\\\ 19.43 & 19.62 & 19.00 & 18.37 & 17.75 & & & \\citet{garcia09}\\\\ \\hline \\end{tabular} \\label{stardata} \\end{table*}% The nature of V39 has remained unclear since its discovery. \\cite{sandage71}, in his analysis of IC1613 based on unpublished previous work by Baade, points out that V39 is the only peculiar variable in the field, because of its inverted $\\beta$-Lyrae light curve. This was already noticed by Baade, who never considered it to be a Cepheid, not only because of this peculiar light curve, but also because it was too bright to fit the P-L relation compared to other Cepheids with the same period. Consequently, in spite of a suggestion by \\cite{sandage71} (that implied a new P-L relation and a change in the distance modulus to bring it into agreement with other Cepheids) most authors did not consider it a genuine Cepheid. \\citet{udalski01} used its position in the I vs (V-I) color-magnitude diagram as a new argument to discard it as a Cepheid. Very recently, \\citet{pietr06} have also shown that V39 is too bright to fit the Cepheids P-L relation in the infrared. \\citet{antonello99} showed that the combined light curve from different observers is the result of the superpositon of two periods: a long period of 1123 days and a short period of 28.699 days. They showed that by subtracting the long period from the light curve, the remaining one did not show a clear difference between the primary and secondary maxima, so that the short period could be halved. Based on these results, \\cite{mantegazza02} proposed V39 to be actually the casual overlapping of a distant Galactic W Vir star and a red supergiant belonging to IC\\,1613. If confirmed, the W Vir star would be at a distance of at least 115 Kpc, being the most distant Galactic star from the Galactic plane, and raising the question about the true extent of the Galactic halo. In this article we show that V39 does actually belong to IC\\,1613 and analyze its possible nature, including the possibility that it is an LBV candidate, the first one known in IC\\,1613. If confirmed, it would be the (resolved) LBV candidate within the lowest metallicity environment. The observations are presented in Sect.~\\ref{obs}, a description of the spectrum is offered in Sect.~\\ref{spec} and the nature of V39 is then discussed in Sect~\\ref{nature}. An estimation of the stellar parameters is offered in Sect.~\\ref{param} and the evolutionary status and other properties are discussed in Sect.~\\ref{discus}. Finally, the conclusions are presented in Sect.~\\ref{conc}. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ", "conclusions": "\\label{conc} We have presented intermediate resolution (R=2050-2150) spectroscopic observations of the variable star V39 in IC\\,1613 secured with VIMOS at the VLT. The final spectrum is the combination of 19 blue spectra and 10 red spectra and is characterized by i) strong P-Cygni profiles in the Balmer lines and \\ion{Fe}{ii} lines; ii) an hybrid character, more consistent with an early type supergiant (B-A) in the blue and a late type (G) in the red. All spectral features are consistent in radial velocity with the star belonging to IC 1613. The individual spectra do not show significant radial velocity variations in spite of the time span covered, which is longer than one photometric period (during which photometric changes are larger than one magnitude in B and decrease with wavelength), and show only moderate spectral line variations. Therefore we have discarded scenarios including a binary system, specially those demanding large radial velocity variations. We also discarded the scenario proposed by \\cite{mantegazza02} (a Galactic W Virginis star plus an IC 1613 red supergiant) because it is not consistent with the system belonging to IC1613 and the observations do not show the spectral characteristics and changes of a W Virginis star. We propose a scenario consisting of an early (B-A) type star surrounded by a hot ($\\sim$ 5000 K) thick disk precessing around it. Although this scenario can potentially explain the photometric changes and its wavelength dependence as well as the lack of significant spectroscopic changes, details have to be worked out (for example, the cause of disk precession) and independent evidence (like for example the presence of double peaks in spectral lines) has still to be obtained. The analysis of the spectrum under the assumption that this is a single star indicates a low-metallicity (Fe $\\sim$ 0.2 Fe$_\\odot$, from the iron P-Cygni lines), hot evolved star with a strong wind. The agreement of this result with other spectroscopic Fe or isochrone fitting global metallicities (\\citet{taut07}, \\citet{skillman03}) on the one hand and the stellar and nebular O abundances on the other (\\citet{bresolin07}, \\citet{lee03}) strongly points to a low [$\\alpha$/Fe] ratio in IC 1613 as compared to solar values, although more work will be needed to confirm this possibility. From its position in the HRD the star lies in the low-luminosity edge of LBV and LBV candidates and its progenitor may have had a mass of about 25 M$_\\odot$ (which would be close to the present mass). The derived mass-loss rate is higher than expected for a star of the given luminosity and metallicity, and it may represent a challenge to theory. The terminal velocity is consistent with other LBVc, and we could not confirm the suggestion by \\citet{izotov09} that it increases at low metallicities. Our result agrees with that of \\citet{drissen01} for NGC 2363-V1. If confirmed, V39 would be the resolved LBV candidate located in the most metal poor environment known to date. Alternatively, the star could be a high-luminosity version of a sgB[e]. This second alternative would imply some significant dust formation around V39 (which we could neither confirm nor reject) and is not inconsistent with the star being an LBV at some stage in its evolution. The combination of the low abundances with the high mass-loss rate raises the issue of how the winds of evolved stars behave at low metallicities (like those of LBV and sgB[e] stars) and whether they are driven by line radiation, pulsations or continuum radiation playing a stronger role than in Galactic stars. Our data are not sufficient to explore this question in detail. Higher resolution spectra, allowing us to dissentangle the contributions of the star from that of possible circumstellar material, disk or faint companion are needed. Whether V39 is an LBVc or a new transient phase of massive star evolution, an LBV impostor, is an important question. If the former, then it will hold the record as the lowest metallicity LBVc known which has important implications for stellar evolution (as described in the introduction). If the latter, then understanding what this phase might represent has great intrinsic interest, in addition to raising the issue of whether or not there may be other LBV impostors which are undiscovered due to a lack of adequate data, like a light curve. Unraveling its nature is therefore of great importance. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%" }, "1003/1003.0007_arXiv.txt": { "abstract": "We present measurements of the dark matter bispectrum in N-body simulations with non-Gaussian initial conditions of the local kind for a large variety of triangular configurations and compare them with predictions from Eulerian Perturbation Theory up to one-loop corrections. We find that the effects of primordial non-Gaussianity at large scales, when compared to Perturbation Theory, are well described by the initial component of the matter bispectrum, linearly extrapolated at the redshift of interest. In addition, we find that, for $\\fNL=100$, the nonlinear corrections {\\em due} to non-Gaussian initial conditions are of the order of $\\sim 3$-$4$\\% for generic triangles up to $\\sim 20$\\% for squeezed configurations, {\\em at any redshift}. We show that the predictions of Perturbation Theory at tree-level fail to describe the simulation results at redshift $z=0$ at scales corresponding to $k\\sim 0.02$-$0.08\\kMpc$, depending on the triangle, while one-loop corrections can significantly extend their validity to smaller scales. At higher redshift, one-loop Perturbation Theory provides indeed quite accurate predictions, particularly with respect to the {\\em relative} correction due to primordial non-Gaussianity. ", "introduction": "In recent years, a significant research activity has been devoted to the effects of a possible small departure from Gaussianity in the primordial cosmological perturbations. While current constraints on primordial non-Gaussianity from measurements of the cosmic microwave background (CMB) and the large-scale structure are still consistent with the Gaussian hypothesis \\citep{KomatsuEtal2010, SlosarEtal2008}, a possible detection in forthcoming experiments would constitute a major discovery, providing crucial information on the early Universe and on the high-energy physics of inflation \\citep[see, for instance,][]{KomatsuEtal2009A}. The effect of primordial non-Gaussianity on the large-scale structure has been assumed, for a long time, to be limited to an additional, {\\em primordial} component to the matter skewness and bispectrum induced by gravitational instability and to a correction to the abundance of massive cluster \\citep[see][for recent reviews]{LiguoriEtal2010, DesjacquesSeljak2010B}. Numerical and analytical studies have indeed shown that a matter density probability distribution initially skewed toward positive values produces more overdense regions and, consequently, collapsed objects while a negatively skewed distribution produces larger voids \\citep[see][for recent work]{GrossiEtal2008, PillepichPorcianiHahn2010, KamionkowskiVerdeJimenez2009, LamSheth2009, LamShethDesjacques2009, MaggioreRiotto2009C}. Moreover, a nonvanishing skewness in the initial conditions corresponds to a primordial component to the matter bispectrum, \\ie the three-point function in Fourier space. For the {\\em local} non-Gaussian model considered here, the primordial matter bispectrum exhibits a scale, redshift and triangle shape dependence distinct from that of the component sourced by the nonlinear growth of structures. This enables us in principle to disentangle the two contributions. In the specific case of equilateral triangular configurations, the primordial contribution to the matter bispectrum scales as $\\sim k^{-2}$ relative to the gravity-induced term, leading to large, potentially observable corrections at low wavenumbers. Measurements of the galaxy bispectrum in future large-volume redshift surveys (such as Euclid or HETDEX) should be able to provide constraints on the local non-Gaussian model competitive with those from CMB observations \\citep{ScoccimarroSefusattiZaldarriaga2004, SefusattiKomatsu2007, SefusattiEtal2009}. In addition to these effects, \\citet{DalalEtal2008} have recently discovered a large correction to the galaxy bias in numerical simulations of {\\em local} primordial non-Gaussianity. Further numerical and theoretical work has confirmed this result \\citep{MatarreseVerde2008, SlosarEtal2008, AfshordiTolley2008, McDonald2008, GrossiEtal2008, TaruyaKoyamaMatsubara2008, PillepichPorcianiHahn2010, DesjacquesSeljakIliev2009, GiannantonioPorciani2009}. The constraints obtained from power spectrum measurement of highly biased objects in current data-sets are already comparable to the CMB results \\citep{SlosarEtal2008, DesjacquesSeljak2010}, and the prospects for detecting local primordial non-Gaussianity with galaxy clustering look exciting \\citep{DalalEtal2008, CarboneVerdeMatarrese2008, Seljak2009, Slosar2009, VerdeMatarrese2009, DesjacquesSeljak2010}. At this point, analyses of the galaxy bispectrum preceeding the work of \\citet{DalalEtal2008} must be updated to account for the non-Gaussian correction to the galaxy bias. In fact, a rigorous {\\em joint} analysis of the galaxy power spectrum and bispectrum in presence of local non-Gaussianity is in order. First steps in this direction have been taken by \\citet{JeongKomatsu2009B} and \\citet{Sefusatti2009} with a preliminary comparison with simulations in \\citet{NishimichiEtal2009}. In this perspective, we will consider the measurement of several triangular configurations of the {\\em matter} bispectrum on mildly nonlinear scales, with both Gaussian and non-Gaussian initial conditions of the local type. Although the matter bispectrum is not directly observable with tracers of the large-scale structure, it is instructive to assess the extent to which perturbation theory describes the shape dependence of the matter three-point function in the presence of non-Gaussianity of the local type. This analysis will be useful when considering the complication brought by biasing, which will be addressed in a forthcoming publication. Measurements of the matter power spectrum with local non-Gaussianity can be found in \\citet{PillepichPorcianiHahn2010, DesjacquesSeljakIliev2009}, where the small corrections at mildly nonlinear scales predicted in the framework of perturbation theory by \\citet{TaruyaKoyamaMatsubara2008} are observed. In the case of the matter bispectrum, measurements in simulations with Gaussian initial conditions are shown in \\citet{ScoccimarroEtal1998, HouEtal2005, PanColesSzapudi2007, SmithShethScoccimarro2008, GuoJing2009A}, with \\citet{SmithShethScoccimarro2008} considering, in addition, redshift space predictions in the context of the halo model. By contrast, the only measurement so far of the matter (and halo) bispectrum in simulations with local non-Gaussian initial conditions can be found in \\citet{NishimichiEtal2009}, where a relatively small subset of isosceles triangular configurations is considered. We will compare our measurements with predictions of the matter bispectrum at the one-loop approximation in Eulerian perturbation theory. A comparison of one-loop results with the bispectrum extracted from simulations with Gaussian initial conditions is shown in \\citet{ScoccimarroEtal1998}, whereas a comparison of the effect of primordial non-Gaussianity with the tree-level prediction of perturbation theory is performed in \\citet{NishimichiEtal2009} for ``squeezed'' isosceles configurations at $z=0$ with $k\\lesssim 0.1\\kMpc$ only. Here, we will extend the analysis to include several triangular configurations covering the range of scales $0.002\\lesssim k \\lesssim 0.3\\kMpc$ and redshifts $z=0$, $1$ and $2$. This will allow us to broadly test the accuracy of one-loop perturbation theory in the mildly nonlinear regime. We will also discuss the validity of two phenomenological prescriptions for the nonlinear bispectrum with Gaussian initial conditions, namely the fitting function of \\citet{ScoccimarroCouchman2001} and the formula of \\citet{PanColesSzapudi2007} based on a scaling transformation. This paper is organized as follows. In section~\\ref{sec:theory} we summarize previous results on the predictions of the matter power spectrum and bispectrum in cosmological perturbation theory for both Gaussian and local non-Gaussian initial perturbations. In section~\\ref{sec:sims} we describe the N-body simulations and the bispectrum estimator employed in our analysis whereas, in section~\\ref{sec:results}, we present our measurements of the matter bispectrum and compare them to one-loop predictions in perturbation theory. Finally, we conclude in section~\\ref{sec:conclusions}. ", "conclusions": "\\label{sec:conclusions} A rather surprising effect of local primordial non-Gaussianity on the large scale clustering properties of biased objects has been observed in various numerical studies over the last years \\citep{DalalEtal2008, DesjacquesSeljakIliev2009, PillepichPorcianiHahn2010, GrossiEtal2009}. These results attracted a great deal of attention as they showed that measurements of the power spectrum of galaxies and quasars from current data sets can lead to constraints on the local non-Gaussian parameter $\\fNL$ comparable to those of CMB observations \\citep{SlosarEtal2008, DesjacquesSeljak2010B}. Previous work assumed that the main effect of primordial non-Gaussianity is limited to an extra contribution to the matter and galaxy bispectrum. Still, even under such incorrect but ``conservative'' assumption, it has been shown that future large-volume redshift surveys will reach a sensitivity to a non-zero $\\fNL$ comparable or better than the CMB bispectrum \\citep{ScoccimarroSefusattiZaldarriaga2004, SefusattiKomatsu2007}. The inclusion of the non-Gaussian bias in the analysis of the galaxy bispectrum or, better, in a combined analysis of the power spectrum and bispectrum, is desirable to reliably assess the potentiality of forthcoming surveys of the large scale structure. As a first step in this direction, we have measured the matter bispectrum for the main classes of triangle shape using a set of large-volume N-body simulations seeded with Gaussian and non-Gaussian initial conditions of the local type. We focused on mildly nonlinear scales, $0.02\\lesssim k\\lesssim 0.3\\kMpc$, presented a wide choice of triangular configurations of different shapes and obtained a determination of the bispectrum with an overall errors of the order of $3$-$4\\%$. Of particular interest in this range of scales are the nonlinear corrections induced by gravitational instability {\\em due} to non-Gaussian initial conditions as they generate an additional non-Gaussian signal on top of the primordial component. For a nonlinear parameter $\\fNL=100$, we found that the amplitude of these corrections range from $3$-$4\\%$ for generic triangle configurations up to $20$-$30\\%$ for ``squeezed'' configurations where we expect most of the signal for local non-Gaussianity. We quantified these corrections with the aid of the ratio and the difference between the non-Gaussian and the Gaussian bispectrum. Our set of eight different realizations of those models ensure that our results are robust to sampling variance. We considered simulations snapshots at redshift $z=0$, $1$ and $2$. Overall, we found that the magnitude of the correction induced by non-Gaussian effects is similar regardless the scale and the redshift. This is due to a compensation between the primordial component that decreases with time on the one hand, and the contribution from nonlinear structure growth that increases with time on the other hand. We compared our results with the predictions of Eulerian perturbation theory, both at tree-level and one-loop \\citep{Sefusatti2009}. As expected, and similarly to what happens for Gaussian initial conditions, the tree-level approximation fails at relatively large scales, $k\\sim0.05$ - $0.1\\kMpc$, even at high redshift. One-loop corrections extend significantly the predictive power of perturbation theory down to mildly non-linear scales $k\\sim 0.3\\kMpc$ at redshift $z\\gtrsim 1$, similarly to the case of the power spectrum analyzed by\\citet{JeongKomatsu2006}. They describe, in fact, the matter bispectrum measured in simulations at the few percent level, with an even better agreement with respect to the ``relative'' effect of primordial non-Gaussianity on the Gaussian bispectrum. Furthermore, they also show a good qualitative agreement with simulations at redshift zero." }, "1003/1003.2691_arXiv.txt": { "abstract": "We discuss the non-thermal leptogenesis in the scheme of 5D orbifold SO(10) GUT with the smooth hybrid inflation. With an unambiguously determined Dirac Yukawa couplings and an assumption for the neutrino mixing matrix of the tri-bimaximal from, we analyze baryon asymmetry of the universe via non-thermal leptogenesis in two typical cases for the light neutrino mass spectrum, the normal and inverted hierarchical cases. The resultant baryon asymmetry is obtained as a function of the lightest mass eigenvalue of the light neutrinos, and we find that a suitable amount of baryon asymmetry of the universe can be produced in the normal hierarchical case, while in the inverted hierarchical case the baryon asymmetry is too small to be consistent with the observation. ", "introduction": " ", "conclusions": "" }, "1003/1003.0970_arXiv.txt": { "abstract": "Using high resolution VLT spectra, we study the multi-component outflow systems of two quasars exhibiting intrinsic Fe II absorption (QSO~2359--1241 and SDSS~J0318--0600). From the extracted ionic column densities and using photoionization modeling we determine the gas density, total column density, and ionization parameter for several of the components. For each object the largest column density component is also the densest, and all other components have densities of roughly 1/4 of that of the main component. We demonstrate that all the absorbers lie roughly at the same distance from the source. Further, we calculate the total kinetic luminosities and mass outflow rates of all components and show that these quantities are dominated by the main absorption component. ", "introduction": "Rest-frame UV spectra of roughly 20\\% of all quasars exhibit blueshifted Broad Absorption Lines (BAL) that are indicative of an outflow. BAL are mainly associated with resonance lines of high ionization species, like \\ion{C}{4}, \\ion{N}{5}, \\ion{O}{6} (HiBAL), and can reach velocities as high as 50,000 km $s^{-1}$ \\citep{weymann, turnshek}. Despite various recent statistical studies of BAL quasars \\citep{hall, trump, ganguly} the relationship between HiBAL and the host galaxy remain illusive as these outflows do not contain distance diagnostics in their spectra. Thus, it is difficult to establish whether the outflows affect only the near AGN environment (0.1-10~pc) or whether they extend to the scales of the entire galaxy (1-10~kpc). A subset of BALQSOs also show absorption features from low ionization species such as \\ion{Mg}{2}, \\ion{Al}{2}, and most importantly for diagnostics, \\ion{Fe}{2} and \\ion{Si}{2}. These absorption features are generally complex, as they are made of multiple components of narrower (of the order of a few hundred km~s$^{-1}$) absorption troughs. AGN with this kind of features are often refer to as FeLoBAL. In these systems the spectra of \\ion{Fe}{2} and \\ion{Si}{2} is valuable because they often include absorption troughs from metastable levels, which serve to determine the distance of the outflow from the central source and thereby beginning to relate the outflows to their host galaxy. Such outflows were studied in the past \\citep{wampler, dekool01, dekool02, hamann, arav01}. However, the presently available Sloan Digital Sky Survey (SDSS) and the development of advanced BAL analysis tools (Arav et al 2002, Gabel 2005, Arav et al 2005) allow for a more accurate and systematic study of these systems. With this in mind, we began a comprehensive study of FeLoBAL outflows that contain distance diagnostics \\citep{arav08, korista08, dunn09, moe09}. Up to the present, though, all studies are based on either global properties of the multi-component absorption troughs or the strongest kinematic component. Little is known about the individual physical properties of the components and their relationship to the whole outflow. In the best studied cases, \\cite{wampler, dekool01, dekool02} were able to obtain some column densities for individual components, but could neither diagnose the physical conditions of the components nor derive their relative distances. Simple assumptions that all components had similar physical conditions and that they have either constant speed or equal monotonic acceleration would lead to the conclusion that either the absorbers are scattered along a wide range of distances from the central source or the absorbers vary in age as a result of episodic ejections. But BAL systems in general are far from simple. Apart from a weak correlations between the absorber bulk velocities and bolometric luminosity of the AGN \\citep{laor, dunn08, ganguly}, no correlations have been found among velocity, width, ionization, or any other observable properties. Recent high spectral resolution eschell observations of two quasars QSO 2359--1241 and SDSS J0318--0600 with the Very Large Telescope (VLT) have allowed us to study in detail the various independent components that compose their FeLoBAL. QSO 2359--1241 (VNSS J235953--124148) is a luminous ($M_B= -28.7$), reddened ($A_V \\approx 0.5$) quasar at redshift $z=0.868$ \\citep{brotherton01}. SDSS J0318--0600 is also a highly reddened bright quasar ($A_V\\approx 1$ and $M_B=-28.15$) at redshift $z=1.967$. Both quasars exhibit a rich multi-component FeLoBAL comprising five and eleven components respectively. The strongest components in the FeLoBAL, which are also the first absorbers from the central engine (see Section 3), for each of these objects were measured and analyzed in \\cite{arav08} and \\cite{korista08} for QSO 2359--1241 and \\cite{dunn09} for SDSS J0318--0600. It was found that the absorbers are located at $\\sim$3~kpc and between 6 and 20 kpc from the central engine, respectively. QSO 2359--1241 and SDSS J0318--0600 are the first to be studied in detail from a sample of $\\sim$80 FeLoBAL quasars with resonance and metastable \\ion{Fe}{2} absorption lines (FeLoBAL quasars) in their spectra. These lines can be used for direct determination of the physical conditions and energy transport of the outflows. The sample was extracted out of 50,000 objects in the SDSS database as part of a major ongoing effort to study the nature of quasar outflows and their effects on the host galaxy \\citep{moe09}. ", "conclusions": "The high spatial resolution and signal-to-noise of the VLT spectra allowed us to study the properties of each of the kinematic components in the FeLoBALS of quasars QSO 2350--1241 and SDSS J0318--0600. From the measurements of column densities for different kinematic components we determined the electron number density for these components. For QSO 2350--1241 we used the ratio of column densities of \\ion{Fe}{2} from the excited level at 385 cm$^{-1}$ and the ground level. In the case of SDSS J0318--0600 we used the ratios of column densities among levels of the ground multiplets of \\ion{Si}{2} and \\ion{C}{2}. Interestingly, there is a clear density contrast between the maint kinematic component in each object and the smaller components. By contrast, all smaller components in each object seems to exhibit roughly the same density. The density contrast between the densest components and the smaller one in each object are $\\sim 0.8$ dex for QSO 2350--1241 and $\\sim 0.5$ dex for SDSS J0318--0600. Next, we determine the ionization parameters characteristic of each of the absorption components in both quasars. To this end, we designed diagnostic plots by which the ionization parameter as well as the total hydrogen column can be uniquely determined. These plots demonstrate that: (1) any given ratio of column densities among medium and low ionization is a smooth function of the column density for a fixed value of the ionization parameter and (2) these curves of column density rations vs. column density are monotonic with $U_H$. There are various consequences of this: (a) for a given pair of measured column densities (and fixed density, chemical composition, and SED) no more than one solution in $U_H$ and $N_H$ can be found; (b) $U_H$ and $N_H$ and their errors are necessarily correlated; (c) these solutions can be found either graphically or numerically in a way much more efficient than though generic numerical optimization techniques; (d) the graphical nature of the diagnostic allows one to set the observations from various different observers on the same page and gain valuable insights. In the determination of $\\vy{U}{H}$ and $\\vy{N}{H}$ traditional plots of predicted column densities vs. $\\vy{N}{H}$ as abscissa have various disadvantages. This is because for every column density measured there is whole family of solutions ($\\vy{U}{H}$, $\\vy{N}{H}$), thus both parameters must be determined simultaneously. A typical approach used is to normalize the predicted column densities to the measured columns and then look for the intersection between various curves. The disadvantages in that is that in dividing theoretical values by measured values mixes up theoretical and observational uncertainties. Further, a solution based on intersections of curves or broad regions, if uncertainties are accounted for in some fashion, offers no simple intuitive understanding of how the results may change in case of systematic effects on either the theoretical modeling or the observations. By contrast, the plots that we propose here clearly allow one to visualize the error bars of the measurements and their significance relative to the predictions of different models. One can also visualize how systematic effects on the calculations, like for example chemical composition or shape of the SED, would shift the results of the diagnostics. Another very important advantage of the proposed plots is the potential to put various kinematic components of the same trough on the same plot and compare them under equal conditions. It is true that the proposed plots do not show the resulting $\\vy{N}{H}$ explicitly, but by fixing $\\vy{U}{H}$ the whole problem is solved and there is a direct correspondence between the observed column density of observed species and $\\vy{N}{H}$. Thus, $\\vy{N}{H}$ can be read directly from the tabulated solutions of the models. We determined relative distances of the various kinematic components, firstly under the assumption that all components see the same unattenuated SED. It becomes immediately clear that the component with the largest column density is always the first in line from the source. Once the first kinematic component in line was identified for each object we include the effect of attenuation of the SED on the distance determination for the remaining components. It was found that distance determinations that ignore attenuation affects are significantly overestimated. By contrast, when attenuation by the innermost component is considered in the distance estimation all the kinematic components in the absorption troughs are found in close proximity to each other, and possibly related. This result, if found generally true in most FeLoBAL, ought to have important consequences in our understanding of the dynamics and energetics of quasar outflows." }, "1003/1003.2144_arXiv.txt": { "abstract": "The Horava\\,-\\,Lifshitz (HL) theory has recently attracted a lot of interest as a viable solution to some quantum gravity related problems and the presence of an effective cosmological constant able to drive the cosmic speed up. We show here that, in the weak field limit, the HL proposal leads to a modification of the gravitational potential because of two additive terms (scaling respectively as $r^2$ and $r^{-4}$) to the Newtonian $1/r$ potential. We then derive a general expression to compute the rotation curve of an extended system under the assumption that the mass density only depends on the cylindrical coordinates $(R, z)$ showing that the HL modification induces a dependence of the circular velocity on the mass function which is a new feature of the theory. As a first exploratory analysis, we then try fitting the Milky Way rotation curve using its visible components only in order to see whether the HL modified potential can be an alternative to the dark matter framework. This turns out not to be the case so that we argue that dark matter is still needed, but the amount of dark matter and the dark halo density profile have to be revised according to the new HL potential. ", "introduction": "Inspired by the Lifshitz theory in condensed matter physics, Horava (2009a,b) has recently proposed a new theory of gravity based on an anisotropic scaling of space and time in the UV limit. Usually referred to as Horava\\,-\\,Lifshitz (hereafter, HL) theory, the HL proposal shows a reduced invariance, dubbed {\\it foliation preserving diffeomorphism invariance}, which however reduces to the standard one in the IR limit where General Relativity is recovered. As an attractive feature, the HL theory turns out to be power counting renormalizable which has motivated the great interest in investigating with great detail its theoretical and cosmological aspects. It is worth remembering that, in its original formulation, two conditions were imposed in order to drive the choice of the field action. First, the {\\it projectability condition} was supposed to hold true. Since time plays a fundamental role from the very beginning, it was assumed that the lapse function (defined when working in the ADM formulation of gravity) has to be a projectable function on the spacetime foliation, that is to say a function of time only. Second, in order to reduce the number of independent terms entering the action, the principle of {\\it detailed balance} was used. Unfortunately, it soon became clear that this second condition leads to problems in the low energy limit \\cite{LMP09} thus motivating the search for modification of the original HL theory where the action breaks the detailed balance condition either softly \\cite{KS09,LKM09,CP09,KK09a,KK09b} or not. In particular, Sotiriou, Visser and Weinfurtner (2009a,b) have worked out a modified HL theory with no detailed balance condition imposing that only parity preserving operators enter the potential. Although the discussion about the foundations and the possible conceptual and phenomenological problems of the HL theory and its modified versions is still open, it is nevertheless worth systematically investigating its consequences at every scale. In particular, it is interesting to study its static spherically symmetric solutions since one can thus derive the gravitational potential generated from a point mass source. Recently, this problem has been addressed by Tang \\& Chen (2009) for the HL theory with the projectability condition and no detailed balance. They argue that only the Minkowski or de Sitter spacetime are solutions, but we will show here that this is actually not the case. As a consequence, we find that the gravitational potential generated by a point mass differs from the Newtonian one because of the presence of additional terms depending on the HL coupling parameters. Such terms have to be taken into account when computing the potential generated by an extended system, such as a galaxy. Therefore, we work out a general formalism to estimate the rotation curve (i.e., the circular velocity $v_c$ as function of the distance $R$ from the centre) for an extended source showing that the HL theory can boost the $v_c(R)$ with respect to the Newtonian value. Motivated by this consideration, we then try to fit the Milky Way rotation curve using visible matter only to both constrain the HL parameters and investigating whether it can work as an effective dark matter component. The plan of the paper is as follows. In Sect.\\,2, we look for static spherically symmetric solutions for the HL theory with projecatibility condition and show that there is indeed a new solution leading to a modified gravitational potential. A general formalism to compute the rotation curve for an extended system is presented in Sect.\\,3 and then used in Sect.\\,4 to work out the predicted rotation curve for the Milky Way. Here, we also present the data and the results of fitting them with our modified potential and no dark matter. Conclusions are finally given in Sect.\\,5. ", "conclusions": "Initially motivated by its attractive features from the point of view of quantum gravity, the HL proposal has soon become one of the most investigated theories of gravity. We have here complemented recent works on its cosmological consequences by addressing its impact on the gravitational potential. Contrary to the claim in TC09, we have demonstrated that static spherically symmetric solutions other than the Schwartzschild\\,-\\,de Sitter one exist. As a consequence, we have found a modified gravitational potential made out of the Newtonian $1/r$ one corrected by the addition of three further terms scaling as $r^2$, which corresponds to the effect of a cosmological constant, and a quickly decreasing term, proportional to $1/r^4$. The importance of these two terms is parametrized in terms of two conveniently defined scaling radii, namely $(r_A, r_D)$, related to the couplings entering the HL Lagrangian. In order to consider astrophysically interesting situations, we have then developed a general formalism to compute the circular velocity curve provided the mass density and the mass function of the system are given. As an application, we have then evaluated the Milky Way rotation curve using as only source of the gravitational field a spheroidal truncated power\\,-\\,law bulge and a double exponential disc. It turns out that the modified rotation curve is unable to fit the data thus demonstrating that the HL theory can not play the role of an alterative solution to the missing mass problem. It is worth noting that Mukohyama (2009) has shown that classical solutions to the infrared (IR) limit of the HL theory can mimic General Relativity plus cold dark matter so that one can argue that it should be possible for the HL theory to provide an effective dark matter halo in apparent contradiction with our finding. Actually, this is not the case. First, Mukohyama only refers to the HL IR limit so that the Lagrangian he has considered does not include any potential term, while we here explicitly include this term through the coupling parameters $g_i$. Second, the cold dark matter term comes out as a consequence of the global Hamiltonian constraint being less restrictive than the local one typically imposed in General Relativity. In our static spherically symmetric metric, the global Hamiltonian constraint reduces to Eq.(\\ref{eq: ss2}). In order to find our solution, we then impose that the integrand in Eq.(\\ref{eq: ss1}) identically vanishes so that we are actually converting the global constraint in a local one thus going back to a situation similar to the General Relativity case. As a consequence, we argue that the possibility to get a matter term as an integration constant is lost in our approach. The next step in our analysis of the HL theory on galactic scale should naturally be the inclusion of a dark halo. There are however some subtle issues making such a logical step forward not so easy to address. First, needless to say, there are few hints on which the dark halo mass density profile should be. All the models used in literature are motivated by the outcome of numerical simulations, such as, e.g., the popular NFW \\cite{NFW97} and Einasto \\cite{E65,Polls}, or evidences from rotation curve fitting, such as the isothermal sphere \\cite{BT87} and the Burkert \\cite{B95,BS00,BS01} profile. All the previous works implicitly assume the validity of the Newtonian potential so that they are no more valid in a modified framework as the one we are using here. As a consequence, we have therefore to explore a wide class of density profile able to mimic most of the models in literature to finally select the most empirically motivated one. Moreover, as Eq.(\\ref{eq: rotcurvegen}) shows, we also need to know the dark matter mass function which is completely unknown. Should the dark matter be composed of pointlike particles all having the same mass, we can adopt a Dirac $\\delta$ leaving the mass of the particle as an unknown or setting it according to some particle physics model. As a final consequence, the fit to the rotation curve should determine the three HL scaling parameters $(r_A, r_D)$ and ${\\cal{N}}_{DM}$ halo parameters so that severe degeneracies among these ${\\cal{N}}_{DM}$ quantities may take place. In order to reduce them, a better way should be to consider external galaxies rotation curves assuming that the MF is the same as the MW one. In such a way, we could take advantage of the data probing the full radial range and not only the outer disc as done here with the MW sample. Moreover, such a test can also probe the universality of the scaling radii $(r_A, r_D)$ thus providing a further mandatory test of the HL model. As a final remark, we want to stress that, should such an analysis be successful, one has still to address a different issue. Let us suppose that we have indeed well fitted the rotation curves of a large sample of spiral galaxies thus determining a halo model and the values of $(r_A, r_D)$. One could then use Eqs.(\\ref{eq: defhlpar}) and (\\ref{eq: defradii}) to infer constraints on the HL coupling parameters $(g_0, \\ldots, g_8)$. Although we have thus five constraints (the two radii plus the two conditions imposed in the derivation of the potential) so being unable to determine all the eight quantities $(g_0, \\ldots, g_8)$, one could nevertheless try to see whether the allowed region of the parameter space is consistent with what is inferred from cosmological analyses \\cite{DS09} or Solar Sytem tests (see e.g. Iorio \\& Ruggiero 2009 and references therein). It is worth noting that such an attempt should give a consistent picture of the HL framework on all physical scales thus allowing to draw a deeper insight into its viability. \\\\ {\\it Acknowledgements.} We warmly thank Mauro Sereno for helpful comments on a preliminary version of this manuscript. MC is supported by Regione Piemonte and Universit\\`{a} degli Studi di Torino." }, "1003/1003.5849_arXiv.txt": { "abstract": "In order to establish whether the unstable r-modes in a rotating neutron star provide a detectable source of gravitational waves, we need to understand the details of the many dissipative processes that tend to counteract the instability. It has been established that the bulk viscosity due to exotic particles, like hyperons, may be particularly important in this respect. However, the effects of hyperon superfluidity have so far not been fully accounted for. While the associated suppression of the reaction rates that give rise to the bulk viscosity has been estimated, superfluid aspects of the fluid dynamics have not been considered. In this paper we determine the r-mode instability window for a neutron star with a $\\Sigma^{-}$ hyperon core, using the appropriate multifluid formalism including, for the first time, the effect of the ``superfluid'' bulk viscosity coefficients. We demonstrate that, even though the extra terms may increase the bulk viscosity damping somewhat, their presence does not affect the qualitative features of the r-mode instability window. ", "introduction": "Comprising one and a half solar masses inside a radius of roughly ten kilometers, neutron stars provide an arena where many extremes of physics meet. A detailed model of neutron star dynamics must account for strong magnetic fields, various superfluid/superconducting components, the interaction between the crust nuclei and the fluid, as well as exotic states of matter that may be present in the neutron star core. Needless to say, it is a formidable task to construct such a model. Especially since it requires an understanding of physics well beyond the laboratory. While the equation of state for matter approaching the nuclear saturation density $n_0 \\approx 0.16$~fm$^{-3}$ (corresponding to $2.48\\times 10^{14}$~g/cm$^3$) is quite well understood, it seems unlikely that laboratory experiments will ever be able to probe the densities expected in the deep core of a neutron star (above several times $n_0$). Despite decades of research into the supranuclear equation of state, considerable uncertainties remain. Furthermore, neutron star (NS) observations have only recently begun to reach the level of precision necessary to constrain the theoretical models in a severe way. Very recent results by \\citet{Ozel} suggest that current measurements of NS masses and radii can be used to rule out several nuclear equations of state and support the notion that the core of a NS should contain exotic particles, such as pions, kaons, hyperons or even a deconfined quark condensate. Establishing observable signatures of the presence of such exotic states of matter is a priority for modelling in this area. During the last decade, the notion that gravitational waves (GWs) may drive the so-called r-modes of a rotating neutron star unstable, and that this may lead to the star spinning down on a timescale of weeks to months, has been discussed in a number of papers, see \\citet{review} and \\citet{anderssonREVIEW} for reviews. The r-mode instability initially attracted attention because it provided a mechanism that could spin a newly born neutron star down dramatically, releasing GWs at a level that might be detectable in the process \\citep{O98}. For the purpose of GW detection, rapidly rotating accreting neutron stars in low-mass X-ray binaries (LMXBs) have also attracted attention. In these systems, the r-modes could provide a mechanism for torque balance \\citep{Bildsten,AKS} and, in some cases, lead to persistent GW emission \\citep{strange, wagoner2004, Nayyar}. This scenario currently holds interesting prospects for detection, albeit with considerable technical difficulties \\citep{watts2008}. Not surprisingly, the details of the supranuclear equation of state (EOS) are key to understanding the r-mode instability. In a real neutron star many (viscous) mechanisms compete with the GW driving of the r-mode. If the star contains exotic particles, such as hyperons, additional dissipation channels may become relevant. At first sight, it would appear that additional damping should reduce the chances of the r-mode GWs being detectable as it would reduce the region of parameter space where the instability is active. However, this is not necessarily the case. As an illustration of this, let us consider the instability in the temperature range $10^7-10^{10}$~K. First assume that the main damping mechanism is due to a viscous Ekman layer at the core-crust interface (for a discussion see \\citet{eck1, eck2}) (the left panel of figure \\ref{windows}). Standard bulk viscosity due to modified URCA reactions is only relevant above $10^{10}$ K and is not included in this example (we shall discuss this in detail in section \\ref{rmode}). In this case a NS in an LMXB, which will heat up to a core temperature of a few times $10^8$~K, would spin up due to accretion and enter the instability window before the star reaches the break-up limit. The shear from the unstable r-mode then heats the star and the emission of GWs spins the star down until it returns to the stable region. At this point the star will cool down and the cycle can begin again \\citep{levin1}. Unfortunately, with the current estimates for the nonlinear saturation amplitude \\citep{Arras,Brink} the duty cycle for this scenario is very low, meaning that the star would emit brief bursts of gravitational radiation and we would observe most systems in quiescence. Let us contrast this model with a system where we have added the effect of hyperon bulk viscosity (the right panel of figure \\ref{windows}). In this case, extra damping leads to a positive slope of the curve in the $10^9$ K region and there are now three possible scenarios. Depending on the mode amplitude and on the exact details of the damping, the star could either i) execute the cycle we have already described, or ii) the heating may be sufficient for the system to evolve horizontally all they way to the positively sloped part of the curve before GW emission has time to spin the star down \\citep{Bondarescu}. Finally, it may be the case that, iii) the heating due to accretion is such that the system becomes unstable in the region with positive slope. In the last two scenarios the system will not be able to evolve away (significantly) from the instability curve, and should become a persistent source of GWs \\citep{strange}. This hypothesis was examined by \\citet{Nayyar}, who found that, in the case of a NS with a hyperon core, persistent emission is possible over a wide range of parameters for the bulk viscosity and the EOS. \\begin{figure} \\centerline{\\includegraphics[height=7cm,clip]{window1.pdf}\\includegraphics[height=7cm,clip]{window2.pdf}} \\caption{The r-mode instability window for stars with $M=1.4 M_{\\odot}$ and $R=13$km. On the left we have a star where the dominant damping mechanism at low temperatures is shear due to electrons in an Ekman layer at the base of the crust, on the right a star where the main damping mechanism above $10^9$ K is bulk viscosity due to hyperons. The rotation rate is expressed in terms of the parameter $\\tilde{\\omega}=\\omega/\\sqrt{G M/R^3}$. The instability curve in the left panel always has a negative slope and the system undergoes a limit cycle as described in the text. This is in contrast with the instability curve on the right which has a positive slope in the $10^9$ K region, which could halt the thermal run-away and lead to persistent emission of gravitational waves. } \\label{windows} \\end{figure} LMXBs contain old NSs which will have cooled well below the temperature at which the hyperons (and other components of the star, such as neutrons and protons) are likely to become superfluid. Superfluidity adds dimensions to the problem as, not only does it reduce the reaction rates for hyperon creation processes, it also increases the dynamical degrees of freedom of the system. In general, we have to work with multifluid hydrodynamics, with each particle species (potentially) leading to an independent flow. This leads to the appearance of new families of oscillation modes \\citep{Epstein, LM94, ac2001} and also has profound consequences for viscous dissipation. Not only are there new dissipation mechanisms, such as mutual friction between the components \\citep{ALS, Ma,Mb,Trev}, but bulk- and shear viscosity can no longer be described by single coefficients. In general, there are many additional viscosity coefficients. In the context of neutron stars, this was first pointed out by \\citet{CQG}. The particular case of a hyperon core was first considered by \\citet{Gusakov2008} and has recently been analysed in detail by \\citet{prep}. In the simplest case, that of a core comprising neutrons, protons, electrons and $\\Sigma^{-}$ hyperons, one can show that the problem is very similar to that for superfluid Helium \\citep{helium} and can be described by three bulk viscosity and one shear viscosity coefficient. One would, of course, also need to account for mutual friction between the various superfluid components. It has, however, been shown in several calculations that in a NS composed of neutrons, protons and electrons, mutual friction is unlikely to have a significant effect on the r-mode instability \\citep{LM,LY,rmode}. As the nature of mutual friction involving hyperons is largely unknown, we shall assume that it can also be neglected (the veracity of this assertion obviously needs to be checked by detailed work in the future). We will focus on the bulk viscosity, which is expected to give the main contribution to the r-mode damping \\citep{PBJones,LO,haensel,Nayyar}. Although the relevance of the new ``superfluid'' bulk viscosity coefficients is well established for superfluid Helium, their effect has mostly been neglected in the study of NS oscillations. The notable exception is the work of \\citet{Gusakovsound} who studied the damping of sound waves in a dense superfluid hyperon core and showed that the additional bulk viscosity terms can play a significant role. It is clearly important to understand the role of the extra damping coefficients and refine our theoretical understanding of the bulk viscosity, as we have seen that the nature of the damping mechanisms can have profound consequences for the r-mode instability and the associated GW emission. In fact, direct GW detection from these systems should allow us to discern whether the star is emitting persistently or is executing a limit cycle. This would provide valuable information on the physics of the NS interior. The purpose of this paper is to study the effect of superfluid hyperon bulk viscosity on the r-mode instability window. The formalism for studying r-modes in multifluid neutron stars has been developed by \\citet{rmode} and we shall extend it to include hyperon bulk viscosity, thus considering for the first time the global dynamics of a multifluid NS with a hyperon core. ", "conclusions": "Our results represent the first investigation into the effect of including the extra superfluid bulk viscosity coefficients in a calculation of the r-mode instability window. We have shown that, even though the additional bulk viscosity coefficients do not alter the qualitative aspects of the instability window, there are regions of parameter space in which they could play a significant role, and may even suppress the instability entirely. In the light of these results we believe it is important to move beyond the qualitative analysis presented here. One should clearly account for the presence of $\\Lambda$ hyperons and use a more realistic equation of state to describe the star. Furthermore, if one is to construct a more realistic model one clearly needs to work in general relativity and include finite temperature effects (such as in the entrainment coefficients calculated by \\citet{Gusakov2}) and dissipative effects such as hyperon bulk viscosity. More detailed theoretical input is also needed from the nuclear physics community, in order to calculate the superfluid reaction rates needed to evaluate the bulk viscosity coefficients. Developing the relevant tools will allow us to make progress on a range of related problems, e.g. involving finite temperature superfluid in the outer neutron star core or exotic phases of deconfined quarks in the deep core. An improved understanding of these systems is crucial if we are to take advantage of the unique opportunity that gravitational-wave detection could offer for the study of matter under extreme conditions." }, "1003/1003.3468_arXiv.txt": { "abstract": "The thermal stability of rotating, stratified, unmagnetized atmospheres is studied by means of linear-perturbation analysis, finding stability, overstability or instability, depending on the properties of the gas distribution, but also on the nature of the perturbations. In the relevant case of distributions with outward-increasing specific entropy and angular momentum, axisymmetric perturbations grow exponentially, unless their wavelength is short enough that they are damped by thermal conduction; non-axisymmetric perturbations typically undergo overstable oscillations in the limit of zero conductivity, but are effectively stabilized by thermal conduction, provided rotation is differential. To the extent that the studied models are representative of the poorly constrained hot atmospheres of disc galaxies, these results imply that blob-like, cool overdensities are unlikely to grow in galactic coronae, suggesting an external origin for the high-velocity clouds of the Milky Way. ", "introduction": "Galaxy clusters and massive elliptical galaxies are embedded in hot atmospheres of virial-temperature gas revealed by X-ray observations. Lower-mass galaxies are also believed to have hot gaseous coronae, difficult to detect because the gas is very rarefied. If the coronal gas were thermally unstable \\citep{Fie65} it could fragment into cold gas clouds with substantial implications for the dynamics of cooling flows \\citep{Mat78,Cow80,Nul86}, but also for galaxy formation and for the origin of the \\hi\\ high-velocity clouds of the Milky Way \\citep{Mal04,Pee08,Kau09}. In fact, linear-perturbation analysis has shown that the X-ray emitting hot atmospheres of massive elliptical galaxies and galaxy clusters are likely stabilized against thermal instability by a combination of buoyancy and thermal conduction \\citep[Malagoli, Rosner \\& Bodo 1987, hereafter MRB;][]{BS89,Tribble89}. The coronae of disc galaxies are expected to have lower gas temperatures and densities, but their detailed properties are poorly known, because---so far---they eluded detection in X-rays. Nevertheless, their physical parameters are not totally unconstrained: for some nearby massive disc galaxies upper limits to the total X-ray luminosity of the hot halos are available \\citep{Ras09}, while in the special case of the Milky Way different pieces of information can be used to constrain the physical properties of the corona \\citep[e.g.][]{Spi56,Fuk06}. Studying model coronae consistent with these constraints, \\citet[][hereafter BNF]{Bin09} showed that also the coronae of disc galaxies are likely thermally stabilized by buoyancy and thermal conduction, at least if rotation of the gas can be neglected. There are no significant constraints on the angular momentum of galactic coronae, but one might expect that they are characterized by at least slow rotation. It is well known that there is an important interplay between rotation and convection \\citep[e.g.][]{Tas78,Bal00,Bal01}, so rotation might influence the stabilizing effect of buoyancy and then the overall problem of the thermal stability of galactic coronae. This paper addresses the question of the thermal stability of rotating galactic coronae. For this purpose one needs to study the stability against non-axisymmetric perturbations of a differentially rotating, stratified gas in the presence of thermal conduction and radiative cooling. This is a special case of the general problem of the stability of a rotating stratified gas, which has been widely studied in the astrophysical literature, with specific applications to rotating stars~\\citep[e.g.][]{Cow51,Gol67,Leb67,Fri68,Lyn67,Tas78,Kno82,Lif93,Men04} and accretion discs~\\citep[e.g.][]{Pap84,BalH91,BalH92,Ryu92,Lin93,Bra06}. The stability-analysis techniques developed in these previous studies can be applied to the problem addressed in the present paper. In particular, we will perform a linear analysis using Eulerian perturbations, considering both axisymmetric and non-axisymmetric disturbances. It is reasonable to expect that galactic coronae are magnetized, though there are almost no constraints on the properties of their magnetic fields. For simplicity, in the present work we neglect magnetic fields, but we must bear in mind that they might have some influence on the thermal stability properties of the galactic hot gas, as suggested by previous studies of thermal instability in non-rotating cooling flows \\citep[][see also BNF for a discussion]{Loe90,Bal91}. The paper is organized as follows. In Section~\\ref{sec:eq} the problem is set up by recalling the relevant hydrodynamic equations. The thermal-stability analysis is then described for axisymmetric perturbations (Section~\\ref{sec:axisymm}) and for non-axisymmetric perturbations (Section~\\ref{sec:nonaxisymm}). The implications for galactic coronae are discussed in Section~\\ref{sec:impl}; Section~\\ref{sec:con} concludes. A list of commonly used symbols is given for reference in Table~\\ref{tab:list}. ", "conclusions": "\\label{sec:con} Motivated by the question of whether cool, pressure-supported clouds can condense out of the hot coronae of disc galaxies, we studied the problem of the thermal stability of rotating stratified fluids, in the presence of radiative cooling and thermal conduction. As in the non-rotating case, the time evolution of a perturbation depends on its wavelength, so it is useful to distinguish short-wavelength perturbations (such that the dissipative frequency $\\omegad$ is positive) from long-wavelength perturbations (such that the dissipative frequency $\\omegad$ is negative). We found that---against either axisymmetric or non-axisymmetric perturbations---a {\\it uniformly rotating}, convectively stable configuration is thermally stable when $\\omegad>0$ (damping by thermal conduction), but is thermally unstable when $\\omegad<0$. Similarly---against axisymmetric perturbations---a {\\it differentially rotating}, convectively stable corona with outward-increasing specific angular momentum is thermally stable when $\\omegad>0$, but is thermally unstable when $\\omegad<0$. Non-axisymmetric perturbations in the presence of differential rotation behave differently from the axisymmetric ones. In the {\\it absence of thermal conduction}, the combination of non-axisymmetry of the perturbation and of differential rotation has a stabilizing effect (turning instability into overstability), and barotropic distributions tend to be more stable than baroclinic distributions. In the {\\it presence of thermal conduction}, stability replaces overstability: in differentially rotating systems, the shear makes conductive damping of non-axisymmetric disturbances particularly effective for both barotropic and baroclinic distributions. These results have been discussed in the context of the problem of the growth of thermal perturbations in galactic coronae. The above calculations allow to address unambiguously this question when applied to a specific model of rotating galactic corona. Unfortunately, given that the physical properties of the hot atmospheres of disc galaxies (in particular, the distribution of specific entropy and angular momentum) are very poorly constrained observationally, it is difficult to be conclusive about whether these systems are prone to thermal stability. Additional uncertainties come from magnetic fields, which are expected to be present in real systems, but have been neglected for simplicity in the studied models. The question of the effect of magnetic fields on the thermal instability is very interesting and would require a full thermal-stability analysis of a rotating, magnetized corona, for which our calculations may be a starting point. Though limited by the mentioned uncertainties, the present work gives further indications against the hypothesis that thermal instability is important for galactic coronae. In particular, though rotation can destabilize against thermal disturbances very elongated in the azimuthal direction, we argued that blob-like thermal perturbations are unlikely to grow in a differentially rotating corona. This finding, combined with previous results on non-rotating coronae (BNF), suggests that the high-velocity clouds of the Milky Way did not form spontaneously from small thermal perturbations in the Galactic corona, but must be of external origin. A possibility is that at least the seeds of these clouds are formed originally as a consequence of stripping of gas-rich satellites or cosmic infall of cold gas. These cool gaseous seeds might then grow while travelling through the hot Galactic corona \\citep{Som06,Ker09}, for instance via turbulent mixing \\citep{Mar10}. Provided that the overdensities associated with these accreted seeds are substantial (i.e. they are non-linear perturbations), such a scenario for the formation of the high-velocity clouds is not in contrast with the results of the present paper, in which only linear perturbations are considered." }, "1003/1003.3697_arXiv.txt": { "abstract": "We study compressible MHD turbulence, which holds key to many astrophysical processes, including star formation and cosmic ray propagation. To account for the variations of the magnetic field in the strongly turbulent fluid we use wavelet decomposition of the turbulent velocity field into Alfv\\'{e}n, slow and fast modes, which presents an extension of the Cho \\& Lazarian (2003) decomposition approach based on Fourier transforms. The wavelets allow to follow the variations of the local direction of magnetic field and therefore improve the quality of the decomposition compared to the Fourier transforms which are done in the mean field reference frame. For each resulting component we calculate spectra and two-point statistics such as longitudinal and transverse structure functions, as well as, higher order intermittency statistics. In addition, we perform the Helmholtz-Hodge decomposition of the velocity field into the incompressible and compressible parts and analyze these components. We find that the turbulence intermittency is different for different components and we show that the intermittency statistics depend on whether the phenomenon was studied in the global reference frame related to the mean magnetic field or it was studied in the frame defined by the local magnetic field. The dependencies of the measures we obtained are different for different components of velocity, for instance, we show that while the Alfv\\'{e}n mode intermittency changes marginally with the Mach number the intermittency of the fast mode is substantially affected by the change. ", "introduction": "\\label{sec:intro} Astrophysical fluids are magnetized and therefore the astrophysical turbulence is magnetohydrodynamic (MHD) in its nature. Compressible MHD turbulence is a key element for understanding star formation \\citep[see][and references therein]{maclow04a,elmegreen04,mckee07} and velocity fluctuations determine many of its properties. For instance, in the modern paradigm of star formation that turbulent velocity sweep up the matter from large expands of the interstellar space to create molecular clouds. Thus, it is important to know the statistical properties of the velocity field, e.g. its spectrum that reflects how much energy is associated with the motions at a particular scale (see below). A further insight into the properties of turbulence, including its generation, consequences and dissipation calls for the use of more sophisticated measures. For example, the processes of magnetic field generation depend on the velocity field vorticity associated with the solenoidal motions, while the processes of compressing gas are determined by the compressible motions. In approaching the problem of decomposing velocity field into solenoidal and dilatational part following the Helmholtz-Hodge decomposition is frequently attempted \\cite[see][]{federrath08,federrath09}. Another approach is the decomposition of the turbulent field in MHD case into basic modes, i.e., Alfv\\'en, slow and fast waves. While this approach is trivial for the case of the strong magnetic field with infinitesimal fluctuations \\cite[see][]{dobrowolny80}, \\cite{cho02a,cho03} proposed a {\\it statistical} decomposition of modes in the Fourier space. The statistical nature of the procedure is clear when one considers its application to strongly perturbed magnetic fields. As the Fourier transform is defined in the reference frame related to the mean magnetic field, while the MHD motions happen in respect to the local magnetic field, there is an inevitable contribution of all types of motion to the decomposed modes. However, studying the cases when the real space decomposition was possible in real space, \\citet[][henceforth CL03]{cho03} showed that the cross-talk between the modes is small for subAlfv\\'enic turbulence. Testing of the results in CL03 and increasing the accuracy of the MHD mode decomposition of turbulence is one of the key goals of the present study. In doing so in the present paper we make use of the wavelet transformations in addition to the Fourier transformations. Wavelets \\cite[see][]{meneveau91a,meneveau91b} present a natural way of describing MHD turbulence. Indeed, while in the representation of the \\citet[][henceforth GS95]{goldreich95} model of turbulence the anisotropy is frequently described in terms of eddies with parallel $k_{\\|}$ and perpendicular $k_{\\bot}$ wave vectors, the actual description calls for choosing for $\\|$ and $\\bot$ in respect to the {\\it local} magnetic field\\footnote{Due to this fact that closure relations used for the model justification in GS95 are doubtful. The importance of the {\\it local} system of reference was clearly stressed in the works that followed the original GS95 study \\citep{lazarian99,cho00,maron01,cho02b,cho03b,lithwick01,cho02a}.}. The latter is really easy to understand, as it is only local magnetic field that influences fluid motions at a given point. Wavelets allow for a local description of the magnetized turbulent eddies. In the paper we decompose the turbulent velocity fields using both wavelets and a more traditional Helmholtz-Hodge decomposition into solenoidal and compressible parts. We feel that the latter decomposition is more justified for the hydrodynamic turbulence than for the MHD turbulence that we study here. However, we feel that the use of the Helmholtz-Hodge decomposition provides an additional, although limited, insight into the properties of compressible motions. The three major properties of the velocity field that we focus our attention in the paper are turbulence spectra, anisotropies and intermittency. These three measures require further description that we provide below. While turbulence is an extremely complex chaotic non-linear phenomenon, it allows for a remarkably simple statistical description \\cite[see][]{biskamp03}. If the injections and sinks of the energy are correctly identified, we can describe turbulence for {\\it arbitrary} $Re$ and $Rm$. The simplest description of the complex spatial variations of any physical variable, $X({\\bf r})$, is related to the amount of change of $X$ between points separated by a chosen displacement ${\\bf l}$, averaged over the entire volume of interest. Usually the result is given in terms of the Fourier transform of this average, with the displacement ${\\bf l}$ being replaced by the wave number ${\\bf k}$ parallel to ${\\bf l}$ and $|{\\bf k}|=1/|{\\bf l}|$. For example, for isotropic turbulence the kinetic energy spectrum, $E(k)dk$, characterizes how much energy resides at the interval $k, k+dk$. At some large scale $L$ (i.e., small $k$), one expects to observe features reflecting energy injection. At small scales, energy dissipation should be seen. Between these two scales we expect to see a self-similar power-law scaling reflecting the process of non-linear energy transfer. We shall attempt to get the power-law scalings for the components of the velocity field. The presence of a magnetic field makes MHD turbulence anisotropic \\citep{montgomery81,matthaeus83,shebalin83,higdon84,goldreich95}\\cite[see][for review]{oughton03}. The relative importance of hydrodynamic and magnetic forces changes with scale, so the anisotropy of MHD turbulence does too. Many astrophysical results, e.g. the dynamics of dust, scattering and acceleration of energetic particles, thermal conduction, can be obtained if the turbulence spectrum and its anisotropy are known \\cite[see][for review]{lazarian09b}. The knowledge of the anisotropy of Alfvenic mode of MHD turbulence determines the extend of magnetic field wandering influencing heat transfer \\citep{narayan01,lazarian06a} and magnetic reconnection \\citep{lazarian99,kowal09}. We would like to stress that in what follows we discuss the properties of {\\it strong} MHD turbulence. This type of turbulence is not directly related to the amplitude of the magnetic perturbations, however. The low-amplitude turbulence can be strong and isotropically driven turbulence with $\\delta B\\ll B$ at the injection scale exhibits only a limited range of scales for which it is weak \\cite[see][]{galtier00}, while at sufficiently small scales it gets strong \\cite[see the discussion in][]{lazarian99}. An anisotropic spectrum alone, say $E({\\bf k)}\\,d{\\bf k}$, cannot characterize MHD turbulence in all its complexity because it involves only the averaged energy in motions along a particular direction. To have a full statistical description, one needs to know not only the averaged spectrum of a physical variable but higher orders as well. The tendency of fluctuations to become relatively more violent but increasingly sparse in time and space as the scales decreases, so that their influence remains appreciable, is called {\\it intermittency}. The intermittency increases with the ratio of the size scales of injection and dissipation of energy, so the very limited range of scales within numerical simulations may fail to reflect the actual small scale processes. The turbulence intermittency can result in an important intermittent heating of the interstellar medium \\citep{falgarone05,falgarone06,falgarone07}. In this article we investigate the scaling properties of the structure functions of velocity and its components for compressible MHD turbulence with different sonic and Alfv\\'{e}nic Mach numbers. In \\S\\ref{sec:models} we describe the numerical models of compressible MHD turbulence. We decompose velocity into a set of components including the incompressible and compressible parts, and MHD waves: Alfv\\'{e}n, slow and fast using methods described in \\S\\ref{sec:decomp}. In \\S\\ref{sec:spectra} we study spectra of velocity and its components. In \\S\\ref{sec:anisotropy} we study the anisotropy of dissipative structures. We show differences in the structures for different components. In \\S\\ref{sec:intermittency} we study the scaling exponents and the intermittency of velocity structures. We show their dependence on the sonic and Alfv\\'{e}n regime of turbulence. In \\S\\ref{sec:discussion} we discuss our results and their relation to the previous studies. In \\S\\ref{sec:summary} we draw our conclusions. ", "conclusions": "\\label{sec:discussion} \\subsection{Major Accomplishments and Limitations of the Present Study} In the paper we have introduced a new procedure of decomposition of MHD turbulence field into Alfv\\'en, slow and fast modes which uses wavelets. Compared to the decomposition procedure based on Fourier transforms described in CL03, the wavelet decomposition is more local, thus it follows better the local magnetic field direction in respect to which the decomposition into modes takes place. As a result, we expect that the wavelet decomposition procedure to be more accurate for larger amplitudes of turbulence, i.e. larger perturbations of magnetic field. Our decomposition of the MHD turbulence confirmed the results in CL03 in terms of spectra, namely, that the Alfv\\'enic and slow modes are anisotropic and consistent with the predictions of the GL95 model of incompressible turbulence, while the fast modes are mostly isotropic and form an acoustic turbulence cascade \\citep{lithwick01,cho02a}. As these results were used in studies of cosmic ray scattering and acceleration \\citep{yan02,yan04a,yan08,cho05,brunetti07} as well as charged dust acceleration \\citep{lazarian02,yan03,yan04b,yan09} this is an encouraging development. At the same time, the intermittency of the different MHD modes were shown to be very different. We clearly see the dependence of high order statistics of compressible motions on the Mach number. We interpret this dependence as the result of shock formation, which eventually changes the nature of the compressible motion cascade compared to the CL03 assumptions. The limitations of the present study arise from the yet unclear nature of the turbulent cascade. For instance, it was shown in \\cite{beresnyak09} that the degree of locality of interactions in hydrodynamic and MHD cascade are different. Thus even largest available MHD simulations may not present the actual inertial range of the cascade, but the measured slope may be strongly affected by the extended bottle-neck effect of the simulations. In addition, the limited range over which Alfvenic turbulence is weak may exhibit a rather different scaling of fast modes as a result of the interactions of the Alfvenic and fast modes \\citep{chandran05}. In addition, within the present study we intentionally do not consider the scaling of magnetic perturbations. The velocity and magnetic perturbations for subAlfvenic turbulence show some differences, which are rather difficult to study reliably with the available numerical simulations. These differences are not a part of the GS95 picture, but may reflect additional yet unclear properties of the MHD cascade \\cite[see][]{mueller00}. In our study we used only the incompressible driving. In the presence of the compressible supersonic driving \\citep{federrath09} the scaling looks different, but the existence of the inertial range is then questionable. \\cite{kritsuk09} claimed that combining the compressible and incompressible driving in the Mach number dependent fashion one can obtain a better power-law inertial range. This issue requires further studies. The turbulence driving in our study is balanced, in the sense that the energy flows in opposite directions are equal. In the presence of sources and sinks of turbulent energy, astrophysical turbulence is expected to be imbalanced. Our numerical studies of imbalanced turbulence in \\cite{beresnyak10} show that the properties of Alfvenic turbulence changes substantially in the presence of imbalance. However, the degree of sustainable imbalance in compressible turbulence is still unclear. One expects the density fluctuation in turbulent fluid to reflect the incoming waves, altering the imbalance. We believe that in high Mach number fluids the imbalance is low due to the existence of substantial density contrasts. \\subsection{Astrophysical Implications of the Turbulence Anisotropy and Intermittency} Depending on driving astrophysical turbulence may be subAlfv\\'enic, if the injection velocity $V_L$ is less than Alfven speed $V_A$, Alfvenic, if $V_L=V_A$, and superAlfvenic, if $V_L>V_A$. This frequently is also described by the Alfven Mach number $M_A=V_L/V_A$. Formally, the GS95 model applies only to incompressible motions with $V_L=V_A$, or equivalently $M_A=1$. Some of the astrophysical applications of the model, indeed, use the original form of the theory, which substantially limits the applications of the theory \\cite[see][]{narayan01}. However, the model can be easily generalized to cover extensive ranges of superAlfvenic and subAlfvenic turbulence \\cite[see][]{lazarian99,lazarian06a}. For subAlfvenic turbulence with isotropic driving at the scale $L$ an initial weak cascade, in which the parallel scale of motions stays the same and the spectrum $E(k_{\\bot})\\sim k_{\\bot}^{-2}$ is applicable, transfers to the regime of strong turbulence at the scale of $LM_A^2$, for which the GS95 critical balance arguments are applicable. For superAlfvenic turbulence, while up to the scale $LM_A^{-3}$ the turbulence is hydrodynamic, it approaches the GS95-type regime for smaller scales. Therefore, the relations obtained for MHD turbulence that we have studied above can be generalized for cases of different intensity of driving. While the GS95 model is a model of incompressible turbulence, our simulations confirm the numerical findings in \\cite{cho02a} and CL03 that the scaling of the Alfvenic mode in the compressible turbulence is very similar to its scaling in the incompressible case. In particular, the GS95 anisotropy of MHD turbulence determines the rate of magnetic field wandering which is important for many astrophysical processes, including the ubiquitous process of magnetic reconnection \\citep{lazarian99}. Additional implications of magnetic field wandering include the diffusion of heat and cosmic rays, MHD acceleration of dust etc. \\cite[see][for a review]{lazarian09b}. The wavelet approach has the potential of increasing accuracy while studying small-scale anisotropy in simulations with strongly perturbed magnetic fields. \\cite{falgarone05,falgarone06,falgarone07} and collaborators \\cite[and refs. therein]{hily-blant07a,hily-blant07b} attracted the attention of the interstellar community to the potential important implications of intermittency. A small and transient volume with high temperatures or violent turbulence can have significant effects on the net rates of processes within the ISM. For instance, many interstellar chemical reactions (e.g., the strongly endothermic formation of CH$^+$) might take place within very intensive intermittent vortices. The aforementioned authors claimed the existence of the observational evidence for such reactions and heating, but a more quantitative approach to the problem is possible. \\cite{beresnyak07} \\cite[see also][]{lazarian09b} used the intermittency scaling and calculated the distribution of the dissipation rate in the turbulent volumes. In doing so they used the fact that \\cite{she94} model of intermittency corresponds generalized log-Poisson distribution of the local dissipation rates \\cite{dubrulle94,she95}. The obtained rates of enhancement were not sufficient to explain the heating required for inducing interstellar chemistry \\citep{beresnyak07}. The same approach was used by \\cite{pan09} who obtained, however, a different result. We believe that one should distinguish shocks from vortical motions while calculating the heating induced by intermittency. Our present study show very different scalings relevant to these types of motions. \\subsection{Studies of Compressible MHD Turbulence in Astrophysical Context} Numerous studies of compressible MHD turbulence are done in the context of star formation \\cite[see reviews by][and references therein]{maclow04b,mckee07}. Most of these simulations are focused on the large-scale appearances of the turbulence, which is determined by the turbulent driving and do not exhibit any extended inertial range of turbulence. Search for the universal relations for compressible turbulence resulted in the rise of interest to the \\cite{fleck83} idea of searching universality not for velocity, but for the combination of the velocity and density in the form $\\rho^{1/3} v$. The numerical study of hydrodynamic compressible turbulence revealed that, indeed, the density modified velocity shows the same Kolmogorov scaling both for low and high Mach number turbulence \\citep{kritsuk07}. Similar effect was confirmed in our MHD simulations \\cite{kowal07b}. However, the physical justification of this universality is unclear and it may result just from the coincidental compensation of the change of velocity and density indexes as shocks develop at high Mach number turbulence. We report steepening of the spectra of compressible motions at high Mach numbers. High resolution hydro simulations \\cite[see][]{kritsuk07} show that the velocity spectrum becomes steeper for high Mach number simulations. This corresponds to the observational studies of the supersonic velocity turbulence in \\cite{padoan06,padoan09} and \\cite{chepurnov06}. These studies are done with the Velocity Channel Analysis (VCA) and Velocity Coordinate Spectrum (VCS) techniques, which are theory-motivated and tested techniques \\citep{lazarian00,lazarian04,lazarian06b,lazarian08,chepurnov08}. The application of these techniques should enhance the range of astrophysical turbulent velocity fields that can be studied observationally\\footnote{Recent examples of the techniques of observational studies of the turbulent density field can be found in \\cite{kowal07a,burkhart09,burkhart10}. The velocity is a more covered statistics, but it is more difficult to study \\cite[see][for a review]{lazarian09a}.} It is comparing numerics, observations and theory that the progress in understanding of turbulence requires. In this article we presented a new technique of decomposition of turbulent MHD motions into Alfven, slow and fast modes. The technique is based on the use of wavelets, which provide a more local decomposition compared to the Fourier approach in CL02 and CL03. This enables one to have better accuracy of the decomposition of MHD turbulence into fundamental modes for higher amplitude of magnetic perturbations. By applying the wavelet decomposition to the results of our simulations of compressible MHD turbulence, we investigated the scaling properties of velocity in compressible MHD turbulence for different sonic $M_s$ and Alfv\\'{e}nic $M_A$ Mach numbers. We analyzed spectra, the anisotropy, scaling exponents and intermittency of the total velocity and its components corresponding to the Alfv\\'{e}n, slow and fast modes. We found that: \\begin{itemize} \\item The amplitude of velocity fluctuations depends on $M_s$ only marginally. The lack of significant dependence of the velocity fluctuations is also observed for its incompressible part, as well as for the Alfv\\'{e}n and slow waves. The compressible part of velocity and the fast wave show a dependence on ${\\cal M}_s$, but only for subsonic turbulence. In the case of supersonic models, the fluctuations of the compressible part and fast mode of the velocity have comparable amplitudes. \\item The spectral indices depend on $M_A$ in turbulence with a strong magnetic field. In the case of turbulence with a weak magnetic field only the indices of spectra of the fast wave change between sub- and supersonic models. For the other components, the spectral indices do not change appreciably with the sonic Mach number. While our conclusions about spectra of fast modes for subsonic turbulence agree with the CL03 conclusion about the acoustic cascade of these modes, we feel that for high Mach number we get the spectrum of shocks. The anisotropy of Alfv\\'enic turbulence and slow modes is in agreement with GS95 theory for both the cases of high and low beta plasmas. The velocity fluctuations of the fast modes demonstrate isotropy. \\item In the global reference frame, we observe stronger changes of the scaling exponents and intermittency for velocity and its all components with ${\\cal M}_s$ in the case of turbulence with a weak magnetic field. The intermittency of structures grows with the values of ${\\cal M}_s$. However, when the external magnetic field is strong, the intermittency for all components depends on the sonic Mach number only marginally. In the local reference frame, the scaling exponents turbulence depend on the direction with respect to the direction of the local mean magnetic field. The dependence is stronger for the subAlfv\\'{e}nic turbulence. \\end{itemize}" }, "1003/1003.4281_arXiv.txt": { "abstract": "The solar F10.7 index is has been a reliable and sensitive activity index since 1947. As with other indices, it has been showing unusual behavior in the Cycle 23/24 minimum. The origins of the solar microwave flux lie in a variety of features, and in two main emission mechanisms: free-free and gyroresonance. In past solar cycles F10.7 has correlated well with the sunspot number SSN. We find that this correlation has broken down in Cycle~23, confirming this with Japanese fixed-frequency radiometric microwave data. ", "introduction": "The F10.7 daily solar index, introduced by Covington from 1947 (see Tapping et al., 2003, for more detail) has become a standard measure of solar activity and as a proxy for many variables not routinely measured. \\nocite{tapping03} It represents the microwave flux density at 10.7~cm wavelength, which corresponds coincidentally to the electron Larmor frequency for $|${\\bf B}$|$~=~10$^3$~G. Shortly after its introduction, routine observations in Japan began at a set of four fixed frequencies (1.0, 2.0, 3.75, and 9.4~GHz), straddling the F10.7 frequency (2.8~GHz) and encompassing the gyroresonance signature in the solar microwave spectrum. \\citet{tanaka73} used these and other data to establish a uniform absolute photometric standard for the microwave range. Microwave emission from the quiet Sun has many sources, all detectable via free-free emission or thermal gyroresonance radiation. Coronal holes may be bright, while filaments are dark. A basic minimum level of emission somewhere below 70~SFU (one solar flux unit is 10$^{-22}$~W/m$^2$Hz) is presumed to come from the quiet background photosphere, and an ``S-component'' due to the other components, especially active regions and the enhanced network, varies on all time scales longer than those of flares. The Canadian (F10.7) and Japanese (four fixed frequency) measurements have continued in an uninterrupted time series up to the present. The only substantial changes in the observing programs appear to have been the observing sites in each case: in 1991 the F10.7 facilities moved from Ottawa to Penticton, and in 1994 the Japanese facilities moved from Toyokawa to Nobeyama. Each of these changes induced small but definite shifts of calibration, as we discuss further below. The two microwave data sets compare extremely favorably with one another, as we show in Figure~\\ref{fig:ts}. An adjusted composite of the data shows that the six minima thus far observed all agree closely. \\begin{figure} \\plotone{timeseries.eps} \\caption{\\label{fig:ts} The full length of the timeseries data for the microwave indices. We have scaled the four Nobeyama fixed frequefncies to F10.7 and overplotted them as monthly averages. Note the stability of the minima (box). } \\end{figure} ", "conclusions": "We confirm the anomalous long-term behavior of the non-flare solar microwave fluxes as represented in the F10.7 index, reported elsewhere by W.~D. Pesnell and by K.~Tapping (personal communications, 2009). Cycle 23 appears to have been anomalous when referred to the sunspot number SSN. We suggest that the anomaly really lies in a secular variation in the nature of sunspots, rather than in the active-region magnetism responsible for the microwave S-component as such. \\bigskip\\noindent {\\bf Acknowledgements:}" }, "1003/1003.3718_arXiv.txt": { "abstract": "{ In this work we study the contribution of magnetic fields to the Sunyaev Zeldovich (SZ) effect in the intracluster medium. In particular we calculate the SZ angular power spectrum and the central temperature decrement. The effect of magnetic fields is included in the hydrostatic equilibrium equation by splitting the Lorentz force into two terms \\(\\hbox{--}\\) one being the force due to magnetic pressure which acts outwards and the other being magnetic tension which acts inwards. A perturbative approach is adopted to solve for the gas density profile for weak magnetic fields (\\( \\le 4 \\mathrm{\\mu G}\\)). This leads to an enhancement of the gas density in the central regions for nearly radial magnetic field configurations. Previous works had considered the force due to magnetic pressure alone which is the case only for a special set of field configurations. However, we see that there exists possible sets of configurations of ICM magnetic fields where the force due to magnetic tension will dominate. Subsequently, this effect is extrapolated for typical field strengths (\\(\\sim 10 \\mu G\\)) and scaling arguments are used to estimate the angular power due to secondary anisotropies at cluster scales. In particular we find that it is possible to explain the excess power reported by CMB experiments like CBI, BIMA, ACBAR at \\(\\ell > 2000\\) with \\(\\sigma_{\\rss 8} \\sim 0.8\\) (WMAP 5 year data) for typical cluster magnetic fields. In addition we also see that the magnetic field effect on the SZ temperature decrement is more pronounced for low mass clusters (\\(\\langle T \\rangle \\sim \\) 2 keV). Future SZ detections of low mass clusters at few arc second resolution will be able to probe this effect more precisely. Thus, it will be instructive to explore the implications of this model in greater detail in future works.} ", "introduction": "It has been known that the intracluster medium has magnetic fields of micro-gauss strength. They affect the evolution of galaxies \\cite{ars09}, contribute significantly to the total pressure of interstellar gas, are essential for the onset of star formation \\cite{sch09}, and control diffusion, confinement and evolution of cosmic rays in the intracluster medium (ICM) \\cite{kus09}. In clusters of galaxies, magnetic fields may play also a critical role in regulating heat conduction (e.g., \\cite{chan98,nar01}), and may also govern and trace cluster formation and evolution. Magnetic fields in the intra-cluster medium have been inferred in various manners using diagnostics such as radio synchrotron relics within clusters, inverse Compton X-ray emissions from clusters, Faraday rotation measures of polarized radio sources within or behind clusters and cluster cold fronts in X-ray (\\cite{clar01,car02}). From these observations it can be inferred that the medium within most clusters is magnetized with typical field strengths at a few \\(\\mu\\)G level distributed throughout the cluster scale. In the cores of \"cooling flow clusters\" (\\cite{eil02,tay02}) and also in cold fronts \\cite{vi01}, the field strengths may reach 10\\(\\mu\\)G-40\\(\\mu\\)G and could be dynamically important. It is thus essential to study the effect of an intra- cluster magnetic field on the gas density distribution in the ICM. A very important observational probe of the cluster gas density distribution is the thermal Sunyaev Zeldovich effect \\cite{sun70,suny72}. This effect occurs because of the re-scattering of primary CMB photons with the hot electrons in the ICM resulting in secondary anisotropies in the CMB at small angular scales (for a detailed review see \\cite{bir99}). At 30 GHz, anisotropies from the thermal SZ effect are expected to dominate over the primary CMB fluctuations for multipoles \\(\\ell \\le 2500\\). The angular power of SZ fluctuation depends sensitively on the integrated cluster abundance and cluster gas distribution (for details see \\cite{ag08}). The SZ effect in clusters depends directly on the density and temperature profile of the ICM. For any particular cluster, it is quantified through the y-parameter which is essentially the temperature decrement along a line of sight through the cluster. This is therefore an important observational probe and can be used to constrain/detect any additional parameter that affects the density of the ICM. In the present work, we study the effect of intracluster magnetic fields on the gas density profile of the ICM and compute the central SZ decrement for different cluster masses as well as the angular power spectrum . To study magnetic field effects in the ICM, earlier authors have incorporated it by introducing the magnetic pressure in the hydrostatic equilibrium condition (\\cite{koch03,cola06, zh04}). Although this is an important effect, there is, in addition to the pressure, a contribution arising from the tension force due to the field. We show in this paper, that the most general treatment involves using both terms arising out of the Lorentz force due to the field. In general, the detailed form of the magnetic field configuration is unknown and moreover it can vary from cluster to cluster. It is only in the specific case of isotropic configurations, that the magnetic pressure is the only contribution. However, there could be other plausible configurations of magnetic fields in clusters where both the pressure and the tension force contribute. In this paper, we classify the different field configurations and study the problem for a near-radial field . The paper is organized as follows: Section 2 presents the formulation of the problem by setting up the hydrostatic equilibrium equation incorporating the magnetic field. Sections 3 and 4 present the mathematical formulation of the y-parameter and the CMB angular power spectrum respectively. Section 5 presents the results for the density profiles and SZ observables which are got by solving the hydrostatic equilibrium equation and finally section 6 summarizes the results and concluding remarks. ", "conclusions": "In this work, we have studied the effect of an intracluster tangled magnetic field on the gas distribution in the ICM. In addition, we have also investigated the effect of cluster magnetic fields on the central SZ decrement for a range of cluster masses. We have also computed the CMBR angular spectrum for different field strengths . In contrast to previous studies, we incorporated the complete radial Lorentz force due to the magnetic field in the hydrostatic equilibrium equations, thereby introducing the effect of magnetic tension as well in addition to magnetic pressure for a generic field configuration. It is only in the special case of an isotropic configuration that magnetic pressure is the sole contribution. However, realistic cases would involve a range of configurations. In particular we presented the results for a nearly radial magnetic field configuration. The results would be more pronounced for a configuration of magnetic fields in which the tension force is dominant. It would be interesting to look at all the other important baryonic effects present in the cluster gas like AGN heating, cosmic-ray heating and the like in addition to this effect of the magnetic field to determine the structure of the intracluster gas in more detail and to disentangle various contributions. Our results can thus be summarized as follows: \\begin{enumerate} \\item For a nearly radial magnetic field, the ICM gas density shows an enhancement in the regions close to the center whereas there is depletion in the outer regions, the crossover scale being dependent on the cluster mass. \\item The gas density enhancement/depletion is large for low-mass clusters. \\item The central SZ decrement is enhanced compared to the default (i.e no magnetic field) case if magnetic field effects are included because of the enhancement in the gas density \\item The CMB angular power is also enhanced in the presence of magnetic fields and future precise observations on these small scales can be used to constrain the strength of such fields and their configurations. \\item For an isotropic field configuration, the effects can be opposite to what we have concluded in this piece of work, as pointed out by earlier authors. \\end{enumerate}" }, "1003/1003.5658_arXiv.txt": { "abstract": "We present a sample of 20 massive galaxy clusters with total virial masses in the range of $6 \\cdot 10^{14}M_{\\odot} \\leq M_{vir} \\leq 2 \\cdot 10^{15}M_{\\odot}$, re-simulated with a customized version of the 1.5. ENZO code employing Adaptive Mesh Refinement. This technique allowed us to obtain unprecedented high spatial resolution ($\\approx 25kpc/h$) up to the distance of $\\sim 3$ virial radii from the clusters center, and makes it possible to focus with the same level of detail on the physical properties of the innermost and of the outermost cluster regions, providing new clues on the role of shock waves and turbulent motions in the ICM, across a wide range of scales. In this paper, a first exploratory study of this data set is presented. We report on the thermal properties of galaxy clusters at $z=0$. Integrated and morphological properties of gas density, gas temperature, gas entropy and baryon fraction distributions are discussed, and compared with existing outcomes both from the observational and from the numerical literature. Our cluster sample shows an overall good consistency with the results obtained adopting other numerical techniques (e.g. Smoothed Particles Hydrodynamics), yet it provides a more accurate representation of the accretion patterns far outside the cluster cores. We also reconstruct the properties of shock waves within the sample by means of a velocity-based approach, and we study Mach numbers and energy distributions for the various dynamical states in clusters, giving estimates for the injection of Cosmic Rays particles at shocks. The present sample is rather unique in the panorama of cosmological simulations of massive galaxy clusters, due to its dynamical range, statistics of objects and number of time outputs. For this reason, we deploy a public repository of the available data, accessible via web portal at http://data.cineca.it. ", "introduction": "\\label{intr} Simulating the evolution of Cosmological Large Scale Structures of the Universe is a challenging task. In the last thirty years different numerical techniques were designed to follow the dynamics of the most important matter/energy components of the Universe: Dark Matter (DM), baryonic matter, and dark energy. In order to account for the great complexity and for the number of details provided by real cluster observations, a number of physical processes in addition to gravitational collapse and non-radiative hydro-dynamics have been implemented in many numerical works in the last few years: radiative gas processes, magnetic fields, star formations, AGN feedback, Cosmic Rays, turbulence, etc.(e.g. Dolag et al.2008; Borgani \\& Kravtsov 2009, and references therein, for a recent review). At present, two main numerical approaches are massively applied to cosmological numerical simulations: Lagrangian methods, which sample both the DM and the gas properties using point-like fluid elements, usually regarded as particles (e.g. Smoothed Particles Hydrodynamics codes, SPH) and Eulerian methods, which reconstruct the gas properties with a discrete space sampling with regular or adaptive meshes and model the Dark Matter properties with a Particle Mesh approach (see Dolag et al.2008 and references therein for a modern review). \\begin{figure*} \\begin{center} \\includegraphics[width=0.45\\textwidth]{E1_lev_cont.ps} \\includegraphics[width=0.45\\textwidth]{E18_lev_cont.ps} \\caption{The hierarchy of refinement levels in our runs. Color maps: level of mesh refinement for slices through the center of cluster E1 (left panel) and E18A (right panel) at $z=0$, from level=0 ($\\Delta=200kpc/h$, in black color) to level=3 ($\\Delta=25kpc/h$, in white color); the contour map shows the gas temperature distribution within the same region (the contours are equally spaced in $\\Delta log(T) \\approx 0.5$). The side of both images is $\\approx 14Mpc/h$.} \\label{fig:levels} \\end{center} \\end{figure*} \\bigskip High resolution, AMR simulations (such as the ENZO simulations presented in this paper) can provide an accurate representation of the cosmic gas dynamics in galaxy clusters, achieving a very large dynamical range. Recent works have shown that the adoption of proper mesh refinement criteria allows to study also the details of chaotic motions in the ICM (e.g. Iapichino \\& Niemeyer 2008, Vazza et al.2009; Maier et al.2009; Vazza, Gheller \\& Brunetti 2010; Paul et al.2010) Vazza et al.(2009, hereafter Va09) recently focused on the re-simulation of galaxy clusters by employing a new mesh refinement criterion, which couples the ``standard'' refinement criteria based on large gas or DM over-densities, to the mesh refinement criterion based on cell to cell 1--D jumps of the velocity field. In Va09 and Vazza, Gheller \\& Brunetti (2010, hereafter VGB10) we showed that the extra-refinement on 1--D velocity jumps opportunely increases the number of resolution elements across the ICM volume, allowing us to achieve a better spectral and morphological representation of chaotic motions in the ICM. Furthermore it reduces the artificial dampening of mixing motions due to the effect of the coarse resolution. Since the above works were focused on the re-simulation of a few intermediate mass systems (e.g. $M<3 \\cdot 10^{14}M_{\\odot}$), it is interesting now to extend the same method to a larger sample of higher mass clusters. Here we present the first results obtained analyzing 20 galaxy clusters, with total masses in the range $6 \\cdot 10^{14}M_{\\odot} \\leq M_{vir} \\leq 2 \\cdot 10^{15}M_{\\odot}$, obtained with the above techniques and designed to reach very high {\\it spatial} resolution around both DM/gas clumps, shocks and turbulent motions. Such rich sample accounts for objects of very different dynamical history and it is characterized by a large dynamical range ($N_{AMR}\\sim 500^{3}$, where $N_{AMR}$ is the number of grid elements at the maximum mesh refinement level) within the clusters volume. This allows us to study a broad variety of multi-scale phenomena associated to cluster growth and evolution. The paper is organized as follows: in Section \\ref{sec:methods} we present the clusters sample, the numerical techniques adopted and the archiving procedure for the data sample; in Section \\ref{subsec:scaling} we present the integrated (e.g. scaling laws) properties of our clusters, in Section \\ref{subsec:profiles} we present the and radial properties of gas density, gas temperature and gas entropy for all clusters in the sample. In Section \\ref{subsec:shocks} we characterize shock waves within the clusters and give estimates on the energy level of injected Cosmic Rays particles. The discussion and the conclusions are reported in Section \\ref{sec:conclusions}. In the Appendix, we report consistency tests for the adopted re-ionization scheme (Sect.\\ref{sec:appendix1}), and present a visual inspection of all clusters of the sample (Sect.\\ref{sec:appendix2}). \\begin{table} \\label{tab:clusters} \\caption{Main characteristics of the simulated clusters at $z=0$. Column 1: identification number; 2: total virial mass ($M_{vir}=M_{\\rm DM}+M_{gas}$); 3: virial radius ($R_{v}$); 4:$X=E_{k}/E_{tot}$ ratio inside $R_{v}$; 5:dynamical classification: RE=relaxing, ME=merging or MM=major merger (with approximate redshift of the last merger event).} \\begin{tabular}{c|c|c|c|c|c|c|c} ID & $M_{vir}$ & $R_{v}$ & $X$ & note \\\\ & [$10^{15}M_{\\odot}$] & [$Mpc$] & [$E_{kin}/E_{tot}$] & \\\\ \\hline E1 & 1.12 & 2.67 & 0.43 & MM(0.1)\\\\ E2 & 1.12 & 2.73 & 0.47 & ME \\\\ E3A & 1.38 & 2.82 & 0.43 & MM(0.2) \\\\ E3B & 0.76 & 2.31 & 0.55 & ME \\\\ E4 & 1.36 & 2.80 & 0.44 & MM(0.5)\\\\ E5A & 0.86 & 2.39 & 0.47 & ME \\\\ E5B & 0.66 & 2.18 & 0.75 & ME \\\\ E7 & 0.65 & 2.19 & 0.45 & ME \\\\ E11 & 1.25 & 2.72 & 0.40 & MM(0.6)\\\\ E14 & 1.00 & 2.60 & 0.23 & RE\\\\ E15A & 1.01 & 2.63 & 0.85 & ME\\\\ E15B & 0.80 & 2.36 & 0.33 & RE\\\\ E16A & 1.92 & 3.14 & 0.36 & RE \\\\ E16B & 1.90 & 3.14 & 0.67 & MM(0.2) \\\\ E18A & 1.91 & 3.14 & 0.37 & MM(0.8) \\\\ E18B & 1.37 & 2.80 & 0.34 & MM(0.5)\\\\ E18C & 0.60 & 2.08 & 0.55 & MM(0.3) \\\\ E21 & 0.68 & 2.18 & 0.40 & RE\\\\ E26 & 0.74 & 2.27 & 0.29 & MM(0.1)\\\\ E62 & 1.00 & 2.50 & 0.63 & MM(0.9) \\\\ \\end{tabular} \\label{tab:char} \\end{table} ", "conclusions": "\\label{sec:conclusions} In this paper, we presented a sample of 20 massive galaxy clusters in the range of total masses $6 \\cdot 10^{14} M_{\\odot} \\leq M \\leq 2 \\cdot 10^{15} M_{\\odot}$, extracted from large scale cosmological simulations and re-simulated with high mass resolution for the DM particles and high spatial resolution for the gas component, up to $\\sim 2-3 R_{v}$ from their centers. We used the ENZO code with Adaptive Mesh Refinement, using a refinement criterion based on gas/DM over-density and 1--D jumps in the velocity field (as in Vazza et al.2009). With this approach, we obtained a statistical sample with unprecedentedly large dynamical range within the virial volume of massive galaxy clusters, which can be used to study in detail the thermal properties, accretion phenomena and chaotic processes in the ICM over $2-3$ decades in spatial scale, for each cluster. We presented the first exploratory statistic study of this sample, showing the properties of gas density, gas temperature, gas entropy and baryon fraction for all clusters in our sample and the radial profiles in the range $0.01 \\leq r/R_{v} \\leq 3$ (Sections \\ref{subsec:scaling}-\\ref{subsec:profiles}). The reported trends are in line with previous studies that used complementary numerical techniques (e.g. SPH or standard AMR simulations), however they make possible to considerably extend the possibility of performing these measurements at much larger radii, thanks to the high spatial resolution in our simulations. The additional mesh refinement scheme adopted in this work (based on Vazza et al.2009) is explicitly designed to focus also on shock features and chaotic motions leading to significant 1--D jumps in the velocity field. This allowed us to characterize the morphologies, frequency and energy distributions of shock waves in these massive systems (Sec.\\ref{subsec:shocks} with unprecedented resolution and to estimate their relative efficiency in accelerating Cosmic Rays particles, by adopting two reference model of diffusive shock acceleration (Kang \\& Jones 2002,07). In agreement with previous studies based on much lower resolution we confirm that the distribution of shock energy flux inside clusters is extremely steep ($\\alpha_{th} \\approx - 4 \\div 5$, with $f_{th}(M)M \\propto M^{\\alpha_{th}}$), that the peak of the thermalisation at shocks is located at $M \\approx 2$, and that the average Mach number inside clusters is small, $M \\sim 1.5$. Only two clusters over 20 are interested by strong shocks inside $R_{v}/2$ (at $z=0$), with $M \\sim 2.7 $ and $\\sim 3.5$ respectively. The rarity of strong shocks found for $r 20 \\sim 30$ \\cite{car88,mat03}, where $M_A$ is the Alfv\\'{e}n Mach number and $\\beta_e = 8 \\pi n T_e / B^2$ the electron plasma beta. If the above condition is satisfied, rapid electron heating with saturation temperature $\\propto M^2_A$ occurs and the ion acoustic instability sets in leading to further electron heating \\cite{pap88, car88}. This condition is easily satisfied in SNR shocks, although it is usually not in the earth's bow shock. The two step instabilities were again shed light on after Ref. 9 by utilizing electromagnetic full PIC (particle-in-cell) simulations. They improved understandings of the process not only in terms of electron heating but also in terms of a production process of non-thermal electrons, and promoted a number of subsequent simulation studies \\cite{hos02,shi04,die04,shi05,mcc05,die06, ama07,ohi07,ohi08,ume08,die09,ama09a,ama09b}. As a result, understandings of electron heating or acceleration processes initiated by the BI in the context of the shock physics have been extensively developed. On the other hand, it has been well-known that various microinstabilities are possible to get excited also in relatively low Mach number supercritical shocks \\cite{wu84,pap85}. Since most of the microinstabilities are inseparable with electron dynamics, their nonlinear evolutions in self-consistently reproduced shock structures in PIC simulations are studied only recently \\citep{sch03,sch04,mus06,mat06a} except for Refs. 29-31. It is revealed that MTSI (modified two-stream instability) becomes dominant in the foot of the earth's bow shock or interplanetary shocks at $\\sim1$ AU \\citep{sch03,sch04,mat06a}. Simulation studies on nonlinear evolutions of the MTSI in periodic systems were extensively studied in the 1970s and 1980s \\cite{mcb72,ott72,tan83}, although these simulations imposed some strong assumptions that the code was electrostatic \\cite{mcb72,ott72} and unmagnetized ions with small ion-to-electron mass ratio were used \\cite{mcb72,ott72,tan83}. Recently, some developement without these assumptions are made \\cite{mat03,mat06b}. Ref. 8 showed in a one-dimensional simulation with realistic mass ratio that long time evolution of the MTSI results in lower cascade of wave spectrum and associated electron heating. Furthermore, it is indicated in two-dimensional simulation that the MTSI can finally survive even if other possible instability like electron cyclotron drift instability is present \\cite{mat06b}. All these simulation studies show electron heating or acceleration in the long time evolution of the system. However, as mentioned already, electron heating in the earth's bow shock is seldom observed in-situ. Such a discripancy may arise because of lack of systematic estimate of electron heating rate in simulation studies with more realistic parameters. For example, all of the above simulation studies on the MTSI assume very low values of squared ratio of electron plasma to cyclotron frequencies, $\\tau = \\omega^2_{pe} / \\Omega^2_e\\leq 10$. Importance of $\\tau$ in association with electron heating in the long time evolution of the MTSI is unresolved, while $\\tau$ is known to be a crucial parameter controlling nonlinear electron heating and acceleration processes in the BI dominant system \\cite{shi04}. Electron plasma beta is also rather small in most of the previous studies ($\\beta_e \\leq 0.1$). This may exaggerate effects of electron trapping in the nonlinear stage of the MTSI \\cite{mcb72,sch03,mat03,sch04,mat06a}. The electron plasma beta in the solar wind is usually several times higher so that importance of the trapping process may be blurred. Hence, in realistic situations expected to be achieved in the typical solar wind conditions an alternative electron heating mechanism and its efficiency are ought to be considered. In this paper, first, electron heating through the MTSI in transition regions of high Mach number quasi-perpendicular shocks is reconsidered in section II. Here, quasilinear diffusion is assumed as an alternative electron heating process. The analysis assumes a broad wave spectrum and small wave amplitudes which are likely to be satisfied in the past simulation studies mentioned above \\cite{mcb72,mat03}. An additional assumption that relative phases among different wave modes are random is also imposed, while this point is currently not evident. Under these assumptions, taking second order velocity moments of the so-called quasilinear equations provides evolution equations of electron kinetic energies. These equations are numerically integrated by imposing the quasilinear assumptions in every instantaneous time steps to obtain saturation kinetic energies as a fuction of the Mach number of shocks. This approach allows to discuss long time evolution of macroscopic quantities including the electron temperature without any restrictions for parameters. It is intriguing whether the Mach number dependence of electron heating throug the MTSI in relatively low Mach number regime is different from that through the BI in extremely high Mach number regime. If that is the case, one may expect presence of a critical Mach number above which electron heating rate runs up due to the switching of the dominant microinstability \\cite{hos02}. In section III, the results are compared with 1D PIC simulations in which the MTSI is dominantly generated in the foot and effective electron temperature just behind of the shock is measured. Furthermore, post-shock effective electron temperature is measured also in cases that the BI gets excited in the foot. In those cases a rough estimate of the electron temperature is given by using the well-known trapping theory. Finally, a summary and discussions including estimate of the critical Mach number are given in section IV. ", "conclusions": "In the previous section it was confirmed that the extended quasilinear analysis for the MTSI and the trapping analysis for the BI give consistent results with the PIC simulations. Here, let us compare the results of these two analysis. The thick gray dashed line in Fig.\\ref{fig03} denotes $\\beta^{eff}_{ex}$ obtained from eq.(\\ref{sat-ehole}) for $\\mu = 1836, \\alpha = 0.5$. This line intersects at $M_A \\sim 24$ with the dotted line with triangles which is $\\beta_{e\\parallel,sat}$ for the MTSI with $\\beta_{0e\\parallel} / \\beta_{0i} = 1/3$. The gray thick solid line based on eq.(\\ref{sat_ql}) likely to intersect with the gray thick dashed line at $M_A \\sim 45$. Depending on parameters, an actual intersection may occur in the range of $20 < M_A < 50$. If this is written as $M^*_A$, the dominant electron heating process possibly switches from the MTSI for $M_A < M^*_A$ to the BI for $M_A > M^*_A$. For relatively low Mach numbers like in the earth's bow shock, $\\beta_{e\\parallel,sat}$ is more or less constant with $10^{0 \\sim 1}$ because the MTSI becomes dominant. If $\\beta_{e,sat}=5$ and $\\tau = 10^4$ are assumed, corresponding electron temperature is $\\sim 100$eV which is consistent with a typical temperature observed downstream of the earth's bow shock. However, addiabatic heating due to increases of magnetic field and cross shock potential also results in the downstream electron temperature similar to this value. This may be the reason why remarkable electron heating was seldom observed in near earth shocks in the past. But one may capture the nonadiabatic parallel electron heating if upstream electron beta is low enough ($\\beta_{e0} < 0.1$), since eq.(\\ref{sat_ql}) is still valid for rather low beta cases where the trapping effects become essential as mentioned in section II. Furthermore, as shown in the upper right panel of Fig.\\ref{fig04}, local temperature anisotropy, $T_{e\\parallel}/T_{e\\perp} > 1$, in a transition region may be observed as a result of strong MTSI in some parameter regimes. On the other hand, the effective electron temperature significantly increases with being proportinal to $M^2_A$ in $M_A > M^*_A$. Although this is consistent with the past studies \\cite{pap88,car88,shi05}, the saturation electron temperature estimated in this study seems to be rather smaller. This should be because of that the heating process in highly nonlinear stage including the second step ion acoustic instability is neglected. Electron temperature seen in Ref. 10 seems to be about one order higher than the estimate given by eq.(\\ref{sat-ehole}). This implies that an actual $M^*_A$ may appear at a little smaller value. Consistency of the results of the extended quasilinear analysis and the PIC simulations for the MTSI dominant cases indicates that the resonant wave-particle interactions are essential in electron heating through the MTSI. In other wards, nonresonant wave-particle interactions which have been neglected in the analysis may not be effective. It should also be noted that using unrealistically small $\\tau$ in PIC simulations is justified as far as the MTSI is concerned because of the weak dependence on $\\tau$ (bottom panel of Fig.2) in contrast to the BI dominant systems. The transition at $M_A = M^*_A$ might not be so drastic in a realistic case. In the analyses presented here, all other possible candidates of microinstabilities have been neglected. For example, ECDI (electron cyclotron drift instability) is one candidate \\cite{mus06} and it might become important around this critical Mach number, although there are some negative indications for the ECDI to become dominant. For instance, it is known that the saturation level of the ECDI is usually not so high \\cite{wu84}. In addition, if the ECDI gets excited simultaneously with the MTSI, the MTSI becomes dominant for wide parameter range \\cite{mat06b}. Furthermore, it is shown that the ECDI can be important only when $\\tau$ is of the order of unity \\cite{shi04}. Hence, the Mach number regime where the ECDI is dominant seems not to be so wide, but we do not remove the possibility that the ECDI gives some contributions to electron heating around $M^*_A$. Because the ECDI is insensitive to the electron Landau damping which may strongly affect to the BI in this Mach number regime. Further, it is known that the MTSI cannot get excited in extremely high Mach number shocks like $M_A > 30 \\sim 40$ \\cite{mat03}. There are also other possible microinstabilities in higher dimensional cases \\cite{wu84}. Contributions from them should be carefully estimated. Some of them have been discussed recently by performing two dimensional PIC simulations \\cite{ohi08,ume08,ama09a}. The multi-dimansionality may give another important contributions to electron heating. Ref. 42 pointed out that electrons are accelerated in the so-called rippled structure which is the ion scale structure along the shock surface. Nevertheless, what should be emphasized here is that the Mach number dependence of the effective electron temperature is systematically different between high and low Mach number regimes. In the present extended quasilinear analysis the damping of the wave energy in the late stage (e.g., $\\Omega_i t > 0.45$ in Fig.\\ref{fig01}) may be an artifact. Because of the assumption that the distribution function is always Maxwellian, the so-called plateau of the distribution function is never produced in this system. The field energy might saturate earlier at a certain level if the plateau is produced. In this regard, it might be better that the saturation levels are defined as the values at the time when the field energy becomes the maximum. However, the resultant saturation levels based on the two definitions, i.e., the saturation temperatures estimated at times corresponding to the maximum and the later sufficiently small field energies, are not so much different from each other. Although it is possible to solve the quasilinear equation, eq.(\\ref{qleq}), directly as done in Refs. 39 and 40, that is the future work. The explorations of the inner heliosphere will be underway through the BepiColombo mission. The perihelion point of the mercury is about 0.3AU where quite high velocity IPSs which have not been decerelated are expected to be observed. Ref. 43 estimated the propagation speed of one of the IPSs observed in August 1972 and concluded that it approached to $\\sim 2000$ km/s at 0.3AU from the sun. More optimistic estimate for the same IPS is given in Ref. 44 as $\\sim 4000$ km/s at 0.3AU. These results imply that the Mach number of this IPS approached to several tens. Therefore, the transition of the electron heating efficiency at $M^*_A$, if present, will be possibly observed in situ \\cite{hos02}." }, "1003/1003.2849_arXiv.txt": { "abstract": "We present a time-dependent and spatially inhomogeneous solution that interpolates the extremal Reissner-Nordstr\\\"om (RN) black hole and the Friedmann-Lema\\^itre-Robertson-Walker (FLRW) universe with arbitrary power-law expansion. It is an exact solution of the $D$-dimensional Einstein-``Maxwell''-dilaton system, where two Abelian gauge fields couple to the dilaton with different coupling constants, and the dilaton field has a Liouville-type exponential potential. It is shown that the system satisfies the weak energy condition. The solution involves two harmonic functions on a $(D-1)$-dimensional Ricci-flat base space. In the case where the harmonics have a single-point source on the Euclidean space, we find that the spacetime describes a spherically symmetric charged black hole in the FLRW universe, which is characterized by three parameters: the steepness parameter of the dilaton potential $n_T$, the U$(1)$ charge $Q$, and the ``nonextremality'' $\\tau $. In contrast with the extremal RN solution, the spacetime admits a nondegenerate Killing horizon unless these parameters are finely tuned. The global spacetime structures are discussed in detail. ", "introduction": "Black holes could have formed in the early stage of the universe as a consequence of primordial density fluctuation~\\cite{PBH1,PBH2}. Since these primordial black holes (PBHs) are of the order of the horizon scale mass, they span a considerably broad mass spectrum ranging from the Planck mass ($10^{-5}$g) to $10^5 M_{\\odot}$, or even much larger. Hence, their presence could leave diverse physical imprints throughout the cosmic history (see~e.g.,~\\cite{PBHs} and references therein). The PBHs with $M>10^{15}$g survive until the present epoch, so that they are plausible candidates of cold dark matter, and likely origin of supermassive black holes and/or the sources of gravitational waves. The PBHs with mass $M\\sim 10^{15}$g are now evaporating via Hawking evaporation~\\cite{Hawking1974}. Such PBHs emit quanta of order 100MeV, which could also contribute to the cosmological $\\gamma $-ray background and generate $\\gamma $-ray bursts. While, the PBHs with mass smaller than $10^{15}$g that have completely evaporated during a first second of big bang could generate large amount of entropy of the universe. Thus, the number density of small PBHs could place strong constraints on the big-bang nucleosynthesis. In these contexts, studying the formation and the evaporation process of black holes in the expanding universe are of great significance as a probe of the early universe, high energy physics, and quantum phase of gravity. Apart from the above astrophysical interests, a study of black holes in the expanding universe has been a principal stirring subject in general relativity. Black holes have played a central role in general relativity, since they encode the essential characteristics of the gravity theory, viz, the nonlinearity and highly curved spacetime. Various kinds of physical and geometrical properties of black holes have been clarified thus far. The most enduring achievements among these is the uniqueness theorem of black holes~\\cite{Carter,uniqueness}, according to which isolated black holes in equilibrium states--these states occur if sufficiently long time passed after the gravitational collapse of massive stars--necessarily belong to the Kerr family. This means that stationary black holes are completely characterized by two conserved charges (mass and angular momentum) without any additional ``hairs.'' By virtue of this theorem, stellar-sized black holes with futile ambient sources are well approximated by Kerr black holes. To the contrary, in the non-isolated and dynamical background, a variety of black holes equipped with much richer properties are expected. However, when one attempts to obtain the black hole spacetime in the dynamical background, a serious difficulty arises. If a black hole is put on the homogeneous and isotropic FLRW universe on which the standard cosmological scenario lays the foundation, the background universe will become inhomogeneous, and at the same time the black hole will continue to grow and/or deform by swallowing ambient matters. Hence, the fact that the spacetime is time-evolving and spatially inhomogeneous enforces us to solve nonlinear partial differential equations for the geometry, as well as for the matter fields. The exact black-hole solutions in the FLRW universe found hitherto have enjoyed high degrees of symmetry. Among them are the Einstein-Straus model~\\cite{Einstein:1945id} (a black hole in the ``Swiss-Cheese Universe''), the Schwarzschild-de Sitter (SdS) and the Reissner-Nordstr\\\"om-de Sitter (RNdS) black hole. Since these solutions maintain the equilibrium due to the timelike Killing field, they are unlikely to capture the envisaged dynamical fact in realistic situations. In recent years, Sultana and Dyer have exploited the conformal technique and have constructed a dynamical black hole~\\cite{SultanaDyer} which is asymptotically Einstein-de Sitter universe and conformally related to the Schwarzschild solution. Unfortunately, the Sultana-Dyer solution suffers from violating energy conditions. Meanwhile, if one imposes a self-similarity (characterized by a homothetic Killing field), it has been proved that a spherically symmetric black hole cannot exist in the asymptotically decelerating FLRW universe~\\cite{HMC}. These examples illustrate the difficulty in constructing a regular black hole immersed in the FLRW universe with desirable physical properties. Recently, a time-dependent ``black hole candidate'' was obtained via the dimensional reduction of the dynamically intersecting branes in 11-dimensional (11D) supergravity~\\cite{MOU}. The solution appears to behave like a charged black hole for small radii, while for large radii it approaches to the FLRW universe filled with stiff matter fluid. In the previous paper~\\cite{MN}, we elucidated the detailed spacetime structure and established that the solution indeed describes a charged black hole that approaches to the flat FLRW universe. It is shown that the metric is an exact solution in Einstein-``Maxwell''-dilaton system with four kinds of U$(1)$-fields (three of them are degenerate) coupled to the dilaton, thereby the system satisfies the dominant energy condition. Although the BPS black hole plays the leading part in string theory, what is intriguing us is that the solution is {\\it not} extremal, i.e., the Hawking temperature does not vanish. Despite the little astrophysical concern because of the brane charges, the intersecting brane picture will enable us to obtain a variety of dynamical black holes with physical matter sources. In~\\cite{GMII}, Gibbons and one of the present authors found the ``black hole candidate'' that asymptotically looks like an FLRW universe with arbitrary power-law expansion by introducing an exponential potential. The solution found in~\\cite{GMII} reduces to the solution in~\\cite{MN,MOU} as a special case and is expected to describe a black hole in the expanding universe. As lessons from the McVittie solution~\\cite{Mcvittie1933,Nolan,Carrera:2009ve}, a great deal of care is needed in order to obtain the global spacetime structure of a time-evolving and spatially inhomogeneous solution.\\footnote{In contrast to the claims in~\\cite{Nolan}, it is recently argued that McVittie's solution includes a regular black-hole horizon if the positive cosmological constant dominates the universe at late time~\\cite{Kaloper:2010ec}. } In this paper, following the previous study~\\cite{MN} we intend to clarify the spacetime structure of their solution, which has been an open issue in~\\cite{GMII}. There are several motivations for studying black holes in the Einstein-``Maxwell''-dilaton system with an exponential potential. In the light of string theory, the dilaton and the form fields are elementary constituents in the theory, and the exponential potential of a dilaton naturally arises in various contexts: the excess central charge (the ``noncritical'' string theory)~\\cite{noncritical_string_theory}, the massive IIA supergravity~\\cite{massiveIIA}, the Kaluza-Klein compactification of a curved internal space~\\cite{Kaluza-Klein_compactification}, the vacuum expectation value of the four-form field strength~\\cite{VEV_form}, and so on. From the general relativistic point of view, these fields obey suitable energy conditions. Thus, it is interesting to see if the black hole exhibits a peculiar aspect under a realistic circumstance. In addition, since the exponential potential may drive the power-law inflation~\\cite{Lucchin:1984yf}, this could give a profound implication for the background geometry of PBHs during inflation. Furthermore, by generalizing the present solution to include multiple black holes one is able to discuss the collision and coalescence of black holes, providing a valuable arena to test the cosmic censorship conjecture~\\cite{Penrose:1969pc}. This paper constitutes as follows. In the ensuing section, we present a $D$-dimensional solution as a simple and honest generalization of the solution given in~\\cite{GMII}. We discuss the matter fields and singularities of the solution in Section~\\ref{sec:matter}, where the system is shown to satisfy the suitable energy conditions. Section~\\ref{sec:spacetime} develops our main argument on the spacetime structure. Analysis of the near-horizon geometry and null geodesic motions leads to the conclusion that the present spacetime can admit a regular event horizon with constant circumference radius, if the parameters of the solution are chosen appropriately. Conclusions and future outlooks are summarized in Section~\\ref{sec:summary}. We shall work in units $c=\\hbar =1$ with $\\kappa^2 =8\\pi G$ and follow the standard curvature conventions ${\\ma R^\\mu}_{\\nu\\rho\\sigma}V^\\nu:=2\\nabla_{[\\rho }\\nabla_{\\sigma ]}V^\\mu$, $\\ma R_{\\mu \\nu }:={\\mathcal R^{\\rho }}_{\\mu \\rho \\nu }$ and $\\ma R:=\\ma R^\\mu_\\mu $. Since the circumference radius will be denoted by $R$, the script notation is used for curvature tensors throughout the paper. ", "conclusions": "\\label{sec:summary} We have presented a family of solutions in a $D$-dimensional Einstein-``Maxwell''-dilaton system with a Liouville potential. In the single mass case, the solution interlies the extremal black hole and the FLRW universe with a power-law expansion. This solution reduces in the special cases to the extremal RN black hole ($n_T=0$), the nonextremal RNdS black hole ($n_S=0$) and the solution of~\\cite{MOU} derived from the intersecting M-branes ($D=4$ and $n_T=1$). The present spacetime is characterized by three parameters: the steepness of the dilaton potential $n_T$, the U$(1)$ charge $Q$ and the ``nonextremality'' (or the ratio of two energy densities at the horizon) $\\tau $. The parameter $n_T$ controls the expanding power $p$ of the background FLRW universe. For any choice of parameters, the system is shown to obey the weak energy condition. The primary aim of this paper is to obtain global structures of the solution. The spacetime can be grouped into the nine cases [Cases I-(i)--III-(iii)] summarized in Table~\\ref{Table}, according to the parameter values. Our consequence is visually captured in Figure~\\ref{PD}, where the 5-types of global structure are obtained. In the case where the background universe is accelerating ($n_T>2$) with large nonextremality parameter ($\\tau\\ge \\tau_{\\rm cr}$), the spacetime indeed describes a charged black hole in the FLRW universe undergoing an accelerated expansion [Case III-(i)]. We have also clarified the global structures for the marginally accelerating and decelerating cases. This has been an open issue in~\\cite{GMII}. The present solution displays various interesting features. At first glance, the solution appears to have a degenerate horizon. We have shown that this is not the case. Taking the near-horizon limit~(\\ref{NHlimit}), the horizon is not degenerate except for $\\tau=\\infty $ and $\\tau=\\tau_{\\rm cr}$. Amazingly, the event horizon of a black hole constitutes a Killing horizon. This means that the matter fields fail to accrete into the black hole, whereby the area of the black hole keeps constant. Hence, the present spacetime is not suitable for describing a growing black hole in FLRW universe. Instead, the solution preserves the analogue of the equilibrium state despite the time-dependence of the metric. The situation is closely analogous to the supersymmetric state. The black hole thermodynamics in the expanding universe is an interesting future work to be discussed. Since the event horizon is generated by a (partial) Killing field, the black hole has a nonvanishing surface gravity. However, it is far from obvious to which observers the physical temperature is assigned, since the surface gravity is sensitive to the normalization of a Killing field. This may be rephrased as what is the mass of the black hole. We hope to visit this issue in a separated paper. Another interesting issue to be explored in the time-dependent black hole spacetime is a black hole merger. Adding multiple point sources and reversing time backwards, the present metric turns out to give the analytic description of black hole coalescence, similar to the Kastor-Traschen solution~\\cite{KT,HH}. In the single-mass Kastor-Traschen spacetime (RNdS), an analogue of ``thermal equilibrium'' is realized due to the fact that the temperature of the event horizon and that of the cosmological horizon become the same. If this viewpoint continues to be valid in the present spacetime, what plays the role of a box containing the black hole which is thermal equilibrium with a bath of radiation? These are interesting questions to be explored." }, "1003/1003.3281_arXiv.txt": { "abstract": "We study models of late-time cosmic acceleration in terms of scalar-tensor theories generalized to include a certain class of non-linear derivative interaction of the scalar field. The non-linear effect suppress the scalar-mediated force at short distances to pass solar-system tests of gravity. It is found that the expansion history until today is almost indistinguishable from that of the $\\Lambda$CDM model or some (phantom) dark energy models, but the fate of the universe depends clearly on the model parameter. The growth index of matter density perturbations is computed to show that its past asymptotic value is given by 9/16, while the value today is as small as 0.4. ", "introduction": "Since the discovery of cosmic acceleration various possibilities have been explored to account for this mystery. Probably the most conservative possibility is that the accelerated expansion arises from a cosmological constant or some unknown dynamical field (e.g., a quintessence)~\\cite{Copeland:2006wr}. Instead of introducing a new component in ``$T_{\\mu\\nu }$'' one may alternatively consider that dark energy is {\\em geometric}, i.e., modification to general relativity (GR) at long distances is responsible for cosmic acceleration. The latter class of models includes scalar-tensor theories, $f(R)$ gravity~\\cite{f(R)}, and the Dvali-Gabadadze-Porrati (DGP) braneworld~\\cite{DGP, DDG}. A difficulty in modifying GR at long distances lies, however, on short distance scales; a new scalar gravitational degree of freedom, which commonly appears in geometric dark energy models, must be tamed carefully in order to pass the stringent tests of gravity in the solar system~\\cite{will, uzan-review}. Two main ways of hiding the scalar degree of freedom are known to exist. The first one is to make the scalar effectively massive in the vicinity of matter. This is called the chameleon mechanism~\\cite{chameleon} and is utilized in viable $f(R)$ models~\\cite{viablefr}. The second way is decoupling the scalar from matter in the vicinity of the matter sources. This is known as the Vainshtein mechanism~\\cite{Vs}, and the DGP braneworld offers a nice example that implements it by a non-linear self-interaction in the kinetic term. See also a recent paper~\\cite{symmetron} for yet another way of suppressing the scalar-mediated force. Inspired by the DGP braneworld, a class of scalar-tensor theories of gravity has been explored that enjoys self-screening of the scalar-mediated force in the vicinity of matter due to non-linear derivative interactions. Such a scalar degree of freedom is called the Galileon field in the original proposal~\\cite{G1}, because the equation of motion for the scalar $\\phi$ is invariant under $\\partial_\\mu\\phi\\to\\partial_\\mu\\phi+b_\\mu$ on Minkowski spacetime. The Galileon scalar-tensor theories have been covariantized and a general form of the Lagrangian in curved spacetime has been considered in~\\cite{G2}\\footnote{Upon covariantization the original symmetry under $\\partial_\\mu\\phi\\to\\partial_\\mu\\phi+b_\\mu$ is lost. However, covariant Galileon Lagrangians are uniquely determined by requiring that the equations of motion have only second derivatives of the fields. Thus, the name ``Galileon'' may be misleading and inappropriate. Nevertheless, we use the name in this paper because no other name is as good as ``Galileon.''} (see also~\\cite{GE1, GE2}). The resultant equations of motion have only second derivatives of the fields, and hence the Galileon theories are in some sense similar to Lovelock gravity~\\cite{Lovelock}. Cosmology based on a Galileon field has been studied in Refs.~\\cite{CK, SK, KTS}. In this paper, we study aspects of Galileon cosmology in terms of a certain wider class of the Lagrangian than previously studied. The action we consider is of the form \\begin{eqnarray} S&=&\\int\\D^4x\\sqrt{-g}\\biggl[ \\phi R -\\frac{\\omega(\\phi)}{\\phi}(\\nabla\\phi)^2 \\cr&&\\qquad\\qquad\\qquad +\\frac{\\lambda^2(\\phi)}{\\phi^2}\\Box\\phi(\\nabla\\phi)^2 +{\\cal L}_{\\rm m} \\biggr],\\label{action} \\end{eqnarray} where $\\phi$ is the Galileon field and the coupling $\\lambda(\\phi)$ has dimension of length. Matter (represented by the Lagrangian ${\\cal L}_{\\rm m}$) is universally coupled to gravity in the Jordan frame. The gravitational field equations derived from~(\\ref{action}) are \\begin{eqnarray} \\phi G_{\\mu\\nu}&=&\\frac{1}{2}T_{\\mu\\nu}+\\nabla_\\mu\\nabla_\\nu\\phi-g_{\\mu\\nu}\\Box\\phi +\\frac{\\omega}{\\phi}\\biggl[\\nabla_\\mu\\phi\\nabla_\\nu\\phi \\cr&& -\\frac{1}{2}g_{\\mu\\nu}(\\nabla\\phi)^2\\biggr] -\\frac{1}{2}g_{\\mu\\nu}\\nabla_\\lambda\\left[\\frac{\\lambda^2}{\\phi^2}(\\nabla\\phi)^2\\right]\\nabla^\\lambda\\phi \\cr&& +\\nabla_{(\\mu}\\left[\\frac{\\lambda^2}{\\phi^2}(\\nabla\\phi)^2\\right]\\nabla_{\\nu)}\\phi -\\frac{\\lambda^2}{\\phi^2}\\nabla_\\mu\\phi\\nabla_\\nu\\phi\\Box\\phi, \\end{eqnarray} and the equation of motion for the Galileon field is given by \\begin{eqnarray} &&R+\\omega\\left[\\frac{2\\Box\\phi}{\\phi}-\\frac{(\\nabla\\phi)^2}{\\phi^2}\\right] +\\frac{\\omega'}{\\phi}(\\nabla\\phi)^2 \\cr&&\\quad +\\frac{2\\lambda^2}{\\phi^2} \\left[\\nabla_\\mu\\nabla_\\nu\\phi\\nabla^\\mu\\nabla^\\nu\\phi -(\\Box\\phi)^2+R_{\\mu\\nu}\\nabla^\\mu\\phi\\nabla^\\nu\\phi \\right] \\cr&&\\quad +4\\left(\\frac{\\lambda^2}{\\phi^2}\\right)'\\nabla_\\mu\\phi\\nabla_\\nu\\phi\\nabla^\\mu\\nabla^\\nu\\phi +\\left(\\frac{\\lambda^2}{\\phi^2}\\right)''(\\nabla\\phi)^2(\\nabla\\phi)^2 \\cr&&\\quad =0, \\end{eqnarray} where a prime denotes differentiation with respect to $\\phi$. The above modified gravity theory may be regarded as the generalization of Brans-Dicke gravity, but the cubic derivative interaction, $[\\lambda^2/\\phi^2]\\Box\\phi(\\nabla\\phi)^2$, plays a key role in manifesting the Vainshtein mechanism. The Vainshtein radius, below which the Galileon-mediated force is screened, is evaluated as $r_V\\sim [r_g\\lambda^2/(1+2\\omega/3)^2]^{1/3}$, where $r_g$ is the Schwarzschild radius of the matter source. In what follows we shall consider the particular case with $\\omega=$ const and $\\lambda^2\\propto \\phi^\\alpha$, where $\\alpha$ may be positive and negative. The self-accelerating de Sitter universe of~\\cite{SK} corresponds to the subclass $\\alpha=0$. We are going to study the modified gravity theory with general $\\alpha$ as a possible alternative to a cosmological constant or dynamical dark energy. Accelerating cosmologies from a scalar field non-minimally coupled to gravity {\\em without} the non-linear derivative interaction have been studied extensively in~\\cite{extquint, Reconst}. Since the background expansion history in successful modified gravity is by definition almost identical to that of the standard $\\Lambda$CDM model or other dynamical dark energy models, it is important to study the growth history of perturbations as a tool to distinguish modified gravity from models with a cosmological constant or dark energy. We compute the growth index of matter density perturbations~\\cite{Peeb,WS} in Galileon cosmology, which is known to be a powerful discriminant among models of cosmic acceleration~\\cite{Linder1,Linder2, P-G}. The paper is organized as follows. In the next section we study the background evolution of Galileon cosmology derived from the above action. Then, in Sec.~III, we discuss the growth density perturbations in Galileon cosmology, paying particular attention to the growth index. The final section is devoted to conclusions. ", "conclusions": "In this paper, we have studied a class of Galileon cosmology, i.e., cosmology in the Brans-Dicke theory with the non-linear derivative interaction of the form $[\\lambda^2(\\phi)/\\phi^2](\\nabla\\phi)^2\\Box\\phi$ with $\\lambda^2(\\phi)\\propto \\phi^{\\alpha}$. Thanks to this term in the Lagrangian, the Vainshtein mechanism works and the scalar-mediated force is screened near matter sources. The Vainshtein mechanism operates in the cosmological context as well, leaving the cosmic expansion history indistinguishable from the usual matter-dominated universe. Moreover, we have illustrated that modification to gravity due to the Galileon field can be responsible for the late-time acceleration of the universe, closely mimicking the $\\Lambda$CDM model or rather phantom dark energy models. We have clarified that the late-time asymptotics depends on $\\alpha$: a future big rip singularity occurs for $\\alpha<0$, while the phantom-like behavior is only temporal and we eventually have $w_{\\rm eff}>-1$ for $\\alpha>0$. The result here is the generalization of the self-accelerating de Sitter universe found for $\\alpha=0$ in~\\cite{SK}. We have also investigated the growth history of density perturbations in Galileon cosmology. The asymptotic growth index was obtained analytically as $\\gamma_\\infty=9/16$, which was confirmed numerically. The growth index today was found to be $\\gamma\\sim 0.4$. The behavior of the growth index is thus a clear discriminant among Galileon scalar-tensor gravity, the $\\Lambda$CDM model, and the DGP braneworld." }, "1003/1003.4718_arXiv.txt": { "abstract": "We perform two-dimensional MHD simulations of the flux emergence from the solar convection zone to the corona. The flux sheet is initially located moderately deep in the adiabatically stratified convection zone ($-20,000\\ {\\rm km}$) and is perturbed to trigger the Parker instability. The flux rises through the solar interior due to the magnetic buoyancy, but suffers a gradual deceleration and a flattening in the middle of the way to the surface since the plasma piled on the emerging loop cannot pass through the convectively stable photosphere. As the magnetic pressure gradient enhances, the flux becomes locally unstable to the Parker instability so that the further evolution to the corona occurs. The second-step nonlinear emergence is well described by the expansion law by \\citet{shiba89}. To investigate the condition for this `two-step emergence' model, we vary the initial field strength and the total flux. When the initial field is too strong, the flux exhibits the emergence to the corona without a deceleration at the surface and reveals an unrealistically strong flux density at each footpoint of the coronal loop, while the flux fragments within the convection zone or cannot pass through the surface when the initial field is too weak. The condition for the `two-step emergence' is found to be $10^{21}-10^{22}\\ {\\rm Mx}$ with $10^{4}\\ {\\rm G}$ at $z=-20,000\\ {\\rm km}$. We have some discussions in connection with the recent observations and the results of the thin-flux-tube model. ", "introduction": "} Solar active regions are generally thought to be the consequence of the flux emergence, i.e., dynamo-generated magnetic fluxes risen from the deep convection zone due to the magnetic buoyancy force \\citep{par55}. Observations indicate that the flux tube should be in a coherent form and strong enough so as not to be disintegrated by the turbulent motion during its emergence through the convection zone. For generating such a strong flux, the sub-adiabatically stratified overshoot region at the bottom of the convection zone has been suggested to be the suitable place \\citep{par75}. Therefore, a flux emergence should be understood as a whole process from the base of the convective layer to the upper atmosphere through the surface. In last two decades, various series of numerical experiments have been carried out to investigate the physics of the flux emergence. For the local evolution above the surface, the pioneering work was done by \\citet{shiba89}, who performed two-dimensional (2D) magnetohydrodynamic (MHD) simulations of flux emergence through the undular mode of magnetic buoyancy instability ($\\mbox{\\boldmath $k$}\\parallel \\mbox{\\boldmath $B$}$, where $\\mbox{\\boldmath $k$}$ and $\\mbox{\\boldmath $B$}$ denote the wavenumber and the initial magnetic field vector, respectively; the Parker instability) to reproduce some dynamical features such as the rise motion of an arch filament system and downflows along the magnetic field lines. Since then, the evolution of an emerging flux and the interaction with pre-existing coronal fields have been studied by 2D and 3D simulations \\citep{noza92,matsu93,yoko95,yoko96,fan01b,che08}. \\citet{noza92} performed the emergence from the convectively unstable solar interior (the convective-Parker instability), while \\citet{yoko95,yoko96} studied the reconnection between the expanding loop and the pre-existing fields in the corona and the subsequent formation of X-ray jets. Three-dimensional calculations by \\citet{matsu93} were performed for the studies of the interchange ($\\mbox{\\boldmath $k$}\\perp \\mbox{\\boldmath $B$}$) and quasi-interchange ($\\mbox{\\boldmath $k$}\\parallel\\mbox{\\boldmath $B$}$ with $kH_{\\rm ph}\\ll 1$, where $H_{\\rm ph}$ is the photospheric pressure scale height) mode instabilities. \\citet{fan01b} compared her 3D simulation results with observed features of newly emerged active regions. \\citet{che08} found that the numerical modeling of emerging flux regions by 3D radiative MHD simulations exhibits photospheric characteristics that are comparable with the observations from the {\\it Hinode}/Solar Optical Telescope (SOT). These simulations assumed that the initial flux is embedded just below the photosphere ($\\ga -2000\\ {\\rm km}$) a priori as a horizontal sheet or a twisted tube. It is because their experiments mainly aimed to clarify the local behaviors in the solar atmosphere. Such an initial structure, however, is not obvious since the initial form depends on its history of pushing through the convection zone. Therefore, a numerical experiment including the evolution from a substantial depth has been needed. The simulations focusing on the flux emergence within the relatively deep solar interior have been done by using the thin-flux-tube approximation \\citep{spr81,dsil93,cal95}. One of the most important conclusions obtained from the various thin-flux-tube simulations is that rising tubes with small magnetic flux (below 10$^{21}$ Mx for 10$^{4}$ G at the base) cannot reach the photosphere because the apices of the loops loose magnetic fields and subsequently `explode' \\citep{mor95}. By assuming the anelastic approximation \\citep{gou69,lan99}, \\citet{fan01a} computed the evolution of the 3D undulatory instability of the horizontal magnetic flux layer, which formed arch-like magnetic tubes with downflows from their apices to the troughs. Note that both types of approximations (thin-flux-tube and anelastic) are not applicable in the upper convection zone ($\\ga -30\\ {\\rm Mm}$) where the diameter of the flux tube exceeds the local pressure scale height and the flow velocity becomes close to the sound speed. \\citet{abb03} calculated a flux emergence by connecting the anelastic MHD convective layer and the fully compressible MHD solar atmosphere from the photosphere to the low corona. However, their full MHD atmosphere did not include the upper convection zone. To investigate the detailed behavior of the emerging flux from the convection zone to the atmosphere, we have to deal with the full MHD numerical box including the convective layer, the photosphere/chromosphere, and the corona. The dynamical behavior of an emerging flux including both the interior and the upper atmosphere is not yet clear, but is thought to obey the following picture, which we name `two-step emergence' model \\citep[cf.][]{matsu93,maga01}. Magnetic fluxes emerged from the bottom of the convection zone are depressed and decelerated by the sub-adiabatic photosphere and extended horizontally around the photosphere/chromosphere. Meanwhile, fluxes are still transported from below to enhance the magnetic pressure gradient. Finally, the fluxes above the photosphere become unstable to the magnetic buoyancy instability so that the further evolution into the corona occurs. By the development of supercomputers, full MHD simulations in the convection zone without approximations have come to be realizable. The purpose of this work is to investigate the `two-step emergence' model numerically. Several experiments have confirmed this `two-step' model. \\citet{maga01} studied the emergence of the magnetic flux tube from the convection zone by means of 2.5-dimensional MHD simulations focused on the cross section of the tube. He found the deceleration of the rising flux tube due to the convectively stable photosphere and the subsequent horizontal outflow. \\citet{arc04} performed 3D simulations using the criterion by \\citet{ach79} to analyze the magnetic buoyancy instability within the photosphere/chromosphere, while \\citet{mur06} did parameter studies of the dependence of the initial magnetic field strength of the tube and its twist, finding that the tube evolves in the self-similar way when varying the field strength and that the magnetic buoyancy instability and the second-step evolution do not occur when the field is too weak. In these studies, the initial flux tubes are located in the uppermost convective region at some 1000 km depth. Recently, Solar Optical Telescope (SOT) aboard the {\\it Hinode} satellite \\cite[e.g.][]{tsune08} observed small-scale magnetic flux emergence. Thanks to high-resolution and high-cadence multi-wavelength observations, \\citet{otsu10} found the deceleration of the apex of the arch filament system in the chromosphere, which can be the possible evidence of the two-step model. In this study, we perform two-dimensional fully compressible MHD simulations to investigate the dynamical evolution from the adiabatically stratified convective layer into the corona through the isothermal (strongly sub-adiabatic, i.e., convectively stable) photosphere/chromosphere. We set the initial magnetic flux sheet in the moderately deep convection zone at 20 Mm depth, not at the bottom of the solar interior, because the emergence from the base is beyond the computation ability. The numerical results reproduce the two-step model well. However, the picture of the two-step emergence is much far from the previous studies \\citep[e.g.][]{maga01,arc04,mur06}. The location of the initial flux is so deep that the typical wavelengths of the Parker instability ($10-20$ times the local pressure scale height) are different between the primary and the secondary emergence; the first-step evolution occurs at $z=-20\\ {\\rm Mm}$ and its typical wavelength is about $100\\ {\\rm Mm}$, while, at the photosphere, the second-step emergence has its wavelength of the order of a few $1000\\ {\\rm km}$. We also discuss the dependence of the flux sheet's behavior (`two-step emergence', `direct emergence' or `failed emergence') on its magnetic field strength and amount of fluxes through the parameter survey. The numerical setup and the assumed conditions used in this study are presented in Section \\ref{model}. We show the typical case of the `two-step emergence' in Section \\ref{general}, while Section \\ref{parameter} gives the results of the parameter survey and make some discussions with preceding studies and recent observations. Finally, in Section \\ref{conclude}, we summarize the study. ", "conclusions": "} We perform the nonlinear two-dimensional simulations to investigate the behavior of emerging flux from moderately deep convection zone ($z=-20,000\\ {\\rm km}$). We set a much wider numerical box ($160\\ {\\rm Mm} \\times 80\\ {\\rm Mm}$) than those of the previous experiments on the Parker instability \\citep[e.g.][]{shiba89}. In the typical case ($B=10^{4}\\ {\\rm G}$ with $\\Phi=10^{21}\\ {\\rm Mx}$ at $z=-20,000\\ {\\rm km}$), the results show the `two-step emergence'. In the middle of the way of the first emergence to the solar surface, the flux loop turns from an acceleration phase to deceleration when approaching the (sub-adiabatically stratified, i.e., convectively stable) photosphere/chromosphere. The emerging flux has a sheet-like shape, thus it is difficult for the mass on the loop to escape from the area between the loop and the convectively stable surface. This mass pile-up causes the loop decelerate. The deceleration of the apex of the expanding flux and the continuous rising of the hillsides make loop flattened, which results in the plasma kept on the flux. This deceleration mechanism is another new one and different from those of the preceding studies \\citep{maga01,arc04,mur06} with magnetic flux tubes in much smaller regions. However, our result predicts the behavior of a flux within the convection zone, provided the flux has a sheet-like structure . As a result of the deceleration and the flattening, the flux spreads sideways just beneath the surface, at which point the rise velocity of the crest of the loop is almost zero. Meanwhile, the flux is continuously transported from below, then the magnetic pressure gradient enhances locally in the photosphere. We found that the further evolution to the corona occurs on the basis of the Parker instability. At the point of the instability, plasma $\\beta$ is calculated to be order of unity ($\\sim 2$) and the magnetic field strength is about $700\\ {\\rm G}$. In the final stage, the flux shows the nonlinear evolution to the corona, which resembles the classical experiments \\citep[e.g.][etc.]{shiba89}. The second-step evolution is described clearly by the expansion law by \\citet{shiba89}. We find that the coronal loop exhibits $80,000\\ {\\rm km}$ width with $40,000\\ {\\rm km}$ height, while the field strength of each footpoint at the surface is about $1200\\ {\\rm G}$. We perform parameter runs by changing the initial field strength $B_{x}$ and the total flux $\\Phi$ to investigate the condition of the `two-step emergence'. The results of the runs under considerations can be divided into three groups: `direct', `two-step', and `failed' emergence. In case of the `direct emergence', the flux do evolve to the corona, but they do not show the deceleration by the isothermal surface due to their strong initial magnetic fields ($10^{23} - 10^{24}\\ {\\rm Mx}$ with $10^{5}\\ {\\rm G}$ at $z=-20,000\\ {\\rm km}$). The coronal loops present irregularly strong flux densities at the footpoints; thus, we conclude that they are not suitable for the formation models of active regions. As for the cases showing the `two-step emergence', two out of five exhibit the favorable values of the photospheric field strength and plasma $\\beta$. The others have so large values that they cannot be regarded as realistic models of active regions. We can say that active regions on the sun are likely to have undergone the deceleration and likely to show the `two-step emergence' mentioned above. The condition for this `two-step' active region is ranging from $10^{21}$ to $10^{22}\\ {\\rm Mx}$ with $10^{4}\\ {\\rm G}$ at $z=-20,000\\ {\\rm km}$ in the convection zone. Some recent observations support this two-step model. The cases with $B\\lesssim 10^{4}\\ {\\rm G}$ reveal `failed' evolutions; they fragment within the convection zone or cannot have sufficient magnetic pressure gradient to trigger the instability that the second-step emergence do not occur although the flux maintains its coherency. We have some discussions in connection with the results of the thin-flux-tube (TFT) model by \\citet{mor95}. The cases which are found to have `exploded' in the deeper point in the TFT scheme also do not show further evolutions in our MHD scheme. However, there are some cases which escape the `explosion' fail the second-step evolutions, one of which is possibly the source of the magnetic field in the quiet sun. The present calculations are in a two-dimensional scheme solving simplified equations. Thus we have to demonstrate the more realistic experiments in 3D. At the same time, advanced observations by helioseismological technique are needed to reveal the detail of the emerging flux in the convection zone." }, "1003/1003.1731_arXiv.txt": { "abstract": "We provide measurements of the integrated galaxy light at 70, 160, 250, 350 and 500\\,$\\mu$m using deep far-infrared and submillimeter data from space ({\\it Spitzer}) and balloon platform (BLAST) extragalactic surveys. We use the technique of stacking at the positions of 24\\,$\\mu$m sources, to supplement the fraction of the integrated galaxy light that is directly resolved through direct detections. We demonstrate that the integrated galaxy light even through stacking, falls short by factors of 2$-$3 in resolving the extragalactic far-infrared background. We also show that previous estimates of the integrated galaxy light (IGL) through stacking, have been biased towards high values. This is primarily due to multiple counting of the far-infrared/submillimeter flux from 24\\,$\\mu$m sources which are clustered within the large point spread function of a brighter far-infrared source. Using models for the evolution of the luminosity function at $z<1.2$ which are constrained by observations at 24\\,$\\mu$m and 70\\,$\\mu$m, and which are consistent with the results from the stacking analysis, we find that galaxies at $z<1.2$, account for $\\sim95-55$\\% of the extragalactic far-infrared background in the $\\sim70-500\\,\\mu$m range respectively. This places strong upper limits on the fraction of dust obscured star-formation at $z>1$, which are remarkably, below the values derived from the extinction corrected ultraviolet luminosities of galaxies. The largest fraction of the total $40-500$\\,$\\mu$m EBL comes from galaxies between L$_{\\rm IR}$=L$(8-1000\\,\\mu m)$=10$^{10-11.5}$\\,L$_{\\sun}$; ultraluminous infrared galaxies with L$_{\\rm IR}>10^{12}$\\,L$_{\\sun}$ contribute only 5\\%. We use the results to make predictions for the nature of galaxies that extragalactic surveys with {\\it Herschel Space Observatory} will reveal. Finally, from our constraints on the far-infrared IGL, we provide evidence for the existence of ice mantle dust, orbiting the sun at a distance of $\\sim$40 AU, which is contributing intensity to both the near- and far-infrared background. The presence of this component which is a tiny fraction of the zodiacal light emission even at high ecliptic latitudes, eliminates any discrepancy between the integrated galaxy light and the diffuse, extragalactic background light at all infrared wavelengths. ", "introduction": "The extragalactic background light (EBL)\\footnote{Also referred to as the Cosmic Infrared Background} at far-infrared (FIR) wavelengths, spanning a wavelength range of $40-1000$\\,$\\mu$m, is the sum total of radiation from all astrophysical processes over cosmic time that is absorbed by dust and re-radiated. It places a strong constraint on the fraction of nucleosynthesis and accretion activity that is obscured by dust. Decomposition of the EBL into the contribution from individual galaxies as a function of redshift, provides a more complete understanding of the rate of build up of stellar mass and the growth of black holes over cosmic time and thereby helps identify the epoch at which most of the stars and black holes in the present day Universe formed. The FIR EBL is thought to have a total intensity between $40-1000$\\,$\\mu$m of $\\sim$25\\,\\cirbu\\ although values as high as $\\sim$55\\,\\cirbu\\ have been suggested by DIRBE observations \\citep{Hauser, Finkbeiner:00}. Much of the difference arises due to the difficulty in subtracting the contribution from zodiacal light emission. The zodiacal light contributes $\\sim$30\\% of the line of sight intensity at 140$\\mu$m in even a relatively high ecliptic latitude ($\\beta\\sim45\\arcdeg$) field like the Lockman Hole, with a rapidly increasing fractional contribution at shorter wavelengths \\citep{Hauser}. Even small errors in the subtraction of the zodiacal light propagate through as large errors in the estimate of the FIR EBL, an issue which has plagued measurements of the EBL at 60\\,$\\mu$m and 100\\,$\\mu$m. Despite the fact that it is measured with a factor of $\\sim$100 less precision than the cosmic microwave background, the FIR EBL is arguably the second most energetic extragalactic background after the cosmic microwave background which has a total intensity of $\\sim$1000\\,\\cirbu. The intensity of the FIR EBL is comparable, if not in excess of the intensity of the optical/near-infrared extragalactic background, whose measurement is affected by similar uncertainties associated with the zodiacal light contribution \\citep{Bernstein, Levenson:07}. Since the optical/near-infrared EBL measures the redshifted contribution of unobscured starlight and AGN activity, integrated over cosmic time, the fact that the two intensities are comparable, illustrates the importance of dust obscuration in quantifying the bolometric luminosity of star-forming galaxies and active galactic nuclei. Ever since the direct measurement of the EBL using the Far-Infrared Absolute Spectrometer (FIRAS) and the Diffuse Infrared Background Experiment \\citep[DIRBE;][]{Puget, Hauser, Fixsen}, various attempts have been made to identify the individual sources contributing to the background. The sum of the contribution of individual galaxies is termed the integrated galaxy light (IGL). Using backward evolution models which fit multiwavelength number counts, \\citet[][hereafter CE01]{CE01} demonstrated that dust obscured infrared luminous galaxies contribute 85\\% of the 140\\,$\\mu$m EBL and that 82\\% of the peak of the FIR EBL arises from galaxies at $z<1.5$. The 140\\,$\\mu$m EBL is $\\sim$60\\% of the total FIR EBL in the CE01 models. This implies that about 50\\% of the total FIR EBL arises at $z<1.5$. \\citet{Elbaz} utilized model spectral energy distributions to estimate the contribution to the FIR EBL of infrared luminous galaxies individually detected by the ISOCAM in a deep 15\\,$\\mu$m survey in the Hubble Deep Field-North. They found that the ISOCAM LIRGs at $z<1$ contribute 16$\\pm$5\\,\\cirbu\\ of the 140$\\mu$m background. More recently, by stacking the far-infrared data, \\citet{Bethermin} showed that galaxies with 24$\\mu$m flux densities $>$35\\,$\\mu$Jy contribute 9.0$\\pm$1.1\\,\\cirbu\\ of the 160\\,$\\mu$m background. Based on photometric redshifts, these galaxies are primarily thought to be infrared luminous galaxies at $z\\sim1$ with a tail extending to $\\sim2.5$ \\citep{Caputi}. In this context, the recent results from Balloon Large Aperture Submillimeter Telescope (BLAST) have been somewhat surprising \\citep{Devlin, Marsden}. Using a technique similar to \\citet{Dole}, the BLAST team estimated that stacking on sources with 24$\\,\\mu$m flux density greater than 13\\,$\\mu$Jy results in an IGL of 8.6$\\pm$0.6\\,\\cirbu\\ at 250\\,$\\mu$m (i.e. the majority of the EBL at this wavelength). Furthermore, the conclusion is that 24\\,$\\mu$m detected galaxies contribute $\\sim$80-90\\% of the total FIR EBL between 250 and 500\\,$\\mu$m. The deep 24\\,$\\mu$m data which have been used for their analysis are only sensitive to galaxies brighter than 10$^{10}$\\,L$_{\\sun}$ at $z\\sim1$ and $>$10$^{11}$\\,L$_{\\sun}$ at $z\\sim2$ \\citep{Chary:07a}. It is well known that there is a whole population of Lyman-break galaxies which are less luminous and which appear to have significant dust obscuration based on their UV-slope. Although they should be contributing to the EBL, the BLAST analysis seems to suggest that the contribution from fainter galaxies is surprisingly small. In this paper, we utilize deep, public {\\it Spitzer} data taken as part of the Far-Infrared Deep Extragalactic Legacy (FIDEL; P.I. M. Dickinson) survey\\footnote{http://data.spitzer.caltech.edu/popular/fidel/} at 70 and 160\\,$\\mu$m and the publicly released BLAST maps at 250, 350 and 500\\,$\\mu$m to evaluate the consistency of these different estimates. We combine direct detections of sources with stacking analysis to provide lower limits to the integrated galaxy light. We utilize recent measurements of the far-infrared luminosity function out to $z\\sim1.3$, in conjunction with the model spectral energy distribution of galaxies based on various spectroscopic and imaging data, to determine the nature of sources contributing to the far-infrared background as a function of redshift and luminosity. We place constraints on dust obscured star-formation at $z>1$ and discuss implications for deep {\\it Herschel} extragalactic surveys. Finally, we propose the existence of a previously unknown cloud of high albedo dust grains orbiting the Sun at a distance of $\\sim$40 AU which contributes significantly to the intensity of the near- and far-infrared background. Throughout this paper we assume a standard cosmology with $H_{0}=71\\,\\rm{km}\\,\\rm{s}^{-1}\\,\\rm{Mpc}^{-1}$, $\\Omega_{\\rm{M}}=0.27$ and $\\Omega_{\\Lambda}=0.73$. ", "conclusions": "We have utilized data from the deepest far-infrared surveys currently undertaken to evaluate the fraction of the extragalactic background light that has been resolved into individual galaxies. By using a combination of direct detections and stacking on the imaging data, we demonstrate that more than $\\sim$50\\% of the background currently remains unresolved at all wavelengths between $100-500$\\,$\\mu$m. The 70\\,$\\mu$m IGL, after large completeness corrections however comprises $\\sim$75\\% of the extragalactic background light derived from models for the evolution of the galaxy luminosity function. We also demonstrate that previous estimates of the integrated galaxy light using stacking have significantly overestimated the values by multiple counting of the flux from clustered sources and provide steps through which these can be corrected. We have then, extrapolated measurements of the infrared luminosity function measured at $z<1.2$ to higher redshifts and use the resolved fraction of the background along with the DIRBE/FIRAS measured EBL to derive our best estimates of the dust obscured star-formation rate at $z>1$. We find that the UV slope appears to overestimate the dust extinction at $z>2$, most likely due to the significant contribution of nebular emission to the broadband ultraviolet photometry. We provide average measures of the extinction at these redshifts and find that the average extinction increases with decreasing redshift. If translated to an increase in dust content, this is consistent with the rate of increase in metallicity with redshift, suggesting that the two processes are co-evolving. We also quantify the fractional contribution to the far-infrared EBL as a function of both redshift and luminosity. We find that ultraluminous infrared galaxies contribute $\\sim$5\\% of the total EBL and that the bulk of the star-formation occurs in galaxies with luminosities $\\sim$10$^{10-11.5}$\\,L$_{\\sun}$. Finally, we use the difference between the EBL derived from our models and the DIRBE measured values to provide tentative evidence for a contribution from ice mantle dust in the outer solar system. This dust appears to have a total intensity of $\\sim$25\\,\\cirbu\\ at high ecliptic latitudes and can account for the discrepancy between the IGL and EBL at both near- and far-infrared wavelengths." }, "1003/1003.3444_arXiv.txt": { "abstract": "We report the discovery of HD\\,156668\\,b, an extrasolar planet with a minimum mass of $M_P\\sin i$\\,=\\,4.15\\,\\mearthe. This planet was discovered through Keplerian modeling of precise radial velocities from Keck-HIRES and is the second super-Earth to emerge from the NASA-UC Eta-Earth Survey. The best-fit orbit is consistent with circular and has a period of $P$\\,=\\,4.6455\\,d. The Doppler semi-amplitude of this planet, $K$\\,=\\,1.89\\,\\mse, is among the lowest ever detected, on par with the detection of GJ\\,581\\,e using HARPS. A longer period ($P$\\,$\\approx$\\,2.3\\,yr), low-amplitude signal of unknown origin was also detected in the radial velocities and was filtered out of the data while fitting the short-period planet. Additional data are required to determine if the long-period signal is due to a second planet, stellar activity, or another source. Photometric observations using the Automated Photometric Telescopes at Fairborn Observatory show that HD\\,156668 (an old, quiet K3 dwarf) is photometrically constant over the radial velocity period to 0.1\\,mmag, supporting the existence of the planet. No transits were detected down to a photometric limit of $\\sim$3\\,mmag, ruling out transiting planets dominated by extremely bloated atmospheres, but not precluding a transiting solid/liquid planet with a modest atmosphere. ", "introduction": "\\label{sec:intro} The search for low-mass planets is driven by a desire to observationally study the full range of planetary systems in order to better understand their formation, dynamics, composition, and diversity. We also seek Earth-like worlds of terrestrial composition as a goal in itself and as targets for future transit and imaging observations. This search has taken several leaps forward recently because of instrumental improvements. The precision of radial velocity (RV) measurements with Keck-HIRES by the California Planet Survey (CPS) group \\citep{Howard09a}, HARPS \\citep{Lovis06}, and the AAT \\citep{OToole09} has now reached 1\\,\\ms or better and has led to the discovery of several super-Earths around nearby, bright stars. Ground-based transit surveys such as MEarth \\citep{Charbonneau09} and HATNet \\citep{Bakos09} have made important discoveries of transiting low-mass planets. Microlensing searches have detected two super-Earths orbiting distant stars \\citep{Beaulieu06,Bennett08} and the statistics of microlensing detections suggest than cold Neptunes are a factor of three more common that cold Jupiters \\citep{Sumi09}. From space, \\textit{CoRoT} has found a system with two super-Earths (one of which transits; \\citealt{Leger09}) and \\textit{Kepler} is poised to detect true Earth analogues in 1\\,AU orbits using transit photometry with a precision of 20\\,ppm in 6.5\\,hr \\citep{Borucki09}. In the next decade, \\textit{SIM Lite} \\citep{Unwin08} will astrometrically characterize essentially all planets down to Earth mass orbiting $\\sim$100 nearby stars, as well as the more massive planets orbiting $\\sim$1000 stars. The Eta-Earth Survey plays a unique role in the study of low-mass exoplanets. The population of 230 GKM stars in the survey is nearly free of statistical bias since the stars were not chosen based on their likelihood of harboring a planet, but rather on proximity, brightness, and chromospheric activity. Each star is observed at least 20 times, insuring a minimum detectability threshold. Thus, the distributions of planet detections and non-detections from the Eta-Earth Survey will yield a wealth of information about the efficiency and mechanisms of planet formation as well as the range of subsequent dynamical histories. The 20 survey observations per Eta-Earth Survey target are nearly complete and we are aggressively re-observing several promising candidate low-mass planets. Current theories of planet formation \\citep{Ida04a,Mordasini09} are consistent with the measured distributions of massive planets (Saturn mass and above), but their predictions for the abundance and properties of low-mass planets are only now being observationally tested. In particular, they predict a dearth of planets below roughly Saturn mass in orbits inside the ice line. Such planets are predicted to be rare because once a planet core grows by planetesimal accretion to a of threshold mass of $\\sim$10\\,\\mearthe, it undergoes runaway gas accretion and becomes a gas giant. While these theories consistently predict such a ``planet desert'', they differ in the contours of the desert in $M$ and $a$ as well as whether the planets below the critical core mass survive Type I inward migration in significant numbers. In this context we announce the discovery of HD\\,156668\\,b, a super-Earth planet with minimum mass $M_P\\sin i$\\,=\\,4.15\\,\\mearth and an orbital period of $P$\\,=\\,4.6455\\,d. This is the second super-Earth (\\msinie\\,$<$\\,10\\,\\mearthe) to emerge from Keck observations explicitly for the NASA-UC Eta-Earth Survey, the first being HD\\,7924\\,b \\citep{Howard09a}. The remainder of this paper is structured as follows. We describe the host star properties in \\S\\,\\ref{sec:props}. The spectroscopic observations and their Doppler reduction are described in \\S\\,\\ref{sec:obs}. In \\S\\,\\ref{sec:orbital}, we describe the detection of the $P$\\,=\\,4.6455\\,d orbit of HD\\,156668\\,b, and the high-pass filtering of the RV data that was necessary to obtain good estimates of the orbital parameters. In \\S\\,\\ref{sec:null} we carefully consider the null hypothesis---the non-existence of HD\\,156668\\,b---using $S_\\mathrm{HK}$ measurements, photometric observations, and FAP analyses. We summarize and discuss the implications of this discovery in \\S\\,\\ref{sec:discussion}. ", "conclusions": "\\label{sec:discussion} We present the detection of HD\\,156668\\,b, a super-Earth planet with minimum mass $M_P\\sin i$\\,=\\,4.15\\,$M_\\earth$ in a $P$\\,=\\,4.6455\\,d orbit around a K3 dwarf. We draw on several lines of evidence to support the existence of HD\\,156668\\,b. We showed in \\S\\,\\ref{sec:fap} that the short-period signal is statistically significant. This signal is apparent in a fit to the unfiltered RVs, and stands out strongly when isolated by high-pass filtering. The host star is middle-aged and quiet, providing a nearly ideal RV target. The planet's short-period signature is not seen in photometric observations or in chromospheric indices. Thus, the evidence strongly points to a planetary interpretation. To estimate the orbital parameters of HD\\,156668\\,b, we found it necessary to apply a high-pass filter to the RVs by subtracting the model of a long-period Keplerian. This model changes slowly over time and clarifies, but does not artificially enhance, the $P$\\,=\\,4.6455\\,d signal of HD\\,156668\\,b. Filtering of this type is common in other areas of time domain astronomy (e.g.\\ transit photometry, astroseismology) and is uncommon, but not unprecedented, in RV planet detection. Similar filtering techniques were employed to disentangle the planetary signals of Gl\\,176\\,b \\citep{Forveille09} and Gl\\,674\\,b \\citep{Bonfils07}, although in those cases the non-planetary signal was clearly due to spots modulated by stellar rotation. \\cite{Queloz09} used a ``pre-whitening'' technique to extract the signatures of two super-Earths from the complicated RV time series of the significantly more active star, Corot-7. Further, many exoplanets are announced with a model that includes a linear or second order RV trend, presumably due to a long period orbital companion. The source of the long-period signal remains unknown and motivates additional measurements. If it is due to a planet, the body has minimum mass $\\sim$45\\,\\mearth and orbits with $P_c$\\,=\\,2.31\\,yr and $a_c$\\,=\\,1.6\\,AU, a cold super-Neptune near the ice line. We see no evidence for transits of HD\\,156668\\,b down to a level of $\\sim$3\\,mmag. However, given the large \\textit{a priori} transit probability of 7\\%, it is instructive to speculate about the transit signatures of various possible planet compositions. Using the models in \\cite{Seager2007}, a 4\\,\\mearth planet composed of pure Fe, MgSiO$_3$, H$_2$O, or H would yield planets of radius $R_{\\mathrm{pl}}$\\,=\\,1.2, 1.5, 2.0, and 4.5\\,\\rearthe, producing transits of depth 0.22, 0.35, 0.61, and 3.1\\,mmag, respectively. These homogeneous planet models are oversimplified, but set the scale for admixtures of those ingredients: transits of planets made of solids and water would have depths of $\\sim$0.2--0.6\\,mmag, while transits of a planet with a significant atmosphere could be much deeper. For comparison, the transiting super-Earth Corot-7\\,b has a transit depth of 0.36\\,mmag implying a radius $R_{\\mathrm{pl}}$\\,=\\,1.68\\,\\rearth \\citep{Leger09}. Using the RV-determined mass $M_{\\mathrm{pl}}$\\,=\\,4.8\\,\\mearth \\citep{Queloz09}, the bulk density is terrestrial, $\\rho_{\\mathrm{pl}}$\\,=\\,5.6\\,g\\,cm$^{-3}$. In contrast, GJ\\,1214\\,b has $M_{\\mathrm{pl}}$\\,=\\,6.6\\,\\mearth and $R_{\\mathrm{pl}}$\\,=\\,2.7\\rearthe, implying $\\rho_{\\mathrm{pl}}$\\,=\\,1.9\\,g\\,cm$^{-3}$, intermediate between Earth and the ice giants of the Solar System. Transits of GJ\\,1214\\,b are unusually deep (15\\,mmag) for a planet of this size because it orbits an M dwarf with $R_{\\star}$\\,=\\,0.21\\,$R_{\\sun}$ \\citep{Charbonneau09}. Several authors (e.g.\\ \\citealt{Ida04a}, \\citealt{Mordasini09}) have argued that super-Earths will have insignificant hydrogen atmospheres (by mass) because during formation their masses failed to reach a critical mass (typically $\\sim$10\\,\\mearthe) when the growth from solid planetesimals is augmented by runaway accretion of gas from the protoplanetary disk. Smaller atmospheres (up to several per cent by mass) could be produced by degassing during impact accretion and geological activity \\citep{Elkins-Tanton08,Kite09}. However, whatever atmosphere is acquired from these processes may be lost to atmospheric escape \\citep{Baraffe06,Valencia09}. Nevertheless, the brief history of exoplanets is replete with observational surprises (hot Jupiters, eccentric orbits, etc.) so we consider the observational consequences of an atmosphere. \\cite{Adams08} find that adding a H/He gas envelope equivalent to 0.2--20\\% of the mass of a solid 5\\,\\mearth exoplanet increases the radius 8--110\\% above the gas-free value. Atmospheres dominated by heavier molecules such as H$_2$O and N$_2$ (as on Earth) would swell the planet less for the same atmospheric mass because of the higher mean molecular weight and reduced scale height. Thus, we conclude that the APT photometric observations rule out transits for HD\\,156668\\,b if the radius is dominated by a H/He atmosphere (tens of per cent by mass), but do not preclude transits if the atmosphere is less massive or composed of heavier elements. HD\\,156668\\,b adds statistical weight to the emerging trends of the properties of super-Earths and their hosts. Like most other such planets, it orbits a K or M dwarf. % In contrast to other super-Earth hosts, HD\\,156668 has a slightly super-solar metallicity \\citep{Howard09a}. The rate of multiplicity in systems with super-Earths and Neptune-mass planets appears to be much higher than for higher mass planet hosts, with HD\\,40307 \\citep{Mayor09}, GJ\\,581 \\citep{Mayor09b}, and HD\\,69830 \\citep{Lovis06} being the standard examples of multiplicity in low-mass systems. The long-period signal seen for HD\\,156668 is suggestive of a second low-mass planet in the system. HD\\,156668\\,b pushes the frontier of RV planet discovery to lower masses and smaller Doppler amplitudes. It is the second lowest minimum mass exoplanet discovered to date by the RV technique, after GJ\\,581\\,e \\citep{Mayor09b}. With a Doppler semi-amplitude of 1.89\\,\\mse, HD\\,156668\\,b is also only the second exoplanet discovered to date with $K < 2.0$\\,\\mse. GJ\\,581\\,e is the other with $K$\\,=\\,1.85\\,\\ms and \\msinie\\,=\\,1.9\\,\\mearth \\citep{Mayor09b}. This progress is remarkable: 51\\,Peg\\,b was discovered \\citep{Mayor95} with \\msinie\\,=\\,0.47\\,\\mjup and $K$\\,=\\,57\\,\\ms while the signal from HD\\,156668\\,b is smaller by factors of 35 and 30, respectively. However, we expect that true Earth analogs---1\\,\\mearth planets in 1\\,AU orbits around G stars---will remain permanently out of reach of the Doppler technique and will instead be discovered by transit photometry, astrometry, and microlensing. The Earth imparts a $K$\\,=\\,0.09\\,\\ms signal on the Sun, a factor of 20 smaller than HD\\,156668\\,b. Our Keck/HIRES observations of HD\\,156668 take $\\sim$5\\,min to achieve 1\\,\\ms precision; scaling to 0.1\\,\\ms precision requires 100 times the photons and impractically long integration times. (A handful of exceptionally bright nearby stars such as $\\alpha$ Cen A/B may be an exception \\citep{Guedes08}.) Moreover, stellar jitter, even when averaged over with long integrations, likely limits Doppler precision to a few tenths of a \\ms for the most quiet stars. Longer period planets suffer from the additional disadvantage that only a few orbits transpire during an observational campaign (compared with $\\sim$100 orbits for the discoveries of HD\\,155668\\,b and GJ\\,581\\,e), reducing their detectability by clock-like coherence. Nevertheless, the future of the Doppler technique is bright and we expect that it will continue to play a prominent role in planet searches. Discoveries of super-Earths and Neptunes from the Eta-Earth Survey will shape our understanding of planet formation and migration. Doppler work in other domains---long period giant planets, subgiants, multi-planet systems, dynamically interacting planets, Rossiter-McLaughlin measurements, etc.---will also continue to inform and enrich our knowledge of the rich astrophysics of exoplanets." }, "1003/1003.5262_arXiv.txt": { "abstract": "We study the effect of the non-linear process of ambipolar diffusion (joint transport of magnetic flux and charged particles relative to neutral particles) on the long-term behavior of a non-uniform magnetic field in a one-dimensional geometry. Our main focus is the dissipation of magnetic energy inside neutron stars (particularly magnetars), but our results have a wider application, particularly to the interstellar medium and the loss of magnetic flux from collapsing molecular cloud cores. Our system is a weakly ionized plasma in which neutral and charged particles can be converted into each other through nuclear beta decays (or ionization-recombination processes). In the ``weak-coupling'' limit of infrequent inter-particle interactions, the evolution of the magnetic field is controlled by the beta decay rate and can be described by a non-linear partial integro-differential equation. In the opposite, ``strong-coupling'' regime, the evolution is controlled by the inter-particle collisions and can be modelled through a non-linear diffusion equation. We show numerically that, in both regimes, ambipolar diffusion tends to spread out the magnetic flux, but, contrary to the normal Ohmic diffusion, it produces sharp magnetic field gradients with associated current sheets around those regions where the magnetic field is weak. ", "introduction": "Ambipolar diffusion is the joint drift of charged particles and the associated magnetic flux with respect to the neutral particles in a partially ionized plasma. \\citet{MestelSpitzer56} first proposed it in order to explain the loss of magnetic flux from the dense cores of molecular clouds, required for the formation of stars, starting an active field of research in this area. Later, it was suggested to also play a role in the decay of the magnetic fields of neutron stars (\\citealp{J-87,H-91,P-92,GR-92}, hereafter GR-92) which became particularly relevant with the identification of ``magnetars'', neutron stars whose main power source appears to be the dissipation of their magnetic field \\citep{DT-92,TD-96,ACT-04}. In a previous paper (\\citealt{HRV-08}; hereafter Paper I), we established a multifluid formalism in which it is possible to study the long-term evolution of magnetic fields in neutron stars (see \\citealt{R-09} for a discussion of the main properties of the magnetic field equilibria and their subsequent long-term evolution). In that work, and following the ideas developed by GR-92, we included the effects of several physical processes that are also relevant for star formation and protoplanetary disks, including {\\it ambipolar diffusion}, {\\it Hall drift} (non-dissipative advection of the magnetic field by the associated electrical current), and {\\it ohmic diffusion} (dissipation of currents through the electrical resistivity). Here we continue this study and concentrate our analysis on the long-term evolution of the magnetic field caused by ambipolar diffusion aided by beta decays. Following the same philosophy of Paper I, and as a first approach to the understanding of our general formalism, we focus on a simplified, one-dimensional configuration in which the magnetic field points in one Cartesian direction $z$ but varies only along an orthogonal direction $x$, i.~e., $\\vec B=B(x,t)\\hat z$. Such models have also been considered in several studies of ambipolar diffusion in the interstellar medium \\citep{MouschoviasPaleologou81,Shu-83,BZ-94}, although some of the assumptions differ from case to case. In our analysis, we consider separately two relevant limits, similar in spirit, though not exactly equivalent to those of \\citet{MouschoviasPaleologou81}: \\begin{itemize} \\item In the {\\it weak coupling} limit, there are few collisions between the particles, and the beta decays proceed slowly. Therefore, the particles can reach the diffusive equilibrium easily, but it takes much longer to reach the chemical equilibrium. \\item In the {\\it strong coupling} limit, there are many collisions between the particles, and the beta decays proceed fast, so the (local) chemical equilibrium is reached much more quickly than the diffusive equilibrium. \\end{itemize} For each of these cases, we find that the long-term evolution of the magnetic field can be modelled by a single equation that gives the time-derivative $\\partial B/\\partial t$ at a given instant $t$ only in terms of the configuration of the magnetic field at the same instant, $B(x,t)$. This makes it easy to carry out numerical simulations of the evolution of some selected non-linear magnetic field profiles, and even find some exact, analytical solutions. In Sect.~\\ref{1dmodel} of this paper, we briefly review the one-dimensional model of neutron star magnetic field evolution introduced in Paper I, paying particular attention to its characteristic evolutionary time-scales. In Sects.~\\ref{asy1} and \\ref{asy2}, we obtain the equations for the long-term, asymptotic magnetic field evolution promoted by ambipolar diffusion in each of the two opposite regimes mentioned above, and we make numerical simulations of the evolution of different initial magnetic field configurations. We show that, in both cases, the magnetic flux of a given sign tends to spread out, but singularities develop at the null points where regions with different signs meet, as previously found by \\citet{BZ-94}. In the weak coupling case, these singularities correspond to current sheets that are dissipated by resistive effects, in this way leading to reconnection. In the strong coupling case, the singularities have a somewhat different character (a smoothly diverging current density) and might lead directly to reconnection even in the case of no ohmic resistivity (but see \\citealp{H-03a,H-03b}). Finally, in Sect.~\\ref{conc}, we give the main conclusions of our study. ", "conclusions": "\\label{conc} We have studied the asymptotic magnetic field evolution promoted by ambipolar diffusion in a one-dimensional geometry for two opposite, limiting regimes. In the {\\it weak coupling limit}, in which neutral and charged particles drift easily with respect to each other, the bottleneck for the evolution is the conversion of one species into another, which is required in order to eliminate the charged-particle pressure gradients caused by the magnetic field, which impede its evolution. In the {\\it strong coupling limit}, conversions are easy, but the inter-particle collisions are the corresponding bottleneck. In molecular clouds, the second regime appears to be generally relevant \\citep{Shu-83,BZ-94}. Neutron stars in their hot, early phase will also be in the strong-coupling regime, and evolve to the weak coupling regime as they cool. In the {\\it weak coupling limit}, the magnetic field evolution is described by a non-linear, partial integro-differential equation, while in the {\\it strong coupling limit} this evolution is described by a non-linear diffusion equation. We made numerical simulations of the evolution of different initial magnetic field profiles in each of these limits and found agreement between our numerical results and some analytic solutions that can be found for these differential equations. From our results, we infer that, in both limits, the ambipolar diffusion process operates in a tendency to spread out the magnetic flux, but contrary to the normal Ohmic diffusion this process asymptotically produces singular points with sharp magnetic field gradients. These sharp gradients develop around those points where the magnetic field is null, and separate regions of magnetic fields with opposite signs. We observe some generic properties of this process, as follows: In the {\\it weak coupling limit}, the resulting discontinuities can be modelled as step solutions [Eqs.~(\\ref{explicit-step}), (\\ref{explicit-step1})]. The asymptotic magnetic field is spatially uniform in each of the regions separated by these singularities, and its absolute value (and thus the magnetic pressure) is the same in each region. Ambipolar diffusion by itself does not cause magnetic flux transfer (and thus reconnection) across the singularities, but the associated current sheets will easily be dissipated by Ohmic diffusion, so reconnection will occur in a realistic system. In the {\\it strong coupling limit}, at the singular points the magnetic field vanishes continuously but with infinite spatial derivative. Ambipolar diffusion acts in a tendency to spread out the magnetic flux, but, contrary to the weak coupling limit, the magnetic flux is not preserved in each one of the regions separated by the magnetic null points. Therefore, there is a transfer of magnetic flux through these null points, i.~e. reconnection without Ohmic resistivity (but see \\citealp{H-03a,H-03b}). The main limitation in applying the present formalism to realistic systems (either neutron stars or molecular cloud cores) is the very restrictive, one-dimensional geometry. An extension to more realistic geometries (i.~e., axial symmetry) will be attempted in further work." }, "1003/1003.2397_arXiv.txt": { "abstract": "As members of the instrument team for the Advanced CCD Imaging Spectrometer (\\ACIS) on NASA's {\\it Chandra X-ray Observatory} and as \\Chandra\\ General Observers, we have developed a wide variety of data analysis methods that we believe are useful to the \\Chandra\\ community, and have constructed a significant body of publicly-available software (the {\\it ACIS Extract} package) addressing important \\ACIS\\ data and science analysis tasks. This paper seeks to describe these data analysis methods for two purposes: to document the data analysis work performed in our own science projects, and to help other \\ACIS\\ observers judge whether these methods may be useful in their own projects (regardless of what tools and procedures they choose to implement those methods). The \\ACIS\\ data analysis recommendations we offer here address much of the workflow in a typical \\ACIS\\ project, including data preparation, point source detection via both wavelet decomposition and image reconstruction, masking point sources, identification of diffuse structures, event extraction for both point and diffuse sources, merging extractions from multiple observations, nonparametric broad-band photometry, analysis of low-count spectra, and automation of these tasks. Many of the innovations presented here arise from several, often interwoven, complications that are found in many \\Chandra\\ projects: large numbers of point sources (hundreds to several thousand), faint point sources, misaligned multiple observations of an astronomical field, point source crowding, and scientifically relevant diffuse emission. ", "introduction": "} Since its launch in 1999, the {\\it Chandra X-ray Observatory} \\citep{Weisskopf02} has revolutionized X-ray astronomy. \\Chandra\\ provides remarkable angular resolution---unlikely to be matched by another X-ray observatory within the next two decades---and its most commonly used instrument, the Advanced CCD Imaging Spectrometer (\\ACIS), produces observations with a very low background \\citep{Garmire03}.\\footnote{See also the \\anchorparen{http://asc.harvard.edu/proposer/POG/pog_pdf.html}{Chandra Proposers' Observatory Guide}.} These two technical capabilities allow detection of point sources with as few as ${\\sim}5$ observed X-ray photons (commonly referred to as ``events'' or ``counts''), a data analysis regime unique among X-ray observatories. Observations of Galactic star clusters and mosaics of nearby galaxies or extragalactic deep fields often produce hundreds to thousands of weak X-ray sources. \\Chandra's excellent sensitivity to point sources and angular resolution also provide a unique capability for studying diffuse emission superposed onto those point sources, since they can be effectively identified and then masked. For many types of \\ACIS\\ ``imaging'' studies,\\footnote{ Our discussion here is limited to data taken in the most common \\ACIS\\ configuration, called Timed Exposure Mode, at either the \\ACIS-I or \\ACIS-S aimpoint. Dispersed data from the \\Chandra\\ gratings are not addressed here. } most observers follow a data analysis workflow that is similar to that outlined in Figure~\\ref{workflow.fig}. Relatively raw data derived from satellite telemetry, known as \\anchorfoot{http://cxc.harvard.edu/ciao/dictionary/levels.html}{``Level 1 Data Products''} (L1), are passed through a variety of repair and cleaning operations to produce ``Level 2 Data Products'' (L2) that are appropriate for analysis. A common workflow for studying point sources (solid boxes and arrows on left side of Figure~\\ref{workflow.fig}) consists of binning the L2 event data into one or more images that are searched for sources. The events and background associated with each point source in the catalog are ``extracted'' and calibrated. Observed sources properties (e.g., count rates, apparent fluxes, spectra, light curves) are estimated, then combined with calibration products to estimate intrinsic astrophysical source properties. A common and similar workflow for studying diffuse emission (right side of Figure~\\ref{workflow.fig}) consists of removing (``masking'') the point sources from the data, constructing images, identifying several regions of diffuse emission to study, and then extracting and analyzing those diffuse sources in a manner similar to that used for point sources. \\begin{figure}[htb] \\centering \\includegraphics[width=1.0\\textwidth]{f1} \\caption{The data analysis workflow described here for a single \\ACIS\\ field containing multiple point sources (PS, left branch) and diffuse sources (DS, right branch). \\label{workflow.fig}} \\end{figure} Many \\Chandra\\ studies exhibit one or more of five characteristics that significantly complicate this familiar workflow. \\begin{enumerate} \\item In many studies hundreds to thousands of point sources can be readily identified; executing the workflow is intractable without significant automation. \\item Numerous sources with very few detected counts are identified in most \\Chandra\\ studies; common statistical methods based on large-N assumptions break down at several points in the workflow. \\item Many studies require covering a large field of view with multiple \\Chandra\\ pointings (e.g., Figure~\\ref{emap.fig}). For many reasons, such mosaicked pointings usually overlap significantly and/or are observed at a variety of roll angles. Thus, many sources are observed multiple times at very different locations on the \\ACIS\\ detector. Analysis of such sources can be very complex, because the \\Chandra\\ point spread function (PSF) exhibits large variations in size and shape across the focal plane.\\footnote{See Figure~8 and Figure~10 in the \\anchorparen{http://cxc.harvard.edu/cal/Hrma/users_guide/}{HRMA User's Guide}. } Analysis of diffuse emission is also complicated by multiple pointings, since a single diffuse region may be only partially covered by a particular \\ACIS\\ observation. \\item Some fields (such as deep exposures of rich star clusters, the Galactic Center, or nearby galaxies) are crowded, with adjacent PSFs in close proximity. In such conditions, simple approaches to source extraction and background estimation are not adequate. \\item Scientifically relevant diffuse emission must often be extracted and studied while avoiding contamination from point sources (e.g., in star forming regions, the Galactic Center, and galaxy clusters in deep fields). \\end{enumerate} \\begin{figure}[htb] \\centering \\includegraphics[width=0.8\\textwidth]{f2} \\caption{ Exposure map for the Chandra Carina Complex Project \\citep{Townsley10} study of the Carina Nebula comprised of 22 ACIS-I pointings (38 observations), with point sources masked (\\S\\ref{bkg_extract.sec}) in preparation for extraction of diffuse emission (\\S\\ref{diffuse.sec}). All point sources are masked in the individual, high-resolution, exposure maps; at this scale only the larger masks (on mostly off-axis sources) are visible. \\label{emap.fig}} \\end{figure} Since our \\Chandra\\ studies often exhibit all five of these issues, we have developed a set of data analysis methods that incorporate a variety of enhancements to standard techniques. Publicly available software implementing these methods---the {\\it ACIS Extract} package---has been cited in publications for at least 50 \\Chandra\\ targets. This paper seeks to describe these methods at a moderate level of detail. Our primary purpose is to encourage other \\ACIS\\ observers to consider whether the various departures from standard techniques described here may be useful in their own projects (regardless of what tools and procedures those observers use to implement those methods). Our secondary purpose is to document the data analysis methods used in our own \\Chandra\\ studies. In many cases, the only documentation available for software and data analysis methods is on-line; thus this paper is liberally footnoted with relevant URLs. We acknowledge that URLs are more ephemeral than journal citations but we believe that they are better than no documentation at all. The high-level structure of our data analysis workflow differs from standard practice in two ways. First, for technical reasons, the optimal data cleaning steps for point sources and diffuse sources differ for \\ACIS\\ data (\\S\\ref{L1_L2.sec}); thus Figure~\\ref{workflow.fig} shows separate ``cleaning'' operations on those two branches of the analysis. Second, since the point source extraction process generates estimates of source position (\\S\\ref{src_positions.sec}) and source validity (\\S\\ref{signif.sec}) that are expected to be better than the estimates made by typical source detection procedures, we choose to adopt an iterative workflow (\\S\\ref{iterative_detection.sec} and dashed boxes and lines in Figure~\\ref{workflow.fig}) in which the point source detection process merely nominates candidates that are then repositioned and possibly discarded as likely noise peaks after extraction results are in hand. Most of this paper describes the individual data analysis tasks implied by Figure~\\ref{workflow.fig}, emphasizing the changes to standard methods that we have adopted. We assume the reader is familiar with \\ACIS\\ data and with standard analysis methods, both of which are well described by the \\anchorfoot{http://cxc.harvard.edu/ciao/threads}{Chandra Science Threads.} Throughout the text, we make liberal use of footnotes that direct the reader to on-line documentation, much of it provided by the Chandra X-ray Center (CXC), that is useful for understanding \\ACIS\\ data analysis and issues. The process of preparing L2 data products is discussed in \\S\\ref{L1_L2.sec}. Point source detection is reviewed in \\S\\ref{source_det.sec}. Extraction of point sources and background estimation are presented in \\S\\ref{extraction.sec}. Section~\\ref{merging.sec} describes our approach to handling multiple observations of a source. Estimation of observed and intrinsic source properties is discussed in \\S\\ref{src_prop.sec}. In \\S\\ref{diffuse.sec} we describe modifications of point source methods for use on diffuse sources. It is critical for the reader to recognize that the methods described here are not unique, definitive, or optimal for all purposes. Alternative approaches to \\ACIS\\ data analysis are provided by the \\anchorfoot{http://asc.harvard.edu/csc/}{Chandra Source Catalog} developed by the CXC, and by analysis tools such as \\anchorfoot{http://xassist.pha.jhu.edu/zope/xassist}{{\\it XAssist}} \\citep{Ptak03} and \\anchorfoot{http://cxc.harvard.edu/contrib/yaxx/}{{\\it yaxx}} \\citep{Aldcroft06} developed by other researchers. ", "conclusions": "} Many \\Chandra-\\ACIS\\ imaging studies face significant data analysis challenges arising from large numbers of weak and sometimes crowded point sources embedded in scientifically relevant diffuse emission, observed with multiple misaligned pointings. We have discussed here a variety of innovations to standard \\ACIS\\ analysis methods that address these challenges; the most important of these are summarized below. \\begin{enumerate} \\item Currently, a single set of data cleaning procedures is not adequate if the observer plans to study both very weak (${\\leq}10$ counts) point sources (or diffuse emission) and bright point sources, because several cleaning algorithms remove legitimate X-ray events from bright sources. Thus, we find that distinct data cleaning procedures are required for different types of analyses (\\S\\ref{L1_L2.sec}). \\item When point sources are to be extracted, we recommend evaluating the existence of sources using those extraction results rather than relying on typical source detection tools for that judgment (\\S\\ref{iterative_detection.sec}). \\item In crowded fields we recommend searching for candidate sources in reconstructed images (\\S\\ref{cand_src.sec}), defining extraction apertures that do not overlap (\\S\\ref{extract_reg.sec}), and defining background regions that seek to model the background contributed by the wings of nearby sources (\\S\\ref{bkg_extract.sec}). \\item We recommend correcting all point source extractions for the energy-dependent fraction of light that lies outside the source aperture (\\S\\ref{aperturecorr.sec}). \\item When a source is observed multiple times, we recommend estimating source validity, position, and photometry using three independent combinations of the extractions, each allowed to discard observations to optimize the accuracy of the corresponding quantity (\\S\\ref{composite_discard.sec}). \\item We offer an algorithm for grouping spectra (\\S\\ref{grouping.sec}) that lessens biases inherent in standard algorithms. \\item We raise concerns about employing under-constrained background models in the context of spectral fitting and offer an alternative (\\S\\ref{XSPEC.sec}). \\item When diffuse X-rays are extracted from a region that contains unobserved areas (e.g., due to detector edges or point source masks) we describe a necessary correction to the calibration provided by the tool {\\em mkwarf} (\\S\\ref{exposure.sec}). \\end{enumerate} Most of the data analysis methods we have discussed are implemented in the \\anchor{http://www.astro.psu.edu/xray/acis/acis_analysis.html}{{\\em ACIS Extract} (AE) software package,} which has been freely available to the community since its development began in 2002.\\footnote { Development of AE is on-going; observers can be notified of new releases by joining the \\anchorparen{http://lists.psu.edu/cgi-bin/wa?A0=L-ASTRO-ACIS-EXTRACT}{AE mailing list}. } \\AEacro\\ is written in the \\IDL\\ language, and makes extensive use of tools in \\CIAO\\ and in several other public packages. Although much of the analysis we perform on our \\ACIS\\ observations has been automated, we believe that the obserer should retain many vital roles in the process. The human eye is often able to spot omissions and spurious entries in the set of candidate sources derived from detection procedures (\\S \\ref{cand_src.sec}). We rely on the observer to judge what effect CCD readout streaks may have on the data analysis and to take appropriate mitigating actions (\\S \\ref{bkg_extract.sec}). Scientific judgement is required to select appropriate levels of smoothing for diffuse sources (\\S \\ref{smooth.sec}) and to define appropriate regions to extract (\\S \\ref{diffuse_extract.sec}). We encourage the observer to visually review (\\S \\ref{visualiz.sec}) data cleaning steps, extraction apertures, catalog pruning and source repositioning proposed by algorithms, spectral fitting results, and multi-wavelength associations asserted by our matching algorithm." }, "1003/1003.0671_arXiv.txt": { "abstract": "We present cosmological hydrodynamical simulations of the formation of dwarf galaxies in a representative sample of haloes extracted from the Millennium-II Simulation. Our six haloes have a $z=0$ mass of $\\sim10^{10}\\Ms$ and show different mass assembly histories which are reflected in different star formation histories. We find final stellar masses in the range $5\\times10^7 - 10^8\\Ms$, consistent with other published simulations of galaxy formation in similar mass haloes. Our final objects have structures and stellar populations consistent with observed dwarf galaxies. However, in a $\\Lambda$CDM universe, $10^{10} \\Ms$ haloes must typically contain galaxies with much lower stellar mass than our simulated objects if they are to match observed galaxy abundances. The dwarf galaxies formed in our own and all other current hydrodynamical simulations are more than an order of magnitude more luminous than expected for haloes of this mass. We discuss the significance and possible implications of this result. ", "introduction": "Dwarf galaxies are by far the most abundant type of galaxy in the Local Group and in the Universe. They span a large range of stellar masses, morphologies and star formation histories. The largest dwarf irregulars such as the large Magellanic Cloud have stellar masses of $\\sim 10^9 \\Ms$, rotationally supported and HI-rich disks, and strong ongoing star formation. In contrast, dwarf spheroidal galaxies have stellar masses from $10^7 \\Ms$ to below $10^3 \\Ms$, they possess no interstellar gas, and they show no sign of rotational support or ongoing star formation. The number of dwarf galaxies observed in the Local Group continues to grow as new, `ultra-faint' satellite galaxies are discovered \\citep[e.g.][]{Martin-2006, Chapman-2008, Belokurov-2010}. Estimates using luminosity functions corrected for completeness and bias predict the total number of faint satellites to be an order of magnitude higher still \\citep{Tollerud-2008, Koposov-2008}. Nevertheless, this is still much smaller than the total number of dark matter subhaloes found in high-resolution simulations of the standard $\\Lambda$CDM cosmology \\citep[e.g.][]{Klypin-1999, Moore-1999, Diemand-2007, Springel-2008}. This difference has become known as the ``Missing Satellites Problem''. It may only be an apparent discrepancy, however, since it can be removed if one accounts for the fact that not all low-mass subhaloes must contain stars, and those that do may have very high mass-to-light ratios. Several astrophysical mechanisms have been suggested that can lead to a number of visible satellite galaxies similar to that observed. Perhaps haloes were able to form a few stars initially, but the baryonic components of all haloes below some critical mass were subsequently destroyed by supernova feedback \\citep[e.g.][]{Larson-1974, Dekel-1986, Ferrara-2000}. Alternatively (or perhaps additionally) photoionisation may have prevented star formation in the smallest haloes \\citep[e.g.][]{Efstathiou-1992, Somerville-2002, Hoeft-2006, Simon-2007}. As \\cite{Sawala-2010} have shown, these two mechanisms can combine to produce very high mass-to-light ratios in haloes of $10^9 \\Ms$ and below, perhaps reconciling the number of very faint dwarf galaxies produced in $\\Lambda$CDM simulations with the observations. In this work, we turn our focus to more massive dwarf galaxies, and follow the evolution of the objects that form in haloes of $10^{10}\\Ms$. Our initial conditions are based on six haloes selected from the Millennium-II Simulation \\citep[MS-II,][]{Boylan-Kolchin-2009}, and resimulated at high resolution using smoothed particle hydrodynamics (SPH). Our simulations include cooling and star formation, supernova feedback, metal-enrichment and a cosmic UV background. Starting at redshift $z=49$, we are able to follow the formation of each individual halo and its central galaxy in their full cosmological context, all the way to $z=0$. On the other hand, the large volume of our parent simulation allows us to verify that our sample of resimulated haloes is representative of haloes of similar mass, and to predict a stellar mass~-- halo mass relation that can be tested against observation. With a box size of 137~Mpc and a mass resolution of $9.4\\times10^6\\Ms$, the MS-II has sufficient dynamic range to capture the statistics of the assembly of dark matter haloes between $10^9$ and $10^{14}\\Ms$. By comparing its halo/subhalo mass function to the observed SDSS stellar mass function of \\cite{Li-2009}, \\cite{Guo-2010} derived a typical mass-to-light ratio for each halo mass. This analysis assumes a monotonic relationship between halo mass and galaxy mass with relatively small scatter, but does not rely on any other assumptions about the processes involved in galaxy formation. We use its result to test the viability of our simulations and the underlying physical model as a description of the formation of ``typical'' $\\Lambda$CDM dwarf galaxies. The present work constitutes the first direct comparison of high resolution, hydrodynamical simulations of individual dwarf galaxies with the observed abundance of such objects. We combine the ability to follow star formation self-consistently in individual objects with the ability to draw conclusions about the general population of dwarf galaxies.\\\\ This paper is organised as follows: We begin in Section~\\ref{sec:previous} by reviewing the current status of simulations of the formation of dwarf galaxies. Section~\\ref{sec:ICs} describes the selection of haloes for resimulation and the generation of our high resolution initial conditions, while the numerical methods of our hydrodynamic simulations are discussed briefly in Section~\\ref{sec:methods}. In Section~\\ref{sec:evolution}, we show results for six haloes of final mass $10^{10}\\Ms$, and compare the properties of the galaxies to previous work, and to observation. In Section~\\ref{sec:sams}, we consider the predictions of our simulations for the stellar mass~-- halo mass relation and discuss the discrepancy with that inferred from comparing the observed stellar mass function to the halo abundance in $\\Lambda$CDM simulations. We conclude with a summary and interpretation of our results in Section~\\ref{sec:summary}. Unless stated otherwise, where we refer to the mass of a {\\it galaxy}, we mean the stellar mass M$_\\star$, whereas the mass of a {\\it halo} includes the total dynamical mass enclosed within r$_{200}$, the radius that defines a spherical overdensity 200 times the critical density of the universe. When quoting the results for our own simulations, we always use physical mass units of $\\Ms$, assuming h~$=0.73$. \\begin{figure*} \\vspace{-.2in} \\begin{center} \\begin{tabular}{lcccc} \\hspace{-.2in} \\includegraphics*[scale = .415]{scatter_resim_211262.eps} & \\hspace{-.5in} \\includegraphics*[scale = .415]{scatter_m2_211262_s.eps} & \\hspace{-.5in} \\includegraphics*[scale = .415]{scatter_m2_211262_l.eps} \\end{tabular} \\end{center} \\vspace{-.3in} \\caption{Comparison of Halo~4 at $z=0$ in a pure dark matter resimulation and in the parent Millennium-II Simulation. The left panel shows the position of $0.5\\%$ of the particles in a box of sidelength 1~Mpc in the resimulation, while the central panel shows the position of all particles within the same region in the MS-II. The panel on the right shows all particles in a box of 5~Mpc in the MS-II, with Halo~4 in the centre. All three panels are centred on the same absolute coordinates for the parent box of sidelength 137~Mpc, showing the position of the halo to be in perfect agreement. The FoF mass of the halo also agrees to within less than $1\\%$. A comparison of the left and the central panel reveals the additional substructure resolved in the resimulation.\\label{fig:compare_m2}} \\end{figure*} ", "conclusions": "\\label{sec:summary} We have performed high-resolution hydrodynamical simulations of six $\\sim 10^{10}\\Ms$ haloes, extracted from a large, cosmological parent simulation. We find that differences in merger histories lead to the formation of dwarf galaxies with different star formation histories and final stellar masses between $4.9\\times 10^7$ and $10^8\\Ms$. These stellar masses agree with previous simulations of similar mass haloes, and the structure of our simulated galaxies resembles that of observed galaxies of similar stellar mass, to the extent which we can resolve structure in our simulations. However, all these simulations imply an efficiency of conversion of baryons into stars which is at least an order of magnitude larger than that which is required to explain the observed abundance of dwarf galaxies in a $\\Lambda$CDM universe. While current hydrodynamical simulations, including our own, are consistent with almost arbitrarily high mass-to-light ratios for the faintest galaxies in haloes of $10^9\\Ms$ or less, they thus appear to be inconsistent with the mass-to-light ratios of larger dwarf galaxies, even when a moderately steep faint-end slope of the stellar mass function is assumed. The current recipes for mechanisms such as UV heating and supernova feedback appear sufficient to remove the ``Missing Satellites Problem'' for the smallest satellites. However, isolated dwarf galaxies with stellar masses of $10^{8}\\Ms$ are still substantially overproduced in current hydrodynamical simulations, even when these mechanisms are included. Our results suggest three possible explanations: The current observational count of dwarf galaxies could be incomplete, underestimating the true number density of $10^8\\Ms$ galaxies by a factor of four or more. In that case, the hydrodynamical simulations could be correct, but the semi-analytical models that produce low abundances of dwarf galaxies have been tuned to incorrect data. If the count of dwarf galaxies is almost complete at $10^8\\Ms$, these galaxies must, in a $\\Lambda$CDM universe, be residing in haloes significantly more massive than $10^{10}\\Ms$, and all current hydrodynamical simulations overpredict the efficiency of star formation by more than a factor of ten. This could be an indication of numerical problems, or, more likely, of incorrect or incomplete assumptions about the relevant astrophysics. Several possible mechanisms may contribute to a star formation efficiency in current simulations that is too high compared to real galaxies: \\begin{itemize} \\item Supernova feedback may be more efficient in ejecting gas from dwarf galaxies than current hydrodynamical simulations predict. For example, \\citet[b]{Guo-SAM} showed that the observed stellar mass function can be reproduced in semi-analytic models by assuming very strong mass-loading of winds in low mass haloes. \\item The full effect of reionisation on the IGM may not be captured in current models. As a result, cooling times may be underestimated, and the fraction of gas-rich mergers overestimated. Local sources of extreme UV and soft X-ray radiation may also ionise the interstellar medium, inducing another self-regulation mechanism for star formation \\citep{Ricotti-2002, Cantalupo-2010}. \\item Low dust content may lead to less efficient cloud formation and shielding at low metallicities. The transformation rate of cold gas into stars, currently assumed to be universal, may therefore be overestimated in dwarf galaxy simulations \\citep{Gnedin-2008}. \\item Processes such as magnetic fields, cosmic rays and the feedback from population-III stars are not included in any of the current models, and may further reduce the star formation efficiency. \\end{itemize} \\noindent Any revised model would however still have to reproduce features of individual galaxies consistent with observations. We also note that a model which substantially decreases the number of $10^8\\Ms$ galaxies would imply that the halo masses of fainter dwarf galaxies would need to be revised upwards, as some of these would now be required to live in $10^{10}\\Ms$ haloes. If the observed stellar mass function is complete, and the hydrodynamical simulations correctly capture the relevant physics of galaxy formation, the Millennium-II Simulation (and similar $\\Lambda$CDM simulations) overpredict the number of $10^{10}\\Ms$ dark matter haloes. This would seem to require the underlying physical assumptions of the $\\Lambda$CDM model to be revised. Warm Dark Matter may offer a possibility, but only for particle masses of $\\sim1$~keV, below the limit apparently implied by recent Lyman-$\\alpha$ observations. Of the three proposed scenarios, it appears that missing astrophysical effects in the simulations are the most likely cause of the discrepancy, and the most promising target in search of its resolution. While the three scenarios differ in nature, none is without significant implications for galaxy formation, which will have to be addressed in the future." }, "1003/1003.0447_arXiv.txt": { "abstract": "\\textit{Spitzer} spectroscopy has revealed that $\\simeq80$\\% of submm galaxies (SMGs) are starburst (SB) dominated in the mid-infrared. Here we focus on the remaining $\\simeq20$\\% that show signs of harboring powerful active galactic nuclei (AGN). We have obtained \\textit{Spitzer}-IRS spectroscopy of a sample of eight SMGs which are candidates for harboring powerful AGN on the basis of IRAC color-selection ($S_{8\\mu\\mathrm{m}}/S_{4.5\\mu\\mathrm{m}}>2$; i.e.~likely power-law mid-infrared SEDs). SMGs with an AGN dominating ($\\gsim50\\%$) their mid-infrared emission could represent the `missing link' sources in an evolutionary sequence involving a major merger. First of all, we detect PAH features in \\textit{all} of the SMGs, indicating redshifts from 2.5--3.4, demonstrating the power of the mid-infrared to determine redshifts for these optically faint dusty galaxies. Secondly, we see signs of both star-formation (from the PAH features) and AGN activity (from continuum emission) in our sample: 62\\% of the sample are AGN-dominated in the mid-infrared with a median AGN content of 56\\%, compared with $<30$\\% on average for typical SMGs, revealing that our IRAC color selection has successfully singled out sources with proportionately more AGN emission than typical SB-dominated SMGs. However, we find that only about 10\\% of these AGN dominate the bolometric emission of the SMG when the results are extrapolated to longer infrared wavelengths, implying that AGN are not a significant power source to the SMG population overall, even when there is evidence in the mid-infrared for substantial AGN activity. When existing samples of mid-infrared AGN-dominated SMGs are considered, we find that $S_{8\\mu\\mathrm{m}}/S_{4.5\\mu\\mathrm{m}}>1.65$ works well at selecting mid-infrared energetically dominant AGN in SMGs, implying a duty cycle of $\\sim15$\\% if all SMGs go through a subsequent mid-infrared AGN-dominated phase in the proposed evolutionary sequence. ", "introduction": "\\label{sec:intro} The era some 3\\,Gyrs after the Big Bang ($z\\sim 2$--2.5) coincides with a peak in activity in two important populations: QSOs, which represent accretion onto supermassive black holes (SMBHs); and a population of extremely luminous, but highly obscured galaxies (e.g.\\ \\citealt{Chapman05}; \\citealt{Wall08}). The bulk of the luminosity of these obscured galaxies is emitted in the rest-frame far-infrared waveband. As a result of redshifting, they are most directly selected through their emission in the submillimeter (submm) or mm wavebands, typically in the atmospheric windows around 850 or $1100\\,\\mathrm{\\mu m}$, and so are termed submm galaxies (SMGs). The infrared luminosities ($L_\\mathrm{IR}$) inferred for SMGs from their submm emission are highly uncertain, but assuming, as appears to be the case, that they follow the far-infrared--radio correlation for local starburst (SB) galaxies, then their typical luminosities will be $L_\\mathrm{IR}\\simeq10^{12}$--$10^{13}\\,\\mathrm{L}_\\odot$ (e.g.~\\citealt{Kovacs06}; \\citealt{Murphy09}). Thus this population may contain some of the most luminous galaxies in the Universe comparable in luminosity to QSOs. The increasing availability of precise redshifts for samples of SMGs (e.g.~\\citealt{Chapman05}; \\citealt{Eales09}) has allowed their properties to be studied in detail. SMGs are strongly clustered \\citep{Blain04}, massive ($M_\\star > 10^{11}$\\,M$_\\odot$, \\citealt{Borys05}; \\citealt{Hainline09}), and gas rich ($f_{gas}\\sim 0.3$, \\citealt{Frayer98}; \\citealt{Greve05}; \\citealt{Tacconi06,Tacconi08}) systems which are known to harbor (apparently) low-luminosity Compton-thin AGN (i.e.~$N_\\mathrm{H}<10^{24}$\\,cm$^{-2}$, \\citealt{Takata06}; \\citealt{Alexander05nat,Alexander05,Alexander08a}) and (apparently) strong star-formation (SF) activity ($\\mathrm{SFR}\\sim 1000$\\,M$_\\odot$\\,yr$^{-1}$; \\citealt{Swinbank04}; \\citealt{Chapman05}). Many of these properties, and the similarity between the redshift distributions of QSOs and SMGs \\citep{Chapman05}, support a link between SMGs, QSOs and the formation phase of massive elliptical galaxies (e.g.~\\citealt{Lilly99}; \\citealt{Archibald02}; \\citealt{Stevens05}; \\citealt{Coppin08b}). In the high-redshift interpretation of the evolutionary sequence first presented by \\citet{Sanders88}, SMGs would trace an infrared ultraluminous phase followed by a short `transition phase' where the galaxy would display a mix of SF and obscured AGN activity before evolving into an optically luminous QSO. Studying the relative SF and AGN activity in these `transition' or `missing link' sources and performing a comparison to typical SF SMGs can thus provide important insight on the validity of this evolutionary sequence. Here we focus on the comparison of the mid-infrared spectral properties between existing samples of typical SMGs and a new sample of candidate `transition phase' SMGs. The benefit of undertaking this energy audit in the mid-infrared is that the spectra can provide measurements of both SF and AGN activity simultaneously (see \\citealt{Pope08} and \\citealt{KMD09}; hereafter \\citetalias{Pope08} and \\citetalias{KMD09}), providing a complementary approach to radio (e.g.~\\citealt{Ibar09b}) and X-ray studies of SMGs (e.g.~\\citealt{Alexander05}; \\citealt{Laird09}). The {\\it Spitzer} InfraRed Spectrograph (IRS; \\citealt{Houck04}) era enabled the study of the energetics of $\\approx45$ SMGs with $24\\,\\mathrm{\\mu m}$ flux densities as low as $\\sim0.1$\\,mJy and with redshifts as high as $z=2.6$, demonstrating that accurate redshifts for dust-enshrouded (and sometimes optically invisible) galaxies can be obtained (e.g.~\\citealt{Lutz05}; \\citealt{Valiante07}; \\citealt{MD07}; \\citetalias{KMD09}; \\citetalias{Pope08}). The majority of the IRS work has confirmed that SMGs are primarily SB-dominated systems, with hot dust continuum from an AGN contributing at most 30\\% of the mid-infrared luminosity. Only about 15\\% of blank-field SMGs appear to be continuum-dominated in the mid-infrared (i.e.~$>50$\\% contribution; \\citetalias{Pope08}), with InfraRed Array Camera (IRAC; \\citealt{Fazio04}) 8.0\\,$\\mu$m to 4.5\\,$\\mu$m color ratios of $S_{8}/S_{4.5}>2$. Mid-infrared continuum-dominated SMGs are potentially an important sub-population of SMGs representing the `missing link' sources in the proposed evolutionary sequence of \\citet{Sanders88}, but being the minority of this luminous population, have not yet been studied in a systematic or statistically robust way. Thus, an important question to address is: \\textit{Are the energetics of these composite SB/AGN objects with submm emission dominated by SBs or by AGN, even when the IRAC colors indicate that an AGN is likely present?} Other relevant IRS samples include near-infrared-selected SBs, X-ray-selected AGN \\citep{Weedman06}, and \\textit{Spitzer}-selected $0.3\\lsim z \\lsim 3$ ULIRGs (\\citealt{Yan07}; \\citealt{Sajina07}; \\citealt{Farrah08}; \\citealt{Dasyra09}; \\citealt{HC09}). AGN-dominated sources in these samples tend to reside in a distinct parameter space in \\textit{Spitzer} color-color diagrams, which is consistent with color-redshift evolution tracks of well-known local AGN (see e.g.~Fig.~\\ref{fig:selection} ; \\citealt{Ivison04}; \\citealt{Ashby06}; \\citealt{Hainline09}). The reason why $S_{8}/S_{4.5}>2$ should locate mid-infrared AGN-dominated sources is simple: seeing an enhanced 8\\,$\\mu$m flux density compared with 4.5\\,$\\mu$m is expected if a source has significant thermal power-law emission from an AGN accretion disk, which can dilute both the PAHs and the H-opacity minimum (1.6\\,$\\mu$m stellar bump, which these channels trace). Although some contamination is expected: SF-dominated sources can also show enhanced $S_{8}/S_{4.5}$ at $z\\gtrsim4$ (when these channels begin to sample over the peak of stellar photospheric emission) aswell as at $z\\lesssim1$ (since these channels are not yet climbing up the restframe 1.6\\,$\\mu$m stellar bump). The color cut is thus appropriate for separating mid-infrared SB- and AGN-dominated SMGs from $z\\simeq1$--4, which is well matched to the known redshift distribution of the SMG population. We have thus selected eight SMGs from the Submillimeter Common-User Bolometer Array (SCUBA; \\citealt{Holland99}) HAlf Degree Extragalactic Survey (SHADES; \\citealt{Mortier05}; \\citealt{Coppin06}; \\citealt{Austermann09}) within this relatively unexplored $S_{8}/S_{4.5}>2$ parameter space that are likely harboring AGN. These were targeted with the IRS in order to determine the relative contribution of power-law/AGN emission versus PAH/SF emission to their power output, enabling a comparison to similar IRS samples of more typical SF-dominated SMGs. We use these data to obtain independent redshift estimates as well as to test if the IRAC-color criterion is a secure means of pre-selecting SMG counterparts with an enhanced AGN component compared to typical SMGs -- the `missing link' sources in the proposed evolutionary sequence we wish to investigate. This paper is organized as follows. The sample selection, {\\it Spitzer}-IRS observations, data reduction, and analysis approach are described in \\S~\\ref{sec:obsdr}. In \\S~\\ref{sec:results} we present the main results of the IRS spectroscopy, including redshifts, spectral decomposition and AGN classification, and full SED fits to determine their total infrared luminosities. We discuss the implications that our results and those of other IRS SMG studies have on the role that SMGs play in galaxy evolution in the framework of the proposed evolutionary sequence in \\S~\\ref{sec:discuss}. Finally, our conclusions are given in \\S~\\ref{sec:concl}. Here, we discuss the observed `mid-infrared' spectral properties of our SMGs probed by the IRS which, at our source redshifts of $z>2.5$, traces $\\sim4$--10\\,$\\mu$m in the rest-frame. All magnitudes in this paper are on the AB system, unless otherwise stated. We adopt cosmological parameters from the \\textit{WMAP} fits \\citep{Spergel03}: $\\Omega_\\Lambda=0.73$, $\\Omega_\\mathrm{m}=0.27$, and $H_\\mathrm{0}=71$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$. \\begin{figure*} \\epsscale{.80} \\plotone{fig1.eps} \\caption{{\\it Spitzer} color-color diagram as an AGN diagnostic (see e.g.~\\citealt{Ivison04}; \\citetalias{Pope08}). The dotted and dot-dashed curves show the positions of Mrk231 (an AGN) and M82 (a SB), respectively, as a function of redshift (redshift is indicated by the numbers along the tracks). The smallest filled and open circles represent robust ($P<0.05$) and tentative ($P>0.05$) 24\\,$\\mu$m SMG counterparts, respectively, for our parent sample of blank-field SHADES SMGs \\citep{Ivison07}. For comparison, we have also plotted the GOODS-N SMG sample (blue squares; \\citealt{Pope06}), and also the \\citet{Chapman05} radio-selected spectroscopically-confirmed sample of SMGs (green triangles). SMGs from these parent samples that have been observed with the IRS (\\citetalias{Pope08}; \\citetalias{KMD09}; this work) are indicated by large versions of the corresponding symbols. Of these IRS-observed SMGs, those with $>50\\%$ contribution from continuum (AGN) emission to the mid-infrared luminosity are circumscribed and tend to lie far from the M82 SB sequence. Although our sample was initially selected on the basis of $S_{8}/S_{4.5}>2$ (thin dashed vertical line), when all the results from the literature are taken into account a selection of $S_{8}/S_{4.5}\\gtrsim1.65$ thus seems to pick out AGN fairly efficiently, with some small amount of scatter over this boundary (thick dashed vertical line; see \\S~\\ref{sec:discuss}). Extrapolating from this result implies that $\\approx15$\\% of blank-field SMGs will likely have an enhanced dominant AGN component in the mid-infrared.\\label{fig:selection}} \\end{figure*} ", "conclusions": "\\label{sec:concl} {\\it Spitzer} IRS spectroscopy has been obtained for a sample of eight SMGs brighter than 200$\\,\\mu$Jy at 24$\\,\\mu$m and selected to possess observed IRAC colors of $S_{8}/S_{4.5}>2$, in order to investigate the level of AGN dominance in the SMG population. Our conclusions are as follows: 1) Taking into account sources in the literature, we find a that a color-selection of $S_{8}/S_{4.5}>1.65$ is a better description overall for defining the boundary between SB and AGN-dominated SMGs, with a small amount of scatter across this division. 2) We see signs of both SF and AGN activity in our sample of SMGs, with a continuous distribution of AGN fractions in the mid-infrared. There are signs of SF in the form of PAH features in \\textit{all} of the SMGs in our sample, from which we derive redshifts between 2.5--3.4, demonstrating the power of the mid-infrared to determine redshifts when the optical counterparts are too faint to study with current facilities. 3) We find a median AGN contribution to the mid-infrared of 56\\% in our sample of SMGs, indicating that overall, SMGs with $S_{24}>0.2$\\,mJy and $S_{8}/S_{4.5}>2$ tend to have more dominant AGN-components in the mid-infrared than typical SMGs ($<30$\\% mid-infrared AGN contribution). Extrapolation to the far-infrared reveals that the AGN is bolometrically important only in two of our SMGs (each with a $>75$\\% mid-infrared AGN contribution). This result suggests that significant contamination from AGN to the far-infrared luminosities affects $\\lsim5$\\% of the SMG population overall. 4) To first order, these AGN range from being fairly low-luminosity types found in typical SMGs to more extreme cases that likely overlap with AGN-dominated $24\\,\\mu$m-selected ULIRGs in \\citet{Sajina07}. In any case, all our AGN appear to be obscured to some degree (at least five have low $L_\\mathrm{X}/L_\\mathrm{6\\mu m}$ and six are optically faint with $R>23$) and a small subset of our sample may be obscured by Compton-thick material. 5) Our results are thus consistent with the \\citet{Sanders88} evolutionary scenario, with all SMGs undergoing a `transitional' AGN-dominated phase with a duty cycle of $\\simeq15$\\%. Our sample of AGN-dominated SMGs could be at a slightly later stage of evolution than SF-dominated SMG systems, with the SF still occurring but where the AGN has now begun to heat the dust appreciably in the SMG as the BH undergoes a period of rapid growth." }, "1003/1003.5112_arXiv.txt": { "abstract": "% Multiline techniques assuming similar line profiles have become a standard tool in stellar astronomy for increasing the signal-to-noise ratio (SNR) of spectropolarimetric measurements. However, due to the widely-used weak field approximation their benefits could not so far be used for solar observations, where a large variety of Stokes profiles emerge from local magnetic fields and measuring weak fields in the quiet Sun remains a challenge. The method presented here permits us to analyze many lines with arbitrary Zeeman splitting and to simultaneously deploy Stokes $IQUV$ spectra to determine a common line profile with the SNR increased by orders of magnitude. The latter provides a valuable constraint for determining separately field strengths for each contributing absorber. This method represents an extension of our recently developed technique of Nonlinear Deconvolution with Deblending \\citep[NDD,][]{sennhauseretal2009}, which accounts for the nonlinearity in blended profiles. Equipped with all those abilities, ZCD is the perfect tool to further increase the informative value of high-precision polarimetric observations. ", "introduction": "Magnetic fields in the quiet Sun being ubiquatuous but extremely weak and mixed in polarities remain a challenge to detect and study. They have drawn recently more attention, first, because they are thought to play an important role in solar dynamo and heating the chromosphere and, second, because the Sun is so unusually quiet during the current minimum. A coherent picture is however still missing. For instance, measurements of the internetwork magnetic field strength by different methods differ by an order of magnitude, from roughly 10$\\,$G up to several hundred G (\\citeauthor{litesetal2008} \\citeyear{litesetal2008}, \\citeauthor{suarezetal2007} \\citeyear{suarezetal2007}). This is obviously due to the limited accuracy and/or spatial resolution of current spectropolarimetric observations, which at best achieve the noise level of $10^{-3}$ at the resolution of 0.3 arcsec, as with the SP instrument on Hinode. Here we propose a method which can boost the signal-to-noise ratio (SNR) of spectropolarimetric observations by orders of magnitude and, thus, is capable of providing a more sensitive constraint on very weak internetwork magnetic fields. Our method is based on a new technique called Nonlinear Deconvolution with Deblending \\citep[NDD,][]{sennhauseretal2009}, which extracts a common Stokes profile from many spectral lines while accounting for the nonlinearity in blends. ", "conclusions": "We have demonstrated by finding the common line pattern for Stokes $I$ and $V$ simultaneously, that our ZCD code is capable of identifying Zeeman signatures which are far below the noise level of the original spectrum. It directly returns reliable magnetic field strengths, as well as a common line profile containing the physical conditions of the line forming atmosphere. A full implementation of the method \\citep{sennhauserberdyuginainprep} enables the ZCD to recover a magentic field \\emph{vector}, i.e., strength and orientation, by inverting a full set of Stokes parameters. There, magnetooptical effects are included as well, in order to take into account linear polarization signals originating from anomalous dispersion. In addition, our method can be successfully applied to molecular bands despite severe blending of many lines \\citep{sennhauseretal2009}. Furthermore, an extension of the ZCD for the (partial) Paschen-Back regime enables investigation of multi-Tesla magnetic objects using both atomic and molecular lines. While being excellent tracers of magnetic activity in cooler regions, molecular lines demonstrate departures from the Zeeman regime at relatively low magnetic field strengths, $\\sim 100$ G (\\citeauthor{berdyuginaetal2005} \\citeyear{berdyuginaetal2005}) which makes it crucial for them to be treated in the Paschen-Back regime." }, "1003/1003.2504_arXiv.txt": { "abstract": "We report for the first time general geometrical expressions for the angular resolution of an arbitrary network of interferometric gravitational-wave (GW) detectors when the arrival-time of a GW is unknown. We show explicitly elements that decide the angular resolution of a GW detector network. In particular, we show the dependence of the angular resolution on areas formed by projections of pairs of detectors and how they are weighted by sensitivities of individual detectors. Numerical simulations are used to demonstrate the capabilities of the current GW detector network. We confirm that the angular resolution is poor along the plane formed by current LIGO-Virgo detectors. A factor of a few to more than ten fold improvement of the angular resolution can be achieved if the proposed new GW detectors LCGT or AIGO are added to the network. We also discuss the implications of our results for the design of a GW detector network, optimal localization methods for a given network, and electromagnetic follow-up observations. ", "introduction": "\\label{intro} Several types of astrophysical sources are expected to be detectable both in gravitational waves (GWs) and in conventional electromagnetic (EM) wavelengths. For example, long gamma-ray bursts have been conjectured to originate from asymmetric core collapse of massive stars, and short gamma-ray bursts might be produced by the coalescence of compact binary objects containing neutron stars. Both of these could emit gravitational waves in the frequency band of ground-based laser interferometer GW detectors (e.g., Ref.~\\cite{grb_gw}). Several large-scale interferometric GW detectors have reached (or approached) their design sensitivity, and are coordinating to operate as a global array. These include the LIGO detectors at Livingston, Louisiana, and Hanford, Washington, US, the Virgo detector in Pisa, Italy, the GEO\\,600 detector in Hannover, Germany, and the TAMA\\,300 detector in Tokyo Japan. Upgrades to existing detectors (Advanced LIGO and Advanced Virgo) have been planned \\cite{ligo,virgo,geo}, while new detectors (LCGT in Japan \\cite{lcgt} and AIGO in Australia \\cite{aigo}) are still being proposed. In case of a strong EM event, follow-up searches for GW signals can be conducted in archived data in the time window of the event (e.g., Ref.~\\cite{ligo_grb}). On the other hand, EM follow-ups to probable GW events require a clear understanding of the angular resolution of an array of GW detectors. The angular resolution of an individual GW detector, arising from its antenna beam pattern, is rather poor \\citep{thorne87}. However, the large baselines of the current GW-detector network facilitate better angular resolution via triangulation. Several localization algorithms have been proposed and the effect of arrival timing uncertainties as well as amplitude information of GWs have been investigated \\cite{tinto89, wen05a, cavalier06, beauville06, acernese07, wen07a, wen_fan, markowitz08,searle08,searle09, steve09}. Quantitative studies of the angular resolution of a network of GW detectors have been conducted by several authors, both for a ground-based detector network and for the future space GW detector LISA \\citep{krolak94,cutler98,Pai01,leor04}. A standard approach is to calculate numerically the Fisher information matrix, which leads to a method-independent lower bound on the statistical errors of estimated parameters (see a review in Ref.~\\cite{krolak05}). On the other hand, explicit analytical expressions for the network angular resolution are rare in the literature largely because of the complexity involved in derivations. Two approximate analytical expressions for the angular resolution can be found in the literature (summarized in \\citep{sylvestre03}) for a network of three GW detectors. One is an elegant approximate geometrical formula for 3 detectors due to Thorne (as cited in Eq.(8.3) of Ref.~\\cite{tinto89}): the solid angle uncertainty is \\begin{equation} \\Delta \\Omega = \\frac{2c^2 \\Delta \\tau_{12}\\Delta \\tau_{13}}{A \\cos \\theta}\\,, \\end{equation} where $c$ is the speed of light, $\\Delta \\tau_{12}$ and $\\Delta \\tau_{13}$, are time-of-arrival accuracy between pairs of detectors, $A$ is the area formed by the three detectors, and $\\theta$ is the angle between the source direction and the normal to the plane of the three detectors. However, the underlying assumptions and derivation of this expression are not available in the literature. The dependence of angular resolution on the signal-to-noise ratio (SNR) was derived by Tinto in Ref.~\\citep{tinto89}, by expressing the above time-of-arrival accuracy as a function of SNR and frequency~\\citep{schutz89} (derived from the Fisher matrix assuming all other information of the waveform as perfectly known). The other formula is based on the numerical result of the angular resolution of the LIGO-Virgo network around a wave incident direction normal to the plane formed by the three detector sites~\\citep{sylvestre03} --- for GWs emitted from neutron star-neutron star (NS-NS) inspirals~\\citep{Pai01}. This particular resolution was then rescaled by the cosine of the wave incidence angle and SNR~\\citep{sylvestre03}. Analytical geometrical expressions or approximate ones for the angular resolution for an array of more than three detectors have not been obtained in the literature. In this paper, we deduce explicit analytical expressions for the angular resolution of an arbitrary GW detector network in terms of observables such as cross-sectional areas of the network and energy flux of the incoming GW. We use only the time-of-arrival information, ignoring additional (usually rather poor) information from the directional-derivatives of antenna beam pattern functions --- and therefore arrive at a conservative estimate. Such an approximation allows us to obtain expressions that have explicit geometrical meanings, further generalizing Thorne's formula to an arbitrary number of detectors, and several particular scenarios. In particular, we consider both {\\it short signals,} during which motion of the detector network is negligible, and {\\it long signals,} for which the trajectory traced by the detectors during the signal determines the effective size of the detector network. We also consider signals with {\\it known} or {\\it unknown} waveforms, but always assume unknown arrival times of the signal. In this paper, the scenarios where signals have completely {\\it known} and completely {\\it unknown} waveforms are termed interchangeably as the {\\it best-case } and the {\\it worst-case} scenarios respectively. We focus on deriving explicit expressions for several situations that will arise in the practice of searching for and localizing GWs. Specifically, we derive general expressions for {\\it short signals} assuming only that arrival time is one of the unknown parameters (summarized in Eq.~\\eqref{Omega} and text thereafter) and for {\\it long signals} assuming known waveform (Eq.~\\eqref{long_general}). Based on these formulae, we show simplified solutions for Eq.~\\eqref{Omega} in several realistic situations: (1) {\\it short signals} in the {\\it worst-case} scenario (Eq.~\\eqref{worst}) and in the {\\it best-case} scenario (Eq.~\\eqref{best}) for an arbitrary network of detectors, (2) special cases of the two- and three- detector networks in the {\\it best case} (Eq.~\\eqref{theta_2}--\\eqref{ang_3det}) and their representations when the wave is short and monochromatic (Eq.~\\eqref{theta_2f}--\\eqref{omega_3}), (3) {\\it long signals} in short and long observations with detectors in circular motion (Eq.~\\eqref{pulsar_short} and Eq.~\\eqref{pulsar_long} respectively). This paper is organized as follows. In section~\\ref{principle}, we explain our notation. In section~\\ref{expression}, we derive analytical expressions of the angular resolution for an arbitrary detector network and for special cases. We show explicit derivations for the {\\it worst-case} and the {\\it best-case} scenario in section~\\ref{sec_worst} and section ~\\ref{sec_best} and then derive a general expression in section~\\ref{short_general}. In section~\\ref{sec_long}, we derive a general expression for long-duration wave and its application to detectors at circular motion. We then discuss the implications for the design of a detector network and localization strategies in section~\\ref{implication}. The astrophysical applications of our results are shown in section~\\ref{astro}. In section~\\ref{sec_antenna}, we discuss the possible errors in our estimation caused by ignoring the directional derivatives of antenna beam patterns. Our results are summarized in section~\\ref{concl}. ", "conclusions": "\\label{concl} In summary, we have derived for the first time analytical expressions for the angular resolution for an arbitrary network of gravitational-wave detectors where the directional derivatives of antenna-beam pattern functions are ignored. Our results demonstrate the explicit dependence of the angular resolution on the geometrical configuration of the network, the total noise-weighted energy flux coupled to the network, its fractional distribution to individual detectors, and correlation of data between detectors. These dependences are intrinsic to the configuration of the detector network and the sources. The results are method-independent and they correspond to the best possible localization any unbiased methods can achieve in the presence of statistical errors. Our estimate is conservative especially for low-frequency sources as the directional derivatives of antenna beam pattern functions are ignored (see sec.~\\ref{sec_antenna}). In reality, careful designs, e.g., to remove multiple local minima or maxima caused by interference when combining data from different detectors, or to break the mirror degeneracy in arrival time delays for the three-detector case (by using the wave amplitude information), are required for localization methods in order to achieve the ``best'' limits. Derivation of the angular resolution for {\\it short signals} including those for a network of two and three detectors are presented in Eq.~\\eqref{theta_2} and Eq.~\\eqref{ang_3det}. Our results are consistent with what was previously known from the diffraction limit: that a larger network yields better angular resolution. We confirm that the angular resolution is poor along the plane formed by current LIGO-Virgo detectors and is better for directions normal to the plane. Numerical results are included to show how a new detector in Japan (LCGT) or in Australia (AIGO) can dramatically improve the angular resolution of the existing network (Fig.~\\ref{LHVC_cum} and Fig.~\\ref{LHVA_cum} and discussions thereafter), by contributing to longer baselines, additional energy flux, extended null signal space, and by breaking the plane-degeneracy formed by three detectors. Compared with previous approximate expressions for 3-detectors (sec.~\\ref{intro}), our results are more rigorous and show the explicit roles of individual detectors. The angular resolution of a detector network depends on the fractional energy flux coupled to individual detectors in the {\\it best case} and on correlation of data between detectors in the {\\it worst case}. In the {\\it best-case} scenario, the angular resolution is largely limited by the least sensitive detector. These have significant implications for the design of a detector network and for the design of an optimal localization method for a fixed network (see sec.~\\ref{implication} for more discussion on implications). Moreover, our results apply to an arbitrary network of any number of detectors. We have also derived a geometrical expression for {\\it long signals} for a simplified case where waveforms are known (Eq.~\\eqref{long_general}). The situation where detectors are in circular motion and the signal is monochromatic is discussed. The dependence of the angular resolution on the latitude of the source is apparent in our formulae. We also demonstrate that the angular resolution improves rapidly with the observing time $\\sim T^4$ initially with short observations and saturates to $\\sim T$ for longer observations (Eq.~\\eqref{pulsar_short}, Eq.~\\eqref{pulsar_long}), consistent with previous knowledge~\\cite{schutz89, prix07}. We have further presented through numerical simulations the distribution of the areas of error ellipses at 95\\% confidence level for two of the most important GW sources for ground-based detectors. (1) We apply our calculations to GWs from coalescing binaries of neutron stars using the sensitivity curves of detectors that are operating at this writing. The actual limit of the angular resolution for these inspiral sources should be closer to the {\\it best-case} scenario since theoretical waveforms are known and essential parameters can be estimated with great accuracy independent of source direction determination \\cite{cutler94}. (2) We apply our method to burst-like GWs using a representing waveform from bar-instability of neutron stars in supernovae for advanced detectors. The actual angular resolution for this type of ``burst'' source fits in our {\\it worst-case} scenario for {\\it short signals}. We shows that, for the existing LIGO-Virgo detector network, assuming uniform distribution of sources, at an optimal network SNR of around 15, 50\\% \\ of inspiral sources can be located within 23 sq-degs ({\\it best case}) at the 95\\% confidence level. For the burst source, without any knowledge of the waveform, at SNR of 10, 50\\% of the sources can be localized within 50 sq-degs ({\\it worst-case}), but it can be reduced to 8 sq-degs if we have predicted waveforms available (e.g., from Ref. \\cite{mario, ott}). Results for the initial or advanced detectors are similar. Our results imply that, for prompt follow-up electromagnetic observations directly using triggers from current GW network, wide-field telescopes are desirable." }, "1003/1003.3925_arXiv.txt": { "abstract": "We propose the most general modified first-order Ho\\v{r}ava-Lifshitz gravity, whose action does not contain time derivatives higher than the second order. The Hamiltonian structure of this theory is studied in all the details in the case of the spatially-flat FRW space-time, demonstrating many of the features of the general theory. It is shown that, with some plausible assumptions, including the projectability of the lapse function, this model is consistent. As a large class of such theories, the modified Ho\\v{r}ava-Lifshitz $F(R)$ gravity is introduced. The study of its ultraviolet properties shows that its $z=3$ version seems to be renormalizable in the same way as the original Ho\\v{r}ava-Lifshitz proposal. The Hamiltonian analysis of the modified Ho\\v{r}ava-Lifshitz $F(R)$ gravity shows that it is in general a consistent theory. The $F(R)$ gravity action is also studied in the fixed-gauge form, where the appearance of a scalar field is particularly illustrative. Then the spatially-flat FRW cosmology for this $F(R)$ gravity is investigated. It is shown that a special choice of parameters for this theory leads to the same equations of motion as in the case of traditional $F(R)$ gravity. Nevertheless, the cosmological structure of the modified Ho\\v{r}ava-Lifshitz $F(R)$ gravity turns out to be much richer than for its traditional counterpart. The emergence of multiple de Sitter solutions indicates to the possibility of unification of early-time inflation with late-time acceleration within the same model. Power-law $F(R)$ theories are also investigated in detail. It is analytically shown that they have a quite rich cosmological structure: early/late-time cosmic acceleration of quintessence, as well as of phantom types. Also it is demonstrated that all the four known types of finite-time future singularities may occur in the power-law Ho\\v{r}ava-Lifshitz $F(R)$ gravity. Finally, a covariant proposal for (renormalizable) $F(R)$ gravity within the Ho\\v{r}ava-Lifshitz spirit is presented. ", "introduction": "Recently, it has become clear that our universe has not only undergone the period of early-time accelerated expansion (inflation), but also is currently in the so-called late-time accelerating epoch (dark energy era). An extremely powerful way to describe the early-time inflation and the late-time acceleration in a unified manner is modified gravity. This approach does not require the introduction of new dark components like inflaton and dark energy. The unified description of inflation and dark energy is achieved by modifying the gravitational action at the very early universe as well as at the very late times (for a review of such models, see \\cite{Nojiri:2006ri}). A number of viable modified gravity theories has been suggested. Despite some indications to possible connection with string/M-theory \\cite{Nojiri:2003rz}, such theories remain to be mainly phenomenological. It is a challenge to investigate their origin from some (not yet constructed) fundamental quantum gravity theory. Among the recent attempts to construct a consistent theory of quantum gravity much attention has been paid to the quite remarkable Ho\\v{r}ava-Lifshitz quantum gravity \\cite{Horava:2009uw}, which appears to be power-counting renormalizable in four dimensions. In this theory the local Lorentz invariance is abandoned, but it is restored as an approximate symmetry at low energies. Despite its partial success as a candidate for fundamental theory of gravity, there are a number of unresolved problems related to the detailed balance and projectability conditions, consistency, its general relativity (GR) limit, realistic cosmological applications, the relation to other modified gravities, etc. Due to the fact that its spatially-flat FRW cosmology \\cite{cosmology} is almost the same as in GR, it is difficult to obtain a unified description of the early-time inflation with the late-time acceleration in the standard Ho\\v{r}ava-Lifshitz gravity. Recently the modified Ho\\v{r}ava-Lifshitz $F(R)$ gravity has been proposed \\cite{Chaichian:2010yi}. Such a modification may be easily related with the traditional modified gravity approach, but turns out to be much richer in terms of the possible cosmological solutions. For instance, the unification of inflation with dark energy seems to be possible in such Ho\\v{r}ava-Lifshitz gravity due to the presence of multiple de Sitter solutions. Moreover, on the one hand, there is the hope that the generalization of Ho\\v{r}ava-Lifshitz gravity may lead to new classes of renormalizable quantum gravity. On the other hand, one may hope to formulate the dynamical scenario for the Lorentz symmetry violation/restoration, caused by the expansion of the universe, in terms of such generalized theory. In the present work (section \\ref{sec:2}) we propose the most general modified first-order Ho\\v{r}ava-Lifshitz-like theory, without higher derivative terms which are normally responsible for the presence of ghosts. The general form of the action in the spatially-flat FRW space-time is found, and the Hamiltonian structure of the action is analyzed in section \\ref{sec:3}. As a specific example of such a first-order action we introduce the modified Ho\\v{r}ava-Lifshitz $F(R)$ theory which is more general than the model of ref.~\\cite{Chaichian:2010yi}. Nevertheless, its spatially-flat FRW cosmology turns out to be the same as for the model \\cite{Chaichian:2010yi} (this is not the case for BH solutions, etc). Therefore it also coincides with the conventional $F(R)$ spatially-flat cosmology for a specific choice of the parameters. The ultraviolet structure of the new Ho\\v{r}ava-Lifshitz $F(R)$ gravity is carefully investigated. It is shown that such models can have very nice ultraviolet behaviour at $z=2$. Moreover, for $z=3$ a big class of renormalizable models is suggested (section \\ref{sec:2}). The Hamiltonian analysis of the modified Ho\\v{r}ava-Lifshitz $F(R)$ gravity is presented in section \\ref{sec:4}. The fixed gauge modified Ho\\v{r}ava-Lifshitz $F(R)$ gravity is analyzed in section \\ref{sec:5}. Section \\ref{sec:6} is devoted to the investigation of spatially-flat FRW cosmology for power-law $F(R)$ gravity. The general equation for the de Sitter solutions is obtained. It acquires an extremely simple form for a special choice of parameters, when de Sitter solutions are roots of the equation $F=0$. The existence of multiple de Sitter solutions indicates the principal possibility of attaining the unification of the early-time inflation with the late-time acceleration in the modified Ho\\v{r}ava-Lifshitz $F(R)$ gravity. The reconstruction technique is developed for the study of analytical and accelerating FRW cosmologies in power-law models. A number of explicit analytical solutions are presented. It is shown by explicit examples that some of the quintessence/phantom-like cosmologies may develop the future finite-time singularity of all the known four types, precisely in the same way as for traditional dark energy models. The possible curing of such singularities could be achieved in a similar way as in the case of traditional modified gravity. Some remarks about small corrections to the Newton law are made in section \\ref{sec:7}. A summary and outlook are given in the last section \\ref{sec:8}. In the appendix \\ref{appendix} we propose a covariant $F(R)$ gravity that is quite similar to the corresponding Ho\\v{r}ava-Lifshitz version but remains to be a covariant theory. It seems that it could also be made renormalizable. ", "conclusions": "\\label{sec:8} We have proposed a first-order modified Ho\\v{r}ava-Lifshitz-like gravity action and studied its Hamiltonian structure. As a large explicit class of such models we considered the modified Ho\\v{r}ava-Lifshitz $F(R)$ gravity that is more general than the one introduced in ref.~\\cite{Chaichian:2010yi}, which for the special choice of parameter $\\mu=0$ coincides with the degenerate model introduced in ref.~\\cite{Kluson:2009xx}. Its ultraviolet properties are discussed and it is demonstrated that such $F(R)$ gravity may be renormalizable for the case $z=3$ in a similar way as the original proposal for Ho\\v{r}ava-Lifshitz gravity. The Hamiltonian analysis of the proposed modified Ho\\v{r}ava-Lifshitz $F(R)$ gravity shows that this theory is generally consistent with reasonable assumptions. The $F(R)$ gravity action has also been analyzed in the fixed gauge form, where the presence of the extra scalar is particularly illustrative. The methods presented in the Hamiltonian analyses of sections \\ref{sec:3} and \\ref{sec:4} can be used to study any action of the general form (\\ref{HLF26}). The spatially-flat FRW cosmology of the modified Ho\\v{r}ava-Lifshitz $F(R)$ gravity is studied. It is shown that it coincides with the one of the earlier model \\cite{Chaichian:2010yi}, but only in the spatially-flat FRW case. For specific choice of the parameters of the theory, its FRW equations of motion coincide with the ones of the traditional $F(R)$ gravity. The presence of the multiple de Sitter solutions shows the principal possibility of the unification of the early-time inflation with the late-time acceleration in the Ho\\v{r}ava-Lifshitz background, which proves that it can have rich cosmological applications. The power-law theories are investigated in detail. A number of analytical FRW solutions is found, including the ones with behavior relevant for the early/late cosmic acceleration. The quintessence/phantom-like cosmologies derived in our work may show all the four possible types of finite-time future singularities like in the case of standard dark energy. The conditions to cure such future singularities are discussed in analogy with the traditional $F(R)$ gravity. It is also interesting that the correction to the Newton law in the $F(R)$ gravity under discussion can be made unobservably small. Finally, a covariant proposal for $F(R)$ gravity in Ho\\v{r}ava-Lifshitz spirit has been made. Despite some successes in the formulation of modified Ho\\v{r}ava-Lifshitz $F(R)$ gravity which can be made renormalizable and in its cosmological applications, a number of unsolved questions remain. What is the appropriate way to introduce matter in the theory? Is the theory itself fundamental (or at least, fully consistent) or does it descend from another more fundamental proposal? Can it comply with all the local tests in the Solar system as well as with cosmological bounds? What is the dynamical scenario for the restoration of the Lorentz invariance at late times? What are the cosmological and astrophysical consequences of the first-order modified Ho\\v{r}ava-Lifshitz gravity when compared with those of the traditional modified gravity \\cite{tradModGrav}. Moreover, the traditional questions about the properties of black holes in such a theory can be straightforwardly investigated. Nevertheless, even at the present stage some surprises can be expected from the theory. While the universe has likely undergone a perioid of inflation in its early moments, it is interesting to note that Ho\\v{r}ava-Lifshitz gravity could produce cosmological perturbations that are almost scale-invariant even without inflation \\cite{Mukohyama:2009gg}. Ho\\v{r}ava-Lifshitz gravity has also been considered in the presence of scalar fields \\cite{Chen:2009ka,Lee:2010iu}. In principle, it is possible to extend our Ho\\v{r}ava-Lifshitz $F(R)$ gravity by including its coupling with scalar fields. We would also like to mention a recent paper \\cite{Kluson:2010aw}, where a new class of Lorentz-invariance breaking non-relativistic string theories, inspired by the Ho\\v{r}ava-Lifshitz gravity, has been presented and analyzed. Using the $F(R)$ version of gravity one can propose even a more general formulation of string theory in the Ho\\v{r}ava-Lifshitz background: for instance, rigid strings, membranes and $p$-branes, etc. On the other hand, it may suggest unusual solutions for the known cosmological problems. There also exists an attempt to explain the homogeneity of our universe in a model with varying speed of light \\cite{Albrecht:1998ir}. Having in mind that in the ultraviolet region the speed of the Ho\\v{r}ava-Lifshitz graviton changes, one may speculate that the homogeneity of the universe may be described without the need for inflation. In any case, such a theory is both theoretically and cosmologically rich and it deserves further study." }, "1003/1003.4734_arXiv.txt": { "abstract": "We present a study of a 20cm selected sample in the Deep SWIRE VLA Field, reaching a 5--$\\sigma$ limiting flux density at the image center of $S_{1.4GHz}\\sim13.5 \\mu$Jy. In a $0.6\\times0.6$ square degrees field, we are able to assign an optical/IR counterpart to 97\\% of the radio sources. Up to 11 passbands from the NUV to 4.5$\\mu$m are then used to sample the spectral energy distribution (SED) of these counterparts in order to investigate the nature of the host galaxies. By means of an SED template library and stellar population synthesis models we estimate photometric redshifts, stellar masses, and stellar population properties, dividing the sample in three sub--classes of quiescent, intermediate and star--forming galaxies. We focus on the radio sample in the redshift range $0.31$. Our conclusion from this analysis is that $\\MDOT_{\\rm w}$ is greater than the rate at which the central engine is fueled. This in turn can make the disk variable. One could even expect this relatively strong wind to cause a recurrent disappearance of the inner disk even if the disk is fed at a constant rate at large radii. Neilsen \\& Lee (2009) suggested such a process while interpreting jet/wind/radiation variability in the microquasar GRS 1915+105. It is beyond the scope of this work to model the effect of the wind on the disk. However, this problem was studied by Shields et al. (1986), who showed that due to viscous processes no oscillations appear for $\\MDOT_{\\rm w}\\, / \\,\\MDOT_{\\rm a}\\, <\\,15$! We conclude that the thermal wind is to weak to cause the oscillation. But the thermal wind is also too weak to account for the observed wind. Therefore, it is possible that in GRO J1655--40, and other sources, e.g., GRS 1915+105, the wind responsible for the observed X absorption will be so strong that $\\MDOT_{\\rm w}\\, / \\,\\MDOT_{\\rm a}\\, >\\,15$ and as such cause disk oscillations and contribute to the observed disk variability. We thank Tim Waters for his comments on the manuscript. We acknowledge support provided by the Chandra awards TM8-9004X and TM0-11010X issued by the Chandra X-Ray Observatory Center, which is operated by the Smithsonian Astrophysical Observatory for and on behalf of NASA under contract NAS 8-39073. \\clearpage" }, "1003/1003.5959_arXiv.txt": { "abstract": "We introduce a new figure of merit for comparison of proposed dark energy experiments. The new figure of merit is objective and has several distinct advantages over the Dark Energy Task Force Figure of Merit, which we discuss in the text. ", "introduction": "\\label{sec:introduction} The measurements of the Cosmic Microwave Background (CMB) anisotropies, most notably by the Wilkinson Microwave Anisotropy Probe (WMAP) mission \\cite{Jarosik:2010iu}, but also by ground based and balloon borne experiments, such as VSA \\cite{2003MNRAS.341L..29S}, CbI \\cite{CBI:04}, AcbaR \\cite{ACBAR:07}, SPt \\cite{Reichardt:2009ys}, QuaD \\cite{Brown:2009uy} are in spectacular agreement \\cite{Melch} by predictions of the standard cosmological model. Measurements of the low-redshift universe including data from large spectroscopic surveys like SDSS \\cite{abazajian:09,Schleg} analysed using a variety of methods \\cite{Percival:2009xn,2006PhRvD..74l3507T}, measurements of the luminosity distance to type Ia supernovae \\cite{riess:09,2008ApJ...686..749K} and the galaxy-galaxy lensing \\cite{2008MNRAS.390.1157R,2006ApJ...640..691V} strengthen the standard picture. These datasets provide constraints on the cosmological model using a variety of techniques with very different systematic issues, which nevertheless converge, within the error-bars, on the standard model of cosmology. Perhaps the most surprising aspect of the standard cosmological model is the overwhelming evidence that the Universe is undergoing accelerated expansion. This expansion can be most easily explained in terms of cosmological constant $\\Lambda$. In fact, from the theoretical perspective, the cosmological constant is perhaps a very natural phenomenon, whose small size only illustrates our poor understanding of the fundamental theory of the Universe \\cite{Bianchi:2010uw,Bousso:2009ks}. However, motivated by the need to establish new and much needed gaps in literature, various authors have considered alternatives to cosmological constant, which most often include new degrees of freedom. Such models generically predict effective equation of state $w=p/\\rho$ for these novel components of the Universe that can deviate from the value for the cosmological constant, namely $w=-1$ by arbitrarily small amounts \\cite{Caldwell:1997ii}. Concurrently with these efforts, an industry of phenomenological models has been established. One of the most popular descriptions for the dynamical dark energy is the $w_0$-$w_a$ parametrisation, in which the equation of state is postulated to evolve with cosmological scale factor $a=(1+z)^{-1}$ (where $z$ is redshift) as \\cite{2003PhRvL..90i1301L} \\begin{equation} w(a) = w_0 + (1-a) w_a. \\end{equation} At the same time, motivated by the need to attract funding, experimentalists have begun proposing various experiments that will measure the value of $w$ and its derivatives with an ever increasing precision. In fact, designing cosmological experiments around measuring neutrino masses from cosmology has a distinct disadvantage in that sum of neutrino mass eigenstates has a lower limit given by the ground-based experiments \\cite{kamland:2008ee,2004PhRvL..92r1301A,2004PhRvL..93j1801A}. The same holds true for constraining theories of inflation by measuring the running of the spectral index, which is expected to be of $O(10^{-4})$ in the simplest inflationary models, given the current limits on the tilt of the spectral index $n_s\\sim0.96$ \\cite{Komatsu:2010fb}. Measuring $w$ poses no such difficulties: since there is a strong theoretical prejudice that $w=-1$, one can hope to improve limits on deviation from the cosmological constant value for decades to come. An interesting question, worth every penny of scientific funding, is the question of comparison of various dark energy experiments. Several methods have been established for this purpose, the most common is the Dark Energy Task Force (DETF) Figure of Merit (FoM) \\cite{2006astro.ph..9591A}. In this paper we propose a new metric, that has several advantages. We discuss the DETF FoM and our new metric in Section \\ref{sec:new-figure-merit}. We conclude in Section \\ref{sec:Conclusions}. Finally, we also note that clusters of galaxies are the most massive gravitationally bound objects in the Universe. ", "conclusions": "\\label{sec:Conclusions} In this paper we have introduced a new figure of merit, $\\mathbf{\\ddot{\\phi}}$, which is proportional to the inverse of the circumference of the error ellipse. As discussed in the text and we discuss it here again, the new figure of merit has several advantages over the old one. You and your dog should use it. If you do not use it and think it is a pointless number, you should nevertheless cite this paper or I will write you hassling emails. If worse come to worse, I'll resort to the crowbar and smash your 30 inch liberal screen. This work opens clear avenues for further research. The quantity $\\mathbf{\\ddot{\\phi}}$ can and should be calculated for many future experiments and further compared to the DETF FoM. Rigorous extension of this work into models of dark energy with more than two parameters remains to be performed. Theoretical basis for similarities and differences between the two figures of merit should be established and elaborated. Different parametrisation of the dark-energy should be integrated into the new figure of merit resulting in a multitude of useful figures of merit. The best in the science of figures of merit has yet to come!" }, "1003/1003.5218_arXiv.txt": { "abstract": "We explore the nature of Infrared Excess sources (IRX), which are proposed as candidates for luminous [$L_X(\\rm 2-10\\,keV) > 10^{43} \\, erg \\, s^{-1}$] Compton Thick ($\\rm N_H>2\\times 10^{24} \\, cm^{-2}$) QSOs at $z\\approx2$. Lower redshift, $z\\approx1$, analogues of the distant IRX population are identified by firstly redshifting to $z=2$ the SEDs of all sources with secure spectroscopic redshifts in the AEGIS (6488) and the GOODS-North (1784) surveys and then selecting those that qualify as IRX sources at that redshift. A total of 19 galaxies are selected. The mean redshift of the sample is $z\\approx1$. We do not find strong evidence for Compton Thick QSOs in the sample. For 9 sources with X-ray counterparts, the X-ray spectra are consistent with Compton Thin AGN. Only 3 of them show tentative evidence for Compton Thick obscuration. The SEDs of the X-ray undetected population are consistent with starburst activity. There is no evidence for a hot dust component at the mid-infrared associated with AGN heated dust. If the X-ray undetected sources host AGN, an upper limit of $L_X(\\rm 2-10\\,keV) = 10^{43} \\, erg \\, s^{-1}$ is estimated for their intrinsic luminosity. We propose that a large fraction of the $z\\approx2$ IRX population are not Compton Thick QSOs but low luminosity [$L_X(\\rm 2-10\\,keV) < 10^{43} \\, erg \\, s^{-1}$], possibly Compton Thin, AGN or dusty starbursts. It is shown that the decomposition of the AGN and starburst contribution to the mid-IR is essential for interpreting the nature of this population, as star-formation may dominate this wavelength regime. ", "introduction": "The composition of the diffuse X-ray background remains a problem for high energy astrophysics. X-ray imaging surveys with the Chandra and XMM-Newton observatories have resolved between 80 to 100\\% of the XRB into discrete sources below about 6\\,keV \\citep[e.g.][]{Worsley2005, Hickox2006, Georgakakis2008_sense}. The vast majority of these sources are Compton thin AGN ($\\rm N_H \\la 10^{24} \\rm \\,cm^{-2}$) at a mean redshift $z\\approx1$ \\citep[e.g.][]{Barger2005, Akylas2006}. At higher energies however, between 20 and 30\\,keV, where the bulk of the XRB energy is emitted \\citep{Marshall1980}, only a small fraction of its total intensity has been resolved into discreet sources \\citep{Sazonov2007}. As a result the nature of the populations that make up the XRB at these energies is still not well known. Population synthesis models use our knowledge on the properties of the X-ray sources below $\\approx10$\\,keV to make predictions on the nature of the X-ray populations close to the peak of the XRB \\citep[e.g.][]{Gilli2007}. These models indicate that Compton thin AGN alone cannot account for the shape of the XRB spectrum at $\\approx20-30$\\,keV. An additional population of heavily obscured, Compton thick ($\\rm N_H \\ga 10^{24} \\rm \\,cm^{-2}$) AGN is postulated to reconcile the discrepancy \\citep[e.g.][]{Gilli2007}. The required number density of such sources is however, under debate \\citep[][]{Treister2009xrb, Draper_Ballantyne2009}. Unfortunately the identification of the heavily obscured AGN population predicted by the models is far from trivial. The X-ray emission of these sources below about 10\\,keV is suppressed by photoelectric absorption and as a result most of them are expected to lie well below the sensitivity limits of the deepest current X-ray observations. Although a handful of Compton thick AGN candidates have been identified in deep X-ray surveys \\citep{Tozzi2006, Georgantopoulos2009}, the bulk of this population remains to be discovered. Selection at the mid-IR ($\\rm 3-30\\,\\mu m$) is proposed as a powerful tool for finding heavily obscured X-ray faint AGN. The UV/optical photons emitted by the central engine are absorbed by the gas and dust clouds and appear as thermal radiation with a broad bump in the mid-IR \\citep[$\\nu f_{\\nu}$ units; e.g. ][]{Elvis1994, Prieto2009}. Diverse selection methods have been developed to identify this AGN spectral signature in the mid-infrared. \\cite{Lacy2004}, \\cite{Stern2005} and \\cite{Hatziminaoglou2005} propose simple colour cuts based on the mid-IR colours of luminous high redshift QSOs and/or type-2 Seyferts. \\cite{Polletta2006} and \\cite{Rowan-Robinson2009} fit templates to the broad-band photometry from UV to the far-IR. \\cite{Alonso-Herrero2006} and \\cite{Donley2007} select sources with power-law Spectral Energy Distributions (SEDs) in the mid-IR. The methods above have merits and shortcomings. Selection by mid-IR colour for example, is simple but suffers contamination from star-forming galaxies if applied to the deepest mid-IR samples currently available \\citep{Georgantopoulos2008, Donley2008}. Template fits are powerful but require high quality photometry over a wide wavelength baseline for meaningful constraints. Power-law SEDs provide the most clean samples of infrared selected AGN, but are sensitive only to the most luminous and hence, rare sources. A much promising method that is believed to be efficient in identifying heavily obscured, possibly Compton Thick, AGN is based on the selection of sources that are faint at optical and bright at mid-IR wavelengths. In its simplest version this method applies a cut in the $\\rm 24\\,\\mu m$ over $R$-band flux density ratio, $f_{\\rm 24\\mu m} / f_{\\rm R} > 1000$ to identify Dust Obscured Galaxies \\citep[DOGs;][]{Dey2008} at a mean redshift $z\\approx2$. In addition to the limit above \\cite{Fiore2008, Fiore2009} also use the colour cut $R-K>4.5$ to select Infrared-Excess sources (IRXs). This colour selection is motivated by the observational result that redder sources include a higher fraction of obscured AGN \\citep[e.g.][]{Brusa2005}. A different approach has been adopted by \\cite{Daddi2007}. They select $BzK$ sources that show excess mid-IR emission relative to that expected based on the rates of star formation measured from shorter wavelengths. Despite differences in the adopted criteria, all the studies above identify a population of galaxies with similar properties in terms of average redshift ($z\\approx2$), mean X-ray properties and mid-IR luminosities. These sources are believed to be massive galaxies \\citep[$M_{star} \\approx 10^{10} - 10^{11} \\, M_{\\odot}$][]{Treister2009, Bussmann2009_hst} that experience intense bursts of star-formation and rapid supermassive black hole growth, possibly triggered by mergers \\citep[e.g.][]{Bussmann2009_hst, Narayanan2009}. The apparently brighter subset of this population ($\\rm S_{24}>300\\, \\mu Jy$) are proposed as descendants of Submillimeter Galaxies on their way to becoming unobscured QSOs \\citep[e.g.][]{Bussmann2009_hst, Bussmann2009_sed, Narayanan2009}, which will eventually evolve into present-day $4\\,L^{*}$ galaxies \\citep{Dey2008, Brodwin2008}. Although these sources are undoubtedly important for understanding the co-evolution of galaxies and SBHs at high redshift, in this paper we focus on their significance to XRB studies by putting into test claims that they include a large fraction of Compton Thick AGN. The two key properties of these sources (hereafter referred to as IRX sources), which are interpreted as evidence for Compton Thick AGN are (i) their hard mean X-ray spectrum and (ii) their faintness at X-ray wavelengths relative to the mid-IR. The first point is demonstrated in Figure \\ref{fig_hr} which shows the that the mean hardness ratio of the IRX sources is consistent with that of the local Compton Thick AGN NGC\\,1068 \\citep{Matt1999}. With respect to the second point, the majority of the IRX sources \\citep[80 per cent;][]{Georgantopoulos2008} are not detected in the deepest X-ray surveys available and their mean X-ray properties can only be studied through stacking analysis. Figure \\ref{fig_hist}, shows the average X-ray to mid-IR luminosity ratio, $L_X({\\rm 2-10\\,keV}) / \\nu L_{\\nu} {\\rm 5.8\\mu m}$, of the X-ray undetected IRX population. Relative to local AGN from the sample of \\cite{Lutz2004} these sources appear underluminous in the mid-IR by 2-3\\,dex. Assuming a narrow range of $L_X({\\rm 2-10\\,keV}) / \\nu L_{\\nu} {\\rm 5.8\\mu m}$ for typical AGN (e.g. Lutz et al. 2004), the position of the IRX population in Figure \\ref{fig_hist} is consistent with a Compton Thick obscuring screen that suppresses the observed X-ray emission relative to the mid-IR. Under the assumption that the mid-IR luminosity is a good proxy of the AGN power, one can estimate intrinsic 2-10\\,keV luminosities for the IRX sources in excess of $10^{43} \\rm erg \\, s^{-1}$ \\citep{Daddi2007,Fiore2008,Fiore2009, Treister2009}. AGN with intrinsic luminosities above this limit are hereafter referred to as QSOs. IRX sources are therefore, excellent candidates for the Compton Thick QSOs needed by population synthesis models to explain the XRB spectrum at 20-30\\,keV. Simulations indeed show that the bulk of these sources (80 per cent, Fiore et al. 2008; 95 per cent, Fiore et al. 2009) are likely to be Compton Thick QSOs, while their estimated space densities are consistent with the predictions of the XRB models. Although there is no doubt that the IRX population includes a fraction of heavily obscured AGN \\citep{Georgantopoulos2009}, the evidence for Compton Thick sources is far from conclusive. The two key properties of the IRX sources, hardness ratio and X-ray to mid-IR luminosity ratio, are also consistent with Compton thin AGN of lower luminosity, hereafter defined as $L_X(\\rm 2-10\\,keV) < 10^{43} \\rm erg \\, s^{-1}$. Figure \\ref{fig_hr} for example, shows that the mean hardness ratio of this population could be due to moderate obscuring column densities of few times $\\rm 10^{23}\\,cm^{-2}$. Moreover, Figure \\ref{fig_hist} demonstrates that lower luminosity Compton Thin AGN (Terashima et al. 2002) have very low X-ray to mid-IR luminosity ratios, i.e. similar to those observed for IRX sources. This is not surprising as both AGN and star-formation contribute to the mid-IR, thereby resulting to a broad $L_X({\\rm 2-10\\,keV}) / \\nu L_{\\nu} {\\rm 5.8\\mu m}$ distribution. One has to isolate the AGN component at the mid-IR part of the SED to get a tight correlation with X-ray luminosity \\citep[e.g. ][]{Prieto2009}. \\cite{Lutz2004} for example, accomplished that by selecting {\\it only} sources where the AGN dominates in the mid-IR, i.e. those {\\it without} a significant starburst component relative to the AGN. For the high redshift IRX sources it is not clear what fraction of the mid-IR emission is from AGN heated dust. Only under the strong assumption that the {\\it bulk} the mid-IR luminosity is associated with reprocessed radiation from accretion on the central SBH, can one infer that these sources are Compton Thick QSOs. A number of studies on the IRX sources have already suggested that this population may include a large fraction of lower luminosity Compton thin AGN or even dusty starbursts \\citep[e.g.][]{Georgantopoulos2008, Donley2008, Pope2008, Murphy2009}. Elucidating the nature of IRX sources, Compton Thick QSOs, Compton Thin lower luminosity AGN or starbursts, requires very deep X-ray data and/or AGN/starburst decomposition at the mid-IR. Unfortunately, IRX samples at $z\\approx2$ are too faint at almost any wavelength, expect the mid-IR, for such a detailed study. In this paper we address this issue by selecting lower redshift ($z\\approx1$) analogues of the distant ($z\\approx2$) IRX population. The advantage of this approach is that the selected sources are apparently bright and therefore their SEDs can be constrained over a wide wavelength baseline, from the far-UV to the far-IR. The sources are selected in fields with some of the deepest X-ray observations available (Chandra Deep Field North, All wavelength Extended Groth strip International Survey: AEGIS), thereby allowing study of their X-ray properties via spectral analysis. An additional advantage of our strategy is that there is only about 2.5\\,Gyr difference in the age of the Universe between $z=1$ and $z=2$ ($\\rm H_0 = 70 \\, km \\, s^{-1} \\, Mpc^{-1}$, $\\rm \\Omega_{M} = 0.3$, $\\rm \\Omega_{\\Lambda} = 0.7$). Therefore, by selecting galaxies at $z\\approx1$, it is more likely to identify systems that are physically similar to the IRX population at $\\approx2$. In this respect it is interesting that in the local Universe there are no known analogues of the IRX sources, which can be used to study in detail the nature of this population \\citep[e.g. ][]{Dey2008}. Throughout this paper we adopt $\\rm H_0 = 70 \\, km \\, s^{-1} \\, Mpc^{-1}$, $\\rm \\Omega_{M} = 0.3$ and $\\rm \\Omega_{\\Lambda} = 0.7$. \\begin{figure} \\begin{center} \\includegraphics[height=0.9\\columnwidth]{hr.pdf} \\end{center} \\caption{Hardness ratio against redshift. The hardness ratio is defined as HR=(H-S)/(H+S), where H and S are the counts in the 1.5-6 and 0.3-1.5\\,keV energy intervals respectively. The (blue) dot is the mean hardness ratio of X-ray undetected sources in the COSMOS survey (Fiore et al. 2009) estimated using stacking analysis. The vertical errorbar is the $1\\sigma$ uncertainty of the HR, while the horizontal errorbar shows the redshift range of the Fiore et al. (2009) IRX population. The dotted line (black) shows the expected HR of the local Compton Thick AGN NGC\\,1068 Matt et al. (1999). The (red) dashed curve corresponds to the HR of a power-law X-ray spectrum with photon index $\\Gamma=1.9$, which is absorbed by a column density of $\\rm N_H= 1.5 \\times 10^{23} \\, cm^{-2}$. The continuous (black) line corresponds to the HR of an unobscured AGN, i.e a power-law X-ray spectrum with photon index $\\Gamma=1.9$. }\\label{fig_hr} \\end{figure} \\begin{figure} \\begin{center} \\rotatebox{0}{\\includegraphics[height=0.9\\columnwidth]{hist_lxlir.pdf}} \\end{center} \\caption{Distribution of X-ray--to--mid-IR luminosity ratio, $L_X({\\rm 2-10\\,keV})/ \\nu L_{\\nu}(\\rm 5.8\\mu m)$. The (black) unshaded histogram is for AGN from Lutz et al. (2004). This sample was selected to include only AGN that dominate in the mid-IR. The vertical dashed lines show the dispersion of the X-ray--to--mid-IR luminosity ratio for the same sample. The low-luminosity AGN of Terashima et al. (2002) are shown with the blue hatched histogram. The mean $L_X({\\rm 2-10\\,keV})/ \\nu L_{\\nu}(5.8\\mu m)$ of IRX sources is plotted with the (red) arrow. This estimate is for X-ray undetected IRX sources in the CDF-North (Georgantopoulos et al. 2008). The mean X-ray luminosity of this population is determined via stacking analysis and $\\nu L_{\\nu}(\\rm 5.8\\mu m)$ is the average mid-IR luminosity of the sample. Both luminosities are estimated assuming $z=2$ for the IRX sources. The $\\rm 5.8\\mu m$ luminosity of the Terashima et al. (2002) sources is determined using the IRAS $\\rm 12\\,\\mu m$ flux density and adopting an average flux density ratio $f_{\\rm 12\\,\\mu m}/f_{\\rm 5.8\\,\\mu m}=0.9$, estimated from the subsample of these sources with $6\\mu m$ flux density measurements from the literature. The 2-10\\,keV X-ray luminosities of the Terashima et al. (2002) AGN are corrected for obscuration, except for a total of 6 sources in that sample, which show a strong Fe\\,Ka 6.4\\,keV line in their X-ray spectra with equivalent width $\\rm EW > 900 \\, eV$ and are therefore Compton Thick candidates. }\\label{fig_hist} \\end{figure} ", "conclusions": "\\label{sec_results} The population of IRX galaxies has attracted much attention recently as they are proposed as good candidates for Compton thick QSOs at $z\\approx2$ with luminosities $L_X(\\rm 2-10\\, keV) > 10^{43} \\, erg \\, s^{-1}$. The low X-ray--to--mid-IR luminosity ratio and the mean hardness ratios of these sources are consistent with this interpretation as long as the bulk of the mid-IR luminosity is associated with the dust heated by the central engine. There are indications however, that this is not the case, at least not for all IRX sources. \\cite{Murphy2009} combined mid-IR spectroscopy with far-IR sub-mm data to show that only 50 per cent of the infrared excess sources selected in a way similar to that described by Daddi et al. (2007), show evidence for obscured AGN activity. The other half of this population have SEDs consistent with star-formation. \\cite{Yan2007} presented Spitzer-IRS spectra of 52 galaxies brighter than 1\\,mJy at $\\rm 24\\mu m$ selected in the Spitzer extragalactic First Look Survey (xFLS). The relative strengths of the AGN and starburst components in those sources have been estimated by \\cite{Sajina2008} using in addition to the mid-IR spectroscopy, multiwavelength photometric observations (UV to far-IR and radio), optical and NIR spectroscopy. There are 21 sources in the \\cite{Yan2007} sample with properties similar to the IRX population at $z\\approx2$, i.e. $f_{\\rm 24\\mu m}/f_R>900$ and spectroscopic redshift determination $z>1.5$. \\cite{Sajina2008} found that about 76 per cent (16/25) of those sources are AGN dominated in the IR, while the remaining 24 per cent have starburst components that either dominate or contribute nearly equally to the AGN at IR wavelengths. If star-formation contributes to the mid-IR luminosity of IRX galaxies then the two key properties of this population, X-ray--to--mid-IR luminosity ratio and mean hardness ratio, are also consistent with either lower-luminosity Compton thin AGN (e.g. Figures 1, 2) or pure starbursts \\citep{Donley2008}. In order to shed more light in the nature of IRX sources we select galaxies at $z\\approx1$ in the AEGIS and CDF-North fields with SEDs similar to the $z\\approx2$ IRX population. The advantage of this approach is that the selected sources have fluxes that are brighter at almost any wavelength compared to IRX galaxies at $z\\approx2$, thereby greatly facilitating their study. It is interesting that about 35 per cent of the sources in the AEGIS (5/14) and 80 per cent in the CDF-North (4/5) are associated with hard (2-7\\,keV) X-ray detections with significance $<4 \\times 10^{-6}$. These fractions increase to 43 (6/14) and 100 (5/5) per cent in the AEGIS and CDF-N respectively, if lower significance sources are included. For comparison only about 4 per cent of red cloud galaxies in the AEGIS have X-ray counterparts (Nandra et al. 2007). This suggests a high AGN identification rate in samples selected using the infrared excess criteria. Excluding X-ray detected starburst candidates (irx-16 and irx-17), the IRX sources with X-ray counterparts are indeed AGN, as indicated by their X-ray luminosities, $L_X(\\rm 2-10\\,keV) > 10^{42} \\, erg \\, s^{-1} \\, cm^{-2}$, which are higher than what can be attributed to star-formation \\citep{Georgakakis2007}. These sources also have hard X-ray spectral properties, which are however consistent with Compton Thin column densities, $N_H \\approx \\rm 10^{22} - 5\\times 10^{23} \\, cm^{-2}$. There is only tentative evidence for Compton Thick obscuration among the X-ray detected IRX sources. For four of them the 90 per cent upper limit of the reflection fraction is higher than $R=10$. Also, the X-ray spectra of two of those four sources either have 90 per cent upper limit in the column density $\\approx1.2\\times10^{24} \\rm \\, cm^{-2}$ or show the FeK$\\alpha$\\,6.4keV line with an equivalent width of $0.5$\\,keV. These properties can be interpreted as evidence for Compton Thick AGN but are also consistent with Compton thin obscuration. More relevant to the deeply buried AGN picture are the infrared excess sources in the sample that are not detected at X-ray wavelengths. The SED modelling shows that their mid- and far-IR is dominated by star-formation and that there is no indication for a hot dust component associated with a powerful AGN that remains undetected at X-ray wavelengths. It is noted the QSO torus template is required to fit the IR SEDs of {\\it all} X-ray detected AGN with intrinsic X-ray luminosities brighter than $L_X(\\rm 2-10\\, keV)\\approx 10^{43} \\, erg \\, s^{-1}$. Therefore, if some of the X-ray undetected sources were associated with heavily obscured and powerful AGN with X-ray emission suppressed by the intervening dust and gas clouds, we would have identified them in the mid-IR as sources with a QSO torus component. We can therefore, place an upper limit of $L_X(\\rm 2-10\\, keV) = 10^{43} \\, erg \\, s^{-1}$ to the intrinsic AGN luminosity of X-ray undetected sources, if they host an active SBH. We conclude that the IRX sources without X-ray counterparts in our sample are either starbursts, lower luminosity AGN, or a combination of the two. This result has implications on the nature of the $z\\approx2$ IRX population detected in deep surveys. There is no doubt that some of these sources are Compton Thick QSOs. \\cite{Georgantopoulos2009} for example, identified Compton Thick QSOs among X-ray sources in the CDF-North through X-ray spectroscopy and showed that some of them satisfy the infrared excess selection criteria. Our analysis however, shows that a potentially large fraction of the IRX population at $z\\approx2$ are not luminous [$L_X(\\rm 2-10\\, keV) > 10^{43} \\, erg \\, s^{-1}$] Compton Thick QSOs but lower luminosity [$L_X(\\rm 2-10\\, keV) < 10^{43} \\, erg \\, s^{-1}$] possibly Compton thin AGN and/or starbursts. Dust enshrouded star-formation has already been proposed as an alternative to Compton Thick QSOs to explain the properties of the $z\\approx2$ IRX sources. \\cite{Donley2008} critically reviewed different methods proposed in the literature for finding AGN at infrared wavelengths, including the infrared excess selection. They showed that the red SED of these sources are consistent with those of local pure starbursts with some moderate amounts of additional reddening at optical wavelengths ($A_V\\sim1$\\,mag). They also cautioned that the stacked X-ray signal of this population may be dominated by few sources. Moreover, they argued that flat mean X-ray spectral properties of the IRX population could be the result of high mass X-ray binaries, which are known to have hard X-ray spectral properties (power-laws with $\\Gamma \\approx 1.2$). Whether binary stars can dominate the integrated X-ray emission of local starbursts however, is still under debate. In any case, based on the arguments above \\cite{Donley2008} suggested that as much as half of the IRX population at $z\\approx2$ are starbursts with only about 20 per cent showing evidence for heavily obscured, possibly Compton Thick, AGN activity. \\cite{Pope2008} analysed the mid-IR spectra of 12 IRX sources with $f_{\\rm 24\\mu m}>300\\, \\rm \\mu Jy$ in the GOODS-North and found that 6 of them are dominated by star-formation at the mid-IR. They also showed that AGN and starburst dominated IRX sources have distinct $\\rm 8.0\\mu m$ over $\\rm 4.5\\mu m$ flux ratios, $f_{\\rm 8.0\\mu m}/f_{\\rm 4.5\\mu m}$. Extrapolating their results to fainter IRX sources $f_{\\rm 24\\mu m}=100-300\\, \\rm \\mu Jy$, for which mid-IR spectroscopy is not available, they found that about 80 per cent of them have $f_{\\rm 8.0\\mu m}/f_{\\rm 4.5\\mu m}$ consistent with star-formation. They also estimate the mean SED of IRX sources in their sample and conclude that less than about 10 per cent of the total infrared luminosity ($\\rm 8-1000\\mu m$) is associated with hot dust, possibly heated by an AGN. For the average infrared luminosity of the \\cite{Pope2008} sample, $L_{IR}=10^{12} \\, L_{\\odot}$, the fraction above translates to an upper limit in the AGN luminosity of $L_{IR}\\approx4\\times 10^{44} \\rm \\, erg \\, s^{-1}$. Adopting the AGN bolometric correction factors $L_{bol}/L_{IR}= 3$ \\citep{Risaliti_Elvis2004} and $L_{bol}/L_{X}(\\rm 2 - 10 \\, keV)= 35$ \\citep{Elvis1994}, we estimate a mean hard X-ray luminosity $L_{X} (\\rm 2-10 \\,keV) < 3\\times 10^{43} \\, erg \\, s^{-1}$. Although there are uncertainties in this calculation, there is broad agreement in the upper limits in $L_X$ estimated for the \\cite{Pope2008} and our sample of IRX sources. The sample presented here has properties similar to moderately bright IRX sources at $z\\approx2$, i.e. $\\rm 5.8\\, \\mu m$ luminosity few times $\\rm 10^{44} \\, erg \\,s^{-1}$, like those found in the CDF-South by Fiore et al. (2008). The simulations of \\cite{Narayanan2009} suggest that IRX sources above $S_{24} \\rm = 300 \\, \\mu Jy$ are dominated by gaseous mergers and include a large fraction of powerful AGN, whereas less luminous systems, like those studied here, are typically secularly evolving galaxies dominated by star-formation and with only weak AGN activity. Our results on the nature of moderately luminous IRX sources are therefore consistent with those simulation. It is likely that IRX sources more luminous than those studied here (i.e. COSMOS field, $S_{24} \\rm > 500 \\, \\mu Jy$, $\\nu L_{\\nu}(\\rm 5.8\\mu m)> 10^{45} \\, erg \\,s^{-1}$; Fiore et al. 2009) include a higher fraction of Compton Thick QSOs. The results of Sajina et al. (2008) on the nature of bright ($S_{24} \\rm > 900 \\, \\mu Jy$) IRX sources selected in the xFLS (Yan et al. 2007) suggests that this may be the case. \\cite{Bauer2010} explored the X-ray properties of a subset of the Yan et al. (2007) sample and found evidence for at least mildly Compton Thick ($N_H\\approx10^{24}\\, \\rm cm^{-2}$) obscuration in a large fraction of the sources that are AGN dominated in the mid-IR. Even at the extreme luminosities of the xFLS sample however, a non-negligible fraction of the IRX sources (about 24 per cent) have a substantial or even dominant starburst component in the mid-IR. This underlines the importance for subtracting the contribution of star-formation to the mid-IR to assess the intrinsic AGN luminosity before concluding whether the X-ray/mid-IR properties of these systems are consistent with Compton Thick obscuration. \\cite{Mullaney2009} have recently found that the X-ray to total infrared luminosity ratios of X-ray AGN with $L_{X} (\\rm 2-10 \\,keV) = 10^{42} - 10^{43} \\, erg \\, s^{-1}$ increases by about 1\\,dex from $z=0$ to $z\\approx1-2$. Although the origin of this trend is unclear, it can be interpreted as an enhancement of the average star-formation rate in lower-luminosity AGN at $z\\approx2$ compared to the local Universe. In this picture, the enhanced star-formation is likely to have an impact on the mid-IR part of the SED and will make lower luminosity AGN at $z\\approx1-2$ appear underluminous at X-rays for their mid-IR luminosity compared to local AGN samples. This would be in agreement with our interpretation of the IRX population and emphasises the need to properly model the mid-IR part of the SED to assess the relative contribution of AGN and star-formation at this wavelength regime. Further progress in the study of IRX galaxies is expected from upcoming Herschel observations, which when combined with the existing Spitzer data will constrain the SEDs of individual sources at $z\\approx2$ over a sufficiently large wavelength baseline (mid to far-IR) to allow decomposition of the AGN and starbursts components through template fits, as done in this paper. This exercise will provide estimates of the level of dusty star-formation activity and the intrinsic AGN luminosity of individual sources to confirm or refute claims that they are Compton Thick QSOs. The determination of the column density of these sources through X-ray spectroscopy has to wait future X-ray missions with large collecting areas, such IXO, which is expected to provide spectra for $z\\approx2$ IRX sources which have typical X-ray fluxes $f_X ( \\rm 2 - 10 \\, keV ) \\approx 10^{-17} \\, erg \\, s^{-1} \\, cm^{-2}$ \\citep{Georgantopoulos2009} and are currently accessible only through stacking analysis." }, "1003/1003.2438_arXiv.txt": { "abstract": "% We report on our Chandra Cycle 9 program to observe half of the 60 (unobserved by Chandra) 3C radio sources at z$<$0.3 for 8 ksec each. Here we give the basic data: the X-ray intensity of the nuclei and any features associated with radio structures such as hot spots and knots in jets. We have measured fluxes in soft, medium and hard bands and are thus able to isolate sources with significant intrinsic column density. For the stronger nuclei, we have applied the standard spectral analysis which provides the best fit values of X-ray spectral index and column density. We find evidence for intrinsic absorption exceeding a column density of 10$^{22}$ cm$^{-2}$ for one third of our sources. ", "introduction": "Extended radio galaxies are classified into two main types (Fanaroff and Riley 1974; Miley 1980; Bridle 1984). The more powerful sources (FRII) tend to have an edge-brightened radio structure dominated by compact bright hot spots. These sources often show either no or one jet which may be relativistic along its entire length (Laing 1988; Garrington et al. 1988). The lower luminosity sources (FRI) tend to have edge-darkened structures which resemble ``plumes'' and usually exhibit two jets. The jets in FRIs may initially be launched relativistically but seem to decelerate on subkpc scales likely through interaction with the environment (e.g., Laing et al. 2008). There are intrinsic differences in the central AGN of these two types of sources (e.g., Baum, Zirbel, O'Dea 1995; Evans et al. 2006, Hardcastle et al. 2009a). Most Narrow Line FRIIs and a few FRIs show evidence for a hidden quasar continuum source and Broad Line Region (BLR) (e.g., Cohen et al. 1999; Tadhunter et al. 2007). The sources with the hidden quasar also produce optical emission line nebulae with high ionization lines (High Excitation Galaxies, HEGs) while those without the hidden quasar produce only low ionization lines (Low Excitation Galaxies, LEGs) (e.g., Hine and Longair 1979, Laing et al. 1994, Rector and Stocke 2001). There are also radio sources with properties intermediate between the FRIs and FRIIs, e.g., the ``Fat Doubles'' (Owen and Laing 1989). The two main unsolved issues concern the origin of the FR~I~/~FR~II dichotomy (how it is related to different acceleration and emission processes), and the nature of the different emission line regions between LEGs and HEGs (see Chiaberge et al. 2002 and Hardcastle et al. 2007). The morphological features of extragalactic radio sources can be described naturally with a small number of components: core, jets, hotspots and lobes. While their radio to optical emission is typically described in terms of synchrotron radiation by relativistic particles, the origin of X-ray emission in extended structures (jets and hotspots) is still unclear, but certainly non-thermal (Harris \\& Krawczynski 2002). The main open question lies in which mechanism, synchrotron or inverse Compton (IC) scattering, dominates the X-ray emission. The former describes emission from low power jets (Harris \\& Krawczynski 2006), while the latter provides a good explanation for high power radio galaxy and quasar jets, in which the seed photons for the IC scattering could be the Cosmic Microwave Background (CMB) (Tavecchio et al. 2000). Only by combining X-ray observations with historical and/or simultaneuos data in other wavebands, is possible to build up the Spectral Energy Distribution (SED) of cores, jets and hotspots and compare them with synchrotron or inverse Compton models to investigate the origin of their emission. During the last few years several snapshot surveys of 3C radio galaxies have been carried out using the Hubble Space Telescope in red, blue, ultraviolet and near-IR continuum and optical spectroscopy which approaches the statistical completeness of the radio catalog: $\\sim 90$\\%. A ground based spectroscopic program for the whole sample with the Galileo Telescope has been completed (Buttiglione et al 2009). We also obtained deep ground based IR $K$-band imaging. Radio images with arcsec resolution are available for most 3C sources from colleagues, the NRAO VLA Archive Survey (NVAS), and the archives of the VLA and MERLIN. VLBA data for some 3C objects with $z<0.2$ have already been obtained (see e.g. Giovannini et al. 2001, Liuzzo et al. 2009 and references therein). Chandra is the only X-ray facility that can offer angular resolution comparable to the optical and radio. Previous X-ray studies are mostly biased towards observations of a special group of X-ray bright sources or objects with well-known interesting features or peculiarities instead of carefully selected samples, unbiased with respect to orientation and spectroscopic classification. The general goals of our program are to discover new jets and hot spots, determine their emission processes on a firm statistical basis, study the nuclear emission of the host galaxy, and derive spectral energy distributions (SED) for an unbiased sample of objects. The resulting dataset will be used to test the unification model and study the nature of nuclear absorption. The current paper consists of all the basic data for this Chandra sample. After a description of the observations and data reduction (\\S\\ref{sec:obs}), we give the general and particular results in \\S\\ref{sec:results}. \\S\\ref{sec:summary} contains a short summary. For our numerical results, we use cgs units unless stated otherwise and we assume a flat cosmology with $H_0=72$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{M}=0.27$ and $\\Omega_{\\Lambda}=0.73$ (Spergel et al., 2007). Spectral indices, $\\alpha$, are defined by flux density, S$_{\\nu}\\propto\\nu^{-\\alpha}$. ", "conclusions": "\\label{sec:summary}% We have presented our analyses of the Chandra 3C Snapshot Survey. Our goal is to obtain X-ray data for all extragalactic 3C sources with z$<$0.3 so as to have a complete, unbiased sample. Our AO9 proposal resulted in the current new data for 30 of the (then) unobserved 60 sources not already in the archive with exposures of 8ks or greater. We will submit a Chandra proposal for AO12 in order to complete the sample. When the remaining 27 are observed, it will then be possible to have a complete sample by recovering many more 3C sources from the archives. We have constructed flux maps for all the observations and given photometric results for the nuclei and radio hotspots. For the stronger nuclei, we have employed the usual spectral analysis, and compared the column densities of intrinsic absorption to those obtained from the ratio of hard to medium fluxes. As expected, we find a sizable fraction (1/3) of our sources showing evidence for significant absorbtion (N$_H~>~5\\times~10^{22}$~cm$^{-2}$). In particular, for 3C~105, we analyzed the archival observations performed by \\xmm~and \\swf~ and we found that its core is absorbed and these spectra are in agreement with our Chandra {\\bf spectral} analysis. We also found a marginal detection of the Fe K$\\alpha$ emission line and the evidence of a soft X-ray excess. We have found several sources worthy of more detailed study, and with collaborators have published papers on 3C~17, 3C~171, and 3C~305. Our decision to waive proprietary rights has permitted others (e.g. Hardcastle et al. 2009) to use relevant data in a timely manner, and we intend to continue this policy in our Chandra AO12 proposal for the remaining unobserved sources." }, "1003/1003.0045_arXiv.txt": { "abstract": "Gamma-ray bursts (GRBs) have been invoked to explain both the $511\\,\\mathrm{keV}$ emission from the galactic bulge and the high-energy positron excess inferred from the ATIC, PAMELA, and Fermi data. While independent explanations can be responsible for these phenomena, we explore the possibility of their common GRB-related origin by modeling the GRB distribution and estimating the rates. For an expected Milky Way long GRB rate, neither of the two signals is generic; the local excess requires a $2\\%$ coincidence while the signal from the galactic center requires a $20\\%$ coincidence with respect to the timing of the latest GRB. The simultaneous explanation requires a $0.4\\%$ coincidence. Considering the large number of statistical ``trials'' created by multiple searches for new physics, the coincidences of a few per cent cannot be dismissed as unlikely. Alternatively, both phenomena can be explained by GRB if the galactic rate is higher than expected. We also show that a similar result is difficult to obtain assuming a simplified short GRB distribution. ", "introduction": " ", "conclusions": "" }, "1003/1003.5660_arXiv.txt": { "abstract": "We report the first detection of asymmetry in a supernova (SN) photosphere based on SN light echo (LE) spectra of Cas~A from the different perspectives of dust concentrations on its LE ellipsoid. New LEs are reported based on difference images, and optical spectra of these LEs are analyzed and compared. After properly accounting for the effects of finite dust-filament extent and inclination, we find one field where the \\ion{He}{1} $\\lambda 5876$ and H$\\alpha$ features are blueshifted by an additional \\about 4000~\\kms\\ relative to other spectra and to the spectra of the Type IIb SN~1993J. That same direction does not show any shift relative to other Cas~A LE spectra in the \\ion{Ca}{2} near-infrared triplet feature. We compare the perspectives of the Cas~A LE dust concentrations with recent three-dimensional modeling of the SN remnant (SNR) and note that the location having the blueshifted \\ion{He}{1} and H$\\alpha$ features is roughly in the direction of an Fe-rich outflow and in the opposite direction of the motion of the compact object at the center of the SNR. We conclude that Cas~A was an intrinsically asymmetric SN. Future LE spectroscopy of this object, and of other historical SNe, will provide additional insight into the connection of explosion mechanism to SN to SNR, as well as give crucial observational evidence regarding how stars explode. ", "introduction": "\\label{sec:intro} Light echoes (LEs) are the scattered light of a transient event that reflects off dust in the interstellar medium. The extra path length of the two-segment trajectory results in LE light arriving at an observer significantly later than the undelayed photons. Such circumstances provide exciting scientific opportunities that are extremely rare in astronomy --- specifically, to observe historical events with modern instrumentation and to examine the same event from different lines of sight (LoS). We have previously employed LEs to take advantage of the time delay, identifying systems of such LEs associated with several-hundred-year old supernova (SN) remnants (SNRs) in the Large Magellanic Cloud and subsequently spectroscopically classifying the supernovae \\citep[SNe;][]{Rest05b, Rest08a}. Our work represented the first time that the spectral classification of SN light was definitively linked with a SNR. The analysis of the X-ray spectrum of that SNR by \\citet{Badenes08} provided confirmation of both its classification (Type Ia) and subclass (high luminosity). Clearly, LE spectroscopy is a powerful technique for understanding the nature of SNe in the Milky Way and Local Group galaxies. Cas~A is the brightest extrasolar radio source in the sky \\citep{Ryle48} and the youngest (age \\about 330~yr) Milky Way core-collapse SNR \\citep{Stephenson02}. Its distance is approximately 3.4~kpc. Dynamical measurements of the SNR indicate that the explosion occurred in year $1681 \\pm 19$ \\citep{Fesen06_expansionasym_age}; we adopt this date for age calculations in this paper. A single historical account of a sighting in 1680 by Flamsteed is attributed to the Cas~A SN \\citep{Ashworth80}, although this has been disputed \\citep{Kamper80}. Cas~A is the youngest of the certain historical CC~SNe and is thus an excellent target for LE studies. \\citet{Krause05} identified a few moving Cas~A features (called ``infrared echoes'') using infrared (IR) images from the {\\it Spitzer Space Telescope}, the result of dust absorbing the SN light, warming up, and reradiating light at longer wavelengths. Their main scientific conclusion, that most if not all of these IR echoes were caused by a series of recent X-ray outbursts from the compact object in the Cas~A SNR, was incorrect because they did not take into account that the apparent motion strongly depends on the inclination of the scattering dust filament \\citep{Dwek08, Rest11_casamotion}. Rather, the echoes must have been generated by an intense and short burst of ultraviolet (UV) radiation associated with the breakout of the SN shock through the surface of the progenitor star of Cas~A \\citep{Dwek08}. The first scattered LEs of Galactic SNe associated with Tycho's SN and the Cas~A SN were discovered by \\citet{Rest07, Rest08b}. Contemporaneously, \\citet{Krause08a} obtained a spectrum of a scattered optical light echo spatially coincident with one of the Cas~A IR echoes, and identified the Cas~A SN to be of Type IIb from its similarity to the spectrum of SN~1993J, the very well-observed and prototypical example of the SN~IIb class \\citep[e.g.,][]{Filippenko93, Richmond94, Filippenko94, Matheson00}. The discovery and spectroscopic follow-up observations of different LEs from the same SN allow us to benefit from the unique advantages of LEs --- their ability to probe the SN from significantly different directions. The only time this technique has been applied previously was by \\citet{Smith01_eta,Smith03_eta}, who used spectra of the reflection nebula of $\\eta$~Carinae to observe its central star from different directions. Dust concentrations scattering SN light lie at numerous different position angles and at different radial distances from the observer. This is illustrated in Figure~\\ref{fig:3Dspecillustration}, which shows the Cas~A SNR (red dots), the scattering dust (brown dots), and the light paths. We denote the light echoes as LE2116, LE2521, and LE3923, where the number represents the ID of the grid tile of the search area in this region of the sky. LE2521 and LE3923 are newly discovered and LE2116 was discovered by \\citet{Rest08b}. The scattering dust of LE3923 is more than 2000~ly in front of Cas~A, much farther than any other scattering dust; thus, we show only part of its light path. \\begin{figure*}[t] \\begin{center} \\epsscale{1.15} \\plotone{f1.pdf} \\caption[]{North is toward the positive-$y$ axis (up), east is toward the negative-$x$ axis (left), and the positive-$z$ axis points toward the observer with the origin at the SNR. The SN to LE cone apex distance, in light years, is half the interval of time since the SN explosion. The red and brown circles indicate the SN and scattering dust, respectively. The black lines show the path of the light scattering from the LE-producing dust concentrations. The scattering dust of LE3923 is more than 2000~ly in front of Cas~A, much farther than any other scattering dust; thus, we show only part of the light path, and do not include its dust location. The top-left panel shows an \\textit{Chandra} X-ray image \\citep{Hwang04}, with the projected light path from SN to scattering dust overplotted (gray arrows). The red arrow indicates the X-ray compact object and its apparent motion. In this false-color image, red corresponds to low-energy X-rays around the Fe~L complex (\\about 1 keV and below), green to mid-energy X-rays around the Si~K blend (\\about 2 keV), and blue to high-energy X-rays in the 4--6 keV continuum band between the Ca~K and Fe~K blends. \\label{fig:3Dspecillustration}} \\end{center} \\end{figure*} The study of a single SN from different LoS is particularly relevant for the Cas~A SNR. With observations of the LEs shown in Figure~\\ref{fig:3Dspecillustration}, we are positioned to {\\it directly} measure the symmetry of a core-collapse SN, and compare it to the structure of the remnant. The thermal X-ray emission as well as the optical emission of the Cas~A SNR is very inhomogeneous, with large Fe-rich and Si-rich outflows being spatially distinct, indicating that the SN explosion was asymmetric \\citep{Hughes00,Hwang04,Fesen06_expansionasym_age}. The nature of these outflows and their individual relevance to the SN explosion is still debated \\citep{Burrows05,Wheeler08,Delaney10}. \\citet{Tananbaum99} detected in \\textit{Chandra} images a compact X-ray source 7\\arcsec\\ from the Cas~A SNR center; it is an excellent candidate for being the neutron star produced by the SN explosion \\citep{Tananbaum99,Fesen06_CO}. The position angle of the X-ray source is off by only $\\sim 30$\\arcdeg\\ from the position angle of the southeast Fe-rich structure \\citep{Wheeler08}, and has a projected apparent motion of 350~\\kms. In \\S \\ref{sec:observations} of this paper, we report two new LE complexes associated with Cas~A, and we show optical spectra of these LEs as well as one of the previously known LEs. With these data we are able to view three distinct directions, where each dust concentration probes different hemispheres of the SN photosphere. In \\citet{Rest11_leprofile} we have introduced an innovative technique for modeling both astrophysical (dust inclination, scattering, and reddening) and observational (seeing and slit width) effects to measure the light-curve weighted window function which is the determining factor for the relative time-weighting of the observed (integrated) LE spectrum. We apply this technique in \\S \\ref{sec:spectempl} to similar well-observed SNe to produce appropriate comparison spectra. Such spectra are necessary to compare the Cas~A spectra to other SNe as well as to compare the LE spectra to each other. We show that the Cas~A SN was indeed very similar to the prototypical Type~IIb SN~1993J, as claimed by \\citet{Krause08a}. In \\S \\ref{sec:speccomp} we demonstrate that, despite the excellent agreement between the Cas~A spectra and SN~1993J, one LE has a systematically higher ejecta velocity than either SN~1993J \\textit{or the other LEs}, revealing that Cas~A was an intrinsically asymmetric explosion; observers from different directions would have viewed a ``different'' SN spectroscopically. In this section we also discuss the implications of our finding for both other historical SNe and for core-collapse SNe and their explosions. ", "conclusions": "\\label{sec:conclusions} We have obtained optical spectra of LEs from three different perspectives of the Cas~A SN, effectively probing different regions of the SN photosphere --- the first time that this technique has been applied to a SN. The spectra are very similar to each other and are all similar to the prototypical SN~IIb~1993J. After accounting for the window function determined by the combination of dust inclination and slit orientation, we are able to precisely compare Cas~A to other SNe as well as compare the LEs to each other. From these comparisons, two of the three directions have spectra which are indistinguishable from that of SN~1993J; however, one direction has \\ion{He}{1} and H$\\alpha$ P-Cygni features that are significantly blueshifted (\\about 4000~\\kms) relative to SN~1993J and the other two directions, indicating a higher ejecta velocity from Cas~A \\textit{in that one direction}. This is direct and independent evidence of an asymmetric explosion. The spectrum for the discrepant LE has an H$\\alpha$ line profile consistent with that of the high-luminosity SN~IIb~2003bg, but its \\ion{He}{1} $\\lambda 5876$ line profile had an even higher velocity than that of SN~2003bg. This may indicate that Cas~A had a very thin hydrogen layer, significant ejecta mixing, or different ionization structure in this direction. All LE spectra have \\ion{Ca}{2} NIR triplet line profiles consistent with each other as well as with those of SNe~1993J and 2003bg. This suggests that the emitting region of the Ca is distributed more spherically than that of the H or He-emitting regions. Even though there seems to be a ``jet-like'' structure in the NE corner and a counterjet in the SW corner \\citep{Hwang04, Fesen06_expansionasym_age}, recent optical and X-ray data from the Cas~A SNR indicate that the dominant Cas~A SN outflow is in the SE at a position angle of \\about 115\\arcdeg, slightly tilted toward the observer, and its counterpart approximately on the opposite side \\citep{Burrows05, Wheeler08, Delaney10}. Our detection of a blueshift looking into the counter outflow in the NW corner is the first direct, unambiguous, and independent confirmation of this outflow. It is also in excellent agreement with the apparent motion of the compact object, which moves at a position angle of $169\\arcdeg \\pm 8.4\\arcdeg$ \\citep{Tananbaum99, Fesen06_CO} away from the center of the SNR. Finally, we note that the existing surveys for LE features in this portion of the Galactic plane are far from complete, and that additional LEs are very likely to be discovered, providing additional perspectives of the SN in three dimensions. The inventory of such features will further illuminate the degree of asymmetry of the SN, but will also serve the purpose of testing the degree of coherence of spectra from similar perspectives. Since the spectrum at a given dust-concentration location is the result of integration over an entire hemisphere of SN photosphere, spectral differences are expected to vary slowly with changes in the perspective angle." }, "1003/1003.0759_arXiv.txt": { "abstract": "Planetary nebulae (PNe) derive from the evolution of $\\sim$1--8 M$_{\\odot}$ mass stars, corresponding to a wide range of progenitor ages, thus are essential probes of the chemical evolution of galaxies, and indispensable to constrain the results from chemical models. We use an extended and homogeneous data set of Galactic PNe to study the metallicity gradients and the Galactic structure and evolution. The most up-to-date abundances, distances (calibrated with Magellanic Cloud PNe), and other parameters have been employed, together with a novel homogeneous morphological classification, to characterize the different PN populations. We confirm that morphological classes have a strong correlation with PN Peimbert's Type, and also with their distribution on the Galactic landscape. We studied the $\\alpha$-element distribution within the Galactic disk, and found that the best selected disk population (i.e., excluding bulge and halo component), together with the most reliable PN distance scale yields to a radial oxygen gradient of $\\Delta$log(O/H)/$\\Delta$R$_{\\rm G}$=-0.023$\\pm$0.006 dex kpc$^{-1}$ for the whole disk sample, and of $\\Delta$log(O/H)/$\\Delta$R$_{\\rm G}$= -0.035$\\pm$0.024, -0.023$\\pm$0.005, and -0.011$\\pm$0.013 dex kpc$^{-1}$ respectively for Type I, II, and III PNe, i.e., for high-, intermediate-, and low-mass progenitors. Neon gradients for the same PN types confirm the trend. Accurate statistical analysis show moderately high uncertainties in the slopes, but also confirm the trend of steeper gradient for PNe with more massive progenitors, indicating a possible steepening with time of the Galactic disk metallicity gradient for what the $\\alpha$-elements are concerned. We found that the metallicity gradients are almost independent on the distance scale model used, as long as these scales are equally well calibrated with the Magellanic Clouds. The PN metallicity gradients presented here are consistent with the local metallicity distribution; furthermore, oxygen gradients determined with young and intermediate age PNe show good consistency with oxygen gradients derived respectively from other young (OB stars, H~II regions) and intermediate (open cluster) Galactic populations. We also extend the Galactic metallicity gradient comparison by revisiting the open cluster [Fe/H] data from high resolution spectroscopy. The analysis suggests that they could be compliant with the same general picture of a steepening of gradient with time. ", "introduction": "It has been assumed over the years that the building of galactic disks is driven inside-out by the continuous accretion of in-falling gas (Larson 1976). Evidences for the validity of the inside-out paradigm are still slim however, and are actively searched for. Inspection of detailed CDM simulations have shown that this picture may not be universal, and may vary from one galaxy to the other, to the point that some simulations show examples where galaxies can build their disk outside-in, at least partly (see e.g Sommer-Larsen et al. 2003, Roberston et al. 2004). In this respect, radial variations of properties of the Galactic disk may provide important observational constraints, but this has been hampered by the lack of tracers for which distance can be measured with the relevant accuracy, explaining why the magnitude of abundance gradients in the disk is still actively debated. Nonetheless, it can be considered that a consensus on the existence of gradients has been reached, and there is little doubt that inner regions are, in the mean, more metal rich than the outer parts of the disk (Davies et al. 2009), and this effect can be seen both in the iron-peak and the $\\alpha$-elements. Also, recent understanding of the local metallicity distribution and kinematics strongly suggest that the solar radius is polluted by wanderers coming from the outer and inner disk, indicating a rather strong radial variation of the metallicity (Haywood 2008). Because radial metallicity gradients in galactic disks may be produced by a number of different processes, the measurement of how these gradients have evolved with time should provide an even stronger constraint (e.~g., Fu et al. 2009), thus the importance of studying metallicity distributions of tracers of different Galactic ages. Planetary nebulae (PNe) represent the Galactic stellar population with progenitors of turnoff mass M$_{\\rm to}$ between $\\sim$1 and 8 M$_{\\odot}$, probing Galactic ages between $\\sim3\\times10^7$ yr and $\\geq$10 Gyr (Maraston 1998). Planetary nebulae are distributed in the Galactic disk, bulge, and halo. Their Galactic distribution and radial velocities offer a first subdivision into populations, which can be studied separately. Furthermore, markers such as nitrogen abundances and morphological types allow to further discriminate between disk PNe of relatively young and old progenitors. It is then sensible to undergo a study of Galactic metallicity for the different PN populations. This has been done in the past. Galactic disk PNe have been studied, among others, by Kingsburg \\& Barlow (1994) and, more recently, by Perinotto et al. (2004, P04), while bulge PNe are the subject of a focused study by Exter et al. (2004). The Galactic disk PNe have been the subject of many studies on metallicity gradients: $\\alpha$-element abundances of PNe trace the original progenitor composition, thus the chemical evolution of the Galactic disk. Perinotto \\& Morbidelli (2006, PM06) reviewed all Galactic metallicity gradients published since the seventies, and found that PNe in the Galactic disk trace a one dimensional, negative oxygen gradient $\\Delta$log(O/H)/$\\Delta$R$_{\\rm G}$=-0.07 (Faundez-Abans \\& Maciel 1987) and -0.03 (Pasquali \\& Perinotto 1993) dex kpc$^{-1}$. The excellent data revision by Perinotto and collaborators provides its own gradient of -0.016 dex kpc$^{-1}$ (PM06), which is shallower than previously known. Almost contemporary Stanghellini et al. (2006) found a shallow gradient of -0.01 dex kpc$^{-1}$ for the oxygen abundances of Galactic disk PNe, consistent with PM06's analysis. The abundance analysis that yield to the gradients published so far might be marginally different, but hardly enough to make the metallicity gradients to differ outside the data uncertainties. A stronger factor for divergence is certainly the method of Galactic PN distance calculation, although PM06 showed that the gradients should result rather flat independent on the distance scale used. The distance scale of PNe used is also a determining factor in the study of the different Galactic PN populations, their distance from the Galactic plane and their belonging to the thick disk, the bulge, the halo. Until a couple of years ago, distances to Galactic PNe where estimated by means of statistical distance scales, calibrated with a few known distances to Galactic PNe. Now, with the availability of a wealth of spatially resolved images of Magellanic Cloud PNe from {\\it HST}, the calibration has been done in a much sounder way (Stanghellini et al. 2008, SSV). There is then purpose to reanalyze both the Galactic distribution of PNe and the metallicity gradients in a homogeneous way, with the distance scales based on the physical parameters of Magellanic Cloud PNe, whose distances are known. In this paper we use the Magellanic Cloud-calibrated Galactic PN distance scale from SSV to examine the population of PNe in our Galaxy. We extend the P04 abundance database with all the abundances published more recently, and calculate the metallicity gradients for the different populations, and homogeneously classify the Galactic PN morphology for further insight on population selection. In $\\S$2 we describe the database built for this study; in $\\S$3 we discuss the Galactic structure based on PNe; in $\\S$4 we present the metallicity gradients from our analysis; sections 5 and 6 present respectively a discussion of our results and the conclusions of this work. ", "conclusions": "A study on PN metallicity gradients employing the most updated abundances and distance scale shows that gradients of $\\alpha$-elements across the Galactic disk are moderate, and depend on the age of the PN population considered. The amplitude of these gradients agree with the metallicity dispersion observed in the solar vicinity, which arises mainly from the radial mixing of stars through the disk. Planetary nebulae with young progenitors show steeper gradient slopes than in the old populations, indicating that $\\alpha$-element gradients are steepening with time. The results are statistically sound, and do not depend very much on the distance scale, as long as it is Magellanic Cloud-calibrated. Gradients of oxygen and iron abundance from young (young stars and H~II regions) and intermediate-age (open clusters) populations agree with those of PNe derived here, giving strength to the scenario of gradients steepening with time. The oxygen gradient is found to increase from near-flat radial abundance distribution in Type III PNe, those whose progenitors are in the lowest mass range for AGB stars, to -0.035 dex kpc$^{-1}$ for Type I PNe, those with the most massive progenitors. These results contradict recent claim that gradients are flattening with time, and strongly suggest that such constraints should be taken with some caution. The data on PNe are still sparse and the new distance scale, although significantly improved, could still be affected by uncertainties. In addition, abundance gradients for older populations of PNe (in particular the Type III sample) might be affected by migration, which tends to flatten such gradients. However, we find it unlikely that gradients in the old PN population could reach values as high as those found in some recent studies (e.g Maciel \\& Costa 2009), and this is confirmed by the old open cluster population with high resolution spectroscopic abundance determinations." }, "1003/1003.4666_arXiv.txt": { "abstract": "We review the evidence for the argument that Rudolf Wolf's calibration of the Sunspot Number is likely to be correct and that Max Waldmeier introduced an upwards jump in the sunspot number in 1945. The combined effect of these adjustments suggests that there has been no secular change in the sunspot number since coming out of the Maunder Minimum $\\sim$1715. ", "introduction": "The Sunspot Record goes back 400 years and is the basis for many reconstructions of solar parameters (e.g. TSI), but, how good is it? And can we agree on which one (Wolf Number, International Number, `Boulder' Number, Group Number, ...)? Are the old values good? Are the new ones? And what is a `good' or `correct' Sunspot Number anyway? Johann Rudolf \\citet{wo1859} defined his Relative Sunspot Number, taking into account both individual spots and their appearance in distinct groups (what we today call `active regions'), as \\(R_W=10\\,Groups+Spots\\). Wolf started his own observations in 1849 and assembled observations from earlier observers back to 1749 and beyond (Figure 1). \\begin{figure}[!ht] \\includegraphics [width=\\textwidth]{fig1.png} \\caption{Rudolf Wolf and excerpts from his 1861 list of published Relative Numbers compared with his latest list (now the official list from SIDC in Brussels).} \\end{figure} As is clear, the earlier values were subsequently adjusted (upwards) as Wolf were struggling with the difficulty of bringing different observers onto the same `scale', compensating for telescope size, counting method, acuity, seeing, and personal bias. Wolf published several versions of his celebrated Relative Sunspot Numbers based on data gathered from many observers from both before and during Wolf's own lifetime (Figure 2). How to `harmonize' data from different observers? \\begin{figure}[!ht] \\includegraphics [width=\\textwidth]{fig2.png} \\caption{Evolution of the Wolf Number from his first 1857 list to the final version, with color coded symbols and curves for each list} \\end{figure} ", "conclusions": "Modern data shows that the diurnal range of the geomagnetic variation is an extremely good proxy for the solar microwave flux. To the extent that we take the flux to be a measure of general solar activity of which the sunspot number was meant to be an indicator, we argue that Wolf's calibration makes his sunspot series essentially an equivalent F10.7 series. Accepting the soundness of Wolf's procedure and correcting for the Waldmeier discontinuity (+20\\%) lead to a picture of solar activity with but little difference between activity levels in the 18$^{th}$, 19$^{th}$, and 20$^{th}$ centuries (Figure 15). \\begin{figure}[!ht] \\includegraphics [width=\\textwidth]{fig15.png} \\caption{Suggested equivalent sunspot numbers calibrated by the geomagnetic record.} \\end{figure}" }, "1003/1003.1275_arXiv.txt": { "abstract": "A 2D particle simulation models the collision of two electron-ion plasma clouds along a quasi-parallel magnetic field. The collision speed is 0.9c and the density ratio 10. A current sheet forms at the front of the dense cloud, in which the electrons and the magnetic field reach energy equipartition with the ions. A structure composed of a solenoidal and a toroidal magnetic field grows in this sheet. It resembles that in the cross-section of the torus of a force-free spheromak, which may provide the coherent magnetic fields in gamma-ray burst (GRB) jets needed for their prompt emissions. ", "introduction": " ", "conclusions": "" }, "1003/1003.3588.txt": { "abstract": "% % Context {Globular clusters with their large populations of millisecond pulsars (MSPs) are believed to be potential emitters of high-energy gamma-ray emission. The observation of this emission provides a powerful tool to assess the millisecond pulsar population of a cluster, is essential for understanding the importance of binary systems for the evolution of globular clusters, and provides complementary insights into magnetospheric emission processes.} % % Aims {Our goal is to constrain the millisecond pulsar populations in globular clusters from analysis of gamma-ray observations.} % % Methods {We use 546 days of continuous sky-survey observations obtained with the Large Area Telescope aboard the \\textit{Fermi} Gamma-ray Space Telescope to study the gamma-ray emission towards 13 globular clusters.} % % Results {Steady point-like high-energy gamma-ray emission has been significantly detected towards 8 globular clusters. Five of them (47~Tucanae, Omega~Cen, NGC~6388, Terzan~5, and M~28) show hard spectral power indices $(0.7 < \\Gamma <1.4)$ and clear evidence for an exponential cut-off in the range $1.0-2.6$~GeV, which is the characteristic signature of magnetospheric emission from MSPs. Three of them (M~62, NGC~6440 and NGC~6652) also show hard spectral indices $(1.0 < \\Gamma < 1.7)$, however the presence of an exponential cut-off can not be unambiguously established. Three of them (Omega Cen, NGC~6388, NGC~6652) have no known radio or X-ray MSPs yet still exhibit MSP spectral properties. From the observed gamma-ray luminosities, we estimate the total number of MSPs that is expected to be present in these globular clusters. We show that our estimates of the MSP population correlate with the stellar encounter rate and we estimate $2600-4700$ MSPs in Galactic globular clusters, commensurate with previous estimates.} % % Conclusions {The observation of high-energy gamma-ray emission from globular clusters thus provides a reliable independent method to assess their millisecond pulsar populations.} ", "introduction": "\\label{sec:intro} With their typical ages of $\\sim10^{10}$ years, globular clusters form the most ancient constituents of our Milky Way Galaxy. These gravitationally bound concentrations of ten thousand to one million stars are surprisingly stable against collapse which implies some source of internal energy that balances gravitation. The potential energy of binary systems is a plausible source of this internal energy, tapped by close stellar encounters that harden the orbits of the systems \\citep{hut92}. Indeed, globular clusters contain considerably more close binary systems per unit mass than the Galactic disk which eventually show up as rich populations of X-ray binaries \\citep{clark75}. This scenario is strengthened by the observation that the number of low-mass X-ray binary systems containing neutron stars is directly correlated with the stellar encounter rate \\citep{gendre03}. Another consequence of this scenario is the presence of many millisecond pulsars\\footnote{See http://www.naic.edu/$\\sim$pfreire/GCpsr.html for an updated list.} \\citep[hereafter MSPs; see e.g.][]{camilo05,ransom08}, also known as `recycled' pulsars, i.e. pulsars that were spun-up to millisecond periods by mass-accretion from a binary companion \\citep{alpar92}. Observations with the {\\em Large Area Telescope} (LAT) onboard the {\\em Fermi Gamma-ray Space Telescope} have confirmed MSPs as gamma-ray sources \\citep{abdo09a,abdo09b}. The spectral energy distribution of millisecond pulsars is characterised by hard ($1.0\\lesssim\\Gamma\\lesssim2.0$) power law spectra with exponential cut-offs in the $1-3$~GeV energy range \\citep{abdo09b}. Recently, \\cite{abdo09c} presented the first detection of a globular cluster (GC) in the gamma-ray domain. This GC, 47 Tuc, has a spectral energy distribution best described by a photon index of 1.3$\\pm$0.3 with a cut-off energy of $2.5^{+1.6}_{-0.8}$ GeV \\citep{abdo09c} typical of the other MSPs detected to date \\citep{abdo09b}. Further, 47 Tuc contains at least 23 MSPs, known from radio and X-ray observations. The lack of variability over days to months is consistent with MSP emission. In addition, folding the data on known ephemerides from the 47 Tuc pulsars reveals no significant detections, thus it appears that the gamma-ray emission is not due to a single MSP but rather attributable to an entire population of MSPs in this globular cluster. Using the observed, average efficiency of converting spin down energy into the observed gamma-ray luminosity, constraints can be placed on the MSP population \\citep{abdo09c}. As the number of neutron star X-ray binaries are correlated with encounter rate and MSPs are the progeny of these systems, it would follow that the number of MSPs per globular cluster scales in a similar way. It is difficult to test such a correlation using radio and X-ray observations as the former are affected by dispersion and scattering by the turbulent ionized interstellar medium, in particular for clusters near the Galactic bulge, while the latter are affected by interstellar absorption rendering the detection difficult due to the low count rates observed. The gamma-ray domain is not affected by interstellar absorption and there is also the added advantage that the gamma-ray beams may be wider than the radio/X-ray beams \\citep[e.g.][]{abdo10a}, which would permit more MSPs to be detected in the gamma-ray domain than those at lower energies, thus making gamma-ray observations ideal for testing such a correlation. In this paper we consider gamma-ray sources that are spatially consistent with GCs and that show the spectral characteristics of MSPs, i.e. that have hard power law spectra with exponential cut-offs in the few GeV regime, that are steady, and that are point-like. We analyse {\\it Fermi} LAT data for 13 globular clusters (see Table \\ref{tab:gc}) and include in our list the 8 globular clusters that have been formally associated with sources in the first year {\\it Fermi} LAT catalogue \\citep[][hereafter named 1FGL sources]{abdo10b}. We add two further globular clusters that lie spatially close to 1FGL sources (Omega Cen and NGC~6624), and include also NGC~6441 due to the high stellar collision rate that is believed to favour the formation of MSPs \\citep{freire08}, NGC~6752 due to its relative proximity of 4 kpc \\citep{damico02}, and M~15 due to its relatively large population of known MSPs \\citep{anderson93}. %%% GC analysed in this work %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{table}[!t] \\footnotesize \\caption{Globular clusters analysed in this work. \\label{tab:gc} } \\begin{center} \\begin{tabular}{llll} \\hline \\hline \\noalign{\\smallskip} Name & Other name & MSPs & Reason for inclusion \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 47 Tucanae & NGC~104 & 23 & 1FGL~J0023.9$-$7204 \\\\ Omega Cen & NGC~5139 & 0 & close to 1FGL~J1328.2$-$4729 \\\\ M~62 & NGC~6266 & 6 & 1FGL~J1701.1$-$3005 \\\\ NGC~6388 & ... & 0 & 1FGL~J1735.9$-$4438 \\\\ Terzan~5 & ... & 33 & 1FGL~J1747.9$-$2448 \\\\ NGC~6440 & ... & 6 & 1FGL~J1748.7$-$2020 \\\\ NGC~6441 & ... & 4 & high collision rate \\\\ NGC~6541 & ... & 0 & 1FGL~J1807.6$-$4341 \\\\ NGC~6624 & ... & 4 & close to 1FGL J1823.4$-$3009 \\\\ M~28 & NGC~6626 & 12 & 1FGL~J1824.5$-$2449 \\\\ NGC~6652 & ... & 0 & 1FGL~J1835.3$-$3255 \\\\ NGC~6752 & ... & 5 & 5 MSPs, nearby \\\\ M~15 & NGC~7078 & 8 & 8 MSPs \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{center} {\\bf Notes.} The known number of MSPs (column 3) has been taken from http://www.naic.edu/$\\sim$pfreire/GCpsr.html. \\end{table} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % Observations %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ", "conclusions": "\\label{sec:conclusion} An analysis of {\\it Fermi} LAT data from 13 globular clusters has revealed 8 significant, point-like and steady gamma-ray sources that are spatially consistent with the locations of the clusters. Five of them (47~Tuc, Omega Cen, NGC~6388, Terzan~5, and M~28) show hard spectral power law indices $(0.7 < \\Gamma < 1.4)$ and clear evidence for an exponential cut-off in the range $1.0-2.6$ GeV, which is the characteristic signature of magnetospheric emission from MSPs. We thus classify these 5 sources as {\\it plausible} globular cluster candidates. Three of them (M~62, NGC~6440 and NGC~6652) also show hard spectral indices $(1.0 < \\Gamma < 1.7)$, however the presence of an exponential cut-off cannot unambiguously be established. More data are required before definite conclusions can be drawn; hence we qualify these 3 sources as {\\it possible} globular cluster candidates. From the 8 globular clusters that are associated with significant gamma-ray sources, 5 are known to harbour MSPs. In Omega Cen, NGC~6388 and NGC~6652, however, no MSPs have so far been detected, neither by radio nor by X-ray observations. The observation of gamma-ray signatures that are characteristic of MSPs provides strong support that these GCs indeed also harbour important populations of MSPs. In particular, we predict from the observed gamma-ray luminosities that the total MSP populations amount to $10-30$ (Omega Cen), $80-300$ (NGC~6388), and $30-80$ (NGC~6652) in these clusters. Deep radio and X-ray follow-up observations may help to unveil first members of these populations. Our predicted number of MSPs shows evidence for a positive correlation with the stellar encounter rate in a similar way to their progenitors, the neutron star low mass X-ray binaries. This correlation allows us to deduce the total number of MSPs in Galactic globular clusters ($2600-4700$) which lies midway between all previous estimates, supporting such a correlation. Such an estimate can be used to derive constraints on the original neutron star X-ray binary population, essential for understanding the importance of binary systems in slowing the inevitable core collapse of globular clusters. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % Acknowledgments %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%" }, "1003/1003.4450_arXiv.txt": { "abstract": "{}{We present for the first time a coherent model of the polarized Galactic synchrotron and thermal dust emissions which are the main diffuse foreground for the measurement of the polarized power spectra of the CMB fluctuations with the Planck satellite mission. }{We produce 3D models of the Galactic magnetic field including regular and turbulent components, and of the distribution of matter in the Galaxy, relativistic electrons and dust grains. By integrating along the line of sight we construct maps of the polarized Galactic synchrotron and thermal dust emission for each of these models and compare them to currently available data. We consider the 408 MHz all-sky continuum survey, the 23 GHz band of the Wilkinson Microwave Anisotropy Probe and the 353 GHz Archeops data.}{ The best-fit parameters obtained are consistent with previous estimates in the literature based only on synchrotron emission and pulsar rotation measurements. They allows us to reproduce the large scale structures observed on the data. Poorly understood local Galactic structures and turbulence make difficult an accurate reconstruction of the observations in the Galactic plane.} {Finally, using the best-fit model we are able to estimate the expected polarized foreground contamination at the Planck frequency bands. For the CMB bands, 70, 100, 143 and 217~GHz, at high Galactic latitudes although the CMB signal dominates in general, a significant foreground contribution is expected at large angular scales. In particular, this contribution will dominate the CMB signal for the B modes expected from realistic models of a background of primordial gravitational waves.} ", "introduction": "\\label{introduction} \\indent The PLANCK satellite mission, currently in flight, should permit the more accurate measurements of the CMB anisotropies both in temperature and polarization. Planck, which observes the sky on a wide range of frequency bands from 30 to 857~GHz, has a combined sensitivity of $\\frac{\\Delta T}{T_{CMB}} \\sim 2 \\frac{\\mu K}{K}$ and an angular resolution from 33 to 5 arcmin~(\\cite{bluebook}). Of particular interest is the measurement of the polarization B modes which implies the presence of tensor fluctuations from primordial gravitational waves generated during inflation. Planck should be able to measure the tensor-to-scalar ratio, $r$, down to 0.1 (\\cite{betoules2009,efstathiou1}) in the case of a nominal mission (2 full-sky surveys) and to 0.05 for the Extended Planck Mission: four full-sky surveys (\\cite{efstathiou2}) . The value of $r$ sets the energy scale of the inflation (\\cite{peiris}) and then provides constraints on inflationnary models (\\cite{baumann}).\\\\ \\\\ \\indent To achieve this high level of sensitiviy it is necessary to accurately estimate the temperature and polarization foregrounds mainly from diffuse Galactic emission components -- synchrotron, thermal and rotational dust, and free-free -- as well as from the Galactic and extra-Galactic point-like and compact sources. Indeed, at the Planck frequency bands these foreground components may dominate the CMB signal and therefore, they need to be either masked or subtracted prior to any CMB analysis. For this purpose, the Planck collaboration plans to use component separation techniques (see \\cite{leach} for a summary) in addition to the traditional masking of highly contaminated sky regions including identified point-like and compact sources. As these component separation techniques will be mainly based on Planck data only, one of the main issues will be to estimate the residual foreground contamination on the final CMB temperature and polarization maps. These residuals will translate into systematic biases and larger error bars on the estimation of the temperature and polarization power spectra of the CMB fluctuations (see \\cite{betoules2009} for a recent study). Thus, they will impact the precission to which cosmological information can be retrieved from the Planck data. \\\\ \\\\ \\indent For polarization the main foreground contributions will come from the diffuse Galactic synchrotron and thermal dust emission. From the WMAP observations, \\cite{page2007} have shown that the radio synchrotron emission from relativistic electrons is highly polarized, up to 70~\\%, between 23 and 94GHz. Furthermore, \\cite{benoit2004a,ponthieu2005} have observed signicantly polarized thermal dust emission, up to 15 \\% at the 353~GHz Archeops channel. By contrast the diffuse free-free emision is not polarized and the anomalous microwave emission has been measured to be weakly polarized, $3 ^{+1.3}_{-1.9}$ \\%, (\\cite{battistelli2006}). Finally, at the Planck frequency bands the polarized contribution from compact and point sources is expected to be weak for radio (\\cite{nolta2009}) and dust (\\cite{desert2008}) sources. The spatial and frequency distribution of both Galactic synchrotron and thermal dust polarized emission at the Planck frequencies are not well known and the only available information comes from micowave and submillimetre observations. For synchrotron, Faraday rotation (\\cite{burn1966}) makes it very difficult to extrapolate the polarized observed radio emission (\\cite{wolleben,wolleben07, carretti2009}) to the microwave domain. For thermal dust, polarized observations are not currently available in the infrared and the current optical data (\\cite{heiles}) are too sparse (\\cite{page2007}) for a reliable extrapolation to lower frequencies. \\\\ \\indent{The diffuse Galactic synchrotron emission is produced by relativistic electrons spiraling around the Galactic magnetic field lines and its polarization is orthogonal both to the line-of-sight and to the field direction (\\cite{ribicki}). Based on these statements, \\cite{page2007} proposed to model the polarized synchrotron Galactic emission observed by the WMAP satellite using a 3D model of the Galaxy including the distribution of relativistic electrons and the Galactic magnetic field structure. Although this model allowed them to explain the observed polarization angle at the 23~GHz band where the synchrotron emission dominates, it was not used for the CMB analysis. Instead they constructed a template of the polarized synchrotron emission from the 23 GHz band and extrapolated it to higher frequencies. Independently, \\cite{han2004,han2006} has also proposed a 3D model of the Galactic free electrons (\\cite{cordes}) and of the Galactic magnetic field including a regular and a turbulent component to explain the observed rotation measurements on known pulsars. Based on previous works \\cite{sun} has performed a combined analysis of the polarized WMAP data and of the rotation measurements of pulsars using the publicly available HAMMURABI code (\\cite{waelkens}) for computing the integrated emission along th line-of-sight. This work has been extended by \\cite{jaffe} for the study of the Galactic plane using a MCMC algorithm for the determination of the parameter of the models and by \\cite{jansson} for the full sky using a likelihood analysis for parameter estimation. \\\\ \\indent Dust grains in the Interstellar Medium (ISM) are heated by stellar radiation and radiate in the form of thermal dust emission (\\cite{desert1998}). They are considered to be oblate and to align with their longitudinal axis perpendicular to the magnetic field lines (\\cite{davis}). When aligned they will end up rotating with their angular moment parallel to the magnetic field direction. As the thermal dust emission is more efficient along the long axis, this generates linear polarization orthogonal to the magnetic field direction and to the line-of-sight. The polarization fraction of the emission depends on the distribution of the size of the grains and is about a few percent at the millimeter wavelengths (\\cite{hildebrand1999, vaillancourt}). \\cite{ponthieu2005} concluded that the polarized emission observed in the 353~GHz Archeops data was associated with the thermal dust emission and proposed a simple magnetic field pattern to explain the measured direction of polarization on the Galactic plane. \\cite{page2007} suggested that part of the observed polarized emission at the 94~GHz WMAP data was also due to thermal dust. They modeled it using the observed polarization direction of stellar light (\\cite{heiles}) which from the above statement must be perpendicular to that of thermal dust. \\\\ \\indent With the prospect of the data analysis of the Planck satellite mission in mind, we present here consistent physical models of the synchrotron and thermal dust emission based on the 3D distribution of relativistic electrons and dust grains on the Galaxy, and on the 3D pattern of the Galactic magnetic field. The paper is structured as follows: Sect.~\\ref{data} describes the 408 MHz all-sky continnum survey (\\cite{haslam}), the five-year WMAP data set (\\cite{page2007}) and the Archeops data (\\cite{ponthieu2005}) used in the analysis. In Sect.\\ref{3dgal_model} we describe in detail the models. Section~\\ref{gal_comp} describes the statistical comparison of data and models. In Sect.~\\ref{gal_bias} we discuss the impact of polarized foreground emission on the measurement of the polarized CMB emission with the Planck satellite. Finally, conclusions are presented in Section \\ref{conc}. ", "conclusions": "\\label{conc} \\indent We have presented in this paper a detailed study of the diffuse synchrotron and thermal dust polarized Galactic foreground emission components. We have constructed coherent models of these two foregrounds based on a 3D representation of the Galactic magnetic field and of the distribution of relativistic electrons and dust grains in the Galaxy. For the Galactic magnetic field we have assumed a large-scale regular component plus a turbulent one. The relativistic electron distribution and dust grain distribution have been modeled as exponentials peaking at the Galactic center. From these analyses we have been able to study the main parameters of the models, the magnetic field pitch angle, $p$, the radial width of the relativistic electron distribution, $h_{er}$, the relative amplitude of the turbulent component, $A_{turb}$ and spectral index of the synchrotron emission $\\beta_{s}$. We have been able to set constraints only on the pitch angle and the synchrotron spectral index. An upper limit on the relative amplitude of the turbulent component is obtained although the data seems to prefer no turbulence at large angular scales. With the current data we are not able to constrain the radial width of the relativistic electron distribution. Notice that our constraints are compatible with those in the literature. \\indent Using the best--fit parameters we have constructed maps in temperature and polarization for the synchrotron and dust thermal emission at 23 and 353~GHz and compare them to the WMAP and ARCHEOPS data at the same frequencies. We find good agreement between the data and the model. However, when comparing the temperature and polarization power spectra for the data and model maps, we observe that synchrotron emission model is not realistic enough. For dust the model seems to reproduce better the data but it is important to realize that the errors on the Archeops data are much larger. \\indent From this, we can conclude that the models presented in this paper can not be used for direct subtraction of polarized foregrounds for CMB purposes. However they can be of great help for estimating the impact of the polarized Galactic foreground emission on the reconstruction of the CMB polarized power spectra. Indeed, we have extrapolated the expected polarized Galactic foreground emission to the Planck CMB frequencies, 70, 100, 143 and 217~GHz and found they dominate the emission at low $\\ell$ values where the signature in the polarized CMB power spectra of important physical processes like reionization is expected. Furthermore, the Galactic polarized foreground emission seems to dominate the B modes for which we expect an unique signature from primordial gravitational waves. Because of this, we propose the use of models like the ones presented in this paper to asses the errors in the reconstruction of the CMB emission when using template subtraction techniques or component separation algorithms." }, "1003/1003.3877_arXiv.txt": { "abstract": "The spin probability distribution of Dark Matter haloes has often been modelled as being very near to a lognormal. Most of the theoretical attempts to explain its origin and evolution invoke some hypotheses concerning the influence of tidal interactions or merging on haloes. Here we apply a very general statistical theorem introduced by Cram\\'{e}r (1936) to study the origin of the deviations from the reference lognormal shape: we find that these deviations originate from correlations between two quantities entering the definition of spin, namely the ratio $J/M^{5/2}$ (which depends only on mass) and the modulus $E$ of the total (gravitational + kinetic) energy.\\\\ \\noindent To reach this conclusion, we have made usage of the results deduced from two high spatial- and mass resolution simulations. Our simulations cover a relatively small volume and produce a sample of more than 16,000 gravitationally bound haloes, each traced by at least 300 particles. We verify that our results are stable to different systematics, by comparing our results with those derived by the \\textsc{GIF2} and by a more recent simulation performed by Macci\\`{o} et al.\\\\ We find that the spin probability distribution function shows systematic deviations from a lognormal, at all redshifts $z \\la 1$. These deviations depend on mass and redshift: at small masses they change little with redshift, and also the best lognormal fits are more stable. The $J-M$ relationship is well described by a power law of exponent $\\alpha$ very near to the linear theory prediction ($\\alpha=5/3$), but systematically lower than this at z$\\la 0.3$. We argue that the fact that deviations from a lognormal PDF are present only for high-spin haloes could point to a role of large-scale tidal fields in the evolution of the spin PDF. ", "introduction": "Within the hierarchical galaxy formation model, Dark Matter (hereafter DM) haloes are thought to play the role of gravitational building blocks, within which baryonic diffuse matter collapses and becomes detectable. On galactic scales, the formation of stars and their evolution provides an important probe of the evolution of the visible content of the Universe \\citep{1978MNRAS.183..341W,1991ApJ...379...52W}, although the subtleties of the stellar formation processes within galaxies, as of today not yet completely understood, hinders an exploitation of these objects as a clean probe of the evolution of DM haloes. On the other extreme of the mass scale, the most massive clusters of galaxies are regarded as one of the most reliable cosmological probes \\citep{1997ApJ...485L..53B}: in particular, their abundance and evolution with redshift is a very sensitive test of the underlying cosmological model \\citep{1993Natur.366..429W,1997ApJ...489L...1H,1997ApJ...485L..53B,1997A&A...317....1O,1998MNRAS.298.1145E,1999ApJ...523L.137D,2000ApJ...534..565H,2002A&A...386..775A,2004ApJ...609..603H,2007ApJ...655..128G}.\\\\ \\noindent The Mass Function (hereafter MF) is an indirect test of the total virialised mass of DM haloes: exact predictions of the latter can be done using the nonlinear spherical collapse model \\citep{1972ApJ...176....1G}, an essential ingredient of the Press-Schechter model for the MF. However, DM haloes also possess angular momentum, and taking into account its effect on the gravitational collapse of rotating shells has been shown to have a detectable consequence on the MF \\citep{2006A&A...454...17D,2009ApJ...698.2093D}. Thus, an investigation of the distribution of angular momentum, and its connection with mass is particularly useful, as high-resolution simulations of increasing resolution will produce MFs with a very small statistical uncertainty.\\\\ \\noindent An exact determination of the shape of the DM halo spin PDF can also have important consequences for the abundance of \\emph{low surface brightness galaxies}, if the latter form preferentially within high-spin DM haloes \\citep{1998MNRAS.299..123J, 2007MNRAS.378...55M}. Finally, it is of great importance also in models of the formation and evolution of central Black Holes. In collapse models where most of the baryons' specific angular momentum is a fixed fraction of that of their host dark matter haloes \\citep{1980MNRAS.193..189F,1998MNRAS.295..319M}, the angular momentum and extent of the gaseous central accretion disc are strongly dependent on their's spin. \\citet{2005ApJ...633..624V} find that the central density of the disc varies as $\\rho_{0} \\simeq\\lambda^{-4}$, thus the initial rate of accretion of the central Black Hole turns out to be a very sensitive function of the spin $\\lambda$ \\citep{1971A&A....11..377P}.\\\\ \\noindent For all these reasons, investigations of the origin of the angular momentum growth and of the spin PDF of DM haloes, and of their evolution, can have an impact on many different open issues in large-scale structure formation and evolution. \\\\ \\noindent Theoretical investigations \\citep{1969ApJ...155..393P} predict that to zeroth-order the angular momentum should have a power-law dependency on the total virialized mass, with exponent $5/3$. They also make a prediction concerning the \\emph{spin}, a dimensionless quantity defined as: \\be \\lambda = \\frac{J E^{1/2}}{GM^{5/2}} \\label{eq:am:0} \\ee where $J\\equiv\\mid\\bmath{J}\\mid$, $E\\equiv\\mid E_{kin} + E_{grav} \\mid$ are respectively the moduli of the angular momentum, and of the total (kinetic plus potential) energy. The PDF of $\\lambda$ has been predicted to have an approximately lognormal distribution. This initial prediction was subsequently found to be basically valid also when higher order effects were taken into account \\citep{1996MNRAS.282..436C,1996MNRAS.282..455C}. More recently, the radial profile of $\\lambda$ has been derived from modified Jeans' equations \\citep{2009ApJ...694..893S}, and found to be in good agreement with results from simulations.\\\\ \\noindent The properties of the spin distribution of DM haloes have recently been considered, taking advantage of the availability of high spatial- and mass-resolution simulations. \\citet{2007MNRAS.376..215B} have analysed a large sample of haloes drawn from the Millennium simulation, and found (among other things) that the \\emph{global} spin distribution is poorly described by a lognormal distribution: \\be P(\\lambda) = \\frac{1}{\\lambda\\sqrt{2\\pi\\sigma_{\\log{\\lambda}}^{2}}}\\exp\\left(-\\frac{{\\rmn \\log}^{2}(\\lambda/\\bar\\lambda)}{2\\sigma_{\\log{\\lambda}}^{2}}\\right) \\label{eq:1} \\ee They alternatively suggest an empirical fit: \\be P(\\lambda) \\propto \\left(\\frac{\\lambda}{\\lambda_{0}}\\right)^{3}\\exp\\left[-\\alpha\\left(\\frac{\\lambda}{\\lambda_{0}}\\right)^{3/\\alpha}\\right] \\ee which has the shape of a lognormal except at very high and very low values of the spin. They also suggest that the actual shape of the distribution depends on the adopted \\emph{numerical definition} of halo.\\\\ \\noindent Theoretical and numerical studies are aimed at understanding the origin of the (almost) lognormal spin distribution, using a limited set of statistical and dynamical assumptions or by performing \\emph{controlled} numerical experiments. \\citet{2002ApJ...581..794C} have shown that the distribution of specific angular momentum (which they define as $j' = J/M^{5/3}$) can be described by a rather complex PDF, which can be approximated by a lognormal in the central part, but deviates significantly from it at low- and high values of $j'$. They also make predictions for the dependence of the peak of the spin distribution for different values of mass, suggesting that it varies linearly with the average $\\sigma_{\\delta}$ of the distribution. \\citet{2008ApJ...678..621K} find a correlation of spin $\\lambda$ with mass, albeit a weak one. Note that they use a set of simulations using boxes of varying sizes, while in the present work we adopt a single simulation, thus minimising the spurious tidal effects. Furthermore, \\citep{2007MNRAS.376..215B,2007MNRAS.377L...5G,2008MNRAS.389.1419L}, studying the spin distributions at low redshifts, have found that more massive haloes show larger values of $\\lambda$. Extensions of these results to higher redshift have been provided by \\citet{2009MNRAS.393.1498D}, who have studied the evolution of the spin distributions at $z > 6$, and showed that more massive haloes ($M\\simeq 10^{7} {\\rm M}_{\\odot}$) tend to have a median $\\lambda$ higher than that of $M\\simeq 10^{6} {\\rm M}_{\\odot}$ haloes. Also, high-$\\lambda$ haloes tend to cluster more (by a factor 3-5) than low-spin haloes, a trend which strengthens with time. However, their simulation is restricted to a small box ($L_{b} = 2.46 h^{-1} {\\rm Mpc}$), thus making their result more prone to uncertainties from cosmic variance.\\\\ \\noindent The steady improvement of the available hardware and software resources makes today possible simulations where the limits imposed by the finite spatial and mass resolution limits are challenged. The simulations we have performed in this work were aimed at providing a \\emph{statistically significant} sample of reasonably well-resolved DM haloes. We have obtained a catalogue of more than 77600 DM haloes, each resolved by more than 20 particles, and we have used only those haloes with more than 300 particles, thus resulting in a catalogue containing more than 16400 haloes. We have chosen a box size of $L = 70 \\, h^{-1}$ Mpc, smaller than the one used in previous papers \\citep[e.g.][]{2006ApJ...646..815S}, and a large number of particles, to maximize the mass resolution.\\\\ \\noindent In this work we present the results of a numerical experiment, where we study the evolution of the angular momentum/spin distributions within a relatively small cosmological volume. We find that the spin distribution shows some clear deviations from a lognormal, and does not seem to relax towards the latter, above statistical uncertainty. These deviations \\emph{do not seem to depend significantly on the redshift or the mass of the halo}, at least for the underlying LCDM model studied in this paper. We also perform a possible check on the role of possible systematics in our analysis by comparing with the results of two different simulations: the GIF2 \\citep{2008MNRAS.386.2135G}, and a simulation recently perfomed by \\citet{2008MNRAS.391.1940M}. Both these works adopted a Spherical Overdensity (SO) \\emph{halo finder}, slightly different from our chosen halo finder (the Amiga Halo Finder, \\textsc{AHF}), recently introduced by \\citet{2009ApJS..182..608K}. We address the problem of the origin of these deviations by exploiting an exact result from statistics \\citep[\\emph{Cram\\'{e}r's theorem},][]{Cramer..MatZeit..41..405..1936} to demonstrate that the deviations from a lognormal distribution are induced from \\emph{correlations} between the total energy and mass of haloes, as should be expected for not completely virialised haloes.\\\\ \\noindent The plan of the paper is as follows: In section~\\ref{sec:sims} we present the details of the simulations we have performed, briefly describing the code adopted, the initial conditions and the halo finder. We then describe the main results concerning $J-M$ relationship (section~\\ref{sec:jm}). In section~\\ref{sec:diss:spin} we concentrate on the spin probability distribution, and we show that the deviations from a lognormal shape are dependent on halo mass. We trace the origin of these deviations to the presence of \\emph{detectable} correlations among the quantities entering its definition, by applying the Kendall test. We discuss our findings in the Conclusion (section~\\ref{sec:concl}).\\\\ In the following, we will denote the natural logarithm using the symbol ``$\\ln$'', and the decimal logarithm as ``$\\log$'' (or ``Log''). ", "conclusions": "particularly at low values of $\\lambda$ the lognormal seems to provide a very good fit of our simulations. One possible reason of this discrepancy could be due to the lies in the fact that our simulations, compared to the \\textsc{Millennium} run analysed by Bett et al., do not have haloes more massive than $\\approx 10^{12} {\\rm M}_{\\sun}$, while in the \\textsc{Millennium} run this upper limit extends to ${\\rm M}\\approx 10^{15} {\\rm M}_{\\sun}$. If high mass haloes tend to have a smaller spin \\citep{2008ApJ...678..621K}, then we conclude that our simulations probe better the mass range typical of low- to intermediate mass galactic haloes, avoiding instead massively galactic and cluster-sized haloes. The latter tend to be not completely relaxed, thus their spin could deviate from the lognormal form expected from Tidal Torque Theory.}\\\\ We have chosen a relatively small cosmological volume ($L_{b} = 70\\, h^{-1}\\, {\\rm Mpc}$), in order to test whether the angular momentum--mass relation and spin distributions show \\emph{statistically detectable} deviations from the predicted analytical forms they have for dynamically relaxed haloes. The large statistics made possible by the high resolution allows us to detect some significant deviations: however, any numerical experiment is subject to some systematics.\\\\ \\noindent We cannot rule out that some residual systematic uncertainties could arise from the adopted halo finder: the actual spin of a numerical halo should also depend on the adopted halo finder, i.e. on the numerical definition. Even in the ''nearest neighbours'' tidal torque theory by \\citet{1969ApJ...155..393P} the distribution of angular momentum within the (spherically averaged) collapsed halo reflects the evolution of the torque during the collapse: the torques on shells placed at larger distances decrease on average with radius. Thus, one could expect that the distribution of $\\lambda$ would depend on the radius at which it is estimated, and this is generally different for halo finders based on critical isodensity contours or on friends-of-friends algorithms.\\\\ \\noindent {Concerning the environmental dependence of the spin distribution, it is now evident that the original claim by \\citet{1999MNRAS.302..111L}, who} did not find any evidence for this dependence, {was due to a lack of sufficent numerical resolution in their simulations. Already} in a previous work \\citep{2002MNRAS.332....7A} we found that the spin's PDF has a significantly smaller dispersion and mean in Voids, a result consistent with the more recent findings that a higher merger frequency tends to produce haloes with higher spin \\citep{2001ApJ...557..616G,2002MNRAS.329..423M,2006MNRAS.370.1905H}. \\\\ Our simulations have sufficient statistics to test the predictions of the linear model concerning the $J$--$M$ relationship. We notice that the index $\\alpha$ of the power law $J\\propto M^{\\alpha}$ at redshift zero is smaller than the standard value $5/3$. A similar fact had been previously noticed by \\citet[][see their Table 6]{2000MNRAS.311..762S}, using simulations on a larger volume, thus it is not probably a peculiarity of our use of a small volume.\\\\ \\noindent We measure average values of the spin at $z\\sim 0.1$ consistent with those found recently by other authors \\citep{2007MNRAS.378...55M}: $\\bar{\\lambda}\\simeq 0.03$. Macci\\`{o} et al. find little dependence of $\\bar{\\lambda}$ on halo mass, even when they restrict their analysis to relaxed haloes. Our results are consistent with their findings, and with theoretical analyses \\citep{2002ApJ...581..794C}. We do however see a dependence of the \\emph{spin dispersion} on mass: more massive haloes tend to have a lower dispersion, although it is a very slight trend which decreases towards low redshifts, consistently with the idea that it is induced by the presence of not fully relaxed haloes. \\\\ \\noindent The deviations of the spin's PDF from a lognormal arise because of statistical correlations between mass and angular momentum {and because of significant deviations of (at least) the total energy $E$ from a lognormal distribution}, as implied by Cram\\'{e}r's theorem. These correlations tend to disappear when haloes collapse, as one should expect for {ideal} completely relaxed haloes$^1$\\footnote{For instance, in the spherical collapse model the final overdensity only depends on the value of the matter overdensity $\\Omega_{m}$ at collapse redshift \\citep{1972ApJ...176....1G}}. Note that this conclusion is true both for merger- {and} collapse-induced relaxation: {in either case, the result is that the halo retains no \"memory\" of its previous state, due to dynamical friction acting on dynamical timescales.}" }, "1003/1003.3100_arXiv.txt": { "abstract": "{} {We present a diagnostic tool to determine the abundance of the crystalline silicate forsterite in AGB stars surrounded by a thick shell of silicate dust. Using six infrared spectra of high mass-loss oxygen rich AGB stars we obtain the forsterite abundance of their dust shells.} {We use a monte carlo radiative transfer code to calculate infrared spectra of dust enshrouded AGB stars. We vary the dust composition, mass-loss rate and outer radius. We focus on the strength of the 11.3 and the 33.6 $\\mu$m forsterite bands, that probe the most recent (11.3 $\\mu$m) and older (33.6 $\\mu$m) mass-loss history of the star. Simple diagnostic diagrams are derived, allowing direct comparison to observed band strengths.} {Our analysis shows that the 11.3 $\\mu$m forsterite band is a robust indicator for the forsterite abundance of the current mass-loss period for AGB stars with an optically thick dust shell. The 33.6 $\\mu$m band of forsterite is sensitive to changes in the density and the geometry of the emitting dust shell, and so a less robust indicator. Applying our method to six high mass-loss rate AGB stars shows that AGB stars can have forsterite abundances of 12\\% by mass and higher, which is more than the previously found maximum abundance of 5\\%.} {} ", "introduction": "Asymptotic Giant Branch (AGB) stars represent a late phase in the evolution of low- and intermediate mass stars (0.8-8.0\\,M$_{\\odot}$). At the end of the Red Giant Branch phase when the helium in the core is depleted a star enters the AGB phase, with a degenerate carbon and oxygen core surrounded by a helium and a hydrogen burning shell. The burning shells are surrounded by a convective envelope and a dynamically active atmosphere \\citep{HO03}. Depending on the carbon-to-oxygen ratio, either oxygen or carbon is locked in the CO molecule. This leads to two possible chemistries. Initially AGB stars exhibit an oxygen-rich chemistry reflecting the original composition of the natal molecular cloud. An AGB star becomes C-rich when sufficient carbon is dredged up from the helium-burning shell where it is formed by nucleosynthesis processes. In this study we will only consider O-rich AGB stars. One of the most important aspects of AGB stars is their stellar wind. AGB stars can have mass-loss rates between $10^{-8}$ \\nolinebreak M$_{\\odot}$ \\nolinebreak yr$^{-1}$ and about $10^{-4}$ M$_{\\odot}$ yr$^{-1}$. The high mass-loss rate automatically constrains the life-time of these objects, and gives rise to dense circumstellar envelopes. Two mechanisms contribute to this mass-loss. First matter gets accelerated outwards due to shocks \\citep{HO03}. When the material reaches distances where the temperature is low enough the material can condense into solid-state particles (from here on called dust). Then a second acceleration mechanism occurs: radiation pressure on the dust grains. The gas is coupled to the dust and dragged along. This leads to expansion speeds of 10 km s$^{-1}$ and higher. Even though amorphous silicate dust is more abundant, crystalline silicate dust is also seen in many astronomical environments like disks around pre-main-sequence stars \\citep{waelkens96, meeus01,spitzer1}, comets \\citep{wooden02}, post-main-sequence stars \\citep{waters96, syl99, mol02} and active galaxies \\citep{kemper07, spoon06}. Theoretical studies predict that the formation of crystalline silicate dust is dependent on the gas density \\citep{tielens98,gailsedl99, sogawa99}. Higher densities are more favorable for the formation of crystals than low densities. \\citet{speck08} also argue that there is a correlation between crystallinity, mass-loss rate and the initial mass of the star. These suggestions are consistent with the detection of crystalline silicates in the high density winds of OH/IR stars and the disks around post-AGB stars. Indeed, observations show a lack of crystalline silicate dust bands in the spectra of low mass-loss rate AGB stars \\citep{waters96, syl99}. However, an alternative explanation for this lack exists. \\citet{kemper01} have shown that a contrast effect hinders the detection of crystalline silicate bands in infrared spectra of low mass-loss rate AGB stars. This contrast effect is further discussed in sect. 3.2.2. \\citet{kemper01} obtained an indication for the crystalline silicate abundance in the outflow of AGB stars by comparing the strength of crystalline silicate emission bands of observations to those of model spectra. For the crystalline silicates enstatite and forsterite they measured a maximum abundance of $\\sim$5\\% by mass. AGB stars are the main contributors of crystalline silicate material to the ISM. The creation of crystalline silicates in the ISM is not possible since the temperatures are too low. Therefore the hypothesis would be that the abundance of crystalline silicate dust in the ISM would be comparable to that found in the dust shells of AGB stars. The crystalline silicate abundance in the diffuse ISM has been investigated by \\citet{kemper04, kemperErr05}. They looked at the line of sight towards the Galactic Centre and found an upper limit of the degree of crystallinity in the diffuse ISM of 2.2\\% by mass. The abundance found in AGB stars by \\citet{kemper01} is therefore higher than that for the diffuse ISM, suggesting the crystalline material may be effectively destroyed or hidden. Possible mechanisms could be amorphization by inclusion of iron atoms in the crystalline lattice, shocks and UV radiation (see \\citet{kemper04, kemperErr05} for a discussion on such mechanisms). Determining the crystalline fraction of silicates in AGB stars, as a function of mass-loss rate, will enhance our understanding in two areas: First, what are the conditions needed for the formation of crystalline silicates? Dust condensation theory suggests that the density is critical in the formation of crystalline silicates \\citep{tielens98,gailsedl99, sogawa99}, but this has still not been verified by observations. Secondly, how much crystalline silicate material is deposited by AGB stars to the diffuse ISM? Do the abundances in the outflow of AGB stars and the ISM match? The most thorough way to determine the forsterite abundance in the outflows of AGB stars is to fit the full spectral range offered by ISO or Spitzer spectra. This is an exhausting endeavor to apply to multiple stars and it does not necessary deliver a good understanding on how the spectral bands depend on dust shell parameters. Therefore it is useful to follow and expand the approach of \\citet{kemper01} who derived diagnostics for the 33.6 $\\mu$m forsterite band. We expand on this study by carefully studying the dependence of the spectral bands on the mass-loss rate, dust composition, thermal contact between grains and we also use two well-chosen bands, one probing the most recent mass-loss (11.3 $\\mu$m) and one probing the older, colder mass-loss (33.6 $\\mu$m). We derive a diagnostic diagram that is fairly easy to use and allows one to derive the forsterite abundance in the AGB star outflows. In sect. 2 we outline how the models were computed and show several examples of computed spectra. The strength determination of the two forsterite bands is discussed in sect. 3. Also the diagnostic diagram to measure the forsterite abundance is introduced and possible limitations are discussed. In sect. 4 we apply the diagnostic diagram to six ISO-SWS spectra of AGB stars. ", "conclusions": "We have shown that for optically thick dust shells the strength of the 11.3 $\\mu$m band is a robust quantity to use as an indicator of the forsterite abundance of the current mass-loss period. Applying the 33.6 $\\mu$m band is practically impossible without knowledge of the mass-loss rate, the size of the dust shell and other parameters that influence the emission of the cold dust in the outskirts of the dust shell of the AGB star. Additionally, for certain mass-loss rates (depending on the dust shell parameters) the 33.6 $\\mu$m band is also in transition from an emission to an absorption band, making the use of this band impossible. As an example and first result we used the 11.3 $\\mu$m band as an indicator of the forsterite abundance for six AGB stars. This showed that these objects can have low (below 2\\%) but also very high forsterite abundances (12\\% and possibly higher), meaning that AGB stars can have a forsterite abundance higher than 5\\% by mass, which was the maximum found by \\citet{kemper01}. It also means that the discrepancy between the amount of crystalline material in the ISM and that in AGB stars could even be higher. If the forsterite abundances found for the six sources in this study represent all AGB stars, the forsterite abundance in AGB stars alone would be higher than the crystalline abundance (upper limit of 2.2\\%) found by \\citet{kemperErr05} for the ISM. To find out if the abundances found for these stars are typical for AGB stars, more spectra have to be analyzed. The number of spectra are also not enough to test any correlations. But this paper opens the way to a more elaborate study similar to the one described in sect. 4 of this paper. Determining the amount of forsterite that is produced by high mass-loss rate AGB stars will help to test the suggested correlations and see if there is really a discrepancy between the produced abundance in AGB stars and the abundance found in the ISM." }, "1003/1003.4516_arXiv.txt": { "abstract": "We present the results of modeling dust spectral energy distributions (SEDs) across the Small Magellanic Cloud (SMC) with the aim of mapping the distribution of polycyclic aromatic hydrocarbons (PAHs) in a low-metallicity environment. Using \\emph{Spitzer} Survey of the SMC (S$^3$MC) photometry from 3.6 to 160 \\micron\\ over the main star-forming regions of the Wing and Bar of the SMC along with spectral mapping observations from 5 to 38 \\micron\\ from the \\emph{Spitzer} Spectroscopic Survey of the Small Magellanic Cloud (S$^4$MC) in selected regions, we model the dust spectral energy distribution and emission spectrum to determine the fraction of dust in PAHs across the SMC. We use the regions of overlaping photometry and spectroscopy to test the reliability of the PAH fraction as determined from SED fits alone. The PAH fraction in the SMC is low compared to the Milky Way and variable--with relatively high fractions (\\qpah$\\sim 1-2$\\%) in molecular clouds and low fractions in the diffuse ISM (average $\\langle$\\qpah$\\rangle = 0.6$\\%). We use the map of PAH fraction across the SMC to test a number of ideas regarding the production, destruction and processing of PAHs in the ISM. We find weak or no correlation between the PAH fraction and the distribution of carbon AGB stars, the location of supergiant H I shells and young supernova remnants, and the turbulent Mach number. We find that the PAH fraction is correlated with CO intensity, peaks in the dust surface density and the molecular gas surface density as determined from 160 \\micron\\ emission. The PAH fraction is high in regions of active star-formation, as predicted by its correlation with molecular gas, but is supressed in H II regions. Because the PAH fraction in the diffuse ISM is generally very low--in accordance with previous work on modeling the integrated SED of the SMC--and the PAH fraction is relatively high in molecular regions, we suggest that PAHs are destroyed in the diffuse ISM of the SMC and/or PAHs are forming in molecular clouds. We discuss the implications of these observations for our understanding of the PAH life cycle, particularly in low-metallicity and/or primordial galaxies. ", "introduction": "Polycyclic Aromatic Hydrocarbons (PAHs) are thought to be the carrier of the ubiquitously observed mid-IR emission bands \\citep[][among others]{allamandola89}. The bands are the result of vibrational de-excitation of the PAH skeleton through bending and stretching modes of C-H and C-C bonds after the absorption of a UV photon. The emission in these bands can be very bright and can comprise a significant fraction, up to 10$-$20\\% \\citep{smith07}, of the total infrared emission from a galaxy. For this reason, PAH emission has been suggested to be a useful tracer of the star-formation rate, even out to high redshifts \\citep{calzetti07}. Making use of PAHs as a tracer, however, requires understanding how the abundance and emission from PAHs depends on galaxy properties such as metallicity and star-formation history. PAHs also play a number of important roles in the interstellar medium (ISM). In particular, these small dust grains can dominate photoelectric heating rates \\citep{bakes94}. In dense clouds, PAHs can alter chemical reaction networks by providing a neutralization route for ionized species \\citep{bakes98,weingartner01} and contribute large amounts of surface area for chemical reactions that occur on grain surfaces. PAHs are a crucial component of interstellar dust so we would like to understand the processes that govern their abundance and physical state. The life-cycle of PAHs, however, is not yet well understood. PAHs are thought to form in the carbon-rich atmospheres of some evolved stars \\citep{latter91,cherchneff92}. Emission from PAHs has been observed from carbon-rich asymptotic giant branch stars \\citep{sloan07} and more frequently in carbon-rich post-AGB stars where the radiation field is more effective at exciting the mid-IR bands \\citep{buss93}. A ``stardust'' origin (i.e. formation in the atmospheres of evolved stars) for the majority of PAH material is controversial, however, because it has yet to be demonstrated that PAHs can be produced in AGB stars faster than they are destroyed in the ISM \\citep[for a recent review, see][]{draine09}, i.e. the timescale for destruction of dust by SNe shocks is shorter than the timescale over which the ISM is enriched with dust from AGB stars \\citep[for example,][]{jones94}. In addition to destruction by supernova shocks, PAH material may be destroyed by UV fields, a process that can dominate near a hot star or in the ISM of low metallicity and/or primordial galaxies. If PAHs are mostly not ``stardust'', they must have formed in the ISM itself, by some mechanism which is not yet characterized. A variety of mechanisms have been suggested \\citep{tielens87,puget89,herbst91,greenberg00} however there is little observational support for any one model as of yet. In recent years, observations with ISO and \\emph{Spitzer} have allowed us to study the abundance and physical state of PAHs in a variety of ISM conditions beyond those we observe in the Milky Way. One of the most striking results is the abrupt change in the fraction of dust in PAHs as a function of metallicity. A deficit of PAH emission from low metallicity galaxies has been widely observed \\citep{madden00,engelbracht05,madden06,wu06,jackson06,engelbracht08}. \\citet{engelbracht05} found that the ratio of the 8 to 24 \\micron\\ surface brightness undergoes a transition from a SED with typical PAH emission to an SED essentially devoid of PAH emission at a metallicity of 12 + log(O/H) $\\sim 8$. The weakness of PAH emission in low-metallicity galaxies has been confirmed spectroscopically \\citep{wu06,engelbracht08}. Using the SINGS galaxy sample \\citet{draine07b} modeled the integrated SEDs and determined that the deficit of PAH emission corresponds to a decrease in the PAH fraction rather than a change in excitation of the PAHs. They found that \\qpah\\ (defined as the fraction of the total dust mass that is contributed by PAHs containing less than $10^3$ carbon atoms) changes from a median of $\\sim 4$\\% \\citep[comparable to the Milky Way PAH fraction of 4.6\\%;][]{li01} in galaxies with 12 + log(O/H) $> 8.1$ to a median of $\\sim 1$\\% in more metal poor galaxies. \\citet{munoz-mateos09} have investigated the radial variation of \\qpah\\ and metallicity in the SINGS sample and find results consistent with \\citet{draine07b}. There have been a number of suggestions as to what in the PAH life-cycle changes at low metallicity leading to the observed deficiency. \\citet{galliano08} suggested that the delay between enrichment of the ISM by supernova-produced dust relative to that from AGB stars could lead to a lower PAH fraction at low metallicity. This model relies on the assumption that supernovae contribute a significant amount of dust to the ISM, an assumption which is controversial \\citep{moseley89,dunne03,krause04,sugerman06,meikle07,draine09}, as well as long timescales for dust production in carbon-rich AGB stars, which may be shorter than previously thought \\citep{sloan09}. Fundamentally, the \\citet{galliano08} model assumes a ``stardust'' origin for PAHs, which may not be the case. Other models explaining the low metallicity deficiency rely on enhanced destruction of PAHs. This can be accomplished through more efficient destruction via supernova shocks \\citep{ohalloran06} or via the harder and more intense UV fields in these galaxies \\citep[e.g.][]{madden06,gordon08}. In order to investigate the PAH life cycle at low metallicity, we performed two surveys of the Small Magellanic Cloud (SMC) with the \\emph{Spitzer} Space Telescope. The SMC is a nearby dwarf irregular galaxy that is currently interacting with the MW and the Large Magellanic Cloud. Its proximity \\citep[61 kpc;][]{hilditch05}, low metallicity \\citep[12 + log(O/H) $\\sim 8$, $Z \\sim 0.2 Z_{\\odot}$;][]{kurt98} and tidally disrupted ISM make it an ideal location in which to study the life cycle of PAHs in an environment very different from the Milky Way. The SMC has a low dust-to-gas ratio, $\\sim 10$ times smaller than in the Milky Way \\citep{bot04,leroy07}, leading to more pervasive UV fields. Because of its proximity we can observe the ISM at high spatial resolution and sensitivity in order to characterize the processes driving the PAH fraction. The PAH fraction in the SMC has been controversial. \\citet{li02}, using IRAS and COBE data, found that the PAHs in the SMC Bar contained only 0.4\\% of the interstellar carbon, corresponding to \\qpah $\\approx $ 0.2\\%, whereas \\citet{bot04} concluded that PAHs accounted for 4.8\\% of the total dust mass in the diffuse ISM of the SMC, similar to the Milky Way. In the following, we present results on the fraction of PAHs in the SMC using observations from the \\emph{Spitzer} Survey of the Small Magellanic Cloud (S$^3$MC). We use spectroscopy in the regions covered by the \\emph{Spitzer} Spectroscopic Survey of the Small Magellanic Cloud (S$^4$MC) to verify that our models for the photometry are correctly gauging the PAH fraction. We defer a detailed analysis of the spectroscopy to an upcoming paper (Sandstrom et al. 2010, in prep) In Section~\\ref{sec:obs} we describe the observations and data reduction, particularly focusing on the foreground subtraction and cross-calibration of the IRAC, MIPS and IRS observations. In Section~\\ref{sec:models} we describe the SED fitting procedure using the models of \\citet{draine07a} and the modifications necessary to incorporate the S$^4$MC spectroscopy into the fit. In Sections~\\ref{sec:results} and~\\ref{sec:discussion} we present the results of the SED modeling and discuss their implications for our understanding of the PAH life cycle both in low metallicity galaxies and in the Milky Way. ", "conclusions": "\\label{sec:conclusions} We present results of fitting the \\citet{draine07a} dust models to SEDs and spectra obtained from S$^3$MC and S$^4$MC. Our major results are as follows: \\begin{enumerate} \\item Comparisons of best-fit dust models in regions with overlapping photometry and spectroscopy demonstrate that the IRAC and MIPS SED estimator for \\qpah\\ does not appear to be appreciably biased despite the absence of information between $8-24$ and $24-70$ \\micron. When $5-38$ \\micron\\ IRS spectroscopy is added to the fitting constraints, the resulting ``photospectrofit'' models yield \\qpah\\ values that are only slightly larger than the ``photofit'' model estimates. The photospectrofit models have more dust at intermediate (T $\\approx 60$K) temperatures, and less dust with T $\\approx 200$ K. The estimate for the average radiation intensity scale factor $\\bar{U}$ and the dust surface density M$_{D}$ are nearly unaffected. \\item The PAH fraction in the SMC is low and variable. As a fraction of the total dust mass, the highest PAH fractions we observe are about half of the Milky Way value, but most of the galaxy has PAH fractions at the lower limit of our models (\\qpah= $0.4$\\%), an order of magnitude lower than the Milky Way value. The average $\\langle$\\qpah$\\rangle \\sim 0.6$\\% in the SMC. This is consistent with the earlier estimate of a very low PAH fraction in the SMC \\citep{li02} based on IRAS and DIRBE photometry. \\item The 8 \\micron\\ emission alone does not trace the \\qpah\\ well in the SMC. We find that the 8/24 ratio is correlated with \\qpah\\ and agrees with what one would predict from the metallicity of the SMC and the observed trends in \\citet{engelbracht05}, but for a given \\qpah\\ there is a wide range of 8/24 ratios, depending on the intensity of the local starlight heating the grains. \\item The metallicity of the SMC places it at the transition where lower metallicity galaxies show a deficiency of PAHs and higher metallicity galaxies show approximately Milky Way level PAH fractions. If the SMC is typical of galaxies in this region, the transition seems to represent a decrease in the filling factor of PAH-rich regions rather than a uniform decrease in the PAH fraction throughout the galaxy, although even the highest PAH fractions are still well below MW levels. \\item Lines-of-sight with molecular gas have an average \\qpah\\ of $\\sim 1$\\%, while lines of sight through the diffuse ISM have \\qpah\\ at least a factor of two less. \\item We evaluate the various proposed drivers of the deficit of PAHs at low metallicity. The distribution of PAHs in the SMC does not follow the carbon AGB star distribution, the regions of high turbulent Mach number, or the location of supergiant shells in the ISM. \\item The low PAH fraction in the diffuse ISM versus the high PAH fraction in molecular regions leads us to propose that PAHs may be forming in molecular clouds and/or a recent event (perhaps related to star formation events in the last $\\sim$ 25 Myr) destroyed a large fraction of the PAHs in the diffuse ISM. \\item We surmise that the global PAH fraction at low metallicities is a reflection of the amount of gas in these systems that is found at high extinction (A$_V > 1$). \\end{enumerate}" }, "1003/1003.5106_arXiv.txt": { "abstract": "{Some X-ray dim isolated neutron stars (XDINS) and central compact objects in supernova remnants (CCO) contain absorption features in their thermal soft X-ray spectra. It has been hypothesized that this absorption may relate to periodic peaks in free-free absorption opacities, caused by either Landau quantization of electron motion in magnetic fields $B\\lesssim10^{11}$~G or analogous quantization of ion motion in magnetic fields $B>10^{13}$~G. Here, I review the physics behind cyclotron quantum harmonics in free-free photoabsorption, discuss different approximations for their calculation, and explain why the ion cyclotron harmonics (beyond the fundamental) cannot be observed. } ", "introduction": "\\label{sect:intro} Thermal radiation from neutron stars can provide important information about their physical properties. Among neutron stars with thermal-like radiation spectra (see, e.g., reviews by \\citealt{kaspietal06} and \\citealt{zavlin09}), there are two classes of objects of particular interest: central compact objects (CCOs; see, e.g., \\citealt{deluca08}) in supernova remnants and X-ray dim isolated neutron stars (XDINSs, or the Magnificent Seven; see, e.g., review by \\citealt{turolla09}). The CCOs are young, radio-quiet isolated neutron stars with relatively weak magnetic fields $B \\sim(10^{10}-10^{11})$~G (e.g., \\citealt{halperngotthelf10} and references therein). The XDINSs are older and are believed to have much stronger fields $B \\gtrsim(10^{12}-10^{13})$~G \\citep{haberl07,vankerkwijkkaplan07,turolla09}. For some CCOs and XDINSs, there are estimates of $B$, and in some cases only upper limits to $B$ are available. In the past decade, broad absorption lines have been detected in the thermal spectra of several isolated neutron stars (see, e.g., \\citealt{vankerkwijk04,haberl07,vankerkwijkkaplan07}, and references therein). In all but one case, the energies $\\Ea$ of the absorption are centered on the range 0.2--0.7 keV and the effective black-body temperatures are $\\Teffbb\\approx0.1$ keV. Here and hereafter, the Boltzmann constant is suppressed, and the superscript ``$\\infty$'' indicates a redshifted value. In particular, it has been found that: (i) the spectrum of \\object{RX J1605.3+3249} with $\\Teffbb\\approx96$ eV has a broad absorption at $\\Ea\\approx0.4$--0.5 keV and a possible second absorption at 0.55 keV \\citep{vankerkwijketal04,vankerkwijk04}; (ii) \\object{RX J0720.4-3125} exhibits an absorption feature at $\\Ea\\approx0.27$ keV \\citep{haberletal04} and a possible second absorption at 0.57 keV \\citep{hambaryanetal09}, while the effective black-body temperature varies over years across the range $\\Teffbb\\approx86$--95 eV \\citep{hohleetal09}; (iii) the spectrum of \\object{RBS1223} (RX J$1308.6+2127$) was reproduced by a model with $\\Teffbb\\approx102\\pm2$ eV and two absorption lines at $\\Ea\\sim0.3$ keV and $\\Ea\\sim0.6$ keV \\citep{schwopeetal07}; and (iv) the spectrum of \\object{RBS1774} (1RXS J$214303.7+065419$) with $\\Teffbb\\approx102$~eV shows indications of a line at $\\Ea\\approx0.3$--0.4 keV and an absorption edge at $0.73$--0.75 keV \\citep{cropperetal07,kaplanvankerkwijk09,schwopeetal09}. The first discovered isolated neutron star with absorption lines, CCO \\object{1E 1207.4-5209}, has two absorption features centered on $\\Ea\\approx0.7$ keV and 1.4 keV \\citep{sanwaletal02} and an effective black-body temperature (which may be nonuniform) of $\\Teffbb\\sim0.16$--0.32 keV \\citep{zavlinetal98,delucaetal04}. For this object, two more harmonically spaced absorption features (at $\\Ea\\approx2.1$ keV and 2.8 keV) were tentatively detected \\citep{bignamietal03,delucaetal04}, but were later shown to be statistically insignificant \\citep{morietal05}. We note that realistic values of the effective temperature $\\Teff^\\infty$, obtained using atmosphere models, can differ from $\\Teffbb$ by a factor $\\lesssim2$--3 (see, e.g., \\citealt{zavlin09}, and references therein). Many authors \\citep[e.g.,][]{sanwaletal02,bignamietal03,delucaetal04} have considered the theoretical possibility that the absorption lines in the thermal spectra of the CCOs and XDINSs may be produced by cyclotron harmonics, formed because of quantum transitions between different Landau levels of charged particles in strong magnetic fields. \\citet{zaneetal01} discussed this possibility prior to the observational discovery of these absorption features. The fundamental cyclotron energy equals \\beq \\hbar\\omce=\\hbar {eB/\\mel c}=11.577~B_{12}\\textrm{ keV} \\label{omce} \\eeq for the electrons and $\\hbar\\omci=\\hbar{ZeB/ \\mion c}=6.35\\,(Z/A) B_{12}$~eV for the ions, where $\\mel$ and $\\mion$ are the electron and ion masses, respectively, $Z$ and $A$ are the ion charge and mass numbers, and $B_{12}\\equiv B/10^{12}$~G. In the following, we consider protons, whose cyclotron energy is \\beq \\hbar\\omcp=\\hbar eB/ \\mpr c = 6.305\\,B_{12} \\textrm{ eV}. \\label{omcp} \\eeq Beginning with the pioneering work of \\citet{GnedinSyunyaev74}, numerous papers have been devoted to the physics and modeling of cyclotron lines in X-ray spectra of accreting neutron stars \\citep[e.g.,][]{DV77,PSY80,WWL93,ArayaHarding,ArayaG-Harding,nishimura05,nishimura08}. These emission lines have been observed in many works following their discovery by \\citet{Truemperetal78}. Cyclotron harmonics have been found in spectra of several X-ray pulsars in binaries \\citep[e.g.,][and references therein]{rodes-rocaetal09,enotoetal08,pottschmidtetal04}, and up to four harmonics were registered for one of them \\citep{santangeloetal99}. In the photospheres of isolated neutron stars, unlike X-ray binaries, the typical energies of charged particles are nonrelativistic. In this case, first-order cyclotron transitions of free charged particles are dipole-allowed only between neighboring equidistant Landau levels and form a single cyclotron resonance with no harmonics. Special relativity and non-dipole corrections at the energies of interest can be estimated to be $\\mathrm{max}(\\Teff,E_\\mathrm{a})/\\mel c^2 \\sim10^{-3}$ for the electrons and $\\mathrm{max}(\\Teff,E_\\mathrm{a})/\\mpr c^2<10^{-6}$ for the protons. Beyond the first order in interactions, transitions between distant Landau states are also allowed in the nonrelativistic theory. They are, in particular, caused by Coulomb interactions between plasma particles. Thus cyclotron harmonics appear in free-free (bremsstrahlung) cross-sections. To obtain $E_\\mathrm{a}\\sim0.1$--1 keV, one may assume either the electron cyclotron harmonics at $B\\sim 10^{10}$--$10^{11}$~G, according to \\req{omce}, or proton cyclotron harmonics at $B\\sim 10^{13}$--$10^{14}$~G, according to \\req{omcp}. \\citet{pavlovshibanov78} presented the calculations of spectra for isolated neutron stars with prominent \\emph{electron} cyclotron harmonics due to the free-free absorption in the atmosphere. \\citet{SulPavWer} performed a similar atmosphere modeling and concluded, in agreement with \\citet{zaneetal01}, that electron cyclotron harmonics could be observed in CCO spectra. Proton cyclotron harmonics cannot be calculated based on the assumption of classical proton motion, used by these authors. In this paper, I review the physics of free-free photoabsorption in strong magnetic fields, discuss restrictions on different published approximations for free-free opacities, and present numerical results that demonstrate the relative strengths of the \\emph{electron and proton} cyclotron resonances under the conditions characteristic of the atmospheres of isolated neutron stars with strong magnetic fields. This gives a graphic explanation of the smallness of the ion cyclotron harmonics. I also demonstrate that the contribution of bound-bound and bound-free transitions to the opacities of neutron stars with $B>10^{13}$~G is much larger than that of the proton cyclotron harmonics. In Sect.\\,\\ref{sect:particle}, quantum mechanical integrals of motion and wave functions of a charged particle in a magnetic field are recalled for subsequent use. Section \\ref{sect:ep} is devoted to the properties of an electron-proton system in a magnetic field that is quantizing for both particles: general equations for calculation of wave functions are given, and the Born approximation is considered in detail. In the same order, general expressions and Born approximation are considered in Sect.\\,\\ref{sect:ff} for photoabsorption matrix elements and cross-sections. Section \\ref{sect:cycl} gives numerical examples of cyclotron harmonics in free-free photoabsorption with discussion and comparison of various approximations. Consequences for the CCOs and XDINSs are discussed in Sect.\\,\\ref{sect:discussion}, and Sect.\\,\\ref{sect:sum} presents our summary. ", "conclusions": "\\label{sect:discussion} \\subsection{Corrections beyond Born approximation} The formulae presented in Sects.\\,\\ref{sect:exact} and \\ref{sect:gen} in principle allow one to perform an accurate calculation of photoabsorption rates in the electron-proton system in an arbitrary magnetic field, taking into account the effects of Landau quantization of the electron and proton motion across the field and the transverse motion of the center of mass. For bound-free absorption, this calculation was presented by \\citet{PP97}. For free-free processes, we apply the first Born approximation and the dipole approximation. We plan to perform calculations of the free-free opacities beyond Born approximation in future work. An approximate estimate of the non-Born corrections can be obtained \\citep{PL07} by introducing correction factors $(1+\\gamma_\\ast^{-1} u^{-2})^{-1/2}(1+\\gamma_\\ast^{-1} (u')^{-2})^{-1/2}$ into the integral of \\req{Lambda0}, where $\\gamma_\\ast=(\\mel/\\mred)^2\\gamma$ and $ \\gamma = \\hbar^3 B/(\\mel^2 c e^3) = 425.44\\,B_{12}. $ The accuracy of the approximation is ensured by the smallness of $\\gamma_\\ast^{-1/2}\\approx0.05\\,B_{12}^{-1/2}$ and the additional condition $T\\gg e^4 m_\\ast/\\hbar^2$, which is the usual applicability condition for a Born approximation without a magnetic field. We have checked that these % corrections are sufficiently small for the electron cyclotron harmonics at $B\\sim10^{11}$~G (relevant to CCOs) and negligible for the proton cyclotron harmonics at $B>10^{13}$~G (relevant to XDINSs). \\subsection{Importance of bound states} \\begin{figure} \\includegraphics[width=\\columnwidth]{ang13_7c2.eps} \\caption{ Opacities for the two normal electromagnetic waves propagating at the angle $\\theta_B=45^{\\circ}$ to the magnetic field direction in a hydrogen atmosphere of a neutron star with $B=5\\times10^{13}$~G and $T=120$~eV at density $\\rho=10$~\\gcc\\ (which is in the middle of the photosphere at these $B$ and $T$). The results are shown for fully ionized (dotted lines) and partially ionized (solid and dot-dashed lines) plasma models. In the latter model, the nonionized atomic fraction equals 0.0066. The solid line shows the opacity obtained with the accurate calculation of the free-free Coulomb logarithm, and the dot-dashed line demonstrates the result of the approximate treatment that corresponds to the dashed line in Fig.~\\ref{fig:ql10mes2p}. } \\label{fig:ang13_7c2} \\end{figure} Free-free absorption contributes only a part of the total opacities in the atmospheres of neutron stars. A second constituent is the familiar scattering, and a third the absorption by bound species \\citep[see, e.g.,][]{CanutoVentura,Pavlov95}. It was realized long ago \\citep{Ruderman71} that in strong magnetic fields the increase in the binding energies of atoms and molecules can lead to their non-negligible abundance even in hot atmospheres. With increasing $B$, the binding energies and abundances of bound species increase at any fixed density $\\rho$ and temperature $T$ \\citep{PCS99,Lai01}, so that even the lightest of the atoms, hydrogen, provides a noticeable contribution to the opacities at the temperatures of interest, if the magnetic field is strong enough. Even a small neutral fraction can be important, because the bound-bound and bound-free cross-sections are large close to certain characteristic spectral energies. For electron cyclotron harmonics to appear at $\\hbar\\omega\\lesssim1$ keV, we should ensure that $B\\lesssim10^{11}$~G according to \\req{omce}. At these relatively weak magnetic fields and the characteristic temperature $T^\\infty\\gtrsim100$ eV, the assumption of full ionization may be acceptable. However, at $B>10^{13}$~G, which is required for ion cyclotron harmonics, the situation is different. An illustration is given in Fig.~\\ref{fig:ang13_7c2}. The solid curves show true absorption opacities for two normal electromagnetic waves propagating at the angle $45^\\circ$ to the magnetic field lines at $B=5\\times10^{13}$~G and $T=120$~eV. The upper and lower curves correspond to the ordinary and extraordinary waves, respectively. The density in this example is chosen to be $\\rho=10$~\\gcc, which is a typical atmosphere density at $B=5\\times10^{13}$~G and $\\Teff^\\infty=100$ eV (at this density the thermodynamic temperature $T$ approximately equals the effective temperature $\\Teff$). According to our ionization equilibrium model \\citep{PCS99}, at these $B$, $T$, and $\\rho$ values, 0.66\\% of protons in the plasma are comprised in the ground-state H atoms that are not too strongly perturbed by plasma microfields so that they contribute to the bound-bound and bound-free opacities (the ``optical'' atomic fraction), and only 0.1\\% of protons are in excited bound states. Even though the ground-state atomic fraction is small, it is not negligible. In Fig.~\\ref{fig:ang13_7c2}, at $\\hbar\\omega\\gtrsim0.4$ keV, the opacities in two normal modes, calculated with allowance for partial ionization (solid and dot-dashed curves), are significantly higher and have more characteristic features than the opacity calculated under the assumption of complete ionization (dotted lines). In particular, the broad feature on the lower curve near 0.4 keV is produced by the principal bound-bound transition between the two lowest bound states ($s_i=0\\to s_f=1$), and the increased value of the opacity at higher energies $\\hbar\\omega$ is due to the transitions to other bound and free quantum states. The wavy shape of the lower solid curve (for the extraordinary mode) at $\\hbar\\omega\\gtrsim0.7$ keV is explained by bound-free transitions to different open channels, each having its own threshold energy. All the bound-bound absorption features and photoionization thresholds are strongly broadened by the effects of atomic motion across the magnetic field lines (``magnetic broadening'', see \\citealt{PP97} and references therein). In the insets, we zoom in on the regions of the first and second proton cyclotron harmonics. Both of them are visible, but negligible compared to the effect of partial ionization on the opacities. \\subsection{Other possibilities for CCOs and XDINSs} Apart from the cyclotron harmonics, a number of alternative explanations of the observed absorption features in CCOs and XDINSs have been suggested in the literature. \\citet{moriho07} constructed models of strongly magnetized neutron star atmospheres with mid-$Z$ elements and compared them to the observed spectra of the neutron stars 1E 1207.4-5209 and RX J1605.3+3249. They demonstrated that the positions and relative strengths of the strongest absorption features in these neutron stars are in good agreement with a model of a strongly ionized oxygen atmosphere with $B=10^{12}$~G and $B=10^{13}$~G, respectively. This explanation seems promising, but unsolved problems remain: the effects of motion across the field have been treated approximately, based on the assumption that they are small, and detailed fits to the observed spectra have not yet been presented. Among other hypotheses about the nature of the absorption features, there was a suggestion that they could be due to bound-bound transitions in exotic molecular ions \\citep{TurbinerLopez06}. However, our estimates show that the abundance of these ions in a neutron star atmosphere would be negligible compared with the abundance of H atoms. \\citet{SPW09} proposed a ``sandwich'' model atmosphere of finite depth, composed of a helium slab above a condensed surface and beneath hydrogen, and demonstrated that this model can produce two or three absorption features in the range of $\\hbar\\omega\\sim0.2$--1 keV at $B\\sim10^{14}$~G, although a detailed comparison with observed spectra was not performed. One cannot also rule out that some absorption lines originate in a cloud near a neutron star, rather than in the atmosphere \\citep[see][]{hambaryanetal09}." }, "1003/1003.0515_arXiv.txt": { "abstract": "{} {To determine the Point Source Location Accuracy (PSLA) for the INTEGRAL/IBIS telescope based on analysis of archival in-flight data.} {Over 40000 individual pointings (science windows) of INTEGRAL/IBIS data were analysed using the latest Off-line Science Analysis software, version 7.0. Reconstructed source positions were then compared against the most accurate positions available, determined from focusing X-ray telescopes. Since the PSLA is a strong function of source detection significance, the offsets from true position were histogrammed against significance, so that the 90\\% confidence limits could be determined. This has been done for both sources in the fully coded field of view (FCFOV) and partially coded field of view (PCFOV).} {The PSLA is found to have improved significantly since values derived from early mission data and software for both FCFOV and PCFOV.} {This result has implications for observers executing follow-up programs on IBIS sources since the sky area to be searched is reduced by $>$50\\% in some cases.} ", "introduction": "The \\textit{INTEGRAL} satellite (\\citealt{winkler}), launched in October 2002, is an ESA space mission specifically designed to study the gamma-ray sky. In particular the IBIS (Imager on Board of the \\textit{INTEGRAL} Satellite) telescope (\\citealt{ubertini03}) is the main hard X-ray/soft gamma-ray coded aperture imaging instrument (\\citealt{goldwurm03}), and is responsible for surveying and cataloguing the sky above 17 keV. Here we discuss the point source location accuracy (PSLA) of the ISGRI low energy detector (15-1000 keV) of IBIS (\\citealt{lebrun03}). Due to the continuing necessity to follow-up the growing unidentified \\textit{INTEGRAL} source population in other wavebands (particularly optical and infrared), it is of great interest to assure the correctness of the IBIS/ISGRI PSLA. In particular, it is hoped that with the release of new and updated Off-line Science Analysis (OSA 7.0) software the PSLA could have improved substantially compared to the already published estimates based on early mission data and software releases \\citep[][constructed the PSLA based upon OSA 3.0]{gros}. Any improvement in the PSLA is important in reducing the chance of random or multiple source associations in other wavebands. ", "conclusions": "This work has shown that the IBIS/ISGRI PSLA (ie the size of the 90\\% error circle) has dramatically improved since the last study of \\cite{gros} and this can be attributed to improvements in the {\\em OSA} software. The main implications of this result will be for follow-up observations, particularly in the optical and infrared wavebands, of the \\textit{INTEGRAL} unidentified source population ($\\approx30\\%$ of objects present in the latest IBIS/ISGRI catalogue release by \\citealt{cat4}). This improved error radius will greatly enhance the identification of \\textit{INTEGRAL} counterparts." }, "1003/1003.2984_arXiv.txt": { "abstract": "We have studied the relationship between the nuclear (high-resolution) radio emission, at 8.4~GHz (3.6~cm) and 1.4~GHz (20~cm), the [O~IV]~$\\lambda$25.89~$\\micron$, [Ne~III]~$\\lambda$15.56~$\\micron$ and [Ne~II]~$\\lambda$12.81~$\\micron$ emission lines and the black hole mass accretion rate for a sample of Seyfert galaxies. In order to characterize the radio contribution for the Seyfert nuclei we used the 8.4GHz/[O~IV] ratio, assuming that [O~IV] scales with the luminosity of the AGN. From this we find that Seyfert 1's (i.e., Seyfert 1.0's, 1.2's, and 1.5's) and Seyfert 2's (i.e., Seyfert 1.8's, 1.9's, and 2.0's) have similar radio contributions, relative to the AGN. On the other hand, sources in which the [Ne~II] emission is dominated either by the AGN or star formation have statistically different radio contributions, with star formation dominated sources more ``radio loud\", by a factor of $\\sim 2.8$ on average, than AGN dominated sources. We show that star formation dominated sources with relatively larger radio contribution have smaller mass accretion rates. Overall, we suggest that 8.4GHz/[O~IV], or alternatively, 1.4GHz/[O~IV] ratios, can be used to characterize the radio contribution, relative to the AGN, without the limitation of previous methods that rely on optical observables. ", "introduction": "Active galactic nuclei (AGN) are among the most powerful objects in the universe with an extraordinary amount of energy released from their unresolved nuclei. Active super-massive black holes (SMBH) in the center of these galaxies are responsible for the tremendous diversity in their observed energy emission \\citep[e.g.,][]{1984ARA&A..22..471R,1989ApJ...347...29S,2000ApJ...540L..13P,2004ApJ...613..682P}. Despite the fact that the unified scheme of AGN has been able to explain the observed type I/II dichotomy \\citep[see][for a review]{1993ARA&A..31..473A}, it is clear that by looking at different wave bands there are additional parameters that can complement our current understanding of AGN classification other than to be attributed solely to a viewing angle difference \\citep[e.g.,][]{1997ApJ...485..552M,2001ApJ...546..845G,2006AJ....132..401B,2008ApJ...689...95M}. Therefore, in order to unveil and understand the complexity of these sources, ones needs to rely on multiwavelength studies. In particular, the radio and mid-infrared emissions are basically unaffected by dust extinction which make these energy bands ideal to investigate the nature of the ``hidden\" nuclear engine. On the other hand, extended radio emission and low-ionization mid-infrared emission lines are excellent tracers of star formation activity when properly calibrated \\citep[e.g.,][]{1992ARA&A..30..575C,2007ApJ...658..314H}. Furthermore, the strength of the observed radio emission in AGN seems to follow a bimodal distribution falling into ``radio loud\" and ``radio quiet\" sources \\citep[e.g.,][]{1990MNRAS.244..207M,1999AJ....118.1169X}. Historically, the dividing criterion for the AGN radio loudness was adopted to be the ratio between the radio emission, at the most commonly used frequency of 5~GHz, and the optical continuum at $\\sim 4400$~\\AA, i.e., the radio-to-optical parameter, ${\\rm R=L_\\nu{\\rm (6{\\rm cm})}/L_\\nu({\\rm B})}$ \\citep[e.g.,][]{1989AJ.....98.1195K,1992ApJ...391..560V}. This convention to quantize the radio loudness of a source is straightforward when the contribution from the host galaxy does not affect the observed radio and optical components of the AGN, for example in quasars, where the active nucleus overpowers the underlying galaxy. In the case of Seyfert galaxies the quantification of their radio loudness could be more subtle. \\citet{2001ApJ...555..650H} found that most Seyfert 1 galaxies, previously classified as radio-quiet, fall into the radio loud regime when the nuclear component for the radio-to-optical parameter is properly measured. However, even with a resolution of few arcseconds, the nuclear component of the optical luminosity could be contaminated by star formation and is strongly affected by dust extinction that may lead to an erroneous AGN characterization. Previous multiwavelength studies have shown statistically significant correlations between the radio core luminosity and purportedly orientation-independent measures of the intrinsic luminosity of the active nuclei, which suggests a common mechanism at nuclear (subkiloparsec) scales, , e.g., accretion onto the black hole \\citep[e.g.,][]{1989MNRAS.240..701R,1999AJ....118.1169X,2002ApJ...564..120H,2009ApJ...698..623D}. On the other hand, the correlations between extra-nuclear radio emission and the far-infrared (FIR) luminosity of the galaxy suggest a circumnuclear starburst as the common cause for these observables \\citep[e.g.,][]{1982ApJ...252..102C,1988AJ.....96...30C,1993ApJ...419..553B,1998MNRAS.301.1019R}. In this regard, early studies suggested two possible scenarios for the origin of the compact radio fluxes observed in radio quiet sources: either the radio emission is associated with a nuclear starburst or generated by relativistic electrons accelerated by a low-power jet (similar to radio loud sources), with high-resolution radio-imaging observations favoring the latter \\citep[e.g.,][]{1998MNRAS.299..165B}. More recently, the good correlation between the X-ray and the radio luminosities raise the possibility that radio emission in radio quiet sources may have a coronal origin \\citep{2008MNRAS.390..847L}. All these exemplify that the radio properties of the AGN are diverse and can be intimately associated with fundamental parameters of the active galaxies, such as the black hole masses, mass accretion rates, and host galaxy morphology \\citep[e.g.,][]{1999AJ....118.1169X,2000ApJ...543L.111L,2001A&A...380...31W}. Recently, \\citet{2008ApJ...682...94M} found a tight correlation in Seyfert 1 galaxies between the [O~IV]~$\\lambda$25.89~$\\micron$ and the X-ray 14-195~keV continuum luminosities from {\\it Swift}/BAT observations \\citep{2005ApJ...633L..77M,2008ApJ...681..113T}. A weaker correlation was found in Seyfert 2 galaxies, which was attributed to the effect of Compton scattering in the 14-195~keV band. These results suggest that [O~IV] is a truly isotropic property of AGNs, given its high ionization potential and that it is basically unaffected by reddening, meaning that the [O~IV] strength directly measures the AGN power. This result has been later confirmed in more complete samples \\citep[][Weaver et al. submitted to ApJ]{2009ApJ...700.1878R}. Moreover, the strong correlation found between [O~IV] and the high-ionization ($\\sim$97~eV) [Ne~V]~$\\lambda$14.32/24.32~$\\micron$ emission is a further argument for [O~IV] being a good indicator for the intrinsic luminosity of the AGN \\citep[e.g.,][Weaver et al. submitted to ApJ]{2009ApJ...691.1501D,2009MNRAS.398.1165G}. It is the primary goal of this paper to examine the physical relationship between the nuclear radio emission and the mid-infrared emission line luminosities in Seyfert galaxies. We use high-resolution radio continuum measurements and the [O~IV] emission line to investigate the radio loudness of the Seyfert nuclei and the black hole mass accretion rate for a sample of Seyfert galaxies. ", "conclusions": "We have investigated the relationship between the radio emission and the ionization state of the emission-line gas in Seyfert galaxies with the aim to understand the connection between the radio loudness and the ionizing luminosities of the AGN. We used the [O~IV] emission line to estimate the intrinsic power of the AGN and high-resolution 8.4~GHz emission to characterize the radio core power. From this we defined a radio-to-mid-infrared parameter, 8.4GHz/[O~IV], to identify the radio contribution from the Seyfert nuclei. We found that Seyfert 1 and Seyfert 2 galaxies are statistically similar in their radio emission, relative to the strength of the AGN. On the other hand, when separated by the dominant source of [Ne~II] emission, we found that two groups, AGN and star formation dominated sources, are statistically different in their radio loudness. We found that sources with strong star formation, relative to the strength of the AGN, or alternatively weak AGNs, tend to be more radio loud, by a factor of $\\sim 2.8$, than AGN dominated sources. Furthermore, in the 8.4GHz-[O~IV] correlation, sources in the star formation dominated region have lower mass accretion rates than the rest of the sample. This result is in agreement with the anticorrelation found by \\citet{2002ApJ...564..120H} between the radio loudness and Eddington accretion rates. Given the sub-Eddington nature of the more radio loud sources in our sample, sources in the star formation dominated region, one can invoke advection-dominated accretion models to explain their bright compact radio emission and a low-luminous nucleus both consistent with their high radio-to-mid-infrared parameter. We also found that high resolution 1.4~GHz emission could also be used as a reasonable proxy for the AGN radio contribution. When using the 1.4GHz/[O~IV] ratios we found that sources in which the [Ne~II] emission is dominated by the stellar activity, or alternatively weak AGNs, tend to be more radio loud, by a factor of $\\sim 3.0$, than AGN dominated sources, in close agreement with the radio loudness characterization using the radio core power at 8.4GHz. When comparing the observed 8.4~GHz and 1.4~GHz high-resolution core emission in our sample of Seyfert galaxies we found a strong linear correlation, suggesting that the core radio emission is dominated by the active nuclei with a spectral index of $f_\\nu \\propto \\nu^{-0.7}$. We found the spectral index distribution, between 1.4~GHz and 8.4~GHz, for Seyfert 1 and Seyfert 2 galaxies and for AGN and star formation dominated sources to be statistically similar within the groups. We suggest that the radio-to-optical parameter, $R$, used to classify galaxy radio loudness may erroneously quantize the radio contribution in galaxies with strong star formation or alternatively, weak AGNs, in which the optical luminosity may be contaminated by star formation and/or suffer from dust extinction. Instead, we found that a radio-to-mid-infrared parameter, 8.4GHz/[O~IV], is a better probe for the radio loudness of the galaxy, as it is basically unaffected by star formation contamination and is less affected by extinction than optically dependent diagnostics." }, "1003/1003.0453_arXiv.txt": { "abstract": "Gravitational hydrodynamics acknowledges that hydrodynamics is essentially nonlinear and viscous. In the plasma, at $z=5100$, the viscous length enters the horizon and causes fragmentation into plasma clumps surrounded by voids. The latter have expanded to $38$ Mpc now, explaining the cosmic void scale $30/h=42$ Mpc. After the decoupling the Jeans mechanism fragments all matter in clumps of ca 40,000 solar masses. Each of them fragments due to viscosity in millibrown dwarfs of earth weight, so each Jeans cluster contains billions of them. The Jeans clusters act as ideal gas particles in the isothermal model, explaining the flattening of rotation curves. The first stars in old globular clusters are formed by aggregation of milli brown dwarfs, without dark period. Star formation also happens when Jean clusters come close to each other and agitate and heat up the cooled milli brown dwarfs, which then expand and coalesce to form new stars. This explains the Tully-Fischer and Jackson-Faber relations, and the formation of young globular clusters in galaxy mergers. Thousand of milli brown dwarfs have been observed in quasar microlensing and some 40,000 in the Helix planetary nebula. While the milli brown dwarfs, i.e., dark baryons, constitute the galactic dark matter, cluster dark matter consists probably of 1.5 eV neutrinos, free streaming at the decoupling. These two types of dark matter explain a wealth of observations. ", "introduction": "It is generally understood that hydrodynamics is in principle important in the early Universe, but its role is small in practice. There is a great impetus from quantum field theory, inflation of a mysterious scalar field explains the structure in the CMB. In this approach, hydrodynamics is linearized before the decoupling of photons and matter. It was Gibson 1996 who questioned this assumption~\\cite{Gibson96}. The conclusion is astonishing: {\\it There is a viscous instability in the plasma, so the concordance approach cannot be correct. Hydrodynamics can explain structure formation without cold dark matter trigger.} Here we summarize our recent work, Ref.~\\citen{NGS09}. ", "conclusions": "" }, "1003/1003.0665_arXiv.txt": { "abstract": "Observational evidence is presented for the merging of a downward-propagating plasmoid with a looptop kernel during an occulted limb event on 2007 January 25. RHESSI lightcurves in the 9--18~keV energy range, as well as that of the 245 MHz channel of the Learmonth Solar Observatory, show enhanced nonthermal emission in the corona at the time of the merging suggesting that additional particle acceleration took place. This was attributed to a secondary episode of reconnection in the current sheet that formed between the two merging sources. RHESSI images were used to establish a mean downward velocity of the plasmoid of 12~km~s$^{-1}$. Complementary observations from the SECCHI suite of instruments onboard STEREO-Behind showed that this process occurred during the acceleration phase of the associated CME. From wavelet-enhanced EUVI, images evidence of inflowing magnetic field lines prior to the CME eruption is also presented. The derived inflow velocity was found to be 1.5~km~s$^{-1}$. This combination of observations supports a recent numerical simulation of plasmoid formation, propagation and subsequent particle acceleration due to the tearing mode instability during current sheet formation. ", "introduction": "\\label{intro} During an eruptive event comprising a solar flare and a coronal mass ejection (CME), energy is believed to be converted into the heating of plasma and the kinetic energy of particles and the CME itself through the process of magnetic reconnection. The standard reconnection models (\\citealt{park57,swee58,pets64}) state that newly connected field lines expel plasma from the reconnection site due to the Lorentz force. The pressure gradient across the diffusion region then forces new plasma inwards, along with the field lines frozen to it where they change connectivity and dissipate energy. \\citet{lin00} stated that these inflows are concurrent with the eruption of a CME, which remains connected to the magnetic neutral point by an extended current sheet. Initially the CME rises slowly until no neighboring equilibrium state is available. After reaching this point the CME begins to rise at an increasing rate. Energy release and particle acceleration continue due to sustained reconnection as the CME accelerates. To date, there has been little observational evidence for the predicted inflows associated with reconnection. The most cited example is that of \\cite{yoko01}, who found inflow velocities of 1.0--4.7~km~s$^{-1}$ by tracing the movement of threadlike patterns above the limb in a series of {\\it SOHO}/EIT 195\\AA~images (\\ion{Fe}{12}; \\citealt{dela95}). Evidence for sustained energy release during CME acceleration has been reported in a recent study of two fast ($>$1000~km~s$^{-1}$) halo CMEs by \\cite{temm08,temm10}, who found a strong correlation between the CME acceleration and flare hard X-ray (HXR) time profiles. The bremsstrahlung hard X-rays are signatures of thick-target collisions between the accelerated electrons and the ambient chromospheric material. The authors interpret this correlation as strong evidence for a feedback relationship occurring between CME acceleration and the energy released by magnetic reconnection in the current sheet formed behind the CME. In cases where the current sheet becomes sufficiently longer than its width, it is possible for multiple X-points to form due to the tearing mode instability which can result in the formation of plasmoids \\citep{furt63}. Both upward- \\citep{ko03,lin05} and downward-directed \\citep{shee02} plasmoids associated with CME eruption have been commonly observed in white light coronagraph images, in agreement with MHD simulations (e.g., \\citealt{forb83}). Further evidence for plasmoid motions has been presented through radio observations of drifting pulsating structures \\citep[DPS;][]{klie00,karl04,rile07,bart08b}. \\begin{figure}[!t] \\begin{center} \\includegraphics[height=8.5cm,angle=90]{f1.eps} \\caption{The CME on 2007 January 25 as seen by the COR1 coronagraph (blue) at 06:53:24~UT as well as the associated EUVI field of view (green). The expanded box shows the coronal loops that form part of active region NOAA 10940 with 40 and 80\\% contours of the 5-10~keV emission seen by RHESSI at the same time (solid line).} \\label{euvi_cor1} \\end{center} \\end{figure} Observational X-ray evidence for the formation of a current sheet has been presented by \\cite{sui03} using data from the Ramaty High Energy Solar Spectroscopic Imager (RHESSI; \\citealt{lin02}). The authors were able to show that an above-the-looptop coronal X-ray source (or plasmoid) increased in altitude as a lower lying X-ray loop decreased in altitude during the initial stages of a solar flare. They concluded that magnetic reconnection occurred between the two sources as the current sheet formed. This interpretation was strengthened by evidence that the mean photon energy decreased with distance in both directions away from the reconnection site (see also \\citealt{liu08}). The authors attribute the downward moving looptop to the collapse of the X-point to a relaxed magnetic loop during the reconfiguration of the magnetic field. The same conclusions were reached by \\cite{sui04} and \\cite{vero06}, who observed similar motions of rising plasmoids concurrent with shrinking looptop sources in other events imaged with RHESSI. A recent numerical simulation by \\cite{bart08a} shows that by invoking variable reconnection rates along the current sheet, {\\it downward} propagating plasmoids should also be visible in X-rays below $\\sim$2~R$_{\\odot}$ (see also \\citealt{rile07}). The condition for this scenario is met when the reconnection rate above the plasmoid is greater than that below resulting in a net downward tension in the newly connected magnetic field lines. Furthermore, this model shows that the interaction of such a plasmoid with the underlying loop system can result in a substantial increase in dissipated energy, more so than during the initial ejection of the rising plasmoid or coalescing plasmoid pairs. To date, there has only been one report of such an interaction by \\cite{kolo07} using Yohkoh/SXT data. They found an increase in X-ray and decimetric radio flux and an increase in temperature at the interaction site. \\begin{figure}[!t] \\begin{center} \\includegraphics[height=8.5cm,angle=90]{f2.eps} \\caption{Lightcurves of the flare in the 3--6, 6--12, and 12--25~keV energy bands as observed by RHESSI, as well as the GOES 1--8~\\AA~lightcurve. The horizontal bars at the top of the plot denote RHESSI's attenuator state (A0, A1), nighttime (N) and SAA passes (S).} \\label{hsi_goes_ltc} \\end{center} \\end{figure} \\begin{figure*}[] \\begin{center} \\includegraphics[height=\\textwidth,angle=90]{f3.eps} \\caption{RHESSI images in the 5-10 keV energy band formed over 60s integrations during the onset of the flare, although only alternate images are shown here. Contours mark the 40\\% and 80\\% levels. The plasmoid (source A) and looptop (source B) sources are labeled. The grey pattern around the sources are the CLEAN residuals and reflect the background noise level of the images.} \\label{multi_hsi_plot} \\end{center} \\end{figure*} \\begin{figure}[!b] \\begin{center} \\includegraphics[height=8.5cm]{f4.eps} \\caption{The two sources observed by RHESSI imaged over 2~keV wide energy bins (3--5, 5--7, 7--9~keV) for a single time interval.} \\label{hsi_ht_vs_en} \\end{center} \\end{figure} In this paper observational evidence is presented for increased hard X-ray and radio emission during the coalescence of a downward-moving coronal source with a looptop kernel at the onset of a flare observed with RHESSI. Coordinated observations from the Sun-Earth Connection Coronal and Heliospheric Investigation (SECCHI; \\citealt{howa08}) suite of instruments onboard the Solar Terrestrial Earth Relations Observatory (STEREO; \\citealt{kais08}) show that this interaction was concurrent with the acceleration phase of the associated CME. Using wavelet enhanced images from the EUV Imager (EUVI), evidence is also presented for inflowing magnetic field lines that persisted for several hours prior to reconnection. Section~\\ref{obs} describes the event as viewed by RHESSI and STEREO and the techniques used to determine the motion of the coronal X-ray sources and the CME. Section~\\ref{conc} discusses the findings in the context of numerical simulations, and summarizes the conclusions. ", "conclusions": "\\label{conc} Rare observations are presented of a downward-propagating X-ray plasmoid appearing to merge with a looptop kernel during an eruptive event seen above the solar limb; the first case observed with RHESSI and perhaps only the second ever. Although the majority of above-the-looptop sources observed (in both white light and X-rays) tend to rise due to weaker magnetic field and decreasing density above the flare loops, in certain instances, conditions can be right for downward-moving plasmoids to form also. Enhanced HXR emission detected with RHESSI and radio emission observed by the Learmonth radio telescope suggest that this merging resulted in a secondary episode of particle acceleration (see Figure~\\ref{hsi_ltc_ht_radio}). Images of the plasmoid formed over finer energy bins (as shown in Figure~\\ref{hsi_ht_vs_en}) show that higher energy emission was observed at higher altitudes. This is consistent with the idea that the reconnection rate above the source was greater than that below, unlike rising plasmoids previously observed with RHESSI which show mean photon energy decreasing with height (e.g. \\citealt{sui03}). Complementary observations from STEREO show that the plasmoid-looptop merging was concurrent with the period of the most rapid acceleration of the associated CME (Figures~\\ref{hsi_cme_ht}$c$ and $d$). These observations are in agreement with a recent numerical simulation that predicts an increase in liberated energy during the merging of a plasmoid with a looptop source \\citep{bart08a}. The formation of plasmoids is attributed to the tearing-mode instability during current sheet formation in the wake of an erupting CME \\citep{furt63,lin00}. \\begin{figure}[!t] \\begin{center} \\includegraphics[width=8.5cm]{f8.eps} \\caption{Summary of the kinematics of both the CME observed with STEREO and the coronal X-ray sources observed with RHESSI. {\\it a}) Height-time plot of the CME front from EUVI (diamonds), COR1 (triangles), and COR2 (squares). {\\it b}) and {\\it c}) The associated velocity and acceleration profiles, respectively. {\\it d}) Height-time plot of the 5--10~keV sources as observed by RHESSI. The downward-moving coronal source is shown as crosses with error bars. The solid line denoted a least-squares fit to the data points and has been extended beyond the data points for clarity. The rising looptop source is represented by diamonds also with error bars. {\\it e}) Observing summary profiles for RHESSI in the 3--6, 6--12 and 12--25~keV energy bands. Horizontal bars marked N and S denote RHESSI nighttimes and SAA passes, respectively.} \\label{hsi_cme_ht} \\end{center} \\end{figure} \\begin{figure}[!t] \\begin{center} \\includegraphics[height=8.5cm,angle=90]{f9.eps} \\caption{{\\it Left}: A wavelet enhanced EUVI image taken at 06:51:47~UT. The solid contours overlaid mark the position of the HXR looptop source at the time of the image. The dashed contour marks the position of the plasmoid at 06:29~UT. The dotted line shows the location of the CME leg at 04:46~UT and the arrow points in the direction of its motion. The solid vertical line denotes the pixel column used for the time-slice plot on the right. {\\it Right}: A temporal evolution of a vertical slice through a series of EUVI images. The dotted line marks the time of the image in the left-hand panel. The structure believed to be the inflowing CME leg is identified between the two parallel lines.} \\label{euvi_time_slice} \\end{center} \\end{figure} \\cite{bart07,bart08a} have shown theoretically that the deceleration of the plasmoid as it collides with the looptop source can lead to significant episodes of energy release. During this deceleration, antiparallel magnetic field lines begin to pile up between the two sources and a secondary current sheet is formed. This in turn leads to a secondary episode of magnetic reconnection that is driven by the magnetic tension of the field lines that govern the plasmoid motion. The authors also claim that the merging process triggers the excitation of large amplitude waves which can carry with them some of the stored magnetic energy. Although it is not possible to detect any acceleration or deceleration from the RHESSI images presented, a mean downward velocity of 12~km~s$^{-1}$ was calculated. This value is commensurate with the previous observation of \\cite{kolo07}, who measured 16~km~s$^{-1}$ during a similar event observed with Yohkoh. However, both these observed values are considerably lower than the value predicted by \\cite{bart08a} of $\\sim$40\\% of the local Alfv\\'{e}n speed (i.e. $\\sim$400~km~s$^{-1}$). Similar values of $\\sim$200~km~s$^{-1}$ were predicted by \\cite{rile07} for downward-moving white-light plasmoids. The low velocity measured here may be attributed to the low value of the reconnection rate as estimated from the inflows observed with EUVI (assuming that these field lines converged above the plasmoid). The value of $M_A \\approx$ 0.001 is an order of magnitude lower than that used in the numerical simulation. As the amount of tension exerted on the plasmoid is sensitive to the net reconnection rate, this would result in a lower tension force and therefore lower downward velocity. This in turn may also affect the amount of energy liberated in the subsequent collision with the looptop. It is possible that the model of \\cite{bart08a} may overestimate the velocity (and subsequent dissipated energy) given that the simulation is two-dimensional and does not take into account 3D structures, such as a twisted flux rope. Similarly the plasmoid detected with RHESSI is observed for more than 10 minutes before merging with the looptop source, whereas the simulations which yielded higher velocities predict that the source should exist for only $\\sim$1 minute before merging. While the simulations of \\cite{bart08a} predict a rebrightening of the loop footpoints in HXRs and/or chromospheric emission, both the analysis presented here and that of \\citealt{kolo07} show a distinct increase in coronal emission. A recent analysis of Miklenic et al. (2010; submitted) appears to refute the idea that plasmoid-looptop interactions could be responsible for chromospheric rebrightenings. These observations provide further evidence that the particle acceleration process occurs in the corona rather than at the footpoints as recently suggested by \\cite{flet08}, although acceleration at the footpoints as recently suggested by \\cite{brow09} cannot be ruled out. While plasmoid-looptop interactions are rarely observed, it is possible that they occur more often but are difficult to observe due to the brighter emission from the flare itself and RHESSI's limited dynamic range. A newly developed technique of deep integrations using RHESSI visibility-based X-ray imaging \\citep{sain09} may help to identify faint X-ray sources in the corona during eruptive limb events. By comparing other similar events it may be possible to determine how great an effect the CME acceleration (and magnetic reconnection rate, if possible) has upon the resulting HXR and radio flux. It would therefore be useful to find events observed at a time when RHESSI's detector calibration was better known in order to perform a more rigorous spectral analysis which was not possible for this event due to poorly known calibration. This could reveal more detailed information on the energetics of the resulting accelerated particles." }, "1003/1003.0386_arXiv.txt": { "abstract": "Mercury is the target of two space missions: MESSENGER (NASA) which orbit insertion is planned for March 2011, and ESA/JAXA BepiColombo, that should be launched in 2014. Their instruments will observe the surface of the planet with a high accuracy (about 1 arcsec for BepiColombo), what motivates studying its rotation. Mercury is assumed to be composed of a rigid mantle and an at least partially molten core. We here study the influence of the core-mantle interactions on the rotation perturbed by the solar gravitational interaction, by modeling the core as an ellipsoidal cavity filled with inviscid fluid of constant uniform density and vorticity. We use both analytical (Lie transforms) and numerical tools to study this rotation, with different shapes of the core. We express in particular the proper frequencies of the system, because they characterize the response of Mercury to the different solicitations, due to the orbital motion of Mercury around the Sun. We show that, contrary to its size, the shape of the core cannot be determined from observations of either longitudinal or polar motions. However, we highlight the strong influence of a resonance between the proper frequency of the core and the spin of Mercury that raises the velocity field inside the core. We show that the key parameter is the polar flattening of the core. This effect cannot be directly derived from observations of the surface of Mercury, but we cannot exclude the possibility of an indirect detection by measuring the magnetic field. ", "introduction": "\\par Mercury is the target of two current space missions (see e.g. \\cite{msgnm04}). The first one, MESSENGER (NASA), already performed three flybys on January 14, October 6, 2008, and September 29, 2009, before orbit insertion in March 2011. The second one, BepiColombo (ESA/JAXA), is planned to be launched in 2014 and to reach Mercury in 2020. The preparation of these two missions motivated an in-depth study of the rotation of Mercury. \\par The rotation of Mercury is a unique case in the Solar System because of its $3:2$ spin-orbit resonance, Mercury performing exactly 3 rotations during 2 revolutions about the Sun \\citep{pd65}. It corresponds to an equilibrium state \\citep{c65} known as Cassini State 1. Recently, radar Earth-based measurements by \\citet{mpjsh07} detected a 88-day longitudinal libration of Mercury with an amplitude $\\phi$ of $35.8\\pm2$ arcsec. This amplitude being nearly twice too high to be consistent with a rigid Mercury, it is the signature of an at least partially molten core. If we consider Mercury as a 2-layered body with a rigid mantle and a spherical liquid core that does not follow the short-period ($\\approx88$ days) excitations and does not interact with the mantle, we can derive from this amplitude the inertia of the mantle plus crust. In particular, naming $C_m$ the inertial polar momentum of the mantle and $A2$~\\mearth, and that no significant amount of freshly made $^{26}$Al was added to the protoplanetary disk after the CAI formation. The maximum amount of $^{26}$Al that could have been synthesized by in situ particle irradiation during the short duration of CAI formation \\citep[$\\sim 10^5$~yr;][]{biz04} can account for the canonical $^{26}$Al/$^{27}{\\rm Al}=5\\times 10^{-5}$ over a rocky reservoir of only $\\sim$0.1~\\mearth~\\citep{dup07}. Thus, the origin of this SLR cannot be related to the nonthermal activity of the young Sun and has to be searched for in a stellar nucleosynthetic event contemporary with the formation of the solar system. An origin of SLRs in an asymptotic giant branch (AGB) star has been proposed \\citep{was94,tri09}, but AGB stars are not associated with star-forming regions and the probability of a chance encounter between an AGB star and a star-forming molecular cloud is very low \\citep{kas94}. It is more likely that the protosolar system was contaminated by material freshly ejected from a massive star, either a Type II SN \\citep[e.g.][]{cam77} or a Wolf-Rayet (WR) star \\citep[e.g.][]{arn97}. Massive stars have a profound influence on the surrounding molecular clouds and the process of star formation \\citep[e.g.][]{lee07}. \\citet{cam77} first suggested that a supernova (SN) responsible for injecting SLRs into the presolar nebula may also have been responsible for triggering the formation of the solar system. Detailed numerical simulations have shown that such simultaneous triggering and injection is possible, but the injection efficiency is lower than required \\citep{bos10}. Alternatively, it has been suggested that a nearby SN ($\\sim$0.3~pc) may have injected SLRs into the already-formed protoplanetary disk of the solar system \\citep[see][and references therein]{oue07}. In this scenario, it is assumed that the Sun was born in a large stellar cluster containing massive stars. But this model is questionable, because (1) protoplanetary disks in the vicinity of massive stars are exposed to a rapid photoevaporation and (2) the main-sequence lifetime of even the most massive stars is too long as compared to the mean lifetime of protoplanetary disks \\citep{gou08}. \\citet{gou09} and \\citet{gai09} recently suggested that the Sun is born in a stellar cluster of second generation, whose formation was triggered by the activity of a neighboring OB association. There are many observations of OB associations divided in spatially separated subgroups of different ages \\citep[e.g.][]{bla64}, as well as observations of young stellar objects located on the border of H~II regions \\citep[e.g.][]{kar09} and superbubbles \\citep[e.g.][]{lee09}. Adopting such an astrophysical context, we study in this Letter how hot stellar debris enriched in $^{26}$Al could be injected into a cold protostellar nebula. We show in Section~2 that $^{26}$Al produced by a population of massive stars in an OB association may not be delivered into molecular cores efficiently enough. We then study in Section~3 a possibility already mentioned in the pioneering work of \\citet{arn97} and more recently by \\citet{gai09} that the presolar nebula was contaminated by $^{26}$Al produced by a WR star that escaped from its parent cluster. ", "conclusions": "" }, "1003/1003.4471_arXiv.txt": { "abstract": "{} {Cold steady-state disk wind theory from near Keplerian accretion disks requires a large scale magnetic field at near equipartition strength. However the minimum magnetization has never been tested with time dependent simulations. We investigate the time evolution of a Shakura-Sunyaev accretion disk threaded by a weak vertical magnetic field. The strength of the field is such that the disk magnetization falls off rapidly with radius.} {Four 2.5D numerical simulations of viscous resistive accretion disk are performed using the magnetohydrodynamic code PLUTO. In these simulations, a mean field approach is used and turbulence is assumed to give rise to anomalous transport coefficients (alpha prescription). } {The large scale magnetic field introduces only a small perturbation to the disk structure, with accretion driven by the dominant viscous torque. However, a super fast magnetosonic jet is observed to be launched from the innermost regions and remains stationary over more than 953 Keplerian orbits. This is the longest accretion-ejection simulation ever carried out. The self-confined jet is launched from a finite radial zone in the disk which remains constant over time. Ejection is made possible because the magnetization reaches unity at the disk surface, due to the steep density decrease. However, no ejection is reported when the midplane magnetization becomes too small. The asymptotic jet velocity remains nevertheless too low to explain observed jets. This is because of the negligible power carried away by the jet.} {Astrophysical disks with superheated surface layers could drive analogous outflows even if their midplane magnetization is low. Sufficient angular momentum would be extracted by the turbulent viscosity to allow the accretion process to continue. The magnetized outflows would be no more than byproducts, rather than a fundamental driver of accretion. However, if the midplane magnetization increases towards the center, a natural transition to an inner jet dominated disk could be achieved.} ", "introduction": "Accretion disks are commonly found in young stars, active galactic nuclei, cataclysmic variables and microquasars. In order to allow material to accrete onto a central object, it is necessary to lose some angular momentum in an efficient way. This is possible in a disk in one of two ways, either by radial outward transport in a disk by turbulent transport \\citep{1973A&A....24..337S,1974MNRAS.168..603L} or spiral waves \\citep{1999A&A...349.1003T}, or vertical transport upwards out of the disk in a jet \\citep{1982MNRAS.199..883B}. Two extreme possible disk structures can then be identified, corresponding to each of these two processes of angular momentum removal. The Jet Emitting Disk (hereafter JED) is threaded by a large scale magnetic field of bipolar topology driving a jet (defined here as super-fast magnetosonic flow). The dominant torque in the JED is magnetic, due to the large braking lever arm of the jet, defined by a length scale equivalent to the Alfv\\'en radius \\citep{1992ApJ...394..117P}. The pioneering jet model by \\citet{1982MNRAS.199..883B} establishes a relationship between the mass loading and the magnetic lever arm of magnetocentrifugally driven outflows. But the magnetic field strength was left unconstrained, so in principal any magnetization at the disk surface could drive a low-enthalpy outflow. The reason lies in the fact that an ideal MHD jet model assumes the mass loss and does not compute it as function of the disk parameters. This was precisely the goal of semi-analytical studies done by e.g. \\citet{1995A&A...295..807F} and \\citet{1997A&A...319..340F}, where the disk structure has been consistently computed: these authors showed that steady-state {\\bf cold} jets can be produced only with a vertical field close to equipartition. A few numerical experiments tested the accretion-ejection connection in a consistent way: axisymmetric magneto-hydrodynamic (MHD) simulations of resistive accretion disks reporting the production of self-confined, quasi-steady super-fast jets \\citep{2002ApJ...581..988C, 2004ApJ...601...90C, 2007A&A...469..811Z, 2009MNRAS.400..820T} confirmed most of the results obtained with semi-analytical models. They were however done with a large disk magnetization\\footnote{The magnetization is related to the usual plasma beta by $\\mu= 2/\\beta$ in gas pressure supported disks. It is however a more general concept as s it is defined with the total pressure $P_\\mathrm{gas} +P_\\mathrm{rad}$.}, in the range $\\mu \\sim 0.2 - 1$. The inner regions of the disk from whence jets are observed to be emitted are expected to be JED-like. In the specific case of outflows from young stars, extrapolation of slitless images of Class II jets have constrained the launching region to be confined to a zone of radial extent $\\sim$ a few AU close to the centre of the disk \\citep{2004ApJ...609..261H,2007LNP...723...21C}. The outer regions are expected to behave more like the well studied standard accretion disk, hereafter SAD \\citep{1973A&A....24..337S,2002apa..book.....F}. Here the characteristic lengthscale over which the viscous torque is exerted is of the order of $\\sim \\alpha_\\mathrm{v} h$, where $\\alpha_\\mathrm{v}$ measures the level of the turbulence. Such a turbulence is assumed to arise from the development of magnetic instabilities that are triggered in the disk whenever a magnetic field is present \\citep{1991ApJ...376..214B}. This field must however be below equipartition strength (namely $B_z^2/\\mu_0 \\ll P$) to avoid the stabilizing effect of the magnetic tension. Therefore, the high magnetization required by a JED impedes the development of disk turbulence which is however required to support a steady launching: that would leave only a very tiny parameter space for stationary ejection to take place. The SAD-JED structure has been put forward in several papers, e.g. \\citet{2006A&A...447..813F,2008A&A...479..481C,2008NewAR..52...42F}. The study of low magnetization accretion regimes has been attempted making use of fully 3D global simulations of accretion disks, threaded by a weak large scale magnetic field. MRI sets in and accretion is quickly established \\citep{2002ApJ...573..738H}. The remarkable result is that outflows are also produced \\citep{2003ApJ...592.1042I}, especially when the imposed field is of bipolar topology \\citep{Beckwith:2009ss}. However, many questions remain open: What controls the mass loss in these simulations? Will the outflowing plasma become a self-confined jet? Is grid resolution enough to properly describe the turbulent cascade? The fact is that it is still impossible to properly follow turbulence while solving for the long term evolution of large scale systems. As a consequence, the question of super-fast magnetosonic jet formation from weakly magnetized disks is still open. In this paper we address this issue using 2.5D numerical MHD simulations based on a mean field approximation. We explore the accretion-ejection processes from a quasi-standard accretion disk where the magnetization is very low (smaller than $10^{-3}$). Since the magnetic field is low, we assume that turbulence triggered by the MRI is indeed present but that it provides mainly anomalous transport coefficients: a viscosity $\\nu_\\mathrm{v}$ and a magnetic diffusivity $\\nu_\\mathrm{m}$. On the other hand, we do not expect to observe any MRI feature (such as channel flows for instance) in our simulation because of the presence of explicit viscosity and magnetic diffusivity effects. While measurements of the turbulent viscosity in MRI induced turbulence have been extensively reported in the literature, it is only very recently that such a work has been done for the turbulent magnetic diffusivity \\citep{Lesur:2009bf, Guan:2009gd}. In particular \\citet{Lesur:2009bf} showed that the turbulent magnetic diffusion scales like a resistivity tensor with dominant diagonal terms. Also, as a first approximation, an isotropic value can be safely used. Finally, the effective Prandtl number $ \\mathcal{P}_\\mathrm{m} = \\nu_\\mathrm{v}/\\nu_\\mathrm{m}$, given by the ratio of turbulent viscosity and diffusivity, has been found to be of order unity. The mean field approximation has been successfully employed in a number of semi-analytical \\citep[e.g.][]{1995A&A...295..807F, 1995ApJ...444..848L, 2000A&A...353.1115C, 2001ApJ...553..158O, 2008ApJ...677.1221R} and numerical applications \\citep[e.g.][]{2002ApJ...581..988C, 2003ApJ...589..397K, 2003A&A...398..825V, 2006A&A...460....1M, 2007A&A...469..811Z, 2009MNRAS.399.1802R} related to the study of magnetized accretion-ejection flows. Beside having a precise control of the diffusive and transport phenomena, the numerical experiments based on this approach provide laminar flow solutions which can be compared to semi-analytical models. In section 2, we describe the numerical method used, the boundary and initial conditions. Section 3 is devoted to the description and discussion of the results obtained. Surprisingly, super-fast jets are indeed obtained from a finite disk region and remain stable for a time span never previously achieved in the literature. Section 4 summarizes our findings and, in a companion paper (Murphy et al., in prep), we will examine the long standing issue of the magnetic field redistribution within the disk on long (accretion) time scales. ", "conclusions": "\\label{ConcludingRemarks} In this paper, we performed four 2.5D numerical MHD simulations of a resistive viscous accretion disk threaded by a weak magnetic field. The initial magnetic field distribution was chosen so that the disk magnetization $\\mu= B_z^2/ P$ decreases radially from the central object. Our reference simulation, done with a maximum value of $\\mu= 2\\times 10^{-3}$, has run for more than 950 Keplerian orbits at the inner radius and is therefore the longest to date. It is shown that the disk structure resembles that of a standard Shakura \\& Sunyaev disk with accretion controlled by the turbulent (alpha) viscous torque only. However, a super fast magnetosonic, self-confined jet is observed to be launched from the inner disk regions. It is first time that (i) steady-state super-FM jets are produced from a weakly magnetized disk and (ii) from a finite disk region that remained constant over time. The power carried away by these jets is tiny and directly related to the negligible torque on the disk. The dynamics of the jet and its propagation into the medium will be studied in a forthcoming paper. Here, we focused on the jet acceleration region where the flow crossed the three MHD critical surfaces (Slow Magnetosonic, Alfv\\'en and Fast Magnetosonic). The critical issues of mass loading and initial jet acceleration (the crossing of the SM surface) are shown to be strongly affected by the unavoidable steep decrease of the density profile at the disk surface. Such an effect has been underestimated in previous simulations. It is the quality of the grid resolution {\\em at the disk surface} that ultimately determines the amount of ejected mass. One way to solve this problem is to use either an enhanced resolution at the disk surface, a less diffusive algorithm, a higher order method or an adaptive grid which refines on the density gradient. We argue however that this feature might mimic some additional heat input at the disk surface, as explored for instance by \\citet{Ogilvie:1998lh}, \\citet{2001ApJ...553..158O}, \\citet{2000A&A...361.1178C}. This aspect is extremely promising as most astrophysical accretion disks do probably have superheated layers due to irradiation by the central object (young stars, cataclysmic variables) and/or some X-ray source (e.g. around black holes). As a consequence, ``cold'' (e.g. isothermal or adiabatic) ejection is probably never achieved in Nature. This allows also to relax the constraint of equipartition fields needed for driving jets as our jets were obtained from a very low magnetized disk (but not too low). This opens a new fascinating topic: the magnetic history of any given object. One might indeed consider accretion disks displaying a whole continuum in ejection efficiency, from jets carrying a sizable fraction (if not most) of the released accretion power to jets that are a mere epiphenomenon of accretion. For any given object, the key parameter would be the disk magnetization. This clearly deserves further investigation." }, "1003/1003.4192_arXiv.txt": { "abstract": "The transition from liquid metal to silicate rock in the cores of the terrestrial planets is likely to be accompanied by a gradient in the composition of the outer core liquid. The electrical conductivity of a volatile enriched liquid alloy can be substantially lower than a light-element-depleted fluid found close to the inner core boundary. In this paper, we investigate the effect of radially variable electrical conductivity on planetary dynamo action using an electrical conductivity that decreases exponentially as a function of radius. We find that numerical solutions with continuous, radially outward decreasing electrical conductivity profiles result in strongly modified flow and magnetic field dynamics, compared to solutions with homogeneous electrical conductivity. The force balances at the top of the simulated fluid determine the overall character of the flow. The relationship between Coriolis and Lorentz forces near the outer boundary controls the flow and magnetic field intensity and morphology of the system. Our results imply that a low conductivity layer near the top of Mercury's liquid outer core is consistent with its weak magnetic field. ", "introduction": "Variations in the physical properties of fluids in planetary dynamos define the character of the observed intrinsic magnetic field (e.g., strength, geometry and time variability). Changes in the electrical conductivity of the fluid as a function of depth may become relevant in the context of terrestrial and gas giant planets. In this paper we explore (with a focus on the terrestrial planets) how the radial variation of electrical conductivity in planetary cores may result in changes to dynamo-generated magnetic fields. \\subsection{Terrestrial planets} The cores of the terrestrial planets are composed principally of iron, with minor but significant amounts of nickel and lighter elements. It has long been known that an iron-nickel core would have too high a density to be compatible with Earth's moment of inertia and seismic data \\citep[e.g.,][]{Birch1952,Poirier1994}. A compatible Earth core density model can result from the inclusion of about 8\\% by weight of one or more light elements. Detailed models of core composition are based primarily on the constraints of seismology, mineral physics, geochemistry, metallurgy, and cosmochemisty. Silicon, sulphur and oxygen are the primary candidates for the light elements. Sulphur is likely to be a significant component but the depletion of light elements in the process of accretion in the inner solar system limits sulphur to about 2~wt\\% in the core. Recent reviews of core differentiation and composition distinguish between models considering Silicon versus those considering Oxygen as the primary light element. Composition models give weight percents of Fe $\\simeq$ 85\\%-88\\%, Ni $\\simeq$ 5\\% Si $\\simeq$ 0-7\\%, O $\\simeq$ 0-4\\% and S $\\simeq$ 2\\% \\citep[e.g][]{McDonough2003,Wood2006}. Solidification of a more or less pure iron-nickel inner core may exclude the lighter elements, which would then be enriched in the outer core. The ratio of the inner core radius (1221~km) to the core radius (3480~km) is 0.35 and the mass of the inner core is only about 5\\% of the total core mass. So the bulk composition of the outer core is only slightly different than that of the whole core. Convection in the liquid outer core is driven by a combination of compositional and/or thermal buoyancy. Thermal buoyancy available to drive convection and dynamo action originates primarily from the latent heat of solidification at the inner core boundary (ICB) and possibly from cooling at the core-mantle boundary (CMB). Secular cooling at the CMB does not guarantee a source of convective instability since heat can be conducted through a stably stratified layer (a super-adiabatic heat flux is needed). Compositional buoyancy originates at the ICB due to light element enrichment in the residual liquid associated with inner core solidification. A source of compositional buoyancy near the CMB could come from precipitation from a silicate enriched layer at the top of the core, leaving a light element depleted, heavy residual liquid and a silicate sediment layer at the CMB \\citep{Buffett2000}. Temperature and pressure dependence of electrical conductivity may lead to sizable variations with depth. \\citep[][]{Stacey2001}. Assuming a well-mixed outer core and a Si concentration of $X_{S_i}=0.25$ in a dual species alloy, the authors found the most extreme variation in the fluid core of Earth results in a factor of 1/2 difference for the electrical conductivity and about 3/4 for the thermal conductivity from the ICB to the CMB. Although the effects of different impurities in the alloy have not been explored experimentally, \\citet{Stacey2001} argue that the effects of impurities other than Si do not significantly change their estimates. A later revision \\citep{Stacey2007} found that the variation in electrical and thermal conductivities had been overestimated due to the assumption of treating Fe as an electronically simple metal. Their later results predict a difference in the electrical conductivity of a factor of 0.78 between that of the CMB when compared to the ICB's. Those estimates were made under the assumption of a well-mixed outer core. However gradients in material properties would be amplified in a compositionally stratified layer. The existence of a thermally and/or compositionally stably stratified layer near the top of the Earth's core has been both, suggested and argued \\citep{Jacobs1975,Fearn1981, Lister1998, Braginsky1993, Braginsky2007}. Here we will focus on the possibility of a compositionally stratified layer near the top of the core since the electrical conductivity in such a layer would decrease with radius (due to the higher light-element concentration when compared to the bottom of the fluid core). Two mechanisms by which a compositionally stratified layer can grow are by chemical diffusion from the mantle directly to the top of the core \\citep{Lister1998} and by buoyant transport of light element enriched residual liquid from inner core solidification \\citep{Moffatt1994}. Chemical diffusion is a slow process and would not result in a layer thickness greater than about 10~km \\citep{Lister1998}. On the other hand, buoyant rise of light element rich residual liquid could be an efficient process to build a layer of greater thickness. A thick layer with a significantly anomalous density would likely be detected as a seismically fast layer. Seismological constraints have been so far inconclusive. However, recent seismological results using outermost core waveforms seem to be consistent with the existence of a low density layer of about 100 km thickness \\citep{Alexandrakis2007,Tanaka2007} It has been confirmed that Mercury has a liquid iron core \\citep{Margot2007}. Furthermore, it is likely that Mercury's weak intrinsic magnetic field is generated by a dynamo in its liquid outer core. Neither the details of Mercury's core composition, nor the size of its inner core is known \\citep{Hauck2004,Solomon2008,Heimpel2008}. Observations of Mercury's contractional lobate crustal faults imply a planetary radial contraction of about 3~km, small compared to an estimated 17~km of contraction that would result from a completely solidified core. This implies that Mercury may have a relatively large outer core, perhaps earth-like proportion. For a thick-shell outer core to exist in Mercury, given its small size, \\citet{Hauck2004} estimates that the light elements (such as S or Si) in the core have a relatively high concentration. This suggests the possibility of a large and stratified Hermean outer core with a thick, and stably stratified outermost layer. Such a layer would be compositionally stratified \\citep[in contrast to thermally stratified models previously proposed, e.g.,][]{Christensen2006a} with a radial increase in the proportion of a light elements, and a radial decrease in electrical conductivity. \\subsection{Gas giant planets} The magnetic field in the gas giants is generated within a metallic hydrogen region. Experimental results have found that the transition from metallic to molecular hydrogen yields a wide range of pressures where the electrical conductivity is non-negligible thus varying slowly with depth in planetary interiors \\citep{Nellis2000}. The internal structure has been deduced to first order based on measurements of the moment of inertia and total mass of each planet \\citep[e.g.,][]{Guillot2005}. However, constraints on hydrogen and helium mixtures at high pressures need to be found in order to better determine the internal structure of the gas giants. The additional effect of helium may complicate further the underdetermined internal layering of gas giants \\citep[e.g.,][]{Stevenson2008}. The metallic to molecular hydrogen transition is of particular interest from the point of view of the internal dynamics. The depth and radial extent of this transition are important for magnetic field generation and in understanding the observable magnetic field morphology. This is a task for future investigation. In this paper, we present results from a numerical dynamo model with radially varying electrical conductivity \\citep{Gomez2007a}. To implement the radially variable conductivity, we modified an existent numerical code \\citep[originally MagIC 2.0,][]{Wicht2002}, that uses the Boussinesq approximation. We focus this study on the effect of the varying conductivity on the dynamo action, and on the generated magnetic field. With this new implementation we performed a set of tests to analyze its consistency with previously published work. In section~\\ref{sec:strat} we include a review of numerical models that worked with stratified liquids. In section~\\ref{sec:method} we present the necessary modification to the dynamo equations, and to the numerical code. The parameters explored for twenty runs studied in this paper are included in section~\\ref{sec:param}. We present the results in section~\\ref{sec:result}, the discussion and conclusions are found in section~\\ref{sec:diss} and~\\ref{sec:con}, respectively. We also included a table of symbols in appendix~\\ref{sec:symbols}. ", "conclusions": "\\label{sec:con} In our models, the dynamo-generated magnetic field morphology and intensity are strongly affected by the relative strength of viscous, Coriolis and Lorentz forces near the outer boundary. The relative scale of Ekman and Hartmann boundary layer thicknesses is determined by an electrical conductivity gradient. For uniform electrical conductivity models Coriolis and Lorentz forces are typically of of the same order, yielding a magnetostrophic balance. Low electrically conductivity near the top boundary can separate the Ekman-Hartman layer into a thin Ekman layer and a thicker Hartmann layer, resulting in changes of the detailed magnetic field morphology, and a large-scale external magnetic field that is relatively weak (see Table~\\ref{tab:averages}). Given the likelihood of a high concentration of light element in Mercury's core, it is plausible that a low electrical conductivity layer is present near the core mantle boundary. This means that the dynamics of the boundary layers obtained in our models may be applicable to conditions in Mercury's core. Our results imply that a low conductivity layer is consistent with Mercury's weak observable magnetic field. However, while radially variable electrical conductivity is the mechanism studied in this paper, it is one of several models that can result in weak magnetic fields. With the anticipated arrival of the MESSENGER spacecraft in orbit, we will soon be in a position to use detailed mapping to better constrain the relative contributions of the various dynamical processes that generate Mercury's global magnetic field." }, "1003/1003.1312_arXiv.txt": { "abstract": "{ Giant low surface brightness (GLSB) galaxies are commonly thought to be massive, dark matter dominated systems. However, this conclusion is based on highly uncertain rotation curves. We present here a new study of two prototypical GLSB galaxies: Malin~1 and NGC~7589. We re-analysed existing \\hi observations and derived new rotation curves, which were used to investigate the distributions of luminous and dark matter in these galaxies. In contrast to previous findings, the rotation curves of both galaxies show a steep rise in the central parts, typical of high surface brightness (HSB) systems. Mass decompositions with a dark matter halo show that baryons may dominate the dynamics of the inner regions. Indeed, a ``maximum disk'' fit gives stellar mass-to-light ratios in the range of values typically found for HSB galaxies. These results, together with other recent studies, suggest that GLSB galaxies are systems with a double structure: an inner HSB early-type spiral galaxy and an outer extended LSB disk. We also tested the predictions of MOND: the rotation curve of NGC~7589 is reproduced well, whereas Malin~1 represents a challenging test for the theory. } ", "introduction": "Low surface brightness (LSB) galaxies span a wide range of galaxy sizes, masses and morphologies, from dwarf spheroidals and dwarf irregulars to medium-size late-type disks (Sc-Sd) and to bulge-dominated early-type spirals (Sa-Sb). According to \\citet{Beijersbergen1999}, LSB spirals form a LSB Hubble sequence, parallel to the classical HSB one. Giant low surface brightness (GLSB) galaxies are usually considered as extreme cases of early-type LSB spirals \\citep{Beijersbergen1999}. They have exceedingly extended LSB disks, with scale lengths ranging from $\\sim$10 to $\\sim$50 kpc, and central luminous components resembling the bulges of HSB spirals (e.g. \\citealt{Bothun1987, Sprayberry1995}). They are massive systems, with $L \\sim L_{*}$. Like the ordinary LSB spirals, GLSB galaxies have low \\hi surface densities \\citep{Pickering1997}, but they are among the most gas rich galaxies known, with $M_{\\rm{\\hi }}\\approx 10^{10} M_{\\odot}$ \\citep{Matthews2001b}. Few dynamical studies exist on GLSB galaxies. \\citet{Pickering1997} studied four GLSB galaxies and found slowly rising rotation curves, similar to those of late-type LSB disks (e.g. \\citealt{deBlok1997,DeNaray2006}). They concluded that GLSB galaxies are ``the first examples of galaxies that are both massive and dark matter dominated''. However, they warned that their rotation curves are highly uncertain, due to the low signal-to-noise ratio and the low spatial resolution of the observations. A slowly rising rotation curve in the presence of an inner concentration of light, as seen in those galaxies (see top panels of figure \\ref{fig:deco1}), is in marked contrast with the rule that there is a close correlation between the concentration of light and the shape of the rotation curve (\\citealt{Sancisi2004} and references therein). For this reason, \\citet{Sancisi2007} started a re-analysis of the 21-cm data from \\citet{Pickering1997}. This paper concludes that preliminary work with the study of the structure and dynamics of two GLSB galaxies: Malin~1 and NGC~7589. ", "conclusions": "" }, "1003/1003.5311_arXiv.txt": { "abstract": "Detection of the radiation emitted from some of the earliest galaxies will be made possible in the next decade, with the launch of the James Webb Space Telescope (JWST). A significant fraction of these galaxies may host Population (Pop) III star clusters. The detection of the recombination radiation emitted by such clusters would provide an important new constraint on the initial mass function (IMF) of primordial stars. Here I review the expected recombination line signature of Pop III stars, and present the results of cosmological radiation hydrodynamics simulations of the initial stages of Pop III starbursts in a first galaxy at $z$ $\\sim$ 12, from which the time-dependent luminosities and equivalent widths of IMF-sensitive recombination lines are calculated. While it may be unfeasible to detect the emission from Pop III star clusters in the first galaxies at $z$ > 10, even with next generation telescopes, Pop III star clusters which form at lower redshifts (i.e. at $z$ < 6) may be detectable in deep surveys by the JWST. ", "introduction": "In this contribution, I will address three key questions pertaining to the observational signatures of the first galaxies, and in particular to the prospects for the detection and identification of Pop III stars. These questions are the following: \\begin{itemize} \\item How can Population (Pop) III stars be identified observationally and their initial mass function (IMF) constrained? \\item How long did Pop III star formation continue after the epoch of the first stars? \\item Will observational facilities in the coming years, and in particular the {\\it James Webb Space Telescope} (JWST), be able to detect and identify Pop III stellar populations? \\end{itemize} The next three sections are devoted to addressing these three questions in the order given above; in the final section I will briefly summarize the main conclusions. \\begin{figure} \\includegraphics[height=.3\\textheight]{escape_fractions.ps} \\caption{The escape fraction of H~{\\sc i}-ionizing photons ({\\it left panel}) and that of He~{\\sc ii}-ionizing photons ({\\it right panel}), from a Pop III star cluster formed in a first galaxy at $z$ $\\sim$ 12. Each line corresponds to a different choice of IMF and total number of stars, as labeled. There is a tight anticorrelation between the escape fraction of H~{\\sc i}-ionizing photons and the emission in hydrogen recombination lines (e.g. Ly$\\alpha$ and H$\\alpha$; see Figs. 2 and 3). For most cases, however, the negligible escape fraction of He~{\\sc ii}-ionizing photons leads to a tight correlation between the luminosity emitted in the He~{\\sc ii} $\\lambda$1640 line and the total mass contained in stars, for a given IMF (see \\protect\\cite{JLJetal2009}).} \\end{figure} ", "conclusions": "In closing, I would like to highlight the following conclusions corresponding to the three key questions addressed in the work presented here: \\begin{itemize} \\item The Pop III IMF can be constrained with detection of helium recombination emission (particularly the He II $\\lambda$1640 line). This emission varies due to both stellar evolution and hydrodynamic evolution of photoionized regions. \\item In rare regions which are reionized at early times, Pop III star formation may extend well beyond the epoch of the first stars and galaxies. \\item If this is so, then planned JWST surveys may detect Pop III stellar clusters at redshifts $z$ < 6 and allow for constraints to be placed on the IMF. \\end{itemize} \\begin{figure} \\includegraphics[height=.26\\textheight]{Figure4.ps} \\caption{The number density $n_{\\rm III}$ of Pop~III star clusters formed in reionized regions of the universe, taking into account enrichment of the IGM by galactic winds, and normalized to a cluster lifetime $t_{\\rm cluster}$ = 2 Myr. The solid lines correspond to a minimum halo circular velocity for star formation of 20 km~s$^{-1}$, while the dashed lines correspond to a minimum of 30 km~s$^{-1}$. For each series of lines, the top ({\\it black}) line is a model which neglects external metal enrichment, while the colored lines correspond to different metal-enriched wind velocities, as labeled. The number density at which one cluster per unit redshift is expected to be within the planned JWST Deep-Wide Survey area is shown by the {\\it dotted} line (see \\protect\\cite{JLJ2010}).} \\end{figure} \\begin{theacknowledgments} I would like to thank the conference organizers, Dan Whalen, Naoki Yoshida, and Volker Bromm, for hosting a most enjoyable and productive event, as well as for allowing me to present this work. I am also grateful for support from the Theoretical Modeling of Cosmic Structures (TMoX) Group at MPE. \\end{theacknowledgments}" }, "1003/1003.5916_arXiv.txt": { "abstract": "In recent work we suggested that photons of energy $>$100 MeV detected from GRBs by the {\\it Fermi} Satellite are produced via synchrotron emission in the external forward shock with a weak magnetic field -- consistent with shock compressed upstream magnetic field of a few tens of micro-Gauss. Here we investigate whether electrons can be accelerated to energies such that they radiate synchrotron photons with energy up to about 10 GeV in this particular scenario. We do this using two methods: (i) we check if these electrons can be confined to the shock front; and (ii) we calculate radiative losses while they are being accelerated. We find that these electrons remain confined to the shock front, as long as the upstream magnetic field is $\\gae 10 \\mu$G, and don't suffer substantial radiative losses, the only condition required is that the external reverse shock emission be not too bright: peak flux less than 1 Jy in order to produce photons of 100 MeV, and less than $\\sim$100 mJy for producing 1-GeV photons. We also find that the acceleration time for electrons radiating at 100 MeV is a few seconds (in observer frame), and the acceleration time is somewhat longer for electrons radiating at a few GeV. This could explain the lack of $>$100 MeV photons for the first few seconds after the trigger time for long GRBs reported by the {\\it Fermi} Satellite, and also the slight lag between photons of GeV and 100 MeV energies. We model the onset of the external forward shock light curve in this scenario and find it consistent with the sharp rise observed in the 100-MeV light curve of GRB080916C and similar bursts. ", "introduction": "The {\\it Fermi} Satellite has detected 18 GRBs (Gamma-ray Burst) at energies $>$100 MeV by LAT (Large Area Telescope). This emission can be described as follows. The first 100 MeV photons arrive $\\sim 1$ s (in the host galaxy frame) after the trigger time, for long GRBs; the trigger time is the time when low energy photons ($\\sim1$ MeV) are first detected by the GBM (Gamma-ray Burst Monitor aboard {\\it Fermi}). The 100 MeV light curve rises fast until it peaks and then it decays as a single power-law for a long duration of time (of order 10$^3$ s) -- much longer than the duration of the lower energy photons detected by GBM -- until it falls below the detector's sensitivity. Radiation above 100 MeV from GRBs has been suggested to be produced via the synchrotron mechanism in the external forward shock (Kumar \\& Barniol Duran 2009, 2010); the external forward shock scenario was first proposed by Rees \\& M\\'esz\\'aros (1992), M\\'esz\\'aros \\& Rees (1993), Paczy\\'nski \\& Rhoads (1993), and since then it has been used widely, see, e.g., M\\'esz\\'aros \\& Rees (1997), Sari, Piran, Narayan (1998), Dermer \\& Mitman (1999), for a comprehensive review see, e.g., Piran (2004) and references therein. After our initial suggestion, many groups have also considered and provided evidence for this origin of the $>$100 MeV radiation (Gao et al. 2009; Corsi, Guetta, Piro 2010; De Pasquale et al. 2010; Ghirlanda, Ghisellini, Nava 2010; Ghisellini, Ghirlanda, Nava 2010). The magnetic field required for this model is consistent with being produced via shock-compressed seed magnetic field in the CSM (circum-stellar medium) of strength of a few tens of micro-Gauss. The peak of the 100 MeV light curve can be attributed to the deceleration time which is the time it takes for the GRB-jet to transfer about half of its energy to the external medium. We investigate in this work whether electrons in the external forward shock can be accelerated to sufficiently high Lorentz factors, even for a small CSM magnetic field of a few tens of $\\mu$G, so that the synchrotron radiation can extend to $\\sim 10$ GeV as seen by {\\it Fermi}/LAT for a number of GRBs. We study the electrons acceleration in the context of diffusive shock acceleration (e.g., Krymskii 1977, Axford, Leer \\& Skadron 1978, Bell 1978, Blandford \\& Ostriker 1978, Blandford \\& Eichler 1987), which was developed for non-relativistic shocks and has now been developed to consider relativistic shocks (semi-) analytically (e.g. Gallant \\& Achterberg 1999, Achterberg et al. 2001) and recently using 2D particle-in-cell simulations (e.g. Spitkovsky 2008a,b, Keshet et al. 2009). We assume that the electrons acceleration proceeds in the Bohm diffusion limit and that the magnetic field downstream is simply shock-compressed upstream magnetic field (other possibilities are considered in, e.g., Milosavljevi\\'c \\& Nakar 2006, Sironi \\& Goodman 2007, Goodman \\& MacFadyen 2008, Couch, Milosavljevi\\'c, Nakar 2008). If the downstream magnetic field is simply the shock-compressed large-scale upstream field, then the field component perpendicular to the shock normal is amplified, while the parallel component is not. In this case, the downstream magnetic field will be mainly pointing to the direction perpendicular to the shock front normal, therefore particles trying to cross the shock front from downstream to upstream will find it difficult to catch up with the shock front, which moves with a speed of $\\sim c/3$ with respect to the downstream medium (see, e.g., Achterberg et al. 2001, Lemoine, Pelletier \\& Revenu 2006, Pelletier, Lemoine \\& Marcowith 2009). One way that the particles might return to the upstream is if there is efficient cross-field diffusion of particles, which might occur if turbulent magnetic field is produced downstream (Jokipii 1987, Achterberg \\& Ball 1994, Achterberg et al. 2001). In principle, the turbulent magnetic field could dominate the shock-compressed field throughout the downstream region. However, it seems that although some turbulence is present just downstream of the shock front it does not persist across the entire downstream region (see recent simulations by Sironi \\& Spitkovsky 2010 that show that magnetic field is amplified only right behind the shock front and returns to the shock-compressed value far downstream). In this case, much of the radiation is produced by particles swept downstream where the turbulence has died out and the magnetic field is consistent with the shock-compressed value. We also note that as long as the thickness of the turbulent magnetic field layer is smaller than the thickness of the shocked fluid divided by $(B_t/B_d)^2$ then the energy loss in the turbulent layer is small; $B_t$ is the turbulent magnetic field strength and $B_d$ is the shock-compressed magnetic field. Therefore, in this work we neglect energy loss in the turbulent magnetic field layer since it persists for a very short distance compared to the thickness of the shocked fluid (see, e.g., Keshet et al. 2009 and references therein). This work is organized as follows. In Section 2 we address the question of high-energy electron confinement upstream and downstream of the shock front, and also radiative losses suffered by electrons in between acceleration. Also, in Section 2, we discuss the lag of the $>$100 MeV light curves observed by {\\it Fermi} LAT for several GRBs in light of our results on electron acceleration. In Section 3, we calculate the rise of the external forward shock light curve, taking into consideration the non-zero time to accelerate electrons to high enough energies so they can radiate at $>$100 MeV. We present our conclusions in Section 4. ", "conclusions": "In this paper we have investigated the acceleration of electrons via diffusion shock acceleration in the external forward shock of GRBs, and its implications for the high-energy photon detection by the {\\it Fermi} Satellite. The external shock model, with a weak magnetic field, has been proposed as the origin of the observed $>$100 MeV emission detected by the {\\it Fermi} Satellite from a number of GRBs (Kumar \\& Barniol Duran 2009, 2010). We find that high-energy electrons of Lorentz factor $\\sim 10^8$, required for producing $\\sim$10 GeV photons via the synchrotron process, can indeed be accelerated in an external shock that is moving through a CSM with a magnetic field of strength a few tens of $\\mu$G; they remain confined to the shock front as long as the upstream magnetic field is $\\gae 10 \\mu$G. We have also calculated the time it takes for electrons to be accelerated to a Lorentz factor $\\sim10^7$ so that they can radiate synchrotron photons at $\\sim$100 MeV. We find this acceleration time to be a few seconds in the observer frame; this calculation took into account radiation losses suffered during the acceleration process. This result offers a straightforward explanation as to why, for most {\\it Fermi} GRBs, 100 MeV photons are not observed right at the trigger time but a little later. This also explains, why 100 MeV photons are observed before GeV radiation: it takes electrons radiating at GeV energies even longer time to accelerate. Taking this acceleration time into consideration while calculating high-energy light curves, we find that the light curve rises very rapidly -- much faster than it does for the external forward shock model with instantaneous electron acceleration for which the flux rises as $t^3$ when the CSM has uniform density (the $t^3$ rise reflects the increasing number of swept-up electrons before the blast wave decelerates). The detection of the first 100 MeV photons at some fraction of the deceleration time, the longer delays in the detection of higher energy photons\\footnote{Note that this possible trend in the data goes in the opposite direction than in the prompt $\\sim 1$ MeV emission, where higher energy photons arrive {\\it earlier} than lower energy photons in long GRBs and there is no lag detected for short GRBs (Norris et al. 1986, Norris \\& Bonnell 2006).} and the fast rise of the 100 MeV light curve, follow the expectation of the external forward shock model when the finite time for electron acceleration is taken into account. Detection of synchrotron photons of different energies provides an upper limit for the radiation flux produced in the reverse shock heated GRB-jet. For instance, the peak flux for the external reverse shock emission --- if the peak of the spectrum is at a few eV --- couldn't have been larger than about 300 mJy close to the deceleration time, for GRB080916C, otherwise it would prevent electrons from accelerating to a Lorentz factor of $\\sim10^7$ so that they can produce synchrotron photons of 100 MeV energy at early times (see Table 2). Similarly, the reverse-shock flux should be $\\lae$ 20 mJy for GRB 080916C in order that electrons in the forward shock are accelerated to a LF so that they produce 1 GeV photons. We speculate that the lack of $>$100MeV emission during the prompt phase of GRBs might be due to the presence of a bright optical source with observed flux larger than about 100 mJy, which would prevent electrons from reaching high Lorentz factors. This, coupled with the fact that GRBs with the largest LFs, which have small deceleration time, are the most likely bursts to be detected by {\\it Fermi} (Kumar \\& Barniol Duran 2009) might explain the detection/non-detection of $>100$MeV radiation from GRBs. We note that the shock-compressed magnetic field scenario requires some cross-field diffusion of particles - presumably generated by turbulence - to allow them to travel back to the upstream (e.g. Achterberg et al. 2001, Lemoine et al. 2006). This turbulent layer probably occupies a small fraction of the downstream region as suggested by recent simulations by Sironi \\& Spitkovsky (2010). Therefore, the picture that seems to emerge from numerical simulations and {\\it Fermi} observations, is that there might be a small region of turbulence behind the shock front that aids in the acceleration of particles across the shock, but that the radiation is mainly produced by particles that are swept downstream where the value of the downstream field is consistent with simple shock-compression of upstream field. There exists also the possibility that the CSM seed field is actually a few $\\mu$G and some instability produced ahead of the shock amplifies it to the value of a few tens of $\\mu$G we infer by our modeling of {\\it Fermi} GRBs (Kumar \\& Barniol Duran 2009, 2010). These instabilities have been studied by, e.g., Milosavljevi\\'c \\& Nakar 2006, Sironi \\& Goodman 2007, Goodman \\& MacFadyen 2008, Couch et al. 2008. However, this {\\it possible} amplification of a factor of $\\sim$10 is much smaller than the amplification customarily invoked to explain afterglow observations. We received a preprint from Piran \\& Nakar (2010) soon after this paper was completed. They have also considered the acceleration of electrons in the external shock." }, "1003/1003.5850_arXiv.txt": { "abstract": "\\noindent It has been known for nearly three decades that the energy spectra of thermonuclear X-ray bursts are often well-fit by Planck functions with temperatures so high that they imply a super-Eddington radiative flux at the emitting surface, even during portions of bursts when there is no evidence of photospheric radius expansion. This apparent inconsistency is usually set aside by assuming that the flux is actually sub-Eddington and that the fitted temperature is so high because the spectrum has been distorted by the energy-dependent opacity of the atmosphere. Here we show that the spectra predicted by currently available conventional atmosphere models appear incompatible with the highest-precision measurements of burst spectra made using the \\textit{Rossi X-ray Timing Explorer}, such as during the \\mbox{4U~1820$-$30} superburst and a long burst from \\mbox{GX~17$+$2}. In contrast, these measurements are well-fit by Bose-Einstein spectra with high temperatures and modest chemical potentials. Such spectra are very similar to Planck spectra. They imply surface radiative fluxes more than a factor of three larger than the Eddington flux. We find that segments of many other bursts from many sources are well-fit by similar Bose-Einstein spectra, suggesting that the radiative flux at the emitting surface also exceeds the Eddington flux during these segments. We suggest that burst spectra can closely approximate Bose-Einstein spectra and have fluxes that exceed the Eddington flux because they are formed by Comptonization in an extended, low-density radiating gas supported by the outward radiation force and confined by a tangled magnetic field. ", "introduction": "\\label{sec:intro} Type~I X-ray bursts (hereafter bursts) are produced by thermonuclear burning of matter accumulated in the surface layers of accreting neutron stars \\citep{woos76,joss77,lamb78}. These bursts have rise times ranging from a fraction of a second to a few tens of seconds, durations ranging from about ten seconds to several thousand seconds, recurrence times $\\sim\\,$10$^3$--10$^6$~s, peak luminosities $\\sim\\,$10$^{38}$~ergs~s$^{-1}$, and total energy releases $\\sim\\,$10$^{39}$--10$^{42}$~ergs \\citep{stro06}. The observed X-ray flux typically increases by a factor of $\\sim\\,$10--100 during a burst. The properties of the large number of bursts that have been observed using the \\textit{Rossi X-ray Timing Explorer} (\\textit{RXTE}) have recently been summarized by \\citet{gall08}. Planck (blackbody) functions are often fit to the energy spectra of bursts \\citep{swan77, hoff77, gall08}. During some, the temperature obtained from such fits drops and the derived emitting area increases. These photospheric radius expansion (PRE) bursts are thought to occur when the radiative flux through the stellar atmosphere exceeds the Eddington critical flux, creating an optically thick wind (see \\citealt{gall08}). The radiative flux is greater than the Eddington flux for any realistic neutron star if the emission has a Planck spectrum with a temperature measured at infinity $kT_\\infty>2.0$~keV (see \\citealt{mars82} and Section~\\ref{sec:max-temp}). Yet fits of Planck functions to burst spectra frequently yield temperatures substantially higher than this expected maximum, even during times when there is no evidence of radius expansion. Neutron stars are not blackbodies, and conventional model atmosphere calculations show that they generally do not produce Planck spectra (see, e.g., \\citealt{lond84, lond86, made04, majc05}). In conventional atmospheres, energy-dependent absorption and scattering cause the spectrum to peak at an energy higher than the peak of a Planck spectrum with the same effective temperature. This effect led to widespread acceptance of the hypothesis that the effective temperature is substantially smaller than the temperature obtained by fitting a Planck function to the burst spectrum and that the radiative flux is sub-Eddington even when the fitted temperature exceeds 2.0~keV (see, e.g., \\citealt{ebis84} and \\citealt{gall08}, section 2.2). In contrast to conventional neutron star atmospheres, low-density atmospheres extensive enough to fully Comptonize free-free and cyclotron photons will produce Bose-Einstein spectra $dN/dE\\propto E^2/[\\exp((E-\\mu)/kT)-1]$ with chemical potentials $\\mu$ that satisfy $|\\mu| \\ll kT$ (see, e.g., \\citealt{illa75}). These spectra have almost the same shape and energy flux as a Planck spectrum with the same temperature, because a Planck spectrum is a Bose-Einstein spectrum with $\\mu=0$. An important aspect of Bose-Einstein spectra is that knowledge of the radiation temperature and the chemical potential is sufficient to determine the radiative flux from the emitting surface; knowledge of the distance to the source or its luminosity is unnecessary. Here we report analyses of \\textit{RXTE} data taken during high-temperature segments of a superburst from \\mbox{4U~1820$-$30} and a long burst from \\mbox{GX~17$+$2}. Such segments provide the best opportunity to test spectral models, because the large number of counts collected allows the spectrum to be measured with exceptionally high precision. We find that the spectra predicted by currently available conventional atmosphere models appear incompatible with the spectra during these segments, whereas Bose-Einstein spectra fit these spectra well. The fits give $|\\mu| \\la kT$ and values of $kT$ substantially greater than 2.0~keV, implying radiative fluxes at the emitting surface more than a factor of three larger than the Eddington flux. There is no evidence that the emitting surface is expanded at these times. We find that the spectra of other bursts from \\mbox{4U~1820$-$30} and \\mbox{GX~17$+$2} and bursts from many other bursters are well-fit by similar Bose-Einstein spectra, suggesting that the radiative flux also exceeds the Eddington flux during these bursts. ", "conclusions": "\\label{sec:discussion-conclusions} As discussed in Section~2, measurement of a Bose-Einstein spectrum with a temperature greater than \\mbox{$\\sim\\,$2~keV} implies that the radiative flux at the emitting surface exceeds the Eddington flux, independent of unknowns such as the distance to the source, the radiating area on the star, the radius of the star, and its surface redshift. The implied fluxes are accurate, because the 2--60~keV bandpass of the PCA captures more than 95\\% of the flux of a 3.0~keV Bose-Einstein spectrum with $|\\mu| \\ll kT$. We have found that intervals with temperatures greater than $\\sim\\,$2~keV occur during most bursts, suggesting that the radiative flux exceeds the Eddington flux during most bursts. When combined with the flux profiles seen during PRE bursts from some of these same stars, these results, and the small effective areas inferred during high-temperature intervals, indicate that most of the emission during these intervals comes from only a fraction (in some cases $\\sim\\,$20\\%) of the stellar surface. These high temperatures and fluxes and small emitting areas raise several important questions: How can the flux be super-Eddington without producing a significant wind? What determines the maximum flux from the emitting area, and how big is it? And how do these results fit with evidence that the emitting surface is sometimes far above the stellar surface? We suggest that the radiative flux can exceed the Eddington flux because the emitting gas is confined by a tangled stellar magnetic field. The sudden nuclear energy release that produces a burst creates a zone of turbulent convection at densities $\\sim\\,$ 10$^{5-7}$~g~cm$^{-3}$ (see, e.g., \\citealt{fush87}). The convective energy flux is $F_t = \\rho_t u_t^3$, where $\\rho_t$ is the density in the convection zone and $u_t$ is the turbulent velocity there. The convection will amplify and tangle the star's weak poloidal magnetic field until the tangled field $B_t$ becomes strong enough to inhibit convection, which occurs when $B_t^2/8\\pi \\approx \\rho_t u_t^2$. The maximum value of $B_t$ is $\\approx (8\\pi)^{1/2} \\rho_t^{1/6} F_t^{1/3}$ and is relatively insensitive to $\\rho_t$ and $F_t$. For typical densities in the convection zone and the highest energy fluxes observed from the emitting surface, which are $\\sim\\,$10$^{26}$~erg~cm$^{-2}$~s$^{-1}$, $B_t({\\rm max})$ is $\\sim{\\rm few}\\times 10^{10}$~G, $\\sim\\,$10--100 times stronger than the dipole components inferred from observations and theoretical modeling (see \\citealt{lamb08}). The tangled field will be strong enough to confine the emitting gas if its tension, $f_{\\rm mag} \\approx (1/4\\pi)(B_t\\cdot\\nabla B_t) \\approx B_t^2/4\\pi \\ell_B$, exceeds the outward radiation force, $f_{\\rm rad} \\approx (F_{\\rm rad}/c) n_e\\sigma$. Here $\\ell_B$ is the characteristic scale of the tangled field and $n_e$ is the electron density in the radiative zone. Assuming $\\ell_B$ is no larger than the depth $\\sim\\,$10$^3$~cm of the burning zone, a field $\\sim\\,$$B_t({\\rm max})$ can confine the atmosphere in the presence of a radiative flux $\\approx 10^{26}$~erg~cm$^{-2}$~s$^{-1}$, which is the flux implied by an effective temperature $\\approx 3$~keV. A neutron star atmosphere supported by a super-Eddington radiative flux and confined by magnetic stresses is likely to be more extended and have a lower density than a conventional atmosphere supported by gas pressure and confined by gravity. In a future paper (Lamb et al., in preparation), we show that such an atmosphere naturally produces a Bose-Einstein photon spectrum with $|\\mu| \\la kT$. Comptonization by the electrons in the atmosphere drives the photon distribution close to a Bose-Einstein distribution while weak free-free and cyclotron emission drive the chemical potential to a small value (see, e.g., \\citealt{illa75}). We expect that a region of very hot, confined gas will heat adjacent areas of the stellar surface, which may not be confined by a strong magnetic field. When these adjacent areas become hot enough, they will expand vertically. If the product of the radiative flux from the very hot, confined gas and its emitting area exceeds the Eddington luminosity, adjacent gas will leave the star as a wind, producing a PRE event. Hence the maximum radiative luminosity will be approximately the Eddington luminosity, just as in the conventional picture, even though heat is flowing from below the atmosphere over only a fraction of the stellar surface. The PRE will end when the \\textit{luminosity} of the very hot, confined gas falls below the Eddington luminosity, even if the local \\textit{flux} from this gas exceeds the Eddington flux. The high temperatures and radiative fluxes and small emitting areas found here, which were first noted nearly three decades ago, have important implications for efforts to determine neutron star masses and radii using bursts. For example, it is often assumed that during high-temperature intervals the entire stellar surface emits exactly the Eddington flux. Our analysis of spectral measurements made using \\textit{RXTE} shows that these assumptions must be reconsidered." }, "1003/1003.0472_arXiv.txt": { "abstract": "Using three-dimensional cosmological simulations, we study the assembly process of one of the first galaxies, with a total mass of $\\sim 10^8~M_{\\odot}$, collapsing at $z\\simeq 10$. Our main goal is to trace the transport of the heavy chemical elements produced and dispersed by a pair-instability supernova exploding in one of the minihalo progenitors. To this extent, we incorporate an efficient algorithm into our smoothed particle hydrodynamics code which approximately models turbulent mixing as a diffusion process. We study this mixing with and without the radiative feedback from Population~III (Pop~III) stars that subsequently form in neighboring minihalos. Our simulations allow us to constrain the initial conditions for second-generation star formation, within the first galaxy itself, and inside of minihalos that virialize after the supernova explosion. We find that most minihalos remain unscathed by ionizing radiation or the supernova remnant, while some are substantially photoheated and enriched to supercritical levels, likely resulting in the formation of low-mass Pop~III or even Population~II (Pop~II) stars. At the center of the newly formed galaxy, $\\sim 10^5~M_{\\odot}$ of cold, dense gas uniformly enriched to $\\sim 10^{-3}~Z_{\\odot}$ are in a state of collapse, suggesting that a cluster of Pop~II stars will form. The first galaxies, as may be detected by the {\\it James Webb Space Telescope}, would therefore already contain stellar populations familiar from lower redshifts.\\\\ ", "introduction": "One of the most important goals in modern astrophysics is to understand the end of the cosmic dark ages, when the first stars and galaxies transformed the simple early universe into a state of ever increasing complexity \\citep{bl04a, cf05, glover05, bl07, bromm09}. The first stars, the so-called Population~III (Pop~III), were the source of hydrogen-ionizing UV photons, thus initiating the extended process of cosmic reionization \\citep{fck06}. They also synthesized the first heavy chemical elements, beyond the hydrogen and helium produced in the big bang, to be dispersed into the pristine intergalactic medium (IGM) through supernova (SN) explosions and winds \\citep[e.g.,][]{mfr01,wa08b}. An intriguing possibility for the first stars is that some of them may have died as a pair-instability supernova (PISN), a peculiar fate predicted for progenitor masses in the range $140~M_{\\odot}\\la M_{*}\\la 260~M_{\\odot}$ \\citep{brs67,hw02}. Current theory proposes a top-heavy initial mass function (IMF) for the first stars, with a characteristic mass $M_{*}\\ga 100~M_{\\odot}$ \\citep{abn02,bcl02,on07,yoh08}. Together with the expectation that mass loss due to radiatively--driven winds is negligible at low metallicities \\citep{kudritzki02}, one arrives at the robust expectation that at least a fraction of Pop~III stars should have died as PISNe. Compared to conventional core-collapse SNe, a PISN is distinguished by not leaving behind a compact remnant. Instead, the exploding star is completely disrupted, and {\\it all} metals produced inside the Pop~III star are released into the surroundings, leading to metal yields of $y=M_{\\rm Z}/M_{*}\\sim 0.5$ \\citep{hw02}. An abundant occurrence of PISNe in the early universe could thus have rapidly established a bedrock of metals, at least locally, of order $Z\\ga 10^{-3}~Z_{\\odot}$ \\citep{greif07}. Recently, the extremely luminous SN~2007bi was observed in a nearby galaxy \\citep{gal-yam09}, with characteristics, such as very large Ni masses, that seem to unambiguously point to a PISN origin. This discovery greatly strengthens the possibility for finding these events at high redshifts as well. We here carry out cosmological simulations tracing the detailed assembly process of a primordial galaxy, taking into account the feedback effects from Pop~III star formation inside the minihalo progenitor systems. Our work extends the study by \\citet{greif08}, which followed the virialization of gas in the galactic potential well under the idealized assumption of no such feedback. Specifically, we include radiative feedback, leading to the build-up of H~{\\sc ii} regions around Pop~III stars \\citep[e.g.,][]{wan04,abs06,yoshida07,greif09b}, as well as the mechanical and chemical feedback from a single PISN that explodes in the earliest minihalo progenitor. The latter feedback refers to the additional cooling that becomes available in metal-enriched gas, allowing the gas to reach lower temperatures, and to possibly fragment into low-mass Population~II (Pop~II) stars. It has been suggested that this transition in the star formation mode, from high-mass Pop~III to normal-IMF Pop~II, occurred once a minimum metal enrichment was in place, the so-called critical metallicity $Z_{\\rm crit}\\sim 10^{-6}-10^{-3.5}~Z_{\\odot}$ \\citep{bl03a,schneider06}. It is therefore important to identify any ``super-critical'' regions within a simulation box to distinguish between Pop~III and Pop~II star formation sites. The overall goal is to predict the properties of the first galaxies, to be observed with the {\\it James Webb Space Telescope (JWST)}, planned for launch in $2014$ \\citep{gardner06}. One key ingredient necessary for such predictions is to quantify the amount and distribution of metals inside the emerging galaxy, just prior to the onset of its initial starburst. This would allow us to constrain the detailed mix of stellar populations, and consequently to arrive at observables such as broad--band colors, emission line signatures, and luminosities \\citep[e.g.,][]{johnson09}. A complementary empirical probe of the first stars is given by ``Stellar Archaeology'', the study of abundance patterns in metal-poor Galactic halo stars \\citep{bc05,fjb07,fjb09}, and of globular clusters in the Milky Way and other galaxies \\citep[e.g.,][]{helmi06,pkg06}. To infer the characteristics of the first SNe, and therefore to constrain the primordial IMF from the observed fossil chemical record, we need to better understand the in situ physical conditions in the regions where the first low-mass stars formed. This again requires us to simulate the transport of metals from the first SNe in a realistic cosmological setting. A related problem is the apparent absence of a PISN abundance pattern in any of the known metal-poor stars \\citep{tumlinson06}. It has been argued that such a PISN signature may be hidden in stars with relatively high metallicities, $Z\\sim 10^{-2.5}~Z_{\\odot}$, due to the extremely high yield from even a single explosion \\citep{kjb08}. Our simulations can assess the viability of this scenario with much improved physical realism. We trace the mixing and dispersal of metals, originating in a PISN inside a Pop~III minihalo at $z\\simeq 30$, all the way to their reassembly into the growing potential well of the first galaxy at $z\\simeq 10$. The physics governing this transport is highly complex. Some aspects, such as the advection of metals during hierarchical structure formation, is reliably modeled by our Lagrangian smoothed-particle hydrodynamics (SPH) approach. Other processes, such as hydrodynamical instabilities and turbulence, are not well resolved in our simulation. We may therefore miss the fragmentation of gas within the dense shell of the SN remnant, which could result in the formation of an intermediate generation of stars \\citep{mbh03,on08,whalen08b,nho09}. However, we have developed an efficient algorithm to approximately model the mixing that results from these sub-grid effects as a diffusion process \\citep{greif09a}. The corresponding diffusion coefficient is evaluated assuming a simple mixing-length approach to turbulent transport even in the supersonic regime \\citep{kl03}, with quantities that are entirely local. This allows us to greatly improve on our earlier, ``ballistic'' transport of discrete metal packets \\citep{greif07}, and consequently to derive meaningful metallicity distribution functions. The structure of our work is as follows. In Section~2, we describe our numerical methodology and the procedure to set up the initial conditions. We then discuss the assembly process of the first galaxy, and the resulting distribution of metals within our simulation volume (Section~3). In Section~4, we investigate the gas properties in the different sites where second-generation star formation could occur. We conclude by summarizing and assessing the cosmological implications of our results. All distances quoted in this paper are in proper units, unless noted otherwise. ", "conclusions": "We have performed a set of highly resolved SPH simulations that allow us to investigate the enrichment of the IGM by a PISN exploding in a high-redshift minihalo. We have incorporated a substantially higher degree of realism compared to our previous work in the form of radiation feedback, explicit chemical mixing, and metal line cooling. In particular, we have employed a physically motivated model for the mixing of metals between individual SPH particles based on diffusion \\citep{greif09a}. We have followed the distribution of metals as they are ejected into the IGM and then recollapse into the larger, $M_{\\rm vir}\\sim 10^{8}~M_{\\odot}$ potential well of a ``first galaxy'' assembling at $z\\simeq 10$. We have performed simulations with and without radiative feedback to assess the effects of photoheating on the assembly of the galaxy and the distribution of metals. The heavy elements ejected by the SN are initially distributed into the IGM by the bulk motion of the SN remnant, until it stalls and mixing is instead facilitated by turbulent motions on smaller scales. These are induced by the dynamics of the underlying DM, and, to a smaller degree, by the additional photoheating. A clear correlation between metallicity and gas overdensity is established, where the densest regions consisting of existing minihalos remain largely pristine, while voids around the SN progenitor become highly enriched. This correlation breaks down once the potential well of the galaxy assembles and metal-rich gas residing in the IGM recollapses to high densities. The metallicity at the center of the galaxy then grows to $Z\\sim 10^{-3}~Z_{\\odot}$, likely resulting in the formation of a stellar cluster with a more normal IMF \\citep{cgk08}. Although the influence of photoheating on the mixing of the gas is limited, an interesting side effect is that it can expel enriched gas out of the potential well of the galaxy into the IGM. Finally, the mechanical feedback exerted by the photoheating and the propagation of the SN remnant is usually quite small, although in some cases nearby minihalos are disrupted and enriched, possibly leading to the formation of Pop~III.2 and Pop~II stars. One of the most promising observational tools to understand the nature of these objects relies on the elemental composition of second-generation stars. Previous searches for PISN signatures have targeted the most metal-poor stars in our Galaxy, with an iron abundance as low as $[{\\rm Fe}/{\\rm H}]\\sim 10^{-5}$ \\citep[e.g.,][]{bc05,frebel05}. Unfortunately, the spectra of these stars do not reveal the strong odd--even effect indicative of PISNe, implying that they were either not very common, or that existing surveys did not target a suitable sample of stars. Indeed, our simulations show that a single PISN is sufficient to enrich a gas cloud quite uniformly to $Z\\sim 10^{-3}~Z_{\\odot}$, implying that its signature may be buried in stars with a significantly higher metallicity \\citep{kjb08}. The ideal search strategy may therefore have to be modified to concentrate on stars with $[{\\rm Fe}/{\\rm H}]\\sim 10^{-3}$. Recent surveys of extragalactic globular clusters lend support to this idea, since the oldest of these have been found to contain large $[\\alpha /{\\rm Fe}]$ ratios, which is indicative of enrichment by PISNe \\citep{pkg06}. The standard picture that minihalos only form Pop~III stars, or remain sterile before being disrupted during subsequent merging \\citep[e.g.,][]{hh03}, may also have to be modified. Whereas that is indeed the case for the majority of minihalos in our simulation, there exists a small number of halos that do contain dense, metal-enriched gas. Although we do not follow its subsequent evolution, it is intriguing to speculate whether these clouds might fragment to form the first globular clusters, or at least slightly less massive precursors \\citep[e.g.,][]{bc02,boley09}. Soon, it should become possible to study the further fate of these second-generation star forming clouds with the required extremely high resolution by resimulating the central halo gas, and by including dust cooling and opacity effects at high density. Similar simulations will be able to address the formation of the first star clusters inside the first galaxies, starting from the initial conditions determined here. The goal of understanding galaxy formation from first principles, at least in the relatively simple case of dwarf-sized systems at the beginning of hierarchical structure formation, might finally have come into reach." }, "1003/1003.4584_arXiv.txt": { "abstract": "We discuss new exact solutions of a three-parameter nonminimal Einstein-Maxwell model. The solutions describe static spherically symmetric objects with and without center, supported by an electric field nonminimally coupled to gravity. We focus on a unique one-parameter model, which admits an exact solution for a traversable electrically charged wormhole connecting two universes, one asymptotically flat the other asymptotically de Sitter ones. The relation between the asymptotic mass and charge of the wormhole and its throat radius is analyzed. The wormhole solution found is thus a nonminimal realization of Wheeler's idea about charge without charge and shows that, if the world is somehow nonminimal in the coupling of gravity to electromagnetism, then wormhole appearance, or perhaps construction, is possible. ", "introduction": "The wormhole concept was invented by Wheeler (see, e.g., \\cite{Wheeler}) to provide a mechanism for having charge without charge, since in a such a spacetime without a center, the field lines seen in one part of the Universe could thread the throat and reappear in other part. The idea, was further extended by Morris and Thorne to allow, not only field lines, but also observers to travel through the throat \\cite{MorTho}. By having invoked an arbitrarily technologically advanced civilization, this work \\cite{MorTho} initiated the engineering of wormhole construction, theoretically, and the study of wormholes as topological bridges joining two different spacetimes has since then attracted extraordinary attention in this modern context (see, e.g., \\cite{VisserBook} and references therein). The main feature of traversable wormhole physics is the fact that the matter threading the wormhole throat should possess exotic properties, one of them being the violation of the null energy condition \\cite{HocVis}. To provide the existence of wormholes one should either include some hypothetical forms of matter into the model, or introduce interactions of a new type. Various models, admitting the required violation of the null energy condition, have been considered in the literature, among them thin shells in a cosmological constant background \\cite{LambdaWH1,LambdaWH2}, scalar fields \\cite{scalarfields}, wormhole solutions in semiclassical gravity \\cite{semiclas}, solutions in Brans-Dicke theory \\cite{Nan-etal}, wormholes on the brane \\cite{wormholeonbrane}, wormholes supported by matter with an exotic equation of state, namely, phantom energy \\cite{phantom}, the generalized Chaplygin gas \\cite{chaplygin}, tachyon matter \\cite{tachyon}, nonlinear electrodynamics \\cite{AreLob}, and other cases. Now, when one considers nonminimal coupling of gravity with vector-type fields, such as Maxwell, Yang-Mills or Proca fields, new interesting possibilities appear. Nonminimal phenomena, i.e., phenomena induced by the interaction between curvature, or gravity, and other fields can be characterized on the one hand by unusual effective energy conditions, and on the other hand, allow one to exclude exotic substrates. In \\cite{BaSuZa} an exact solution of the nonminimal Einstein--Yang-Mills model was obtained which demonstrates that an $SU(2)$ symmetric gauge field nonminimally coupled to curvature can support a traversable wormhole. This is a nonminimal Wu-Yang magnetic wormhole, and the throat joins two asymptotically flat regions. The corresponding spacetime has no center and this model could be an illustration of Wheeler's idea about ``charge without charge'' \\cite{Wheeler}. Following this idea we now intend to consider an electrically charged object in the context of a spacetime without a center. A few wormhole models are known in which the electric charge is spread on the spherical shell \\cite{thinshell}, or the throat is filled with some nonminimal and ghostlike scalar field \\cite{charge}. Here, we find an exact solution of a nonminimal traversable wormhole, which contains neither electric charge on spherical shells, either thin or thick, nor scalar fields, nevertheless, possesses a static spherically symmetric electric field which is charged from the point of view of a distant observer. Thus, our goal is twofold. We present explicitly a nonminimal realization of Wheeler's idea about charge without charge, and we show that, if the world is somehow nonminimal in the coupling of gravity to electromagnetism, then wormhole appearance, or perhaps construction by an absurdly advanced civilization \\cite{LambdaWH1}, is possible. For this purpose we address a nonminimal Einstein-Maxwell theory, which has been earlier elaborated in detail in both linear (see, e.g., \\cite{FaraR,HehlObukhov}) and nonlinear (see \\cite{BL05}) versions. The model linear in the spacetime curvature and quadratic in the Maxwell tensor \\cite{BL05} contains three nonminimal coupling constants $q_1$, $q_2$, and $q_3$. These quantities have the dimensionality of area and characterize the cross terms in the Lagrangian linking the Maxwell field $F_{ik}$ and terms linear in the Ricci scalar $R$, Ricci tensor $R_{ik}$, and Riemann tensor $R_{ikmn}$, respectively. These parameters are a priori free ones but can acquire specific values in certain effective field theories. The first example of a calculation of the three coupling parameters was based on one-loop corrections to quantum electrodynamics in curved spacetime, a direct and nonphenomenological approach considered by Drummond and Hathrell \\cite{Drum}. In another instance, Buchdahl \\cite{Buchdahl} and then M\\\"uller-Hoissen \\cite{MH} obtained a nonminimal Einstein-Maxwell model from dimensional reduction of the Gauss-Bonnet action. This model contains one coupling parameter. Based on these works, and specially on \\cite{BL05}, solutions, which described nonminimal electrically charged objects with and without centers, characterized by two sets of relations for nonminimal coupling parameters, namely, $q_1+q_2+q_3=0$, $2q_1+q_2=0$, and $q_1+q_2=0$, $q_3=0$, were obtained \\cite{BBL08}. However, none of these solutions can be used to construct traversable wormholes. In this paper we formulate a new nonminimal model, for which the coupling parameters satisfy the relations, $3q_1+q_2=0$, $q_3=0$. Exact solutions of this model are shown to admit the existence of nonminimal traversable electrically charged wormholes. Comparing these new results with the ones obtained in \\cite{BaSuZa} we would like to emphasize the following features. First, here we deal with an electric field instead of a magnetic one; second, in \\cite{BaSuZa} we used the basic conditions $12q_1+4q_2+q_3=0$, $q_3\\neq0$; third, the spacetime is now nonsymmetric, i.e., the throat joins one asymptotically flat region to another asymptotically de Sitter region. In \\cite{BronMD} the authors have suggested the name ``black universes'' to nonsymmetric wormhole spacetime configurations for which the first asymptotics is flat and the second one is of cosmological type. The paper is organized as follows: In Section \\ref{main}, we quote the details of the three-parameter nonminimal Einstein-Maxwell model and formulate the corresponding key equations for the electric and gravitational fields. Furthermore, we introduce a new one-parameter nonminimal model, given by the conditions $3q_1+q_2=0$, $q_3=0$, derive a key (cubic) equation for the electric field $E$ and obtain the metric functions $\\sigma$, $N$ in terms of $E$. In Section III, we discuss three solutions of the equation for the electric field. We focus on the Coulombian-type solution: This is the solution to which the others should be compared. In Section IV, we obtain exact solutions with a center. In Section V, we thoroughly discuss wormhole solutions. Section VI contains the conclusions. ", "conclusions": "\\label{V} We have presented new exact solutions of a nonminimal Einstein-Maxwell model with coupling parameters satisfying the relations $q_1= -q$, $q_2=3q$, $q_3=0$, with $q$ arbitrary, which describe spherically symmetric electrically charged objects with and without center. The electric field of the objects is characterized by the following features: {\\it (i)} it satisfies a cubic key equation and splits into three branches, one of them has Coulombian asymptotics, $E(r) \\to Q/r^{2}$; {\\it (ii)} when the guiding nonminimal dimensionless parameter $a=4q/\\kappa Q^2$ is positive, the electric field is described by a smooth function finite everywhere; {\\it (iii)} when $a \\geq 256/243$ the Coulombian branch of the electric field conjugates with the non-Coulombian branch with constant asymptotics $E=1/\\sqrt{\\kappa q}$, the position of the junction point depends on the parameter $a$ and coincides with geometrical center for $a=256/243$. There are wormhole solutions. Indeed, when the nonminimal guiding parameter $a$ is in the interval $1.0720$ vs $\\Delta m^2_{31}<0$, and acquire information about $\\theta_{13}$. This method can be applied to neutrino detectors of various types. Examples are the liquid scintillator detector of the Daya Bay experiment \\cite{Guo2} under construction, the decommissioned heavy water detector of the previous Sudbury Neutrino Observatory (SNO) \\cite{SNO1}, etc. The other main purpose of the present work is to propose a method to obtain some information about neutrino masses from SN neutrinos. Since neutrinos have small masses, their transmit velocities are smaller than that of light. Therefore, a massive neutrino will have an energy and mass dependent delay relative to a massless neutrino after traveling over a long distance \\cite{Beacom1}. We will define another ratio, for the neutral-current reactions, as the event number of the delayed SN neutrinos over the total event number. With the relation between this ratio and neutrino masses, we will propose a possible method to acquire information about neutrino masses. This method will also be applied to Daya Bay and SNO for examples. The paper is organized as follows. In Sec. II, we review the detection of SN neutrinos on the Earth. In Sec. III, we propose a possible method to identify the mass hierarchy and acquire information about $\\theta_{13}$ and apply the method to the Daya Bay and SNO experiments. In Sec. IV, a method to acquire information about neutrino masses is given and also applied to Daya Bay and SNO. A summary is given in section V. ", "conclusions": "" }, "1003/1003.3647_arXiv.txt": { "abstract": "The origin of the extragalactic gamma-ray background is a pressing cosmological mystery. The {\\it Fermi} Gamma-Ray Space Telescope has recently measured the intensity and spectrum of this background; both are substantially different from previous measurements. We present a novel calculation of the gamma-ray background from normal star-forming galaxies. Contrary to longstanding expectations, we find that numerous but individually faint normal galaxies may comprise the bulk of the {\\it Fermi} signal, rather than rare but intrinsically bright active galaxies. This result has wide-ranging implications, including: the possibility to probe the cosmic star-formation history with gamma rays; the ability to infer the cosmological evolution of cosmic rays and galactic magnetic fields; and an increased likelihood to identify subdominant components from rare sources (e.g., dark matter clumps) through their large anisotropy. ", "introduction": "The {\\em Fermi} Gamma-Ray Space Telescope has unveiled the high-energy cosmos with unprecedented clarity and depth. The gamma-ray sky has been known \\citep[e.g.,][]{hunter,sreek} to be dominated by diffuse emission from the Galactic plane, while at high Galactic latitudes a diffuse extragalactic gamma-ray background (EGB) has an important, and at some energies dominant, contribution. However, before {\\em Fermi}, the processes dominating the diffuse emission from the Galaxy, especially above 1 GeV, were unclear and highly debated--cf.~discussion on the {\\em GeV excess} reported by the Energetic Gamma-ray Experiment Telescope (EGRET) \\citep[e.g.,][and references therein]{smr00}. {\\em Fermi} has clarified \\citep{FermiMidLat09} that the dominant emission mechanism is cosmic-ray interactions with interstellar gas, which leads to gamma rays mostly from pion decay in flight, i.e., $p_{\\rm cr} + p_{\\rm ism} \\rightarrow p p \\pi^0$ then $\\pi^0 \\rightarrow \\gamma \\gamma$ \\citep{stecker}. Moreover, the {\\em Fermi} EGB differs from previous EGRET estimates: the intensity is {\\em fainter} and the spectrum {\\em steeper}, consistent with a power law of spectral index $2.41 \\pm 0.05$ and integrated intensity $I(>100{\\rm MeV})=(1.03\\pm0.17)\\times10^{-5}{\\rm cm^{-2}s^{-1}sr^{-1}}$ \\citep{fermi-EGB}. Our theoretical understanding of the EGB must therefore be substantially revised in light of the new and smaller {\\em Fermi} signal. Pioneering studies investigating the origin of the EGB first considered the collective emission form star-forming galaxies (like the Milky Way), but found this to give a small EGB signal \\citep{sww,bignami}. \\citet{pf02} first incorporated observations of the cosmic star-formation rate, while ~\\citet{pf06} made the first estimates of the pionic contribution from star-forming galaxies. In both cases, the predicted intensity was below the then-measured EGB. Blazars, the brightest extragalactic sources, have been the favored candidates \\citep{ss96}. However, subsequent estimates of their contribution to the EGB \\citep[e.g.,][]{mucke,cm98,nt07,dermer07} have consistently fallen short. {\\em Fermi} point-source observations suggest that unresolved blazars contribute at most $\\sim 23\\%$ of the EGB; and thus the mystery has become acute \\citep{fermicounts}. Here we present a more realistic model for the EGB from star-forming galaxies, constructed to use as much as possible of our substantially improved multiwavelength observational understanding of these sources, and isolating the signal from normal star formation (as opposed to starburst galaxies). ", "conclusions": "Our full numerical calculation uses a Milky Way pionic source spectrum whose {\\em shape} is derived from \\citet{pfrommer}, calibrated to observations by normalizing the $> 100$ MeV photon emission per hydrogen atom to the {\\em Fermi} result at intermediate Galactic latitudes \\citep{FermiMW09}. The cosmic SFR is from ~\\citet{horiuchi}. For the Milky Way SFR, used to normalize the cosmic-ray flux/SFR ratio, we use the recent estimate of \\citet{robitaille} ($\\psi_{\\rm MW} = 1 {M_\\odot/\\rm yr}$, a factor of 3 lower than earlier work). Figure~1 shows our results for the normal galaxy contribution to the EGB. We plot predictions for the limiting cases of pure luminosity and of pure density evolution. The uncertainties in the model inputs, summed in quadrature, propagate into the displayed error band that applies to each curve, which we estimate to be a factor of $10^{\\pm 0.3}$, resulting from uncertainties of: 30\\% in pionic emissivity \\citep{FermiMW09}, 40\\% in the normalization of the Galactic star-formation rate \\citep{robitaille}, 40\\% in the cosmic star-formation rates \\citep{horiuchi}, and 25\\% in the luminosity scaling in eq.~(\\ref{eq:mavg}). The true systematic uncertainty would also reflect the idealizations in our model (universal cosmic-ray spectra and confinement). These errors are hard to estimate but in any case imply that the uncertainty range in Figure~1 is a lower bound to the error budget. Within errors, our predictions for both limiting models fall at or below the level of the {\\em Fermi} data, where the data seem to support the pure luminosity evolution case that explains nearly the entire signal. Comparing central values, this model gives $\\approx 50\\%$ of the {\\em Fermi} EGB $\\la 10$ GeV. Thus, unresolved normal galaxies make a substantial and likely dominant contribution to the observed EGB, without overpredicting the signal. Even the pure density evolution case accounts for a minimum of 20\\% of the EGB around 0.3 GeV; this provides a {\\em lower} limit to the normal-galaxy signal. Thus, any {\\em other} EGB sources \\citep{ss96,dermer07,fermicounts,mspulsars} must contribute no more than the remaining $80\\%$ of the data. Indeed, the LAT team {\\em upper} limit to the blazar EGB contribution shown in Figure~1 is comparable to our {\\em lower} limit \\citep{fermicounts}. The spectral shapes of the two limiting cases are very similar: the peak in $\\eobs^2 dI/d\\eobs$ lies at $\\sim 0.3$ GeV because the bulk of the signal comes from $z \\sim 1$. These models predict that the EGB turns over for $\\eobs \\la 0.3$ GeV, a testable prediction of our model. For hadronic emission, the high-energy spectral index is the same as the underlying proton spectral index, here $s_{\\gamma,\\rm had} = s_{\\rm cr} = 2.75$; this is somewhat steeper than the {\\em Fermi} single-power-law fit $s_{\\rm obs} = 2.41 \\pm 0.05$. Consequently, our predictions at high energies ($\\ga 10$ GeV) fall below the data. If normal galaxies had a {\\em distribution} of cosmic-ray spectral indices, the resulting EGB spectrum would steepen at high energies where the hardest sources would dominate, developing a {\\em convex tail}. Indeed, the {\\em Fermi} EGB data suggests a slight flattening of slope around $E \\ga 10$ GeV, which might hint at such a transition. A galaxy with characteristic \\halpha\\ luminosity $\\tracer_*$ has $L_\\gamma^*(>100 {\\rm MeV}) = 1.4 \\times 10^{43} \\ {\\rm s^{-1}}$. Such objects have flux $F$ if they lie at distances $r_* = (L_\\gamma^*/4\\pi F)^{1/2} = 11 \\, {\\rm \\ Mpc} \\ ({10^{-9} \\ \\rm cm^{-2} \\ s^{-1}}/F)^{1/2}$. Thus {\\em Fermi} should eventually resolve \\beq N(>F) \\sim 4\\pi r_*^3 \\ n_*(0)/3 = 5 \\ \\pfrac{{10^{-9} \\ \\rm cm^{-2} \\ s^{-1}}}{F}^{3/2} \\eeq normal galaxies, consistent with 2--3 detections to date \\citep[the LMC, SMC, and perhaps M31;][]{fermi-LMC,fermi-SMC, fermi-M31}. Our results do not account for starburst galaxies, nor for inverse-Compton emission from any star-forming galaxies; these must contribute to the star-forming EGB, and could have hard spectra dominating $\\ga 10$ GeV. We have also neglected gamma-ray attenuation by extragalactic background light \\citep[important at $E \\ga 30$ GeV; e.g.,][and references therein]{sms}. We will address these issues in future work. \\begin{figure} \\epsscale{.80} \\includegraphics[angle=-90, width=\\linewidth]{SFEGB_normalgals.eps} \\caption{ The normal galaxy contribution to the extragalactic gamma-ray background. Curves represent two limiting cases of cosmic star formation: pure luminosity and pure density evolution. Colored error band illustrates the factor $10^{\\pm 0.3}$ uncertainty in the normalization of both theory curves. {\\em Fermi} data are from \\citep{fermi-EGB}. Dotted line: unresolved blazar EGB upper limit \\citep{fermicounts}. } \\end{figure} The amplitude and configuration of magnetic fields in a galaxy have an additional effect on the scaling of cosmic-ray flux with SFR. Confirmation that normal galaxies comprise the bulk of the Fermi signal would constitute a unique probe of the evolution of these magnetic fields between the redshift of peak star formation and today. Because of their ubiquity, normal galaxies produce the smallest anisotropies in the EGB, far less than blazars or other proposed sources. Thus, by studying the observed EGB anisotropy as a function of energy, it may be possible to disentangle the spectrum and amplitude of the normal-galaxy contribution from that of other sources \\citep{vaso09}. Moreover, because normal galaxies seem to dominate the {\\em Fermi} EGB, their small contribution to anisotropies will fortuitously provide the optimal chance of finding smaller and more exotic sources in the observed signal \\citep{sgp09}." }, "1003/1003.1368_arXiv.txt": { "abstract": "{With more and more exoplanets being detected, it is paid closer attention to whether there are lives outside solar system. We try to obtain habitable zones and the probability distribution of terrestrial planets in habitable zones around host stars. Using Eggleton's code, we calculate the evolution of stars with masses less than 4.00 \\mo. We also use the fitting formulae of stellar luminosity and radius, the boundary flux of habitable zones, the distribution of semimajor axis and mass of planets and the initial mass function of stars. We obtain the luminosity and radius of stars with masses from 0.08 to 4.00 \\mo, and calculate the habitable zones of host stars, affected by stellar effective temperature. We achieve the probability distribution of terrestrial planets in habitable zones around host stars. We also calculate that the number of terrestrial planets in habitable zones of host stars is 45.5 billion, and the number of terrestrial planets in habitable zones around K type stars is the most, in the Milky Way. ", "introduction": "Nowadays, more and more exoplanets have being indirectly detected, and they were even directly detected last year (Marois et al., 2008). And theories about planetary evolution point out that most of stars own planets. It is paid closer attention to whether there are lives around other stars, not only by astronomy filed, but also by many other realms. A planet which is suited to life survive must satisfy two factors, one is that it is a terrestrial planet, and the other is that it must be in the habitable zone (HZ) of its host star. Terrestrial planets are rocky planets from one to ten Earth masses with the same chemical and mineral composition as the Earth (Valencia et al., 2006). However, others argued that the mass of terrestrial planet could be down to about 0.3 Earth mass, to retain its atmosphere over long geological timescales and to sustain tectonic activity as required for the carbon-silicate cycle to operate (Williams et al., 1997; Menou \\& Tabachnik, 2003). Typically, stellar HZ is defined as region near the host star, where water at the surface of a terrestrial planet is in liquid phase (eg., Hart, 1978; Kasting et al., 1993; Franck et al., 2000; Noble et al., 2002; Jones et al., 2006). Previously, it is thought that the boundary flux of HZ depends on luminosity (eg., Kasting et al., 1993). However, it is pointed out that this flux also depends on stellar effective temperature $T_{\\rm eff}$ (Forget \\& Pierrehumbert, 1997; Williams \\& Kasting, 1997; Mischna et al., 2000; Jones 2004; Jones et al. 2006). As the greenhouse effect can raise the mean temperature of terrestrial planets (Clube et al., 1996), and the greenhouse effect is different for different radiation. For example, the lower $T_{\\rm eff}$, the more the infrared fraction in luminosity, and the more this fraction, the more the greenhouse effect for a given stellar flux (Jones et al., 2006). Thus, the distances at both inner and outer HZ boundaries are farther to host star, with lower $T_{\\rm eff}$, than they would have been if the $T_{\\rm eff}$ effect is not taken into consideration. And the distances are closer to host star with higher $T_{\\rm eff}$. It is well known that HZ widths for different stars are different. For example, the HZ width for a M type star is only about one fifth to one fiftieth of the HZ width for a G type star (Tarter et al., 2007). However, it does not mean that the probability of terrestrial planet in HZ around a M type star is also about one fifth to one fiftieth of the probability around a G type star. Fortunately, planetary evolution has been calculated, which matched well with the observation data considering selective effects (Ida \\& Lin, 2004, 2005; Schlaufman et al., 2009). Ida \\& Lin (2005) gave the distribution of semimajor axis and mass of planets predicted by the Monte Carlo simulations. The masses of host stars are 0.20, 0.40, 0.60, 1.00 and 1.50 \\mo, respectively. And this distribution can be used to calculated the probability of terrestrial planets in HZs around different host stars. Jones et al. (2006) gave the HZ flux at both inner and outer boundaries, as a function of $T_{\\rm eff}$, but did not give the flux for host stars. Using Eggleton's code, we calculate stellar evolution and achieve the relationship between stellar mass and $T_{\\rm eff}$. We get the HZ flux, at both inner and outer boundaries, for host stars at zero age main sequence (ZAMS) and at the terminal of main sequence (TMS). And the HZ flux for host star with mass 4.00 \\mob is about five times more than that of host star with mass 0.40 \\mob at inner boundary and about four times at outer boundary. Then we obtain the HZs of host stars with masses from 0.08 to 4.00 \\mob at ZAMS and at TMS, taken the $T_{\\rm eff}$ effect into consideration. Using the distribution of semimajor axis and mass of planets (Ida \\& Lin, 2005), we obtain the probability distribution of terrestrial planets in HZs around host stars. Then, we calculate that the number of terrestrial planets in habitable zones of host stars is 45.5 billion in the Milky Way. And the number of terrestrial planets in HZs for M, K, G and F type stars are respectively as 11.548 billion, 12.930 billion, 7.622 billion and 5.556 billion, in the Milky Way. The outline of the paper is as follows: we present some input descriptions in Section 2, give our results in Section 3, present some discussions in Section 4, and then finally in Section 5 we give our conclusions. ", "conclusions": "Firstly, we obtain the HZs of host stars with masses form 0.08 to 4.00 \\mob at ZAMS and at TMS, considering the $T_{\\rm eff}$ effect. Secondly, we give the probability distribution of terrestrial planets in HZs of host stars. Thirdly, we calculate the number of terrestrial planets in HZs of host stars is 45.5 billion in the Milky Way, and find that the number of terrestrial planets around K type stars is the most in the Milky Way. Finally, we present discussions about the host stars with masses about 0.10 \\mob and the effects of type I migration, and introduce our future work. One may also send any special request to \\it{guojianpo1982@hotmail.com} \\normalfont{or} \\it{guojianpo16@163.com}\\normalfont{.}" }, "1003/1003.3537_arXiv.txt": { "abstract": "The continuation of resonant periodic orbits from the restricted to the general three body problem is studied in a systematic way. Starting from the Keplerian unperturbed system we obtain the resonant families of the circular restricted problem. Then we find all the families of the resonant elliptic restricted three body problem which bifurcate from the circular model. All these families are continued to the general three body problem, and in this way we can obtain a global picture of all the families of periodic orbits of a two-planet resonant system. We consider planar motion only. We show that the continuation follows a scheme proposed by Bozis and Hadjidemetriou (1976) for symmetric orbits. Our study includes also asymmetric periodic orbits, which exist in cases of external resonances. The families formed by passing from the restricted to the general problem are continued within the framework of the general problem by varying the planetary mass ratio $\\rho$. We obtain bifurcations which are caused either due to collisions of the families in the space of initial conditions or due to the vanishing of bifurcation points. Our study refers to the whole range of planetary mass ratio values ($\\rho \\in (0,\\infty)$) and, therefore we include the passage from external to internal resonances. In the present work, our numerical study includes the case of the 2/1 and 1/2 resonance. The same method can be used to study all other planetary resonances. ", "introduction": "A good model to study the motion of three celestial bodies considered as point masses, e.g. a triple stellar system or a planetary system with two planets, is the famous {\\em three-body problem} (TBP), whose study goes back to Poincar\\'e. In the present study we consider the case where only one of the three bodies is the more massive one, and the other two bodies have much smaller masses. This model is useful in the study of the motion of small bodies in a planetary system (for example asteroids and comets in our Solar System) or planetary systems with two planets. The simplest model is the {\\em circular restricted TBP}. Although much work has been carried out for this model, (see e.g. Bruno, 1994; H\\'enon, 1997), new interesting results continue to appear in the literature (Maciejewski and Rybicki 2004; Papadakis and Goudas, 2006; Bruno and Varin 2006,2007). In this model we consider two bodies with non zero mass, called {\\em primaries}, for example the Sun and Jupiter, moving in circular orbits around their common center of mass, and a third body with negligible mass, for example an asteroid, moves under the gravitational attraction of the two primaries. A more realistic model is the {\\em elliptic restricted TBP}, where the two primaries move in elliptic orbits, with a finite eccentricity. However, an important aspect of the dynamics is missing in these two models. It is the gravitational interaction between the small body and the two primaries, which is not taken into account in the restricted models. When we introduce this gravitational interaction, we have a more realistic model, the {\\em general TBP}. Within this framework we study a system consisting of a Sun and two small bodies, which we call {\\em planets}. In the study of a dynamical system, the topology of its phase space plays a crucial role. The topology is determined by the position and the stability properties of the periodic orbits, or equivalently, of the fixed points of the Poincar\\'e map on a surface of section. This makes clear the importance of the knowledge of the families of periodic orbits in a dynamical system. Particularly, in a planetary system, many families of periodic orbits are associated with resonances, which are mean motion resonances between the two planets. Since in our study of planetary systems only one body is the more massive one (the sun), a good method is to start from the simplest model, which is the circular restricted problem and find all the basic families of periodic orbits. Then we extend the model to the elliptic restricted model, and find all the families of resonant periodic orbits that bifurcate from the circular to the elliptic model. Finally, we give mass to the massless body and continue all these families to the model of the general problem. This is the method that we shall use in the present study to describe the topology of the phase space. The existence of periodic orbits of the planetary type in the general TBP, as a continuation from the restricted problem, has been studied by Hadjidemetriou (1975, 1976) and recently this method found a fruitful field of applicability in the dynamics of resonant extrasolar systems (e.g. Rivera and Lissauer,2001; Hadjidemetriou, 2002; Ji et al, 2003; Haghighipour et al., 2003; Ferraz-Mello et al., 2003, 2005; Psychoyos and Hadjidemetriou, 2005; Voyatzis and Hadjidemetriou, 2005, 2006; Voyatzis, 2008). An alternative method for determining resonant periodic orbits is the computation of stationary solutions of an averaged model (Beauge et al. 2003; Michtchenko at al. 2006). It is interesting to mention at this point that the situation is not the same in all resonances. It may happen that in some resonances there are not new families in the elliptic model, bifurcating from the circular model, or may exist several resonant families bifurcating from the circular model. This has, evidently, important consequences on the topology of the phase space and differentiates the behavior of the model in different resonances (see Tsiganis et al. 2002a, 2002b). Although in this study we present numerical computations of the 2/1 (or 1/2) resonance, our approach and results have a more general applicability. In the next section we discuss briefly known aspects on the families of periodic orbits in the circular unperturbed and restricted problem. However, these issues are fundamental for continuing our study in the elliptic restricted and in the general problem. In section \\ref{ERTBP} we present the continuation of resonant periodic orbits from the circular to the elliptic model and then, in section \\ref{SectionCon}, we consider the continuation in the general problem. In section \\ref{GTBPbif} we study the bifurcation of families of periodic orbits within the framework of the general problem and finally, in section \\ref{CD}, we discuss the generality of our results and conclude. ", "conclusions": "\\label{CD} The method of continuation of periodic orbits has been applied in the present work in order to study the generation and the structure of families of periodic orbits in the general TBP, starting from the restricted problem. We considered the planar TBP of planetary type, referred to a rotating frame and studied resonant planetary motion. In particular, we studied the 2/1 resonant periodic orbits, but the method we used can be applied to all other resonances. We considered planar motion only. We started from the circular restricted TBP and computed the basic families of 2/1 resonant periodic orbits, both for the inner orbits (inside Jupiter) and the outer orbits (outside Jupiter). Along these families the resonance is almost constant, but the eccentricity of the small body increases and may take high values. The basic families are symmetric with respect to the rotating $x$-axis, but asymmetric families also exist in cases of resonances of the form $1/q$. On these families, we found the periodic orbits with period equal to $2\\pi$ which, in the normalization we are using, are the bifurcation points to resonant 2/1 (or 1/2) families of periodic orbits of the elliptic restricted model, along which the eccentricity of Jupiter varies, starting from zero values. Two such families bifurcate from each of these critical points. In particular, we have the following cases: a) an asymmetric periodic orbit of period $2\\pi$ of the circular restricted problem is continued to the elliptic restricted problem as asymmetric also, b) a symmetric periodic orbit of period $2\\pi$ of the circular restricted problem can be continued to the elliptic problem as an asymmetric one - this exceptional case is verified only up to the accuracy of the numerical computations and c) a symmetric periodic orbit of the elliptic restricted problem, which is of critical stability, can be continued to the elliptic problem as an asymmetric one. The last case indicates the existence of asymmetric periodic orbits in the elliptic problem and in particular to resonances which are not necessarily of the form $1/q$. In this way we obtain a clear picture of all the resonant families of the restricted problem, and in this way we can make a complete study of the resonant families of the general problem, by giving mass to the massless body of the restricted model. The continuation of the families of periodic orbits of the restricted problem to the general problem follows the scenario described in the paper of Bozis and Hadjidemetriou (1976). Particularly, the families of the general problem originate from two families of the restricted problem, one of the circular problem and one of the elliptic problem. There is no essential difference in the continuation between symmetric and asymmetric periodic orbits. The asymmetric orbits, which bifurcate from symmetric orbits of the elliptic problem, are continued smoothly in the general problem. The stability of periodic orbits is preserved after the continuation from the restricted to the general problem. In the paper mentioned above, a case of continuation without the formation of a gap is described. We found this case only at bifurcation points of the families of the elliptic restricted problem (see Fig.\\ref{FRG}b). After the continuation of the periodic orbits of the restricted problem to the general problem, by giving a very small mass to the massless body, we studied how these families evolve when we increase the mass of the small body. Particularly we studied such an evolution by considering as a parameter the planetary mass ratio $\\rho=m_2/m_1$, keeping $m_i\\ll 1, i=1,2$. In our numerical computations we considered $\\max(m_1,m_2)=10^{-3}$ and we varied $\\rho$ by varying only the mass of the smaller body. Starting from $\\rho=0$ (external resonances of the restricted problem) and increasing its value, we found that the characteristic curves of two different families can collide in the space of initial conditions. At these points a bifurcation takes place ({\\em collision--bifurcation}) causing a topological change in the structure of the colliding families and the formation of new families. As $\\rho \\rightarrow \\infty$ we approximate the families of the internal resonances of the restricted problem. Since asymmetric periodic orbits are not known for the internal resonances, it would be interesting for a future work a study of the evolution and the bifurcation of all asymmetric families up to values of $\\rho>1$ and, furthermore, as $\\rho \\rightarrow \\infty$. In general, the continuation and the evolution the families of the external resonances ($\\rho<1$) in the general TBP show more rich dynamics (number and structures of families and bifurcations) in comparison with the case of internal resonances. It is a fact that the complexity of the dynamics of continuation and evolution of families of periodic orbits depends on the existence of critical orbits of period $T=2k\\pi$, orbits of critical stability and asymmetric periodic orbits. For example the family $S'_1$ of the 2/1 internal resonance evolves smoothly as $\\rho$ changes and no any bifurcations are obtained in the domain $(\\bar\\rho_3,\\infty)$. Also we can show that a similar smooth continuation and evolution holds for the family $I_{(e)}$ (parts $I_{(ea)}$ and $I_{(eb)}$ in Fig. \\ref{FRG}). As $\\rho$ starts from zero and tends to infinity, the family evolves smoothly and tends to the family $I_(i)$ (parts $I_{ia}$ and $I_{ia}$). No any bifurcations take place, while the collision orbit, which separates the parts (a) and (b) of the families, exists for all values of $\\rho$. This family is equivalent to the family of ($\\pi,\\pi$)- corotations indicated by Beauge et al (2006) and has a similar evolution as that of the 3/1 family $S_2$ of ($\\pi,0$)- corotations shown in Voyatzis (2008). Apart from collision--bifurcations, structural changes in the characteristic curves of periodic orbits occur when their bifurcation points disappear. We showed how the asymmetric family $A_{123}$, which starts and ends at bifurcation points on the symmetric family $S_1$, shrinks and finally disappears after the collapse of both bifurcation points. An interesting case occurs when the starting and the ending points of a family are bifurcation points that belong to different families. In this case it is possible, for a critical value of $\\rho$, one of the points to disappear or to become a bifurcation point for another family. Such a bifurcation has been observed in the 3/1 resonance (Voyatzis, 2008) and cause global changes in the structure of the characteristic curves. Although we presented results that are associated with the 2/1 resonance, we can claim that similar features of continuation are present in other resonances. Our numerical study can not be extended efficiently up to high eccentricities ($e_1 \\rightarrow 1$ or $e_2 \\rightarrow 1$), due to computation restrictions. We believe that the features of the families in these regions can be revealed by considering the rectilinear restricted problem beside the circular and the elliptic one." }, "1003/1003.3701_arXiv.txt": { "abstract": "We investigate the effect of the stochastic gravitational wave (GW) background produced by kinks on infinite cosmic strings, whose spectrum was derived in our previous work, on the B-mode power spectrum of the cosmic microwave background (CMB) anisotropy. We find that the B-mode polarization due to kinks is comparable to that induced by the motion of the string network and hence the contribution of GWs from kinks is important for estimating the B-mode power spectrum originating from cosmic strings. If the tension of cosmic strings $\\mu$ is large enough i.e., $G\\mu \\gtrsim 10^{-8}$, B-mode polarization induced by cosmic strings can be detected by future CMB experiments. ", "introduction": "Cosmic (super)strings can be produced in the early Universe at the phase transition associated with spontaneous symmetry breaking~\\cite{Vilenkin}, the end of supersymmetric hybrid inflation~\\cite{Jeannerot:1997is, Lyth:1997pf}, or the end of the brane inflation~\\cite{Sarangi:2002yt, Dvali:2003zj}. They can be a clue to particle physics beyond the standard model and the history of the early Universe, which is difficult to obtain in terrestrial experiments. How to find signatures of cosmic strings in cosmic microwave background (CMB) experiments has been extensively discussed for decades. Especially, B-mode polarization of the CMB induced by the cosmic string network was investigated in many papers~\\cite{Seljak:1997ii,Benabed:1999wn,Pogosian:2003mz,Bevis:2007qz,Pogosian:2007gi}. B-mode polarization, which has not been detected yet, is polarization of the parity-odd type. It cannot be produced by the primordial scalar perturbation from the inflationary era, which is widely believed to be the main origin of the present structure of the Universe. On the other hand, the tensor perturbation can be a source of B-mode polarization. Some inflation models can produce the intense tensor perturbation enough to generate detectable B-mode, while others cannot. Cosmic strings can also induce B-mode. Cosmic strings move in the Universe in a very complicated and nonlinear way, constantly generating all types of perturbations, scalar, vector and tensor ones. Therefore, dynamics of the cosmic string network induces B-mode and it reaches an observable level if the tension of cosmic strings, $\\mu$, is large enough, say, $G\\mu \\gtrsim 10^{-7}$~\\cite{Pogosian:2007gi}. Here, $G$ denotes the Newton constant. In this paper, we point out that there is an additional source of B-mode when the cosmic string network exists. It is the stochastic gravitational wave (GW) from kinks on infinite strings.\\footnote{% The effects of the kinks on infinite strings are partially reflected in the calculation in Ref.~\\cite{Bevis:2007qz} based on the lattice simulation. However, such a simulation covers only the limited period of the evolution of the string network, so the correct kink distribution on infinite strings cannot be taken into account by this method. Moreover, it is impossible to completely separate the GWs from kinks from those emitted at the phase transition. Therefore, our calculation based on the kink distribution derived analytically is complementary to the calculation in Ref.~\\cite{Bevis:2007qz}. } In the previous paper~\\cite{Kawasaki:2010yi}, we investigated GWs emitted from kinks on infinitely long strings, and found that GWs with a wavelength comparable to the Hubble horizon scale are generated. These long wavelength GWs can produce an observable B-mode in the CMB. As a result, we find that the contribution of such a GW background to a B-mode is comparable to that due to dynamics of the cosmic string network and detectable by future CMB experiments, such as PLANCK or CMBpol. This paper is organized as follows. In section 2, we present the formalism which we adopt in order to compute the BB power spectrum. In section 3, we briefly review the result of Ref.~\\cite{Kawasaki:2010yi} for the GW spectrum from kinks which will be used in the following analysis. In section 4, we show the resulting BB power spectrum and discuss its observational implications. Section 5 is devoted to summary. ", "conclusions": "In this paper, we have studied the effect of the stochastic GW background induced by kinks on the infinite cosmic strings on the BB power spectrum of CMB polarization. Using the GW background obtained in Ref.~\\cite{Kawasaki:2010yi}, we have estimated the resulting BB power spectrum. We found that this effect is comparable to that of the vector and tensor modes induced by motion of the cosmic string network and may leave observable signatures in the spectrum. If the cosmic string tension is large enough, the BB power spectrum by cosmic strings will be detected by future/on-going satellite experiments such as PLANCK and CMBpol. If it is discovered by the CMB experiments, then the direct detection of GWs from cosmic strings by pulsar timing arrays or space-laser interferometers may further confirm the existence of the cosmic string~\\cite{Kawasaki:2010yi}. \\appendix \\def\\thesection{}" }, "1003/1003.3009_arXiv.txt": { "abstract": "We study cosmological constraints on metric $f(R)$ gravity models that are designed to reproduce the $\\Lambda$CDM expansion history with modifications to gravity described by a supplementary cosmological freedom, the Compton wavelength parameter $B_0$. We conduct a Markov chain Monte Carlo analysis on the parameter space, utilizing the geometrical constraints from supernovae distances, the baryon acoustic oscillation distances, and the Hubble constant, along with all of the cosmic microwave background data, including the largest scales, its correlation with galaxies, and a probe of the relation between weak gravitational lensing and galaxy flows. The strongest constraints, however, are obtained through the inclusion of data from cluster abundance. Using all of the data, we infer a bound of $B_0<1.1\\times10^{-3}$ at the 95\\% C.L. ", "introduction": "Cosmic acceleration can either be explained by introducing large amounts of dark energy or considering modifications to gravity such as the addition of a suitable function $f(R)$ of the Ricci scalar to the Einstein-Hilbert action~\\cite{carroll:03, nojiri:03, capozziello:03}. In fact, one may interpret the cosmological constant as being of this kind rather than attributing it to vacuum energy. It has been argued that a valid $f(R)$ model should closely match the $\\Lambda$CDM expansion history~\\cite{hu:07a, brax:08}. We specialize our considerations to functions $f(R)$ that exactly reproduce this background and parametrize the class of solutions in terms of its Compton wavelength parameter $B_0$~\\cite{song:06}. Such $f(R)$ modifications affect gravity at solar-system scales, which are well tested and impose stringent constraints on deviations from general relativity. However, the chameleon effect~\\cite{khoury:03, navarro:06, faulkner:06} provides a mechanism that allows certain $f(R)$ models to evade solar-system tests (e.g.,~\\cite{hu:07a}). The transition required to interpolate between the low curvature of the large-scale structure and the high curvature of the galactic halo sets the strongest bound on the cosmological field ($B_0\\lesssim10^{-5}$~\\cite{hu:07a, smith:09t}). Independently, strong constraints can also be inferred from the large-scale structure alone. The enhanced growth of structure observed in $f(R)$ gravity models manifests itself on the largest scales of the cosmic microwave background (CMB) temperature anisotropy power spectrum~\\cite{song:06}, where consistency with CMB data places an upper bound on $B_0$ of order unity~\\cite{song:07}. Cross correlations of the CMB temperature field with foreground galaxies serve as another interesting test of $f(R)$ gravity models~\\cite{song:06, song:07}, tightening the constraint on the Compton wavelength parameter by an order of magnitude (e.g.,~\\cite{giannantonio:09}). However, the currently strongest constraints on $f(R)$ gravity from large scale structures are inferred from the analysis of the abundance of low-redshift X-ray clusters, yielding an improvement over the CMB constraints on the free field amplitude of the Hu-Sawicki~\\cite{hu:07a} ($n=1$) model of nearly four orders of magnitude~\\cite{schmidt:09}. In this paper, we perform an independent analysis of $f(R)$ gravity constraints, focusing on cosmological data only. Our analysis differs from previous studies partially in terms of the theoretical model, the parametrical approach, or data sets implemented. We compare and discuss our results to the constraints of previous analyses. Here, we conduct a Markov chain Monte Carlo (MCMC) study of metric $f(R)$ gravity models that are designed to reproduce the $\\Lambda$CDM expansion history using data from CMB anisotropies, supernovae distances, the baryon acoustic oscillation (BAO) distances, and the Hubble constant. For observables in the linear regime, we adopt the parametrized post-Friedmann (PPF) framework~\\cite{hu:07b, hu:08} and its implementation into a standard Einstein-Boltzmann linear theory solver~\\cite{fang:08} for the theoretical predictions. This framework allows us to include information from the near horizon scales. We also utilize information from the cross correlation between high-redshift galaxies and the CMB through the integrated Sachs-Wolfe (ISW) effect. We further use a probe of the relation between weak gravitational lensing and galaxy flows, as well as data from the abundance of clusters that are identified by overdensities of bright, uniformly red galaxies. Latter yields the tightest constraints on the cosmological parameters, particularly on $B_0$. We compare these constraints to the results of~\\cite{schmidt:09} derived for the Hu-Sawicki model. In \\textsection{\\ref{sec:theory}}, we review metric $f(R)$ gravity theory. We present the results of our MCMC study in \\textsection{\\ref{sec:constraints}} and discuss them in \\textsection{\\ref{sec:discussion}}. Finally, the PPF formalism for $f(R)$ gravity~\\cite{hu:08} and details about the modifications to the {\\sc iswwll} code~\\cite{ho:08, hirata:08} used for the galaxy-ISW (gISW) cross correlation observations are specified in the Appendix. ", "conclusions": "\\label{sec:discussion} We have performed a MCMC analysis on metric $f(R)$ gravity models that exactly reproduce the $\\Lambda$CDM expansion history. In addition to geometrical probes from supernovae, BAO distance, and Hubble constant measurements, which were used to fix the background, we utilized all of the CMB data, including the lowest multipoles, its correlation with galaxies, the comparison of weak gravitational lensing to large-scale velocities, and the abundance of clusters. We report a constraint on the Compton wavelength parameter of $B_0<1.1\\times10^{-3}$ at the 95\\% C.L. from using all of the measurements. This result is substantially driven by data from cluster abundance. However as the data improve, the limits will saturate due to the chameleon effect in massive haloes. gISW measures in combination with the CMB, supernovae, BAO distance, and Hubble constant probes, yield a constraint of $B_0<0.42$ (95\\% C.L.), which is an order of magnitude improvement over using the CMB alone as probe of the growth of large-scale structure in combination with the geometrical measures. This highlights the power of gISW measurements as a linear theory probe to constrain infrared modifications of gravity. The $E_G$ measurement of the relationship between weak gravitational lensing and galaxy flows does not improve bounds on $f(R)$ gravity on its own. However, when used as a complementary probe to cluster abundance, it contributes substantially to our constraints. This can be attributed to the slow convergence of its prediction toward $\\Lambda$CDM when $B_0\\rightarrow0$. It is likely that with improved data, the $E_G$ probe will become an important discriminator for gravity models." }, "1003/1003.4135_arXiv.txt": { "abstract": "Twenty six Asymptotic Giant Branch (AGB) variables are identified in the Local Group galaxy Leo~I. These include 7 Mira and 5 semi-regular variables for which periods, amplitudes and mean magnitudes are determined. The large range of periods for the Miras, $158

-1$) globular clusters, which have ages greater than 10\\,Gyr. Within these clusters, their periods and therefore of course their magnitudes are proportional to the metallicity of the parent cluster (Feast, Whitelock \\& Menzies 2002). Miras are not found in the more metal-deficient clusters. Although we have yet to confirm spectroscopically that these two Miras are carbon-rich, their colours certainly suggest it. Most models do not produce carbon stars or high-mass-loss objects at ages of 10 Gyr, although recent work suggests it might happen. Karakas (2010) modelled a $\\rm 1 M_{\\odot}$ star with Z=0.0001 and found that it experienced 26 thermal pulses and a small amount of third-dredge-up. In an envelope with such a low metallicity, even a small amount of dredge-up was enough to make $\\rm C/O> 1$ and produce a carbon star. This is clearly an area where more work is needed and these stars are worth a more detailed investigation. For the future, with the next generation of large telescopes working in the infrared, Mira variables will prove vital distance indicators for studying populations of old and intermediate age stars, where they will be amongst the most luminous objects, easily identified via their large amplitude variations." }, "1003/1003.3838_arXiv.txt": { "abstract": "Hot Jupiters, with atmospheric temperatures $T \\gsim 1000$~K, have residual thermal ionization levels sufficient for the interaction of the ions with the planetary magnetic field to result in a sizable magnetic drag on the (neutral) atmospheric winds. We evaluate the magnitude of magnetic drag in a representative three-dimensional atmospheric model of the hot Jupiter HD 209458b and find that it is a plausible mechanism to limit wind speeds in this class of atmospheres. Magnetic drag has a strong geometrical dependence, both meridionally and from the day to the night side (in the upper atmosphere), which could have interesting consequences for the atmospheric flow pattern. By extension, close-in eccentric planets with transiently heated atmospheres will experience time-variable levels of magnetic drag. A robust treatment of magnetic drag in circulation models for hot atmospheres may require iterated solutions to the magnetic induction and Saha equations as the hydrodynamical flow is evolved. ", "introduction": "Hot Jupiters are close-in, presumably tidally-locked gaseous giant planets orbiting only a few stellar radii away from their Sun-like host star. By virtue of their slow rotation (synchronous with their orbital periods, $\\sim$ a few days), high atmospheric temperatures ($T \\gsim 1000$~K) and permanent day-side hemispheric forcing, hot Jupiters are laboratories for the study of atmospheric dynamics in a regime that is absent from the Solar System (see Showman et al. 2008 and Showman et al. 2010 for reviews). In recent years, considerable progress has been made in observationally characterizing the atmospheres of hot Jupiters via a combination of secondary eclipse, transmission spectrum and phase curve measurements (see Deming 2008 and Charbonneau 2009 for reviews). In parallel with this observational progress, theoretical modeling of hot Jupiter atmospheres has expanded greatly, in an attempt to provide robust interpretations of the growing data set (see, e.g., Burrows \\& Orton 2010, Showman et al. 2010 and Baraffe et al. 2010 for reviews). One of the main interests in studying hot Jupiter atmospheres lies in understanding their thermal and dynamical responses to the unusual forcing conditions they are experiencing, with an atmospheric circulation pattern that is likely different from anything known in the Solar System. However, with this new regime also comes the possibility that new physics is at play in these extreme atmospheres (e.g., Menou \\& Rauscher 2010). In this work, we investigate the possibility that magnetic drag on atmospheric motions provides an effective frictional mechanism limiting the asymptotic speeds of winds in hot Jupiter atmospheres. This is particularly important as the fast (transonic) speeds reached by winds in a variety of drag-free atmospheric models for this class of planets (Dobbs-Dixon \\& Lin 2008, Dobbs-Dixon et al. 2010; Showman et al. 2009; Rauscher \\& Menou 2010) raise issues about compressibility, shocks and associated energy conservation for the models (Goodman 2009; Rauscher \\& Menou 2010). Atmospheric motions are driven by pressure-gradient forces arising from differential heating of the atmosphere. A small fraction of the atmospheric \"available\" enthalpy is continuously converted into kinetic energy of the atmospheric motions, which is itself continuously dissipated by friction\\footnote{In this context, friction typically refers to a dissipative process that is much more efficient than the microscopic viscosity of the atmospheric gas.} (Lorenz 1955, Pearce 1978, Marquet 1991, Goodman 2009). In steady-state, asymptotic wind speeds are thus reached through a detailed balance between continuous thermal forcing and sustained friction. While the source of wind friction on the Earth, and other Solar System terrestrial planets by extension, is understood to be largely associated with surface drag, the origin of friction in the atmospheres of gaseous giant planets remains a major open question in atmospheric science, even in the Solar System (e.g., Schneider \\& Liu 2009; Liu et al. 2008; Showman et al. 2010). Identifying dominant sources of internal friction in gaseous giant planet atmospheres can thus be as important as adequately modeling their sources of thermal forcing. We show here that magnetic drag on weakly-ionized winds in the predominantly neutral atmospheres of hot Jupiters, which arises from wind interaction with the magnetic field generated in the planet's bulk interior, is a plausible source of sizable friction, that may need to be accounted for in atmospheric circulation models of hot Jupiters. We note that, while this manuscript was being prepared, Batygin \\& Stevenson (2010) completed a study of the closely related ohmic dissipation process associated with the currents induced by magnetic drag. Here, we focus on the role of magnetic drag on atmospheric winds and defer a study of ohmic dissipation to future work. ", "conclusions": "Circulation models for hot Jupiter atmospheres have gained substantial interest in recent years thanks to the direct detections of such atmospheres, which indicate a possible role for winds in shaping the emergent properties of these planets. All existing calculations have been purely hydrodynamical in nature, neglecting the possibility that interactions between the flow and the planetary magnetic field could influence the circulation. However, given the high atmospheric temperatures of these planets, the small fraction of the fluid that is ionized may be sufficiently large to effectively couple the mostly neutral flow to the planetary magnetic field. In this work, we have evaluated the typical magnitude of the magnetic drag exerted on hot Jupiter atmospheric winds, using the flow pattern obtained in the 3D circulation model of Rauscher \\& Menou (2010). Our results indicate that typical magnetic drag stopping times are a strong function of depth, longitude and latitude. In the upper model atmosphere, temperature differences between the day and night side cause extreme variations of the magnitude of the drag, which could be dynamically important only on the day side. At all levels, zonal drag also varies strongly with latitude in the vicinity of the equator, by virtue of the aligned dipole field geometry adopted. By adopting a simplified treatment of ionization balance, focusing on the zonal component of drag only, ignoring upper boundary currents or the role of ambipolar diffusion, and neglecting horizontal variations in resistivity, the calculations presented in this work are clearly of limited scope. Nevertheless, they appear sufficient to make the case that dynamically interesting levels of magnetic drag may have to be accounted for in circulation models for hot Jupiter atmospheres. This conclusion is supported by a simple comparison between a drag-free circulation model and a few additional models with Rayleigh drag applied at levels commensurate with that expected from magnetic drag, which show dynamically interesting consequences for the flow, especially for large magnetic field values. A better assessment of the role of magnetic drag in hot Jupiter atmospheres may require iterated solutions to the multi-dimensional induction equation, together with Saha's equation, as the time-dependent hydrodynamical flow is evolved. Such work will be important to re-evaluate the role of deep ohmic dissipation in possibly inflating hot Jupiters (Batygin \\& Stevenson 2010) and to elucidate the plausible diversity of atmospheric behaviors expected from a variety of field strengths and geometries. The strong dependence of resistivity on atmospheric temperature could, in principle, lead to different classes of atmospheric behaviors as a function of mean orbital separation, and it will likely lead to an interesting phenomenology for eccentric close-in planets, as they experience time-variable levels of magnetic drag (and ohmic dissipation)." }, "1003/1003.1656_arXiv.txt": { "abstract": "{} {A recent research shows that particles with a spectrum of a relativistic Maxwellian plus a high-energy tail can be accelerated by relativistic collisionless shocks. We investigate the possibility of the high-energy particles with this new spectrum injected in pulsar wind nebulae (PWNe) from the terminate shock based on the study of multiwavelength emission from PWNe.} {The dynamics of a supernova remnant (SNR) and multiband nonthermal emission from the PWN inside the remnant are investigated using a dynamical model with electrons/positrons injected with the new spectrum. In this model, the dynamical and radiative evolution of a pulsar wind nebula in a non-radiative supernova remnant can be self-consistently described.} {This model is applied to the three composite SNRs, G0.9+0.1, MSH 15-52, G338.3-0.0, and the multiband observed emission from the three PWNe can be well reproduced.} {Our studies on the three remnant provide evidence for the new spectrum of the particles, which are accelerated by the terminate shock, injected into a PWN.} ", "introduction": "PWNe, which are prominent sites of high-energy emission in the Galaxy, are powered by pulsars associated with them. A pulsar inside a PWN loses its rotational energy through a pulsar wind composed of magnetic flux and high-energy particles \\citep[][]{GJ69,Get07,GSZ09}. The ultra-relativistic wind flows relativistically into a non-relativistic ejecta of the ambient supernova remnant (SNR), which results in the PWN and a termination shock (TS), where the plasma is decelerated and heated \\citep[][]{RC84,Vet08}. High-energy particles are injected into the nebula from the TS, and multiband nonthermal photons with energies ranging from radio, X-ray to $\\gamma$-ray bands are emitted during the evolution of the PWN. Usually, multi-wavelength observational results of a PWN cannot be well reproduced by the radiation of particles injected with a single power-law spectrum, and a broken power-law must be employed to better explain the observations \\citep[e.g.,][]{AA96,A05a,ZCF08}. However, the physics behind the broken power-law is unclear. For a typical PWN, the wind from the pulsar can flow relativistically with a lorentz factor of $\\sim 10^6$. The resulting TS can accelerate particles to relativistic energy. On the other hand, based on the long-term two dimensional particle-in-cell simulations, \\citet[][]{Sp08} found that the particle spectrum downstream of a relativistic shock consists of two components: a relativistic Maxwellian and a high-energy power-law tail with an index of $-2.4\\pm0.1$. Based on this finding, with the assumption that the high-energy particles in a PWN are injected with a spectrum of a relativistic Maxellian plus a power-law high-energy tail, we investigate the multiwavelength emission from PWNe to test the possibility of particles with the new spectrum injected into the PWNe. The dynamical evolution of the PWN is calculated basically according to the model in \\citet[][]{GSZ09}, which can self-consistently describe the dynamical and radiative evolution of a pulsar wind nebula in a non-radiative supernova remnant. Different from \\citet{GSZ09}, in which a single power-law injection spectrum for the electrons/positrons is employed to discuss the radiative properties during different phase of the PWN, we argue in this paper that the high-energy particles are injected with the new spectrum of a relativistic Maxellian plus a power-law high-energy tail during the evolution in this paper, and a kinetic equation is used to obtain the energy distribution of the particles. We apply the model to the PWNe in the composite SNRs, G0.9+0.1, MSH 15-52 and G338.3-0.0, which have been observed in radio, X-rays and very high-energy (VHE) $\\gamma$-rays. G0.9+0.1 has a 2$'$ PWN inside a 8$'$ shell in the radio band \\citep[][]{HB87}. The PWN is powered by an energetic pulsar PSR J1747-2809, which was recently discovered in G0.9+0.1 with the NRAO Green Bank Telescope at 2GHz \\citep[][]{Cet09}. A jet-like feature of the nebula was revealed in the high-resolution observation with {\\it Chandra} \\citep[][]{GPG01}. The observation in the TeV band with HESS indicates that the PWN is a weak emitter in VHE $\\gamma$-rays \\citep[][]{A05a}. MSH 15-52 is a complex SNR with a ragged shell in the radio observations \\citep[][]{Cea81}. An elongated PWN powered by an energetic pulsar was found in the remnant in the X-ray observations with ROSAT \\citep[][]{Tea96}, BeppoSAX \\citep{M01} and Chandra \\citep[][]{G02}. The remnant has also been well detected in VHE $\\gamma$-rays with HESS \\citep[][]{A05b} and CANGAROO \\citep[][]{Nea08}, and the significant VHE $\\gamma$-rays are identified to be produced from the PWN. A VHE $\\gamma$-ray source HESS J1640-465 was discovered by HESS in the survey of the inner Galaxy \\citep[][]{A06}, and it is spatially coincident with the composite SNR G338.3-0.0. An extended X-ray source was found to be located at the center of the VHE $\\gamma$-ray source with XMM-Newton \\citep[][]{Fea07}. Very recently, \\citet[][]{Lea09} presented the high resolution X-ray observations on the PWN, and a point like source as a putative pulsar appears in the X-ray observations. In this paper, we investigate the possibility of the particles with the new spectrum injected in PWNe based on applications with the spectrum to the three composite SNRs. In Section \\ref{sec:model}, the model is simply described, and the results from the applications to the three SNRs G0.9+0.1, MSH 15-52 and G338.3-0.0, are presented. The main conclusions and discussion are given in Section \\ref{sec:discussion}. ", "conclusions": "\\label{sec:discussion} Motivated by the finding that the spectrum of the particles downstream of a relativistic shock consists of two components: a relativistic Maxwellian and a power-law high-energy tail with an index of $-2.4\\pm0.1$ \\citep[][]{Sp08}, we investigate the possibility of particles with this new spectrum injected in PWNe from the TS based on the studies of multiband emission from PWNe. Following the dynamical method proposed in \\citet[][]{GSZ09}, we study the dynamical and multi-band radiative properties of the three composite SNRs G0.9+0.1, MSH 15-52 and G338.3-0.0. With appropriate parameters, we find that a typical PWN is an important $\\gamma$-ray emitter during its evolution although the non-thermal radiation from the radio to the X-ray band is insignificant sometimes. The multiband observations of the three PWNe in the remnants can be well reproduced with the new spectrum of the injected particles. Therefore, our studies on the dynamical and multiwavelength radiative properties of PWNe provide evidence of high-energy electrons/positrons can be injected into a PWN with a Maxwellian plus a power-law high-energy tail from the TS of the PWN. In modeling the multiband nonthermal emission from a PWN detected in the radio, X-ray and $\\gamma$-ray bands, particles injected with a spectrum of a broken power-law are widely used to reproduce the observed multiwavelength emission \\citep[e.g.,][]{Vd06,S08,ZCF08}. Of course, for the three PWNe discussed in this paper, the multiband observed spectra of them can also be explained if the particles are injected with a broken power-law. However, it is unclear why the broken power-law spectrum is valid when using it to reproduce the multiwavlength emission from a PWN. From our calculations, we have found out that the energy distribution of the electrons/positrons in the nebula can be approximated as a broken power-law with an index $\\sim 1$ in the lower-energy band and an index of $\\sim2.5$ in the higher-energy part before the PWN undergos significant compression, which is most likely the physical explanation of the broad usage of a broken power law in modeling the multi-band non-thermal emission from PWNe. In this paper, high-energy electrons/positrons are injected into the PWN from the TS, and the main energy of the nebula is contained in these particles. These particles undergo radiative and adiabatic losses when the nebula evolves in the host SNR. Our study indicates that, for a typical PWN with the parameters similar as G0.9+0.1, the adiabatic loss of the particles in the nebula is significant after an age of $\\sim1000$ yr (see Fig.\\ref{Epwn} and Fig.\\ref{Epower}). Multiwaveband nonthermal emission from a PWN has been investigated using a simplified time-dependent injection model, in which high-energy electrons/positrons are injected into the PWN \\citep[e.g.,][]{Vd06,S08,ZCF08}. The pulsar inside the PWN transfers a part of its spin-down power to the particles with a spectrum of a broken power-law. In the simplified time-dependent injection model in \\citet[][]{ZCF08}, synchrotron loss of the particles is taken into account, whereas the adiabatic one is ignored. As a result, either a relatively smaller initial spin-down power of the pulsar or a smaller efficiency of the power to the kinetic energy of the accelerated electrons/positrons is employed in the model. Moreover, note that in \\citet[][]{ZCF08}, an initial spin-down power of $1\\times10^{38}$ erg s$^{-1}$ for MSH 15-52 was used to investigate the multiband emission from the PWN, which is a factor of 15 smaller than that used in this paper. Besides the above reasons, the another main one is a relatively big spin-down time scale of $\\sim 5000$ yr, which is $\\propto \\dot{E_0}^{-1}$ in the paper, used by \\citet[][]{ZCF08}, whereas in this paper it is adopted to be $500$ yr. The energy released by the pulsar is mainly determined by $E_0 \\min \\{T_{\\rm age}, \\tau_0 \\}$, and the value in this paper is not much bigger than that in \\citet[][]{ZCF08}. Therefore, the multiband observed spectra for MSH 15-52 can be reproduced within the two scenarios even the initial spin-down power of the pulsar is significantly different." }, "1003/1003.6076_arXiv.txt": { "abstract": "{The identification of the solar-like oscillation modes, as measured by asteroseismology, is a necessary requirement in order to infer the physical properties of the interior of the stars. Difficulties occur when a large number of modes of oscillations with a low signal-to-noise ratio are observed. In those cases, it is of common use to apply a likelihood-ratio test to discriminate between the possible scenarios. We present here a statistical analysis of the likelihood-ratio test and discuss its accuracy to identify the correct modes. We use the AsteroFLAG artificial stars, simulated over a range of magnitude, inclination angle, and rotation rate. We show that the likelihood-ratio test is appropriate up to a certain magnitude (signal-to-noise ratio). } ", "introduction": "The recent development of space-based instruments such as the Convection, Rotation, and planetary Transits mission \\citep[CoRoT;][]{Michel08} and the NASA's Kepler mission \\citep{Borucki09} as well as the organized ground-based campaigns \\citep{Arentoft08} are producing a huge volume of asteroseismic observations of unprecedented quality. Thus, it is now possible to correctly detect individual pressure (p) driven modes and to describe the oscillation modes with Lorentzian profiles, as it is commonly done for the Sun \\citep[e.g.,][]{Chaplin06}. However, in the stellar case, there is a substantial difference: while in the Sun, the inclination of the rotation axis and the surface rotation rate are known, in most of the stars analyzed up to now, these two parameters are unknown. Moreover, as these two parameters are highly correlated \\citep{Gizon03,Ballot06}, in order to improve the stability of the fits, the traditional pair-by-pair fitting methodology followed in the Sun is changed to a global strategy in which all the modes are fitted at the same time along with the inclination angle and one value for the rotational splitting. The first application of this strategy was introduced in asteroseismology by \\citet{Appourchaux08} and it has been since succesfully applied to the CoRoT solar-like targets in which the signal-to noise ratio was high enough \\citep{Barban09, Garcia09, Deheuvels10}. Indeed, some of the measured CoRoT stars were too faint to perform such kind of analyses \\citep{Mosser09,Mathur10a}. Moreover, the global fitting has also been applied to the first Kepler solar-like targets \\citep{Chaplin10}, as well as the latest ground-based observational campaign of Procyon \\citep{Bedding10b}. Everything would be great if we were able to identify the individual modes (or mode-tagging) before performing the peak-bagging. Indeed, for most of the observations, a clear identification of the modes (or of the ridges) appears to be a difficult task and it is then necessary to perform the peak-fitting using the two possible identifications. The possible scenarios (or taggings) are discriminated by comparing a posteriori the likelihoods of the minimization, the highest likelihood being chosen as the correct mode identification. \\citet{Benomar09} using a longer time series -- which means a better overall signal-to-noise ratio -- demonstrated by comparing the likelihoods that the first p-mode identification done by \\citet{Appourchaux08} of the CoRoT target HD49933 was wrong. However, we present in this work that the so-called {\\it likelihood-ratio test} has certain limits depending on both the magnitude of the star and the length of the observations (i.e., the overall signal-to-noise ratio of the modes in the power spectrum). It is important to remember that an incorrect mode tagging would have very bad consequences on the inferences of the stellar properties \\citep[e.g.,][]{creevey07,Stello09} . ", "conclusions": "The identification of the oscillations modes (or mode tagging) is a necessary step before we can infer the physical properties of the star interiors. Even when using global peak-fitting techniques, the mode tagging remains a difficult task to achieve, when for instance, a large number of oscillation modes are observed with a low signal-to-noise ratio. The likelihood-ratio test can help to discriminate between the possible scenarios (i.e. taggings) and it has been already successfully applied in asteroseismology with the CoRoT, Kepler, and Procyon observations. Nevertheless, by analyzing the AsteroFLAG artificial star Pancho, we showed in this work that the likelihood-ratio test has certain limits depending on both the star magnitude and the length of the observations. For example, for a star with a rotational splitting of $5\\mu$Hz and an inclination angle of $60^{\\circ}$, observed during 31 days, the likelihood-ratio test will statistically return the incorrect mode identification 25\\% of the time for a star magnitude up to M$_v$ = 11. This percentage decreases as the length of observation increases and for time series of 365 days, the likelihood-ratio test will return 100\\% of the time the correct identification up to M$_v$=11. However, for stars with slow rotation, the likelihood-ratio test is more {\\it likely} to return the incorrect mode tagging. The use of Bayesian conditions help mostly when the star characteristics (for example, slow rotation) make the mode identification difficult. With the long-term observations that will be collected by the Kepler mission, we hope to have enough signal-to-noise ratio to unambiguously determine the correct identification of the modes for both solar-like stars \\citep{Chaplin10} and red giants \\citep{Bedding10a}, as well as the stars in open clusters \\citep{Stello10}. However the relative faintness of these later stars will probably require the use of the likelihood-ratio test to disentangle between the tagging of the modes." }, "1003/1003.3186_arXiv.txt": { "abstract": "Using the \\kmeans\\ cluster analysis algorithm, we carry out an unsupervised classification of all galaxy spectra in the seventh and final Sloan Digital Sky Survey data release (SDSS/DR7). Except for the shift to restframe wavelengths, and the normalization to the $g$-band flux, no manipulation is applied to the original spectra. The algorithm guarantees that galaxies with similar spectra belong to the same class. We find that 99\\% of the galaxies can be assigned to only 17 major classes, with 11 additional minor classes including the remaining 1\\%. The classification is not unique since many galaxies appear in between classes, however, our rendering of the algorithm overcomes this weakness with a tool to identify borderline galaxies. Each class is characterized by a template spectrum, which is the average of all the spectra of the galaxies in the class. These low noise template spectra vary smoothly and continuously along a sequence labeled from 0 to 27, from the reddest class to the bluest class. Our Automatic Spectroscopic K-means-based (ASK) classification separates galaxies in colors, with classes characteristic of the red sequence, the blue cloud, as well as the green valley. When red sequence galaxies and green valley galaxies present emission lines, they are characteristic of AGN activity. Blue galaxy classes have emission lines corresponding to star formation regions. We find the expected correlation between spectroscopic class and Hubble type, but this relationship exhibits a high intrinsic scatter. Several potential uses of the ASK classification are identified and sketched, including fast determination of physical properties by interpolation, classes as templates in redshift determinations, and target selection in follow-up works (we find classes of Seyfert galaxies, green valley galaxies, as well as a significant number of outliers). The ASK classification is publicly accessible through various websites. ", "introduction": "{\\em The nebulae\\footnote{{\\rm The galaxies.}} are so numerous that they cannot be studied individually. Therefore, it is necessary to know whether a fair sample can be assembled from the most conspicuous objects and, if so, the size of the sample required} \\citep[][ Chapter II]{hub36}. Even though these arguments are from the outset of extragalactic astronomy, and they refer to the morphological classification of galaxies, the reasons put forward by Hubble remain valid today. The need to sort out and simplify justify all recent efforts to classify the spectra of galaxies (\\S~\\ref{intro.intro}), including the present work. Such attempts are now more significant than ever since we have never had the large catalogs of galaxy spectra available today. The seventh and final Sloan Digital Sky Survey data release (SDSS/DR7) provides spectra of some 930000 galaxies \\citep[][ and also the SDSS Web site\\footnote{\\tt http://www.sdss.org/dr7}]{sto02,aba09}. This uniform data set offers a unique opportunity to comprehensively classify the different spectra existing among nearby galaxies. Our paper presents the results of an unsupervised spectral classification of all the catalog. Unsupervised implies that the algorithm does not have to be trained. It is autonomous and self-contained, with minimal subjective influence. Thus, we deliberately avoid the use of physical constraints, or other a priori knowledge. We classify all galaxies simultaneously, requesting that galaxies with similar rest-frame spectra belong to the same class. This approach is in the vein of the rules for a good classification discussed by \\citet{san05}, where he points out that physics must not drive a classification. Otherwise the arguments become circular when the classification is used to drive physics. The \\kmeans\\ algorithm that we implement is commonly employed in data mining, machine learning, and artificial intelligence \\citep[e.g.,][]{eve95,bis06}, but it has been seldom applied in astronomy \\citep[see, however,][]{san00}. From the point of view of the algorithm, the galaxy spectra are vectors in a high-dimensional space, where they are distributed among a number of cluster centers. Each vector is assigned to the cluster whose center is nearest, and the center is the average of all the points in the cluster. It works iteratively. Starting from guess cluster centers, the spectra are assigned to their nearest centers, and then the centers are re-computed until convergence is reached. (Further details are given in \\S~\\ref{algorithm}.) We choose it because of its extreme computational simplicity, as required to deal with large data sets (\\S~\\ref{algorithm}), and because it turned out to work very well in the first case we attempted. To our surprise, the algorithm managed to separate spectra of galaxies in the {\\em green valley} within a collection of dwarf galaxies encompassing the full range of spectral types \\citep[][\\S~3.1]{san09}. Therefore, we found it natural to test the ability of \\kmeans\\ to distinguish among all kinds of galaxy spectra, and the success of this follow-up exercise is precisely the work reported here. In addition to the above virtues, the \\kmeans\\ method provides a prototypical high signal-to-noise spectrum for each class of galaxy, being the spectra of the galaxies in a class similar to the associated prototypical spectrum. These few representative spectra can be studied and characterized in detail as if they were individual galaxies, and then their properties can be attributed to all class members (\\S~\\ref{applications}). Other popular classification methods lack this powerful and convenient feature (see \\S~\\ref{intro.intro}). The acronym ASK stands for Automatic Spectroscopic K-means-based, and it is used throughout the text to denote our classification. The paper is structured as follows. \\S~\\ref{intro.intro} provides an overview of the main spectral classification methods employed so far. It also summarizes systematic trends resulting from the application of those methods. Our \\kmeans\\ classification algorithm is examined in \\S~\\ref{algorithm}, where we test the class recovery upon known classes (\\S~\\ref{test_qbcd}), we analyze the repeatability of the classification (\\S~\\ref{recovery}), and we assign probabilities to class membership (\\S~\\ref{assigning}). The SDSS/DR7 dataset is briefly introduced in \\S~\\ref{data_set}. The actual classification of SDSS/DR7 is described in \\S~\\ref{final_imp}. The \\ask\\ is compared with Principal Component Analysis (PCA) classification in \\S~\\ref{pca_class} (see also \\S~\\ref{intro.intro}). The self-consistency of the \\ask\\ is discussed in various sections dealing with specific results; relationship between ASK class and Hubble type (\\S~\\ref{ask_class_vs_morph}), ASK class and color sequence (\\S~\\ref{colorscolors}), ASK class and AGN activity (\\S~\\ref{ask_vs_agn}), and ASK class and redshift (\\S~\\ref{ask_vs_redshift}). Further applications of the classification procedure are sketched in \\S~\\ref{applications}. The \\ask\\ is publicly available as we explain in \\S~\\ref{conclusions}. This section also outlines ongoing works based on ASK. \\subsection{Spectral classification of galaxies}\\label{intro.intro} The first spectral classifications of galaxies are almost coeval with the discovery of the Hubble sequence. \\citet{hub36} discusses how the spectral types and colors systematically vary within the morphological sequence, being ellipticals the reddest and open spirals the bluest \\citep[see also][]{hum31}. One of the early attempts to set up an spectroscopic classification of galaxies is that by \\citet{mor57}. They assign the blue part of the visible spectrum (3850\\,\\AA\\ -- 4100\\,\\AA ) to stellar classes from A to K. They find a clear relationship between spectral class and shape, with the most concentrated galaxies (E, S0) belonging to class K, and the most diffuse galaxies (Sc, Irr) included in class A. The relationship applies to some 80\\% of the galaxies, a percentage probably larger for the targets of highest luminosity. \\citet{aar78} shows how the visible and IR colors of galaxies along the Hubble sequence can be understood as a one parameter family, in terms of the superposition of spectra of A0V dwarf stars and M0III gigant stars. \\citet{ber95} points out that a simple model consisting of two stellar spectral types can reproduce the observed broad band colors, but only if the spectral types are allowed to vary. Five primary spectral types result from this modeling. Similar conclusions are also reached by \\citet{zar95} using stellar spectrum fitting. Principal Component Analysis (PCA) is probably the most popular classification method employed so far. Each spectrum is decomposed as a linear superposition of a small number of eigenspectra, so that a few coefficients in this expansion (eigenvalues) fully describe the spectrum. It is fairly fast and robust, and a solid mathematical theory supports it \\citep[][]{eve95}. To the best of our knowledge, the first applications of PCA in this field have to do with stellar classification \\citep[e.g.,][]{dee64,whi83}, then moved to quasar spectra \\citep[e.g.,][]{mit90,fra92}, and finally arrived to the spectral classification of regular galaxies \\citep[e.g.,][]{sod94,con95}. PCA is the method of reference, and we compare it in \\S~\\ref{pca_class} with our \\kmeans . Two general results are common to all PCA analyses. Spectrumwise, galaxies can be characterized and distinguished by means of a single parameter that links the coefficients of the two or three first eigenspectra. Then different classes are obtained by splitting (somewhat artificially) this 1-dimensional family into pieces. The approach holds for 2dF galaxies \\citep{fol99,mad03}, for galaxies in \\citet{ken92} \\citep{con95,sod97}, for Las Campanas Redshift Survey galaxies \\citep{bro98}, for DEEP2 galaxies \\citep{mad03b}, for IUE galaxies \\citep{for04}, and for SDSS \\citep{yip04}. The second common result is the correspondence between spectral sequence and Hubble type. Even though ellipticals tend to be red and spirals tend to be blue, such relationship has a large intrinsic scatter \\citep{con95,sod97,fer06b}, which augments towards the UV \\citep{for04}. Sometimes elliptical galaxies with blue colors are found in the local universe \\citep[e.g.,][]{kan09}, and this deviation from the trend is expected to grow even further with increasing redshift if, as \\citet{con06} argues, it is a coincidence that Hubble types correlate with colour in the nearby universe. At higher redshifts morphologically classified ellipticals are often blue in colour and actively forming stars \\citep{con06,hue09}. Despite the advantages mentioned above, PCA presents a clear drawback. It does not provide prototypical spectra to characterize the classes. The PCA eigenspectra do not resemble any member of the data to be classified and, in general, eigenspectra are of difficult physical interpretation \\citep[e.g.,][]{cha03,for04,yip04}. The advantage of having classes characterized by prototypical spectra is clear. These few spectra can be studied in detail using standard diagnostic techniques developed for individual galaxies through the years. Then the attributes of the prototypical spectra can be passed on to the class members, or they can be used as intermediate grid-points to interpolate the properties of the class members (see \\S~\\ref{applications}). Moreover, the differences between a particular galaxy and its class prototype allow for precise relative measurements. In an attempt to complement PCA with this feature, \\citet{cha03} developed an {\\em archetypal analysis} algorithm. As the authors explain, it is like PCA but the eigenspectra are required to be members or mixtures of members of the input data set. However, by construction, the eigenspectra are extreme data points lying on the data set outskirts. Although physically meaningful, the eigenspectra are outliers, and it may be difficult to connect their physical properties with those of typical galaxy spectra. Other improvements on the basic PCA technique are local linear embedding \\citep{van09}, and ensemble learning independent component analysis \\citep{lu06}. These extensions are computationally expensive, and so far they have been introduced only as proof-of-concept works. In addition to the superposition of stellar spectra and the PCA techniques described above, galaxies have been classified using neuronal networks \\citep{fol96,mad03}, massive lostless data compression \\citep{rei01}, information bottleneck \\citep{slo01}, and probably others. The algorithms have flourished in response to the availability of new large spectral databases. We are still in an expanding phase, which should lead to a final convergence of the various techniques. The different methods seem to roughly coincide in the global picture, but it is so far unclear whether they agree in the details. ", "conclusions": "We present an automatic unsupervised classification of all the galaxies with spectra in the final SDSS data release (DR7). It uses the \\kmeans\\ algorithm, which separates the 930000 galaxies into 17 major classes containing 99\\% of the galaxies, plus another 11 minor classes with the rest. The algorithm guarantees that the galaxies in a class have similar spectra, independently of their luminosities. The algorithm does guarantees that the classes represent true clusters in the classification space, nor that all existing clusters are identified. Each ASK class\\footnote{Acronym for Automatic Spectroscopic K-means-based class.} is characterized by an extra-low noise spectrum resulting from averaging all the spectra in the class. These template spectra vary smoothly and systematically among the classes, labeled according to their $g-r$ color from ASK~0, the reddest, to ASK~27, the bluest (Fig.~\\ref{classification} and Table~\\ref{tab2}). The classes are well separated in the color sequence, with a class that collects most of the red sequence galaxies (ASK~2), a set of classes lying along the blue cloud (ASK 9 and larger), and a class that seems to be characteristic of the green valley (ASK~5); see \\S~\\ref{colorscolors}. Usually the classes of red galaxies do not present emission lines, however, when they do, their excitation is characteristic of AGN activity. In contrast, all the galaxies in the classes on the blue cloud seem to present emission lines, but they are typical of star formation regions. The classes in between (i.e., in the green valley) show AGN activity (see \\S~\\ref{ask_vs_agn}). The ASK classification has been compared with the morphological Hubble type. Although the number of galaxies involved in this comparison is rather limited, it clearly shows how the red classes tend to have early morphological types, whereas the blue classes are morphologically late types (Fig.~\\ref{assign_kennicutt} and \\S~\\ref{ask_class_vs_morph}). The relationship has a large intrinsic scatter, as previous studies also find (see \\S~\\ref{intro.intro}). We have confronted the ASK classes with the PCA-based spectroscopic classification also existing for SDSS/DR7 (\\S~\\ref{pca_class}). The two of them are consistent in the sense that ASK classes have well defined PCA eigenvalues. However, the ASK classes are finer. We note that the scatter between these two purely spectroscopic classifications is much smaller than the scatter in the relationship with Hubble type (\\S~\\ref{ask_class_vs_morph}). The distribution of classes with redshifts is studied in \\S~\\ref{ask_vs_redshift}, and it reveals that the bluer classes contain galaxies of lower redshift, indicating that they are made of galaxies less luminous (and so smaller) than the red classes. The same preliminary analysis also suggests a trend for the red classes to be more clustered than the blue classes. All the above properties prove the consistency of the ASK classification. We have not found obvious contradictions between the physical properties of the classes, and the present understanding of galaxy properties. However, one should not forget the limitations of the analysis. The classification is not unique. We know that the borders between classes are not well defined, and that the actual number of classes is somewhat arbitrary (\\S~\\ref{final_imp}; \\S~\\ref{test_qbcd}). The galaxy spectra seem to have a continuous distribution of properties, and we still ignore the reasons why \\kmeans\\ puts the borders between classes where it does. Moreover, the classification rely on a number of reasonable but otherwise subjective hypotheses (e.g., the spectral bandpasses entering into the classification, or the normalization to the g-band; see \\S~\\ref{data_set} and Table~\\ref{tab1}). Alternative hypotheses would render classifications differing from ASK in a way difficult to foretell. All these caveats notwithstanding, the classes inferred by ASK have different spectra that reflect a systematic difference in the gas and stars present in the galaxies. Understanding the physical reasons causing the systematic differences between spectra would help understanding the classification itself. The study of the physical causes responsible for the observed diversity represents a major task that clearly goes beyond the scope of our introductory paper. However, a number of follow-up works dealing with the physical interpretation of classes are underway. As a result, some of the classes may need to be joined or split. For example, spectra of similar objects with different degrees of extinction may have ended up in different classes (remember that we do not correct for extinction to attain a purely empirical classification; \\S~\\ref{data_set}). Similarly, some of the classes may contain unidentified clusters. Sub-divisions can be achieved by applying \\kmeans\\ to selected spectral windows \\citep[e.g., the region around H$\\alpha$ may help us separating emission line galaxies according to their metallicity;][]{pet04}. Deriving the star formation history of the classes is fundamental to understand whether the differences between spectra are tracing different star formation histories, AGN activities, merging histories, or something else. (Is the ASK classification revealing some sort of evolutive sequence for galaxies?) Fortunately, inversion codes able to constrain the star formation history are available \\citep[e.g.,][]{cid05,toj09}, and we plan to use them. We are also carrying out a comparison between morphological types and spectroscopic types in a way that completes the introductory exercise in \\S~\\ref{ask_class_vs_morph}. We try to understand what causes the scatter in the relationship between morphology and spectroscopy. Is it the different characteristic time of evolution of morphological changes (on short timescales) and spectroscopic changes (on long timescales)? Is it the environment? The morphological classification will be based on the automatic procedure by \\citet{hue08} using support vector machines, which will allow us to afford comparing morphology and spectral type for a sizeable fraction of the SDSS/DR7 spectroscopic catalog. Work to derive the luminosity function for the classes is pending, i.e., to characterize the number density of galaxies of each luminosity and class. It is needed to quantify the tendency for high ASK classes to contain dwarf galaxies, as suggested in \\S~\\ref{ask_vs_redshift}. In addition to understanding the physical mechanisms responsible for the diversity among spectra, we foresee other applications of the ASK classification. It provides a crude but fast way of estimating some physical properties of a galaxy once its ASK ascription is known (\\S~\\ref{applications}). The classification is also useful as target selection. For example, ASK~6 is formed by Seyfert galaxies. This class provides an ideal homogeneous sample of some 5000 Seyferts with similar spectra for in-depth AGN studies \\citep[e.g., extending to low mass the relationship between supermassive black-hole mass and bulge mass; see][and references therein]{fer06}. The classification supplies classes of galaxies in the green valley (ASK~5). These targets allow us addressing the question of what characterizes a green valley galaxy, and one can do it in a statistically significant way. Is the green valley a short period during the life of any galaxy, or does it represent a genuine class of galaxies separated from the rest? The qualities assigned to each galaxy provide a simple way to find unusual objects. Low quality galaxies are outliers of the classification and, therefore, abnormal objects that deserve specific follow-up work. The average spectra of the classes can be used as template for redshift determinations. They represent a unique set comprising all spectral types. This application of ASK requires extending the template spectra to the UV, but this upgrade can be done in successive steps as outlined in \\S~\\ref{applications}. In order to facilitate these and possibly other applications, we have made the ASK classification freely available though the ftpsite {\\tt ftp://ask:galaxy@ftp.iac.es/}. We explain how it can be directly employed in SQL queries that use the CasJob facility of SDSS/DR7. We also provide it as ASCII csv tables suitable for uses external to SDSS. In addition, the template spectra are included. \\bigskip \\bigskip \\noindent{\\bf Acknowledgments.} We are indebted to J.~Betancort, I. G.~de la Rosa, and M.~Moles, that contributed with discusions on their area of expertise. Thanks to an anonymous referee we include the discussion on the degree of clustering of the classes at the end of \\S~\\ref{final_imp}. The work has been partly funded by the Spanish Ministries of science, technology and innovation, projects AYA~2007-67965-03-1, AYA~2007-67752-C03-01, and CSD2006-00070. Funding for the Sloan Digital Sky Survey (SDSS) and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, and the Max Planck Society, and the Higher Education Funding Council for England. The SDSS is managed by the Astrophysical Research Consortium (ARC) for the Participating Institutions. The Participating Institutions are the American Museum of Natural History, Astrophysical Institute Potsdam, University of Basel, University of Cambridge, Case Western Reserve University, The University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, The Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington. \\appendix" }, "1003/1003.3465_arXiv.txt": { "abstract": "We have used a precision calibrated photodiode as the fundamental metrology reference in order to determine the relative throughput of the PanSTARRS telescope and the Gigapixel imager, from 400 nm to 1050 nm. Our technique uses a tunable laser as a source of illumination on a transmissive flat-field screen. We determine the full-aperture system throughput as a function of wavelength, including (in a single integral measurement) the mirror reflectivity, the transmission functions of the filters and the corrector optics, and the detector quantum efficiency, by comparing the light seen by each pixel in the CCD array to that measured by a precision-calibrated silicon photodiode. This method allows us to determine the {\\it relative} throughput of the entire system as a function of wavelength, for each pixel in the instrument, without observations of celestial standards. We present promising initial results from this characterization of the PanSTARRS system, and we use synthetic photometry to assess the photometric perturbations due to throughput variation across the field of view. ", "introduction": "\\label{sec:intro} Multiband ground-based photometry plays a central role in a variety of forefront topics in astrophysics. Examples include mapping the expansion history of the Universe with type Ia supernovae, determining redshifts to galaxies and clusters using photometric redshifts, monitoring the spectral evolution of gamma ray bursts, and detecting and characterizing extrasolar planets with transits. A specific contemporary scientific objective that requires a clear understanding of system sensitivity across filter bands is using type Ia supernovae to probe the history of cosmic expansion. Imagine that the spectrum of supernova emission were concentrated in a narrow wavelength range, $\\delta \\lambda$. Observing supernovae at increasing redshift would move the center of emission into redder passbands. In order to precisely compare the luminosity distance vs.~redshift, we must clearly have a careful calibration of the {\\it relative} sensitivity across the passbands of interest, with a well-understood metrology foundation. Ascertaining the equation of state parameter $w=P/\\rho$ of the dark energy, and searching for any variation with redshift, will require supernova flux measurements with percent or better {\\it precision} ({\\it e.g.}~\\cite{w}). The SNLS collaboration has assessed photometric calibration as the dominant source of systematic uncertainty in their program to measure the equation of state parameter $w$ of the dark energy \\citep{SNLS09}. Another opportunity for exploiting high accuracy astronomical flux measurements will arise with the launch of the Gaia astrometric satellite mission, which should provide $\\sim10^5$ stars with distances accurate to better that 0.1\\% \\citep{gaia}. Providing a commensurate {\\it accuracy} for photometric measurements will allow for detailed comparisons between observed stellar luminosities and model atmospheres. These ground-based photometric measurements face a number of calibration challenges. One is knowing the instrumental sensitivity as a function of wavelength. Another is accounting for absorption and scattering in the Earth's atmosphere. The traditional methodology of photometric calibration uses some combination of (1) observations of standard sources at the same airmass as the program objects, or (2) measurements of extinction coefficients in the passbands of interest. In this approach, the spectra of celestial sources constitute (often implicitly) the standard for the measurement. These spectrophotometric standards are either based on ground-based blackbody measurements made decades ago \\citep{Hayes75,meg95}, or in the case of DA white dwarfs, on theoretical model atmospheres ({\\it e.g.} \\cite{Holberg06} ). Since broadband CCD photometry of celestial sources takes an integral of photon flux across a broad optical passband ($u,g,r,i,z$ or $y$, for example), it is an ill-posed problem to infer the system sensitivity function vs. wavelength from broadband photometric data alone, especially through a variable atmosphere. The determination of the effective system throughput has often been done in a piecemeal fashion, taking benchtop measurements of filter transmission, optical catalog values for reflectivity and glass transmissions, and vendor descriptions of typical detector quantum efficiency (QE). The cumulative systematic errors in this procedure limit the precision that can be achieved, and this approach lacks resolution at the pixel scale. Attempting to adjust passbands determined in this fashion to match observations to synthetic photometry is ambiguous: should the passbands be broadened, shifted up or down in wavelength, or should a grey multiplicative scaling be applied? In order to address these concerns, and to push towards an improvement in photometric precision and accuracy, \\cite{Calib06} advocated breaking the spectrophotometric calibration problem into two distinct measurements: 1) the determination of instrumental sensitivity, using laboratory-calibrated detectors as the fundamental metrology standard, and 2) measuring directly the optical transmission of the atmosphere. We will focus here on our implementation of a technique for measuring, {\\it in situ}, the relative throughput of the entire PanSTARRS apparatus, relative to a calibrated photodiode. We can forgo the use of a celestial calibration {\\it source} in favor of a well calibrated {\\it detector} as the fundamental metrology reference for astronomical photometry. National Institute of Standards and Technology (NIST) has calibrated Silicon photodiode quantum efficiencies at the 10$^{-3}$ level. This is an order of magnitude more precise than any celestial spectrophotometric source, either empirical or theoretical. In fact, the primary SI metrology reference for electromagnetic flux is now detector-based rather than source-based \\citep{NIST}. There is ongoing progress at NIST in extending both the wavelength range and accuracy of detector calibration for metrology applications. We assert that no celestial spectrophotometric source currently offers a photon spectral distribution that is known at the $10^{-3}$ level, and the prospect of achieving this precision is one of the main attractions of the method described here. This paper presents promising initial results from an integrated measurement of the total system throughput. We use full-pupil illumination with monochromatic calibration light to measure the entire optical train of the apparatus, including the mirror, corrector optics, filter and detector. The spatial resolution is at the pixel scale, with 1 nm spectral resolution. An earlier realization of this technique, with the Mosaic imager on the 4 meter Blanco telescope at CTIO, was described in \\cite{CTIOpaper}. The determination of and correction for variation in atmospheric transmission is a significant challenge. A companion paper \\citep{Burke10} describes progress in the precise determination of atmospheric transmission. We also refer the interested reader to \\cite{Calib06} and \\cite{Atmos} for a discussion of this issue. The Pierre Auger collaboration has implemented \\citep{Auger10} a comprehensive suite of atmospheric monitoring instruments for the calibration of optical transmission in atmospheric fluorescence detection of high energy cosmic rays, and we advocate taking a similar approach for optical and infrared astronomical observations. The merits of achieving improved photometric accuracy are spelled out in \\cite{Kent09}, and there is considerable work under way to achieve improved accuracy and precision. Many of these are described in \\cite{ASP}. The ACCESS project \\citep{ACCESS} plans to conduct precise spectrophotometry from a sounding rocket. \\cite{Bohlin07} describes spectrophotometric measurements at the 1\\% level using HST. Adelman and colleagues \\citep{Adelman07} are pursuing ground-based precision spectrophotometry. \\cite{SLR} describe the use of the stellar locus in color-color space for precise color determinations. \\cite{Albert09} discussed the merits of a satellite-based calibration source, which would provide an opportunity to determine both atmospheric transmission and apparatus throughput in the same measurement, providing the satellite could be accurately tracked by the telescope. The Sloan Digital Sky Survey has used monochromatic light to measure and monitor the spectral sensitivity function of the SDSS camera when removed from the telescope \\citep{SDSS10}. Our approach uses full aperture illumination of the entire optical system rather than measuring only the camera's response. We are not currently attempting to establish an absolute calibration of the PanSTARRS system, in units of photons cm$^{-2}$ s$^{-1}$ nm$^{-1}$, but rather to establish the {\\it relative} system throughput across wavelength, in arbitrary units. A specific example of the importance of understanding the relative photometric zeropoints across filters is when type Ia supernovae are used to map out the history of cosmic expansion. A complete understanding of the effective system passband, and the associated zeropoints, is essential. Type Ia supernova cosmology is one of the many science objectives planned for the PanSTARRS survey. A single measurement of absolute detection efficiency, at one wavelength, would suffice to place the data presented here onto an absolute flux scale. An approach to full-aperture determination of absolute detection efficiency is described by the Auger collaboration in \\cite{Auger04}. The calibration technique described here can be used in conjunction with more traditional photometric calibration methods. Using a diversity of calibration methodologies provides an opportunity to assess consistency and to quantify potential sources of systematic error. This approach is not limited to imaging instruments. Spectrophotometric sensitivity functions for dispersive instruments can be acquired in a similar fashion. A specific motivation for our calibration program is the requirement of knowing the system throughput well enough to perform high confidence color determination for the PanSTARRS survey. As described below, this technique also allows us to monitor filter transmission curves (to check for variation due to hygroscopic effects, for example), and any other changes in instrumental performance. This paper first presents an overview of the experimental method and apparatus in Section \\ref{sec:methods}, and the observations we obtained in Section \\ref{sec:observations}. Data processing is outlined in Section \\ref{sec:processing}, and results are shown in Section \\ref{sec:results}. Section \\ref{sec:systematics} presents a preliminary assessment of potential sources of systematic error, followed by our conclusions in Section \\ref{sec:discussion}. ", "conclusions": "\\label{sec:discussion} We have shown the ability to obtain a succession of monochromatic calibration images that exhibit short term repeatability at the level of a few parts per thousand. Our long term goal is to provide, for each survey image and associated photometric catalog, the effective passband through which the image was acquired. This would include both the instrumental response function at the location of the source, as well as the estimate of atmospheric transmission along the appropriate line of sight. These results encourage us to further pursue this technique to both measure and monitor system sensitivity over the course of the PanSTARRS survey. We also intend to use the system described here to measure the dependence of the response function on CCD device temperature, in the phonon-assisted photoconversion regime at wavelengths near 1 $\\mu$m. Another interesting measurement would be to compare the response function before and after mirror cleaning or dust deposition events. We can also use these instrumental sensitivity data in conjunction with measurements of atmospheric transmission to compare synthetic photometry with observations, as described in \\cite{Calib06}, across different regions of the focal plane. Table \\ref{tab:sowhat} is an initial exploration of this approach. We took model atmospheres from \\cite{Kurucz96} and converted them to photon spectral energy (PSED) distributions $\\Phi_\\gamma(\\lambda)$. We then integrated\\footnote{We have omitted the effect of atmospheric attenuation since our interest here is only to explore the effect of focal plane position-dependence of the sensitivity function.} these photon distributions over the seven distinct annular $r$ band response functions, to assess the photometric perturbation that would arise from the observed radial variation in the sensitivity functions. We computed for each radial annulus $a$ the ``raw'' synthetic magnitude $$ m(a)_{raw}=\\int T(a,\\lambda) \\Phi_\\gamma(\\lambda) d\\lambda,$$ and also a normalized magnitude that corresponds to constructing a photon-flat $$ m(a)_{normalized}= \\frac{\\int T(a,\\lambda) \\Phi_\\gamma(\\lambda) d\\lambda} {\\int T(a,\\lambda) d\\lambda}, $$ where $a$=1,2...7. The normalized magnitudes correct for the radial variation in the integrated area under the sensitivity curve. Table \\ref{tab:sowhat} shows the resulting magnitude differences, $\\Delta m=m_{edge}-m_{center}$, between the edge ($a$=7) and the center ($a$=1) of the PanSTARRS field, for different stellar types. The second column in Table \\ref{tab:sowhat} shows that the raw flux at the field edge ranges from 3.7\\% to 4.6\\% higher than the value at the center, with the variation being monotonic with decreasing stellar temperature. This is consistent with the sensitivity function behavior shown in Figure~\\ref{fig:ratioplot}: stars at the edge appear brighter, and red stars more so than blue ones. If we compare instead the ``normalized'' magnitudes (the third column in Table \\ref{tab:sowhat}) the flux excess diminishes by a factor of 5 to 20, depending on stellar temperature. However, even after flux normalization at each radial annulus by $\\int T(a,\\lambda) d\\lambda$, a scalar quantity that implicitly weights each wavelength equally, the third column in the Table still shows up to $\\sim$1\\% perturbations in stellar photometry ({\\it e.g.} $\\Delta$m$_{normalized}$(O5V) - $\\Delta$m$_{normalized}$(M6V)=0.0094) that depend on both stellar temperature and separation across the field. \\begin{table}[htdp] \\caption{Influence of Observed $r$ band Center-to-Edge Variations on Synthetic Stellar Photometry. } \\begin{center} \\begin{tabular}{ccc} \\hline \\hline Stellar & Edge-Center & Edge-Center \\\\ Type & ~~~$\\Delta$m$_{raw}$~~~& $\\Delta$m$_{normalized}$ \\\\ \\hline O5V & -0.0367 & ~0.0068 \\\\ O7V & -0.0368 & ~0.0068 \\\\ O9V & -0.0368 & ~0.0067 \\\\ B1V & -0.0371 & ~0.0065 \\\\ B3V & -0.0374 & ~0.0061 \\\\ B5V & -0.0376 & ~0.0059 \\\\ B8V & -0.0379 & ~0.0056 \\\\ A1V & -0.0385 & ~0.0050 \\\\ A3V & -0.0388 & ~0.0047 \\\\ A5V & -0.0393 & ~0.0042 \\\\ F0I & -0.0391 & ~0.0044 \\\\ F0V & -0.0391 & ~0.0044 \\\\ F2V & -0.0411 & ~0.0025 \\\\ F5V & -0.0417 & ~0.0019 \\\\ F8V & -0.0420 & ~0.0015 \\\\ G2V & -0.0425 & ~0.0011 \\\\ G5I & -0.0425 & ~0.0011 \\\\ G5V & -0.0426 & ~0.0009 \\\\ G8V & -0.0429 & ~0.0007 \\\\ K0V & -0.0434 & ~0.0001 \\\\ K4V & -0.0450 & -0.0015 \\\\ K7V & -0.0460 & -0.0025 \\\\ M2V & -0.0457 & -0.0021 \\\\ M4V & -0.0461 & -0.0026 \\\\ M6V & -0.0461 & -0.0026 \\\\ \\hline \\end{tabular} \\end{center} \\label{tab:sowhat} \\end{table}% We conclude from this synthetic photometry exercise that we can expect half-percent level $r$ band zeropoint shifts between the center and the edge of the field, even after flat-fielding, for a given stellar type. Looking across stellar types (from O to M class stars) we can expect systematic zeropoint differences at the percent level, between the center and the edge of the field. As outlined in \\cite{Calib06}, knowing the system throughput for each pixel and for each wavelength will allow us to account for these effects in detail. This suggests that pushing the photometric precision beyond the 1\\% level will likely require a more sophisticated flat-fielding approach, that takes full advantage of the system's measured spectral response function (as well as a determination of atmospheric transmission). Also, the determination of ``illumination corrections'', by rastering a source around the focal plane and requiring consistent results, will be PSED-dependent at the percent level. We are now in a position to quantitatively assess various flat-fielding schemes, such as performing a joint analysis of photometric data from all bands, that take into account the position-dependence (and also, if needed, time-dependence) of the system's response function. Finally we note that a full-aperture calibration taken of a distant source of known radiance, even at just one wavelength, would suffice to determine the single overall multiplicative term we need to extend these measurements to an absolute calibration of the PanSTARRS system. This would require appropriate knowledge and correction for atmospheric transmission, but the single wavelength could be judiciously chosen (such as $\\lambda \\sim 808$~nm, also convenient for laser diode sources) to coincide with good detector QE as well as high and stable atmospheric transmission. {\\it Facilities:} \\facility{PanSTARRS}" }, "1003/1003.5596_arXiv.txt": { "abstract": "The SCUBA polarized $850$ $\\mu$m thermal emission data of the region OMC-2 in Orion A are added to and homogeneously reduced with data already available in the region OMC-3. The data set shows that OMC-2 is a region generally less polarized than OMC-3. Where coincident, most of the $850$ $\\mu$m polarization pattern is similar to that measured in $350$ $\\mu$m polarization data. Only $850$ $\\mu$m polarimetry data have been obtained in and around MMS7, FIR1 $\\&$ FIR2, and in the region south of FIR6. A realignment of the polarization vectors with the filament can be seen near FIR1 in the region south of OMC-3. An analysis shows that the energy injected by CO outflows and H$_{2}$ jets associated to OMC-2 and OMC-3 does not appear to alter the polarization patterns at a scale of the $14''$ resolution beam. A second order structure function analysis of the polarization position angles shows that OMC-2 is a more turbulent region than OMC-3. OMC-3 appears to be a clear case of a magnetically dominated region with respect to the turbulence. However for OMC-2 it is not clear that this is the case. A more in-depth analysis of five regions displayed along OMC-2/3 indicates a decrease of the mean polarization degree and an increase of the turbulent angular dispersion from north to south. A statistical analysis suggests the presence of two depolarization regimes in our maps. One regime including the effects of the cores, the other one excluding it. ", "introduction": "It is generally believed that magnetic fields play an important role in star formation but the evolution of their role from the scale of molecular clouds to that of young stellar objects is still not well understood. One of the best methods for their study is polarimetry of the continuum radiation emitted by aligned dust grains at far-infrared (FIR) to submillimeter (submm) wavelengths \\citep[e.g.][]{dot00,hil99,mat09}. Due to fast rotating mechanisms, elongated grains pervading the dense ISM should mostly have their longer axis preferentially oriented perpendicular to the local magnetic field. Recent advances on grain alignment theory can be found in reviews by \\citet{laz03} and \\citet{laz07}. The Orion ``Integral-Shaped filament'' (ISF) was mapped at $850$ $\\mu$m by \\citet{joh99}. It contains the well-studied OMC-1 region behind the Orion nebula at $\\approx 414$ pc \\citep[see][]{men07}, and the two filamentary structures OMC-2 and OMC-3 located $\\approx 15'$ and $\\approx 25'$ north of OMC-1, respectively. Submm continuum and polarimetry observations in OMC-1 were intensively analysed and discussed by \\citet{kee82}, \\citet{hil84}, \\citet{ait97}, \\citet{lis98}, \\citet{rao98}, \\citet{sch98}, \\citet{val99}, \\citet{cop00} and \\citet{val07}. The counterpart of these studies in OMC-2 and more particularly in OMC-3 were discussed by \\citet{chi97} with $1.3$ mm dust emission observations in both regions; by \\citet{mw00} and \\citet{mwf01} with $850$ $\\mu$m polarimetry in OMC-3; by \\citet{hou04} with polarimetry at $350$ $\\mu$m in both filaments; and more recently by \\citet{mat05} with $1.3$ mm high resolution observations in the MMS6 protostellar core in OMC-3. The 850 $\\mu$m polarization pattern of OMC-3 measured with SCUBA and studied by \\citet{mwf01} shows that a helical magnetic field \\citep[see][]{fp00d} could thread the filament. The change in orientation between the polarization data and the filamentary axis to the south of OMC-3 could be accounted for by either a bend of the filament or the presence of a second filament, oriented roughly orthogonal to the primary ISF. If the latter interpretation is correct, then polarization vectors realigning with the filamentary axis are expected south of OMC-3. If the filament continues to bend, then any orientation of vectors to the filament axis is possible. Our initial motivation for this project was to measure the 850 $\\mu$m polarization pattern of OMC-2 with the aim to better understand the structure of the magnetic field south of OMC-3. Then, with the availability of 350 $\\mu$m polarization data by \\citet{hou04} and estimates of the orientation of the magnetic field relative to the Line-of-Sight (LOS) at several places along the filament, a global comparison with the 850 $\\mu$m polarization vectors observed in both regions can be made. Abundances of several molecular species and clumping were also observed in these regions \\citep[e.g.][]{bat83,cas95,chi97} suggesting an evolutionary effect from north to south along OMC-3/2. In contrast, \\citet{tak08} show that some intermediate-mass (IM) star-forming objects in OMC-3 can be at a more evolved stage than some less active IM star-forming objects in OMC-2. Observations of H$_{2}$ jets \\citep[e.g.][]{sta02,par06} and of molecular outflows \\citep[e.g.][]{wil03} allowed the detection of some progenitors along the chain of active cores embedded in the filament with relatively good certainty. Thus we can compare their respective orientations on the plane of the sky (POS) with the direction of polarization vectors in order to statistically study the impact of outflow and jet activity on the region. The distribution of the offsets in position angles (P.A.s) obtained in this way can help to test detectability of turbulence in the framework of our observations. In addition, a coherent structure function analysis method was recently put forth by \\citet{hil09}. Combined with other analysis tools, this promising approach should help to make fruitful comparisons with available and future models. As a consequence we can use this method along the OMC-2 and OMC-3 filamentary molecular clouds and compare it with steady state and turbulent magnetic field models. All in all, these approaches should help to better understand the impact of magnetic fields with respect to turbulence and gravity in star-forming processes and cloud evolution. In this work we present 850 $\\mu$m SCUBA polarization data of OMC-2. Observations and data reduction techniques are presented and discussed in $\\S$ \\ref{OBS}. Results are shown and analyzed in $\\S$ \\ref{RES}. A discussion and comparisons with models follow in $\\S$ \\ref{DISCU}. ", "conclusions": "\\label{DISCU} \\subsection{A Possible Impact of Jets and Outflows on Submm Polarization Patterns} \\label{joimpact} Based on the statistical tests presented in section \\ref{jetsout}, we conclude that no correlation is evident between the relative orientation of jets or outflows and polarization vectors on the POS. This conclusion is similar to that of \\citet{men04} in the Taurus molecular cloud complex, an active star-forming region containing no massive stars. In their study, using absorption polarimetry data, these authors show that T Tauri stars as a group are apparently oriented randomly with respect to the local magnetic field. This indicates that even if the cloud's magnetic field is dominant at large scale, its influence largely decreases on the much smaller scale of individual objects unless the orientation of these individual objects has changed since birth. The situation seems to be the same in the OMC-3/2 regions. This scenario also implies that if turbulence generated by the outflows is able to misalign grains in the envelopes of the cores, this effect can not be seen on the POS with the resolution of the JCMT. For a distance $d \\approx 414$ pc, this means that, independently of their orientation relative to the line of sight, the energy injected by CO ouflows and H$_{2}$ jets into the clouds appears not to have any impact on the polarization patterns which are observed with the presence of jets and outflows on scales of $\\approx 7700$ AU. In addition, detections of polarized CO emission are now available. A good introduction about the subject and a summary of some results is given by \\citet{for08}. Some detections were established by \\citet{gir99} in NGC 1333 IRAS 4A \\citep[see also][]{gir06}, and by \\citet{gre99} toward the Galactic center and in the molecular clouds S140 and DR21. When the optical depth, $\\tau$, and the spatial distribution of the gas and of the magnetic fields are favorable, it is possible to detect CO polarization vectors perpendicular or parallel to the magnetic field \\citep[see][]{kyl83}. Except in the ``2 pc ring'' where the optical depth $\\tau$ is relatively high, positive detections found in the other regions are consistent with orientations of magnetic fields inferred by submm dust polarimetry. Thus, at the scales of their spatial resolution, these works suggest that the energy injected by jets and flows should have no substantial impact on the net polarization produced by dust grains aligned in the clouds. This appears to be consistent with the OMC-3 north region where the well-defined polarization pattern suggests that whatever the inclination angle relative to the line of sight of the jets and flows, as well as their orientation on the POS, the energy injection rate has no influence on the alignment of dust grains seen at the scale of the observations. In MMS7 the outflow is oriented perpendicular to the 850 $\\mu$m polarization vectors covering this area. This means that if the situation is the same as in NGC 1333 IRAS 4A \\citep{gir99}, polarized CO emission vectors should be seen perpendicular to these vectors. In MMS8, 9 \\& 10, the analysis of the situation is different since in this part of the sky the polarization pattern orientation may be due to another superimposed crossing filament or by a bend of the filament \\citep[see][]{mwf01}. Since the association of jets and outflows with their probable progenitors is not always certain because of overlapping effects on the sky, a better understanding of the magnetic field structure producing the submm polarization pattern observed in the southern part of OMC-3 and in the northern part of OMC-2 added to polarized CO emission observations could help to constrain these associations. The analysis of the situation in OMC-2 is not exactly the same since some of the high intensity regions are devoid of detections. However, in regions where polarization vectors have a S/N ratio such that $p/\\sigma_{p}>3$ the general conclusion about the relative orientation of the jets/outflows with the submm polarization vectors discussed above is still valid. \\subsection{Turbulent Angular Dispersions} \\label{TAD} The second-order structure function of the polarization angles is defined as the average of the squared difference between the polarization angle measured at two points separated by a distance $l$ \\citep[see equation (5) given by][]{falceta08}. Once applied to a grid of pixels containing polarization position angle information, a fit to the square root of this function, namely, the Angular Dispersion Function (ADF), gives a method to estimate the turbulent contribution to the total angular dispersion. First applications of the method and results obtained in regions OMC-1, DR21 Main and M17 are given by \\citet{hil09}. Figure \\ref{DM2OMC23} shows the ADF obtained after application of the method on the 18.6$\\arcsec$ square pixel grid to the maps of OMC-3 and OMC-2 shown in Figure \\ref{polmap}. The dashed line shows the maximum dispersion value, $b_{\\rm max}=\\frac{180^{\\circ}}{\\sqrt[]{12}}$, that would be reached in the case of a purely random polarization angle distribution \\citep[see][]{ser62}. All the maps obtained in several regions of the filament are not shown here but the turbulent angular dispersion fitting parameter, $b$, was estimated for each region displayed in Table \\ref{subregions}. Estimates of $b$ obtained with the correlation method for OMC-2/OMC-3, OMC-3, OMC-2 and the five subregions selected along the filament (see discussion in section \\ref{hole}) are displayed in column (5) of Table \\ref{subregions}. These values can be compared to the angular dispersion values, $s(\\theta)$ obtained with the square root of the variance of the data sets displayed in column (4) of the Table. Figure \\ref{bvsdisp} shows the variation of $b$ with $s(\\theta)$ for the five subregions. The dotted-line expresses equality between these two quantities. A linear fit applied to the data and shown by the dashed-line returns the relation, $b = (0.60 \\pm 1.24) \\rm deg + (0.75 \\pm 0.04) \\times s(\\theta)$, meaning that the angular dispersion values estimated with the correlation method are statistically about $25\\%$ lower than the ones calculated about the mean polarization position angle. The ratio of turbulent to large scale magnetic field is estimated by using equation (7) given by \\citet{hil09}: \\begin{equation} \\label{equ7} \\frac{^{1/2}}{B_{0}}=\\frac{b}{\\sqrt[]{2-b^2}} \\end{equation} and values are displayed in column (6) of Table \\ref{subregions}. In cases where the turbulent component of the field is very small compared to the nonturbulent component, $B_{t} \\ll B_{0}$, \\citet{hil09} show that the uniform component of the field can be approximated by the following equation: \\begin{equation} \\label{equ8} B_{0} \\simeq \\sqrt[]{8 \\pi \\rho} \\frac{\\sigma_{v}}{b} \\end{equation} We use this approach by assuming a density of $10^{4} \\rm cm^{-3}$ and a mean molecular weight of 2.3. Estimates of $B_{0}$ are given as a function of $\\sigma_{v}$ in the last column of Table \\ref{subregions} except for regions MMS1 to MMS7 and FIR3 to FIR5 since line width measurements from H$^{13}$CO$^{+}$ $J = 3 \\to 2$ were made in OMC-3 MMS6 and OMC-2 FIR4 by \\citet{hou00}. We use the values shown in their Table 1 to directly make estimates of $B_{0}$. The line width measurements are reported in column (7) of Table \\ref{subregions}. Given the smooth and well defined polarization patterns observed in regions MMS1 to MMS7 and MMS8 to MMS10 we are confident that the method is suitable to estimate, within a factor of a few, the global mean field component. In the three other subregions, the higher dispersion polarization patterns suggest that the relation $B_{t} \\ll B_{0}$ may not be as well satisfied. However, all in all, the results shown in Table \\ref{subregions} are a first step to intercompare the regions. The estimates of the turbulent angular dispersion components, $b$, displayed in columns (5) of Table \\ref{subregions} for OMC-3 and OMC-2 are all higher than those estimated by \\citet{hil09} in their Table 1. We point out that the turbulent component dispersion estimates obtained in OMC-3 are of the same order than those found in regions OMC-1 and M17. To the contrary, the high value obtained in OMC-2 distinguishes this region from OMC-3, OMC-1, M17 and DR21(Main). \\citet{hil09} find that the dispersions obtained about the mean field orientation are about a factor of 3 times higher than the ones estimated by using the ADF fitting method. Doing the same comparisons we find a factor of about 2 to 1.5 for regions OMC-3/OMC-2, OMC-2 and OMC-3. Abundances of molecular species and clumping studied by \\citet{bat83}, \\citet{cas95} and \\citet{chi97} suggest an evolutionary effect from north to south along OMC-2/3. The two regions could have different ages and OMC-3 could be younger than OMC-2. This could explain why OMC-2 and OMC-3 are so different from the point of view of submm polarization data. On the other hand, the situation is not so clear since Takahashi et al. (2008) show that some intermediate mass objects can be at a more evolved stage in OMC-3 than in OMC-2. Another explanation could invoke some effects of the radiation field. Since OMC-3 is at an higher distance to OMC-1 than is OMC-2, the local ISRF dominated by the bright stars located in the Nebulae in front of OMC-1 could be more efficient to shape the OMC-2 region than the OMC-3 region. The two regions, OMC-2 and OMC-3, could have the same age but the erosion produced by the local ISRF could be stronger on OMC-2 than on OMC-3. \\subsection{Ordered vs Turbulent Magnetic Field Components along OMC-2 and OMC-3} \\label{compar} \\subsubsection{Depolarization and Turbulence along the Filaments} \\label{depol} Table \\ref{subregions} suggests that depolarization is present toward the OMC-2 sources as much as toward OMC-3. To understand how the decrease of polarization with the increase of intensity could be related to turbulent arguments, we show in Figure \\ref{bvsslopes} the variations of the power indices of the $p-I$ relation, $\\gamma$, with the turbulent angular dispersion components, $b$, for the five subregions in Table \\ref{subregions}. We find no specific correlation meaning that the decrease of the polarization degree observed appears independent of turbulent flows or turbulent effects that could be present into the densest regions of the cloud. On the other hand, the variation of the the mean polarization degrees, $

$, with the turbulent angular dispersion estimates, $b$, of the regions considered in Table \\ref{subregions} and shown in Figure \\ref{pmeanvsb} suggests an anti-correlation between the two parameters. If this trend is real, it would mean that, independently of the ordered magnetic field structure component, the mean polarization degree observed in a given region could be a function of the turbulent magnetic field component. To test this hypothesis we did a linear fit of the form $

= + c_{1} \\times b$, to the eight points displayed in Figure \\ref{pmeanvsb}. The results are $=4.2 \\% \\pm 0.2\\%$ and $c_{1}=-0.07 \\% \\rm deg^{-1} \\pm 0.01 \\% \\rm deg^{-1}$, where the fit is constrained by taking into account the errors on $p$. The errors are estimated by dividing the standard deviation of each data set by the square root of the number of data. In addition, a fit of the variations of $b$ with $

$ (not shown here) of the form $b = b_{0} + c_{2} \\times

$, taking into account the uncertainties on $b$, gives the results $b_{0}=54.2^{\\circ} \\pm 1.5^{\\circ}$ and $c_{2} = -11.8 \\rm deg \\%^{-1} \\pm 0.6 \\rm deg \\%^{-1}$. Given the errors on $p$, the first fit constrains reasonably the variations of $

$ with $b$. Within the uncertainties the intersect of the fit with the abscissa axis gives a value of $b_{0}$ which is hardly consistent with the expected value, $b_{\\rm max}$ of about $52^{\\circ}$, that would be found in the case of a purely random polarization position angle distribution \\citep[][]{ser62}. On the contrary, the second fit returns a value $b_{0}$ relatively close to that of $b_{\\rm max}$ suggesting that if such a region was observable the mean polarization degree could be about zero. To test if the trend observed in OMC-2/OMC-3 is consistent with the polarization properties of OMC-1 we use the turbulent dispersion parameters, $b$, estimated by \\citet{hil09} in their Table 1 and Figure 3 of \\citet{vai08} where median polarization ratios are displayed for different wavelengths and for several regions. The mean polarization percentage observed at 350 $\\mu$m by \\citet{hou04} is $

= 2.7 \\%$ with a dispersion $\\sigma_{p}=1.4\\%$. Using Figure 3 shown by \\citet{vai08} and a wavelength ratio of 1.3, a mean polarization of about 3.5 $\\%$ should be representative of the mean degree of polarization observed at 850 $\\mu$m in OMC-1 with 12$^{\\arcsec}$ resolution. Neglecting the effects that could produce a slightly lower resolution of 14$^{\\arcsec}$ and with the assumption that the turbulent ratios derived with 350 $\\mu$m and 850 $\\mu$m polarization maps would be similar this yields a point $(b=8.3\\pm0.3,

\\approx 3.5 \\%)$ that appears to be consistent with our fit. A similar trend was observed along the Pipe Nebula by \\citet{alv08} with $R$-band visible polarization data collected for about 12000 stars. In this case the dispersion of several subsets of polarization position angles is used to estimate the turbulence of regions having a mean density $n(\\rm H_{2}) \\approx 10^{3}$ cm$^{-3}$ and is compared to the mean polarization degrees of the subsets of data. The regions probed in the Pipe Nebula have densities about an order of magnitude lower than the density considered in our calculations where submm data would mostly probe the inner fields embedded in the filaments. The estimates of the mean field strength displayed in Table \\ref{subregions} for regions MMS1 to MMS7 and FIR3 to FIR5 are of the same order. This result dismisses the hypothesis that the dust grain alignment efficiency could decrease with the magnetic field strength. In addition, mean inclination angles of the magnetic field from the LOS were estimated by \\citet{hou04} to vary from about $73^{\\circ}$ to about $80^{\\circ}$ in three dense regions displayed along OMC-2/OMC-3. If these inclination angles are representative of the mean inclination angle of the magnetic field along the filaments and if the grains are aligned with roughly the same efficiency everywhere then the degree of polarization would not be very sensitive to this parameter. These two points do not dismiss the possibility that the mean polarization degree of some regions could be a function of the level of turbulence in the region, but nor do they rule out the possibility of complex structures of an ordered magnetic field component. In cases where the large scale magnetic field dominates over the turbulence it could be that some superposition effects along the LOSs would produce a decrease of the polarization degree with an increase of the polarization position angle dispersion over some regions. These aspects are discussed in the following sections. \\subsubsection{Effects of Steady-State Magnetic Field Models} \\label{steadystate} The Turbulent-to-Mean magnetic field strengths ratios displayed in Table \\ref{subregions} imply intensities of the turbulent magnetic component about 6 to 7 times lower than the mean magnetic field component into the OMC-3 filaments. This ratio supports the steady state magnetic field approach proposed by \\citet{fp00d} and applied to the OMC-3 region as discussed by \\citet{mwf01}. In the work of \\citet{fp00d} the maximum polarization degree is calibrated with submm observations and dust grains are aligned perpendicular to the magnetic fields. The combination of toroidal with poloidal magnetic fields components leads to superposition effects along the same line-of-sight. This geometrical effect decreases the net polarization on the POS and renders uniform magnetic fields indistinguishable from a helical field geometry. The model can reproduce the $r^{-2}$ density profile\\footnote{Here $r$ is the radial distance, in cylindrical coordinates.} observed by \\citet{joh99} and could explain the depolarization observed along the spine of the filament. A similar approach was followed by \\citet{gon05} with a focus on molecular cloud cores and shows that tangling of the magnetic field orientations by the effects of gravity could produce polarization maps showing an significant angular dispersion from the mean uniform field. Since it is not clear that OMC-2 is a clear case of a magnetically dominated region with respect to the turbulence, it would be interesting to test under which conditions steady-state magnetic models could reproduce the polarization properties observed in this region. \\subsubsection{Variation of MHD models} The statistical results shown in Figure \\ref{pmeanvsb} discussed above in section \\ref{depol} suggest that turbulence could be a parameter regulating the mean degree of polarization observed in a given polarization map. One open question would be to understand to what extent MHD models could reproduce such a decrease of polarization. \\citet{falceta08} present results for turbulent, isothermal, three-dimensional simulations of sub/supersonic and sub/super-Alfv\\'enic cases. Dust grains are assumed perfectly aligned and perpendicular to the magnetic field. Depolarization is due to the dispersion increase of the polarization position angles and is a function of the MHD regime considered as well as of the orientation of the initial uniform magnetic field with respect to the LOS. The dust total intensity is assumed to be proportional to the column density. The high resolution of the simulations provides less homogeneous magnetic field structures and higher density contrasts than previous models. On the other hand, \\citet{cho05} show that under peculiar conditions depolarization could occur if grains embedded in dark clouds are aligned by radiative torques (RATs) such that their long axis is perpendicular to the magnetic field. In their model, the nonturbulent field is about 2 times stronger than the fluctuating magnetic field, a condition encountered in OMC-2 (see Table \\ref{subregions}) and the ordered component is assumed to be uniform and in the POS. This condition may exist along the length of the OMC-2 and OMC-3 filaments \\citep[][]{hou04}. Additionally, the results proposed by \\citet{cho05} are valid for clouds without embedded massive stars, a condition valid in OMC-2/3 where forming stars are of intermediate \\citep[see][]{tak08}, or lower mass. Two interesting extensions to the \\citet{cho05} work were proposed by \\citet{pel07} and by \\citet{bet07} but the analysis of the simulations is focused on the effects of RATs on depolarization rather than on the effects of the turbulent regimes on depolarization. Alignment by RATs is not considered by \\citet{falceta08} but an anti-correlation between the polarization degree and the column density, with exponent $\\gamma \\sim -0.5$ is predicted, due to random cancellation of polarization vectors along the LOS. This value is close to some values obtained on larger scales in OMC-2 and OMC-3 under conditions which are discussed in more detail in the following section. Comparisons done by \\citet{falceta08} between four different MHD regimes show a degeneracy between the Alfv\\'enic Mach number and the angle between the mean magnetic field and the LOS. They discuss the effects of different resolutions on the structure function of the polarization angle and the applicability of these structure functions to the determination of turbulent cut-off scales. A comparison of the $14^{\\arcsec}$ resolution 850 $\\mu$m P.A. histogram displayed in Figure \\ref{anghistot} with the results of Figure 4 of \\citet{falceta08} would reject the presence of super-Alfv\\'enic modes in OMC-3 and probably in OMC-2 as well. This inference would be consistent with the lack of correlation between jets/outflows and polarization vectors discussed in section \\ref{joimpact}. But, as suggested in other regions by \\citet{pad04}, this would not rule out the presence of super-Alfv\\'enic modes at smaller scale, in cores of about 0.25 pc in size in the OMC-2/3 regions for a distance to the clouds of about $d = 414 $ pc \\citep[see][]{men07}. Finally, because of the resolution of our data, a direct comparison of our results with the structure function of the polarization position angle calculated for different MHD regimes is not done. \\subsubsection{Variations of Power Index with Column Density Structure} \\label{mhdrats} Figure \\ref{pmeanvsb} illustrates that the mean polarization degree of any region could be controlled by the MHD regime but it does not give any indication about the possible effect of the MHD regime on the level of depolarization. Figure \\ref{bvsslopes} shows an absence of correlation between the power index of the $p-I$ relation and the turbulence parameter, $b$, however, and suggests that one or more phenomena other than turbulence should be considered. To inform this discussion, we consider possible variations in column density structure and the power law index, our best quantifier of depolarization. Estimates of the power indices of the five subregions from Table \\ref{subregions} are shown as a function of the column density contrast, $CDC=(Flux_{\\rm max} - Flux_{\\rm min})/Flux_{\\rm max}$, in Figure \\ref{slopesvsdc}. On a statistical basis, the variation of the power index with the $CDC$ is estimated by considering the several maps obtained by masking high intensity pixels above specific cutoffs. These variations of the power indices with the $CDC$s are shown by the solid lines in Figure \\ref{slopesvsdc}. The highest $CDC$ values are derived from our original maps (Figure \\ref{polmap}). To establish the power index of lower column material, for which we assume lower fluxes to be a proxy, we methodically masked the highest pixels by using a step of 33 $mJy$/18.6$^{\\arcsec}$ in column density and then recalculated the $CDC$ and power index of the resulting $p-I$ relation. By repeating this process, we are able to assess the impact of lower and lower column material on $\\gamma$. We stopped at the level of $26\\%$ of the peak in OMC-3 (71 of 116 pixels were used in the process) and 30$\\%$ of the peak in OMC-2 (108 of 135 pixels). Below these values, divergences were observed in the estimates of $\\gamma$, likely due to the small sampling statistics. In OMC-3, the power index shows almost a linear decrease from $\\gamma=-0.8$ down to $\\gamma=-0.4$. The upper value suggests a first regime reflecting the statistical level of depolarization obtained when the cores and the high density flux regions are included and well represented. The lower value suggests another regime corresponding to the statistical level of depolarization obtained when the high density flux regions are avoided. Along OMC-2, only a small range of variation is observed with $\\gamma \\approx -0.8$, perhaps reflecting the fact that polarization was not detected in many of the brightest regions. The two regimes are illustrated by the two horizontal dashed lines shown in Figure \\ref{slopesvsdc}. To test the possibility of two regimes, we applied the same process in reverse, preferentially masking the lowest flux pixels until reaching 31$\\%$ of the peak in OMC-3 (31 of 116 pixels were used in the process) and 30$\\%$ of the peak in OMC-2 (25 of 135 pixels). The subsequent variations of the power index of the $p-I$ relation with the $CDC$ are plotted on Figure \\ref{slopesvsdcup}. Here again the effect of the cores seems to be under represented on larger scale in OMC-2 and the power index is constant and about -0.6 in the density contrast range $CDC$ $\\approx 84 \\%$ to $93 \\%$. Once the lowest density structures are avoided, however, a second regime appears and the power index is constant and about -0.4 in the density contrast range $CDC$ $\\approx 73 \\%$ to $84 \\%$. In OMC-3 the power index is constant and about -0.4 in the density contrast range $CDC$ $\\approx 87 \\%$ to $95 \\%$. Below a $CDC$ $\\approx 87 \\%$ the power index increases from about -0.4 to -0.2 showing the effects a reduction of the coverage mapping centered around high density regions could produce. Finally, when comparisons are done between the five subregions, the effect of the presence of the cores on the measured power index is clearer. The two regions FIR3 to FIR5 and MMS1 to MMS7 are consistent with the first regime where the cores are included since the maximum intensity pixel used to normalize the intensity of the $p-I$ relation are the same as those used in OMC-2 and OMC-3, respectively. Interestingly, the relatively low density region FIR6 and South of FIR6 appears consistent with the second regime where the effects of the cores are avoided. Departures from the two regimes are observed in the MMS8 to MMS10 and particularly FIR1 \\& FIR2 regions. The maximum intensity pixels used to normalize the intensity of the $p-I$ relation are small in these regions compared to the one used in region FIR6 and South to FIR6. This makes the sizes of the samples of regions MMS8 to MMS10 and FIR1 \\& FIR2 too small to be representative of the density structure of one of the two apparent regimes observed on larger scale. In conclusion, if representative of the column density structure of the molecular cloud, the presence of cores can lead to two distinct regimes of depolarization: one with a shallower power index reflecting the cores and a steeper value representative of lower column density material. We note, as pointed out in observational works \\citep[see][]{and07,whit08}, that dust alignment by RATs could be a promising ingredient for understanding depolarization. In this framework, a possible explanation for the shallower power index in cores could be the growth of larger, non-spherical grains in dense, cold condensations. As a consequence the upper cutoff of the power-law distribution of grain sizes will be higher in maps including cores than without cores and the power index of the $p-I$ relation be lower \\citep[see Figure 6 in the work of][]{cho05}. An alternative explanation could be an enhancement of gravity high enough to distort the magnetic field orientations in the cores. This mechanism will tend to decrease the net polarization and the values of the power index will change accordingly \\citep[see][and details given in section \\ref{steadystate}]{gon05}. We note that our analysis could be subject to a bias if the resolution of the instrument is too small to properly sample the column density structure of the clouds. Because the \\citet{wil03} interferometric mapping in OMC-2/3 precludes any multi-core scenario, the analysis of our data should not be subject to a bias due to intercept of cores along the same LOS, a possibility discussed by \\citet{pel07}." }, "1003/1003.2463_arXiv.txt": { "abstract": "Cygnus~OB2 is the nearest example of a massive star forming region, containing over 50 O-type stars and hundreds of B-type stars. We have analysed the properties of young stars in two fields in Cyg~OB2 using the recently published deep catalogue of {\\it Chandra} X-ray point sources with complementary optical and near-IR photometry. Our sample is complete to $\\sim$1~M$_{\\odot}$ (excluding A and B-type stars that do not emit X-rays), making this the deepest study of the stellar properties and star formation history in Cyg~OB2 to date. From \\citet{sies00} isochrone fits to the near-IR color-magnitude diagram, we derive ages of $3.5^{+0.75}_{-1.0}$ and $5.25^{+1.5}_{-1.0}$~Myrs for sources in the two fields, both with considerable spreads around the pre-MS isochrones. The presence of a stellar population somewhat older than the present-day O-type stars, also fits in with the low fraction of sources with inner circumstellar disks (as traced by the $K$-band excess) that we find to be very low, but appropriate for a population of age $\\sim$5~Myrs. We also find that the region lacks a population of highly embedded sources that is often observed in young star forming regions, suggesting star formation in the vicinity has declined. We measure the stellar mass functions in this limit and find a power-law slope of $\\Gamma = -1.09 \\pm 0.13$, in good agreement with the global mean value estimated by \\citet{krou02}. A steepening of the slope at higher masses is observed and suggested as due to the presence of the previous generation of stars that have lost their most massive members. Finally, combining our mass function and an estimate of the radial density profile of the association suggests a total mass of Cyg~OB2 of $\\sim 3 \\times 10^4$~M$_{\\odot}$, similar to that of many of our Galaxy's most massive star forming regions. ", "introduction": "Stars do not form individually but in groups, with sizes ranging from small star forming regions (SFRs) to the massive superstar clusters seen in merging galaxies \\citep{lada03}. Our current understanding of star formation is mainly derived from observations of nearby small star forming regions such as Taurus or Ophiuchus. However the conditions of star formation vary considerably between these regions and the massive SFRs that contain hundreds to thousands of OB stars and millions of low-mass stars. High stellar densities, strong stellar winds from OB stars and a large UV flux are likely to influence the products of the star formation process such as the initial mass function (IMF), the binary fraction and the evolution of protostellar disks. Our understanding of the influence of environment on these products is currently poor, and needs to improve if we are to develop a global understanding of star formation that can relate the small-scale properties with the evolution of galaxies and the survival of planetary systems. The majority of massive SFRs are found in external galaxies, or close to the center of our own Galaxy. Their rarity and great distances therefore make them difficult objects to study. The one exception to this is Cygnus~OB2, which, at a distance of only 1.45~kpc \\citep{hans03} is the closest known massive SFR. Early surveys of the region revealed the presence of $\\sim$300 OB stars \\citep{redd67}, a number that has grown as deeper surveys have penetrated the high extinction in its direction \\citep[e.g.][]{mass91,come02,hans03}. There are currently 65 known O-type stars in the region making it the largest known concentration of spectroscopically confirmed O-type stars. From a near-IR photometric study \\citet{knod00} estimated the region contained over 100 O-type stars, $\\sim$2600 OB stars, and had a total mass of 4--10~$\\times 10^4$~M$_{\\odot}$. However, the complications of background subtraction in the Galactic plane, combined with strongly variable extinction in the vicinity of Cyg~OB2, give cause to doubt such an estimate. Furthermore the question of the extent of Cyg~OB2 is made more complex by the growing evidence for stars older than the 1-3~Myr age appropriate to the OB stars: \\citet{drew08} found a concentration of 5-7~Myr old A stars in 1~sq.~deg just to the south of Cyg~OB2, while \\citet{come08} identified an unclustered spread of evolved stars up to $\\sim$10~Myrs old over a wider area, 2~degrees in radius. In this study we obtain further evidence that would favour the present-day OB stars as the products of the latest but most certainly not the only phase of star formation. Recent deep photometric surveys of the Galactic plane in the optical and near-IR have penetrated much of the extinction toward the region and are advancing our understanding of this massive association \\citep[e.g.][]{drew08,vink08}. The combination of data from these surveys with that at other wavelengths is vital for studying the star formation process, particularly in the dense and highly obscured regions of the Galactic plane, and one that this work makes particular use of. One of the objectives of this work is to measure the IMF down to low masses. The IMF is one of the most important measurable quantities in star formation studies. Combined with the time-variable star formation rate, the IMF determines the evolution of stellar systems ranging from small SFRs to entire galaxies. Measurements of the IMF have shown it to be highly uniform at high and intermediate masses \\citep{luhm00,mass98} with a power law slope above 1~M$_{\\odot}$ with an index of $\\Gamma = -1.35$ \\citep{salp55} for intermediate mass stars, possibly steepening to $\\Gamma = -1.7$ for the most massive stars \\citep{scal86}. At lower masses the IMF is less well constrained, but appears to flatten below 0.5~M$_{\\odot}$ and exhibits a `turnover' around 0.1~--~0.3~M$_{\\odot}$, with less stars of the lowest masses \\citep{krou01a}. Deviations from this uniformity are becoming increasingly rare as deeper and more complete observations overturn previous measurements and the majority of remaining observations have been shown to be within the statistical uncertainties of the seemingly universal mean \\citep[e.g.][]{krou02}. While some theoretical models predict a shallower IMF in regions of high density or pressure, only a small number of observations in the vicinity of the Galactic Center \\citep[e.g.][]{stol06} or in starburst regions \\citep[e.g.][]{smit01a} support this and more measurements of the IMF in massive SFRs are necessary to resolve this. The most recent measurement of the IMF in Cyg~OB2, due to \\citet{kimi07}, suggests that at high masses it is in fact steeper than the canonical value (though they estimate their spectroscopic sample becomes incomplete at masses below $\\sim$15~M$_{\\odot}$). This further supports the need for a thorough and complete measurement of the IMF in Cyg~OB2 down to lower masses. In this paper we present an analysis and discussion of the stellar properties of sources in Cygnus~OB2 taken from the X-ray derived catalogue presented by \\citet[][hereafter Paper~1]{ wrig09a}. In Section~2 we discuss the source catalogue and the various cuts applied to it to remove foreground and background contaminants. In Section~3 we present the near-IR properties of these sources, in Section~4 we study their X-ray properties and in Section~5 we derive mass functions for the sources. Finally, in Section~6 we discuss the implications of our findings on our understanding of Cygnus~OB2 and compare it to other massive SFRs. ", "conclusions": "We have analysed the integrated stellar properties of young stars in the massive SFR Cygnus~OB2 using the recently published catalogue of {\\it Chandra} X-ray point sources detected in the region. We find that our sample is complete down to $\\sim$1~M$_{\\odot}$, excluding A and B-type stars that are not expected to emit X-rays, making this one of the deepest surveys of the region. Ages for sources in the two fields studied were estimated from pre-MS isochrone fits to the near-IR CMD and found to be $3.5^{+0.75}_{-1.0}$ and $5.25^{+1.5}_{-1.0}$~Myrs, both with significant spreads around the isochrones that cannot solely be explained by photometric errors. Given that the most massive stars in the region have an age of $2 \\pm 1$~Myrs \\citep{mass91,hans03}, we suggest that this supports recent evidence for an older generation of stars in the region \\citep[e.g.][]{drew08,come08} that either represent a significant age spread or multiple epochs of star formation. Additionally, the lack of sizable population of highly embedded sources suggests that star formation in the region studied has declined significantly. Assuming that a significant fraction of sources in Cyg~OB2 are older than the previously accepted age of 2~Myrs, many of the unique properties of the region may be explained. As noted by \\citet{drew08}, the distinct lack of a bright H{\\sc ii} region surrounding Cyg~OB2 is perfectly reasonable if the clearing timescale has been greater than 5~Myrs. Furthermore, we propose that the low fraction of sources with $K$-band excesses found by \\citet[][and confirmed from our larger sample]{alba07} is less abnormal if a significant fraction of sources are older than previously believed. We measure the stellar mass functions for stars within our completeness limits and find a slope of $\\Gamma = -1.09 \\pm 0.13$, in good agreement with the potentially universal value estimated by \\citet{krou02}. This is the deepest and most accurate measurement of the mass function in Cyg~OB2, building on previously measurements that have found either steeper \\citep{knod00} or shallower \\citep{mass91} slopes. A steepening of the mass function is observed at high masses, as found by \\citet{kimi07}, which we attribute to the demise of the highest mass members of the older generations in the region. An estimation of the total size of the region suggests that the stellar density drops off faster than estimated by \\citet{knod00}. This, combined with our mass function estimate, reveals a total mass of $\\sim 3 \\times 10^4$~M$_{\\odot}$. This makes Cyg~OB2 similar in mass to many of the recently discovered massive clusters in our Galaxy such as the Arches and Quintuplet clusters. However, we note that comparing other properties of these massive star forming regions, such as their central densities, reveals differences that may hint at different formation mechanisms. Past studies of Cyg~OB2 have clearly been complicated by an uncertain star formation history and a large spatial extent in a crowded region of the Galactic Plane. These are likely to be inherent properties of massive star forming regions in galactic spiral arms where star formation continues for longer periods and across larger areas, propagated through entire GMCs by continued rejuvenation and feedback triggering. These findings support the need to fully understand the full star formation history of a region before if its observed properties are to be correctly interpreted. Cygnus~OB2 is clearly a very important region to study, not only because of its proximity and size, but also because of the insights it will give us into galactic-scale star formation and the evolution of GMCs." }, "1003/1003.1130_arXiv.txt": { "abstract": "{The AKARI, a Japanese infrared space mission, has performed an All-Sky Survey in six infrared-bands from 9 to 180 $\\mu$m with higher spatial resolutions and better sensitivities than the Infrared Astronomical Satellite (IRAS).} {We investigate the mid-infrared (9 and 18 $\\mu$m) point source catalog (PSC) obtained with the Infrared Camera (IRC) on board the AKARI, in order to understand the infrared nature of the known objects, as well as to identify previously unknown objects.} {Color-color diagrams and a color-magnitude diagram have been plotted, using the AKARI-IRC PSC and other available all-sky survey catalogs. We combine the Hipparcos astrometric catalog, and the 2MASS all-sky survey catalog with the AKARI-IRC PSC. We further searched literatures and SIMBAD astronomical database for object types, spectral types and luminosity classes. We identify the locations of representative stars/objects on color-magnitude and color-color diagram scheme. The properties of unclassified sources can be inferred on the basis of their locations on these diagrams.} {We found that the ($B-V$) v.s. ($V-S9W$) color-color diagram is useful to identify the stars with infrared excess emerged from circumstellar envelopes/disks. Be stars with infrared excess are well separated from other types of stars in this diagram. Whereas ($J-L18W$) v.s. ($S9W-L18W$) diagram is a powerful tool to classify several object-types. Carbon-rich asymptotic giant branch (AGB) stars and OH/IR stars form distinct sequences in this color-color diagram. Young stellar objects (YSOs), pre-main sequence (PMS) stars, post-AGB stars and planetary nebulae (PNe) have largest mid-infrared color-excess, and can be identified in infrared catalog. Finally, we plot $L18W$ v.s. ($S9W-L18W$) color-magnitude diagram, using the AKARI data together with Hipparcos parallaxes. This diagram can be used to identify low-mass YSOs, as well as AGB stars. We found that this diagram is comparable to the $[24]$ vs $([8.0]-[24])$ diagram of Large Magellanic Cloud sources using the \\textit{Spitzer} Space Telescope data. Our understanding of Galactic objects will be used to interpret color-magnitude diagram of stellar populations in nearby galaxies which \\textit{Spitzer} Space Telescope has observed. } {Our study of the AKARI color-color and color-magnitude will be used to explore properties of unknown objects in future. In addition, our analysis highlights a future key project to understand stellar evolution with circumstellar envelope, once the forthcoming astronometrical data with GAIA are available.} ", "introduction": "More than a quarter-century passed since the pioneering infrared whole sky survey of the InfraRed Astronomical Satellite (IRAS), which covered more than 96\\% of the whole sky in four photometric bands at 12, 25, 60 and 100 $\\mu$m (IRAS Explanatory Supplement \\cite{iras}). The IRAS point source catalog (PSC) has shown that mid- and far-infrared census is essential for studying dust embedded objects, such as star-forming regions, debris disks around main sequence stars, evolved stars, and distant galaxies. However, the spatial resolution was not so good to study sources in the crowded retions. After the IRAS, the COsmic Background Explore (DIRBE/COBE; Hauser et al. \\cite{hauser1998}) has mapped the whole sky in 1.25--240\\,$\\mu$m with 10 photometric-bands. It intended to accurately obtain the intensity of diffuse background radiation and did not have high sensitivity for point sources. In the meanwhile, ultra-violet, optical and near-infrared large area surveys have carried out using ground based telescopes (e.g. GALEX (Martin et al. \\cite{martin2005}); SDSS (York et al. \\cite{york2000}); 2MASS (Skrutskie et al. \\cite{skrutskie2006}). Their counterparts are missing or hard to be uniquely identified in mid-infrared and far-infrared catalogs, which prevents us from studying objects surrounded by dust. The demand of new mid- and far-infrared whole sky survey with better sensitivity and higher spatial resolution has been increased. To fulfill the expectations, the AKARI, a Japanese infrared satellite was launched at 21:28 UTC on February 21st, 2006 from the Uchinoura Space Center (Murakami et al. \\cite{murakami2007}). Sharing the time with the pointed observations, AKARI has mapped the whole sky in mid- and far-infrared using two instruments on board; the InfraRed Camera (IRC; Onaka et al. \\cite{onaka2007}) and the Far-Infrared Surveyor (FIS; Kawada et al. \\cite{kawada2007}). The FIS swept about 94\\% of the whole sky more than twice at 65, 90, 140, and 160 $\\mu$m wavebands. Also, the IRC swept more than 90\\% of the whole sky more than twice using two filter bands centered at 9 ($S9W$, 7 -- 12 $\\mu$m) and 18 ($L18W$, 14 -- 25 $\\mu$m) $\\mu$m (Ishihara et al. \\cite{ishihara2010}). These abbreviated filter band names are used throughout this article. The cut-in and cut-off wavelengths indicated in the parentheses correspond to those where the throughput becomes a half of the peak. See Figure~\\ref{filter} for the normalized spectral response function of the IRC bands. In this paper, we use the IRC mid-infrared All-Sky Survey data to study Galactic stellar objects. Compared to the IRAS survey, the sensitivities at 9 and 18 $\\mu$m bands are more than 15 and 5 times better than those of the IRAS's 12 and 24 $\\mu$m bands, and the spatial resolution is more than 100 times finer for the IRC survey. van der Veen \\& Habing (\\cite{veen1988}) utilized IRAS mid- and far-infrared combined colors as a tool to diagnose the nature of IRAS sources. They defined color criteria to classify sources into several groups, and since then those criteria have been used by many authors. Taking this as a role model, we cross-identify the AKARI IRC All-Sky Survey point source catalog with the Hipparcos astrometric catalog (van Leeuwen \\cite{vanLeeuwen2007}), and the 2MASS all-sky survey catalog (Skrutskie et al. \\cite{skrutskie2006}), to classify sources on color-color and color-magnitude diagrams. The main objective of this paper is to study overall infrared characteristics of galactic stellar sources. The new All-Sky Survey catalog should be useful for wide variety of astronomical studies. Practical applications of the catalog are, search for hot debris disk (Fujiwara et al. \\cite{fujiwara2010}), extragalactic objects (Oyabu et al. in preparation) and study on the young stellar objects (Takita et al. \\cite{takita2010}). Refer to the papers for discussions on each topic. In the next section we show general characteristics of the AKARI IRC All-Sky Survey point source catalog. Refer to Ishihara et al. (\\cite{ishihara2010}) for the complete description of the All-Sky Survey, its data reduction processes, the point source catalog compilation processes, and the catalog characteristics. \\begin{figure} \\centering \\includegraphics[angle=0,scale=0.7]{filters.jpg} \\caption{The normalized spectral response function of the AKARI IRC bands and the IRAS bands. As references, the ISO SWS spectra of three rep resentative Galactic AGB stars (T Cep as O-rich AGB with AlOx feature, VX Sgr as O-rich AGB with silicate feature, and IRC+50096 as C-rich AGB with SiC feature) with circumstellar dust features are shown.} \\label{filter} \\end{figure} \\begin{figure} \\centering \\includegraphics[angle=-90,scale=0.36]{galdist.jpg} \\caption{A projection of the IRAS sources without AKARI detection (upper panel) and the 2MASS-drop AKARI sources (lower panel) onto the galactic coordinate map. The blue lines show the equatorial coordinate.} \\label{galdist} \\end{figure} ", "conclusions": "AKARI's mid-infrared All-Sky Survey increased the number of known mid-infrared sources drastically, mainly because of better spatial resolution than previous mid-infrared surveys. We combined the first release version ($\\beta$-1) of the AKARI IRC All-Sky survey point source catalog with the existing all-sky survey catalogs, namely the Hipparcos and the 2MASS. Two-color diagrams are made with an aim to classify sources. We found that oxygen-rich giants and carbon stars are well separated by adding AKARI's new data. Also, we showed that Be stars and Wolf-Rayet stars with strong infrared excesses can be effectively selected by using optical and AKARI's combined colors. Combined with Hipparcos parallax measurements, we plot an infrared color-magnitude diagram. We uncovered the properties of redder, fainter red giants in the LMC, by comparing their galactic counterparts. This work will be greatly expanded in the forthcoming GAIA era, when we have good parallax measurements for tens of millions of stars. AKARI's new All-Sky Survey data reveal not only the mid-infrared characteristics of known objects, but also the existence of many yet-unidentified infrared sources. The color-color and color-magnitude diagrams we presented can be used to extrapolate the properties of the unidentified objects, invoking a follow-up campaign." }, "1003/1003.4653_arXiv.txt": { "abstract": "The structure of electric current and magnetic helicity in the solar corona is closely linked to solar activity over the 11-year cycle, yet is poorly understood. As an alternative to traditional current-free ``potential field'' extrapolations, we investigate a model for the global coronal magnetic field which is non-potential and time-dependent, following the build-up and transport of magnetic helicity due to flux emergence and large-scale photospheric motions. This helicity concentrates into twisted magnetic flux ropes, which may lose equilibrium and be ejected. Here, we consider how the magnetic structure predicted by this model---in particular the flux ropes---varies over the solar activity cycle, based on photospheric input data from six periods of cycle 23. The number of flux ropes doubles from minimum to maximum, following the total length of photospheric polarity inversion lines. However, the number of flux rope ejections increases by a factor of eight, following the emergence rate of active regions. This is broadly consistent with the observed cycle modulation of coronal mass ejections, although the actual rate of ejections in the simulation is about a fifth of the rate of observed events. The model predicts that, even at minimum, differential rotation will produce sheared, non-potential, magnetic structure at all latitudes. ", "introduction": "\\label{sec:intro} The number of sunspots has long been known to vary with a period of about 11 years. Since the early 20th century it has been recognised that sunspots carry magnetic flux originating from the Sun's interior. Once this magnetic flux emerges into the corona, it dominates both the structure and evolution of the plasma there. This is evidenced by 11-year variations in many phenomena, including solar flares \\citep{aschwanden2005}, prominences \\citep{dazambuja1955,hansen1975}, X-ray flux \\citep{pevtsov2001}, coronal mass ejections \\citep[CMEs;][]{webb1994,cremades2007}, and the solar wind \\citep{richardson2001}. Despite the connection to observed phenomena, the coronal magnetic field cannot usually be measured directly due to the tenuous nature of the plasma, so only isolated measurements exist \\citep[{\\it e.g.},][]{lin2000,casini2003,tomczyk2007}. Usually, it must be extrapolated from the routinely observed magnetic field in the solar photosphere, perhaps using images of high-temperature loops to constrain the field topology. This topology is often found to contain non-zero electric currents \\citep[{\\it e.g.} in X-ray sigmoids;][]{canfield1999}, and extrapolation of such magnetic fields is not yet robust, even when limited to a single active region \\citep{derosa2009}. On the global scale, which concerns us in this paper, it is usual to assume a current-free magnetic field (a potential field). This is uniquely determined in the corona given an observed radial component on the photosphere and vanishing horizontal components at an upper ``source surface'' \\citep[usually placed at $r=2.5R_\\odot$;][]{altschuler1969}. Even the more realistic models that solve for MHD equilibria in the corona usually assume a potential magnetic field which is then perturbed to be consistent with a given thermodynamic structure and solar wind \\citep{riley2006,cohen2007}. Whilst the potential field model has been successful in describing certain large-scale aspects of coronal magnetic structure---such as the locations of coronal holes \\citep{wang1996}---it does not include the development of currents and free magnetic energy that are needed to model solar eruptions. In fact, the highly sheared magnetic fields observed in long-lived filament channels all over the Sun \\citep{martin1994} show that, even outside active regions, the assumption of a potential field can be inadequate. As a first step toward understanding the structure of currents and magnetic helicity in the global corona, we have investigated an alternative model for the coronal magnetic field, based also on observational input of the photospheric magnetic field. In this model, the large-scale mean magnetic field in the corona evolves in time through a quasi-static relaxation, in response to flux emergence and to surface motions \\citep{vanballegooijen2000,yeates2008,yeates2009b}. The consequence is that, unlike in a sequence of potential field extrapolations, current and magnetic helicity are generated in and transported through the corona. Flux cancellation above photospheric polarity inversion lines (PILs) then leads to the concentration of helicity in either sheared arcades or twisted magnetic flux ropes \\citep[through the mechanism described by][]{vanballegooijen1989}. Comparison of the simulated magnetic field direction at these locations with observations of filament chirality has demonstrated that, despite its simplicity, the model is able to reproduce the general structure of the magnetic field at most of these locations \\citep{yeates2008}. In this paper, we use the quasi-static model to simulate the coronal magnetic field during six distinct periods over cycle 23. Our aim is to consider how the magnetic field structure predicted by this model varies over the 11-year solar activity cycle. We focus on a particular aspect of the field structure: the formation and ejection of magnetic flux ropes. These are not only observed in the real corona, but their ejection is suggested to give rise to CMEs \\citep[{\\it e.g.},][]{gibson2006}. By their very nature, flux ropes cannot be modelled by potential field extrapolations, so their locations and ejections have not been considered in previous models of the global corona. The formation of flux ropes in the present quasi-static model has been described for a simple two-bipole configuration by \\citet{mackay2006}, and for the global corona over a particular 5-month period by \\citet{yeates2009b}. Here we extend the latter study to simulate different phases of the solar cycle. ", "conclusions": "\\label{sec:conclusions} As a step toward removing the potential-field assumption in modelling the global coronal magnetic field, we have investigated a quasi-static model allowing for electric currents. This model is driven by photospheric magnetic observations and follows the build-up of magnetic helicity as magnetic flux emerges and is transported by surface motions. In this paper, we simulate six distinct periods during solar cycle 23, so as to study how the magnetic field structure predicted by this model varies over the solar cycle. A key feature of this model is the transport of magnetic helicity and its concentration into twisted magnetic flux ropes above photospheric PILs. The solar cycle variation of these flux ropes in the model may be characterised as follows: \\begin{enumerate} \\item The number of flux ropes present at any one time doubles between minimum and maximum of the solar cycle. This follows the total length of photospheric PILs. \\item The rate of flux rope ejections---or losses of equilibrium---increases by a factor of eight between minimum and maximum. This is lower than the relative increase in free magnetic energy, but similar to the relative increases in total magnetic energy and total parallel current, both of which follow the number of emerging active regions. \\end{enumerate} The difference between the relative increases in ejection rate and in the number of flux ropes implies that, at maximum, the flux ropes have shorter lifetimes before losing equilibrium. This is a consequence of the higher rate of active region emergence. Although there are fewer flux rope ejections at minimum, it is not the case that the corona in our model is everywhere close to potential. Even in the absence of many emerging bipoles, shearing of pre-existing coronal field by differential rotation generates current on a large scale at all latitudes. A weak background field originating in earlier decayed active regions is present even in the recent extended minimum (modelled by our period F in 2008). Our model assumes that differential rotation is equally effective at shearing this weaker field, so that, for example, we find systematic sheared arcades on the high-latitude polar crowns. Interestingly, although sheared arcades are formed, the low reconnecting flux over these PILs in 2008 prevents detached helical flux ropes from forming at many locations, unlike at solar maximum. The true nature of the magnetic field above the polar crown PILs remains uncertain. Our simulations predict that differential rotation will develop positive and negative helicity over such east-west PILs in the northern and southern hemispheres respectively. However, the (few) existing observations of the magnetic field in polar crown prominences imply the opposite sign of helicity in each hemisphere \\citep{leroy1983}. In our quasi-static model, the correct sign of helicity is only obtained at such PIL locations if additional emergence of axial flux is included \\citep{vanballegooijen2000}. This was not included in the present paper because it lacks observational justification. A concerted effort to obtain further observations of the magnetic field structure at high latitudes is needed in order to resolve this issue and hence test whether additional sources of helicity are required in addition to those in our existing model. An important question that may be addressed by a time-dependent non-potential model of the coronal magnetic field is the initiation of CMEs. Our model allows us to test whether the loss of equilibrium of flux ropes formed by quasi-static shearing and flux cancellation is sufficient to account for the observed CME rate. For the present form of the model the answer is negative: the ratio of simulated to observed ejection rate (from CDAW) remains about $0.2$ over the whole solar cycle. Even in our earlier (less realistic) simulations where all emerging bipoles were given the same sign of helicity in each hemisphere, the ratio increases only to about $0.25$. This suggests strongly that the formation of flux ropes by quasi-static shearing of the large-scale magnetic field is incapable of initiating all CMEs. We should point out, however, that the CDAW rate for period F (in 2008) may be too high (as discussed in Section \\ref{sec:ejections}), in which case the model may account for a greater proportion of CMEs in the recent very quiet period. Nevertheless, it is clear that the present model lacks the detail to adequately describe initiation of all CMEs. More detailed comparison with observed low-coronal CME source regions \\citep[described by][]{yeates2009} reveals that many of the additional observed CMEs are located in active regions, with individual active regions frequently producing multiple CMEs in the same day. Our global quasi-static simulations cannot produce such dynamic events with the present form of input data (updated once per Carrington rotation). The quasi-static approach could be developed in future to incorporate driving from photospheric magnetograms on much shorter spatial and time scales, in order to test whether the energetic restructuring of the corona following flux emergence may be adequately described by a quasi-static model, or whether fully dynamic MHD modelling is required. From the global perspective, this more detailed modelling of active regions should allow their helicity content, and hence the global transport of helicity, to be better constrained by observations. This has important implications both for models of the coronal magnetic structure and, ultimately, for the origin of the Sun's magnetic field in the solar dynamo \\citep{seehafer2003}. \\begin{acks} We thank A.A. van Ballegooijen, D.H. Mackay, A.N. Wilmot-Smith, D.I. Pontin and an anonymous referee for useful suggestions, and D.H. Mackay also for the use of parallel computing facilities obtained through a Royal Society research grant. ARY acknowledges financial support from NASA contract NNM07AB07C at SAO, and from the UK STFC at Dundee. The visit of JAC to SAO was supported by NASA grant NNX08AW53 and by NSF grant ATM-0851866 for ``REU site: Solar Physics at the Harvard-Smithsonian Center for Astrophysics.'' Magnetogram data from NSO/Kitt Peak were produced cooperatively by NSF/NOAO, NASA/GSFC and NOAA/SEL. SOLIS data used here are produced cooperatively by NSF/NSO and NASA/LWS. \\end{acks}" }, "1003/1003.4379_arXiv.txt": { "abstract": "The very first stars in the Universe can be very massive, up to $10^6M_\\odot$. They would leave behind massive black holes that could act as seeds for growing super massive black holes of active galactic nuclei. Given the anticipated fast rotation such stars would end their live as super massive collapsars and drive powerful magnetically-dominated jets. In this paper we investigate the possibility of observing the bursts of high-energy emission similar to the Long Gamma Ray Bursts associated with normal collapsars. We show that during the collapse of supercollapsars, the Blandford-Znajek mechanism can produce jets as powerful as few$\\times10^{51}$erg/s and release up to $10^{56}$erg of the black hole rotational energy. Due to the higher intrinsic time scale and higher redshift the initial bright phase of the burst can last for about $10^5$ seconds whereas the central engine would remain active for about 10 days. Due to the high redshift the burst spectrum is expected to be soft, with the spectral energy distribution peaking at around 60keV. The peak total flux density is relatively low, few$\\times 10^{-7}\\mbox{erg}\\, \\mbox{cm}^{-2} \\mbox{s}^{-1}$, but not prohibitive. The such events should be rear 0.03 year$^{-1}$, the observations needs long term program and could be done in future. ", "introduction": "\\label{intro} The very first stars in universe were borned in the lack of heavy elements. It could leads to born of the super massive stars (SMS) \\citep{bcl02,bl03,ss06,dv09}. Such stars can be as massive as $1000 M_{\\odot} < M < 50000 M_{\\odot}$. The stars with masses about $ 1000 M_{\\odot}$, which will be referred to as Very Massive Stars (VMSs) are expected to collapse into black holes with very little mass loss \\citep{F01}. They would leave behind massive black holes (MBHs), which could play the role of seeds for the super massive black holes (SMBHs) of Active Galactic Nuclei (AGNs). The collapse of VMS were discussed recently in the paper \\citep{KB09c}. Even more massive $3\\times 10^4 M_{\\odot} < M < 10^6 M_{\\odot}$ stars, which will be referred to as Super Massive Stars (SMSs) could be formed in more massive dark matter haloes with total mass $M\\simeq 10^8M_{\\odot}$ collapsed at $z\\simeq10$ \\citep{bl03,BVR06}. This stars do not reach the instability limit mass which is near $10^6 M_{\\odot}$ and can be formed \\citep{hf63,zn65,bzn67,w69}. The collapse of SMS can provide an alternative way of producing SMBHs. There are two crucial differences between a normal collapsar and a supercollapsar. One is that instead of a proto-neutron star of solar mass the supercollapsars develop proto-black holes of tens of solar masses, within which the neutrinos from electron capture are trapped \\citep{F01,Su07}. The other is that the accretion disks of supercollapsars are far too large and cool for the neutrino annihilation mechanism. This has already been seen in the numerical simulations of supercollapsar with mass $M=300M_\\odot$ \\citep{F01}. Utilising the study of hyper-accreting disks by \\citet{bel08} we find that at best the rate of heating due to this mechanism is \\begin{equation} \\dot{E} \\simeq 6\\times10^{43} \\dot{M}_0^{9/4} M_{h,6}^{-3/2} \\mbox{erg s}^{-1}, \\end{equation} where $\\dot{M}$ is the accretion rate and $M_h$ is the black hole mass. (Here and in other numerical estimates below we use the following notation: $\\dot{M}_k$ is the mass accretion rate measured in the units of $10^k M_{\\odot} \\mbox{s}^{-1}$ and $M_k$ is the mass measured in the units of $10^kM_\\odot$.) Such low values have lead \\citet{F01} to conclude that the magnetic mechanism is the only candidate for producing GRB jets from supercollapsars. In the following we analyse one particular version of the mechanism where the jets are powered by the rotational energy of the black hole via the Blandford-Znajek process \\citep{BZ77,BK08}. This paper is an extension our previous work \\citep{KB09c} here we investigate properties of collapse of VMS. In this paper we estimate temporal structure of the SMS collapse and predict the observational evidence of this process. ", "conclusions": "\\label{conclusions} In spite of the significant progress in the astrophysics of Gamma Ray Bursts, both observational and theoretical, it may still take quite a while before we fully understand both the physics of the bursts and the nature of their progenitors. At the moment there are several competing theories and too many unknowns. Similarly, we know very little about the star formation in the early Universe. For this reason, the analysis presented above is rather speculative and the numbers it yields are not very reliable. Further efforts are required to develop a proper theory of supercollapsars and to make firm conclusions on their observational impact. The collapse of SMS could be one of the most powerful event in the universe. This event can destroy the seed cluster and form single SMBH. The expected very long duration of bursts and their relatively low brightness imply that a dedicated search program using the image trigger will be required. Such search would be useful even in the case of non-detection as this would put important constraint on the models of early star formation, GRB progenitors, and SMBHs." }, "1003/1003.4186_arXiv.txt": { "abstract": "We study gravity mediated supersymmetry breaking in ${\\cal F}$-$SU(5)$ and its low-energy supersymmetric phenomenology. The gaugino masses are not unified at the traditional grand unification scale, but we nonetheless have the same one-loop gaugino mass relation at the electroweak scale as minimal supergravity (mSUGRA). We introduce parameters testable at the colliders to measure the small second loop deviation from the mSUGRA gaugino mass relation at the electroweak scale. In the minimal $SU(5)$ model with gravity mediated supersymmetry breaking, we show that the deviations from the mSUGRA gaugino mass relations are within 5\\%. However, in ${\\cal F}$-$SU(5)$, we predict the deviations from the mSUGRA gaugino mass relations to be larger due to the presence of vector-like particles, which can be tested at the colliders. We determine the viable parameter space that satisfies all the latest experimental constraints and find it is consistent with the CDMS II experiment. Further, we compute the cross-sections of neutralino annihilations into gamma-rays and compare to the first published Fermi-LAT measurement. Finally, the corresponding range of proton lifetime predictions is calculated and found to be within reach of the future Hyper-Kamiokande and DUSEL experiments. ", "introduction": "As we initiate the era of the Large Hadron Collider (LHC), we await with anticipation the expected discovery of supersymmetry and the Higgs states required to break electroweak symmetry and stabilize the electroweak scale. On the other hand, there is thus far no concrete model that can explain all observed physics in a comprehensive mathematical framework. Unique predictions that can be tested at the LHC, future International Linear Collider (ILC), and other forthcoming experiments are necessary if string theory is to be substantiated as the correct fundamental description of nature. Following a top-down approach, it may be feasible to derive all known observable physics from a fundamental theory such as string theory. In contrast, the bottom-up approach offers the possibility to infer the framework of the fundamental theory at high-energy from a low-energy signal at the experiments. In the spirit of this bottom-up approach, our goal here is to study Grand Unified Theories (GUTs) from F-Theory, which have seen exciting progress the past two years, and present F-Theory GUT low-energy physics observable at current and future experiments. F-Theory can be considered as the strongly coupled formulation of ten-dimensional Type IIB string theory with a varying axion ($a$)-dilaton ($\\phi$) field $S=a+ie^{-\\phi}$~\\cite{Vafa:1996xn}. GUTs were first locally constructed in F-theory two years ago~\\cite{Donagi:2008ca, Beasley:2008dc, Beasley:2008kw, Donagi:2008kj}, and subsequently, model building and phenomenological consequences have been studied extensively~\\cite{Heckman:2008es, Heckman:2008qt, Font:2008id, Heckman:2008qa, Jiang:2009zza, Blumenhagen:2008aw, Heckman:2008jy, Hayashi:2009ge, Chen:2009me, Heckman:2009bi, Donagi:2009ra, Jiang:2009za, Li:2009cy, Marsano:2009gv, Cecotti:2009zf, Li:2009fq, Li:2010dp, Marsano:2009wr, Leontaris:2009wi, Heckman:2010xz, Li:2010mr}. In F-theory model building, the gauge fields are on the observable seven-branes that wrap a del Pezzo $n$ ($dP_n$) surface for the extra four space dimensions, the SM fermions and Higgs fields are localized on the complex codimension-one curves (two-dimensional real subspaces) in $dP_n$, and the Standard Model (SM) fermion Yukawa couplings arise from the triple intersections of the SM fermion and Higgs curves. A brand new feature is that the $SU(5)$ gauge symmetry can be broken down to the SM gauge symmetry by turning on the $U(1)_Y$ flux~\\cite{Beasley:2008dc, Beasley:2008kw, Li:2009cy}, and the $SO(10)$ gauge symmetry can be broken down to the $SU(5)\\times U(1)_X$ and $SU(3)\\times SU(2)_L\\times SU(2)_R\\times U(1)_{B-L}$ gauge symmetries by turning on the $U(1)_X$ and $U(1)_{B-L}$ fluxes, respectively~\\cite{Beasley:2008dc, Beasley:2008kw, Jiang:2009zza, Jiang:2009za, Font:2008id, Li:2009cy}. Intriguingly, the $SO(10)$ models have both gauge interaction unification and SM fermion unification, and consequently we believe the $SO(10)$ models are more interesting to study than the $SU(5)$ models. Moreover, we can break the $SO(10)$ gauge symmetry down to the flipped $SU(5)\\times U(1)_X$ gauge symmetry by turning on the $U(1)_X$ flux, wherein the doublet-triplet splitting problem can be solved elegantly~\\cite{smbarr, dimitri, AEHN-0}. The flipped $SU(5)\\times U(1)_X$ models with TeV scale vector-like particles~\\cite{Jiang:2006hf}, which have been dubbed ``${\\cal F}$-$SU(5)$'', have been constructed systematically in F-theory~\\cite{Jiang:2009zza, Jiang:2009za}. Models of this nature can be also obtained in the free fermionic string constructions~\\cite{Lopez:1992kg}. In ${\\cal F}$-$SU(5)$, the $SU(3)_C\\times SU(2)_L$ gauge symmetries are unified at about $10^{16}$ GeV, and the $SU(5)\\times U(1)_X$ gauge symmetries are unified above $10^{17}$ GeV. Thus, we can solve the monopole problem. On top of this, the TeV scale vector-like particles are potentially observable at the Large Hadron Collider (LHC), the lightest CP-even Higgs boson mass can be lifted~\\cite{HLNT-P}, and the predicted proton decay~\\cite{Ellis:1995at,Ellis:2002vk,Li:2009fq,Li:2010dp} is within the reach of the future Hyper-Kamiokande~\\cite{Nakamura:2003hk} and Deep Underground Science and Engineering Laboratory (DUSEL)~\\cite{DUSEL} experiments. Recently, the Cryogenic Dark Matter Search (CDMS) collaboration observed two candidate dark matter events in the CDMS II experiment~\\cite{Ahmed:2009zw}. The recoil energies for these two events were 12.3 keV and 15.5 keV, respectively, and the data set an upper limit on the dark matter-nucleon elastic-scattering spin independent cross section around $10^{-8}-10^{-7}$ pb. The probability of observing two or more background events is $23\\%$, so the CDMS II results cannot be considered statistically significant evidence for dark matter interactions, although there is some strong possibility that they do in fact represent an authentic signal. In particular, the favored dark matter mass from the CDMS II data is about 100 GeV. Interestingly, the CDMS II experiment can be explained in the supersymmetric Standard Model with $R$ parity where the lightest supersymmetric particle (LSP) neutralino is the dark matter candidate~\\cite{Bernal:2009jc, Bottino:2009km, Feldman:2009pn, Allahverdi:2009sb, Endo:2009uj, Holmes:2009uu, Hisano:2009xv, Gogoladze:2009mc, Maxin:2009kp, Gao:2010pg}. In this paper, we first briefly review ${\\cal F}$-$SU(5)$, and then subsequently study the supersymmetry breaking soft terms in the framework of gravity mediation. Although the gaugino masses are not unified at the traditional GUT scale, we still have the same gaugino mass relation at the electroweak scale as the minimal supergravity (mSUGRA) scenario at one loop~\\cite{mSUGRA}. We also introduce two parameters testable at the colliders to measure the small two-loop deviations from the mSUGRA gaugino mass relations at the electroweak scale. Next, we incorporate TeV scale vector particles and generate regions of the parameter space that can satisfy all current experimental constraints and are consistent with the CDMS II experiment~\\cite{Ahmed:2009zw}. In addition, we compute the annihilation cross-section of two neutralinos into two gamma-rays and evaluate the results in light of the first published Fermi-LAT measurement. Then we compute the new parameters to measure the two-loop deviations for both ${\\cal F}$-$SU(5)$ and mSUGRA. We find that the deviations from the mSUGRA gaugino mass relation in mSUGRA are smaller than 5\\%, while the deviations in the ${\\cal F}$-$SU(5)$ are larger as a result of the TeV scale vector-like particles. An analytical comparison of the deviation between mSUGRA and ${\\cal F}$-$SU(5)$ is illustrated. Next, the expected observable final states at the LHC are given for all viable regions of the parameter space. Lastly, the proton lifetime is calculated for the experimentally allowed regions of the ${\\cal F}$-$SU(5)$ parameter space, and we show that the rate of predicted decay is indeed within the reach of the future Hyper-Kamiokande and DUSEL experiments. ", "conclusions": "We have considered gravity mediated supersymmetry breaking in ${\\cal F}$-$SU(5)$. The gaugino masses are not unified at the traditional grand unification scale, though we do indeed obtain the mSUGRA one-loop gaugino mass relation at the electroweak scale. However, the gaugino mass relation will have a small two-loop deviation, and this deviation may be measurable at the LHC and ILC. There is a considerable need for TeV scale vector-like particles in ${\\cal F}$-$SU(5)$, while unflipped GUTs, such as mSUGRA or F-$SU(5)$, require no such vector-like particles. In light of this key distinction between flipped and unflipped GUTs, we introduced a parameter to measure the small two-loop deviation from the mSUGRA gaugino mass relation at the electroweak scale. To implement a numerical analysis, we modified a popular and well-established public RGE code to incorporate the effects of TeV scale vector-like particles. In this work, we employed two-loop RGEs for the gauge couplings, but only considered one-loop RGEs for the Yukawa couplings and soft-supersymmetry breaking terms, though we look to extend this to all two-loop supersymmetry breaking RGEs in our future work. The results lead us to conclude there is a clear disparity in the extent of the deviation between GUTs with vector-like particles, such as ${\\cal F}$-$SU(5)$, and GUTs without vector-like particles, such as mSUGRA or F-$SU(5)$. The predicted correlation between largeness of the deviation from the mSUGRA gaugino mass relation and the existence of vector-like particles can be tested at the colliders. Furthermore, we have determined the viable parameter space of this model which simultaneously satisfies all the current experimental constraints and is consistent with the findings of CDMS II. The cross-section of two neutralinos into two gamma-rays for the experimentally allowed regions of the parameter space was computed and assessed against the first published Fermi-LAT measurement. The results showed the ${\\cal F}$-$SU(5)$ parameter space is consistent with the recent Fermi-LAT findings. Finally, we have calculated the proton lifetime for these experimentally allowed regions, and found it to be within the reach of the future Hyper-Kamiokande and DUSEL experiments. A wealth of experimental data is on the horizon, so it is imperative that phenomenologically appealing GUTs, for instance ${\\cal F}$-$SU(5)$, be researched so that unambiguous experimental predictions may be presented. These predictions will be key milestones in deducing the underlying theory at high energies as we progress through the next few exciting years of LHC, direct dark matter detection, Fermi-LAT, and proton decay experiments. We have supplemented the conventional bottom-up analysis of traditional GUTs to include TeV scale vector-like particles, and our results feature encouraging prospects for the experimental determination of whether high-energy theory indeed admits these proposed multiplets. We believe that in the next few years experiment will certainly have something key to say about ${\\cal F}$-$SU(5)$ in particular and string theory in general." }, "1003/1003.0841_arXiv.txt": { "abstract": "We present forecasts for constraints on deviations from Gaussian distribution of primordial density perturbations from future high--sensitivity X--ray surveys of galaxy clusters. Our analysis is based on computing the Fisher--Matrix for number counts and large-scale power spectrum of clusters. The surveys that we consider have high--sensitivity and wide--area to detect about $2.5\\times 10^5$ extended sources, and to provide reliable measurements of robust mass proxies for about $2\\times 10^4$ clusters. Based on the so-called self-calibration approach, and including Planck priors in our analysis, we constrain at once nine cosmological parameters and four nuisance parameters, which define the relation between cluster mass and X--ray flux. Because of the scale dependence of large--scale bias induced by local--shape non--Gaussianity, we find that the power spectrum provides strong constraints on the non--Gaussianity $\\fnl$ parameter, which complement the stringent constraints on the power spectrum normalization, $\\sigma_8$, from the number counts. To quantify the joint constraints on the two parameters, $\\sigma_8$ and $\\fnl$, that specify the timing of structure formation for a fixed background expansion, we define the figure-of-merit $\\fom=\\left( \\det\\left[Cov(\\sigma_8,\\fnl)\\right]\\right)^{-1/2}$. We find that our surveys constrain deviations from Gaussianity with a precision of $\\Delta \\fnl\\simeq 10$ at $1\\sigma$ confidence level, with $\\fom \\simeq 39$. We point out that constraints on $\\fnl$ are weakly sensitive to the uncertainties in the knowledge of the nuisance parameters. As an application of non--Gaussian constraints from available data, we analyse the impact of positive skewness on the occurrence of XMMU-J2235, a massive distant cluster recently discovered at $z\\simeq 1.4$. We confirm that in a WMAP-7 Gaussian $\\Lambda$CDM cosmology, within the survey volume, $\\simeq 5\\times 10^{-3}$ objects like this are expected to be found. To increase the probability of finding such a cluster by a factor of at least 10, one needs to evade either the available constraints on $\\fnl$ or on the power spectrum normalization $\\sigma_8$. ", "introduction": "The standard inflationary scenario, based on the single scalar field slow-roll paradigm, predicts primordial density perturbations to be virtually indistinguishable from a Gaussian distribution. However, a number of variants of inflation have been proposed which are able to generate a certain amount of non-Gaussianity \\citep[e.g.,][]{bartolo04,Chen10}. Therefore, testing to what precision we can measure possible deviations from Gaussianity with available and future observations has important implications on our understanding of the mechanism that seeded density fluctuations in the early Universe. Analyses of the Cosmic Microwave Background (CMB) provide at present the tighest constraints on the amount of allowed non-Gaussianity. A number of analyses based on the WMAP data converge to indicate consistency with the Gaussian assumption (e.g., \\citealt{komatsu10} and references therein; cf. also \\citealt{Yadav_Wandelt07}). While data from the Planck satellite are expected to further tighten such constraints \\citep[e.g.,][]{Yadav_etal07,Liguori_etal10}, it is worth understanding whether non-Gaussianity can be probed by large--scale structure observations \\citep[e.g.,][]{SL08.1,verde10}.\\\\ Non-Gaussian perturbations are expected to leave their imprint also on the pattern of structure growth at least in two different ways. First, we expect that a positively skewed distribution provides an enhanced probability of finding large overdensities. This translates into an enhanced probability of forming large collapsed structures at high redshift, thereby changing the timing of structure formation and the shape and evolution of the mass function of dark-matter halos. After the first pioneering studies of the effect of non-Gaussiantiy on the mass function \\citep[e.g.][]{Matarrese_etal86,Colafrancesco_etal89,Borgani_etal90}, a number of analyses have been carried out since the beginning of 2000s \\citep[e.g.,][]{matarrese00,mathis04,kang07,sefusatti07,grossi07, maggiore09}. Furthermore, the realization of large-scale cosmological simulations with non-Gaussian initial conditions have recently provided a validation of the non-Gaussian correction to be applied to the Gaussian mass function \\citep[e.g.,][]{grossi07, DA08.1,DE08.1,grossi09,Giannatonio_Porciani09,pillepich10}. More recently non-Gaussianity effects on the large-scale distribution of collapsed halos were studied: it has been demostrated \\citep[e.g.,][]{DA08.1,matarrese08,valageas09,Lam_Sheth09} that non-Gaussianity affects the large-scale clustering of halos in such a way that the linear biasing parameter acquires a scale dependence. This modifies, in a detectable way, the power spectrum of the distribution of any tracer of cosmic structures at small wavenumbers and offers a unique way of testing the nature of primordial fluctuations. The evolution of the mass function of galaxy clusters identified in X--ray surveys has been extensively used in the past to constrain cosmological models \\citep[e.g.,][]{borgani01,rosati02,schuecker03,voit05}. These studies have recently attracted renewed interest, thanks to detailed follow-up observations of clusters selected from ROSAT observations, they have been carried out either to constrain the Dark Energy equation of state within the Gaussian paradigm \\citep[e.g.,][]{vikhlinin09c,mantz09I}, or to test possible deviations from standard gravity \\citep[e.g.,][]{Schmidt_etal09,rapetti09}. \\citet{jimenez09} have recently analyzed the effect of non-Gaussianity on the population of massive high-redshift clusters, like the one recently discovered by \\citet{jee09} at $z\\simeq 1.4$. It is important to remember that the cosmological constraints obtained from clusters so far have been derived from small ROSAT--based samples, containing $\\sim 100$ clusters at $z < 1$. It is therefore easy to imagine the vast margin of improvements offered by next generation X-ray surveys which will detect $\\sim 10^5$ clusters to $z \\sim 2$. These surveys should cover a large enough volume at high redshift to test non-Gaussianity in the regimes where its effects are clearer, namely the high-mass tail of the mass function and the large-scale power spectrum of the cluster distribution. \\citet{fedeli09} and \\citet{roncarelli10} presented predictions for the number counts and clustering of galaxy clusters expected from the eROSITA X--ray survey and from the Sunyaev--Zeldovich SPT survey. While these analyses confirmed the potential of these surveys to provide interesting constraints on non-Gaussian models, they did not include detailed forecasts on the constraints on non-Gaussian models and a detailed assessment of the effect of uncertainties in the scaling relations between cluster masses and observables. \\citet{oguri09} followed the self--calibration approach by \\citet{lima05} (see also \\citealt{majumdar04,battey03}) to forecast the capability of future optical cluster surveys to constrain non--Gaussian models. This study showed that combining number counts and clustering of galaxy clusters can potentially provide quite strong constraints on deviations from Gaussianity. The aim of this paper is to derive forecasts, based on the Fisher--Matrix approach, on the capability of future X--ray cluster surveys to constrain deviations from Gaussian perturbations. Besides focusing on the characteristics of X-ray, rather than optically selected samples, our analysis differs from that by \\citet{oguri09} for the method to include information from large--scale clustering. \\citet{oguri09} adopted the approach by \\citet{lima05} where clustering is included by accounting for fluctuations of cluster counts within cells having a fixed angular size. This implies that, at each redshift, clustering information is restricted to one physical scale. In our analysis, we follow the approach originally presented by \\citet{tegmark97} (see also \\citealt{feldmann94,majumdar03}), in which the clustering Fisher-Matrix is computed for the allowed range of wavenumbers, by weighting them according to the effective volume covered by the surveys. Another distinctive aspect of our analysis is that it is based on X-ray surveys of next generation, whose sensitivity and angular resolution are high enough to warrant an accurate measurement of robust mass proxies, related to the cluster gas mass and X--ray temperature for a large number of clusters. As we will discuss in the following, surveys with these characteristics can be provided by an already proposed X--ray telescope, which joins a large collecting area to a large field-of-view and a high angular resolution over the entire field of view \\citep[e.g.,][]{giacconi09,vikhlinin09b}. The great advantage of having a similar survey is that there is no need to assume any external follow-up observation for a subset of identified clusters. Moreover, the possibility to define a flux-limit down to which measuring accurate mass proxies for all clusters allows one to set robust priors on the scaling relations between cluster mass and observables, which is one of the main source of uncertainty in the cosmological application of galaxy clusters \\citep[e.g.,][]{albrecht09}. In principle, the method used in our analysis can be applied to any cluster surveys, including optical and SZ ones. Although so far X-ray surveys have been mostly used for cosmological applications of clusters, upcoming large optical and SZ surveys promise to provide an important contribution to this field. Our method only requires a well defined selection function and calibrated mass proxies. Since cluster surveys at different wavelengths have different efficiencies probing different mass ranges at different redshifts, they will ultimately provide complementary approaches to the derivation of cosmological constraints. The plan of this paper is as follows. In Section 2 we summarize the formalism to compute non--Gaussian corrections to the mass function and the linear bias parameter of collapsed halos. In Section 3 we describe our approach to compute the Fisher Matrix for both the number counts and the power spectrum of galaxy clusters. In Section 4 we first describe how we compute the selection function and the redshift distribution expected for the X--ray surveys, then we present the results in terms of constraints on the parameter space defined by the non--Gaussianity parameter, $\\fnl$, and the power spectrum normalizations, $\\sigma_8$. Section 5 is devoted to the discussion of these results. In this section we will also discuss the competing effects of non-Gaussianity and normalization of the power spectrum on the expected number of clusters at $z>1.4$, which have a mass of, at least, $5\\times 10^{14}M_\\odot$, as the one recently studied by \\citet{jee09}. We summarize our main conclusions in Section 6. ", "conclusions": "We presented forecasts on the capability of future high-sensitivity X--ray surveys of galaxy clusters to provide constraints on deviations from Gaussian primordial perturbations. Our analysis is based on computing the Fisher Matrix (FM) for the information given by the evolution of mass function and power spectrum of galaxy clusters. Following the approach by \\citet{tegmark97} to compute the power-spectrum FM \\citep[see also][]{majumdar04,rassat08,stril09}, we include in the analysis the information related to the possible scale--dependence of the linear bias, which represents a unique fingerprint of non-Gaussianity \\citep[e.g.,][and references therein]{verde10}. According to the self--calibration approach, the model parameters entering in the FM estimate are nine cosmological parameters and 4 nuisance parameters, the latter defining the relation between cluster mass and observable upon which cluster selection is based. Our analysis is based on assuming an observational strategy designed for the Wide Field X--ray Telescope (WFXT, e.g. \\citealt{giacconi09}), in which a {\\em Wide} Survey covering most part of the extragalactic sky is complemented by a {\\em Medium} and by a {\\em Deep} Survey (see Table \\ref{t:sur}). The latter provides mass proxies down to the flux limit for cluster identification in the Wide Survey (see Table 1). We showed forecasts for the two parameters that, for a fixed expansion history, define the timing of cosmic structure formation, namely $\\sigma_8$ and $\\fnl$, while marginalizing over all the remaining parameters. Informations on such constraints are quantified by introducing the figure-of-merit for structure formation timing of Eq.(\\ref{eq:fom}). The main results obtained from our analysis can be summarized as follows. \\begin{description} \\item[(a)] Power spectrum and number counts of galaxy clusters are highly complementary in providing constraints: while the former is sensitive to deviations from Gaussianity, through the scale dependence of the bias, the latter is mostly sensitive to $\\sigma_8$. \\item[(b)] Most of the constraining power for these two parameters lies in the {\\em Wide} Survey, while the {\\em Medium} and the {\\em Deep} Surveys play an important role for the estimate of X--ray mass proxies for $\\simeq 2\\times 10^4$ clusters out to $z\\sim 1.5$. \\item[(c)] Combining number counts and power spectrum information for the three surveys turns into $\\Delta \\fnl\\simeq 10$ for the $1\\sigma$ uncertainty with which a deviation from Gaussianity associated to a ``local shape'' model can be constrained. Correspondingly, we find $\\fom \\simeq 39$ for the figure-of-merit of structure formation timing. \\item[(d)] Quite interestingly, while the value of $\\fom$ significantly worsens when assuming more conservative priors on the nuisance parameters, the above constraint on $\\fnl$ is weakly sensitive on such priors. \\item[(e)] The presence of a cluster as massive as XMMU-J2235.3 at $z\\simeq 1.4$ \\citep{jee09} turns out to be a rather unlikely event, even allowing for an amount of non-Gaussianity consistent with current CMB \\citep[e.g.,][]{komatsu10} and LSS \\citep{SL08.1} constraints. This further demonstrates the strong constraining power of detecting an even small number of massive high-$z$ clusters. \\end{description} Our analysis lends support to the important role that future cluster surveys will play in constraining deviations from the Gaussian paradigm, with far reaching implications on the primordial mechanisms which seeded density inhomogeneities. The reliability of our forecasts relies on the possibility of calibrating to high precision a universal expression for mass function and large--scale bias. A number of independent groups \\citep{grossi07,DA08.1,DE08.1,grossi09,Giannatonio_Porciani09,pillepich10} carried out large N--body simulations with non--Gaussian initial conditions, finding in general a quite good agreement for the calibration of both mass function and bias. However, the precision required for the calibration of such quantities, for them not to spoil the constraining power of large surveys, is probably higher than what reached at present. First assessments of the impact of uncertainties in the mass function calibration on DE constraints have been already presented \\citep[e.g.][]{wu09} and indicate that such uncertainties may not be negligible. There is no doubt that larger suites of non--Gaussian simulations are required to calibrate mass function and large-scale bias also for a range of models beyond the local non--Gaussian models that we considered in the analysis presented here. A few days after our paper, \\citet{cunha10} also subbitted a paper regarding the study of constraints on non-Gaussian parameter $f_{NL}$ from cluster surveys. They use the Fisher matrix approach applied to the count-in-cell method to extract information on evolution of cluster number density and clustering. Differently from the previous paper by \\cite{oguri09}, \\citet{cunha10} also included the contribution from covariance between counts in different cells. This allowed them to sample the power spectrum over a large scale range, thus obtaining tighter constraints on $f_{NL}$ than \\cite{oguri09}. By specialising their analysis for the mass selection and sky coverage expected for the DES\\footnote{https://www.darkenergysurvey.org/} optical survey, they forecast a precision of $\\sigma_{f_{NL}}\\simeq 1$--5. A detailed comparison between our and their analysis is not straightforward. Besides using different survey specifications, our and their analyses also uses different prescriptions for the mass function and the bias. Just as an example of how sensitive the choice of the bias model is, we verified that excluding the bias correction suggested by \\citet{grossi09} (the factor $q=0.75$ that we introduce after Eq.\\ref{eq:bias}), the expected errors in $f_{NL}$ would decrease by approximately a factor of 2. This further emphasizes the need for a precise calibration of the model mass function and bias through extended sets of non-Gaussian N--body simulations." }, "1003/1003.5655_arXiv.txt": { "abstract": "The cratering history of main belt asteroid (2867) Steins has been investigated using OSIRIS imagery acquired during the Rosetta flyby that took place on the 5$^{th}$ of September 2008. For this purpose, we applied current models describing the formation and evolution of main belt asteroids, that provide the rate and velocity distributions of impactors. These models coupled with appropriate crater scaling laws, allow the cratering history to be estimated. Hence, we derive Steins' cratering retention age, namely the time lapsed since its formation or global surface reset. We also investigate the influence of various factors -like bulk structure and crater erasing- on the estimated age, which spans from a few hundred Myrs to more than 1~Gyr, depending on the adopted scaling law and asteroid physical parameters. Moreover, a marked lack of craters smaller than about $0.6$~km has been found and interpreted as a result of a peculiar evolution of Steins cratering record, possibly related either to the formation of the 2.1~km wide impact crater near the south pole or to YORP reshaping. ", "introduction": "The European Space Agency's (ESA) Rosetta spacecraft passed by the main belt asteroid (2867) Steins with a relative velocity of 8.6~km/s on 5 September 2008 at 18:38:20 UTC. The Rosetta-Steins distance at closest approach (CA) was 803~km. During the flyby the solar phase angle (sun-object-observer) decreased from the initial 38$^{\\circ}$ to a minimum of 0.27$^{\\circ}$ two minutes before CA and increased again to 51$^{\\circ}$ at CA, to reach 141$^{\\circ}$ when the observations were stopped. A total of 551 images were obtained by the scientific camera system OSIRIS, which consists of two imagers: the Wide Angle Camera (WAC) and the Narrow Angle Camera (NAC) \\citep{kel07}. The best resolution at CA corresponded to a scale of 80~m/px at the asteroid surface.\\\\ Steins has an orbital semi-major axis of about 2.36~AU, an eccentricity of 0.15 and an inclination of 9.9$^{\\circ}$. It is therefore orbiting in a relatively quiet region of the main belt, far from the main escape gateways, namely the secular $\\nu_6$ and mean motion 3:1 resonances. Its shape can be fitted by an ellipsoid having axis of $6.67 \\times 5.81 \\times 4.47$~km \\citep{kel10}.\\\\ Previous space missions have visited and acquired detailed data for a total of 5 asteroids, namely three main belt asteroids \\citep[951 Gaspra, 243 Ida, 253 Mathilde;][]{vev99a,bel92,bel94} and two near-Earth objects \\citep[433 Eros, 25143 Itokawa;][]{vev99b,sai06}. Itokawa is the smallest of them, with dimensions of $0.45 \\times 0.29 \\times 0.21$~km. The other asteroids have average sizes ranging from about 12~km to 53~km. In this respect, Steins with its 5.3~km size lies between Itokawa and the large asteroids. It is therefore the smallest main belt asteroid ever imaged by a spacecraft (except for Ida's satellite Dactyl with a diameter of roughly a km). Moreover, Steins is a member of the relatively rare E-type class (composed by igneous materials), while other asteroids visited by spacecraft are members of the most common S- and C-type classes. Previous spacecraft observations opened a new field of investigation, namely the cratering of asteroidal surfaces. A number of interesting processes were therefore studied with unprecedented detail, like the cratering on low gravity bodies, regolith formation, seismic shaking \\citep[e.g.][]{cha02}.\\\\ This paper analyzes some of the highest resolution OSIRIS images with the aim to study the crater size distribution and derive the chronology of the impacts on the surface of the asteroid. This will also provide clues on the Steins bulk structure, evolution, and give new insights on the above mentioned processes.\\\\ ", "conclusions": "In this work we used crater counting, morphological analysis and impactor population modeling to constrain Steins' cratering retention age. The derived ages vary according to the SLs used. In particular, using NSL and crater erasing the derived age is $0.154 \\pm 0.035$~Ga, while using HSL and crater erasing the age ranges from 0.49 to 1.6~Ga, according to the values of strength used. Moreover, the modeling of the crater erasing processes shows that the observed crater density is not saturated (at least for the parameters adopted here). The mean collisional age of Steins is estimated to be $\\sim2.2$~Ga \\citep[e.g.][]{mar06}. Interestingly, similar numbers apply also for Gaspra. Analogue conclusions might be also valid for the near-Earth objects Eros, whose mean collisional age is $\\sim1.7$~Ga \\citep{mar06} while cratering age using NSL gives 0.12~Ga \\citep{obr06} or, using HSL, 1-2~Ga \\citep{mic09}. The larger bodies Ida and Mathilde have crater populations either close to the saturation or saturated, and consequently their cratering age estimate is less constrained. Despite the low number statistics, the cratering age of the main belt asteroids smaller than $\\sim20$~km seems to be systematically younger than their collisional age. This result, if confirmed by further studies, would have important implications on main belt collisional models. \\\\ Notably, the shape of Steins crater size distribution shows a kink for diameters smaller than 0.5-0.6~km, which may require a recent episode of intense erasing, although seismic shaking could have potentially played a role in producing the observed distribution as well. We also attempt to constrain the epoch of such episode, possibly associated to the Ruby crater formation. Focusing on the small diameter end ($D<0.35$~km) of the crater cumulative distribution, and adopting the MPF fitting we obtain an age of $0.032 \\pm 0.004$~Ga using NSL, from 0.072 to 0.237~Ga using HSL. Under the assumption that the formation of the Ruby crater erased all craters $<0.5-0.6$~km, the above ages could possibly indicate the time since the occurrence of the Ruby event. A lower limit to the age of this event, can be derived considering the time required to accumulate the observed craters in the range $0.2-0.3$~km. This produces an age from $\\sim2$~Ma to $\\sim10$~Ma, for NSL and HSLs respectively.\\\\ The derived ages vary up to a factor of ten depending on the SL and the tensile strength used. In the present work we investigated the effects of a relatively large range of SLs and $Y$. However, in the light of our global understanding of Stein properties, we favor a crater retention age ranging from $\\sim0.15$ to $\\sim0.5$~Ga, and a kink related event that could be as young as a few Ma up to some tens of Ma.\\\\ As a final remark, note that the conclusions derived in this paper are based on the bona fide crater distribution. The general scenario outlined (age, depletion of small craters) remains valid also if considering all crater-like features (see fig. \\ref{steins}, lower panel). In the latter case, however, the age of the reset event is about a factor of two higher. \\\\ Acknowledgment We thank A.~Morbidelli for providing us with the impact probability file for Steins. We are also grateful to the referee D.~O'Brien for many helpful comments and for providing us the published data of Itokawa and Eros cratering record. Thanks also to a second anonymous referee for insightful suggestions. Finally, we thank E.~Simioni and E.~Martellato for discussions. \\newpage" }, "1003/1003.1981.txt": { "abstract": "We explore the shape of the galaxy luminosity function (LF) in groups of different mass by creating composite LFs over large numbers of groups. Following previous work using total group luminosity as the mass indicator, here we split our groups by multiplicity and by estimated virial (group halo) mass, and consider red (passive) and blue (star forming) galaxies separately. In addition we utilise two different group catalogues (2PIGG and Yang et al.) in order to ascertain the impact of the specific grouping algorithm and further investigate the environmental effects via variations in the LF with position in groups. Our main results are that LFs show a steepening faint end for early type galaxies as a function of group mass/ multiplicity, with a much suppressed trend (evident only in high mass groups) for late type galaxies. Variations between LFs as a function of group mass are robust irrespective of which grouping catalogue is used, and broadly speaking what method for determining group `mass' is used. We find in particular that there is a significant deficit of low-mass passive galaxies in low multiplicity groups, as seen in high redshift clusters. Further to this, the variation in the LF appears to only occur in the central regions of systems, and in fact seems to be most strongly dependent on the position in the group relative to the virial radius. Finally, distance-rank magnitude relations were considered. Only the Yang groups demonstrated any evidence of a correlation between a galaxy's position relative to the brightest group member and its luminosity. 2PIGG possessed no such gradient, the conclusion being the FOF algorithm suppresses the signal for weak luminosity--position trends and the Yang grouping algorithm naturally enhances it. ", "introduction": "It is well known that galaxy properties, such as their morphology, star formation rate and colour, vary with the density of their local environment \\citep{dres80, lewi02, kauf04}. Unsurprisingly, this is also seen as a variation of the galaxy luminosity function (LF) with environment \\citep[e.g.][]{ferg91, driv98, phil98, zabl00, tren02, hogg03, prac05, crot05}. These variations are likely due to some combination of the suppression of star formation once a threshold local density is reached \\citep[e.g.][]{balo04, tana05} and the significant differences in the interaction and/or merger histories between environments \\citep[e.g.][]{moor98, bell06}. In a previous paper \\citep{robo06} we explored the variation of galaxy populations across the mass range from small groups up to clusters, in terms of the LF parameters for `composite' groups. Using groups from the 2PIGG catalogue \\citep{eke04a}, we showed that the characteristic magnitude $M^*$ brightened and the faint end slope $\\alpha$ steepened with increasing group luminosity. That is, assuming that the total luminosity of the group members is a reasonable proxy for the total mass of the (group or cluster sized) halo containing them, more massive halos contain both brighter galaxies and larger fractions of faint galaxies. This, and the variations seen separately in the blue and red galaxy sub-populations, appeared consistent with a generic picture of the build up of the giant cluster galaxies through mergers and the amplification of the dwarf end of the LF as `field' irregulars `fall in' to larger groups and cluster and are transformed, through truncation of star formation, into dwarf ellipticals. Most of this transformation seemed to take place in quite low luminosity groups, i.e.\\ quite early in the hierarchical build-up of the larger systems \\citep[e.g.][]{hash00,wilm05}. In the current paper, we extend this work \\citep[and that of][]{eke04b} by using different LF estimators, several different ways of dividing the groups by `size' and two different group catalogues. In addition we explore the `environment' in more detail by exploring radial LF variations within groups/clusters . Section 2 describes the input group catalogues, noting in particular the differences in grouping algorithms and the influence this may have on the subsequent results. Section 3 discusses the various (group halo) mass proxies and derives the LF parameters for each of our composite groups obtained from the 2PIGG catalogue. Section 4 presents the equivalent discussion for groups in the catalogue of \\citet{yang05b}, Section 5 considers the question of the positional dependence of galaxy luminosities within groups and Section 6 summarises our results. We assume a standard cosmology with $\\Omega_M = 0.3$, $\\Omega_{\\Lambda} = 0.7$ and a Hubble constant $H_0 = 100h\\,$km$\\,$s$^{-1}$. ", "conclusions": "In this paper a detailed comparison between the grouping characteristics of the 2PIGG and Yang group catalogues was made, with particular reference to the group luminosity functions and how these are affected by the different grouping algorithms. We have extended the work of Eke et al. (2004a) and Robotham et al. (2006) on the luminosity function of galaxies in groups of different masses by (a) using a maximum likelihood method for determining the LF parameters, (b) using different indicators of the group halo mass, viz. a mass direct virial mass estimate and the number of galaxies in the group. For 2PIGG, virial masses were calculated as in Eke et al. (2004b) from the velocity dispersion and the rms distance from the group centre of the member galaxies. The same trends as previously found in \\citet{robo06} are still generally present, but are somewhat less well defined for $M^*$, especially in the case of blue galaxies. The faint end slope, $\\alpha$, for the total population changes in a similar way to that seen in previous work, but the subsets behave significantly differently: red galaxies now showing a large steepening, blue ones little or no change until we reach the highest masses sampled. In terms of ordering by multiplicity, it appears that only the groups with very few (giant) members differ substantially in $M^*$ from the larger objects. This is also the case for $\\alpha$, though the blue galaxies in low multiplicity groups have the least negative slope of any blue galaxies sampled. This appears to imply that a system being numerically poor has more effect on the (lack of) blue dwarfs than the group's mass. This effect is also seen if we consider the relative numbers of galaxies brighter than and fainter than $M_{b_{j}} = -18.25$ (for groups sufficiently nearby for the latter to be visible). A surprising number of the lowest multiplicity groups have few or no galaxies in the fainter range (typically $M_{b_{j}}$ between $-18.25$ and $-17.25$). The Yang-2dF groups provide a useful counterpoint to the 2PIGG groups as they were constructed via an algorithm which essentially works via the binding of a galaxy to a pre-existing group core. In general terms, the Yang-2dF groups demonstrate the same LF parameter trends as before. However, there are a number of distinct differences in detail. Firstly, even for groups clearly centred at the same position in the two catalogues (i.e. representing the same halo), the LF shapes appear to diverge somewhat for the fainter galaxies, though the LF parameters are closely similar. Second, the 2PIGG groups {\\em not} closely matched by a Yang-2dF group have a fainter $M^*$ and flatter $\\alpha$ than the rest, whereas unmatched Yang-2dF groups have brighter $M^*$ and steeper $\\alpha$, though we should also note that the LFs for these subsets are not, in fact, well fitted by a Schechter function. A two component fit would be preferred for the Yang-2dF groups. In terms of the LF shapes, it appears that the matched and unmatched group LFs differ primarily through an excess (shortfall) of intermediate luminosity galaxies in the unmatched 2PIGG (Yang-2dF) groups. These unmatched groups are important as they are the ones which most clearly reflect the behaviour of the grouping algorithms. It appears that 2PIGG will quite generously link objects where there are numerous medium luminosity galaxies, whereas the Yang algorithm does not (presumably because they do not add sufficient mass to make the group encompass further objects). Physically, this indicates that groups with dominant central galaxies or dense central cores, of the sort the Yang algorithm is designed to find, are relatively deficient in medium mass objects, consistent with them being more dynamically evolved. 2PIGG, on the other hand, may tend to link subgroups, adding additional moderate mass objects which are/were the central galaxy of their own (infalling) halo. Among the Yang-2dF groups themselves, while the errors are larger than for 2PIGG, the mass (here really just a scale version of luminosity) ordered composites show the same {\\em trends} as the 2PIGG groups, with $M^*$ brightening and $\\alpha$ steepening for the more luminous groups. However, as implied by the comparison above, the actual {\\em values} differ, sometimes substantially. In particularly $M^*$ is $\\sim 0.5$ magnitudes or more brighter at all masses for Yang-2dF groups than 2PIGG groups. Red galaxies have consistently steeper $\\alpha$ in the Yang-2dF groups, while, given the errors the blue galaxies are consistent with the same slopes as those in 2PIGG (or, indeed, with no change with group mass). The multiplicity ordered Yang-2dF groups show the same changes relative to 2PIGG multiplicity ordered groups as seen for the mass ordered groups, and as in 2PIGG the least negative $\\alpha$ is seen for red galaxies in low multiplicity groups, blue galaxies having a similar $\\alpha$ at all multiplicities. Thus the Yang-2dF groups contain a larger fraction of red early type galaxies at the same group mass or multiplicity than seen in 2PIGG and this difference is largest for the highest multiplicity groups. Generally, this deficit of low-mass passive galaxies can be directly interpret as a build up of the red sequence as a function of halo mass, and indirectly as a build up of the red sequence with redshift. LFs calculated for galaxies within 0.5 $r_{rms}$ and outside 1 $r_{rms}$ show that the outer regions have similar LF shapes for all group multiplicities for both red and blue galaxies. The central LF for late type galaxies changes slightly with multiplicity, but the major change is for centrally located early type galaxies; $\\alpha$ is significantly steeper than for the outer regions at all multiplicities and steepens even more for higher multiplicities. Looked at in a slightly different way, we have seen that the Yang-2dF groups tend to be denser and probably more evolved, and the change in population with radius for such systems is reflected in a strong rank correlation between luminosity and group-centric distance within these groups. Thus we should finally conclude that the distribution of luminosities of galaxies in groups depends {\\em both} on the mass of the group, as determined via any suitable (and consistent) manner {\\em and} on the (density) structure of the group, so that different grouping algorithms will generate different luminosity distributions for the grouped galaxies. Finally distance-rank magnitude relations were considered. Only the Yang groups demonstrated any evidence of a correlation between a galaxy's position relative to the centre and its luminosity (i.e.\\ brighter galaxies have a weak tendency to be nearer the group centre). 2PIGG possessed no such gradient, the conclusion being the non-prescriptive nature of luminosity ordering in the FOF algorithm destroys any weakly observable trend. A comparison was made to the semi-analytic galaxy catalogue of De Lucia et al. (2006) built on top of the Millennium Simulation. Interestingly this theoretical work demonstrated much stronger gradients than present even in Yang-2dF groups, and may indicate a semi-analytic formalism that is too stringent." }, "1003/1003.4293_arXiv.txt": { "abstract": "We analyze the behavior of the parsec-scale jet of the quasar 3C~454.3 during pronounced flaring activity in 2005-2008. Three major disturbances propagated down the jet along different trajectories with Lorentz factors $\\Gamma>$10. The disturbances show a clear connection with millimeter-wave outbursts, in 2005 May/June, 2007 July, and 2007 December. High-amplitude optical events in the $R$-band light curve precede peaks of the millimeter-wave outbursts by 15-50~days. Each optical outburst is accompanied by an increase in X-ray activity. We associate the optical outbursts with propagation of the superluminal knots and derive the location of sites of energy dissipation in the form of radiation. The most prominent and long-lasting of these, in 2005 May, occurred closer to the black hole, while the outbursts with a shorter duration in 2005 Autumn and in 2007 might be connected with the passage of a disturbance through the millimeter-wave core of the jet. The optical outbursts, which coincide with the passage of superluminal radio knots through the core, are accompanied by systematic rotation of the position angle of optical linear polarization. Such rotation appears to be a common feature during the early stages of flares in blazars. We find correlations between optical variations and those at X-ray and $\\gamma$-ray energies. We conclude that the emergence of a superluminal knot from the core yields a series of optical and high-energy outbursts, and that the mm-wave core lies at the end of the jet's acceleration and collimation zone. We infer that the X-ray emission is produced via inverse Compton scattering by relativistic electrons of photons both from within the jet (synchrotron self-Compton) and external to the jet (external Compton, or EC); which one dominates depends on the physical parameters of the jet. A broken power-law model of the $\\gamma$-ray spectrum reflects a steepening of the synchrotron emission spectrum from near-IR to soft UV wavelengths. We propose that the $\\gamma$-ray emission is dominated by the EC mechanism, with the sheath of the jet supplying seed photons for $\\gamma$-ray events that occur near the mm-wave core. ", "introduction": "During the past four years, the quasar 3C~454.3 (z=0.859) has displayed pronounced variability at all wavelengths. In spring 2005 it returned to the night sky with unprecedented brightness, $R\\sim$12.0~mag, a level not seen at optical wavelengths over at least 50 years of observations \\citep{VIL06}. An increase in activity occurred at X-ray and radio wavelengths as well, with the 230 GHz radio variations having a delay of $\\sim$2~months with respect to the optical variability \\citep{RAI08b}. This prominent outburst was followed by a more quiescent period at all wavebands from spring 2006 to spring 2007. During this interval, the optical spectrum possessed characteristics typical of radio-quiet active galactic nuclei (AGN), such as a ``big blue bump'' and ``little blue bump,'' attributed to thermal emission from the accretion disk and broad emission lines from surrounding clouds, respectively \\citep{RAI07}. After the quiescent state, the quasar underwent a new stage of high optical activity \\citep{RAI08a} that continued to the end of 2008. During this time span, very bright $\\gamma$-ray emission was detected \\citep{VER08,TOSTI08}, with an excellent correlation between the $\\gamma$-ray and near-infrared/optical variations \\citep{BON09}. Models proposed to explain the observed variability and spectral energy distribution (SED) of 3C~454.3 across the electromagnetic spectrum involve processes originating in the radio jet of the quasar. \\citet{VIL07} suggest that the very high optical flux in spring 2005 was connected with a disturbance (e.g., a shock) propagating along a curved trajectory in the jet, with optical synchrotron photons emitted over a different volume than the longer-wavelength radiation. As the emission zone of a given wavelength passes closest to the line of sight, the flux peaks at that wavelength. This occurs first at optical and later at longer wavelengths. \\citet{GHIS07} have found that the behavior of 3C~454.3 in 2005-2007 is consistent with the model suggested by \\citet{KG07}, in which the dissipation site of an outburst depends on the bulk Lorentz factor and compactness of the perturbation propagating down the jet. Outbursts occurring closer to the black hole (BH) should have a more compact emitting region with a lower bulk Lorentz factor, $\\Gamma$, and a stronger magnetic field, $B$. Greater compactness of the emission region intensifies the synchrotron flux as well as the high energy component produced via the synchrotron self-Compton (SSC) mechanism, while the external Compton (EC) high-energy component (inverse Compton radiation with seed photons from outside the jet) is suppressed owing to a weaker Doppler factor, $\\delta$, resulting from the lower value of $\\Gamma$. \\citet{GHIS07} model the outburst in 2005 as an event that occurred closer to the BH ($\\Gamma\\sim$8, $\\delta\\sim$13, $B\\sim$15~G, and size of the emitting region $a\\sim$5.5$\\times$10$^{-3}$~pc) than the outburst in 2007 ($\\Gamma\\sim$16, $\\delta\\sim$16, $B\\sim$9~G, and $a\\sim$8$\\times$10$^{-3}$~pc). \\citet{SMM08} argue that the optical, X-ray, and millimeter light curves during the outburst in 2005 require a release of a significant fraction of the jet energy when the jet becomes transparent at millimeter wavelengths at the millimeter-wave ``photosphere.'' These authors conclude that this photosphere is located at $\\sim$10~pc from the BH, coinciding with the expected location of a torus of hot dust. \\citet{SMM08} infer that the X-ray and $\\gamma$-ray emission is most likely produced via the EC mechanism, with seed photons emitted by the hot dust scattered by relativistic electrons in a plasma with bulk Lorentz factor $\\Gamma\\sim$20. Interpretations of $\\gamma$-ray observations with {\\it AGILE} during autumn 2007 and the densely sampled light curve provided by the {\\it Fermi} Gamma-ray Space telescope starting in 2008 August, combined with simultaneous observations at longer wavelengths, involve higher-energy electrons that emit synchrotron radiation at near-IR and optical wavelengths as well as scatter external photons from the broad line region to energies up to $\\sim$100~GeV \\citep{VER09,BON09}. \\citet{J05} monitored the quasar 3C~454.3 at 43~GHz with the Very Long Baseline Array (VLBA) bimonthly from March 1998 to April 2001, and determined parameters of the parsec-scale jet during this period based on the apparent speed of superluminal knots (also referred to as ``components'' of the jet) and time scale of their flux variability. They found that the jet of 3C~454.3 during this time span had physical parameters as follows: $\\Gamma$=15.6$\\pm$2.2, $\\Theta_\\circ$=1.3$^\\circ\\pm$1.2$^\\circ$, $\\delta$=24.6$\\pm$4.5, and $\\theta$=0.8$^\\circ\\pm$0.2$^\\circ$, where $\\Theta_\\circ$ is the viewing angle and $\\theta$ is the opening half-angle of the jet. The Boston University group resumed VLBA monitoring of the quasar in June 2005, and has continued to monitor the source within a program of roughly monthly VLBA imaging of bright $\\gamma$-ray blazars at 43 GHz. In this paper, we analyze disturbances seen in the quasar jet during the period of high optical activity in 2005-2008, and connect events in the jet with prominent variability at different wavebands. Observations cover the range of frequencies from 10$^{10}$~GHz to 10$^{23}$~GHz, which provide significantly different angular resolutions at which an event is observed - 10-15~arcminutes at high energy frequencies, $\\sim$5~arcseconds at millimeter and sub-millimeter wavelengths, $\\sim$1~arcsecond in the optical and IR bands, and $\\sim$0.1~milliarcsecond (mas) with the VLBA at 43~GHz. This renders interpretation of multifrequency behavior challenging, and makes analysis of variability at different wavelengths along with VLBI images the main tool for understanding processes and mechanisms involved in the physics of blazars. ", "conclusions": "We have found a connection between mm-wave and optical outbursts and disturbances propagating down the radio jet in 3C~454.3. At least three major optical outbursts, $O_{K1}, O_{K2}$, and $O_{K3}$, observed in 2005-2008 are related to superluminal knots identified in the VLBA images. \\citet{J07} have found that in blazars the polarization properties of superluminal knots at 43~GHz are consistent with weak shocks and the optical polarized emission originates in shocks, most likely situated between the 86~GHz and 43~GHz VLBI cores. It appears that shocks and their interaction with underlying jet structure might be generally responsible for optical and high energy outbursts. \\subsection{Location of the Optical Emission in the Jet} Comparison of the timing of optical outbursts and epochs of ejection of superluminal knots from the core can be used to determine the location of the sites where kinetic and internal energy is dissipated in the form of radiation. The delay between the peak of an outburst and the time of the ejection of the associated superluminal component, along with the apparent speed of the knot, determines the location of the optical outburst in the jet with respect to the core, $\\Delta r_{\\rm opt} = \\beta_{\\rm app}c\\Delta T_{\\rm opt}/\\sin\\Theta_\\circ$. According to Tables~\\ref{Kparm} and \\ref{Delparm}, the peaks of outbursts $O_{\\rm K1}, O_{\\rm K2}$, and $O_{\\rm K3}$ occurred 12.6$\\pm$5.6~pc upstream, 3.5$\\pm$3.7~pc downstream, and 2$\\pm$18~pc upstream of the core, respectively. The average location for $O_{\\rm K2}$ and $O_{\\rm K3}$ is 2.8$\\pm$3.4~pc downtream of the core. In the case of a conical structure of the jet, \\citet{J05} have estimated the half opening angle of the jet in 3C~454.3 to be $\\theta=0.8^\\circ\\pm0.2^\\circ$. We have obtained an average size of the core during our period of observations using modelling by circular Gaussians, $a_{\\rm core}$=0.068$\\pm$0.011~mas. This value of $a_{\\rm core}$ gives an upper limit for the transverse size of the core, which, along with the opening angle of the jet, provides an estimate of the distance of the mm-wave core from the central engine, $R_{\\rm BH}\\le$18$\\pm$3~pc. Therefore, if the perturbation that created $K1$ is reponsible for $O_{K1}$, then the maximum of the outburst took place when $K1$ was located at a distance 5.4$\\pm$5.2~pc from the central engine, while the maxima of $O_{\\rm K2}$ and $O_{\\rm K3}$ outbursts occurred at a larger distance from the BH, 21$\\pm$4~pc, perhaps when knots $K2$ and $K3$ passed through the mm-wave core. This implies that the dissipation site for $O_{\\rm K1}$ was closer to the central engine (BH) than that for outbursts $O_{\\rm K2}$ and $O_{\\rm K3}$. These results are consistent with the recently advanced ideas that the dissipation sites of optical and high energy outbursts can occur at different locations and involve different sources of seed photons \\citep{GHIS07,MAR08,MAR10,RITABAN08}. \\subsection{The Structure of the mm-Wave Core} The most powerful optical outburst, $O_{K1}$, preceded the passage of $K1$ through the 43~GHz core (ejection) by $\\sim$50~days. In addition, there are optical flares (RJD: 3537 and RJD: 3561, see Table~\\ref{Outparm}) that occurred within 1$\\sigma$ uncertainty of the time of ejection of $K1$. Outbursts $O_{K2}$ and $O_{K3}$ were simultaneous (within 1$\\sigma$ uncertainty) with the passage of $K2$ and $K3$, respectively, through the core. We isolate segments of the optical light curve within the time interval $-\\sigma T_\\circ$ and +2$\\sigma T_\\circ$ of each of knots $K1$, $K2$, and $K3$ (Fig.~\\ref{multi}, the values of $\\sigma T_\\circ$ are given in Table~\\ref{Kparm}). Each ejection is accompanied by a series of flares (denoted in Fig.~\\ref{multi} as 1, 2, and 3) that starts within 1$\\sigma$ of $T_\\circ$. The duration of the series (time between the first and third peaks) is comparable, 50, 50, and 65~days, for events $K1$, $K2$, and $K3$, respectively. The profile of event $O_{\\rm K1}$ is very different from the profiles of $O_{\\rm K2}$ and $O_{\\rm K3}$. The average of the rise and decay times for $O_{\\rm K1}$ has a width of $\\sim$47~days (Table~\\ref{Outparm}), with a rise time of 68~days and a decay of only 26~days. Outbursts $O_{K2}$ and $O_{K3}$, as well as the flares connected with the passage of $K1$ through the core, have much shorter durations and similar rise and decay times (see Fig.~\\ref{multi}). A possible interpretation is that, in the case of $O_{\\rm K1}$, located deeply within the ACZ, the energization of relativistic electrons is gradual and continuous, while for outbursts occurring in the VLBI core, relativistic electrons are energized abruptly owing to interaction between the disturbance and the core. The latter implies that the three-flare structure of optical outbursts seen in Figure~\\ref{multi} is related to the physical structure of the mm-wave core. Indeed, the super-resolved images of 3C~454.3 (Fig.~\\ref{hCore}) support the idea that the core at 43~GHz might have a three-component structure. The core might be a system of alternating conical shocks and rarefactions, as suggested by \\citet{GOM97}, \\citet{MAR06}, and \\citet{CAW06}. In this case, the distance between constrictions, $z_{\\rm max}$, can be calculated for an ultra-relativistic equation of state according to \\citet{DM88}: $z_{\\rm max}\\approx 3.3\\times\\Gamma~a_{\\rm core}^{\\rm t}/\\eta$, where $a_{\\rm core}^{\\rm t}$ is the transverse radius of the core in mas and $\\eta=p_{\\rm ext}/p_{\\circ,{\\rm jet}}$ is the ratio of the external pressure to initial pressure (i.e., at the upstream boundary of the core) in the jet. \\citet{J05} have found that the relation between the viewing angle and Lorentz factor in blazars is consistent with $\\eta\\sim$3. If the core of 3C~454.3 is a system of three conical shocks, then the longitudinal size of the core in projection on the sky can be determined as $a_{\\rm core}^{\\ell,{\\rm theor}}=3~z_{\\rm max}\\sin{\\Theta_\\circ}$. On the other hand, if the three-flare structure of the optical outbursts observed in the vicinity of the ejection is caused by a disturbance moving through this system of conical shocks, then $a_{\\rm core}^{\\ell,{\\rm obs}}=\\Delta T_{\\rm s}\\mu$, where $\\Delta T_{\\rm s}$ is the duration of the optical outbursts and $\\mu$ is the proper motion of the knot. Table~\\ref{Zmax} lists values of $a_{\\rm core}^{\\ell,{\\rm theor}}$ and $a_{\\rm core}^{\\ell,{\\rm obs}}$ for events $K1$, $K2$, and $K3$, which are consistent with the proposed association between a propagating disturbance and with the multi-component structure of the VLBI core providing a viable explanation for the three-flare pattern observed in the optical light curve. From standard gas dynamics, the presence of multiple standing shocks in three dimensions requires a high level of azimuthal symmetry of the pressure at the boundary of the jet in order for the structure to remain intact. \\subsection{Location of Millimeter-Wave Outbursts in the Jet} We have derived the locations of 230~GHz outbursts $M_{\\rm K1}$, $M_{\\rm K2}$, and $M_{\\rm K3}$ in the same manner as for optical outbursts $O_{\\rm K1}$, $O_{\\rm K2}$, and $O_{\\rm K3}$ owing their association with superluminal knots. The maxima of $M_{\\rm K1}$, $M_{\\rm K2}$, and $M_{\\rm K3}$ occurred when a disturbance was located 4.6$\\pm$6.1~pc upstream, 8.2$\\pm$4.8~pc downstream, and 15$\\pm$37~pc downstream of the core, respectively, with the average value 3.8$\\pm$4.5~pc downstream of the 43~GHz core. This implies that the passage of a disturbance through the 230~GHz core should occur during the rising branch of a mm-wave outburst. The sharp mm-wave outbursts $M1$ and $M2$ (see Fig.~\\ref{mainLC}) seem unrelated to the appearance of new knots although we cannot reject the possibility that new knots were ejected near the time of $M1$ and $M2$ but could not be resolved on the images because they were weaker than knots $K1$ and $K3$, which were ejected before outbursts $M1$ and $M2$, respectively. However, the VLBA observations support the idea that the outbursts are caused by curvature in the inner jet that decreases the angle between a disturbance ($K1$ or $K3$) and the line of sight so that the flux peaks when the disturbance passes through the minimum viewing angle, as suggested by \\citet{VIL07}. Indeed, Figure~\\ref{trajKs} shows curvature in the trajectories of $K1$ and $K3$ at a distance $\\sim$0.05-0.07~mas from the core. This is the location of knots $K1$ and $K3$ when outbursts $M1$ and $M2$ peaked. Comparison of Tables~\\ref{Kparm} and \\ref{Outparm} indicates that $M1$ peaked at a time close to the maximum flux of knot $K1$, and that the maximum of $M2$ occurred when knot $K3$ had the highest flux. The mm-wave spectral index during outbursts $M1$ and $M2$ corresponds to optically thin emission (see \\S 8), as expected for radiation coming from downstream of the core. In addition, Table \\ref{Outparm} shows that the timescale of $M1$ is $\\sim$4-5 times shorter than $w_{\\rm mm}$ of $M2$, consistent with $K1$ moving closer to the jet axis (a smaller radius of curvature) and $K3$ moving along periphery of the jet (a larger radius of curvature), which agrees with the viewing angles of the knots (Table \\ref{Kparm}). We use the delay between the times of ejection of $K1$ and peak of $M1$, along with the kinematic parameters of $K1$, to derive the location of the curvature of the jet, which is 54$\\pm$16~pc downstream the core. \\subsection{Location of the High Energy Emission in the Jet} The analysis of spectral characteristics of outbursts in 3C~454.3 reveals that neither during outbursts nor in quiescent states is the X-ray emission a continuation of the optical synchrotron spectrum. The most likely mechanism of X-ray production is inverse Compton scattering of low-energy photons by relativistic electrons, as suggested by \\citet{SMM08} and \\citet{RAI08b}. Figure~\\ref{OptHE} displays normalized high-energy light curves, with the optical light curve superposed, during outburst $O_{\\rm K1}$, the period covering $O_{\\rm K2}$ plus $O_{\\rm K3}$, and $M2$. Table~\\ref{CCFparm} indicates that the strongest correlation between optical and X-ray variations was during $O_{\\rm K1}$, with a delay $\\lesssim$1~day. Inspection of Figure~\\ref{OptHE} reveals that this correlation was dominated by the first $\\sim$15 days of the decaying branch of the outburst. (Unfortunately, the rising branch was not observed at X-rays and has a seasonal gap at optical frequencies; see Fig.~\\ref{mainLC}). According to the discussion above, the optical outburst took place upstream of the 43~GHz core. The X-ray outburst could have been produced via the EC mechanism, with the seed photons coming from the dusty torus, as argued by \\citet{SMM08}. Similarity in the decay at optical and X-ray frequencies implies that the primary cause of the outburst was an increase in the number of relativistic electrons. However, the EC model predicts that the X-ray spectrum should be rather flat \\citep{GHIS07}, whereas $\\alpha_{\\rm X}$ is steepest during $O_{\\rm K1}$ (Table~\\ref{SpecInd}). Although this is comfortably less than $\\alpha_{\\rm opt}$, it is striking that the X-ray spectrum was flatter during the quiescent state. This can be understood if the X-ray emission during the quiescent state is dominated by the EC process while the SSC mechanism contributes significantly in the X-ray production during $O_{\\rm K1}$. A possible delay of X-ray variations by $\\sim$1~day relative to optical variations (Table~\\ref{CCFparm}) supports such a hypotheses \\citep{McH99}. We have suggested that the optical flares observed within 1$\\sigma$ uncertainty of the time of ejection of $K1$ (Fig.~\\ref{multi}) originated in the VLBI core as the result of compression of the disturbance by standing shocks in the core. Figure~\\ref{OptHE} shows that two X-ray flares (designated as $X1$ and $X2$) occurred during the same time span, which implies that these X-ray flares originated in the mm-wave core as well, probably via the SSC mechanism. The SSC model is supported by tentative delays of several days relative to the optical peaks for flares $X1$ and $X2$, as can be inferred from Figure~\\ref{OptHE}. The higher amplitude of the X-ray flares relative to the synchrotron flares (cf.\\ Fig.~\\ref{OptHE}) is also expected in the SSC case if an increase in the number of relativistic electrons causes the events \\citep{RITABAN08}. The measurements during the $O_{\\rm K2}$+$O_{\\rm K3}$ period (middle panel of Fig.~\\ref{OptHE}) show an increase in X-ray activity during $O_{\\rm K2}$ and possibly $O_{\\rm K3}$. \\citet{VER09} have found a good correlation between $\\gamma$-ray and optical variations at the beginning of $O_{\\rm K3}$, and argue that the $\\gamma$-ray emission is dominated by EC scattering of photons from the broad-line region by relativistic electrons in the jet. Although we agree that EC is probably the dominant mechanism for $\\gamma$-ray production, we argue that the optical outburst $O_{K3}$ appears related to the 43~GHz core located $\\sim$20~pc from the broad-line region. We propose that seed photons for the $\\gamma$-ray production are local to the mm-VLBI core, arising from a slower sheath surrounding the spine, although synchrotron radiation from the disturbance itself plus the mm-wave core should contribute to the production of $\\gamma$-rays as well \\citep{MAR10}. The presence of such a sheath is supported by the wide range of apparent speeds observed in the jet, from 0.12 to 0.53~mas~yr$^{-1}$\\citep{J05}. The prominent mm-wave outburst $M2$ (see Fig.~\\ref{mainLC}) featured moderate $\\gamma$-ray, X-ray, and optical flares. This high-frequency activity was especially pronounced during the early portion of the mm-wave outburst, but tapered off during its later stages. It is possible that a new superluminal knot was ejected during outburst $M2$, but could not be separated from $K3$ on the VLBA images. In this case, the multiple peaks of similar amplitude at 230~GHz, as well as the secondary optical maxima, can be explained by a knot passing through several standing shocks in the core region, as occurred during the passage of $K1$, $K2$, and $K3$ (Fig.~\\ref{multi}). In the absence of a new ejection, the mm-wave outburst is probably not associated with the optical variability, rather it is caused by a change of the viewing angle of $K3$ when the knot was $\\sim$50~pc from the core (see \\S 10.3). The optical emission could originate in the core according to polarization measurements, in a similar manner as proposed for the quasar 0420-014 \\citep{FRANI07}. The $\\gamma$-ray light curve during event $M2$ (bottom panel of Fig.~\\ref{OptHE}) is well correlated with the variations in optical flux, and the X-ray variations correlate with both the optical and $\\gamma$-ray light curves, albeit less strongly. The $\\gamma$-ray and optical variations were simultaneous within a day (Table~\\ref{CCFparm}), while the X-ray variations preceded the optical variations by 3$\\pm$2~days, although the X-ray measurements are sparse. The value of the X-ray spectral index during $M2$ is similar to $\\alpha_{\\rm X}$ during $O_{\\rm K1}$. We suggest that the X-ray emission during $M2$ was produced via the SSC mechanism by the scattering of a wide spectrum of synchrotron photons --- from far-IR to optical frequencies --- by relativistic electrons with Lorentz factors $\\gamma\\sim$10$^{2-3}$. This can explain the moderate correlation between the X-ray and optical variations and smoother X-ray variability. If the delay of the optical with respect to the X-ray variations is real, it can be understood in the manner proposed for the quasar 3C~279 \\citep{RITABAN08}, in which relativistic electrons are accelerated gradually. In this case, the acceleration mechanism should be different from that for $O_{\\rm K2}$ and $O_{\\rm K3}$, and probably involves turbulence in the core. This is consistent with a significantly lower degree of polarization in the 43~GHz relative to optical polarization if the two emission regions are co-spatial but occupy different volumes (Table~\\ref{O7evpa}). Our analysis of the $\\gamma$-ray and optical light curves confirms a strong correlation between the two wavelengths found by \\citet{BON09}. The correlation persists beyond $M2$, with $\\gamma$-ray and optical variations having similar amplitudes of variability (see Fig.~\\ref{OptHE}). \\citet{BON09} have concluded that the $\\gamma$-ray outburst during $M2$ was dominated by EC scattering of IR/optical photons from the accretion disk and/or broad-line region by electrons with $\\gamma\\sim$10$^{3-4}$. Although this may seem a reasonable explanation for the observed strong correlation between $\\gamma$-ray and optical variations, polarization observations (see Fig.~\\ref{Poldat} and Table~\\ref{O7evpa}) indicate that during $M2$ (RJD: 4550-4840) the position angle of the optical polarization tended to align with $\\varphi_{\\rm 43}$ in the core. This implies that the variable optical emission arose in the vicinity of the core, as discussed above. For this reason, we favor the alternative scenario, wherein the $\\gamma$-rays are produced by scattering of synchrotron seed photons from the sheath of the jet. We can be more specific regarding the nature of the seed photons by considering the steepening of the $\\gamma$-ray spectrum above 2~GeV reported by \\citet{Abdo09b}, from a spectral index of 1.3 to 2.5, along with steepening of the spectral index of synchrotron emission from 1.4 at near-IR to 2.3 at soft UV wavelengths (Table \\ref{SpecInd}). We can explain the steepening of the $\\gamma$-ray spectrum, as well as the rapid fall-off in the synchrotron emission toward higher ultraviolet frequencies by a steepening of the electron energy distribution above an energy $\\gamma \\sim 10^4$ in rest-mass units, where we assume a magnetic field strength $\\lesssim 1$ G. Recent theoretical calculations by \\citet{REY09} show that, for an inhomogeneous source, more pronounced steepening than by 0.5 in the source's integrated spectral index is possible due to a combination of synchrotron losses and geometrical effects. If the 2~GeV $\\gamma$-rays are produced by the EC process, the seed photons should have observed frequencies near $10^{12}(\\Gamma/25)^{-2}(\\delta/25)^{-2}$ Hz. This favors either the putative dust torus \\citep{SMM08} or sheath of the jet as the source of the seed photons. The similarity of the $\\gamma$-ray and optical light curves, including very rapid variations (see \\S 7), can then be explained as a consequence of the variations being caused by changes in the number of electrons with $\\gamma \\sim 3\\times 10^3$-$10^4$ over a limited volume within a given disturbance." }, "1003/1003.1680_arXiv.txt": { "abstract": "I consider the question of possible observability of the total number of $e$-folds accumulated during the epoch of inflation. The total number of observable $e$-folds has been previously constrained by the de Sitter entropy after inflation, assuming that the null energy condition (NEC) holds. The NEC is violated by upward fluctuations of the local Hubble rate $H$, which occur with high probability in the fluctuation-dominated regime of inflation. These fluctuations lead at late times to the formation of black holes and thus limit the observability of inflationary evolution. I compute the average number $\\left\\langle \\Delta N\\right\\rangle $ of $e$-folds accumulated during the last NEC-preserving fragment of the inflationary trajectory before reheating. This is the maximum number of inflationary $e$-folds that can be observed in principle through measurements of the CMB at arbitrarily late times (if the dark energy disappears). The calculation also provides a reasonably precise definition of the boundary of the fluctuation-dominated regime, with an uncertainty of a few percent. In simple models of single-field inflation compatible with current CMB observations, I find $\\left\\langle \\Delta N\\right\\rangle $ of order $10^{5}$. This upper bound on the observable $e$-folds, although model-dependent, is much smaller than the de Sitter entropy after inflation. The method of calculation can be used in other models of single-field inflation. ", "introduction": "Inflation produces primordial metric fluctuations that may be observed indirectly through CMB measurements such as WMAP~\\cite{Bennett:2003bz}. An observation of CMB at present corresponds to the measurement of the inflaton evolution about 60 $e$-folds before reheating~\\cite{Liddle:2003as}. Assuming that CMB measurements will be possible at indefinitely late times, one might hope to deduce information about arbitrarily early stages of inflation. Of course, late-time acceleration (persistent {}``dark energy'') can make it impossible to observe CMB at very late times~\\cite{Krauss:2007nt}. There is, however, another limit on our ability to see towards the past. This limit is caused by violations of the null energy condition (NEC) during inflation. To make the following arguments more specific, let us consider a model of inflation driven by a canonical, minimally coupled scalar field $\\phi$ such that the field evolves from a large initial value $\\phi_{\\text{in}}$ (perhaps near the Planck boundary $\\phi_{\\text{Pl}}$) to the reheating point $\\phi=\\phi_{*}$. A typical model of this type has the inflaton action\\begin{equation} \\int d^{4}x\\sqrt{-g}\\left(\\frac{1}{2}\\phi_{,\\mu}\\phi^{,\\mu}-V(\\phi)\\right).\\label{eq:inflaton action}\\end{equation} We assume that the inflaton potential $V(\\phi)$ grows monotonically with $\\phi$ and that the slow-roll approximation is valid. In models of chaotic type, e.g.~$V(\\phi)\\propto\\phi^{2n}$, we then expect that $\\phi_{*}\\ll\\phi_{\\text{in}}\\lesssim\\phi_{\\text{Pl}}$, and that $\\phi_{\\text{in}}$ is deep in the fluctuation-dominated regime. The evolution of $\\phi$ during inflation can be pictured as a random walk superimposed on a deterministic drift towards $\\phi=\\phi_{*}$ \\cite{Vilenkin:1983xq,Starobinsky:1986fx,Linde:1986fd}. The random walk can be modeled as {}``diffusion'' in $\\phi$ space with the diffusion coefficient \\begin{equation} D(\\phi)\\equiv\\frac{H^{3}}{8\\pi^{2}},\\quad H(\\phi)\\equiv\\frac{8\\pi G}{3}V(\\phi)\\equiv\\frac{8\\pi}{3M_{\\text{Pl}}^{2}}V(\\phi),\\end{equation} while the mean drift velocity is the time derivative $\\dot{\\phi}$ of the slow-roll evolution,\\begin{equation} \\dot{\\phi}=v(\\phi)\\equiv-\\frac{H^{\\prime}}{4\\pi G}=-\\frac{H^{\\prime}}{4\\pi}M_{\\text{Pl}}^{2}.\\label{eq:phi dot sr}\\end{equation} During the last stages of inflation before reheating, the trajectory $\\phi(t)$ is monotonic ($\\dot{\\phi}<0$) with nearly unit probability, although there is always a small probability of an upward fluctuation ($\\dot{\\phi}>0$). On the other hand, in the fluctuation-dominated regime an upward fluctuation of $\\phi$ has around 50\\% probability because random fluctuations dominate over the slow-roll motion. We now note that an upward fluctuation of $\\phi$ corresponds to an upward fluctuation of the local Hubble rate $H$ and thus to a local violation of the null energy condition (NEC) \\cite{Winitzki:2001fc,Vachaspati:2003de}. A violation of the NEC due to an upward fluctuation of $H$ on a distance scale $L\\sim H^{-1}$ leads to the formation of a Hubble volume that looks like a black hole from the outside~\\cite{Blau:1986cw}, and to a real black hole when the overdensity enters the local Hubble horizon at late times after inflation~\\cite{Linde:1988zp,Bousso:2006ge}. This can be understood qualitatively by noting that a Hubble-size region of approximately de Sitter spacetime with local Hubble parameter $H$ has exactly the energy density $\\Lambda=\\frac{3}{8\\pi}M_{\\text{Pl}}H^{2}$ that corresponds to a black hole with the Schwarzschild radius $H^{-1}$. An upward fluctuation of $H$ therefore leads to an increase of the energy density beyond the Schwarzschild limit. In this way, an NEC violation in the far inflationary past will limit the lifetime of any future observers who may be trying to perform CMB observations at very late times. The main focus of this paper is an investigation of this limiting effect of NEC violations. For the sake of this consideration, I will assume that dark energy eventually decays, so that the late-time universe is not expanding with acceleration and the local Hubble radius grows without limit, permitting (in principle) observations of the primordial density fluctuations on arbitrarily large scales. It is interesting to determine the time range within which the NEC can be violated during inflation. For each random inflationary trajectory $\\phi(t)$ there exists a well-defined time of the \\emph{last} NEC violation before reheating, i.e.~a time $t_{\\text{N}}$ such that the NEC is violated around $t=t_{\\text{N}}$ but then is not violated any more. The evolution of $\\phi(t)$ before $t=t_{\\text{N}}$ is thus not observable even in principle. On the other hand, the evolution of $\\phi$ after $t=t_{\\text{N}}$ is in principle observable: If primordial fluctuations on a distance scale $L$ are produced at $t>t_{\\text{N}}$, these fluctuations will be observed through the CMB fluctuations at sufficiently late times when the scale $L$ reenters the Hubble horizon (we are assuming that the dark energy does not prevent such observations). Therefore, it is only the statistics of the last NEC-preserving% \\footnote{I talk about {}``NEC-preserving'' rather than about {}``monotonically decreasing'' trajectories $\\phi(t)$ because in models with several fields ($\\phi_{1},...,\\phi_{n}$), an NEC violation does not necessarily entail an upward fluctuation of a particular field $\\phi_{k}(t)$. % } segment of the inflaton trajectory $\\phi(t)$ that is --- even in principle --- accessible to observations. In the following sections we will compute (within an adequate approximation) the mean duration of the last NEC-preserving portion of the trajectory $\\phi(t)$ until reheating. When fluctuations are negligible, the field evolves according to the slow-roll equation~(\\ref{eq:phi dot sr}), and so the time needed for evolving from $\\phi=\\phi_{1}$ to reheating (taking into account that $\\phi_{*}<\\phi_{1}$) is\\begin{equation} \\Delta t(\\phi_{1},\\phi_{*})=\\int_{\\phi_{*}}^{\\phi_{1}}\\frac{d\\phi}{-v(\\phi)}.\\end{equation} (We write $-v$ because the value of $\\phi$ decreases with time, so $v(\\phi)<0$.) This formula, however, cannot be used directly to compute the duration of the last NEC-preserving portion of the trajectory, for two reasons: First, the NEC-preserving portion of the trajectory depends on chance and is not confined within a fixed interval, say $\\left[\\phi_{*},\\phi_{1}\\right]$. Second, the initial stages of the trajectory $\\phi(t)$ belong to the fluctuation-dominated regime where the evolution $\\phi(t)$ is not well described by the deterministic slow-roll equation $\\dot{\\phi}=v(\\phi)$. Below we will compute the \\emph{average} duration $\\left\\langle \\Delta t_{\\text{NEC}}\\right\\rangle $ of the last NEC-preserving portion of the trajectory $\\phi(t)$, where the average is performed over the ensemble of all comoving trajectories.% \\footnote{Thus we compute the {}``comoving'' average rather than a {}``volume-weighted'' average, which would require more complicated calculations left for future work. See, e.g., Ref.~\\cite{Winitzki:2006rn} for a review of comoving and volume-weighted averaging prescriptions. % } Heuristically, we may attempt to determine a value $\\phi=\\phi_{q}$ such that the duration of the slow-roll trajectory between $\\phi=\\phi_{q}$ and $\\phi=\\phi_{*}$ is precisely equal to $\\left\\langle \\Delta t_{\\text{NEC}}\\right\\rangle $, i.e.~we first compute $\\left\\langle \\Delta t_{\\text{NEC}}\\right\\rangle $ and then \\emph{define} $\\phi_{q}$ such that\\begin{equation} \\Delta t(\\phi_{q},\\phi_{*})=\\int_{\\phi_{*}}^{\\phi_{q}}\\frac{d\\phi}{-v(\\phi)}=\\left\\langle \\Delta t_{\\text{NEC}}\\right\\rangle .\\end{equation} This value $\\phi_{q}$ can then be interpreted as the boundary between the fluctuation-dominated and the fluctuation-free regimes. One of the main results of this paper is a method of computing $\\left\\langle \\Delta t_{\\text{NEC}}\\right\\rangle $; thus, the boundary $\\phi_{q}$ between the fluctuation-dominated and the fluctuation-free regimes becomes a well-defined quantity. A precise definition of this boundary is relevant, e.g., for certain measure prescriptions for regulating eternal inflation~\\cite{Linde:2007nm,Linde:2008xf} as well as for attempts to count the observable degrees of freedom after inflation~\\cite{ArkaniHamed:2007ky,Linde:2009ah}. The boundary between the fluctuation-dominated and the fluctuation-free regimes can be characterized in terms of the dimensionless ratio of $-v\\delta t$ (the change of the field $\\phi$ due to slow roll during one Hubble timestep $\\delta t\\equiv H^{-1}$) to $\\sqrt{2D\\delta t}$ (the typical fluctuation during the same time),\\begin{equation} b(\\phi)\\equiv\\frac{-v\\delta t}{\\sqrt{2D\\delta t}}=\\frac{-2\\pi\\dot{\\phi}}{H^{2}}=\\sqrt{\\frac{3}{8}\\varepsilon_{1}\\frac{M_{\\text{Pl}}^{4}}{V(\\phi)}},\\label{eq:b def 0}\\end{equation} where \\begin{equation} \\varepsilon_{1}\\equiv\\frac{M_{\\text{Pl}}^{2}}{16\\pi}\\frac{V^{\\prime2}}{V^{2}}\\label{eq:epsilon1 def}\\end{equation} is the first slow-roll parameter. We note that $b^{2}(\\phi)$ coincides, for inflationary models of the type~(\\ref{eq:inflaton action}), with the inverse magnitude of the power spectrum of the primordial scalar fluctuation mode that crossed the Hubble scale at that time,\\begin{equation} P_{S}\\approx\\frac{1}{4\\pi^{2}}\\left(\\frac{H^{2}}{\\dot{\\phi}}\\right)^{2}=\\frac{8}{3\\varepsilon_{1}}\\frac{V(\\phi)}{M_{\\text{Pl}}^{4}}=\\frac{1}{b^{2}(\\phi)}.\\label{eq:PS def}\\end{equation} (In the equation above, we neglected the slow-roll corrections since our final result will only depend logarithmically on $P_{S}$.) The fluctuation-free regime is characterized by $b(\\phi)\\gg1$ and the fluctuation-dominated regime by $b(\\phi)\\lesssim1$ (i.e.~by primordial fluctuations of order 1 or larger). However, this qualitative characterization cannot provide a sharply defined boundary value $\\phi_{q}$ separating the two regimes. We also note that the order parameter $\\Omega$, which was used in Ref.~\\cite{Creminelli:2008es} to characterize the transition between the presence and the absence of eternal inflation, is related to $b$ by\\begin{equation} \\Omega\\equiv\\frac{2\\pi^{2}}{3}\\frac{\\dot{\\phi}^{2}}{H^{4}}=\\frac{\\pi}{3}b^{2}.\\end{equation} The fluctuation-dominated regime was characterized by the condition $\\Omega<1$ in Ref.~\\cite{Creminelli:2008es}, which again corresponds qualitatively to $b\\lesssim1$. It will be shown below that the rigorously computed value $\\left\\langle \\Delta t_{\\text{NEC}}\\right\\rangle $ can be approximated as \\begin{equation} \\left\\langle \\Delta t_{\\text{NEC}}\\right\\rangle \\approx\\int_{\\phi_{*}}^{\\phi_{\\text{Pl}}}\\frac{d\\phi}{-v(\\phi)}f(0;\\phi),\\end{equation} where $f(0;\\phi)$ is the probability of the event that an inflationary trajectory $\\phi(t)$ starting at $t=0$ with the given value of $\\phi$ will \\emph{never} violate the NEC. It will be found that $f$ is close to being a step function, $f(0;\\phi)\\approx\\theta(\\phi_{q}-\\phi)$, effectively cutting the integration at $\\phi=\\phi_{q}$. The relevant value of $\\phi_{q}$ will be determined from an explicit analytic approximation for $f(0;\\phi)$. If we use the number of $e$-foldings, $N\\equiv\\ln a$, as the time variable $t$, the same method yields the average number of $e$-foldings, $\\left\\langle \\Delta N_{\\text{NEC}}\\right\\rangle $, during the last NEC-preserving portion of the trajectory before reheating. Below we will compute $\\phi_{q}$ and $\\left\\langle \\Delta N_{\\text{NEC}}\\right\\rangle $ explicitly for inflationary models with a power-law potential. In a specific example, we will show that the average number of NEC-preserving $e$-foldings in the model with $V(\\phi)=\\frac{1}{2}m^{2}\\phi^{2}$ is of order $10^{5}$, if one assumes the model parameters that fit the WMAP data. In this model, the value of $\\phi_{q}$ turns out to be such that $b^{2}(\\phi_{q})\\approx14$. Thus, $\\phi_{q}$ is well outside the fluctuation-dominated regime. These results can be compared with the upper bound on the $e$-folds of inflation obtained in Ref.~\\cite{ArkaniHamed:2007ky}. Assuming that the fluctuations are never dominant (equivalently, that the NEC always holds), it was found that the number of observable $e$-folds of inflation must be smaller than the entropy $S_{dS}$ of the final de Sitter state \\emph{after} inflation. The latter is an extremely large number, of order $10^{120}$ (if we use the current value of the dark energy density). The present calculation shows that the limit on the number of observable $e$-folds is much more stringent. Therefore, the number of observable degrees of freedom, if it is expressed through the entropy of the final de Sitter state, is in any case not directly related to the total number of observable $e$-folds of inflation. It must be noted that the value $\\left\\langle \\Delta N_{\\text{NEC}}\\right\\rangle $ is highly model-dependent. In the present paper, we perform computations only for models with $V(\\phi)\\propto\\phi^{2n}$ and derive a formula for $\\left\\langle \\Delta N_{\\text{NEC}}\\right\\rangle $ {[}Eq.~(\\ref{eq:DeltaN ans 01})] that shows a sensitive dependence on $n$. However, the method of calculation developed in this paper is sufficiently general so that the number of total observable $e$-folds can be computed in any other model of single-field slow-roll inflation. ", "conclusions": "" }, "1003/1003.3961_arXiv.txt": { "abstract": "Between July 5th and July 7th 2004, two intriguing fast coronal mass ejection(CME)-streamer interaction events were recorded by the Large Angle and Spectrometric Coronagraph (LASCO). At the beginning of the events, the streamer was pushed aside from their equilibrium position upon the impact of the rapidly outgoing and expanding ejecta; then, the streamer structure, mainly the bright streamer belt, exhibited elegant large scale sinusoidal wavelike motions. The motions were apparently driven by the restoring magnetic forces resulting from the CME impingement, suggestive of magnetohydrodynamic kink mode propagating outwards along the plasma sheet of the streamer. The mode is supported collectively by the streamer-plasma sheet structure and is therefore named `` streamer wave'' in the present study. With the white light coronagraph data, we show that the streamer wave has a period of about 1 hour, a wavelength varying from 2 to 4 solar radii, an amplitude of about a few tens of solar radii, and a propagating phase speed in the range 300 to 500 km s$^{-1}$. We also find that there is a tendancy for the phase speed to decline with increasing heliocentric distance. These observations provide good examples of large scale wave phenomena carried by coronal structures, and have significance in developing seismological techniques for diagnosing plasma and magnetic parameters in the outer corona. ", "introduction": "Wave phenomena represent the most fundamental and straightforward response of a system with plasmas and magnetic fields to perturbations arising from either interior or exterior. The solar atmosphere, serving as a good example, is very dynamic { by nature on} all relevant temporal-spatial scales, {and} is therefore expected to be able to support various wave modes with different observational manifestations. Indeed, with the development of observational techniques, many types of wave or wavelike phenomena have been discovered in the solar atmosphere. For instance, compressible density perturbations moving outwards are detected inside coronal plumes (Ofman et al., 1997, 1999; DeForest {\\&} Gurman, 1998), propagating longitudinal waves are found in coronal loops (Berghmans {\\&} Clette, 1999), and many other phenomena driven by nearby solar eruptions, including coronal loop oscillations (Aschwanden et al., 1999; Nakariakov et al., 1999), coronal shocks (Sime {\\&} Hundhausen, 1987; Sheeley et al., 2000), and the so-called Moreton (Moreton {\\&} Ramsey, 1960) and EIT waves (Thompson et al., 1998; Wills-Davey {\\&} Thompson, 1999), are observed. Extensive observational and theoretical studies have been conducted to investigate the nature of these dynamical phenomena (e.g., Aschwanden, 2004; Nakariacov {\\&} Verwichte, 2005; Ofman, 2009; and references therein). These studies, generally speaking, provide valuable information on the coronal medium through which the waves propagate. Helmet streamers are the most conspicuous large-scale quasi-steady structures extending from the lower to outer corona. In the white light images observed by a coronagraph, a well developed streamer is delineated by a sharp brightness boundary. The boundary separates the streamer from its surroundings. Besides the boundary, a typical streamer also includes a bunch of closed field arcades, a streamer cusp, and a high density plasma sheet (also called the streamer stalk or streamer belt) within which a long thin current sheet {is} embedded (see, e.g, Pneuman {\\&} Kopp, 1971; Suess {\\&} Nerney, 2006). On the other hand, coronal mass ejections (CMEs), representing the largest and most energetic dynamical process in the corona, may cause global perturbations with a timescale of minutes to hours. Therefore, close interactions between CMEs and streamers can frequently occur, especially during the active phase of a solar cycle when CMEs and streamers are present at virtually all heliolatitudes. In general, CME-streamer relevant events can be classified into two groups. One comprises those events of CMEs originating and erupting from the streamer interior, like the so-called streamer blowouts (Howard et al., 1985; Hundhausen 1993) or streamer puffs (Bemporad et al. 2005). On the other hand, the events in the second group result from the streamers being hit on the sides by either CMEs with expanding structures or by CME-driven disturbances like shock waves. The collision may cause apparent deflections or kinks of streamer rays tracing the passage of CME disturbances (Sheeley et al. 2000). In some cases, the collision may have triggered reconnections across the streamer current sheet as indicated by the observed streamer disconnection (e.g., Bemporad, et al., 2008), the release of plasma blobs along the streamer stalk, or the formation of streamer in/out pairs (e.g., Sheeley {\\&} Wang, 2007). Given the fundamental role played by wave excitations in a disturbed plasma-magnetic field system, one natural question arises, can the streamer respond in the form of observable waves or wavelike motions to a strong impact from a CME ejecta? If yes, what modes are they? In the following text, an answer to the above questions will be provided with two observational examples of streamer wavy motions driven by CMEs. Their overall details as revealed from the white-light coronagraph data will be described in Section 2. In the 3rd section we present our data manipulation method to extract the profile of the wavy motion, and give the resultant physical analysis on one of the two events in the 4th section. In Section 5, we discuss briefly on the CME-streamer sources and the other observational event. The final section presents our conclusions and discussion. ", "conclusions": "In this paper, we conduct an observational study on the phenomena of streamer wave, which is excited by the CME impact and represents one of the largest wave phenomena ever discovered in the corona. The wave is mostly MHD kink mode propagating outwards along the thin plasma sheet. The restoring force supporting the wavy motion is provided by the magnetic field of the streamer structure, which is generated by the large streamer deflection upon the CME impact. The energy received from the impact is carried outwards by the wave perturbation. Consequently, the amplitude of the wave near the sun declines rapidly with time, and only a few periods of the wave are observable. The wave period is estimated to be about 1 hour, the wavelength varies from 2 to 4 R$_\\odot$, the wave amplitude is a few tens of solar radii, and the phase speed is about 300 to 500 km s$^{-1}$. There exists a general trend for the phase speed to decrease with increasing heliocentric distance. Interactions between CME and streamers are frequently observed, especially during the active phase of solar cycles. Usually, such interactions result in apparent deflections of interacting streamers (e.g., Hundhausen et al., 1987; Sime {\\&} Hundhausen, 1987; Sheeley et al., 2000). We emphasize that the streamer wavy motion, reported in the present study, is a direct consequence of a streamer deflection. Nevertheless, as revealed from a preliminary overview of the long-term LASCO observations, in only a very small fraction of the deflection events the streamer exhibits wavelike phenomena. In other words, most CME-driven deflections, even very fast and strong, are not followed by a streamer wavy motion. Therefore, there exist certain strict conditions for streamer waves to be excited by a CME-streamer deflection. Two observational features of the July 6 event can help us evaluate the relevant conditions. Firstly, it is found that the CME source region lies on the flank side of the closed loops comprising the streamer, that means the CME does not originate from beneath the streamer structure, and the ejecta can collide with the streamer from the flank side. Secondly, the CME is a fast eruption with a speed of $\\ge 1300$~km~s$^{-1}$, which has two consequences favoring the excitation of the streamer wave. One is that a faster eruption results in a stronger impingement on the nearby streamer and a consequent larger deflection of the streamer structure from its equilibrium position, the other is that the ejecta moves out of the corona in a relatively short time, and leaves enough time for the streamer wave to develop. Otherwise if the eruption is not fast enough, the deflected streamer may simply moves backwards along with the ejecta, and no wavy motions result. To observe one example of such a case, one may check the online LASCO observations of the interaction event between a CME and a streamer in the northeastern quadrant dated on July 9th, 2004. Sheeley et al. (2000) also presents LASCO examples of strong streamer deflection events without accompanying apparent streamer wavy motions. It should be noted that a more complete understanding of the excitation conditions of the streamer wave can only be obtained from observational investigations on much more similar events and from elaborate theoretical modelling endeavors. As mentioned in the introduction section, a well developed typical streamer consists of the main body, which is a bunch of closed field arcades confining high density coronal plasmas, and a dense plasma sheet within which a long thin current sheet is embedded. The intersection of the closed streamer main body and the open plasma sheet gives the streamer cusp, which is generally thought to be below 2 to 2.5 $R_\\odot$, very close to the bottom of the LASCO C2 FOV. After the impact from a CME, the streamer deflects away from its original equilibrium position. The consequent restoring motion may excite the wavelike oscillations propagating along the plasma sheet. Therefore, the geometry supporting the discussed streamer wave motion can be simplified as a long slender plasma slab extending to infinity with the lower end attaching to the streamer cusp which bounces back and forth in a quasi-periodic manner. The oscillations are observed to be generally transverse to the nominal direction of the magnetic field. The manifestation and the geometry of the phenomena are very similar to that of the well-known kink mode deduced from a slender magnetic slab except being in a spherical expanding geometry (Roberts, 1981; Edwin {\\&} Roberts, 1982). It is therefore suggested that the wave phenomenon discussed in this study represents the kink mode, which is, in a more general sense, a type of fast magnetosonic waves propagating in an inhomogeneous magnetized plasma environment. It is interesting to notice that the morphology of the streamer wave discussed above is very similar to a traditional Chinese dance named as 'Colored Belt Dance' which is performed by dancers holding one end of a long belt in color. An important extension to the coronal wave study is to develop diagnostic techniques of plasmas and magnetic fields through which the wave propagates, i.e., to conduct the study of coronal seismology. In our case, the period and phase speed of the streamer wave which has been regarded as the propagating kink mode carried by the thin plasma sheet, if well resolved from observations, can be used to provide information on magnetic properties of streamers. Generally speaking, the phase speed for the wave phenomenon investigated in this study is given by the sum of two components. The first one is the speed of the solar wind along the plasma sheet, the medium carrying the mode outwards. The other is of course the phase speed of the wave mode in the plasma rest frame. The phase speed for the kink mode under thin plasma sheet geometry can be tentatively described with available MHD theory developed for a plasma-slab configuration in cartesian geometry (Roberts, 1981; Edwin {\\&} Roberts, 1982). Substituting nominal parameters in the slow-wind plasma sheet region above the streamer cusp into the dispersion relation given by Edwin {\\&} Roberts (1982), we find that the phase speed of the relevant fast kink body mode $c_k$ is smaller than yet rather close to the external Alfv\\'en speed $v_{Ae}=B_e / \\sqrt{\\mu_0 n m_p}$, where $n$ is the proton number density and $m_p$ the proton mass. The difference between the deduced $c_k$ and $v_{Ae}$ is generally less than one third of $v_{Ae}$. Therefore, to implement a preliminary seismological study on the magnetic field strength $B_e$, we take $v_{Ae}$ to be equal to the kink mode phase speed $c_k$ estimated from our observations. Regarding the solar wind conditions in the concerned region, the readers are referred to relevant observational studies (Sheeley et al., 1997; Wang et al., 2000; Strachan et al., 2002; Song et al., 2009) and theoretical modelings (e.g., Wang et al., 1998; Suess et al., 1999; Chen et al., 2001, 2002; Hu et al., 2003; Li et al., 2006). In this short discussion, we simply make use of the solar wind conditions obtained by Chen {\\&} Hu (2001). Only two distances are considered, (1) at 5 $R_\\odot$, the solar wind velocity $v_{sw} = 100$ km s$^{-1}$ and $n = 1\\times10^5$ cm$^{-3}$, and (2) at 10 $R_\\odot$, $v_{sw} = 200$ km s$^{-1}$ and $n = 2\\times10 ^4$ cm$^{-3}$. With these assumptions, it is straightforward to deduce $c_k$ and thus $v_{Ae}$ at the plasma rest frame, and then calculate the value of the magnetic field strength $B_e$ at the above two distances in the region surrounding the plasma sheet. Here we only present our calculations of $B_e$ with the measurements associated with the second phase point P2, whose speeds are 410 km s$^{-1}$ at 5$R_\\odot$ and 360 km s$^{-1}$ at 10 $R_\\odot$, as read from Figure 5. It is found that the magnetic field strength declines from 0.045 G at 5 $R_\\odot$ to 0.01 G at 10 $R_\\odot$, indicating a slightly super-radial expansion of the magnetic flux tube from 5 to 10 $R_\\odot$. These values are consistent with the results given by recent corona and solar wind models (e.g., Li et al., 2006). A more complete seismological study, together with sophisticated numerical MHD simulations of CME-streamer interactions to shed more light on the excitation and propagation of the waves, should be conducted in future." }, "1003/1003.1225_arXiv.txt": { "abstract": "The recent results from the PAMELA, ATIC, FERMI and HESS experiments have focused attention on the possible existence of high energy cosmic ray $e^+ e^-$ that may originate from dark matter (DM) annihilations or decays in the Milky Way. Here we examine the morphology of the $\\gamma$-ray emission after propagation of the electrons generated by both annihilating and decaying dark matter models. We focus on photon energies of 1 GeV, 10 GeV, 50 GeV (relevant for the FERMI satellite) and consider different propagation parameters. Our main conclusion is that distinguishing annihilating from decaying dark matter may only be possible if the propagation parameters correspond to the most optimistic diffusion models. In addition, we point to examples where morphology can lead to an erroneous interpretation of the source injection energy. ", "introduction": "Results from recent cosmic ray experiments (PAMELA \\cite{pamela}, ATIC\\cite{atic}, FERMI \\cite{fermi}, HESS\\cite{hess}) have raised the question of the origin of an ``anomalous'' population of high energy positrons in the Milky Way and motivated many studies. Correlation of the positron flux measured by PAMELA with the $\\gamma$-ray spectrum obtained by FERMI LAT is expected to give insight into the injection energy of the high energy electron and positron ($e^+, e^-$) population, and should also probe their spatial and energy distribution. In scenarios where high energy $e^+, e^-$ are emitted by dark matter (DM), the spatial and energy distribution of this ``additional'' cosmic ray population is expected to follow the DM halo distribution at the injection energy $E=E_{inj}$. This implies (assuming a spherical DM halo) that they should be spherically distributed with an energy density that is maximal near the Galactic Centre. However, this picture could be modified if the high energy $e^+, e^-$ spatially propagate and lose energy in the galaxy owing to inverse Compton and synchrotron losses. As a consequence of propagation, not only will the spatial and energy distributions of the high energy $e^+, e^-$ be modified but their final energy will be smaller than $E_{inj}$. The $\\gamma$-ray spectrum obtained after propagation could therefore differ significantly from that obtained at injection. The issue of the $\\gamma$-ray flux associated with DM annihilations or decays into leptons has been addressed in several papers. For example, both the $\\gamma$-ray flux and $\\gamma$-ray spectrum in decaying and annihilating scenarios have been discussed in ref.~\\cite{cirelli1,cirelli2} but propagation was actually neglected. More recently, the authors of ref.\\cite{sigl} have predicted the expected $\\gamma$-ray flux in a decaying DM model, taking into account $e^+, e^-$ propagation. Although computation of the flux is important, exploiting its value will be difficult owing to large uncertainties due to astrophysical sources at these energies (Ref.\\cite{stefano}). Other papers have considered specific positions on the sky (e.g. intermediate galactic latitudes, \\cite{bertoneGC,ullio2}), or rely on very large-scale anisotropies (\\cite{sigl}). The work in ref. \\cite{ullio1} raised the question of the morphology of the $\\gamma$-ray emission but mainly focused on the spectrum; however the propagation parameters adopted are not those favoured by MCMC studies (\\cite{antje}). The question we raise in this Letter is whether the morphology of the $\\gamma$-ray emission alone (rather than the flux) can actually help to discriminate between the different DM scenarios. To address this issue, we compute $\\gamma$-ray maps originating from the interactions of $e^+ e^-$ with the Interstellar Radiation Field (ISRF) spectra after propagation. We assume that the dark matter only annihilates or decays into $e^+ e^-$ pairs and focus on $\\gamma$-ray energies that are accessible by the FERMI satellite, namely $E_{\\gamma}= 1,10,50$ GeV. {We have neglected prompt $\\gamma$ emission which could arise from internal bremsstrahlung because its spectrum is model-dependent, however one should keep in mind that for some models, this emission could modify our conclusions.} Given our assumptions, the injection energy of the $e^+$ and $e^-$ corresponds to either the DM mass $m_{dm}$ or half the DM mass ($E_{inj} = m_{dm}$ or $m_{dm}/2$), depending on whether DM is annihilating or decaying respectively. We will consider three values of the injection energy: $E_{inj} = 100, 500, 1000$ GeV. At given $E_{inj}$, the comparison between decaying and annihilating scenarios is immediate. In addition, since there are quite large uncertainties in the propagation parameters of cosmic rays, we will use three different sets of parameters referred to as (MIN,MED,MAX) (cf Ref.~\\cite{2001ApJ...555..585M,2001ApJ...563..172D}), which give a fair idea of the related uncertainty. We present difference maps of the $\\gamma$-ray contributions that highlight how morphology could help discriminate between the competing models. ", "conclusions": "} We have generated maps of $\\gamma$-ray emission associated with $e^+, e^-$ population originating from DM annihilations or decays. We show that propagation is important for both DM scenarios, but although the propagation features differ, they are difficult to distinguish if the propagation parameters correspond to MIN (and perhaps MED) rather than to MAX. This is, in fact, surprising, as one might have expected these two scenarios, which involve distinct powers of the dark matter density, to differ significantly. Actually, in the MIN case, detection would be extremely challenging since most of the signal would be hidden by galactic sources. In some cases, the IRSF can make the Galactic Centre bright enough to be misinterpreted as $e^+ e^-$ with a lower injection energy.\\\\ We have verified that changing the energy density of the ISRF has little effect as the increase of the $\\gamma$-ray emissivity is partially compensated by the electron density decrease due to increased energy losses. Varying the intensity of the magnetic field within reasonable values has also little impact as synchrotron emission is not the main energy loss term in most cases. In both cases the impact is mainly on the intensity and not on the ellipticity. However a full spatial description of both the ISRF and the magnetic field could have effects that are beyond the scope of our analytical approach. \\vspace{0.2cm} We would like to thank P. Salati and A. Fiasson for very useful discussions. and acknowledge fundings from the \"low energy electron and positron propagation\" PICS." }, "1003/1003.0962_arXiv.txt": { "abstract": "In this article I review model--independent procedures for extracting properties of Weakly Interacting Massive Particles (WIMPs) from direct Dark Matter detection experiments. Neither prior knowledge about the velocity distribution function of halo Dark Matter particles nor about their mass or cross sections on target nucleus is needed. The unique required information is measured recoil energies from experiments with different detector materials. ", "introduction": "Different astronomical observations and measurements indicate that more than 80\\% of all matter in our Universe are ``dark'' and this Dark Matter interacts at most very weakly with ordinary matter. Weakly Interacting Massive Particles (WIMPs) $\\chi$ arising in several extensions of the Standard Model of particle physics with masses roughly between 10 GeV and a few TeV are one of the leading candidates for Dark Matter% \\cite{SUSYDM96,Bertone05,Steffen08,Bergstrom09}. Currently, the most promising method to detect different WIMP candidates is the direct detection of the recoil energy deposited by elastic scattering of ambient WIMPs off the target nuclei% \\cite{Smith90,Lewin96}. The differential event rate for elastic WIMP--nucleus scattering is given by% \\cite{SUSYDM96}: \\beq \\dRdQ = % \\afrac{\\rho_0 \\sigma_0}{2 \\mchi \\mrN^2} \\FQ \\int_{\\vmin}^{\\vmax} \\bfrac{f_1(v)}{v} dv \\~. \\label{eqn:dRdQ} \\eeq Here $R$ is the direct detection event rate, i.e., the number of events per unit time and unit mass of detector material, $Q$ is the energy deposited in the detector, $\\rho_0$ is the WIMP density near the Earth, $\\sigma_0$ is the total cross section ignoring the form factor suppression, $F(Q)$ is the elastic nuclear form factor, $f_1(v)$ is the one--dimensional velocity distribution function of the WIMPs impinging on the detector, $v$ is the absolute value of the WIMP velocity in the laboratory frame. The reduced mass $\\mrN$ is defined by \\beq \\mrN \\equiv \\frac{\\mchi \\mN}{\\mchi + \\mN} \\~, \\label{eqn:mrN} \\eeq where $\\mchi$ is the WIMP mass and $\\mN$ that of the target nucleus. Finally, \\mbox{$\\vmin = \\alpha \\sqrt{Q}$} is the minimal incoming velocity of incident WIMPs that can deposit the energy $Q$ in the detector with the transformation constant \\beq \\alpha \\equiv \\sfrac{\\mN}{2 \\mrN^2} \\~, \\label{eqn:alpha} \\eeq and $\\vmax$ is the maximal WIMP velocity in the Earth's reference frame, which is related to the escape velocity from our Galaxy at the position of the Solar system, $\\vesc~\\gsim~600$ km/s. The total WIMP--nucleus cross section $\\sigma_0$ in Eq.~(\\ref{eqn:dRdQ}) depends on the nature of WIMP couplings on nucleons. Through e.g., squark and Higgs exchanges with quarks, WIMPs could have a ``scalar'' interaction with nuclei. % The total cross section for the spin--independent (SI) scalar interaction can be expressed as\\cite{SUSYDM96,Bertone05} \\beq \\sigmaSI = \\afrac{4}{\\pi} \\mrN^2 \\bBig{Z f_{\\rm p} + (A - Z) f_{\\rm n}}^2 \\~. \\label{eqn:sigma0_scalar} \\eeq Here $\\mrN$ is the reduced mass defined in Eq.~(\\ref{eqn:mrN}), $Z$ is the atomic number of the target nucleus, i.e., the number of protons, $A$ is the atomic mass number, $A-Z$ is then the number of neutrons, $f_{\\rm (p, n)}$ are the effective scalar couplings of WIMPs on protons p and on neutrons n, respectively. Here we have to sum over the couplings on each nucleon before squaring because the wavelength associated with the momentum transfer is comparable to or larger than the size of the nucleus, the so--called ``coherence effect''. In addition, for the lightest supersymmetric neutralino, and for all WIMPs which interact primarily through Higgs exchange, the scalar couplings are approximately the same on protons and on neutrons: \\( f_{\\rm n} \\simeq f_{\\rm p} . \\) Thus the ``pointlike'' cross section $\\sigmaSI$ in Eq.~(\\ref{eqn:sigma0_scalar}) can be written as \\beq \\sigmaSI \\simeq \\afrac{4}{\\pi} \\mrN^2 A^2 |f_{\\rm p}|^2 = A^2 \\afrac{\\mrN}{\\mrp}^2 \\sigmapSI \\~, \\label{eqn:sigma0SI} \\eeq where $\\mrp$ is the reduced mass of the WIMP mass $\\mchi$ and the proton mass $m_{\\rm p}$, and \\beq \\sigmapSI = \\afrac{4}{\\pi} \\mrp^2 |f_{\\rm p}|^2 \\label{eqn:sigmapSI} \\eeq is the SI WIMP--nucleon cross section. Here the tiny mass difference between a proton and a neutron has been neglected. On the other hand, through e.g., squark and Z boson exchanges with quarks, WIMPs could also couple to the spin of target nuclei, an ``axial--vector'' interaction. The spin--dependent (SD) WIMP--nucleus cross section can be expressed as\\cite{SUSYDM96,Bertone05}: \\beq \\sigmaSD = \\afrac{32}{\\pi} G_F^2 \\~ \\mrN^2 \\afrac{J + 1}{J} \\bBig{\\Srmp \\armp + \\Srmn \\armn}^2 \\~. \\label{eqn:sigma0SD} \\eeq Here $G_F$ is the Fermi constant, $J$ is the total spin of the target nucleus, $\\expv{S_{\\rm (p, n)}}$ are the expectation values of the proton and neutron group spins, and $a_{\\rm (p, n)}$ are the effective SD WIMP couplings on protons and on neutrons. Some relevant spin values of the most used spin--sensitive nuclei are given in Table 1. For the SD WIMP--nucleus interaction, it is usually assumed that only unpaired nucleons contribute significantly to the total cross section, as the spins of the nucleons in a nucleus are systematically anti--aligned% \\footnote{ However, more detailed nuclear spin structure calculations show that the even group of nucleons has sometimes also a non--negligible spin (see Table 1 and e.g., data given in Refs.~\\refcite{SUSYDM96,Tovey00,Giuliani05,Girard05}). }. Under this ``odd--group'' assumption, the SD WIMP--nucleus cross section can be reduced to \\beq \\sigmaSD = \\afrac{32}{\\pi} G_F^2 \\~ \\mrN^2 \\afrac{J + 1}{J} \\expv{S_{\\rm (p, n)}}^2 |a_{\\rm (p, n)}|^2 \\~. \\label{eqn:sigma0SD_odd} \\eeq And the SD WIMP cross section on protons or on neutrons can be given as \\beq \\sigma_{\\chi {\\rm (p, n)}}^{\\rm SD} = \\afrac{24}{\\pi} G_F^2 \\~ m_{\\rm r, (p, n)}^2 |a_{\\rm (p, n)}|^2 \\~. \\label{eqn:sigmap/nSD} \\eeq \\begin{table}[t!] \\tbl{ List of the relevant spin values of the most used spin--sensitive nuclei. More details can be found in e.g., Refs.~1, 7, 8, 9.% }{ \\begin{tabular}{|| c c c c c c c c ||} \\hline \\hline \\makebox[1 cm][c]{Isotope} & \\makebox[0.5cm][c]{$Z$} & \\makebox[0.5cm][c]{$J$} & \\makebox[1 cm][c]{$\\Srmp$} & \\makebox[1 cm][c]{$\\Srmn$} & \\makebox[1.2cm][c]{$-\\Srmp/\\Srmn$} & \\makebox[1.2cm][c]{$\\Srmn/\\Srmp$} & \\makebox[3 cm][c]{Natural abundance (\\%)} \\\\ \\hline \\hline $\\rmXA{F}{19}$ & 9 & 1/2 & 0.441 & \\hspace{-1.8ex}$-$0.109 & 4.05 & $-$0.25 & 100 \\\\ \\hline $\\rmXA{Na}{23}$ & 11 & 3/2 & 0.248 & 0.020 & $-$12.40 & 0.08 & 100 \\\\ \\hline $\\rmXA{Cl}{35}$ & 17 & 3/2 & \\hspace{-1.8ex}$-$0.059 & \\hspace{-1.8ex}$-$0.011 & $-$5.36 & 0.19 & 76 \\\\ \\hline $\\rmXA{Cl}{37}$ & 17 & 3/2 & \\hspace{-1.8ex}$-$0.058 & 0.050 & 1.16 & $-$0.86 & 24 \\\\ \\hline $\\rmXA{Ge}{73}$ & 32 & 9/2 & 0.030 & 0.378 & $-$0.08 & 12.6 & 7.8 / 86 (HDMS)\\cite{Bednyakov08a}\\\\ \\hline $\\rmXA{I}{127}$ & 53 & 5/2 & 0.309 & 0.075 & $-$4.12 & 0.24 & 100 \\\\ \\hline $\\rmXA{Xe}{129}$ & 54 & 1/2 & 0.028 & 0.359 & $-$0.08 & 12.8 & 26 \\\\ \\hline $\\rmXA{Xe}{131}$ & 54 & 3/2 & \\hspace{-1.8ex}$-$0.009 & \\hspace{-1.8ex}$-$0.227 & $-$0.04 & 25.2 & 21 \\\\ \\hline \\hline \\end{tabular}} \\end{table} Due to the coherence effect with the entire nucleus shown in Eq.~(\\ref{eqn:sigma0SI}), the cross section for scalar interaction scales approximately as the square of the atomic mass of the target nucleus. Hence, in most supersymmetric models, the SI cross section for nuclei with $A~\\gsim~30$ dominates over the SD one\\cite{SUSYDM96,Bertone05}. ", "conclusions": "In this article I reviewed the data analysis procedures for extracting properties of WIMP--like Dark Matter particles from direct detection experiments. These methods are model--independent in the sense that neither prior knowledge about the velocity distribution function of halo Dark Matter nor their mass and cross sections on target nucleus is needed. The unique required information is measured recoil energies from experiments with different target materials. Once two or more experiments observe a few tens recoil events (in each experiment), one could in principle already estimate the mass and the SI coupling on nucleons as well as ratios between different cross sections of Dark Matter particles. All this information (combined eventually results from collider and/or indirect detection experiments) could then allow us to distinguish different candidates for (WIMP--like) Dark Matter particles proposed in different theoretical models and to extend our understanding on particle physics." }, "1003/1003.5892_arXiv.txt": { "abstract": " ", "introduction": "\\label{sec1} The electron-photon scattering in a magnetized plasma provides computable source terms for the evolution of the brightness perturbations \\cite{pera} which can be studied with diverse initial conditions, under various kinds of approximations and in different physical systems ranging from the classic problem of line formation of a normal Zeeman triplet \\cite{unno} to sychrotron emission \\cite{sync1,sync2,sync3,sync4}. In conventional Cosmic Microwave Background (CMB) studies the initial conditions of the temperature and polarization anisotropies are provided by the standard adiabatic mode \\cite{bert1,bert2}. The latter requirement implies that the initial radiation field lacks a specific degree of linear polarization. By initial radiation field we simply mean, in the present context, the initial data for the four Stokes parameters prior to matter-radiation equality. In the adiabatic scenario, the collision terms of the brightness perturbations for electron-photon scattering (see, e.g. \\cite{pera}) allow for the generation of linear polarization after photon decoupling while the circular polarization is, comparatively, not affected, i.e. it is vanishing both prior to matter-radiation equality and after photon decoupling. The study of the linear polarization (and of its potential generation via electron-photon scattering) has a long history \\cite{h1,h2,h3,h4,h5} dating even before the hypothesis of adiabatic initial conditions triggered by the formulation of inflationary scenarios. The preliminary detection of the polarization autocorrelations \\cite{QUAD1,QUAD2} (i.e. the EE power spectrum) as well as of the temperature-polarization correlations \\cite{WMAP7a,WMAP7b,WMAP7c,WMAP7d,WMAP7e,WMAP7f} (i.e. the TE power spectrum) confirms that the initial data, prior to matter-radiation equality, are predominantly adiabatic and lacking any specific degree of linear polarization which arises, to leading order in the tight-coupling expansion, because of the quadrupole of the temperature fluctuations. In a series of recent papers \\cite{mg1,mg2}, the idea has been to assume that the initial data of the radiation field are dictated by the conventional adiabatic mode but that the electron-photon scattering takes place in a magnetized environment. An amount of circular polarization is then produced because of the properties of magnetized electron-photon scattering. Analogies with the latter phenomenon can be found in the physics of the magnetized sunspots as well as in complementary astrophysical situations (see, for instance, \\cite{TS1,TS2}). The circular polarization arising from unpolarized initial conditions has then been computed and, to lowest order in the tight-coupling expansion, is proportional to the monopole of the intensity of the radiation field \\cite{mg1,mg2} (see also \\cite{BHR} for a different kind of derivation). The situation explored in \\cite{mg1,mg2} does not exhaust the possible sets of initial data for the system of brightness perturbations. The statement that the initial radiation field does not possess a specified (linear) polarization does not forbid, however, the presence of a primordial (circular) polarization. If the pre-decoupling plasma is not magnetized, then, the circular polarization will evolve independently both from the temperature fluctuations as well as from the linear polarization. Conversely, if the plasma is magnetized, the circular polarization can directly affect both the temperature anisotropies as well as the E-mode polarization. In the present paper we want to address a situation which is opposite to the one investigated in \\cite{mg1,mg2}. The idea will be to assume that the initial radiation field has an unknown amount of circular polarization. The purpose will then be to constrain the primordial V-mode by using the magnetized plasma as a polarimeter rather than as a polarizer. In \\cite{mg1,mg2} the magnetic field acted as a polarizer: the circular polarization was assumed to vanish initially and the problem was to compute the resulting V-mode signal. In the present paper the magnetic field will act as a polarimeter: the initial circular polarization does not vanish and the presence of the magnetic field is used as a diagnostic on the initial radiation field. The layout of the paper will therefore be the following. In section \\ref{sec2} the relevant governing equations will be introduced. In section \\ref{sec3} the V-mode contribution to the temperature autocorrelations (TT correlations for short) will be computed and constrained. In section \\ref{sec4} the same exercise will be repeated in the case of the polarization autocorrelations (EE correlations for short). The bounds obtained in sections \\ref{sec3} and \\ref{sec4} will be used in section \\ref{sec5} with the aim of inferring constraints on the autocorrelations of the circular polarization (VV power spectra for short). Section \\ref{sec6} contains the concluding remarks and the perspectives for future studies. \\renewcommand{\\theequation}{2.\\arabic{equation}} \\setcounter{equation}{0} ", "conclusions": "\\label{sec6} A primordial degree of circular dichroism, uncorrelated with the standard adiabatic mode, can be constrained if, prior to photon decoupling, the plasma is magnetized. The obtained results suggest the following considerations: \\begin{itemize} \\item{} if a V-mode power spectrum (not correlated with the adiabatic mode) is present prior to matter-radiation equality both the TT and the EE power spectra are affected in a computable manner; \\item{} constraints can then be inferred on the amplitude and spectral index of the V-mode power spectrum; \\item{} improved direct experimental limits on the VV correlations could be used for setting a limit on the magnetic field intensity. \\end{itemize} For experimental devices operating in the GHz range, direct limits on the circular dichroism imply constraints on pre-decoupling magnetic fields in the $10$ nG range. Conversely, the current limits on large-scale magnetic fields derived from the distortions of the TT and TE correlations (in the $0.1$ nG range) are compatible with current bounds on the primordial dichroism. Improved bounds on the V-mode polarization are not only interesting in their own right but they might have rewarding phenomenological implications. Direct limits on the V-mode power spectrum in the range ${\\mathcal O}(0.01\\, \\mathrm{mK})$ imply limits on ${\\mathcal A}_{\\mathrm{V}}$ ranging from ${\\mathcal O}(10^{-8})$ to ${\\mathcal O}(10^{-4})$ depending on the value of the spectral index and for large angular scales, i.e. larger than ${\\mathcal O}(1\\, \\mathrm{deg})$. \\newpage \\begin{appendix} \\renewcommand{\\theequation}{A.\\arabic{equation}} \\setcounter{equation}{0}" }, "1003/1003.0686_arXiv.txt": { "abstract": "Hierarchical models predict that present-day massive early-type galaxies (mETGs) have finished their assembly at a quite late cosmic epoch ($z\\sim 0.5$), conflicting directly with galaxy mass-downsizing. In \\citet{2010arXiv1002.3537E}, we presented a semi-analytical model that predicts the increase by a factor of $\\sim 2.5$ observed in the number density of mETGs since $z\\sim 1$ to the present, just accounting for the effects of the major mergers strictly-reported by observations. Here, we describe the relative, coordinated role of wet, mixed, and dry major mergers in driving this assembly. Accordingly to observations, the model predicts that: 1) wet major mergers have controlled the mETGs buildup since $z\\sim 1$, although dry and mixed mergers have also contributed significantly to it; 2) the bulk of this assembly takes place during the $\\sim 1.4$\\,Gyr time-period elapsed at $0.70.6$, becoming dry mergers dominant at $z<0.3$ \\citep[see][LSJ09a hereafter]{2003ApJ...597L.117K,2008ApJ...679..260M,2008ApJ...683L..17T,2009MNRAS.394.1956C,2009A&A...498..379D,2009MNRAS.396.2003L,2009A&A...501..505L}. As the observational properties of the most massive ETGs (mainly, ellipticals) are better reproduced by simulations of dry major mergers than of wet events \\citep{2003ApJ...597..893N,2006ApJ...636L..81N,2006ApJ...641...90R,2007MNRAS.379..401E}, some authors conclude that dry mergers at $z<0.3$ must have driven the final assembly of elliptical galaxies \\citep{2005AJ....130.2647V,2009arXiv0911.0044R}. This idea is also supported by the low number detected of blue galaxies bright enough to be the gas-rich progenitors of ellipticals \\citep[][]{2004ApJ...608..752B,2007ApJ...665..265F}. However, the significant drop registered in the number density of star-forming galaxies since $z\\sim 1$ and the characteristics of post-starburst galaxies entering in the red sequence at $z\\sim 0.7$ point more directly to wet than to dry events for explaining the stellar mass migration observed from the blue galaxy cloud to the red sequence in the last $\\sim 8$\\,Gyr \\citep{2006A&A...455..879Z,2007ApJS..172..406S,2010ApJ...709..644I,2009MNRAS.395..144W}. Summarizing, while some authors question not only the relevance of dry mergers in the final buildup of the mETGs \\citep{2006ApJ...644...54M,2007ApJS..172..494S}, but of all major mergers \\citep{2006ApJ...648..268B}, others claim that major merging must have been essential in it \\citep{2005ApJ...625..621B,2009ApJ...697.1369B,2006ApJ...640..241B,2007A&A...476..137A,2009MNRAS.395..144W,2010arXiv1001.4560L}. Definitely, the relative role of wet vs.\\,dry vs.\\,mixed mergers in the recent assembly of mETGs is still unsettled. In order to reconcile all these observational facts, mixed scenarios for the formation of mETGs have been proposed, in which blue galaxies have their SF quenched in gas-rich mergers, migrate to the red sequence, and merge further through mixed and dry mergers \\citep[][]{2007ApJ...665..265F,2008ApJS..175..390H,2008ApJS..175..356H}. A direct verification of the feasibility of this scenario accounting \\emph{strictly} for the major mergers reported by observations has not been carried out yet. So, we have approached this question directly through semi-analytical modelling, studying how the present-day mETGs would have evolved backwards-in-time under the hypothesis that each observed major merger gives place to an ETG. Results are being published in a series of papers. In the first paper of this series (\\citealt{2010arXiv1002.3537E}, \\citetalias{2010arXiv1002.3537E} hereafter), we showed that it is completely feasible to reproduce the observational buildup of $\\sim 50$-60\\% of present-day mETGs just accounting for the effects of the major mergers strictly reported by current observations since $z\\sim 1$. In the present-paper (Paper II), we analyse in detail how the coordinated action of wet, mixed, and dry mergers since $z\\sim 1$ explains this buildup, showing that many observational results that are apparently against the hierarchical scenario can be reconciled with it. The present paper is organized as follows. In \\S\\ref{sec:model}, we give a brief outline of the model. Section \\S\\ref{sec:results} is devoted to the presentation of results. In \\S\\ref{sec:numberevolution}, we analyse the model predictions on the number evolution experienced by mETGs at $z\\lesssim 1$ through major mergers. The relative role of each merger type in the recent buildup of mETGs is analysed in \\S\\ref{sec:roleofmergers}. Section \\S\\ref{sec:comparison} compares the model predictions with different observational and theoretical estimates at different redshifts. In \\S\\ref{sec:gaspoormergers}, we quantify the contribution of gas-poor mergers (dry and mixed) to the recent mETGs buildup. Section \\S\\ref{sec:mETGinplace} discusses the model predictions on the fraction of present-day mETGs that are really in place since $z\\sim 1$. The discussion of results and a brief summary of them can be found in \\S\\S\\ref{sec:discussion} and \\ref{sec:conclusions}, respectively. We will use a $\\Omega_M = 0.3$, $\\Omega_\\Lambda = 0.7$, $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$ concordant cosmology throughout the paper. All magnitudes are given in the Vega system. ", "conclusions": "\\label{sec:conclusions} In this paper, we analyse the relative role of wet, mixed, and dry major mergers in the recent buildup of mETGs, as derived from the model presented in \\citetalias{2010arXiv1002.3537E}. The model traces back-in-time the evolution of the local galaxy populations considering the number evolution derived from observational merger fractions and the L-evolution of each galaxy type due to typical SFHs. Several relevant results concerning to the recent buildup of mETGs have been derived:\\\\[-0.2cm] \\indent 1. The model demonstrates the feasibility of the mixed scenario proposed by \\citet{2007ApJ...665..265F}, proving that the coordinated effects of the wet, mixed, and dry major mergers strictly reported by current observations since $z\\sim 1$ can explain the increase by a factor of $\\sim 2.5$ observed in the number density of mETGs since then. \\\\[-0.3cm] \\indent 2. There have been $\\sim 2$ major mergers per local mETG since $z\\sim 1$, $\\sim 1$ corresponding to wet mergers and the another one to dry$+$mixed mergers. Therefore, although wet major mergers have played the dominant role in the recent buildup of mETGs, the contribution of dry and mixed mergers has also been essential in it.\\\\[-0.3cm] \\indent 3. The bulk of this mETGs assembly is predicted to take place during a $\\sim 1.4$\\,Gyr time-span at $0.7L^*$ galaxy amounts to $\\sim 65(^{110}_{40})$\\%. According to the model, there have been $\\sim 0.35$ wet major mergers, $\\sim 0.2$ mixed, and only $\\sim 0.1$ dry events per local $L>L^*$ galaxy. \\\\[-0.3cm] \\indent iv. There have been $\\sim 1.8(^{2.6}_{1.2})$ major mergers since $z\\sim 1$ per local mETG. Approximately one of these $\\sim 2$ mergers takes place during a period of $\\sim 0.5$\\,Gyr elapsed at $0.9100$~kpc distances from the host galaxy indicate the importance of jets interactions with the environment on many different physical scales. Morphology of X-ray clusters indicate that the radio-jet activity of a cD galaxy is intermittent. This intermittency might be a result of a feedback and/or interactions between galaxies within the cluster. Here we consider the radiation pressure instability operating on short timescales ($<10^5$ years) as the origin of the intermittent behaviour. We test whether this instability can be responsible for short ages ($ < 10^4$ years) of Compact Symmetric Objects measured by hot spots propagation velocities in VLBI observations. We model the accretion disk evolution and constrain model parameters that may explain the observed compact radio structures and over-abundance of GPS sources. We also describe effects of consequent outbursts. ", "introduction": "% The idea of intermittency is not new. Some observational evidence was given already in early 60-ties. For example \\cite{burbidge1965} suggested intermittent outbursts of NGC1275 the cD galaxy in the center of Perseus A cluster. \\cite{kellerman1966} derived the intermittency timescales of $\\sim 10^4-10^6$~years required for the production of relativistic particles responsible for the observed synchrotron spectra in radio sources and quasars. Signatures of the past recurrent activity in nuclei of normal galaxies were presented by \\cite{bailey1978}, but their paper was not really noticed and has only 24 citations to date. \\cite{shields1978} examined quasar models and suggested that their accumulate the mass during quiescent periods and then through the instability they transfer the mass onto a central black hole during a short period of an outburst of the activity. These are just a few examples of the AGN intermittency that has been considered since the early days of studies of the nuclear emission in galaxies. It is now that we are looking closely at this behaviour as it becomes evident that it is an important component to our understanding of the evolution of structures in the universe. However, there are still many open questions about the origin of the intermittent behaviour. Is it related to the unsteady fuel supply, or accretion flow? Are there many quiescent and outburst phases? What is the mechanism regulating the intermittent behaviour? Observations show a range of timescales for the AGN outbursts. Quasar lifetimes estimated based on large samples of SDSS quasars are of order of $10^7$~years \\citep{martini2001}. Signatures of outbursts in recent observations of X-ray clusters indicated similar timescales \\citep[see][for the review]{mcnamara2007}. However, episodes of activity on timescales shorter than $<10^5$~years have been also observed for example in compact radio sources \\citep{owsianik1998,reynolds1997} or as light echos in nearby galaxies \\citep{lintott2009}. In this review we will focus on the short timescales of the intermittent jet activity. We discuss the radio source evolution, observational evidence for the intermittent activity and present the model for the origin and nature of the short term activity based on the accretion disk physics. ", "conclusions": "Observations indicate a complex behaviour of radio sources: continuous jets, signatures of repetitive outbursts in separated radio components, or the statistic of radio sources. The source complex behaviour may reflect different regimes of the accretion flow that depend on the black hole mass and accretion rate. We also note that for some parameters the radio source may never leave the host galaxy. Large samples of radio sources are needed for statistical studies. We should be able to determine the number of sources, their lifetimes and sizes. We also need more sources with measurements of their age as well as indications for the intermittency in the sources radio morphology. Such data should be available in the future with the new radio surveys that probe fainter, low power compact sources that might be in the fading phase." }, "1003/1003.5936.txt": { "abstract": "In this paper, we aim to estimate the vertical gradients in the rotational velocity of the Galaxy. This is carried out in the framework of a global thin disc model approximation. The predicted gradient values coincide with the observed vertical fall-off in the rotation curve of the Galaxy. The gradient is estimated based on a statistical analysis of trajectories of test bodies in the gravitational field of the disc and in an analytical way using a quasi-circular orbit approximation. The agreement of the results with the gradient measurements is remarkable in view of other more complicated, non-gravitational mechanisms used for explaining the observed gradient values. Finally, we find that models with a significant spheroidal component give worse vertical gradient estimates than the simple disc model. In view of these results, we can surmise that, apart from the central spherical bulge and Galactic halo, the gross mass distribution in the Galaxy forms a flattened rather than spheroidal figure.\\\\ \\medskip \\hrule \\flushleft \\textbf{The definitive version is available at \\\\ \\texttt{http://onlinelibrary.wiley.com/\\\\doi/10.1111/j.1365-2966.2010.16987.x/abstract}} \\medskip \\hrule ", "introduction": "Recently, apart from other local and global characteristics of the rotation speed of the Galaxy, such as rolling motion, \\citet{bib:gradients} determined the fall-off rate $\\gamma=-22\\pm6\\gu$ in the rotation speed from the Galactic mid-plane. This gradient estimation was obtained by fitting a linear profile\\footnote{Throughout this paper we use cylindrical coordinates $(r,\\varphi,z)$ distinguished by the galactic mid-plane $z=0$ and the axis of rotation $r=0$.} \\begin{equation}\\label{eq:model_rot}v_{\\varphi}(r,z)=v_{\\varphi}(r,0)+\\gamma \\abs{z}+\\delta z\\end{equation} to rotation measurements in the vicinity of the Galactic mid-plane: $\\abs{z}<0.1\\,\\kpc$, $r\\in\\br{3,8}\\,\\kpc$. The principal purpose of our work is to reconstruct the vertical gradient magnitude in a simple model of the Galaxy and to find out to what extent the gradient behaviour is dependent on the geometry of mass distribution. In particular, the gradient value can be very well reconstructed in the global thin disc model, provided the disc comprises gross dynamical mass ascertained from the rotation curve of the Galaxy. \\subsection{Motivation of the present work} There is a suggestive heuristics behind the above simple ansatz for the rotation speed. This naturally leads us directly to the global disc model approximation as a mean of determining the vertical gradient of the rotation speed in flattened galaxies, in particular, in the Galaxy. It is a simple matter to note that every axisymmetric function $f(r,z)$ in the cylindrical coordinate frame can be written (for $z\\ne0$) as \\begin{equation}\\label{eq:f}f(r,z)=f(r,0)+ \\Gamma(r,z) \\abs{z}+ \\Delta(r,z) {}z\\end{equation} Here, $\\Gamma$ and $\\Delta$ are $z$-symmetric functions defined as \\[\\begin{array}{@{}l@{\\,\\,\\,}r@{}}\\Gamma(r,z)=\\frac{f(r,\\abs{z})+f(r,-\\abs{z})-2f(r,0)}{2\\abs{z}},& \\Delta(r,z)=\\frac{f(r,\\abs{z})-f(r,-\\abs{z})}{2\\abs{z}}\\end{array}.\\] Consider now the azimuthal velocity $v_{\\varphi}$ in place of the function $f$, $f(r,z)=v_{\\varphi}(r,z)$. We expect $v_{\\varphi}(r,z)$ to be nearly $z$-symmetric. In this case, the difference $f(r,\\abs{z})-f(r,-\\abs{z})$ is small compared to the symmetric part $f(r,\\abs{z})+f(r,-\\abs{z})\\approx 2f(r,z)$. From the Lagrange mean value theorem, it then follows that $\\Gamma\\approx\\frac{f(r,z)-f(r,0)}{\\abs{z}}=f_{,\\zeta}(r,\\zeta)|_{\\zeta=hz} \\sgn{z}$, $00 \\] However, the absolute gradient value would be much lower for a spherically symmetric model with the same rotation curve \\[ -\\frac{3}{2r}\\sqrt{\\frac{G M}{r}} \\frac{z/r}{\\br{1+(a^2+z^2/r^2)}^{{7}/{4}}},\\quad z>0. \\] and it would vanish at the disc plane. The gradients would become comparable only at some height off the disc. In general, the contribution to the vertical gradient of rotation from a spherical potential also vanishes in the galactic mid-plane, whereas a similar contribution from the gravitational field of a disc remains non-zero. From this, it also follows that to attain the actual gradient values at larger heights off the galactic mid-plane (outside the main concentration of masses, where the 'thin disc gravitation' is comparable with the 'wide disc gravitation'), the disc model gradient must decrease in absolute value, while the spherical model gradient must increase in absolute value. Qualitatively, this explains why the vertical fall-off of rotation is, on average, weaker when the gravity of a flattened, disc-like component is dominated by the gravity of a spheroidal component. The presence of a massive spheroidal component reduces the overall vertical gradient in two ways: (i) because (as we have already seen) the contributions to vertical gradients from spheroidal components are small; (ii) because the presence of a spherical component uses up masses that would otherwise be present in the disc component, enhancing the disc-like contributions to the overall gradient. \\subsection{Outline of the results} In the framework of a global thin disc model approximation for the Milky Way, we use various methods for estimating the vertical gradient in rotation speed. All the methods lead to comparable results. In particular, we find the gradients for $0.22\\leq\\abs{z}\\leq2.62\\,\\kpc$ by using an averaging method, mimicking realistic gradient measurements in galaxies other than the Milky Way, and for $0<\\abs{z}<3.6\\,\\kpc$ by analysing the motion of test bodies in the gravitational potential of the Galaxy described in the thin disc model approximation. The gradient values are compared with the observed value determined by \\citet{bib:gradients} from measurements of the rotation of the Milky Way in the region $\\abs{z}<0.1\\,\\kpc$. Although our various gradient estimates are carried out at larger $z$, we consider this comparison possible. One argument for this is that the value of the velocity fall-off is almost constant with the distance off the mid-plane we encountered in our analysis. Surely, it would be best to measure the gradient values in the Galaxy at heights larger than $0.1\\,\\kpc$; however, such data are not available. However, we can utilize the observational fact that the gradient is constant in other galaxies. In addition, we also calculated the gradient analytically in the region $\\abs{z}<0.1\\,\\kpc$, again obtaining results consistent with the observed gradient value in this region. The gradient values we predicted using our simple model agree amazingly well with those obtained by \\cite{bib:gradients} even though the Galaxy rotation curve we utilized is unrelated to that paper. This agreement is striking in view of the fact that the gradients are usually considered to originate from more complex physics than gravitational alone, to mention only radial pressure gradients, magnetic tension, or galactic fountains. For completeness, we compare our results with those predicted for the gradient by a customary three-component model or a maximum halo model. It turns out that the gradient features are crucially dependent on whether the gross mass distribution is more flattened or more spheroidal, and we have already encountered this feature when considering a simple example in Section \\ref{sec:efficiency}. We find that the gradient estimates for the Galaxy predicted in the disc model are in better agreement with observations than analogous predictions of models with a significant spheroidal mass component. Finally, we apply the disc model and estimate the vertical gradients in the galaxy NGC 891, where rotation was measured far from the galactic mid-plane, again obtaining results consistent with measurements. These results show that the simple thin disc model approximation performs well in reconstructing the vertical gradient properties. Of course, this does not mean that galactic discs are thin. However, it does provide a strong argument for the fact that gross mass distribution in the Milky Way and in NGC 891 might be flattened rather than spheroidal. \\subsection{A model (SHO) rotation curve of Milky Way Galaxy} Apart from the main reference \\citep{bib:gradients}, the observational basis for our analysis is a unified set of data representing the rotation velocity of the Galaxy collected by \\citeauthor{bib:sofue_unified} (2009, hereafter SHO). Based on these data, we derive a substitute thin disc surface mass density accounting for the Galaxy rotation and representing a sort of projection of Galaxy masses on to the Galactic mid-plane. If gross Galaxy matter forms a flattened structure, as seen in other spiral galaxy pictures (apart from the galactic luminous halo and the central spheroidal bulge, the latter, nevertheless, commonly described by an equivalent de Vaucouleurs disc), then the actual Galaxy gravitational potential should be better approximated by the disc model than by a model with a significant spheroidal component comprising most of the dynamical Galaxy mass inferred from the Galaxy rotation. The data set is a compilation of several independent measurements of the rotation, suitably adjusted to each other. Unfortunately, the rotation data are poor at larger distances where the measurement points are scattered randomly and are determined with large errors. This leaves too much room for arbitrariness, as neither declining nor even rising rotation could be excluded at outer radii (however, this inaccurately measured, remote Galaxy region is not very important to our analysis). Therefore, instead of the unified measurements, we decided to use a published model rotation curve that accounts for the data satisfactorily, mainly in the $\\,3-8\\,\\kpc$ region in which we are interested. This curve was obtained by SHO for the unified measurement data by fitting a three-component model, including in addition a wavy ring pattern superposed on the exponential disc component. For brevity, we refer to this particular fit as the SHO rotation curve and denote it by $v_{\\sigma}$. The SHO rotation curve agrees, as a constrained least-squares fit, with the Galaxy rotation curve almost perfectly out to large radii (however, any other fit, such as a linear combination of Hermite polynomials, a Fourier series, etc., anything one wishes, would also be acceptable, if only the fit was sufficiently smooth and close to the data in the sense of some norm, such as the least-squares method norm). The rings allowed the authors to reproduce the observed dips in rotation. Leaving the SHO mass model aside, the SHO fit alone can be treated as an actually measured rotation curve of the Galaxy, as it agrees very well with real rotation measurements in the internal Galaxy region, accounting for qualitative characteristics of the rotation. This curve agrees also with additional constraints imposed on the curve in SHO and known from separate precise single rotation measurements.\\footnote{For clarity, it should be noted that a square root in the rotation law for the assumed dark halo mass distribution was unwittingly missed in SHO. This is probably not a mere misprint as a plot of the dark halo contribution to the overall rotation curve presented in their paper is parabolic at the centre, despite being linear for such a halo. Obviously, this mistake is not important to our analysis and does not affect our results.} ", "conclusions": "\\label{sec:summary} Throughout this paper, we have given several arguments to suggest that the global thin disc model of flattened galaxies is naturally suited for describing the large vertical gradients of rotation speed observed in the neighbourhood of the galactic mid-plane. The gradients were estimated in the quasi-circular orbit approximation, which was established in Section \\ref{sec:diskgrad} to give trustworthy results. The gradient values predicted for the Galaxy in this approximation agree very well with measurements when the disc comprises gross dynamical mass. Independent estimates of the gradient in this model (both analytical and from analysing the motion of test bodies) give consistent results. The vertical fall-off in the rotational velocity in this model is not very dependent on the height, at least out to $3\\,\\kpc$ above the mid-plane. This result is consistent with observations of the rotation speed in other galaxies. In other words, the distance from the mid-plane is not crucial for the gradient determination. Thus, more realistic mass models of the Galaxy, such as a finite width disc comprising the whole dynamical mass, should give similar results. We have also found that Galaxy models with significant spheroidal component are not consistent with the gradient measurements, and the discrepancy grows with the mass of the component. With the aid of our model, we have also studied the vertical gradient in NGC 891. The gradients in this galaxy were measured at relatively large heights from the galactic disc compared to the measurements in our Galaxy. Nevertheless, we again obtained predictions in accordance with observations. Based on all these results, we can hypothesize, contrary to what is implied by dark halo models, that gross mass distribution in our Galaxy is more flattened, disc-like, rather than spheroidal." }, "1003/1003.0259_arXiv.txt": { "abstract": "A new two-parametric family of mass distribution for spherical stellar systems is considered. It generalizes families by \\citet{KV72} and by \\citet{AE06}. Steady velocity dispersions are found for these models by solving an equation of hydrostatic equilibrium. Axisymmetric generalizations of the model are discussed. ", "introduction": "In this study we suggest a new family of spherical mass distribution models that generalizes models by An \\& Evans (2006, hereafter AE) and models by Kuzmin et al. \\citep{KV72,KM69}. The family depends on two structural parameters. It includes Plummer's spheres \\citep{pl11}, H\\'{e}non's isochrones \\citep{Henon59} and the model by Hernquist \\citep{lars90} as special cases. \\medskip ", "conclusions": "" }, "1003/1003.2295_arXiv.txt": { "abstract": "{Shock modelling predicts an electron density enhancement within the magnetic precursor of C-shocks. Previous observations of SiO, H$^{13}$CO$^{+}$, HN$^{13}$C and H$^{13}$CN toward the young L1448-mm outflow showed an over-excitation of the ion fluid that was attributed to an electron density enhancement in the precursor.} {We re-visit this interpretation and test if it still holds when we consider different source morphologies and kinetic temperatures for the observed molecules. To do this, we use updated collisional coefficients of HN$^{13}$C and SiO with electrons in our excitation model. We also aim to give some insight on the spatial extent of the electron density enhancement around L1448-mm.} {We estimate the opacities of H$^{13}$CO$^+$ and HN$^{13}$C by observing the $J$=3$\\rightarrow$2 lines of rarer isotopologues. To model the excitation of the molecules, we use the large velocity gradient (LVG) approximation with updated collisional coefficients to i) re-analyse the observations toward the positions where the over-excitation of H$^{13}$CO$^+$ has previously been observed [i.e. toward L1448-mm at offsets (0,0) and (0,-10)], and ii) to investigate if the electron density enhancement is still required for the cases of extended and compact emission, and for kinetic temperatures of up to 400~K. We also report several lines of SiO, HN$^{13}$C and H$^{13}$CO$^+$ toward new positions around this outflow, to investigate the spatial extent of the over-excitation of the ions in L1448-mm.} {From the isotopologue observations, we find that the emission of H$^{13}$CO$^+$ and HN$^{13}$C from the precursor is optically thin if this emission is extended. Using the new collisional coefficients, an electron density enhancement is still needed to explain the excitation of H$^{13}$CO$^+$ for extended emission and for gas temperatures of $\\leq 400$~K toward L1448-mm (0,-10), and possibly also toward L1448-mm (0,0). For compact emission the data cannot be fitted. We do not find any evidence for the over-excitation of the ion fluid toward the newly observed positions around L1448-mm.} {The observed line emission of SiO, H$^{13}$CO$^{+}$ and HN$^{13}$C toward L1448-mm (0,0) and (0,-10) is consistent with an electron density enhancement in the precursor component, if this emission is spatially extended. This is also true for the case of high gas temperatures ($\\leq$400$\\,$K) toward the (0,-10) offset. The electron density enhancement seems to be restricted to the southern, redshifted lobe of the L1448-mm outflow. Interferometric images of the line emission of these molecules are needed to confirm the spatial extent of the over-excitation of the ions and thus, of the electron density enhancement in the magnetic precursor of L1448-mm.} ", "introduction": "C-shock waves associated with molecular outflows are belived to develop a thin region of enhanced fractional ionisation known as the `magnetic precursor' \\citep{Draine80,Flower96,Flower03}. Inside this region, the magnetic field is gradually compressed, forcing the ions to stream through the neutral gas, accelerating, compressing and heating this fluid before the neutral one. The subsequent ion-neutral velocity decoupling leads to the sputtering of dust grains, injecting large amounts of molecular material into the gas phase \\citep[e.g.][]{Caselli97,Schilke97,J-S08} The electron density is predicted to be enhanced within the magnetic precursor by a factor of $\\sim100$ due to the fluorescence UV photons generated after the collisional excitation of H$_2$ molecules \\citep{Flower96,Flower03}. It has been proposed that a narrow (line-width $\\sim 0.5$~km~s$^{-1}$) and very low velocity component of SiO detected towards the very young \\object{L1448-mm} outflow, is the signature of the interaction of the magnetic precursor (Jim{\\'e}nez-Serra et al. \\citeyear{J-S04}, hereafter JS04). The derived fractional abundance of SiO is of the order $10^{-11}$ in the velocity component of the precursor, an enhancement by a factor of 10 with respect to the upper limits measured in the quiescent gas of dark clouds \\citep[$\\leq$10$^{-12}$ in L183 and L1448;][]{Ziurys89,RT07}. Although the origin of this narrow SiO emission is still unclear, it has been suggested that the sputtering of dust grains at the precursor stage, is efficient enough to inject a considerable fraction of the grain mantles into the gas phase. As a consequence, bright and narrow SiO line emission is expected to arise from material whose central radial velocities are very similar to that of the ambient cloud \\citep{J-S08,J-S09}. In correlation with the detection of narrow SiO emission, Jim{\\'e}nez-Serra et al. (\\citeyear{J-S06}, hereafter JS06) reported differences in the excitation of molecular ions, such as H$^{13}$CO$^{+}$, with respect to neutral molecules, such as HN$^{13}$C and H$^{13}$CN, in the precursor component. In particular, the high-$J$ line emission of H$^{13}$CO$^+$ is substantially brighter than that of HN$^{13}$C toward those regions where the precursor has been detected. JS06 proposed that these differences, which cannot be accounted for by considering {\\it only} molecular excitation by collisions with H$_2$ and a single H$_2$ density of few $\\times$10$^5$$\\,$cm$^{-3}$ for the molecular gas, could be produced by the {\\it selective} excitation of molecular ions by collisions with electrons within the precursor. This study established that the over-excitation of the ions is consistent with an electron density enhancement by a factor of $\\sim$500 in the precursor component toward L1448-mm (0,0) and (0,-10), which is similar to that predicted by C-shock modelling at this dynamical time \\citep{Flower03}. However, the observed differences in excitation between the ion and the neutral molecules could be alternatively explained by opacity effects. Indeed, \\citet{Frerking79} and \\citet{Cernicharo84} showed anomalies in the {\\it large-scale} line intensity emission of the hyperfine components of HNC and HCN toward the molecular dark clouds TMC-1 and L134N. The equal intensities measured for these components could be produced by the absorption of the emission arising from the dense cores by the less dense foreground material. This also applies to other high gas density tracers with high dipolar moment such as HCO$^+$, since they also show similar behaviors for the line intensity ratios between its isotopologue species \\citep{Langer78}. JS06 derived the electron density enhancement in the precursor component of L1448-mm by assuming that the H$^{13}$CO$^{+}$, the HN$^{13}$C and the H$^{13}$CN emission had thin to moderate optical depths. Unfortunately, the lack of even rarer isotopologue observations toward this outflow prevented, first, to determine the actual optical depths of this emission, and second, to clearly establish the origin of the over-excitation of the ions in L1448-mm. In this paper we present new observations of $J$=3$\\to$2 lines of the rare isotopologues HC$^{18}$O$^+$, H$^{13}$C$^{18}$O$^+$, H$^{15}$NC and H$^{15}$N$^{13}$C, toward the regions where the over-excitation of the ions has been reported. In addition, we present new observations of the $J$=3$\\to$2 lines of HN$^{13}$C and H$^{13}$CO$^+$ toward extra positions around the L1448-mm source, which we combine with previous observations toward these positions, observed by JS04, to derive the spatial extent of the over-excitation of H$^{13}$CO$^+$ associated with the precursor. These observations will allow to prove that the anomalous excitation of the ions is not due to a large-scale scattering effect, but to a real enhancement of the local density of electrons at the early stages of the interaction of very young C-shocks. We also re-visit the analysis of JS06, using the same excitation model as JS06 but including new collisional data of HNC and SiO with electrons \\citep[][respectively]{Faure07,Varambhia09}, to test if the conclusions of JS06 still hold. The excitation model employed uses the large velocity gradient (LVG) approximation. We note that the L1448-mm outflow is likely the best object where the effects of collisional excitation by electrons on the molecular excitation can be directly tested. Furthermore, in our re-analysis we fully explore the effects of source morphology and kinetic gas temperature on our results. In Sections~\\ref{obs_description} and \\ref{results}, we present the observations carried out with the JCMT telescope and describe the results. In Section~\\ref{opacities_calc}, we estimate the opacities of H$^{13}$CO$^+$ and HN$^{13}$C toward the positions where the over-excitation of the ions has been reported, assuming both compact and extended source emission. In Section~\\ref{section:excitation}, we present the re-analysis of the LVG calculations of JS06 using the new collisional coefficients, considering different source morphologies and temperatures up to 400~K. In this section we also analyse the data measured toward the new positions in L1448-mm to provide some insight into the spatial extent of the over-excitation of H$^{13}$CO$^+$. Finally, discussion and conclusions are presented in Section~\\ref{obs_conclusions}. ", "conclusions": "\\label{obs_conclusions} Anomalies in the intensities of the hyperfine components of the strongly polar HCN molecule were reported on large-scales toward the TMC-1 molecular dark cloud by \\cite{Cernicharo84}. These authors showed that these anomalies are due to large opacity effects and are produced by the re-emission of radiation coming from the densest part of the core, by a more extended and less dense envelope. If this applies to the L1448-mm case, then we should expect to detect such differences in the excitation of other strongly polar molecular species such as HCO$^{+}$ and HNC, on large-scales around the L1448-mm core. In Section~\\ref{opacities_calc}, from isotopologue observations we have shown that the precursor component of H$^{13}$CO$^+$ is optically thin toward (0,0), and if the emission is extended, this component is also optically thin toward (0,-10). This implies that the over-excitation of H$^{13}$CO$^+$ in the precursor toward these positions cannot be due to a large-scale optical depth/scattering effect. In view of the recent release of new collisional coefficients of SiO and HNC with electrons \\citep{Faure07,Varambhia09}, we have re-analysed the multi molecular line observations toward L1448-mm (0,0) and (0,-10) of JS06, but taking into account different possibilities for the source size extent and temperature. Assuming extended emission and kinetic temperatures of up to 400~K, and considering that the molecular excitation is due to collisions with only H$_2$, a higher H$_2$ density is required to match the observations of H$^{13}$CO$^+$ than that needed for the neutral species. For temperatures of $100-400$~K, toward L1448-mm (0,0) it is possible that, considering the $1\\sigma$ errors in the observations, the observations can be explained by a single H$_2$ density for the ions and the neutrals, but toward (0,-10) these errors are not enough to explain the difference in the required H$_2$ densities. This leads to the need for an extra excitation mechanism for H$^{13}$CO$^+$ in the precursor component of L1448-mm. This extra excitation can be provided by an electron density enhancement of up to a factor of $\\sim10^4$. As in JS06, the derived electron density enhancement exceeds those predicted by MHD shock modelling by up to two orders of magnitude \\citep{Flower96,Flower03}. However, simply considering the $1\\sigma$ errors in the observations can reduce the electron density estimates by an order of magnitude. Due to these very large uncertainties, we can conclude that an electron density enhancement can explain the observations, but we cannot quantify it with an accuracy better than one order of magnitude. If the emission is compact, the gas kinetic temparture must be $\\gtrsim50$~K. However, we find that the HN$^{13}$C and H$^{13}$CO$^+$ lines cannot be fitted in this case, so we conclude that the emitting region associated with the precursor must be $>5''$. The observations toward the L1448-mm (0,20), (20,0) and (-20,0) positions, do not show evidence for an over-excitation of H$^{13}$CO$^{+}$ with respect to the neutral HN$^{13}$C molecules in the precursor component of L1448-mm. This suggests that the extra-mechanism responsible for the over-excitation of the ions has its origin in local phenomena likely linked to the recent interaction of very young C-shocks. \\cite{Girart01} have indeed reported very young shocks (with a dynamical age of $\\leq 90$~yr) toward the southern position L1448-mm (0,-10), making the probability to detect the magnetic precursor of C-shocks not negligible. Since our results show that the over-excitation of ions is confined to this region, there seems to be strong evidence to suggest that this over-excitation is produced by the electron density enhancement at the magnetic precursor stage of young shocks. Furthermore, the confinement of the over-excitation of the ions also implies that this effect is unlikely to be caused by external UV illumination, as it should be observed tracing the inner regions of the outflow cavity, and on larger spatial scales. In order to fully understand the extent of the magnetic precursor of L1448-mm, high angular resolution observations of the high-$J$ transitions of H$^{13}$CO$^+$, HN$^{13}$C and SiO carried out with interferometers such as the Submillimeter Array (SMA), are therefore strongly required to measure the source sizes of this molecular emitting regions and to confirm whether or not this emission is optically thin. Observations of very high-$J$ ($J>10$) transitions are also required to verify whether or not the gas has reached the high temperatures predicted by C-shock models, as these lines should be detected for such temperatures. It is also necessary to search for signs of the electron density enhancement toward other young outflows where narrow SiO has been detected, such as L1448-IRS3, NGC1333-IRS4 and NGC1333-IRS2, in order to investigate whether the electron density enhancement is a common phenomenon in this type of objects, and if it is correlated with the detection of very narrow SiO emission centred at ambient cloud velocities." }, "1003/1003.5468_arXiv.txt": { "abstract": "The time delays between gamma-rays of different energies from extragalactic sources have often been used to probe quantum gravity models in which Lorentz symmetry is violated. It has been claimed that these time delays can be explained by or at least put the strongest available constraints on quantum gravity scenarios that cannot be cast within an effective field theory framework, such as the space-time foam, D-brane model. Here we show that this model would predict too many photons in the ultra-high energy cosmic ray flux to be consistent with observations. The resulting constraints on the space-time foam model are much stronger than limits from time delays and allow for Lorentz violation effects way too small for explaining the observed time delays. ", "introduction": " ", "conclusions": "" }, "1003/1003.0883_arXiv.txt": { "abstract": "{Magnetars are neutron stars possessing magnetic field of about $10^{14}$-$10^{15}$~G at the surface. Thermodynamic properties of neutron star matter approximated by pure neutron matter are considered at finite temperature in strong magnetic fields up to $10^{18}$~G which could be relevant for the interior regions of magnetars. In the model with the Skyrme effective interaction, it is shown that a thermodynamically stable branch of solutions for the spin polarization parameter corresponds to the case when the majority of neutron spins are oriented opposite to the direction of the magnetic field (negative spin polarization). Besides, the self-consistent equations, beginning from some threshold density, have also two other branches of solutions corresponding to positive spin polarization. The influence of finite temperatures on spin polarization remains moderate in the Skyrme model up to temperatures relevant for protoneutron stars. In particular, the scenario with the metastable state characterized by positive spin polarization, considered at zero temperature in Phys. Rev. C {\\bf 80}, 065801 (2009), is preserved at finite temperatures as well. It is shown that above certain density the entropy for various branches of spin polarization in neutron matter with the Skyrme interaction in a strong magnetic field demonstrates the unusual behavior being larger than that of the nonpolarized state. By providing the corresponding low-temperature analysis, it is clarified that this unexpected behavior should be addressed to the dependence of the entropy of a spin polarized state on the effective masses of neutrons with spin up and spin down, and to a certain constraint on them which is violated in the respective density range. } ", "introduction": "Magnetars are strongly magnetized neutron stars (\\cite{DT}) with the magnetic field strength at the surface of about $10^{14}$-$10^{15}$~G (\\cite{TD,IShS}). Such huge magnetic fields can be inferred from observations of magnetar periods and spin-down rates, or from hydrogen spectral lines. Among possible classes of various neutron stars, soft gamma-ray repeaters and anomalous X-ray pulsars are believed to be most probable candidates for these ultrastrong magnetized astrophysical bodies (\\cite{WT}). Magnetars are relatively frequent objects in the Universe and comprise about $10\\%$ of the whole population of neutron stars (\\cite{K}). In the interior of a magnetar the magnetic field strength may be even larger, reaching values of about $10^{18}$~G (\\cite{CBP,BPL}). Therefore, magnetars provide a unique playground for studying properties of neutron star matter under extreme conditions of density and magnetic field strength (\\cite{CBP,BPL,CPL,PG,IY09}), which are inaccessible in the terrestrial laboratories. In the recent study by Perez-Garcia~(2008), neutron star matter was approximated by pure neutron matter in a model with the effective nuclear forces. It has been shown that the behavior of spin polarization of neutron matter in the high density region in a strong magnetic field crucially depends on whether neutron matter develops a spontaneous spin polarization (in the absence of a magnetic field) at several times nuclear matter saturation density, or the appearance of a spontaneous polarization is not allowed at the relevant densities (or delayed to much higher densities). The first case is usual for the Skyrme forces~(\\cite{R,S,O,VNB,RPLP,ALP,MNQN,TT,BPM,KW94,I,IY04a,RPV,I06}), while the second one is characteristic for the realistic nucleon-nucleon (NN) interaction~(\\cite{PGS,BK,H,VPR,FSS,KS,BB}). In the former case, a ferromagnetic transition to a totally spin polarized state occurs while in the latter case a ferromagnetic transition is excluded at all relevant densities and the spin polarization remains quite low even in the high density region. If a spontaneous ferromagnetic transition is allowed, it was shown in the subsequent model consideration with the Skyrme effective forces (~\\cite{IY09}) that the self-consistent equations for the spin polarization parameter at nonzero magnetic field have not only solutions corresponding to negative spin polarization (with the majority of neutron spins oriented opposite to the direction of the magnetic field) but, because of the strong spin-dependent medium correlations in the high-density region, also the solutions with positive spin polarization. In the last case, the formation of a metastable state with the majority of neutron spins oriented along the magnetic field is possible in the high-density interior of a neutron star. In this study, we extend our previous consideration of neutron matter with Skyrme forces in a strong magnetic field, given at zero temperature~(\\cite{IY09}), to finite temperatures up to a few tens of MeV being relevant for protoneutron stars. As a framework for consideration, we use a Fermi liquid approach for the description of nuclear matter~(\\cite{AKPY,AIP,IY3}). One of the goals of the research is to study the impact of finite temperatures on spin polarized states in neutron matter in a strong magnetic field, and, in particular, to clarify whether the formation of a metastable state with positive spin polarization is possible in the high-density region of neutron matter. Besides, we provide a fully self-consistent finite temperature calculation of the basic thermodynamic functions of spin polarized neutron matter in a strong magnetic field. It will be shown that beginning from some density the entropy for various branches of spin polarization in neutron matter in a strong magnetic field demonstrates the unusual behavior being larger than that of nonpolarized neutron matter. We relate this unexpected result to the dependence of the entropy on the effective masses of neutrons with spin up and spin down, and to the violation of a certain constraint on them in the corresponding density range. Note that we consider thermodynamic properties of spin polarized states in neutron matter in a strong magnetic field up to the high density region relevant for astrophysics. Nevertheless, we take into account the nucleon degrees of freedom only, although other degrees of freedom, such as pions, hyperons, kaons, or quarks could be important at such high densities. ", "conclusions": "We have studied the impact of finite temperatures on the spin structure in the magnetar interior, approximating the neutron star matter by pure neutron matter and taking the Skyrme effective interaction as a potential of NN interaction (SLy7 parametrization). According to the scalar virial theorem, strong magnetic fields up to $10^{18}$~G can be relevant for the interior regions of magnetars. It has been shown that, together with the thermodynamically stable branch of solutions for the spin polarization parameter corresponding to the case when the majority of neutron spins are oriented opposite to the direction of the magnetic field (negative spin polarization), the self-consistent equations, beginning from some threshold density, have also two other branches of solutions corresponding to positive spin polarization. The influence of finite temperatures on spin polarization remains moderate in the Skyrme model, at least, up to temperatures relevant for protoneutron stars. In particular, a thermodynamic analysis, based on the calculation of the free energy for different branches of spin polarization, shows that the scenario with the metastable state characterized by positive spin polarization, considered at zero temperature~(\\cite{IY09}), is preserved at finite temperatures as well. This is one of the important conclusions of the given research. The possible existence of a metastable state with positive spin polarization will affect the neutrino opacities of a neutron star matter in a strong magnetic field, and, hence, will influence the cooling history of a neutron star~(\\cite{RPLP}). We have also shown that above certain density the entropy for various branches of spin polarization in neutron matter with the Skyrme interaction in a strong magnetic field demonstrates the unusual behavior being larger than that of the nonpolarized state. To clarify this point, we have provided the corresponding low-temperature analysis. It has been shown that this unexpected behavior should be addressed to the dependence of the entropy of a spin polarized state on the effective masses of spin-up and spin-down neutrons and to a certain inequality constraint on them which is violated in the respective density range. It is worthy to note that in the given research a neutron star matter was approximated by pure neutron matter. This should be considered as a first step towards a more realistic description of neutron stars taking into account a finite fraction of protons with the charge neutrality and beta equilibrium conditions. In particular, some admixture of protons can affect the onset densities of enhanced polarization in a neutron star matter with the Skyrme interaction~(\\cite{IY04a}). Nevertheless, at such strong magnetic fields, one can expect that the proton fraction is relatively small and even can completely disappear in the dense interior of a magnetar~(\\cite{MKI})." }, "1003/1003.4002_arXiv.txt": { "abstract": "Beginning with a historical account of the spectral classification, its refinement through additional criteria is presented. The line strengths and ratios used in two dimensional classifications of each spectral class are described. A parallel classification scheme for metal-poor stars and the standards used for classification are presented. The extension of spectral classification beyond M to L and T and spectroscopic classification criteria relevant to these classes are described. Contemporary methods of classifications based upon different automated approaches are introduced. ", "introduction": "\\label{sec:1} \\label{1.Historical Account } In 1866, Fr Angelo Secchi a Jesuit astronomer working in Italy observed prismatic spectra of about 4000 stars visually and divided stars in four broad spectral classes using common absorption features of hydrogen. During 1886-97, Henry Draper Memorial Survey at Harvard carried out a systematic photographic spectroscopy of stars brighter than 9th magnitude covering entire sky using telescopes at Harvard and Arequipa, Peru under the leadership of E.C.Pickering. The Henry Draper Catalog was published in 9 volumes of the Annals of Harvard College Observatory between 1918 and 1924. It contains rough positions, magnitudes, spectral classifications for 225,300 stars. Earlier work by W. Fleming essentially subdivided the previously used Secchi classes (I to IV) into more specific classes, by giving letters from A to N. The strength of hydrogen lines being the main classifier, the spectral type A was assigned to stars with strongest hydrogen lines followed by B,C with weaker hydrogen lines. This system was found to be unsatisfactory since the line strengths of other lines varied irregularly and so did the B-V color. This system was improvised by A.Maury, A.J.Cannon and E.Pickering who re-arranged the spectral sequence taking into consideration the changes in other lines and this new spectral sequence was also a sequence according to the color of the stars. But well-known stars had been already assigned the older spectral classes for long hence it was not possible to change them. We, therefore, have a spectral sequence essentially temperature dependent but goes like OBAFGKM. At the cool end the classification becomes more complex with parallel branches of R,N,S stars. While the M stars have TiO bands, S stars display ZrO bands, while R and N are carbon stars showing strong bands of various molecules with carbon. These have more recently been merged into a unified carbon classifier C scheme of C1, C2 etc with the old N0 starting at roughly C6. Another subset of cool carbon stars are the J-type stars, which are characterized by the strong $^{13}$CN molecules in addition to those of $^{12}$CN. Each of the above mentioned spectral classes OBAFGKM have been subdivided into ten subclasses e.g A0, A1 ... A9. \\subsection{Luminosity Effects in Stellar Spectra} E. Hertzsprung suggested in 1905 that spectral line widths were related to the luminosity of the stars. He pointed out that at a given apparent magnitude, the low proper motion stars would be at larger distance from us than the high proper motion stars of the same apparent magnitude and hence of higher intrinsic luminosity. These low proper motion stars were found to be exhibit narrower spectral lines so Hertzsprung concluded that these narrow line stars have larger intrinsic luminosity than the broad line stars. In 1943, W.Morgan, P.C. Keenan and E.Kellman introduced luminosity as second classification parameter. Morgan noticed the near constancy of the gravity along the main sequence in HR diagram and luminosity class parameter was an attempt to identify stars of different gravities and hence radii at nearly constant temperature. The above mentioned system also known as Yerkes Spectral Classification. Within the system, six luminosity classes are defined on the basis of standard stars over the observed luminosity range. The Six classes are: Ia: most luminous supergiants Ib: less luminous supergiants II: luminous giants III: normal giants IV: subgiants V: main sequence stars The main sequence class (dwarfs) are the stars at the main sequence, sustaining themselves through the conversion of hydrogen to helium by nuclear fusion in the stellar core, giant is a post main sequence star which is no longer burning hydrogen at the core but has H-burning shell outside the core. Giants as well as supergiant have comparable mass to dwarfs but have expanded to a much larger radii resulting in a decrease in their surface gravities. The spectral lines appear broad in the dense atmospheres of dwarfs primarily due to pressure broadening and Stark broadening while the same line would appear narrow in the low gravity atmospheres of supergiants. The luminosity effects are not restricted to the narrowing of strong lines. The line strengths and ratios of line strengths of neutral and ionized species also show remarkable variations over spectral classes and luminosity types and have been used for defining the subclasses and luminosity types. In addition, there are well-known luminosity indicators such as the emission components in the lines of CaII H and K in late type stars which are related to the luminosity (absolute magnitude) of the stars and the calibration of this relationship has been carried out by O. Wilison and M.K.V.Bappu in 1957. The strength of near IR OI triplet at 7771-75$\\AA$ has been used by Osmer (1972), Arellano Ferro, Giridhar and Goswami (1991) Arellano Ferro, Giridhar and Rojo Arellano (2003) for absolute magnitude calibration of A-G stars. In the next section, we will describe the line strengths and their ratios which are used to define the spectral classes and luminosity types. \\begin{figure} \\centering \\includegraphics[height=8cm]{figure4.eps} \\caption{ The luminosity effect in late O stars from Gray 2003 with author's permission} \\label{fig:1} % \\end{figure} \\clearpage \\section {Classification Criteria for various spectral types } \\label{sec:2} \\label{2.Classification Criteria for various spectral types } Beginning from O stars which are hottest, we will briefly describe the spectral characteristic for each spectral class and also list the spectral type and luminosity class indicators. Most of the information for this tutorial section is taken from Jascheck and Jascheck (1987). The figures used in this section illustrating the different spectral types and also the luminosity effects at various spectral types are taken from \"A Digital Spectral Classification Atlas v1.02 \" by R.O.Gray with his permission. More figures can be found on the website http://nedwww.ipac.caltech.edu/level5/Gray/frames.html. \\subsubsection{ O-type Stars } These are hot stars with temperature range of 28,000K to 50,000K. These stars exhibit strong lines of neutral and ionized helium. The strength of HeII lines decrease and HeI as well HI lines increasei in strength towards later O-types. Over the spectral range O4 to B0 the line strength of HeII line at 4541\\AA~ decreases from 800m\\AA~ to 200m\\AA~, while HeI line at 4471\\AA~ increases from 100m\\AA~ to 1000m\\AA~ and that of HI line at 4341\\AA~ increases from 1.5\\AA~ to 2.5\\AA~. The O-type spectra also exhibit the features of doubly and triply ionized carbon, nitrogen, oxygen and silicon. The line ratios such as CIII$\\lambda$ 4649/ HeII$\\lambda$ 4686 are used for luminosity classification. Similarly, the ratio of Si IV $\\lambda$ / HeI$\\lambda$4143 serves as good luminosity indicator in late O type stars. Wolf-Royet stars are a special family of O-type stars that are characterized by broad emission lines of ionized helium, carbon (WC sequence) or nitrogen (WN sequence). The WC stars exhibit emission lines of HeII such as HeII$\\lambda$ 4686, ionised carbon such as CII$\\lambda$4267, CIII$\\lambda$ 3609,4187,4325,4650 etc CIV$\\lambda$ 4441,4658,4758 etc and lines of OII,OIII, OIV and OV. These star are subdivided into WC classes from WC2-WC10 based upon the line ratios of CIII$\\lambda$5696/OV $\\lambda$5592, CIII$\\lambda$5696 / CIV$\\lambda$5805 etc. The WN stars exhibit emission lines of HeII , NIII$\\lambda$4097, 4640,5314, NIV$\\lambda$ 3483,4057 and NV $\\lambda$4605,4622. These stars are also subdivided into WN subclasses using the line ratios such as NIII$\\lambda$4640 / HeII$\\lambda$4686. \\begin{figure} \\centering \\includegraphics[height=8cm]{figure9.eps} \\caption{ The luminosity effect in B stars is illustrated above. The figure is from Gray 2003 with author's permission } \\label{fig:2} % \\end{figure} \\subsubsection{ B-type Stars } The B type spectra contains lines of HeI, HI, CII, CIII, NII, NIII, OII, SiII, SiIII, SiIV, MgII. The lines of higher ionized states of C,N,O are present in early B stars. The maximum strength of HeI line reaches near B2. Many B stars are fast rotators and a emission lines are present in some of them. The line ratios SiIII / SiIV , SiII$\\lambda$ 4128-30 / HeI$\\lambda$ 4121 and SiII$\\lambda$ 4128-30 / HeI$\\lambda$ 4144 are used for spectral class determination. The luminosity criteria used include the ratios of features at $\\lambda$4119 (SiIV+HeII) / $\\lambda$4144 (HeI), $\\lambda$4481 (MgII) / $\\lambda$4471 (HeI) which increase with luminosity. Profiles of Balmer lines become narrower with luminosity. These are blue white stars with temperature range of 10,000K to 28,000K. \\begin{figure} \\centering \\includegraphics[height=8cm]{figure16.eps} \\caption{ The luminosity effect in A5 to G0 stars are illustrated the figure is taken from Gray 2003 with author's permission } \\label{fig:3} % \\end{figure} \\subsubsection{ A-type Stars } These are white stars with temperature range of 7,500K to 10,000K. The A type stars exhibit strong hydrogen lines of Balmer series. The hydrogen lines are strongest at spectral type A2. The H$_\\gamma$ has a strength of 13.6$\\AA~$ at A0, 17$\\AA~$ at A2 and decreases to 13 $\\AA~$ at A7 and to 8$\\AA~$ at F0. The similar pattern is followed by other lines of Balmer series. The metallic lines gradually increase in strength from A0 to A9. The helium lines are absent. To assign subclasses the line ratios CaI$\\lambda$4227 / MgII $\\lambda$4481, FeI $\\lambda$ 4045 / FeII $\\lambda$4173, MgII $\\lambda$4481 / FeI $\\lambda$4485 MgII $\\lambda$4481/ FeII $\\lambda$4416 are useful. However other line ratios are also used. The luminosity criteria used are blend ratios such as FeI,FeII $\\lambda$4383-85 / MgII $\\lambda$4481, FeII $\\lambda$4417 / MgII $\\lambda$4481 which become stronger towards higher luminosity. The hydrogen lines also become narrower towards higher luminosity. The additional luminosity indicators are SrII $\\lambda$ 4215 / CaI $\\lambda$ 4226, FeII $\\lambda$4351 / MgII $\\lambda$4481 which increase at higher luminosity. The near infrared OI triplet at $\\lambda$ 7771-75 is a very good indicator of luminosity for A-F stars. Additional sub-classification of A type stars such as Am. Ap, is done based on their chemical peculiarities, magnetic fields and rotation and the presence of emission lines in their spectra. \\begin{figure} \\centering \\includegraphics[height=8cm]{figure17.eps} \\caption{ The luminosity effect in F stars is illustrated, the figure is taken from Gray 2003 with author's permission } \\label{fig:4} % \\end{figure} \\subsubsection{ F-type Stars } These are white-yellow stars with temperature range of 6,000K to 7,500K. The F type spectra have large number of metallic lines, Ca II H and K lines are very strong which become stronger than hydrogen lines of Balmer series. The hydrogen lines although present, are on decline from F0- F9. The G band of CH molecule makes its appearance around F3. The Ca II at $\\lambda$ 3933 increases from 6.5$\\AA~$ at F0 to 17.0$\\AA~$ at G0, CaI at $\\lambda$ 4226 increases from 0.25$\\AA~$ at F0 to 1.1$\\AA~$ at G0 while H$_{\\gamma}$ decreases from 8 $\\AA~$ at F0 to 3$\\AA~$ at G0. Due to large number of lines present, at the resolution used for classification, most feature are blended. For the spectral class the criteria used are FeI$\\lambda$ 4045 / H$_{\\delta}$, CaI $\\lambda$ 4226 / H$_{\\gamma}$, MnI $\\lambda$ 4030-34 / Si II $\\lambda$ 4128-32 etc. The CaII lines show positive luminosity effect, ratios Ti II $\\lambda$ 4444 / MgII $\\lambda$4481, SrII $\\lambda$ 4077 / FeI$\\lambda$4045 , SrII$\\lambda$4077 / H$_{\\delta}$ are used for luminosity classification. \\subsubsection{ G-type Stars } These are yellow stars with temperature range of 4,900K to 6,000K. In G type stars the hydrogen lines are further weakened and become comparable to the strength of metal lines. Metal lines are stronger and more numerous towards later G, molecular bands of CH and CN become visible. The spectral types are obtained by taking the ratio of metal lines with those of hydrogen lines e.g. using the raios such as FeI $\\lambda$ 4384/ H$_{\\gamma}$, FeI $\\lambda$ 4143 / H$_{\\delta}$, FeI $\\lambda$ 4045 / H$_{\\delta}$, Ca I $\\lambda$ 4226 / H$_{\\delta}$ etc. For Spectral types later than G5 the CaI $\\lambda$ 4226 / H$_{\\delta}$ can be used. The line ratios such as CrI$\\lambda$ 4254 / FeI $\\lambda$ 4250 or CrI$\\lambda$ 4742 / FeI $\\lambda$ 4271 are recommended for stars showing compositional anomalies such as weak metal line stars or weak G band stars. Luminosity effects at low dispersion can be seen through CN bands. The ratio of the SrI+FeI blend at $\\lambda$ 4216 / Ca I $\\lambda$ 4226 is known to be luminoity sensitive. Ratios of YII+FeI at $\\lambda$ 4376 / FeI $\\lambda$4383, SrII $\\lambda$ 4077/ H$_{\\delta}$ is also known to be luminosity sensitive but will not be suitable for stars with anomalous s-process abundances. MgI triplet at 5167-83 $\\lambda$ are luminosity sensitive for spectral type range G8-K5. \\begin{figure} \\centering \\includegraphics[height=8cm]{figure21.eps} \\caption{ The luminosity effect in G stars is illustrated the figure from Gray 2003 with author's permission } \\label{fig:5} % \\end{figure} \\subsubsection{ K-type Stars } These are orange stars with temperature range of 3,500K to 4,900K. In these stars the hydrogen lines are very weak but strong numerous metal lines are seen. The CaII lines are very strong and CH molecular band like G band becomes very strong. In late K stars TiO and VO bands are also seen. The line ratios for spectral type used are CrI$\\lambda$ 4254 / FeI $\\lambda$ 4250, CrI$\\lambda$ 4254 / FeI $\\lambda$ 4260, CrI$\\lambda$ 4254 / FeI $\\lambda$ 4271. Additional qualifiers are TiI$\\lambda$ 3999 / FeI $\\lambda$ 4005, FeI$\\lambda$ 4144 / H$_{\\delta}$, CaI $\\lambda$ 4226 / FeI $\\lambda$ 4250. The TiO band becomes visible at K7 and MgH at K5. The CN band increases with luminosity, so does H$\\lambda$ 4101 / FeI $\\lambda$ 4071. SrII $\\lambda$ 4077 / FeI $\\lambda$ 4063, TiII $\\lambda$ /4400,08 /FeI $\\lambda$ 4405 also increase with luminosity. \\subsubsection{ Carbon Stars } A carbon star is a late type giant with strong bands of carbon compound but no metallic oxide band. In their spectra, very intense bands of C$_{2}$, CN and CH are present but no bands of TiO, VO are seen. The carbon stars as a group have been studied by Secchi(1868) although in older classification they were given two different types R and N. The R type stars were similar to late G and early K but exhibited very strong Swan C$_{2}$ band around $\\lambda$ 4700 and $\\lambda$4395 which is as strong as G-band of CH. In N stars the the Swan band is so strong that the spectrum appears chopped up into the section of different intensities. Although based upon CN band (at $\\lambda$ 4216 and 3833) strengths a sequence R0,R1 to R10=N0 and subsequently upto N7 can be defined, but the same pattern is not shown by Swan C$_{2}$ bands, which become strongest at R5 and weakest at N0 and further strengthening towards later N types. A more detailed study by Shane(1928) revealed that temperature variations over the subclasses are not large and branching of these stars R and N stars is caused by abundance difference in C and O. If oxygen is more abundant than carbon then the spectrum is dominated by oxides like TiO (M stars). Keenan and Morgan (1941) used the term \"carbon stars\" instead of R and N wherein the overabundance of carbon varied from the star to star. They established the sequence of carbon stars by means temperature index of CrI$\\lambda$ 4254 / FeI $\\lambda$ 4250 and the strength of resonance lines of NaI at $\\lambda$ 5890-96 which are good temperature indicators for the late spectral type stars. The temperature sequence C0 to C7 covers full range of R0 to N7 which in temperature is very similar to the sequence of G4-M4. J stars are another group of carbon stars characterized by unusually strong isotopic bands of carbon which implies very low C$^{12}$/C$^{13}$ ratio. \\begin{figure} \\centering \\includegraphics[height=8cm]{figure24.eps} \\caption{ The luminosity effect in M2 stars is illustrated the figure is from Gray 2003 with author's permission } \\label{fig:6} % \\end{figure} \\clearpage \\subsubsection{ M-type Stars } These are red stars with temperature range of 2,000K to 3,500K. The spectra of M stars are dominated by strong bands of TiO, VO, LaO. The bands used for classification upto M2 are those of TiO at$\\lambda$ 4584, 4761, 4954 5448. After M3 these bands saturate, those at $\\lambda$ 5759, 5810 saturate after M5. The VO bands at $\\lambda$ 5737, 7373, 7865, 7896 become conspicuous after M7. The luminosity effect in early M type stars can be seen through decrease of CaI $\\lambda$ 4226 at higher luminosity; a similar negative luminosity dependence is exhibited by Cr I feature at $\\lambda$ 4254-74-90. Ratios SrII $\\lambda$ 4077 / FeI $\\lambda$ 4263 and (YII +FeI )$\\lambda$ 4376 / FeI $\\lambda$ 4386 increase with luminosity. The KI $\\lambda$ 7699, NaI $\\lambda$ 8183 and $\\lambda$ 8185 decrease in intensity from dwarfs to giants. The Ca II triplet at $\\lambda$ 8498, 8542 and 8662 is a very important luminosity indicator in these stars. It is very weak in dwarfs but becomes very strong in giants and supergiants. \\subsubsection{ S-type Stars } These stars are similar in temperature to K5 to M stars but exhibit strong bands of ZrO at with band heads at $\\lambda$ 4640, 5551 and 6474. ", "conclusions": "" }, "1003/1003.3308_arXiv.txt": { "abstract": "We study in detail the motions of three planets interacting with each other under the influence of a central star. It is known that the system with more than two planets becomes unstable after remaining quasi-stable for long times, leading to highly eccentric orbital motions or ejections of some of the planets. In this paper, we are concerned with the underlying physics for this quasi-stability as well as the subsequent instability and advocate the so-called \"stagnant motion\" in the phase space, which has been explored in the field of dynamical system. We employ the Lyapunov exponent, the power spectra of orbital elements and the distribution of the durations of quasi-stable motions to analyze the phase space structure of the three-planet system, the simplest and hopefully representative one that shows the instability. We find from the Lyapunov exponent that the system is almost non-chaotic in the initial quasi-stable state whereas it becomes intermittently chaotic thereafter. The non-chaotic motions produce the horizontal dense band in the action-angle plot whereas the voids correspond to the chaotic motions. We obtain power laws for the power spectra of orbital eccentricities. Power-law distributions are also found for the durations of quasi-stable states. All these results combined together, we may reach the following picture: the phase space consists of the so-called KAM tori surrounded by satellite tori and imbedded in the chaotic sea. The satellite tori have a self-similar distribution and are responsible for the scale-free power-law distributions of the duration times. The system is trapped around one of the KAM torus and the satellites for a long time (the stagnant motion) and moves to another KAM torus with its own satellites from time to time, corresponding to the intermittent chaotic behaviors. ", "introduction": "More than $300$ exoplanets have been discovered so far \\footnote{http://exoplanet.eu for the latest information.} and, interestingly, some of them have a quite different appearance from that of our solar system. The existence of so-called \"eccentric planets\", that is, planets with high orbital eccentricities, for example, attracts attentions of many researchers \\citep{obs}. Stimulated by these observations, the planetary formation theory including the origin of the eccentric planets has made substantial progress over the years \\citep{kki06}. In the standard theory, terrestrial planets are thought to be formed through giant impacts of proto-planets (planets with sub-Earth masses) in crossing orbits \\citep{cw98}. Since N-body simulations suggest that proto-planets are formed in nearly circular orbits separated by several Hill radii \\citep{ki95}, some destabilizing processes are expected to operate and make proto-planets originally in the well-separated circular orbits collide with each other and grow up to terrestrial planets. The stability of the system with two planets around a central star has been thoroughly investigated in celestial mechanics and it is known that there exists a critical orbital separation between the planets, beyond which the planets never experience close encounters and the system remains stable forever \\citep[]{mb80,g93}. The situation changes drastically, however, if another planet is added to the system. Using numerical simulations, \\citet{cwb96} demonstrated that the systems with more than two planets become unstable even for large orbital separations. Although the planetary motions remain regular for some time at first, one of the planets eventually comes close enough to another (that is, within the Hill radius of the latter), leading to subsequent orbital crossings. This is a good news for the terrestrial formation theory and may also account for the formation of the eccentric planets. In fact, \\citet{mw02} numerically integrated the motions of three Jupiter-mass planets and found in most of their simulations that one of the planets is ejected from the system and the others are left in the system with high eccentricities. Several attempts \\citep[]{jt07,fr07} have been made to reproduce the observed eccentricity distribution by the orbital instability. Their results seem to be consistent with the observations although the latter itself may be somewhat biased \\citep{st08}. The planetary motions in these systems are interesting in their own right. As mentioned above, we commonly observe a long period of quasi-regular motions that look like independent Keplerian motions before the eventual orbital crossings. The switch is sudden and quick. The duration of the quasi-regular motions is sensitive to the initial orbital separations and eccentricities \\citep[]{cwb96,ykm99}. Using numerical simulations and simplified analytical models, \\citet{zls07} claimed that the gradual deviation from the Keplerian motions can be regarded as a random walk process. These efforts notwithstanding, the underlying physics behind the phenomena such as the long period of quasi-regular motions followed by the sudden transition to chaotic states is remaining to be revealed and is the main concern of this paper. We attempt to understand this phenomenon as the so-called ``stagnant motion'' in the phase space, which will be described below. We pay attention to a similar phenomenon known in the field of dynamical system. In nearly integrable systems, a sudden transition from a regular motion sustained for a long time to chaotic motions is often observed. In the nonlinear lattice problems, for example, \\citet{hs68} found that an initially imposed normal mode experiences sudden energy exchanges among several other modes after long regular oscillations. They called this \"the induction phenomenon\" and referred to the duration of the regular motion as \"the induction period\" \\citep{soah70}. \\citet{aizawa} constructed a so-called \"stagnant motion\" model for this induction phenomenon. According to the KAM theorem \\citep[]{kolmogorov,moser,arnold}, the phase space of a nearly integrable system retains tori, which exist in the integrable system, if perturbations to the integrable system are sufficiently weak. It is generally expected that the so-called KAM tori will survive even for not so small perturbations. In the stagnant motion model, it is assumed that the KAM torus exists in \"the chaotic sea\", the region corresponding to chaotic motions of the system, being surrounded by a thin layer called \"the stagnant layer\", in which smaller tori are distributed in a self-similar manner (see Fig.~\\ref{schematic1}). The system shows nearly regular behavior when the phase space orbit is trapped in the stagnant layer whereas the sudden transition to chaotic motions occur when the orbit escapes out of the layer. This model is successful in reproducing the scale-free power spectra and the distribution of the induction period. In this paper we show some evidence to support the interpretation of the motions of three-planet system as one of the induction phenomena and attempt to understand them in the frame work of the stagnant motion model. The organization of the paper is as follows. We summarize the numerical models in section 2. In section 3, we describe the methods of analysis. The results are presented in section 4 and the summary and discussions are given in section 5. ", "conclusions": "In this paper we have numerically investigated the dynamics of the three-planet system and inferred its phase space structure from the obtained Lyapunov exponents, power spectra of orbital elements and distributions of induction periods based on the stagnant motion model. What we have found are the followings: 1. The global and local Lyapunov exponents show that the system is almost non-chaotic until the first close encounter between two planets and it then turns into chaotic motions with intermittent non-chaotic periods. This suggests that the phase space consists of KAM tori surrounded by the stagnant layer and immersed in the chaotic sea. In fact, the dense bands are formed in the action-angle plot, corresponding to the quasi-regular motions. The lines composing a band are undulating with frequencies that obey a power-law and this may represent the motions around the satellite tori in the stagnant layer, which have a self-similar distribution. The phase space orbit goes from one system of KAM torus and stagnant layer to another through the chaotic sea. 2. The power spectra of the orbital eccentricities of planets can be approximated by the power law, $\\propto 1/f^{\\nu}$, in general. Such power-law spectra are known to be one of the characteristic features of the stagnant motions although the power-law index is not predicted by the theory. In our models the power-law index is $\\nu \\simeq 1$ for the pre-encounter phase whereas it becomes $\\nu \\simeq 2$ after the encounter. The spectrum in the post-encounter phase is similar to that of the Brownian motions or the random walks. On the other hand, the spectrum in the pre-encounter phase might be originated from the fractional Brownian motions. 3. The distributions of the duration of the pre-encounter phase that was referred to as the induction periods obey the power law in the long-duration part. The power-law indices are substantially different between models. It is stressed that the stagnant motion model predicts the power law for the distribution of the induction periods as a consequence of the self-similar distribution of smaller tori in the stagnant layer around a KAM torus. The distributions are deviated from the power law in the short-duration part and has a peak in between. Connecting the peaks for various models with different initial orbital separations, we have obtained the relation similar to what \\citet{cwb96} found. It is also shown that the duration of the pre-encounter phase has actually a considerably broad distribution. 4. For the Jupiter system, the distribution of the time from the first encounter to the ejection of a planet from the system also obeys a power law, which was not expected initially. From the analogy to the stagnant motion model, we might be able to infer the phase space structure as shown schematically in Fig.~\\ref{schematic2}: Many KAM tori with its own stagnant layer and satellite tori in it are distributed self-similarly in the chaotic sea. Although we expect the phase space structure depicted in Fig.~\\ref{schematic2} is true both of the proto-planet system and the Jupiter system, the difference between them should be also mentioned. In general, the number of KAM tori in the phase space becomes smaller and the stagnant layers around them get thinner as the perturbation to the integrable system is greater. In the system of our concern, the perturbation is the interactions among the planets and hence it is larger for more massive planets. The pre-encounter phase is an exception, though. In this phase, the planets have nearly circular orbits separated by several Hill radii. Then the strength of the interactions between the planets depends only weakly on the planetary mass thanks to the definition of the Hill radius given in Eq.~(\\ref{eq:hill}). These facts suggest that the KAM torus corresponding to the initial regular motion and its stagnant layer are robust and similar for the proto-planet and Jupiter systems while the number of other KAM tori is smaller and their stagnant layers are thinner for the Jupiter system than for the proto-planet system. This difference in the phase space structures is supposed to be responsible for the observed difference in the orbital evolutions of the two systems: one of the planets is ejected in short times for the Jupiter system whereas no ejection occurs at least for $\\sim 10^7$ years in the proto-planet system. The results obtained in this paper appear to be consistent with our interpretation that the dynamics of three-planet system is a stagnant motion at least in the pre-encounter phase. It is also suggested from the results for the Jupiter system that even the post-encounter phase may be described by some extension of the stagnant motion model. It is true, however, that a more direct capture of satellite tori in the phase space is certainly desirable. Although we have attempted to do this with the so-called Poincare mapping, but in vain so far. We are afraid that the degree of freedom of our system is just too large to find an appropriate two dimensional section in the 12-dimensional phase space. Maybe other approaches should be pursued in the future work. In so doing, the number of realizations should be increased and other initial settings should be tried. Not to mention, we are also interested in how the results will change as the number of planets are varied." }, "1003/1003.4691_arXiv.txt": { "abstract": "Compared with ordinary spirals, the ISM in ring galaxies experiences markedly different physical conditions and evolution. As a result, ring galaxies provide interesting perspectives on the triggering/quenching of large scale star formation and the destructive effects of massive stars on molecular cloud complexes. We use high resolution radio, sub-mm, infrared, and optical data to investigate the role of gravitational stability in star formation regulation, factors influencing the ISM's molecular fraction, and evidence of peculiar star formation laws and efficiencies in two highly evolved ring galaxies: Cartwheel and the Lindsay-Shapley ring. ", "introduction": "Ring galaxies are dramatic examples of galaxy transformation caused by a remarkably simple interaction. Observations and numerical models (\\citealp{lyndstoomre}) argue persuasively that these objects are formed by the near central passage of a companion through a spiral along the rotation axis. The brief additional gravitational force induces epicyclic motions throughout the disk, which act to form radially propagating orbit-crowded rings of gas and stars. The concentration of the ISM into the expanding ring (at the disk's expense) is nearly total and can last for $\\approx400$ Myrs. It is this radical rearrangement of the spiral's ISM that is responsible for their interesting star forming properties. ", "conclusions": "The distribution and gravitational stability of a ring galaxy's ISM changes dramatically as it evolves. Further, the long confinement in the dense star forming ring is expected result in a fragmented and largely atomic ISM, which may explain the peculiar star formation laws, efficiencies, and $f_{\\rm mol}$ observed in the two largest and most evolved ring galaxies, Cartwheel and {\\em L-S}. Future progress will require sensitive and high resolution assays of the molecular ISM in these and other ring galaxies, which will be possible with ALMA and CARMA." }, "1003/1003.1282_arXiv.txt": { "abstract": "This presentation reviews Chandra's major contribution to the understanding of nearby galaxies. After a brief summary on significant advances in characterizing various types of discrete X-ray sources, the presentation focuses on the global hot gas in and around galaxies, especially normal ones like our own. The hot gas is a product of stellar and AGN feedback --- the least understood part in theories of galaxy formation and evolution. Chandra observations have led to the first characterization of the spatial, thermal, chemical, and kinetic properties of the gas in our Galaxy. The gas is concentrated around the Galactic bulge and disk on scales of a few kpc. The column density of chemically-enriched hot gas on larger scales is at least an order magnitude smaller, indicating that it may not account for the bulk of the missing baryon matter predicted for the Galactic halo according to the standard cosmology. Similar results have also been obtained for other nearby galaxies. The X-ray emission from hot gas is well correlated with the star formation rate and stellar mass, indicating that the heating is primarily due to the stellar feedback. However, the observed X-ray luminosity of the gas is typically less than a few percent of the feedback energy. Thus the bulk of the feedback (including injected heavy elements) is likely lost in galaxy-wide outflows. The results are compared with simulations of the feedback to infer its dynamics and interplay with the circum-galactic medium, hence the evolution of galaxies. ", "introduction": "X-ray observations are playing an increasingly important role in the study of galaxies. With its arcsecond spatial resolution, \\chandra\\ in particular has made a significant impact on our understanding of discrete X-ray source populations, which mostly represent various stellar end-products [e.g., low- and high-mass X-ray binaries (LMXBs and HMXBs) and supernova remnants (SNRs)] as well as active galactic nuclei (AGNs). The resolution also allows for a clean excision of such sources from the data so low-surface brightness emission (e.g. from diffuse hot gas) can be mapped out. An under-appreciated aspect of \\chandra\\ is its spectroscopic capability in the study of diffuse hot gas when the low- and high-energy grating instruments are used. Although the sensitivities of the instruments are quite limited, the existing observations of a dozen or so bright objects (AGNs and LMXBs) have yielded data of high enough quality for unprecedented X-ray absorption line spectroscopic measurements of the global hot gas in and around our Galaxy. Useful constraints have also been obtained on the overall content of hot gas around galaxies within certain impact distances of the sight lines toward the AGNs. These measurements, compared with physical models and simulations of the hot gas, are shedding important insights on its relationship to the feedback from stars and AGNs. These aspects of the \\chandra's legacy are reviewed in the following. ", "conclusions": "As shown above, the diffuse soft X-ray emission observed in and around galaxies clearly traces the global hot gas that is heated by the stellar and possibly AGN feedback. By comparing the X-ray observations directly with physical models and/or simulations of the feedback, we may further study its dynamics, which is so far hardly known. The following discussion is concentrated on the modeling of the feedback in galactic spheroids for their relative simplicity. The presence of substantial cool gas and star formation, as in typical disk galaxies, would certainly add complications. \\subsection{Comparison of feedback models with X-ray observations} The missing feedback problem is particularly acute in so-called low $L_X/L_K$ spheroid-dominated galaxies. Empirically, the Ia SN rate is $\\approx 0.019 {\\rm~SN~yr^{-1}}) 10^{- 0.42(B-K)} [L_K/(10^{10}L_{\\odot K})]$, where $ L_{K}$ is the K-band luminosity of a galaxy \\cite{man06}. Adopting the color index $B-K \\approx 4$ for a typical spheroid, we estimate the total mechanical energy injection from Ia SNe (assuming $10^{51} {\\rm~ergs}$ each) is $L_{Ia} \\approx (1.3 \\times 10^{40} {\\rm~ergs~s^{-1}}) [L_K/(10^{10}L_{\\odot K})]$. The mechanical energy release of an AGN can also be estimated empirically \\cite{bes06, dav06} as $L_{AGN} = (1.1 \\times 10^{39} {\\rm~erg~s^{-1}}) [L_K/(10^{10} L_{K,\\odot})]^2.$ Therefore, averaged over the time, Ia SNe are energetically more important than the AGN in a galactic spheroid with $L_K \\lesssim 10^{11}L_{K,\\odot}$. Furthermore, SN blastwaves provide a natural distributed heating mechanism for hot gas in galactic spheroids \\cite{tw05}. Additional continuous heating is expected from converting the kinetic energy of stellar mass loss from randomly moving evolved stars to the thermal energy. On the other hand, the AGN feedback, likely occurring in bursts with certain preferential directions (e.g., in form of jets), can occasionally result in significant disturbances in global hot gas distributions, as reflected by the asymmetric X-ray morphologies observed in some elliptical galaxies (e.g., \\cite{die08}). It is not yet clear as to what fraction of the AGN feedback energy is converted into the heating of the hot gas observed. In general, the lack of distributed cool gas in galactic spheroids makes it difficult to convert and release the thermal energy into radiation in wavelength bands other than the X-ray. Not only the mechanical energy, the gas mass from the stars and Ia SNe, especially heavy elements, is missing as well. The mass injection rate is expected to be $0.026[L_K/(10^{10} L_{K,\\odot})] M_\\odot {\\rm~yr^{-1}}$ with a mean iron abundance $Z_{Fe} \\approx Z_{*,Fe}+ 4(M_{Fe}$$/0.7$$ M_\\odot)$, where $Z_{*,Fe}$ is the iron abundance of the stars while $M_{Fe}$ is the iron mass yield per Ia SN (e.g., \\cite{man06,kna92,cio91}). The above empirical estimates of the energy and mass feedback rates, which should be accurate within a factor of $\\sim 2$, are typically a factor of $\\sim 10^2$ greater than what are inferred from the diffuse X-ray emission. Naturally, one would expect that the missing feedback is gone with a wind (outflow), spheroid-wide or even galaxy-wide. The notion that Ia SNe may drive galactic winds is not new (e.g., \\cite{cio91, mb71, bre80}). But could such winds explain the diffuse X-ray emission of galactic spheroids? It is easy to construct a 1-D steady-state supersonic wind model, assuming that the specific energy of the feedback (per mass) of a galaxy is large enough to overcome its gravitational bounding and that the CGM pressure (thermal or ram) is negligible (e.g., \\cite{tan09b}). This supersonic wind model depends primarily on two feedback parameters: the integrated energy and mass input rates. However, the model in general fails miserably: it predicts a too low luminosity (by a factor of $\\sim 10^2$), a too high temperature (a factor of a few), and a too steep radial intensity profile to be consistent with \\chandra\\ observations of low $L_X/L_B$ galactic spheroids, in particular those in M31 and M104 \\cite{lw07,liz07}. Only few very low-mass and gas-poor spheroids, hence very faint in X-ray emission, still seem to be consistent with the 1-D wind model (e.g., \\cite{tri08}). \\begin{figure}[tbh] \\centerline{\\includegraphics[width=0.9\\textwidth,angle=0]{f4.jpg}} \\caption{ Sample snapshots of hydrodynamic simulations of the stellar feedback. {\\bf (a)} the 3-D simulated gas density distribution in an M31-like galactic spheroid. The slice is cut near the spheroid center and the units are in atoms~cm$^{-3}$, logarithmically. {\\bf (b)} Simulated 2-D large-scale gas density distribution in the $r-z$ plane (in units of kpc). A nearly vertical magnetic field, similar to what is observed in the inter-cloud medium of the Galactic center, is included in the simulation to test its confinement effect on the spheroid wind; a reverse-shock is clearly visible. Also apparent are instabilities at the contact discontinuity between the shocked spheroid wind gas and the ejected materials from the initial starburst. The density is scaled in units of M$_{\\odot} {\\rm~kpc^{-3}}$. } \\label{fig:f4} \\end{figure} To compare with the X-ray emission, in fact one needs to account for 3-D effects. X-ray emission is proportional to the emission measure and is thus sensitive to the detailed structure of hot gas in a galactic spheroid. To realistically generate the inhomogeneity in the heating and chemical enrichment processes, Tang \\& Wang have developed a scheme to embed adaptively selected 1-D SNR seeds in 3-D spheroid-wide simulations of supersonic winds or subsonic outflows (e.g., \\cite{tw05,tw09}). These 3-D simulations, reaching a resolution down to adaptively refined scales of a few pc (e.g., Fig.~\\ref{fig:f4}a), show several important 3-D effects \\cite{tan09b}: \\begin{itemize} \\item Soft X-ray emission arises primarily from relatively low temperature and low abundance gas shells associated with SN blastwaves. The inhomogeneity in the gas density and temperature substantially alters the spectral shape and leads to artificially lower metal abundances (by a factor of a few) in a spectral fit with a simplistic thermal plasma model. \\item Reverse shock-heated SN ejecta, driven by its large buoyancy, quickly reaches a substantially higher outward velocity than the ambient medium. The ejecta is gradually and dynamically mixed with the medium at large galactic radii and is also slowly diluted and cooled by {\\sl insitu} mass injection from evolved stars. These processes together naturally result in the observed positive gradient in the average radial iron abundance distribution of the hot gas, even if mass-weighted. \\item The average 1-D and 3-D simulations give substantially different radial temperature profiles; the inner temperature gradient in the 3-D simulation is positive, mimicking a ``cooling flow''. (Thus the study of nearby galactic spheroids may provide important insights into the behavior of the intragroup or intracluster gas around central galaxies, for which sporadic AGN energy injections somewhat mimic the SN heating considered here.). \\item The inhomogeneity also enhances the diffuse X-ray luminosity by a factor of a few in the supersonic wind case and more in a subsonic outflow. In addition, a subsonic outflow can have a rather flat radial X-ray intensity distribution. \\item The dynamics and hence the luminosity of a subsonic flow further depend on the spheroid star formation and feedback history. This dependence, or the resultant luminosity variance, may then explain the observed large dispersion in the X-ray to K-band luminosity ratios of elliptical galaxies. \\end{itemize} All considered, subsonic outflows appear to be most consistent with the X-ray observations of diffuse hot gas in typical intermediate-mass galactic spheroids and elliptical galaxies. \\subsection{The interplay between the feedback and galaxy evolution} Whether a galactic outflow is supersonic or subsonic depends not only on the specific energy of the ongoing feedback, but also on the properties of the CGM, which is a result of the past interplay between the feedback and the accretion of a galaxy or a group of galaxies from the intergalactic medium (IGM). Tang et al. have illustrated how this interplay may work, based on several 1-D hydrodynamic simulations in the context of galaxy formation and evolution \\cite{tan09a}. They approximate the feedback history as having two distinct phases: (1) an early starburst during the spheroid formation (e.g., as a result of rapid galaxy mergers) and (2) a subsequent long-lasting and slowly declining injection of mass and energy from evolved low-mass stars. An energetic outward blastwave is initiated by the starburst (including the quasar/AGN phase) and is sustained by the long-lasting stellar feedback. Even for a small galactic spheroid such as the one in the Milky Way, this blastwave may heat up the CGM on scales beyond the present virial radius, thus the gas accretion from the IGM into the galactic halo could be largely reduced (see also \\cite{kim09} for similar results from 3-D cosmological structure formation simulations). The long-lasting stellar feedback initially drives a galactic spheroid wind (Fig.~\\ref{fig:f4}b). As the mass and energy injection decreases with time, the feedback may evolve into a subsonic and quasi-stable outflow. This feedback/CGM interplay scenario provides a natural explanation to various observed phenomena: \\begin{itemize} \\item It solves the missing feedback problem as discussed earlier; the energy is consumed in maintaining the hot CGM and preventing it from fast cooling. \\item The very low density of the CGM explains the dearth of the chemically enriched hot gas observed around galaxies; much of the CGM or the intragroup gas has been pushed away to larger scales (see also \\cite{opp08}). \\item The predicted high temperature is consistent with observations of the large-scale hot intragroup medium, which seems always higher than $10^{6.5}$ K \\cite{sun09}. \\item The cooling of the material ejected early (e.g., during the initial starbursts; \\cite{tan09a}) provides a natural mechanism for high-velocity clouds with moderate metal abundances (Fig.~\\ref{fig:f4}b). Such clouds may be seen in 21 cm line emission and in various absorption lines (e.g., Ly$\\alpha$, Mg II, Si III; \\cite{sto06,ws09,shu09}). O VI line absorptions can then arise from either photo-ionization of the clouds or collisional ionization at their interfaces with the pervasive hot CGM. \\end{itemize}" }, "1003/1003.0976_arXiv.txt": { "abstract": "The overall classification of X-ray jets has clung to that prevalent in the radio: FRI vs. FRII (including quasars). Indeed, the common perception is that X-ray emission from FRI's is synchrotron emission whereas that from FRII's may be IC/CMB and/or synchrotron. Now that we have a sizable collection of sources with detected X-ray emission from jets and hotspots, it seems that a more unbiased study of these objects could yield additional insights on jets and their X-ray emission. The current contribution is a first step in the process of analyzing all of the relevant parameters for each detected component for the sources collected in the XJET website. This initial effort involves measuring the ratio of X-ray to radio fluxes and evaluating correlations with other jet parameters. For single zone synchrotron X-ray emission, we anticipate that larger values of fx/fr should correlate inversely with the average magnetic field strength (if the acceleration process is limited by loss time equals acceleration time). Beamed IC/CMB X-rays should produce larger values of fx/fr for smaller values of the angle between the jet direction and the line of sight but will also be affected by the low frequency radio spectral index. ", "introduction": "\\label{sec:prob} The paradigm for the classification of extragalactic jets is based on the Fanaroff-Riley (\\cite{fan:1974}) distinction between powerful FRII radio galaxies (and radio quasars) on the one hand, and low power FRI radio galaxies (and BL Lac objects) on the other. The (perhaps overly simplified) notion is that the jets of FRI sources are lossy, often displaying a number of brightness enhancements (a.k.a. \"knots\") whereas FRII jets in radio galaxies (\"RG\" hereafter) are often difficult to detect and deposit the bulk of their energy at large distances from the parent galaxy in the radio 'hotspots'. Generally speaking, FRI jets often terminate within the confines of their host galaxy (e.g. M87), or at least the well collimated section of the jets are of similar size as their host (e.g. 3C~31). Obviously this simple interpretation has a number of caveats; chief amongst them is the effect of relativistic beaming which means that we must make model-based assumptions to imagine what the jets of Cyg A would look like if one of them was close to our line of sight (l.o.s.). In the unification hypothesis (\\cite{urr:1995}) radio quasars and FRII RG are distinguished solely on the basis of their viewing angle. If we were to view the jet of Cyg A close to the l.o.s. would it be a knotty jet like that of 3C~273? Our primary interest in this problem arises from the question of the emission processes for the generation of X-rays from jets. The idea that kpc scale jets have substantial bulk Lorentz factors, $\\Gamma \\simgt 5-10$ that would engender copious production of X-rays from inverse Compton scattering of low energy relativistic electrons on photons of the cosmic microwave background (\"IC/CMB\", \\cite{tav:2000, cel:2001, hk:2002}) came to be the accepted explanation for X-ray emission from quasar and FRII RG jets. For FRI jets, synchrotron emission from very high energy electrons (i.e. electrons with Lorentz factors $\\gamma~\\geq~10^7$) is thought to provide a consistent suite of physical parameters and simple models were able to produce the observed spectral energy distributions (SED). Thus the classification of X-ray emitting jets tended to follow the FRI - FRII separation, with the additional notion that FRI jets would generally have the X-ray spectral index ($\\alpha_x$, defined by flux density, S$_{\\nu}\\propto\\nu^{-\\alpha}$) $>$1 whereas $\\alpha_x$ for quasars and FRII RG's would be $<$1. These values of $\\alpha_x$ are consistent with expectations that at the high end of the electron distribution (N($\\gamma$)), we may expect steep spectra because we may be observing close to an exponential cutoff and/or the initial N($\\gamma$) has been strongly affected at these energies by synchrotron and IC losses which go as $\\gamma^2$; and at the low end of N($\\gamma$), the emitted spectra (either synchrotron or IC) would be flatter. A number of doubts and counter examples have been raised about the IC/CMB model of X-ray emission of quasar and FRII jets. Amongst these are: (a) the inability to independently demonstrate that $\\Gamma\\geq$5 on kpc scales; (b) the assumption that one can extrapolate N($\\gamma$) from the radio regime ($\\gamma\\approx~10^4$) down to $\\gamma\\approx$~100 with a power law of slope -p with p=2$\\times\\alpha_{radio}$~+~1 (i.e. are there really enough low energy electrons?); and (c) the evidence that the X-ray jet of 3C~273 comes from a second synchrotron component rather than IC/CMB (\\cite{uch:2006}). These concerns have prompted us to attempt to devise a different classification scheme for X-ray jets, hoping thereby to achieve a better understanding of the role of IC and synchrotron emissions at X-ray frequencies. The present contribution reports on the first step of this project: an investigation of the ratio of X-ray to radio flux for X-ray detected knots and hotspots. To do this, we have embarked on a standard reduction scheme of all X-ray jets detected by the CXO (Chandra X-ray Observatory). The finding list for these sources comes from the XJET compilation\\footnote{http://hea-www.harvard.edu/XJET/}, see Massaro et al., this volume. \\begin{figure} \\centerline{\\includegraphics[height=8cm,width=0.55\\linewidth,origin=c,angle=-90]{ne.ps} \\includegraphics[height=8cm,width=0.65\\linewidth]{lx_tot.ps} } \\caption{Left: An idealized, broken power law electron spectrum. N1 represents the number of electrons contributing to IC/CMB X-rays; N2 the number contributing to the radio emission; and N3, the number responsible for synchrotron X-rays. Right: A scatter plot of R versus X-ray luminosity ($\\Omega_{\\Lambda}$=0.73, $\\Omega_M$=0.23, and $H_{\\rm 0} = 71$ km s$^{-1}$ Mpc$^{-1}$). FRI knots are black diamonds; FRII knots are blue filled triangles, and quasar knots are red filled circles. Hotspots are open triangles (blue FRII) and circles (red Q). The green line is not a fit to anything; it has a slope of +1 (see text).}\\label{fig:lx} \\end{figure} ", "conclusions": "From our preliminary analysis using the ratio of X-ray to radio fluxes, we have found that hotspots tend to have a lower ratio than knots. Features classified as 'quasar hotspots' generally have smaller values of R than quasar knots, and most FRII hotspots have smaller R values than FRI and quasar knots. Most FRI and quasar knots have R values between 1 and 100. We have failed to find any significant distinction between FRI and quasar knots except for the obvious difference in apparent x-ray luminosity. \\begin{theacknowledgments} We are indebted to a number of colleagues who have generously donated radio maps for our use, and for public access via the XJET webpage. Other radio maps were downloaded from the NRAO VLA Archive Survey. The work at SAO was partially supported by NASA grant AR6-7013X. FM acknowledges the Foundation BLANCEFLOR Boncompagni-Ludovisi, n'ee Bildt for the grant awarded him in 2009. \\end{theacknowledgments} \\small" }, "1003/1003.2446.txt": { "abstract": "We use the Spitzer Wide-area InfraRed Extragalactic Legacy Survey (SWIRE) to explore the specific star-formation activity of galaxies and their evolution near the peak of the cosmic far-infrared background at 70 and 160\\micron. We use a stacking analysis to determine the mean far-infrared properties of well defined subsets of galaxies at flux levels well below the far-infrared catalogue detection limits of SWIRE and other Spitzer surveys. We tabulate the contribution of different subsets of galaxies to the far-infrared background at 70$\\micron\\ $ and 160$\\micron$. These long wavelengths provide a good constraint on the bolometric obscured emission. The large area provides good constraints at low $z$ and in finer redshift bins than previous work. At all redshifts we find that the specific far-infrared luminosity decreases with increasing mass, following a trend $L_{\\rm FIR}/M_* \\propto M_* ^\\beta$ with $\\beta =-0.38\\pm0.14$. This is a more continuous change than expected from the \\cite{Delucia2007} semi-analytic model suggesting modifications to the feedback prescriptions. We see an increase in the specific far-infrared luminosity by about a factor of $\\sim100$ from $0$4.5$\\sigma$ during its first 11 months survey (Abdo et al. 2010, A10 hereafter). This is the 1FGL catalogue\\footnote{http://fermi.gsfc.nasa.gov/ssc/data/access/lat/1yr\\_catalog/} and contains both galactic and extragalactic $\\gamma$--ray sources. In a recent paper Abdo et al. (2010a, A10a hereafter) classified 831 out of 1451 1FGL sources as blazars. These form the 1LAC catalogue\\footnote{http://heasarc.gsfc.nasa.gov/W3Browse/fermi/fermilac.html}. Considering the 796 sources in the 1LAC associated with a single counterpart, there are 37 sources generically classified as AGN, 314 BL Lac and 285 Flat Spectrum Radio Quasars (FSRQ), classified according to their optical spectrum. The remaining 160 sources are candidate blazars but of ``unknown\" class because of lacking an optical spectrum or, if available, because of its poor quality for the optical classification. However, for these sources a counterpart could still be found in one of the radio catalogues adopted for the source identification by A10a (CRATES -- Haeley et al. 2007; CGRaBS -- Healey et al. 2008; the BZCat -- Massaro et al. 2009). Thirty five sources in the 1LAC sample have more than one associated counterpart. The 1FGL sample is equally populated in the northern (713 sources) and southern sky (738 sources). However, only 50\\% of the 1FGL sources in the south sky are classified as blazars in the 1LAC catalogue with respect to $\\sim$70\\% of the 1FGL northern sources which are classified as blazars in the 1LAC catalogue. The recently published Australia Telescope 20 GHz Survey (AT20G, Murphy at al. 2010) represents the largest catalogue of high frequency radio sources detected with the Australia Telescope Compact Array (ATCA) in a survey conducted from 2004 to 2008 and covering the whole sky south of declination 0$^{\\circ}$. It contains 5890 radio sources with flux at 20 GHz exceeding 40 mJy and it is complete at a level of 91\\% above 100 mJy beam$^{-1}$. For several sources the fluxes at 8 and 5 GHz are also measured. The AT20G survey, being conducted at relatively high radio frequencies (20 GHz), most likely detects the radio emission from the compact cores of AGN. Several sources in this survey have flat or inverted radio spectra between 5 and 20 GHz. We correlate the AT20G radio catalogue with the 1FGL in search for possible counterparts. This allows us to study the radio versus $\\gamma$--ray properties of these associations. A possible correlation between the radio and the $\\gamma$--ray emission in radio loud AGN can shed light on the physical link between the emission processes in these two energy bands. The $\\nu F_{\\nu}$ spectra of both FSRQ and BL Lac objects show two peaks which are widely interpreted as due to the synchrotron and inverse Compton emission, respectively. The Spectral Energy Distribution (SED) of blazars form a sequence (so--called ``blazar sequence\", e.g. Fossati et al. 1998) with the most powerful sources (FSRQs) having the synchrotron and Inverse Compton peaks in the IR and MeV energy range respectively, while the less powerful BL Lacs have the synchrotron and Inverse Compton peak in the optical/UV (or even X--ray) and GeV--TeV energy range respectively. The blazar sequence was built by dividing blazars into bins of radio luminosity, thought to be a proxy for the bolometric one, and establishes a link between the radio and the $\\gamma$--ray emission. Before \\fe, no conclusive claim could be made on the existence of a possible correlation between radio and $\\gamma$--ray properties for EGRET sources (e.g. Mucke et al. 1997). Taylor et al. (2007), by studying the radio properties of EGRET detected blazars with the VLBA Imaging and Polarimetry Survey (VIPS -- extending down to 85 mJy) found that at low radio fluxes the radio flux density does not directly correlate with the $\\gamma$--ray flux. Further investigation of this correlation was possible with the source list of the \\fe\\ first three months survey (LBAS, Abdo et al. 2009; Giroletti et al. 2010): while a correlation between the radio flux at 8.4 GHz and the peak flux (above 100 MeV) appears when considering BL Lacs and FSRQ together, the statistical evidence is only marginal for FSRQs (chance correlation probability of 8\\%), being more robust only for BL Lacs (0.5\\%). Such correlation analysis should consider the different redshift range spanned by these two populations (e.g. Mucke et al. 1997). Abdo et al. (2009) found a separation in the $\\gamma$--ray spectral index versus radio luminosity among FSRQs, BL Lacs and misaligned AGN which suggests that more luminous radio sources have softer $\\gamma$--ray spectra. Recent studies of the radio--$\\gamma$ flux correlation in the LBAS sources (Kovalev et al. 2009; 2009a) was conducted using the MOJAVE sample of extragalactic sources (with a flux limit of 1.5 Jy at 15 GHz). Kovalev et al. (2009) find that the parsec--scale radio emission and the $\\gamma$--ray flux are strongly related in bright $\\gamma$--ray objects, suggesting that \\fe\\ selects the brightest objects from a flux--density limited sample of radio--loud sources. In this paper we search for the possible counterparts of the \\fe\\ 1FGL sample in the AT20G survey (\\S 1). By intersecting the association results with the 1LAC sample we study the radio vs $\\gamma$--ray properties of these associations (\\S 2). The 1FGL--AT20G correlation also reveals 8 new associations which are unclassified sources in the 1FGL sample. For two of these there are \\sw\\ observations which, combined with the radio and \\fe\\ data, allow to build the broad band SED. One of these two sources is classified here as a new candidate FSRQ (\\S 3). We also present, as a first result of an on--going study of the properties of the 46 unknown sources of the 1LAC sample, the SED of the brightest radio source of the UIS class that we associate with one 1FGL source. We classify this object as a BL Lac due to the similarity of its SED to that of OJ 287 (\\S 4). We discuss our findings in \\S 5. ", "conclusions": "In this paper we cross correlated the \\fe\\ 1FGL catalog (A10) containing the sources detected above 100 MeV during the 11 months survey with a complete sample of radio sources selected by the AT20G survey in the southern emisphere with 20 GHz flux density larger than 40 mJy. The cross correlation led to identify highly probable (association probability $>$80\\%) radio counterparts for 230 1FGL sources. 222 of these are already classified as blazars in the 1LAC catalog (A10a) and 8 sources are new associations. Using the sample of the associated sources and considering the uniform radio flux measurement of the AT20G survey, we have studied the radio to $\\gamma$--ray flux correlation among different blazar sub--classes finding that there is a significant correlation between these two fluxes both for BL Lacs and for FSRQs. Such a correlation has a slope of $\\sim$0.85 for both blazars sub--classes. If this correlation is further confirmed in the future by the extension of the sample of \\fe\\ blazars with measured radio fluxes, it would help to correctly estimate the contribution of blazars to the $\\gamma$--ray background (e.g. Stecker et al. 1993). The radio properties of the 230 associations show also that they are typically flat spectrum radio sources. In particular, steep spectrum radio sources are not the radio counterparts to \\fe\\ sources, but also radio sources with extremely hard radio spectra (i.e. $\\alpha$(5--20 GHz)$>$+0.7) which are the 1.2\\% of the AT20G sample are not preferentially detected by \\fe. Indeed, we find only three associations with such hard radio sources. We have presented the spectral energy distribution of three sources: two (1FGL J0904.7--3514 and 1FGL J1656.2--3257) are among the eight new associations that we have found by cross correlating the \\fe\\ sample and the AT20G survey. Their spectrum, obtained by combining radio, \\sw\\ and \\fe\\ data resembles that of a typical FSRQ. The third source (1FGL J1327--5257) is presented as a first interesting result of an on--going program aimed at studying and classify the unknown blazar candidates of the 1LAC sample. This source % is among the brightest radio sources in the AT20G survey and it is characterised by a rapidly variable radio emission. Its SED is remarkably similar to that of the prototypical BL Lac OJ 287, but located at the slightly greater redshift $z\\sim$0.4. \\begin{table} \\centering \\begin{tabular}{|l|l|l|l|l|l|} \\hline \\hline \\small RA &Dec &P &Sep &1FGL &AT20G \\\\ deg ° & &arcmin & & \\\\ \\hline 116.613 &--07.164 &0.97 &2.70 &J0746.5--0711 &J074627--070951 \\\\ 136.176 &--35.239 &0.99 &1.07 &{\\bf J0904.7--3514} &J090442--351423 \\\\ 188.529 &--57.598 &0.99 &1.18 &J1234.0--5736 &J123407--573552 \\\\ 194.020 &--59.328 &0.98 &3.06 &J1256.1--5922 &J125604--591943 \\\\ 198.767 &--53.576 &0.98 &3.77 &J1314.9--5338 &J131504--533436 \\\\ 254.070 &--33.035 &0.94 &5.10 &{\\bf J1656.2--3257} &J165616--330207 \\\\ 262.945 &--30.052 &0.90 &7.07 &J1732.0--2957 &J173146--300309 \\\\ 275.910 &--34.903 &0.99 &0.93 &J1823.5--3454 &J182338--345412 \\\\ \\hline \\hline \\end{tabular} \\caption{ The eight new associations found by the cross-correlation of the 1FGL sample and the AT20G survey. Positions (RA and Dec) of the associated radio source, probability of the association and angular separation between the AT20G counterpart and the 1FGL centroid are reported. The two sources (J0904 classified as FSRQ in this paper) whose SED is studied in this work are marked in boldface. } \\label{tab2} \\end{table}" }, "1003/1003.4509_arXiv.txt": { "abstract": "We have searched for variable sources in the core region of M\\,80, using far ultra-violet data taken with the Advanced Camera for Surveys on board the Hubble Space Telescope. We found three sources that exhibit strong signs of variability in our data. Among these is source TDK\\,1, which we believe to be an RR\\,Lyrae star that reached maximum brightness during our observations. The light curve shows a $\\gtrsim3$ mag FUV brightening over the course of $\\approx5$ hours, with an estimated peak brightness of $\\approx16.7$ mag, followed by a decrease to $\\approx20$ mag. Archival optical data obtained with WFPC2 confirm that TDK\\,1 is variable in all wavebands. TDK\\,1's SED is reasonably fit by a star with temperature $T_{eff}\\simeq6700$\\,K and radius $R\\simeq4.2\\,R_{\\odot}$, consistent with the suggestion that it is an RR\\,Lyrae. Based on the photometric and variability characteristics of the other two variables, we suggest that TDK\\,2 is likely to be an SX\\,Phoenicis star with $\\simeq55$ minutes period, and TDK\\,3 is likely another RR\\,Lyrae. Finally, we briefly discuss the FUV counterparts to two previously known variables in M\\,80, the classical nova T\\,Sco and the dwarf nova DN\\,1. ", "introduction": "M\\,80 is one of the densest Globular Clusters (GCs) in the Milky Way, and is famous for the classical nova T~Scorpii which was discovered in 1860 when it outshone the rest of the cluster. Despite the many observations that resulted from the discovery of the nova, only a few variable sources are known in M80 (Wehlau 1990, Clement, private communication, Kopacki 2009). Here, we report the results of a variability study based on our far ultra-violet (FUV) survey of M\\,80 (Dieball et al. 2010; henceforth Paper~1). We have discovered three stars which exhibit significant variability, the most interesting of which is an RR\\,Lyrae star in the core of the cluster. This star, which we call TDK\\,1, was observed around the peak of the light curve in our observations, manifesting a high-amplitude ($> 3$ mag), luminous ($L_{\\rm{UV}}\\simeq6\\times10^{34}\\ \\rm{erg}\\ \\rm{s}^{-1}$), short-duration ($t\\lesssim5\\ \\rm{hours}$) FUV brightening. Further investigation using archive data show that it is also variable in optical wavebands but the data coverage is not sufficient to determine the period of the variation. The other two new variables discovered in our survey include another possible RR\\,Lyrae and an SX\\,Phoenicis star. Finally, we discuss the FUV counterparts to two known variables: the famous nova T~Scorpii and the known dwarf nova DN\\,1 (Shara \\& Drissen 1995). Our paper is structured as follows: In Section~\\ref{obs}, we describe the observations and data analysis, including the FUV survey and analysis of archive HST observations of TDK\\,1. The variable sources identified in the cluster are discussed in Section~\\ref{variables}. Our results are summarized in Section~\\ref{summary}. ", "conclusions": "\\label{summary} We used 32 individual FUV images from our UV survey of the core region of M\\,80 (Dieball et al. 2010; referred to throughout this paper as Paper~1) to search for variable sources in our FUV catalogue. Three sources exhibit strong evidence for variability. TDK\\,1 (Paper~1's source no. 2817) is an RR\\,Lyrae in the core of the cluster. The FUV light curve shows that it was observed from around 40 minutes before to 4.5 hours after maximum brightness, and further investigation using archival WFPC2 optical data showed that it is clearly variable in all wavebands. Its SED is reasonably well described by a star of temperature $T_{eff}\\approx6700$\\,K and radius $R\\approx4.2\\,R_{\\odot}$, consistent with expected parameters for an RR\\,Lyrae star. More specifically, we show that TDK\\,1 is a type~ab RR\\,Lyrae, based on the asymmetry in the FUV light curve. This is only the third cluster in which an RR\\,Lyrae star has been identified based on UV observations (others were found in NGC\\,1851 by Downes et al. (2004) and M\\,15 by Dieball et al. (2007)). UV surveys can be useful tools in identifying RR\\,Lyraes and similar objects, particularly in the cores of (dense) GCs where optical surveys are seriously hampered by crowding. TDK\\,2 (source no.~2238 in Paper~1) is likely an SX\\,Phoenicis star with a period of $55.42 \\pm 0.66$ minutes and amplitude of $\\approx1$\\,mag. TDK\\,3 (source no.~2324) might be another RR\\,Lyrae or a Cepheid. Finally, we discussed two well known variable sources, T\\,Sco and DN\\,1, the FUV counterparts of which were recovered in our survey. T\\,Sco exhibited surprisingly little flickering in our FUV data, possiby because it was in a high state compared with the NUV observations a month later. DN\\,1 is a very faint UV source, so photometric errors dominate over any possible intrinsic flickering or other variations. After this paper was completed, we found that TDK\\,1 and TDK\\,3 are included in Kopacki's variability survey of M80 (Kopacki, private communication). Kopacki agrees with our classification of these two sources as RR\\,Lyrae stars. A preliminary summary of his results, including periods but not including coordinates or finder charts for the sources, is given in Kopacki (2009)." }, "1003/1003.4744.txt": { "abstract": "We use $1\\,{\\rm kpc}$ resolution cosmological AMR simulations of a Virgo--like galaxy cluster to investigate the effect of feedback from supermassive black holes (SMBH) on the mass distribution of dark matter, gas and stars. We compared three different models: (i) a standard galaxy formation model featuring gas cooling, star formation and supernovae feedback, (ii) a \"quenching\" model for which star formation is artificially suppressed in massive halos and finally (iii) the recently proposed AGN feedback model of \\cite{Booth:2009p501}. Without AGN feedback (even in the quenching case), our simulated cluster suffers from a strong overcooling problem, with a stellar mass fraction significantly above observed values in M87. The baryon distribution is highly concentrated, resulting in a strong adiabatic contraction (AC) of dark matter. With AGN feedback, on the contrary, the stellar mass in the bright central galaxy (BCG) lies below observational estimates and the overcooling problem disappears. The stellar mass of the BCG is seen to increase with increasing mass resolution, suggesting that our stellar masses converges to the correct value from below. The gas and total mass distributions are in better agreement with observations. We also find a slight deficit ($\\sim$10\\%) of baryons at the virial radius, due to the combined effect of AGN-driven convective motions in the inner parts and shock waves in the outer regions, pushing gas to Mpc scales and beyond. This baryon deficit results in a slight {\\it adiabatic expansion} of the dark matter distribution, that can be explained quantitatively by AC theory. ", "introduction": "Galaxy clusters are ideal laboratories to study galaxy formation in a dense environment. Galaxies in clusters are observed in many morphological types, from blue extended spirals to red massive elliptical spheroids. The origin of the morphological evolution of galaxies in clusters is still poorly understood. Tidal and ram pressure stripping trigger fast evolution in the properties of accreted satellites, while complex gas cooling and heating processes control the amount of stripped gas that can actually reach the central region of the cluster. In this context, the origin of bright cluster galaxies (BCG) still raises many questions. In the current cosmological framework, the formation of galaxies at the bright end of the luminosity function is still affected by the so--called \"overcooling problem\": using both semi-anaytical models and computer simulations, it has been shown that the massive galaxies are predicted to be too bright and too blue when compared to massive galaxies in the nearby universe (see the recent review by \\cite{Borgani:2009p728}). As a consequence, these models find a stellar content in massive clusters that is significantly above the observed values \\citep{Kravtsov:2005p702}, even including rather extreme supernovae feedback recipe \\citep{Borgani:2004p1066}. In order to overcome this issue, feedback from supermassive black holes (SMBH) have been proposed as a mechanism to prevent gas from accumulating in the cluster core. The so--called AGN feedback scenario has received support from theoretical considerations \\citep{Tabor:1993p1080, Ciotti:1997p1087, Silk:1998p941}, but the strongest evidence comes from the correlated observations of X-ray cavities and radio blobs in massive clusters. These features are usually interpreted as buoyantly rising bubbles of high entropy material injected in the cluster core by jets of relativistic particles. Detailed models of bubbles \\citep{Churazov:2001p2079, Ruszkowski:2004p2156, Bruggen:2005p2162} and jets \\citep{Reynolds:2001p2143, Omma:2004p2154, Cattaneo:2007p420} have been proposed in the context of cluster cores heating, usually based on spherically or planar symmetric, idealized cluster configurations. Based on simple energetic arguments, it is possible to relate the growth of SMBHs at the centre of massive galaxy spheroids to the energy required to unbind the overcooling gas \\citep{Silk:1998p941}. The idea of star formation being regulated by AGN feedback at the high mass end of the galaxy mass function has been applied first quite successfully to hydrodynamical simulations of galaxy merger \\citep{DiMatteo:2005p1020} and then to semi-analytical models of galaxy formation \\citep{Croton:2006p5348, Bower:2006p1007, Cattaneo:2006p990, DeLucia:2007p1656, Cattaneo:2009p946}. AGN feedback models have been included only recently in cosmological simulations of galaxy groups and clusters \\citep{Sijacki:2007p1032, Puchwein:2008p767,McCarthy:2009p747}. Although the detailed physical modeling of SMBHs growth usually differ \\citep{Sijacki:2007p1032, Booth:2009p501}, these simulations, exclusively based on the GADGET code \\citep{Springel:2005p1064}, have obtained very encouraging results regarding the global properties of the simulated clusters \\citep{Puchwein:2008p767, McCarthy:2009p747, Puchwein:2010p763, Fabjan:2010p787}. In this paper, we report on the first AMR high-resolution simulation of a Virgo-like cluster, following both SMBH and star formation, with the associated feedback. A companion paper is exploring the properties of a jet--based, AGN feedback model \\citep{Dubois:2010p5834}. We use the AMR code RAMSES, which differs significantly from the other code used so far to model AGN feedback in a fully cosmological simulation, namely the GADGET code. Although AMR schemes suffer from larger advection errors than SPH, which is a strictly Galilean invariant method, there are some definitive advantages of using AMR in this context: although the total energy (kinetic + thermal + potential) is conserved only at the percent level, it is strictly conserving for the fluid energy (kinetic + thermal), which is very important in presence of strong shocks, and it captures hydrodynamical instabilities more realistically \\citep{Agertz:2007p415, Wadsley:2008p1093, Mitchell:2009p4163}, which is very important in presence of convective motions of buoyant gas. It also captures gas stripping of infalling satellites due to hydrodynamical instabilities more realistically than standard implementations of SPH \\citep{Agertz:2007p415}. We have adapted the SMBH growth model of \\cite{Booth:2009p501} to the sink particle method for AMR presented by \\cite{Krumholz:2004p1079}. With respect to the previous work of \\cite{Booth:2009p501} and follow-up papers, we have made significant improvements over the OWL simulations suites in term of mass and spatial resolution, but only for one single zoom-in simulation of a Virgo--like cluster. Recently, \\cite{Puchwein:2008p767} and \\cite{Puchwein:2010p763} have also used the GADGET code, but a different AGN feedback model, to perform zoom-in simulations of groups and clusters of galaxies, with however a mass resolution and a gravitational softening length not as good as the one we used in our high resolution case. Note also that the minimum smoothing length in these SPH simulations is usually much smaller than the gravitational softening length. In this paper, we have specifically chosen a Virgo-like cluster, in order to compare our results with the very detailed observations that are available for this well-known astronomical object. We will focus here on the mass distribution of the three main components, namely dark matter, gas and stars. The paper is organized as follows: the first section is dedicated to numerical methods and physical ingredient (cooling, star formation and AGN feedback), while our second section presents our results, comparing our three models. The final section is left for discussion. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ", "conclusions": "\\label{sec:summary} We have simulated the formation of a Virgo--sized galaxy cluster to study the effects of feedback on the overcooling problem. The impact of AGN feedback on the distribution of the baryonic mass is strong, and in good agreement with previous SPH simulations: star formation in massive galaxies is drastically reduced. At the same time, left-over gas is very efficiently removed from the core of the parent halos, where it would have otherwise accumulated. In order to quantify the effect of AGN feedback, we have run two other reference simulations: one model with only star formation and supernovae feedback (the standard scenario) and one model for which we have artificially prevented star formation to occur in massive enough spheroids (the quenching scenario). A detailed comparison of the three models clearly demonstrates that AGN feedback is needed to control star formation in the central BCG, but also to unbind the overcooling gas from the cluster core. We also clearly identify the effect of the baryon dynamics on the dark matter mass distribution on large scale. Interestingly enough, in case of AGN feedback, we observe the adiabatic expansion of the dark halo, an effect well modeled by the AC theory of \\cite{Gnedin:2004p569}. A comparison of our simulation results with observational data for Virgo and its central galaxy M87 rules out the standard model, but also the quenching model. On the contrary, our simulation with AGN feedback, although not fully converged yet, shows a much better agreement with M87 data in term of mass distribution. In particular, we obtain a significantly reduced baryon fraction within the virial radius, in agreement with observations compiled by \\cite{Lin:2003p1820} and \\cite{Gonzalez:2007p916}. We clearly identify in our simulation that gas is removed from the core of the cluster by convective motions and/or strong shocks, and accumulates in a region just outside the virial radius. When compared to gas profiles inferred from X-ray data, our AGN model produces too shallow gas distribution, suggesting that we probably need even more powerful feedback processes. Our cluster formation simulations with AGN feedback have not fully converged yet - as we increase the resolution, we find a stellar mass profile for the BCG that is in better agreement with observations, but it is still too low by about a factor of two. We are still missing the lowest mass galaxy population, which could provide the missing stellar mass in the central elliptical galaxy. We also note that, in the current picture, AGN feedback is a morphologically dependent process: it only directly affects galaxies with SMBHs, i.e. galaxies with a significant bulge/spheroid component. Higher resolution studies would be needed, in order to reliably model this distinction, so that star formation in disky galaxies is not artificially suppressed." }, "1003/1003.1430_arXiv.txt": { "abstract": "To understand the origin of hard X-ray emissions from black hole X-ray binaries during their low/hard states, we calculate the X-ray spectra of black-hole accretion flow for the following three configurations of hot and cool media: (a) an inner hot flow and a cool outer disk (inner hot flow model), (b) a cool disk sandwiched by disk coronae (disk corona model), and (c) the combination of those two (hybrid model). The basic features we require for successful models are (i) significant hard X-ray emission whose luminosity exceeds that of soft X-rays, (ii) high hard X-ray luminosities in the range of $(0.4 - 30)\\times 10^{37}$ erg s$^{-1}$, and (iii) the existence of two power-law components in the hard X-ray band with the photon indices of $\\Gamma_{\\rm s}\\sim 2 > \\Gamma_{\\rm h}$, where $\\Gamma_{\\rm s}$ and $\\Gamma_{\\rm h}$ are the photon indices of the softer ($< 10$ keV) and the harder ($> 10$ keV) power-law components, respectively. Contribution by non-thermal electrons nor time-dependent evolution are not considered. We find that Models (a) and (b) can be ruled out, since the spectra are always dominated by the soft component, and since only one power-law component, at most, can be reproduced. Only Model (c) can account for sufficiently strong hard X-ray emissions, as well as the existence of the two power-law components, for a large ratio of the accretion rate in the corona to that in the thin disk. The outer disk corona (where the Compton $y$-parameter is smaller, $y < 1$) produces the softer power-law component with photon index of $\\Gamma_{\\rm s} \\sim 2$, whereas the inner hot flow (where $y \\gtrsim 1$) generates the harder component with $\\Gamma_{\\rm h} < 2$. This model can also account for the observed relationship between the photon index and the reflection fraction. ", "introduction": "It is well known that black hole X-ray binaries show two basic spectral states: the high/soft state appearing at high luminosities (typically 10\\% $L_{\\rm Edd}$, where $L_{\\rm Edd}$ is the Eddington luminosity) and the low/hard spectral state at low luminosities, typically a few percent of $L_{\\rm Edd}$, by factors of a few smaller than that of the high/soft state (McConnell et al. 2002). In the high/soft state X-ray spectrum is dominated by a thermal soft X-ray component of temperature, $kT_{\\rm eff} \\sim 1$ keV, associated with a weak steep power-law component with photon index $\\Gamma>$2. In the low/hard state, on the other hand, spectra are mainly composed of Comptonized emission components in the hard X-ray band, soft excess components, and reflection components from the optically thick cold medium (for recent reviews, see McClintock \\& Remillard 2003; Done et al. 2007). While the high/soft state is reasonably well understood in terms of the standard disk model (see, e.g. Makishima et al. 1986), the nature of the accretion flow in the low/hard state is still a matter of debate. Historically, two classes of models have been considered: the inner hot flow model and the disk corona model (Thorne \\& Price 1975; Shapiro et al. 1976; Liang \\& Price 1977; for a recent review of theoretical models, see Kato et al. 2008). In the inner hot flow model, an optically thick disk is radially truncated at some radius, $r_{\\rm tr}$, and the central region, $r$~--15.7) galaxies, a metallicity of 12~+~log~(O/H)~$\\sim$~7.65 (i.e., the threshold metallicity for XMD galaxies) is consistent with what one would expect from the luminousity--metallicity relation. \\item The lower bound to the 95~per cent confidence interval, around the L--Z relation, is given by Z~=~--0.177~M$_B$~+~4.87. We suggest that this would be a more appropriate definition for XMD galaxies. As per this definition, 17 of the 31 galaxies, in our sample, would be identified as XMD. \\item As shown in previous work, we find that the L--Z relation of BCGs is shifted to higher luminousities at a given metallicity, compared to dIs. This has been considered as an evidence for the fact that the luminousity of BCGs is temporarily enhanced due to ongoing starburst. Consistent with this suggestion, we find that the M--Z relations of our BCG and dI samples are better matched. \\item The effective chemical yields of XMD BCGs are systematically lower than BCGs and dI galaxies, with similar gas fractions and baryonic masses. We suggest that this is because of better mixing of the ISM in the case of XMD galaxies, as compared to other gas-rich dwarf galaxies. Motivated by the observation that many XMD galaxies are interacting (Ekta et al. 2008) we suggest that the better mixing is due to tidal interactions. We propose that XMD galaxies deviate from the L--Z and M--Z relations because of a combination of having low effective chemical yields and higher gas fractions. \\end{enumerate} ~\\\\ {\\bf ACKNOWLEDGMENTS} \\\\ We thank the anonymous referee for useful comments, which helped in increasing the clarity of the paper. This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England. The SDSS Web Site is http://www.sdss.org/. The SDSS is managed by the Astrophysical Research Consortium for the Participating Institutions. The Participating Institutions are the American Museum of Natural History, Astrophysical Institute Potsdam, University of Basel, University of Cambridge, Case Western Reserve University, University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington." }, "1003/1003.1789_arXiv.txt": { "abstract": "We present the analysis of an \\xmm observation of the Seyfert galaxy NGC~2992. The source was found in its highest level of X-ray activity yet detected, a factor $\\sim 23.5$ higher in 2--10~keV flux than the historical minimum. NGC~2992 is known to exhibit X-ray flaring activity on timescales of days to weeks, and the \\xmm data provide at least factor of $\\sim 3$ better spectral resolution in the Fe~K band than any previously measured flaring X-ray state. We find that there is a broad feature in the $\\sim 5-7$~keV band which could be interpreted as a relativistic \\fekalfa emission line. Its flux appears to have increased in tandem with the 2--10~keV continuum when compared to a previous \\suzaku observation when the continuum was a factor of $\\sim 8$ lower than that during the \\xmm observation. The \\xmm data are consistent with the general picture that increased X-ray activity and corresponding changes in the \\fekalfa line emission occur in the innermost regions of the putative accretion disk. This behavior contrasts with the behavior of other AGN in which the \\fekalfa line does not respond to variability in the X-ray. ", "introduction": "\\label{intro} Relativistically broadened \\fekalfa emission lines in active galactic nuclei (AGNs) and Galactic Black-Hole Binaries (GBHB) are potentially a powerful probe of accreting black-hole systems. In principle, modeling the \\fekalfa line profile, supplemented with assumptions about its radial and angular emissivity function in the accretion disk, can yield information on the inclination angle of the disk, its radial extent, the proximity of its inner edge to the black-hole, and the angular momentum, or spin of the black hole (e.g. Dov\\v{c}iak \\etal 2004a; Beckwith \\& Done 2004; Brenneman \\& Reynolds 2006). Whilst there are plenty of studies in the literature that address the {\\it occurrence} and modeling of the time-averaged broad \\fekalfa line profiles (e.g. Guainazzi \\etal 2006; Nandra \\etal 2007; Miller \\etal 2007, and references therein), observational data on the {\\it variability} of the relativistic lines in AGNs is remarkably sparse (e.g. see Fabian 2006, and references therein). Variability of the line shape and intensity, especially in relation to the variability of the continuum, would provide stronger and more model-independent information on the accretion-disk and black-hole parameters (in particular on the unknown radial emissivity function). It would also lay the groundwork for future reverberation mapping of the black-hole metric in the strong gravity regime (e.g. Reynolds \\etal 1999). However, there are only two AGNs for which it has been possible to assess the variability of the broad \\fekalfa line and in both of those cases it is found that the line intensity does not respond to the observed X-ray continuum (MCG~$-$6-30-15: e.g. Fabian \\etal 2002; NGC~4051: Ponti \\etal 2006a). {Strictly speaking the NGC~4051 results refer to the reflection continuum appearing to correlate with the intrinsic continuum at small fluxes, but then remaining roughly constant as the intrinsic continuum flux increases. However, this interpretation is model-dependent because the intrinsic continuum shape and flux is determined indirectly.} This behavior has been interpreted in terms of the so-called ``light-bending'' model in which physical motion of the X-ray source in the strong gravity field of the black hole causes variability in the X-ray continuum for an observer ``at infinity'' but the illumination of the disk (and therefore the \\fekalfa line flux) remains relatively unaffected. It is the variable amplification effect of light-bending that causes the X-ray continuum variability--any intrinsic variability of the luminosity of the X-ray source must be suppressed or absent, otherwise the \\fekalfa line intensity would once again vary. In a few other AGN, although the variability of the \\fekalfa line cannot be isolated, the overall behavior of the disk-reflection and fluorescence spectrum (of which the emission line is a part), has been interpreted as being consistent with this scenario (e.g. 1H~0707$-$495, Fabian \\etal 2004; 1H~0419$-$577, Fabian \\etal 2005; IRAS~13224$-$3809, Ponti \\etal 2006b). Similar non-variability of the relativistic \\fekalfa line has been observed in a GBHB (XTE~J16550--Rossi \\etal 2005), and this has also been interpreted in the context of the light-bending model. The origin of the physical motion of the X-ray source (relative to the black-hole) and it's geometry remain unspecified in the light-bending model. We note that alternative interpretations of the apparent broad \\fekalfa emission lines in the above-mentioned (and other) sources in terms of absorption-only models remain viable and are still a matter of considerable debate (e.g. Miller \\etal 2008, 2009; Turner \\etal 2009; Turner \\& Miller 2009 and references therein). The reason why so little observational data exists on the variability of the relativistic \\fekalfa line, in AGN at least, is because several criteria must be met by the source at once, and such objects are not common. Firstly, the current sensitivity of X-ray detectors require that the source is bright enough that the time-sliced spectra have sufficient signal-to-noise. Secondly, there must be be a sufficient quantity of data (commensurate with the timescale of variability) in order to meaningfully test variability models. Thirdly, the source must have a well-measured broad \\fekalfa line with a large enough equivalent width that its contrast against the continuum is high enough to constrain the line profile parameters with sufficient accuracy. Fourth, the amplitude of variability of the {\\it continuum} should be large enough in order to be able to test whether the {\\it line} flux varies in response to it. Although there are plenty of AGN that satisfy one of these criteria, sources that satisfy all four are very rare. Indeed MCG~$-$6-30-15 and NGC~4051 are the only two AGN for which a definitive result has been reported on the non-response of the broad \\fekalfa to continuum variations (as opposed to response of the broadband reflection spectrum). In the present paper we show that non-variability of the relativistic \\fekalfa line in NGC~2992 in response to an increase in the X-ray continuum level is ruled out, in contrast to the known behavior of the other AGN discussed above. NGC~2992 is a nearby ($z=0.00771$) Seyfert galaxy that has been observed by every major X-ray mission over the past $\\sim30$ years, and has shown hard X-ray flux variability by more than a factor of $20$ (Piccinotti \\etal 1982; Turner \\& Pounds 1989; Turner \\etal 1991; Nandra \\& Pounds 1994; Weaver \\etal 1996; Gilli \\etal 2000; Matt \\etal 2003; Beckmann \\etal 2007; Murphy \\etal 2007; Yaqoob \\etal 2007). In the present paper we describe the results of an \\xmm observation of NGC~2992 made in 2003, May. Some results from this observation from spectral fitting over a restricted bandpass ($2.5-10$~keV) have been presented by Brenneman \\& Reynolds (2009). However, Brenneman \\& Reynolds (2009) did not account for the heavy photon pile-up in the \\xmm CCD data, which considerably distorts the spectrum. Here we examine the data in detail, over the full \\xmm bandpass, and in the context of the historical X-ray behavior of NGC~2992. During the \\xmm observation NGC~2992 was at the top end of its historical range of activity in terms of the 2--10~keV flux. We detected a strong, broad emission-line feature in the Fe~K band which implies extreme variability in line flux compared to historical \\suzaku data. This paper is organized as follows. In \\S\\ref{obs} we describe the \\xmm observation and data reduction. Spectral analysis of the data is described in \\S\\ref{fitting}. In \\S\\ref{results} we discuss the results in detail and in \\S\\ref{historicaldata} we compare the \\xmm results with those from historical data. In \\S\\ref{conclusions} we discuss some general implications of our findings and present our conclusions. A cosmology with $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Lambda=0.73$, $\\Omega=1$ is assumed throughout. ", "conclusions": "\\label{conclusions} In the past three decades, the Seyfert galaxy NGC~2992 has shown X-ray continuum flux variations of more than a factor of 20. We have reported the results of an \\xmm observation of NGC~2992 performed in 2003, during which the source was found to have the highest 2--10~keV flux ($9.4 \\times 10^{-11} \\ \\rm erg \\ cm^{-2} \\ s^{-1}$) compared to previous historical values. No previous X-ray astronomy mission to date has observed NGC~2992 in a high state (2--10~keV flux greater than $\\sim 7 \\times 10^{-11} \\ \\rm erg \\ cm^{-2} \\ s^{-1}$) with CCD spectral resolution ($\\sim 7000 \\ \\rm km \\ s^{-1}$ FWHM at 6.4~keV). The best spectral resolution available for any previous high-state data set for NGC~2992 was a factor of $\\sim 3$ worse than that for the \\xmm data in the Fe~K band. The rather unique \\xmm pn spectrum of NGC~2992 has a broad feature in the $\\sim 5-7$~keV band that can be interpreted as relativistic \\fekalfa line emission. Its flux is an order of magnitude larger than that found in \\suzaku data (obtained in 2005), when the 2--10~keV continuum flux of NGC~2992 was a factor of $\\sim 8$ less than that during the \\xmm observation. Although the detailed \\fekalfa line parameters obtained from the \\xmm data are model-dependent, it appears that the absolute luminosity of the broad \\fekalfa line and the associated reflection continuum increase as the continuum luminosity increases. The observation of {\\it variable} broad, relativistic \\fekalfa line emission in AGN is rare, and the observation of such variability in response to X-ray continuum variability is even rarer. In the light-bending model (Fabian \\& Vaughan 2003; Miniutti \\& Fabian 2004) invoked to account for the {\\it non-variability} of the broad \\fekalfa line in MCG~$-$6-30-15, the X-ray continuum variability for an observer at infinity is due entirely to relativistic effects as the X-ray source physically changes position relative to the black hole. The disk, being much closer to the X-ray source and black hole than the distant observer, is then not subject to large variability in illumination by the X-ray continuum and therefore produces an \\fekalfa emission line that is not variable. In this model there can of course be no intrinsic variability of the X-ray source, otherwise the \\fekalfa line would be variable. The \\xmm data for NGC~2992, when considered in the context of historical data, therefore suggest that the light-bending scenario is not relevant for this AGN, implying that the large-amplitude X-ray continuum variability is {\\it intrinsic} to the X-ray source. Further, if there were X-ray continuum variability that had {\\it both} an intrinsic origin, and one due to relative motion of the X-ray source and black hole, the broad \\fekalfa line and continuum variability would not be related in a simple way. The high level of the intrinsic continuum during the \\xmm observation reported here swamped the features that originate in circumnuclear matter that is more extended than the X-ray source because these features did not respond to the large-amplitude change in the intrinsic continuum that must have occurred prior to the start of the \\xmm observation. These features (namely the optically-thin thermal emission, the scattered intrinsic continuum, and the narrow core of the \\fekalfa line) have been better studied in historical spectra taken during low-continuum states (e.g. Weaver \\etal 1996; Gilli \\etal 2000; Yaqoob \\etal 2007). We thank the referee for his/her useful comments. The authors thank the \\xmm instrument teams and operations staff for making the observation of NGC~2992. This research made use of the HEASARC online data archive services, supported by NASA/GSFC. The work was supported by Chinese NSF through Grant 10773010/10825312, and the Knowledge Innovation Program of CAS (Grant No. KJCX2-YW-T05). \\newpage" }, "1003/1003.3745_arXiv.txt": { "abstract": "{The sterile neutrino is a weakly-interacting particle that emerges in the framework of certain extensions of standard particle physics and that fits naturally with the properties of a warm dark matter particle candidate. We present an analysis of a deep archival \\cha\\ observation of Willman 1, one of the darkest ultra-faint dwarf galaxies up to date, to exclude the presence of sterile neutrinos in the 1.6--16 keV mass range within 55 pc of its center down to the limiting flux of the observation. Spectral analysis of the \\cha\\ data fails to find any non-instrumental spectral feature possibly connected with the radiative decay of a dark matter particle. Accordingly, we establish upper bounds on the sterile neutrino parameter space and discuss it in the context of previous measurements. Regarding the point source population, we identify a total of 26 sources within the central $5\\arcmin$ to a limiting $0.5-2.0$ keV X-ray flux of $6\\times 10^{-16}$ erg cm$^{-2}$ s$^{-1}$. While some of these sources could be formal members of Willman 1, we find no outstanding evidence for either an unusual population of bright X-ray sources or a densely populated cluster core. In fact, the entire X-ray population could be explained by background active galactic nuclei and/or foreground stars unrelated to Willman 1. Finally, possible associations of the X-ray point like population with optical sources from the {\\it SDSS DR7} catalogue are presented.} ", "introduction": "There is strong evidence for the existence of dark matter (DM) in the Universe. In the concordance cosmological model (CCM), 83\\% of the mass density in the Universe cannot be explained with ordinary baryonic matter and requires an additional non-baryonic component \\citep{Komatsu:2010a}. Without a doubt, understanding DM is one of the most important topics of physics today \\citep{Gaitskell:2004a,Bertone:2005a}. In the search for DM, three different approaches coexist: direct production in collider experiments \\citep{Kane:2008a}, direct detection through scattering off ordinary matter \\citep{Cerdeno:2010a}, and indirect detection based on the search for secondary particles produced by the annihilation or decay of DM particles \\citep{Bertone:2005b}. Within the last category, indirect searches for supersymmetric particles at center-of-mass energies between 10 GeV and a few TeV from pair annihilations or particle decays are being performed through gamma-ray observations of astrophysical objects with high dark matter density \\citep{Aliu:2009a,Abdo:2010b}. While no significant dark matter signal has been verified using gamma-ray photons as messenger particles \\citep{Anderson:2010a} some intriguing claims have emerged from measurements of the cosmic-rays positron fraction \\citep{Chang:2008a,Adriani:2009a} and electron spectrum in the same energy range \\citep{Abdo:2009a,Aharonian:2009a}. There are several theories offering DM particle candidates which can annihilate into standard model particles and eventually produce photons and leptons at energies above 10 GeV. A neutralino in the supersymmetric extensions of the standard model \\citep{Haber:1985a} or the Kaluza-Klein particle emerging from universal extra dimension theories \\citep{Bergstrom:2005a} are two well known examples of cold dark matter (CDM) particle candidates. However, since the nature of the DM particle remains uncertain, one must broaden the range of hypothesised candidates/signatures taking in consideration options outside the gamma-ray regime \\citep{Feng:2010a}. If we move to the X-ray band, a possible candidate emerges in the form of the sterile neutrino. Sterile neutrinos are weakly-interacting fermions which arise as the right-handed counterparts of the standard neutrinos in some extensions of the standard model of particles. The lightest of these might lie in the keV range and would be compatible with a warm dark matter (WDM) candidate \\citep{Dodelson:1994a}. The sterile neutrino, besides qualifying as a good WDM particle candidate and resolving the neutrino mass problem, overcomes the puzzle of baryon asymmetry in some of these extensions of the standard model \\citep{Asaka:2005a} and may mitigate some of the shortcomings of cold dark matter cosmologies including the apparent dearth of dwarf galaxies around the Milky Way \\citep[but see][]{deNaray:2009a}. Sterile neutrinos decay in a two-body final state composed by a standard neutrino and a photon. Therefore, compact regions with significant accumulations of sterile neutrinos could ``shine'' in X-rays, producing a detectable X-ray flux line in the 0.1--100 keV energy range \\citep{Abazajian:2001a,Dolgov:2002a,Yuksel:2008a}. Ultimately, if the sterile neutrino is ever found in collider experiments, a clear identification of this particle in dense astrophysical regions would have to follow \\citep{Gelmini:2009a}. As we have discussed above, an unidentified X-ray line could be an invaluable DM smoking gun. Since we are dealing with a decay, the subsequent line flux is directly proportional to the density of the DM region. Hence, astrophysical objects with the highest DM density are the favoured targets. Another condition that one should take into account when defining the best targets is whether the expected X-ray signal would be accompanied by any unrelated source of background that could eventually curtain a sterile neutrino signal. Accordingly, the most prominent targets to date are the recently discovered ultrafaint dwarf galaxies around the Milky Way \\citep{Willman:2005a,Belokurov:2007a}. These objects, whose kinematic properties are still under study \\citep{Simon:2010a}, could be the closest and densest dark matter structures in the Local Group, making them excellent targets to conduct DM searches \\citep{Strigari:2008a,Simon:2010a}. Additionally, their low baryonic content could minimise any external X-ray background. Nonetheless these predictions must be verified through careful study of the diffuse and point-like X-ray emission coming from these objects. Indirect searches for sterile neutrinos in dwarf galaxies have been conducted in Fornax \\citep{Boyarsky:2010a}, Ursa Minor \\citep{Loewenstein:2009a}, Willman 1 \\citep{Loewenstein:2010a}, and Segue 1 \\citep{Mirabal:2010a}. Here, we perform such a search using a deep archival \\cha\\ observation of the enigmatic Willman 1 \\citep{Willman:2005a}, an ultra-faint object discovered in the {\\it Sloan Digital Sky Survey} ({\\it SDSS}) and originally classified as a dwarf galaxy \\citep{Willman:2005a,Willman:2006a}. The arganization of the paper is as follows. \\S 2 describes the X-ray observations. In \\S 3 we discuss the diffuse component of Willman 1 and derive the parameter space for sterile neutrinos allowed by the data. The X-ray point source population and spectral analyses are explained in \\S 4. Summary and conclusions are presented in \\S 5. ", "conclusions": "Adopting an estimated distance of 38 kpc, it might be the case that the bulk of sources connected with Willman 1 lies below our luminosity limit of $10^{32}$ erg s$^{-1}$ in the 0.5--2.0 keV energy band, as in the case of globular cluster GLIMPSE-C01 \\citep{Pooley:2007a}. However, our encounter rate calculations indicate a very low probability of finding candidate binaries within Willman 1. Combined with existing optical imaging, the \\cha\\ observation of Willman 1 could help purge possible member stars of any AGN contaminants/stellar interlopers. Furthermore, further analysis might help us investigate how will the X-ray point source population influence the prospects of detecting a dark matter signal from Willman 1 with gamma-ray measurements. Given the existing sterile neutrino limits, it might be wise to concentrate on the detectability of the sterile neutrino with future X-ray experiments \\citep{Herder:2009a} as it will be difficult to improve current measurements with the existing instrumentation \\citep{Abazajian:2009b}. In particular, making the case for future calorimeter experiments that could definitely detect it or exclude it as a DM candidate (F. Paerels, private communication). We have shown here that any viable indirect search for DM in this energy range must deal not only with possible overlap with instrumental lines in the spectrum \\citep{RiemerSorensen:2006a}, but also with possible contamination from intervening plasma lines that may mask a DM origin. We caution that extraordinary claims related to DM should be backed by outstanding observational evidence \\citep{Sagan:1980a}. With the Large Hadron Collider (LHC) and other laboratory experiments coming online, it might be possible to achieve a direct detection of the particles responsible for DM. However, any direct detection in the laboratory will then have to be followed by confirmation or dismissal of sameness in an astrophysical context. Based on the present study, Willman 1 continues to hold a select place among targets to conduct such searches. It is worth noting that within the inner $5\\arcmin$ radius circle of Willman 1, there are no radio sources at 1.4 GHz in the NRAO VLA Sky Survey (NVSS) source catalog \\citep{Condon:1998a}. Moreover, we have shown that the point source population to a limiting 0.5--2.0 keV X-ray flux of $6\\times 10^{-16}$ erg cm$^{-2}$ s$^{-1}$ is consistent with background AGN and/or foreground stars. Pending a final verdict regarding the kinematic distribution of Willman 1 \\citep{Willman:2010a}, the available data from radio through X-rays thus make Willman 1 a notable candidate for the eventual astrophysical verification of a DM particle." }, "1003/1003.3573_arXiv.txt": { "abstract": " ", "introduction": "\\label{sec:intro} Gamma-Ray Bursts (GRBs) are the most powerful explosions in the Universe and, arguably, represent the most significant new astrophysical phenomenon since the discovery of quasars and pulsars. As their name suggests, GRBs are detected as brief, intense and totally unpredictable flash of high-energy gamma rays, thought to be produced during the core collapse of massive stars (long-soft bursts, T$_\\gamma$$>$2~s) or the merger of two compact objects such as two neutron stars or a neutron star and a stellar-mass black hole (short-hard bursts, T$_\\gamma$$<$2~s). Although discovered through their $\\gamma$-ray emission~\\cite{klebesadel73}, they are now known to emit non-thermal radiation detectable across the electromagnetic spectrum \\cite{costa97,vanparadijs97,frail97}. However, despite their enormous luminosity, their unpredictability and short duration limit rapid, accurate localisation and observability with traditional telescopes. Consequently, new ground and space-based facilities have been developed over the past decade; dedicated satellites optimised for GRB detection and followup, such as Swift~\\cite{gehrels04}, are revolutionizing GRB studies by locating $\\sim$100 bursts per year with $\\gamma$-ray positions accurate to $\\sim$3$'$ and X-ray positions accurate to 5$\"$ within seconds or minutes of the burst. Here we describe the automatic ground-based follow-up of GRBs with the world's largest robotic optical telescopes that use intelligent software and innovative instruments. \\paragraph{The Era of Rapid Follow-up: Predictions and Outcomes } Before the launch of current satellites such as Swift, Integral and Fermi, significant progress in understanding GRBs had been made since their discovery, in particular the general $\\gamma$ and X-ray properties. The first crucial step in disseminating real-time GRB positions to ground observers was triggered by BATSE on the CGRO~\\cite{paciesas99} through the GRB Coordinates Network (GCN) \\cite{barthelmy08} via internet socket connection (no humans-in-the-loop). This drove development of the first generation of wide-field robotic followup ground-based facilities, such as GROCSE, ROTSE, and LOTIS, culminating with the discovery of the optical flash associated with GRB~990123 \\cite{akerlof99}. BATSE provided an invaluable catalogue of prompt $\\gamma$-ray profiles, whose isotropic sky distribution and inhomogeneous intensity distribution suggested a cosmological origin \\cite{paciesas99}, and BeppoSAX \\cite{boella97} revolutionised the cosmological study of GRBs by providing sub-arcmin ($\\sim$50$\"$) localisation of X-ray afterglows that enabled late-time ($\\sim$ hours) optical followup with traditional ground-based telescopes and redshift determinations. Collimation of the ejecta (i.e. jets) was inferred from temporal breaks - steepening - of optical light curves at $\\sim$1 day post-burst and the concept of a universal central engine and the use of GRBs as standardisable cosmological candles was introduced \\cite{frail01,ghirlanda04}. \\begin{figure} \\centering \\includegraphics[height=6cm]{fig1.eps} \\caption{Redshift distribution of {\\em Swift} GRBs detected to-date.} \\label{fig:z} % \\end{figure} The possibility for great advances with the launch of Swift was fully recognised. Optical counterparts were expected to be found for all GRBs with many GRBs expected to exhibit bright optical flashes from reverse shock emission at early times, similar to GRB 990123 \\cite{akerlof99}. An increase in the number of GRBs detected would lead to many jet breaks being identified, short GRBs would be easily observed and understood and identification of GRBs at very high redshift would be routine. Instead, 50\\% of GRBs remain optically dark, despite deep, rapid followup \\cite{melandri08,cenko09,oates09,yost07}; there is a dearth of bright reverse-shock optical emission \\cite{roming06}; light curves are complex in all bands with a variety of chromatic and achromatic breaks and flares observed (e.g., \\cite{tagliaferri05,burrows05,obrien06,chincarini07,falcone07,melandri08,klotz09,rykoff09}). Jet breaks have proven elusive, short bursts remain technically challenging \\cite{graham09} and only 3 GRBs have been identified to lie a $z>6$ (Fig. \\ref{fig:z}) \\cite{kawai06,greiner09a,salvaterra09,tanvir09}. ", "conclusions": "Exploring the extreme physics exhibited by GRB explosions is technically challenging due to (a) the unpredictability of their occurrence (b) their short-lifetime rapidly fading emission (c) the wide range of observed brightnesses of optical counterparts, ranging from R~=~5 to $>$22~mag within minutes of the burst itself - all drivers of autonomous follow-up technology. Deep, fast and multi-filter observations are crucial to identify the counterparts to these events that represent the brightest stellar objects observed out to the epoch of reionization. With the advent of Swift, discoveries such as the canonical early X-ray decay, the X-ray flares, the detection of the afterglows of short-duration GRBs, and the recognition that GRB-producing stars exist out at least to z=8.2 keep GRB studies at the forefront of astrophysics. Efforts continue to understand the complexity of the X-ray versus optical afterglow temporal evolution, the circumburst environment properties, in particular the dust versus gas content around the GRB progenitor and along the line of sight through the host galaxy, and the origin(s) of optical flares and their possible interpretations (e.g. GRB~080129 \\cite{greiner09b}: e.g residual collisions in the GRB outflow or hot spots in strongly magnetized ejecta?\\cite{lyutikov06}). Questions remain on the fundamental nature of the relativistic ejecta, the underlying radiation mechanisms and the role of magnetic fields. Observational surprises such as the relative lack of GeV emission from many bursts detected by Fermi and the rich variety in optical properties of GRB counterparts continue to drive developments in GRB modelling and observational technology. The 2-m robotic telescopes described here are proving decisive in tackling many of these issues: the variety of light curves has been investigated both on statistical grounds \\cite{melandri08} as well as in individual cases of special interests: the dark GRB population as well as the luminosity function distribution have been characterised over a broad range of apparent brightnesses and the presence and lack of reverse shock emission in specific cases was investigated. The interplay between forward and reverse shocks within the standard fireball model, as determined by the magnetic properties of the outflow, may yet succeed in explaining the dearth of reverse shocks previously expected from pre-Swift observations. The use of the RINGO polarimeter on the LT, capable of measuring the polarisation of optical counterparts to GRBs as early as a few minutes after the onset of the prompt $\\gamma$-ray emission, has provided the earliest measurements and detection of GRB polarisation, setting important direct constraints on the magnetic field structure of the fireball and on the jet configuration \\cite{mundell07_sci,steele09}. Further progress made in understanding the magnetic field structure of the fireball - large-scale ordered fields are currently preferred over locally tangled fields in the shock layer- will contribute to our knowledge of the nature of the outflow along the jets. Time-resolved early time polarisation light curves (\\% and PA) from the newly commissioned RINGO2 promise to provide unprecedented diagnostic information on the structure and evolution of the outflow and its magnetic field for a statistically significant sample of GRBs down to R$<$17~mag and thus allow powerful discrimination between predictions of the hydrodynamical versus magnetised jet models and ultimately constrain the physics of GRB central engines. {\\it Acknowledgment.} We thank the Liverpool Telescope group for technical, scientific and artwork support. CGM acknowledges financial support from Research Councils U.K. The Liverpool Telescope is located at the Observatorio del Roque de Los Muchachos, La Palma, Canary Islands, Spain. It was designed and built by Telescope Technologies Ltd and is owned and operated by Liverpool John Moores University with financial support from the U.K. PPARC. The Faulkes telescopes, owned by Las Cumbres Observatory, are located in Hawaii (FTN) and Siding Spring, Australia (FTS) and are operated with support from the Dill Faulkes Educational Trust." }, "1003/1003.0485_arXiv.txt": { "abstract": "Spectacular breakthroughs in numerical relativity now make it possible to compute spacetime dynamics in almost complete generality, allowing us to model the coalescence and merger of binary black holes with essentially no approximations. The primary limitation of these calculations is now computational. In particular, it is difficult to model systems with large mass ratio and large spins, since one must accurately resolve the multiple lengthscales which play a role in such systems. Perturbation theory can play an important role in extending the reach of computational modeling for binary systems. In this paper, we present first results of a code which allows us to model the gravitational waves generated by the inspiral, merger, and ringdown of a binary system in which one member of the binary is much more massive than the other. This allows us to accurately calibrate binary dynamics in the large mass ratio regime. We focus in this analysis on the recoil imparted to the merged remnant by these waves. We closely examine the ``antikick,'' an anti-phase cancellation of the recoil arising from the plunge and ringdown waves, described in detail by Schnittman et al. We find that, for orbits aligned with the black hole spin, the antikick grows as a function of spin. The total recoil is smallest for prograde coalescence into a rapidly rotating black hole, and largest for retrograde coalescence. Amusingly, this completely reverses the predicted trend for kick versus spin from analyses that only include inspiral information. ", "introduction": "\\subsection{Modeling binary systems in general relativity} After roughly three decades of effort, numerical relativity can now model nearly arbitrary binary black hole configurations. Following Pretorius' pioneering ``breakthrough'' calculation {\\cite{frans1}}, and then the successes of the Brownsville and Goddard groups using techniques that required only modest modifications to the methods they used before the breakthrough {\\cite{brownsville,goddard}}, the past few years have seen an explosion of activity. Recent work has studied the impact of the many physical parameters which describe binaries, such as mass ratio {\\cite{massratio1,massratio2}}, spin and spin alignment {\\cite{spin1, spin2, spin3, spin4}}, and eccentricity {\\cite{ecc1,ecc2}}. As numerical models have improved, analytic tools for modeling binary systems {\\cite{blanchet06}} and connecting numerics and analytics have likewise matured. In particular, the {\\it effective one-body} (EOB) {\\cite{bd1, bd2, djs, d01}} approach to binary dynamics, which maps the dynamics of a binary to that of a point particle moving in an ``effective'' spacetime corresponding to a deformed black hole, has been found to outstandingly describe the outcome of numerical relativity calculations after some adjustable parameters in the EOB framework are calibrated to numerical calculations {\\cite{eob_num1, eob_num2, eob_num3, eob_num4}}. Our understanding of the two-body problem in general relativity has never been better. These efforts are largely motivated by the need for accurate models of coalescing black holes to detect and measure merger signals in the data of gravitational-wave (GW) detectors. Black holes with masses of roughly $10^6 - 10^9 \\,M_\\odot$ indisputably reside at the cores of essentially every galaxy with a central bulge {\\cite{kg01,ferrarese02}}. In the hierarchical growth of structure, these black holes will form binaries as their host galaxies merge and grow {\\cite{bbr1980}}; estimates of how often such binaries form indicate that the proposed space-based detector LISA {\\cite{lisa_nasa,lisa_esa}} should be able to measure at least several and perhaps several hundred coalescences over a multiyear mission lifetime {\\cite{bbh_rate_lisa}}. There is already a catalog of candidate binaries in this mass range, such as active galaxies with double cores \\cite{komossa2003, maness2004, rodrigues2006}, systems with doubly-peaked emission lines \\cite{zhou2004, gerke2007}, and systems that appear to be periodic or semi-periodic, such as the blazar OJ287 \\cite{valtonen2007}. The last year or so of the binary's life will generate GWs at frequencies to which LISA is sensitive; measuring those waves will make it possible to precisely map the distribution of cosmic black hole masses and spins, opening a new observational window onto the high-redshift growth of cosmic structure. Less massive black hole binaries (several to several hundred $M_\\odot$) will be targets for the ground-based GW detector network, currently including LIGO {\\cite{ligo}}, Virgo {\\cite{virgo}}, and GEO {\\cite{geo}}, and hopefully including the proposed detectors LCGT {\\cite{lcgt}}, AIGO {\\cite{aigo}}, and the ``Einstein Telescope'' {\\cite{et}} in the future. Formation scenarios and event rate estimates in this band are much less certain, since the demographics of the relevant black holes and scenarios for them to form binaries are not as well understood as in the supermassive range. However, scenarios involving dynamic binary formation in dense clusters suggest that this network can plausibly expect an interesting event rate {\\cite{bbh_rate_ligo1, bbh_rate_ligo2, bbh_rate_ligo3, bbh_rate_ligo4, bbh_rate_ligo5}}, strongly motivating the construction of binary merger models for these detectors. Finally, moving back to the LISA band, binaries in which one member is much less massive than the other are expected to be an important source. Such extreme mass ratio binaries are created when a stellar mass secondary ($\\sim 1 - 100\\,M_\\odot$) is scattered through multibody interactions onto a highly relativistic orbit of a roughly $10^6\\,M_\\odot$ black hole in the center of a galaxy. Though rare on a galaxy-by-galaxy basis, enough galaxies will be in the range of LISA that the number measurable is expected to be several dozen to several hundreds {\\cite{emri_rate}}. The waves from these binaries largely probe the quiescent spacetimes of their larger (presumably Kerr) black hole, making possible precision tests of the strong-field nature of black hole spacetimes {\\cite{emri_science}}. In short, astrophysical binary black holes will come in a wide range of mass ratios. Computational models must be able to handle systems with mass ratios ranging from near unity, to millions to one. Each mass $m$ sets a lengthscale $Gm/c^2$ which the code must be able to resolve. Large mass ratios require codes that can handle a large dynamic range of physically important lengthscales. Perturbation theory is an excellent tool for modeling binaries with very large mass ratios. In this limit, the binary's spacetime is nearly that of its largest member, with the smaller member acting to distort the metric from the (presumably) exact Kerr solution of that ``background.'' It is expected that tools based on perturbation theory will be crucial for modeling extreme mass ratio systems described above (mass ratios of $10^4:1$ or larger). Even for less extreme systems, perturbative approaches are likely to contribute important wisdom, working in concert with tools such as numerical relativity and the effective one-body approach. The foundational examples of such an analysis are the papers of Nagar, Damour, and Tartaglia {\\cite{ndt07}} and of Damour and Nagar {\\cite{dn07}}. In that work, the EOB framework is used to construct the quasi-circular late inspiral and plunge of a small body into a non-rotating black hole. Regge-Wheeler-Zerilli methods {\\cite{rw57,z70}} are then used to compute the GWs that arise from a small body that follows that trajectory into the larger black hole. Those authors use this large mass ratio system as a ``clean laboratory'' for investigating binary dynamics, and advocate using these techniques as a tool for probing delicate issues such as the form of the waves which arise from the plunge, and the matching of the final plunge waves to the late ringdown dynamics of the system's final black hole. Our goal here is to develop a similar toolkit based on perturbation theory applied to spinning black holes. We have developed two perturbation theory codes which we use to model different aspects of binary coalescence. Both codes solve the Teukolsky equation {\\cite{teuk}}, computing perturbations to the curvature of a Kerr black hole. One code works in the frequency domain {\\cite{h2000, dh2006}}, which works well for computing the averaged flux of quantities such as energy and angular momentum carried by GWs. The other code works in the time domain {\\cite{skh07, skhd08}}, which is excellent for calculating the aperiodic GW signature of an evolving source. As originally proposed in Ref.\\ {\\cite{hdff05}}, we have developed a hybrid approach which uses the best features of both the time- and frequency-domain codes to model the full coalescence process. (Although our ultimate goal is to develop a set of tools similar to those developed by Damour, Nagar, and Tartaglia, we note that our techniques are the moment largely numerical, as opposed to the mixture of numerical and analytic techniques developed in Refs.\\ {\\cite{ndt07,dn07}}. It would be worthwhile to connect the work we present here to the body of EOB work, but have not yet begun doing so in earnest.) As we were completing this paper, a perturbation-theory-based analysis of binary merger was presented by Lousto et al.\\ {\\cite{lnzc}}. Their analysis does not use the Teukolsky equation, but is otherwise very similar in style and results to what we do here. In particular, they note as we do here that the perturbation equations terminate the merger waveform in a set of ringdown waves in a very natural way, thanks to the manner in which the equation's source redshifts away as the infalling body approaches the large black hole's event horizon. This behavior was also pointed out and exploited by Mino and Brink {\\cite{mb08}} in their (largely analytic) perturbative analysis of recoil from waves from the late plunge. We expand on this point in more detail at appropriate points later in the paper. As our use of the Teukolsky equation requires, we assume that a binary can be well-described by a small body moving in the spacetime of a (much larger) Kerr black hole. We first build the worldline that the smaller body follows as it slowly inspirals and then plunges into the black hole. We assume that, early in the coalescence, the small body moves on a geodesic of the background Kerr spacetime. Using the frequency-domain perturbation theory code to compute their rates of change, we allow the energy $E$, angular momentum $L_z$, and Carter constant $Q$ of this configuration to evolve. (In fact, we confine ourselves to equatorial orbits in this analysis, so $Q = 0$ throughout the binary's evolution.) This drives the smaller body in an adiabatic inspiral through a sequence of orbits, until we approach the last stable orbit of the large black hole\\footnote{At present, we do not include the conservative impact of self forces. These forces {\\it are} included in the EOB-framework analyses of Damour, Nagar, and Tartaglia.}. We then make a transition to a plunging orbit, using the prescription of Sundararajan {\\cite{s2008}} which in turn generalized earlier work by Ori and Thorne\\footnote{A similar approach to the transition from inspiral plunge, but valid for arbitrary mass ratios and presented using the EOB framework, was developed by Buonanno and Damour {\\cite{bd2}}, and appeared in press before Ref.\\ {\\cite{ot2000}}; we thank T.\\ Damour for clarifying this chronology to us. This approach is used in Refs.\\ {\\cite{ndt07,dn07}} to compute the transition from the slow, adiabatic inspiral to plunge.} {\\cite{ot2000}}. By properly connecting the adiabatic inspiral to a plunge, we make a full worldline describing the small body's coalescence with the larger black hole. This worldline gives us the source for our time-domain perturbation theory code, from which we compute the GWs generated by the system as the small body evolves from the (initially near geodesic) inspiral through the plunge and merger. The waves which we compute in this way have qualitatively the same ``inspiral, merger, ringdown'' structure seen in numerical relativity simulations, though much work remains to quantify the degree of overlap. As an illustration of the utility of our perturbative toolkit, we focus in this paper on the problem of GW recoil. Studies of GW recoil have been particularly active in recent years; we review this problem and its literature in the next subsection. \\subsection{Gravitational-wave recoil} The asymmetric emission of GWs from a source carries linear momentum. The system then recoils to enforce global conservation of momentum. Early work demonstrated the principle of this phenomenon \\cite{br1961,peres1962}; Bekenstein \\cite{b1973} appears to have been the first to appreciate the important role it could play in astrophysical problems. Much recent work has focused on the recoil imparted to the merged remnant of binary black hole coalescence. The first estimates of binary black hole kick were made by Fitchett \\cite{fitchett}. He treated the gravitational interaction as Newtonian and included the lowest order mass and current multipoles needed for GW emission to compute the recoil velocity. This early calculation predicted that recoil velocities could approach thousands of km/s, which is greater than the escape velocity for many galaxies. Because of his restriction to low-order radiation formulas, and his use of Newtonian gravity to describe binary dynamics, it was clearly imperative that Fitchett's calculations be revisited; a prescient analysis by Redmount and Rees {\\cite{redmountrees}} particularly argued for the need to account for the effect of black hole spins in the coalescence. Over the past several years, quite a few calculations have substantially improved our ability to model the recoil in general relativity. The various approaches can be grouped as follows: \\begin{itemize} \\item {\\it Black hole perturbation theory}. As discussed extensively above, black hole perturbation theory is a good tool for describing binaries involving a massive central black hole (of mass $M$) and a much less massive companion (of mass $\\mu$). Shortly after Fitchett's pioneering binary calculation, Fitchett and Detweiler examined whether strong-field gravity changed the conclusions using perturbation theory {\\cite{fd84}}. Twenty years later, Favata, Hughes, and Holz \\cite{FHH} argued that, properly extrapolated, reasonable results can be obtained for quantities such as the integrated black hole kick up to a mass ratio $\\mu/M \\sim \\mathcal O(0.1)$. Unfortunately, the Favata et al.\\ analysis has a rather large final error since the frequency-domain tools they use do not work well at modeling the GWs arising from the final plunge of the smaller body into the large black hole. One of our goals in this analysis is to revisit that calculation and reduce those substantial error bars. Another application of perturbation theory is the ``close-limit approximation,'' {\\cite{closelimit}} which describes the last stages of a merging binary as the dynamics of a distorted single black hole. Sopuerta, Yunes, and Laguna \\cite{syl06} applied the close-limit approximation to describe the final waves from unequal mass binaries, obtaining results that compare very well with those that have since been computed within ``full'' numerical relativity. Finally, Mino and Brink {\\cite{mb08}} used perturbative techniques to model the waves from the plunge, quantifying the manner in which the geometry of the final infall impacts the kick imparted to the binary. As already mentioned, their analysis also took advantage of the manner in which the source redshifts away as the infalling body approaches the larger black hole's event horizon. \\item {\\it Post-Newtonian (PN) theory}. PN theory describes the spacetime and the motion of bodies in the spacetime as an expansion in the Newtonian gravitational potential $Gm/rc^2$ (where $m$ is a characteristic system mass, and $r$ a characteristic black hole separation). Blanchet, Qusailah, and Will \\cite{bqw05} used an approach based on this expansion to substantially improve estimates from the recoil from the final plunge and merger; though consistent with the results from {\\cite{FHH}}, they were able to reduce the error bars by a substantial factor. More recently, Le Tiec, Blanchet, and Will {\\cite{ltbw10}} combined a post-Newtonian inspiral with a close-limit computation of the merger and ringdown to compute the recoil for the coalescence of non-spinning black holes. This analysis is quite similar in spirit to the one we present here, though it does not use perturbation theory throughout. \\item {\\it Numerical relativity}. Not long after it first became possible to model the coalescence of two black holes in numerical relativity, this became the technique of choice for computing black hole recoil. No other technique is well-suited to computing wave emission and spacetime dynamics for very asymmetric, strong-field configurations which are likely to produce strong GW recoils. Numerical relativity was needed to discover the so-called ``superkick'' configuration: an alignment of spin and orbital angular momentum which results in a kick of several thousand kilometers per second {\\cite{superkick0,superkick1,superkick2}}. In most configurations, the kick tends to be substantially smaller, peaking at a few hundred kilometers per second \\cite{NR1,NR2,NR3}. \\item {\\it Effective one-body}. As already described, EOB describes a binary as a test body orbiting in the spacetime of a ``deformed'' black hole, with the deformation controlled by factors such as the mass ratio of the binary. Damour and Gopakumar {\\cite{dg06}} first examined the issue how to compute recoil within the EOB framework, analytically identifying the major contributions to the recoil that accumulates over a coalescence, including the importance of the final merger and recoil waves in providing an ``antikick'' contribution. By calibrating some parameters of the EOB framework with results from numerical relativity, EOB has had great success generating waveforms and recoil velocities that match well with those from numerical relativity \\cite{EOB_orig,EOB}. \\end{itemize} With the exception of the ``superkick'' configuration, all of these techniques predict recoils that peak at roughly a few hundred kilometers per second (depending on mass ratio, spins, and spin orbit orientation; see {\\cite{statistics}} for detailed discussion and statistical analysis). This is substantially lower than the peak predicted by Fitchett's original calculation; his overestimate can be ascribed to neglect of important curved spacetime radiation emission and propagation effects. In addition to their potential astrophysical applications, recoil computations serve another important purpose: They are a common point of comparison for these four approaches to strong field gravity. The recoil velocity from a merging binary is calculated by integrating the emitted radiation over some number of orbits. Any significant systematic error in the approach used used will tend to magnify the error in the estimated recoil velocity. Thus, the evaluated recoils for a range of BH spins and mass ratios serve as a good platform for comparing various approaches to strong field binary models. \\subsection{This paper} Our goal is to revisit and improve the estimate of black hole recoil via black hole perturbation theory that was originally developed in Ref.\\ {\\cite{FHH}}. That analysis predicts upper and lower bounds which are rather widely separated. This is because the analysis of {\\cite{FHH}} could not accurately model wave emission from the final plunge and merger. Using the time-domain perturbation theory code developed and presented in Refs.\\ {\\cite{skh07,skhd08}}, we can now compute the contribution of those waves. As we describe in more detail in Sec.\\ {\\ref{sec:results}}, doing so completely reverses the conclusions of Ref.\\ {\\cite{FHH}} regarding how the kick behaves as a function of spin. In particular, including the plunge and merger is crucial to correctly computing the ``antikick,'' the out-of-phase contribution to the recoil that arises from the merger's final GWs. This contribution to a binary's total recoil was first identified and characterized by Schnittman et al.\\ {\\cite{antikick}}. We find that the inability to include this contribution in Ref.\\ {\\cite{FHH}} is largely responsible for the large error bars in that analysis. We begin by reviewing in Sec.\\ {\\ref{sec:trajectory}} how we construct the worldline which the smaller member of our binary follows as it spirals into the larger black hole. As briefly described above, we break this trajectory into a slowly evolving ``inspiral'' (Sec.\\ {\\ref{sec:inspiral}}) followed by a transitional regime (Sec.\\ {\\ref{sec:transition}}) that takes the binary into a final plunge and merger (Sec.\\ \\ref{sec:plunge}). This review is left general, so that in principle one could describe this dynamics for generic orbital geometry. We specialize in our analysis here to the simplest circular and equatorial orbits (Sec.\\ {\\ref{sec:circeq}}). We next briefly review how we compute gravitational radiation from a body moving on this trajectory. As mentioned above, our approach is based on finding solutions to the Teukolsky equation {\\cite{t73}} for Kerr black hole perturbations. We review this equation's general properties in Sec.\\ {\\ref{sec:radiation}}, and then discuss the principles behind solving it in the frequency domain (Sec.\\ {\\ref{sec:freqdomain}}) and in the time domain (Sec.\\ {\\ref{sec:timedomain}}). Section {\\ref{sec:recoil}} summarizes how one computes the radiation's linear momentum and the recoil of a merged system. Section {\\ref{sec:results}} presents the results of our analysis. We begin in Sec.\\ {\\ref{sec:massratio}} with general considerations on how our results scale with mass ratio. Because we work strictly within the context of linearized perturbation theory, all of our results can be easily scaled to different mass ratios, provided that the scaling does not change the system so much that the validity of perturbation theory breaks down. Reference {\\cite{FHH}} argued that a modified scaling would allow us to estimate with reasonable accuracy quantities related to the recoil even out of the perturbative regime. Although those arguments are valid during the adiabatic inspiral, they break down when the members of the binary merge. In Sec.\\ {\\ref{sec:waves}}, we then discuss in some detail the gravitational waveform we find for binary coalescence in the large mass ratio limit. We examine the different multipolar contributions to the last several dozen cycles of inspiral, followed by the plunge and merger. These examples illustrate the manner in which the coalescence waves very naturally evolve into a ``ringdown'' form. As discussed in some detail in Sec.\\ {\\ref{sec:timedomain}}, this behavior arises by virtue of how the Teukolsky equation's source term goes to zero, so that its solutions transition to their homogeneous form, as the infalling body approaches the large black hole's event horizon. Mino and Brink {\\cite{mb08}} first appear to have exploited this behavior, which was also seen in recent work by Lousto, Nakano, Zlochower, and Campanelli {\\cite{lnzc}}. This demonstrates the power of perturbative methods at modeling physically important aspects of the merger waves. Section {\\ref{sec:spin}} examines the recoil that arises from these waves, focusing on how it depends (for the circular, equatorial case that we study) on the spin of the larger black hole. This analysis demonstrates very clearly the impact of the ``antikick'' first reported by Schnittman et al.\\ {\\cite{antikick}}. With the antikick taken into account, the smallest recoils come from the largest spins when the merger is a prograde sense; the largest spins come from retrograde mergers with large spins. The waves which give the system its antikick come from those produced by the final plunge and merger, demonstrating very clearly the substantial impact these waves have on the system. We conclude this section by briefly discussing the convergence of our recoil results as a function of black hole spin. Interestingly, we find that the number of modes we must include in order for our results to converge is a strong function of the black hole's spin --- rapid spin, prograde cases need more modes than do slow spin cases, which in turn need more modes than rapid spin, retrograde cases. We conclude the paper by discussing in Sec.\\ {\\ref{sec:conclude}} how these tools may be used to expand the reach of two-body modeling in general relativity, and our future plans. Throughout our analysis we generally use units in which $G = c = 1$. We sometimes use $c = 3\\times10^5\\,{\\rm km/sec}$ in order to present kicks in ``physical'' units. ", "conclusions": "\\label{sec:conclude} Now that numerical relativity has effectively solved the two-body problem in general relativity, a major task for researchers has become to explore the parameter space of binary coalescence. This will insure that wave models constructed as templates for GW data analysis fully encompass the range of behaviors that are likely in real binary mergers, and allow us to more fully understand the phenomenology of binary black hole merger astrophysics. In this analysis, we have demonstrated that perturbation theoretical techniques based on the Teukolsky equation are an excellent tool for extending the reach of our computations, allowing us to model large mass ratios that are challenging for 3+1 numerical simulations, but may be of astrophysical significance. Our analysis joins previous work by Damour and colleagues {\\cite{ndt07,dn07}}, Mino and Brink {\\cite{mb08}}, and by Lousto and colleagues {\\cite{lnzc}} which likewise used perturbation theory to model large mass ratio binaries. By using the Teukolsky equation, we can explore how the larger black hole's spin impacts the analysis, exemplified by our demonstration of how the previously identified ``antikick'' {\\cite{antikick}} strongly depends on this spin. Two directions for future analysis strike us as particularly noteworthy. First, the major motivation for this work is that perturbation theory makes exploring parameter space computationally fast and simple. As such, it would be worthwhile to continue this exploration, examining how the waveform varies as a function of spin-orbit alignment, and exploring (for example) how the antikick evolves as one varies the inclination smoothly from the prograde to the retrograde geometry. Preliminary calculations of this behavior indicate that the antikick rapidly evolves with spin-orbit alignment, consistent with the results of Mino and Brink {\\cite{mb08}} which demonstrate a strong dependence on the final kick with the plunge geometry. Second, as Damour and Nagar have emphasized {\\cite{dn07}}, particularly useful application comes by including input from the effective one-body formalism in our description of the small body's motion; input from perturbation theory can likewise be used to calibrate certain parameters in the EOB framework. Now that the spin-augmented Hamiltonian for binary systems is understood {\\cite{djs08,brb09}}, we expect that work to extend EOB to more broadly include the impact of spin will become very active. We hope that the tools we have presented here will be useful for further refining what has already proved to be a valuable tool for modeling coalescing binaries." }, "1003/1003.0020_arXiv.txt": { "abstract": "We present a detailed study of a peculiar source detected in the COSMOS survey at z=0.359. Source CXOC~J100043.1+020637, also known as CID-42, presents two compact optical sources embedded in the same galaxy. The distance between the two, measured in the HST/ACS image, is 0.495\"$\\pm$0.005\" that, at the redshift of the source, corresponds to a projected separation of 2.46$\\pm$0.02 kpc. A large ($\\sim$1200 km/s) velocity offset between the narrow and broad components of H$\\beta$ has been measured in three different optical spectra from the VLT/VIMOS and Magellan/IMACS instruments. CID-42 is also the only X-ray source in COSMOS having in its X-ray spectra a strong {\\it redshifted} broad absorption iron line, and an iron emission line, drawing an {\\it inverted} P-Cygni profile. The \\chandra\\ and \\xmm\\ data show that the absorption line is variable in energy by $\\Delta$E=500 eV over 4 years and that the absorber has to be highly ionized, in order not to leave a signature in the soft X-ray spectrum. That these features, the morphology, the velocity offset and the {\\it inverted} P-Cygni profile, occur in the same source is unlikely to be a coincidence. We envisage two possible explanations, both exceptional, for this system: (1) a gravitational wave recoiling black hole (BH), caught 1-10 Myr after merging, (2) a Type 1/ Type 2 system in the same galaxy where the Type~1 is recoiling due to slingshot effect produced by a triple BH system. The first possibility gives us a candidate gravitational waves recoiling BH with both spectroscopic and imaging signatures. In the second case, the X-ray absorption line can be explained as a BAL-like outflow from the foreground nucleus (a Type 2 AGN) at the rearer one (a Type 1 AGN), which illuminates the otherwise undetectable wind, giving us the first opportunity to show that fast winds are present in obscured AGN, and possibly universal in AGNs. ", "introduction": "Double super-massive black holes (SMBHs) within a single galaxy are predicted from the combination of hierarchical models of structure formation (see Colpi \\& Dotti 2009 for a recent review), and the observed link between black hole (BH) mass and host galaxy bulge mass in the local Universe (e.g., Magorrian et al. 1998, Kormendy \\& Gebhardt 2001). Kiloparsec scale double nuclei should be common at redshift z$\\sim$2 where the merging process is very efficient (Springel et al. 2005) but at this high z they are hard to resolve. At later epochs (z$<$0.7) they are easier to find, but the merger rate is lower (Hopkins et al. 2008) and binary SMBHs should be relatively scarce. For these reasons, observational evidence for binary active SMBHs at kpc separations remains sparse. A few low redshift examples have been found. The most famous double active SMBHs are the highly obscured pair in the Ultra-luminous Infrared galaxy (ULIRG) NGC 6240 (Komossa et al. 2003; z=0.0245). \\chandra\\ imaging and spectroscopy of NGC 6240 show the presence of two active galactic nuclei (AGNs) separated by just 1.2 kpc. Only few other AGN pairs at kiloparsecs separations have been discovered so far, all by chance: Arp 299 (Ballo et al. 2004; z=0.01), Mrk 463 (Bianchi et al. 2008; z=0.05), 3C 75 (Hudson et al. 2006; z=0.02), and 0402+379 (Rodriguez et al. 2006; z=0.05). A candidate sub-parsec binary system at z=0.38 has been recently discovered in the SDSS through the detection of two emission line systems in the optical spectrum (Boroson \\& Lauer 2009), but this interpretation is still under debate (e.g., Chornock et al. 2010). Once the SMBH binary tightens, its merger is characterized by the emission of gravitational waves (GW) and the merged BH under favorable conditions (of galaxy mass ratio and BH spins) can recoil with respect to the center of the galaxy (Peres 1962, Bekenstein 1973) at substantial velocities (few$\\times$10$^2$-10$^3$ km/s, Campanelli et al. 2007a). If not enough mass is driven to the SMBH binary before the merger, it can stall until the arrival of a third galaxy induces the hardening of the binary which ejects the just arrived BH for gravitational slingshot effect (Saslaw et al. 1974). Recently, observational searches were initiated for single quasars which are displaced from their host galaxy spatially or spectroscopically due to GW recoil after the merger (e.g., Bonning et al. 2007). Only a few candidates have been reported so far (Komossa et al. 2008, Shields et al. 2009), all of them discovered through the detection of multiple line systems in optical spectra, but more prosaic interpretations for these candidates are not yet ruled out. The quasar HE 0450-2958 is candidate to be a BH ejected by the nucleus of a companion galaxy during a major merger (Magain et al. 2005, Hoffman \\& Loeb 2006), but questions remain about the nature of this source (e.g., Merritt et al. 2006, Kim et al. 2007, Jahnke et al. 2009a). A survey for binary SMBH systems or the discovery of many recoiling BH would be helpful to test evolutionary merger models and to put more robust constraints on our understanding of BH and galaxy co-evolution (e.g., Hopkins et al. 2008). The Cosmic Evolution Survey (COSMOS; Scoville et al. 2007a) is a deep multiwavelength survey of 2 deg$^2$, based on a 590-orbit ACS program (Koekemoer et al. 2007), the largest mosaic ever made by the HST. Of the $10^6$ galaxies to $i_{AB}\\sim 27$ detected in the full field, the $i$-band (F814W) ACS data provide morphological information down to $i_{AB} < 23$ for 70,000 galaxies and enables detailed structural analysis out to $z \\sim 0.8$ for $\\sim$16,000 galaxies. The large number of spectra ($>$20,000, Lilly et al. 2007, 2009, Trump et al. 2007) available for galaxies and quasars allow to search for sources with more than one line system, although the typical resolution available ($R=600-700$) only allows sources with high velocity shifts to be found. The rich multiwavelength COSMOS database allows detailed study of the source SEDs at all frequencies from radio to X-ray. These properties make COSMOS an excellent database in which to look for sources just before or after the SMBHs binary merging. According to the prediction of Volonteri \\& Madau (2008), the number of detectable off-center AGN (up to 30 kpc away from the center of its host) in the HST COSMOS area should be $\\sim$30, for the best case of large kicks (spinning holes), long decay timescales (no bulge), and long active phase, while considering those X-ray emitting the expected number is $\\sim$1. In contrast, the short lifetime of binary BH at small separations implies a much lower number of sub-parsec candidates in the whole area, even at low redshift ($<$1/deg$^2$, Volonteri et al. 2009). SMBHs are easier to find when they are active, and X-ray surveys are the most efficient means for finding AGNs (e.g., Brandt \\& Hasinger 2005). Hence, we started our survey for pre- and post-binary candidates with the optical counterparts of X-ray sources within COSMOS. A visual inspection of the optical counterparts of all 2600 X-ray sources, detected in the \\chandra-COSMOS (C-COSMOS, Elvis et al. 2009) and XMM-COSMOS (Hasinger et al. 2007) surveys, using the HST/ACS F814W images (Koekemoer et al. 2007; Leauthaud et al. 2007), led to the discovery of just one candidate recoil SMBH in the source CXOC~J100043.1+020637 (CID-42, Elvis 2009) at z=0.359. At the same time, Comerford et al. (2009) reported the discovery of dual AGNs in the same source, through the detection of two emission line systems in a Keck/DEIMOS spectrum. We report here a detailed study of CID-42, including the SED analysis and optical and X-ray imaging and spectroscopy. We discuss several models for the nature of this system. In this paper, magnitudes are reported in the AB system (Oke \\& Gunn 1983) and a WMAP 5-year cosmology\\footnote{ http://lambda.gsfc.nasa.gov/product/map/dr3/params/lcdm\\_sz\\_lens\\_wmap5.cfm} with H$_0$ =71~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_M$ = 0.26 and $\\Omega_{\\Lambda}$=0.74 is assumed. ", "conclusions": "We have reported the multiwavelength properties of CID-42, which presents three unusual features: two close sources, embedded in the same galaxy, resolved in the optical ACS image, but unresolved in the X-ray one, a high velocity offset, measured between the broad and narrow H$\\beta$ lines in 3 different optical spectra and a {\\it redshifted} absorption iron line in the X-ray spectrum. Thanks to the rich database of the COSMOS survey, the analysis of the properties of CID-42 has been performed at all wavelengths. The overall analysis suggests two possible explanations for the source: a GW recoiling BH, caught $\\sim$few Myr after the kick from the center of the galaxy or a Type 1/Type 2 AGN system where the Type 1 is recoiling due to a slingshot effect. Analysis of the ACS image shows the presence of a SE point-like source, playing the role of the Type 1 AGN (or SMBH with disc and BLRs) in both scenarios, confirmed by the X-ray flux variability across the time and by the offset in velocity measured between the narrow and the broad components of the H$\\beta$ line. The NW source, being less compact than the SE one, could be either the naked core of the galaxy from which a BH has been kicked out or a Type 2 AGN. The presence of a {\\it redshifted} absorption iron line changing its energy centroid in the different observations, allows to explain the geometry of the system by inflow of material into the BH in the GW recoil BH case or with the ``Backlit Wind'' model in the second scenario. Monitoring observations in the X-ray would be suitable to study the moving X-ray absorption feature and its variability. The ``Backlit Wind'' implies that fast BAL-like winds are present in Type 2 AGNs, an otherwise untestable hypothesis. The variability of the iron absorption is, in this special case, a new tool to study BAL flows. This model predicts corresponding high ionization UV absorption lines (e.g. OVI), making CID-42 a suitable target for HST/COS observations. \\chandra\\ high-resolution observations with HRC could resolve the two sources to see if both are X-ray emitting or not, while optical and IR spectroscopy with higher spatial resolution, which currently is only possible with HST, could confirm the still uncertain velocity shift measured in H$\\beta$ in other lines. If future observations allow us to confirm that the GW recoiling BH is the best explanation of the system, it will be the first time in which this phenomenon is observed via both spatially offset from the galactic core and a velocity offset in the optical spectrum. Although models for this system have been sketched qualitatively in this paper, a numerical analysis will be the subject of a forthcoming paper. Upcoming Keck Integral Field Unit (IFU) observations with OSIRIS will help us to disentangle major competing models mapping the inner region of CID-42. IFU observations on larger scales could map the velocity of gas and stars in the galaxy and in the tails, in order to constrain the theoretical models on the geometry of the merger and the inclination of the galaxies in the first encounter (D'Onghia et al. in prep.). Calibrating the number of recoiling BH, produced either in close triple encounters and by gravitational waves ejection could be of great importance for the proposed Laser Interferometer Space Antenna (LISA) mission (see discussion in Loeb 2009). CID-42 is a unique source with a cluster of rare features. Whichever is the best explanation, in CID-42 we are witnessing a runaway BH." }, "1003/1003.3486_arXiv.txt": { "abstract": "The nonlinear memory effect is a slowly-growing, non-oscillatory contribution to the gravitational-wave amplitude. It originates from gravitational waves that are sourced by the previously emitted waves. In an ideal gravitational-wave interferometer a gravitational-wave with memory causes a permanent displacement of the test masses that persists after the wave has passed. Surprisingly, the nonlinear memory affects the signal amplitude starting at leading (Newtonian-quadrupole) order. Despite this fact, the nonlinear memory is not easily extracted from current numerical relativity simulations. After reviewing the linear and nonlinear memory I summarize some recent work, including: (1) computations of the memory contribution to the inspiral waveform amplitude (thus completing the waveform to third post-Newtonian order); (2) the first calculations of the nonlinear memory that include all phases of binary black hole coalescence (inspiral, merger, ringdown); and (3) realistic estimates of the detectability of the memory with LISA. ", "introduction": "We often think of gravitational-wave (GW) signals as having an oscillatory amplitude that starts small at early times, builds to some maximum, and then decays back to zero at late times. For example, this is the standard picture of a waveform from a coalescing compact-object binary. However, this picture is incomplete. In reality, \\emph{all gravitational-wave sources} possess some form of \\emph{gravitational-wave memory}. The GW signal from a `source with memory' has the property that the late-time and early-time values of at least one of the GW polarizations differ from zero: \\begin{equation} \\label{eq:memdef} \\Delta h_{+,\\times}^{\\rm mem} = \\lim_{t\\rightarrow +\\infty} h_{+,\\times}(t) - \\lim_{t\\rightarrow -\\infty} h_{+,\\times}(t), \\end{equation} where $t$ is time at the observer. When a GW \\emph{without memory} passes through a detector, it causes oscillatory deformations but eventually returns the detector to its initial state. After a GW \\emph{with memory} has passed through an \\emph{idealized} detector (one that is truly freely-falling), it causes a permanent deformation---leaving a `memory' of the waves' passage. High-frequency detectors like bars or LIGO are rather insensitive to the memory from most sources because the detector response timescale is generally much shorter than the \\emph{rise-time} of typical memory signals (the characteristic time for the non-oscillatory piece of the GW signal to build up to its final value). A detector like LISA is better able to detect the memory because of its good sensitivity in the low-frequency band where typical memory sources are stronger.\\footnote{It has recently been realized that pulsar timing arrays \\cite{sazhin-1978-pulsartiming,detweiler-ApJ1979-pulsartiming,hobbs-etal-amaldi8review} (which operate at much lower frequencies than LISA) could also be sensitive to the memory effect \\cite{seto-memory-MNRAS09,levin-vanHaasteren-memory,pshirkov-etal-memory}, although the expected event rates from supermassive black hole binaries are small.} Note also that bars and ground-based interferometers like LIGO are incapable of `storing' a memory signal because they are not truly free: internal elastic forces push a bar back to its equilibrium shape, and magnets on the LIGO test masses (as well as its pendulum wires) push them back to their equilibrium positions. Because its proof-masses are truly freely-floating, a detector like LISA could maintain a permanent displacement. However, this late-time displacement caused by the GW memory is not directly observable without information on the prior state of the detector: the spacetime metric near a detector long after a GW with memory has passed is equivalent to flat spacetime in non-standard coordinates (see Sec.~VD of \\cite{favata-pnmemory} for details). Rather, it is the build-up of the memory (the difference in the metric between late and early times) that \\emph{is} observable. The memory effect has been known since the 1970's \\cite{braginskii-grishchuk,zeldovich-polnarev,braginskii-thorne} in its linear form. The linear memory generally arises in systems with unbound components: a binary on a hyperbolic orbit (two-body scattering) \\cite{turner-unbound}, matter or neutrinos ejected from a supernova \\cite{epstein-neutrinomemory,turner-neutrinomemory,burrows-hayesPRL96}, or gamma-ray burst jets \\cite{sago-GRBmemory}. In the 1990's a nonlinear form of memory was discovered independently by Blanchet \\& Damour \\cite{blanchet-damour-hereditary} and Christodoulou \\cite{christodoulou-mem}. The nonlinear memory arises from the contribution of the emitted GWs to the changing quadrupole and higher mass moments. As discussed by Thorne \\cite{kipmemory}, the nonlinear memory can be described in terms of a linear memory in which the unbound masses are the individual radiated gravitons. This implies that nearly all GW sources are sources with memory (even if the component objects remain bound). Because the memory is a non-oscillatory effect with poor observational prospects for LIGO and other ground-based interferometers, it has been largely ignored by the GW community. But for the most important (or most studied) GW source---quasi-circular compact binaries---the nonlinear memory has quite a large contribution to the time-domain waveform amplitude: in a post-Newtonian (PN) expansion of the waveform polarizations, the memory effect enters at leading-(Newtonian)-order! That the memory enters at such low PN order is related to the fact that it is a \\emph{hereditary} effect---the memory amplitude at any retarded time depends on the entire past motion of the source (and not just on the source's instantaneous retarded-time configuration). In addition, the nonlinear memory is a unique nonlinear effect because its non-oscillatory nature makes it distinctly visible in the waveform. For these reasons the memory should be studied further and its prospects for detection reassessed. In the rest of this article I will briefly review the linear and nonlinear memory effects. I will then summarize my previous work in three areas: (i) computing the post-Newtonian memory corrections to the GW polarizations; (ii) calculating the evolution and saturation of the memory during the merger and ringdown of coalescing black hole (BH) binaries; and (iii) estimating the ability of LISA to detect the memory from supermassive BH binary mergers. Except for the presentation of the linear memory from hyperbolic binaries in Sec.~\\ref{sec:linmem}, this conference proceeding concisely summarizes the results of references \\cite{favata-pnmemory,favata-memory-saturation} as presented at the $8^{\\rm th}$ Edoardo Amaldi Conference on Gravitational Waves. Readers are referred to those references for a more detailed exposition. ", "conclusions": "" }, "1003/1003.3347_arXiv.txt": { "abstract": "A wideband analog correlator has been constructed for the Yuan-Tseh Lee Array for Microwave Background Anisotropy. Lag correlators using analog multipliers provide large bandwidth and moderate frequency resolution. Broadband IF distribution, backend signal processing and control are described. Operating conditions for optimum sensitivity and linearity are discussed. From observations, a large effective bandwidth of around 10 GHz has been shown to provide sufficient sensitivity for detecting cosmic microwave background variations. ", "introduction": "Interferometric observations have gained much popularity in the study of the cosmic microwave background (CMB) anisotropy (White et al. 1999; Carlstrom et al. 2002; Padin et al. 2002; Leitch et al. 2002; Scaife et al. 2008), due to their advantage in stability and phase preserving characteristic via the heterodyne scheme for correlation and polarization observations. The cross correlations used in an interferometric array can effectively suppress many systematics. To achieve high brightness sensitivity, broadband low noise receivers and matching correlators are the two most important components for a continuum interferometer. The advance in millimeter and microwave detector technologies in recent years has produced very broadband components. Low-noise amplifiers (LNAs) with bandwidths of 10 GHz or more are easily accessible with noise performance comparable to bolometric direct detectors, e.g. ACBAR (Runyan et al. 2003) with bandwidths of 30 - 40 GHz and an equivalent noise temperature or RJ temperature loading $T_{RJ}$ between 40 K and 70 K. The CMB interferometers are therefore very competitive at millimeter wavelengths as compared to bolometers which are usually preferred at shorter wavelengths. The Yuan-Tseh Lee Array for Microwave Background Anisotropy (AMiBA) (Lo et al. 2001; Ho et al. 2009) is a radio interferometric array for the study of the CMB at 3mm wavelength. AMiBA detects the minute deviations of the nearly constant CMB temperature over the sky, and can study the spatial variation of this temperature fluctuation. In particular, AMiBA is imaging galaxy clusters via the Sunyaev-Zel'dovich effect (SZE) (Sunyaev and Zel'dovich 1970, 1972; Birkinshaw 1999) for the first time at 3 mm wavelength. The array presently consists of 13 elements of 1.2 m reflectors distributed on a 6-m diameter platform. The receiver attached to each reflector is dual-polarization, and equipped with cryogenically cooled LNAs operating in the 84 to 104 GHz frequency range. The intermediate frequency (IF) is from 2 to 18 GHz, and is matched with a wide-band analog correlator. The strong interest in CMB observations has motivated the development of very broadband correlators with a limited spectral resolution. Utilizing a filter bank and complex correlators, Padin et al (2002) built an analog correlator with 10 GHz bandwidth. Harris and Zmuidzinas (2001) worked on a different approach toward broadband operations by adopting an analog lag correlation scheme to build an auto correlator with a 4 GHz bandwidth. The AMiBA correlator is also based on the concept of the lag correlator, and is designed to correlate the entire 16 GHz bandwidth. Recently Holler et al. (2007) also reported a lag correlator for AMI with a 6 GHz bandwidth. In principle, analog correlators can achieve better sensitivity over their digital counterparts due to the absence of the digitization process. Analog multipliers can easily achieve high sensitivity over multi-octave frequency ranges. In comparison the use of analog-to-digital converters (ADCs) is expensive and has limited bandwidth. Thus the analog approach is preferred for observations which require high sensitivity but modest spectral resolution. However, the major difficulty in making a broadband system lies in the distribution and processing of the multi-octave signals coming from the antennas. Broadband impedance matching between components presents a major technical challenge for integrating a large scale microwave system. Due to the non-linear responses of an analog system, applying appropriate drive power levels and modulation/demodulation techniques to minimize the effects of the spurious terms are also very important. The AMiBA correlator is our attempt to address these technical issues. The 4-lag analog correlator has a nominal 16 GHz bandwidth. This is currently the only correlator in operation with an effective correlation bandwidth of around 10 GHz. The scientific goals and design philosophy of AMiBA are presented in Ho et al. (2009). A broader description of the AMiBA detection system is given in Chen et al. (2009) while a detailed description of the AMiBA hexapod mount can be found in Koch et al. (2009). From 2007 to 2008, observations were carried out with the 7-element array equipped with 60 cm dishes (Koch et al. 2006). Details on the observations and analysis of six massive galaxy clusters are presented in Wu et al. (2009). Subaru weak lensing data of 4 galaxy clusters were analyzed with the SZE data to derive the baryon fraction (Umetsu et al. 2009). This paper describes in detail the instrumental design and testing of the AMiBA correlator. We provide a system overview in Section 2 and major components of the system are described in Sections 3 to 7. In particular, we discuss various aspects of the correlator module in Section 4. In Section 8, we outline the testing and data processing. Finally, a conclusion on the system is given in Section 9. ", "conclusions": "From observations, the AMiBA correlator has proven to have the sensitivity required for CMB detection. The analog multipliers used in the lag correlators provide the wide bandwidth required for high sensitivity. The inherent noise rejection of the interferometer is also very beneficial. Compared with the filterbank scheme used by CBI (Padin et al. 2002), the lag correlator design is simpler and more compact. Large bandwidth with a small number of lags does not present a significant challenge for the lag-to-visibility transformation after proper calibrations. At the moment, the effective bandwidth we can achieve is limited by bandpass variations due to wideband impedance matching. However, a significant portion ($60\\%$) of the nominal bandwidth has been achieved. To improve on the bandpass variations, it is possible to operate the correlator at a higher IF, since the response of the analog multipliers is not limited to frequencies of few GHz. Thus, with a similar fractional bandwidth, a larger bandwidth could be achieved. Wideband complex correlators in conjunction with bandpass filters can also be considered. For AMiBA, complex correlators with a bandwidth of 8 GHz are suitable. By interleaving a number of them in frequency, a large bandwidth is feasible. The filters would also improve the isolation between channels. Similar analog interferometric systems can be constructed for high sensitivity, high angular resolution, and moderate frequency resolution observations. \\begin{table}[H] \\caption{AMiBA 4-Lag Correlator Specifications} \\begin{center} \\begin{tabular}{|l|l|} \\hline Input Frequency Range & 2 - 18 GHz \\\\ \\hline Responsivity & 80 $V_{rms}$/W minimum\\\\ \\hline Responsivity Variation vs. frequency & $<$ 3 $dB_{pp}$\\\\ \\hline Phase Response & $<$ 30 degree peak-to-peak deviation from linear\\\\ \\hline Delay Increment Accuracy & 25.0 +/- 2.5 ps \\\\ \\hline Input 1 dB Compression Point & $>$ -5.0 dBm\\\\ \\hline Squared Term Contribution & $<$ 5 $\\%$ \\\\ \\hline Output Impedance & $<$ 100 k$\\Omega$ \\\\ \\hline \\end{tabular} \\end{center} \\end{table} \\begin{table}[H] \\caption{ A typical result of the correlator output fluctuations or RMS, measured under 4 conditions. When both receivers were off, the output noise could be attributed to the backend noise from the DC amplifier and the ADC. When only one antenna was on, e.g. Ant 2, the output of the multiplier would include various terms from $V_1$. The situation was similar when only Ant 3 was on. For an ideal multiplier, the $v_1^m$ and $v_2^n$ terms are considered spurious. From the measurements, these terms also contributed to the output noise so that correlator output RMS increased. Assuming that noise from each term can be summed in quadrature, we can estimate the output fluctuation when there is no $v_1^m$ or $v_2^n$ terms at the output and no backend noise either, so as to determine how much the S/N has been degraded.} \\begin{center} \\begin{tabular}{| l | l | l | l | l | l | l |} \\hline Measurement & 1 & 2 & 3 & 4 & &\\\\ \\hline Ant 2 ($v_1$) & off & off & on & on & &\\\\ \\hline Ant 3 ($v_2$) & off & on & off & on & &\\\\ \\hline YY lag 3 RMS (counts)& 95 & 783 & 805 & 1850 & 1120 & 1473\\\\ \\hline noise & backend & $v_1^m$ terms, & $v_2^n$ terms, & overall & $v_1^m$ terms, & $v_1^m v_2^n$ terms\\\\ composition & & backend & backend & & $v_2^n$ terms, & \\\\ & & & & & backend & \\\\ \\hline \\end{tabular} \\end{center} \\end{table} \\begin{figure}[H] \\begin{center} \\includegraphics[width=4.8 in]{F1.eps} \\caption{Block diagram of the AMiBA correlator. Signal flow for one baseline (1X2X) is presented. Following the correlator module, the signal flow for a particular lag (lag 4) is depicted.\\label{fig:rxscheme}} \\end{center} \\end{figure} \\begin{figure}[H] \\begin{center} \\includegraphics[width= 6.4 in]{F2.eps} \\caption{IF power profile from the receiver output to correlator module input (prior to 2009).\\label{fig:rxscheme}} \\end{center} \\end{figure} \\begin{figure}[H] \\begin{center} \\includegraphics[width = 6.4 in]{F3.eps} \\caption{Schematic of the AMiBA receiver, including the RF, LO, and IF sections. \\label{fig:rxscheme}} \\end{center} \\end{figure} \\begin{figure}[H] \\begin{center} \\includegraphics[width=6.4 in]{F4.eps} \\caption{Schematic of the correlator IF sections. The gain, insertion loss (I.L.), noise figure (N.F.), and output 1 dB compression point (P1dB) of each component are listed where applicable. \\label{fig:rxscheme}} \\end{center} \\end{figure} \\begin{figure}[H] \\begin{center} \\includegraphics[width=4.8 in]{F5.eps} \\caption{Modular layout of the correlator 3rd section power dividers, correlator modules, DC amplifiers, and readout boards. The central portion consists of an array of lag correlators and DC amplifiers. The correlators are fed by power dividers at both ends. Readout boards underneath the horizonal power dividers are responsible for VFC ADC and integration.} \\end{center} \\end{figure} \\begin{figure}[H] \\begin{center} \\includegraphics[width=3.2 in]{F6.eps} \\caption{Picture of the MMIC IF amplifiers and power dividers used in the 3rd section of correlator IF. Bias circuit boards are not shown. \\label{fig:rxscheme}} \\end{center} \\end{figure} \\begin{figure}[H] \\begin{center} \\includegraphics[width=4.8 in]{F7.eps} \\caption{Assembly diagram for the IF amplifiers used in the 3rd section of the correlator IF as recommended by the manufacturer, except the 100 pF bypass capacitors used to remove the oscillation at high frequencies.\\label{fig:rxscheme}} \\end{center} \\end{figure} \\begin{figure}[H] \\begin{center} \\includegraphics[width=4.26 in]{F8a.eps} \\includegraphics[width=3.2 in]{F8b.eps} \\caption{Layouts of the custom 4-way power dividers\\label{fig:rxscheme}} \\end{center} \\end{figure} \\begin{figure}[H] \\begin{center} \\includegraphics[width=4.8 in]{F9.eps} \\caption{{\\it Upper Left}: Recovered spectrum for a 4-lag correlator with CW input signals at the center of 2 channels. {\\it Upper Right}: Spectrum for a 8-lag correlator {\\it Bottom Left}: Bandwidth pattern for a 4-lag correlator and delay range for 1.2m dishes with 6m baseline {\\it Bottom Right}: Bandwidth pattern for a 8-lag correlator\\label{fig:rxscheme}} \\end{center} \\end{figure} \\begin{figure}[H] \\begin{center} \\includegraphics[width=3.2 in]{F10.eps} \\caption{The 4-lag correlator module manufactured by Marki Microwave Inc. The power dividers and multipliers can be seen.} \\end{center} \\end{figure} \\begin{figure}[H] \\begin{center} \\includegraphics[width=4.0 in]{F11a.eps} \\includegraphics[width=4.0 in]{F11b.eps} \\caption{{\\it Upper}: With $v_1$ and $v_2$ as the large un-correlated noises, the output from a multiplier can be represented as the summation of various possible terms $v_1^m$, $v_2^n$, and $v_1^m v_2^n$, where $m$ and $n$ are integers. Each output terms can contribute to correlator output fluctuations. On the other hand, $v_3$ and $v_4$ present the small correlated signals, and $v_3 v_4$ is the expected product. A current noise source is used to represent the backend noise from the DC amplifier and the VFC ADC at the input of the DC amplifier. {\\it Bottom}: In the simulation of a double-balanced mixer with 4 tones, when the large signals have power below a certain level (-2 dBm), the product of 2 small signals (red curve) drops dramatically. As the input power of the large signals increase above the threshold, products of both small signals and large signals (black curve) increase linearly. As the input power keeps increasing (above 8 dBm), eventually both small-signal and large-signal products get compressed. Since a standard diode model is used during the simulations, the input power levels of the large signals where the diodes are sufficiently pumped or compressed are different from our measurements. } \\end{center} \\end{figure} \\begin{figure}[H] \\begin{center} \\includegraphics[width=3.2 in]{F12a.eps} \\includegraphics[width=3.2 in]{F12b.eps} \\includegraphics[width=3.2 in]{F12c.eps} \\includegraphics[width=3.2 in]{F12d.eps} \\caption{To determine the IF power for optimum S/N at correlator outputs, drift scans of Jupiter were taken with different input power levels to the correlator. The input power were varied by adjusting the variable gain amplifiers along the IF paths. The input S/N was fixed and could be referred to the values at the receiver inputs. Output from the 4-lag correlator module, usually referred to as the lag output, can be fitted to estimate the signal strength. During data processing, each of 168 lag outputs of the 7-element array is designated with a trace number. {\\it Upper Left}: One lag output (trace \\# 41, labelled as tr41) is plotted (green crosses) and curve-fitted (red line), in terms of counts after the VFC ADC. {\\it Upper Right}: High-pass filtering is applied to the data to remove the low-frequency terms and the signal. The noise is estimated from the RMS of the remaining fluctuations (green crosses). The responses of the lag correlators vary. {\\it Lower Left}: the S/N of several lag outputs (circles), are plotted against the RMS values. Lines are drawn through data points of each output, and different colors are used for distinction. {\\it Lower Right}: For 4 lag outputs of one baseline, output fluctuations (RMS in counts), including the backend noise, are plotted (circles) as a function of the input power.} \\end{center} \\end{figure} \\begin{figure}[H] \\begin{center} \\includegraphics[width=4.8 in]{F13.eps} \\caption{Schematic of the low frequency or \"DC\" amplifier\\label{fig:rxscheme}} \\end{center} \\end{figure} \\begin{figure}[H] \\begin{center} \\includegraphics[width=6.4 in]{F14a.eps} \\includegraphics[width=6.4 in]{F14b.eps} \\caption{{\\it Upper}: The timing diagram of the control signals and the schematic of the VFC/counter ADC used in the correlator for the 13-element array. Signals (dump, scan, reset, and demod - demodulation) are applied to the counter only when counting is stopped (\"freezed\"). \"Freeze\" needs to be buffered before being applied to the counter to avoid glitches. Since AD7741 is a synchronous VFC, i.e. the output pulse is initiated by the edge of the clock (CLKIN), \"freeze\" signal is buffered with a flip-flop triggered by the clock of the VFC. {\\it Bottom}: Schematic of the scan block used to transform 24-bit data into series of bits for the scan-out process. \\label{fig:rxscheme}} \\end{center} \\end{figure} \\begin{figure}[H] \\begin{center} \\includegraphics[width=6.4 in]{F15.eps} \\caption{AMiBA correlator control block diagram.\\label{fig:rxscheme}} \\end{center} \\end{figure} \\begin{figure}[H] \\begin{center} \\includegraphics[width=3.2 in]{F16a.eps} \\includegraphics[width=3.2 in]{F16b.eps} \\caption{ The cleaned images of Jupiter. {\\it Left}: The image was formed with uncalibrated visibilities directly transformed from correlator outputs. {\\it Right}: The constituent visibilities of the image have been calibrated by another set of Jupiter data. Both images are plotted with the same residual noise after CLEAN. The green circle indicates the FWHM of the primary beam, and the blue shaded ellipse at the bottom right corner represents the synthesized beam of the 7-element array in the compact configuration. \\label{fig:jupiter} } \\end{center} \\end{figure} {\\bf Acknowledgments.} We thank the administrative staff for their support over the years. We thank A. Harris for useful discussions and notes. We thank the Ministry of Education, the National Science Council, and the Academia Sinica for their support of this project. We thank the NOAA for accomodating the AMiBA project on their site on Mauna Loa. We thank the Hawaiian people for allowing astronomers to work on their mountains in order to study the Universe." }, "1003/1003.4798_arXiv.txt": { "abstract": "We investigate the critical core mass and the envelope growth timescale, assuming grain-free envelopes, to examine how small cores are allowed to form gas giants in the framework of the core accretion model. This is motivated by a theoretical dilemma concerning Jupiter formation: Modelings of Jupiter's interior suggest that it contains a small core of $< 10M_\\oplus$, while many core accretion models of Jupiter formation require a large core of $> 10M_\\oplus$ to finish its formation by the time of disk dissipation. Reduction of opacity in the accreting envelope is known to hasten gas giant formation. Almost all the previous studies assumed grain-dominated opacity in the envelope. Instead, we examine cases of grain-free envelopes in this study. Our numerical simulations show that an isolated core of as small as $1.7 M_\\oplus$ is able to capture disk gas to form a gas giant on a timescale of million years, if the accreting envelope is grain-free; that value decreases to $0.75 M_\\oplus$, if the envelope is metal-free, namely, composed purely of hydrogen and helium. It is also shown that alkali atoms, which are known to be one of the dominant opacity sources near $1500 \\mathrm{K}$ in the atmospheres of hot Jupiters, have little contribution to determine the critical core mass. Our results confirm that sedimentation and coagulation of grains in the accreting envelope is a key to resolve the dilemma about Jupiter formation. ", "introduction": "} Core masses provide clues to unveiling the origins of gas giants. In the core accretion scenario for gas giant formation, a forming solid protoplanet (i.e., a proto-core) experiences a rapid gas capture from a protoplanetary disk to form a massive gas envelope, when its mass exceeds a critical mass \\citep{hm80,bp86}. That critical core mass must be reached within the lifetime of the disk gas of several million years \\citep[e.g.][]{h01}, which places a limit on the core mass of a formed gas giant. Because of the slow increase in the critical core mass with core accretion rate \\citep{ds82,mi00}, faster formation in general results in larger core mass. Indeed, in many core-accretion models that are successful in forming Jupiter within several Myr \\citep[][etc.]{jbp96,iiw03,alibert05}, the resultant core mass is as large as $\\sim 10M_\\oplus$ or more. In contrast, the mass of Jupiter's present core is inferred to be small. \\citet{sg04} made an extensive investigation of the interior structure of Jupiter, finding successful models that are consistent with the observed values of its gravitational moments and equatorial radius by using a variety of equations of state (EOSs) for hydrogen and helium. They demonstrated that the possible core mass of Jupiter is smaller than $\\sim 10~M_\\oplus$. This is also supported by recent calculations with an {\\it ab initio} EOS derived in the first-principle approach \\citep{nn08}. While a more massive core of $>10 M_\\oplus$ is reported by \\citet{bm08} who used their own EOS of hydrogen-helium mixtures based on density functional molecular dynamics, \\citet{fn09} pointed out the difference in the mass fraction of helium used by the two groups is responsible for this discrepancy in the derived value of Jupiter's core mass. Although this pending problem about Jupiter's core may arise from the uncertainty of EOS, the core mass suggested by interior modeling is, on an average, smaller than that derived by formation theories. This fact motivates us to know how small a core can start the rapid gas accretion to form a massive envelope within several Myr. Reduction of opacity in the protoplanet's envelope has the potential to make a small core possible. As the opacity becomes small, the critical core mass decreases and the post-critical-mass gas accretion becomes fast, because low opacity in the envelope makes it difficult to maintain the envelope's hydrostatic structure without gravitational energy released by contraction of the envelope \\citep{hm80,ds82,mi00}. Since the opacity sources are dust grains and gaseous components in the envelope, a minimum critical core mass is achieved in the case of grain-free envelopes. All the previous studies except one calculation done by \\citet{hm80} (see Section~\\ref{sec: gas opacity}) assumed grain-dominated opacity in the outer envelope. Thus, in this paper, we consider grain-free envelopes and make an extensive investigation of the critical core mass and timescale for gas accretion. Results from the recent works by Podolak and his colleagues \\citep[][]{mp03,mp08} are encouraging. They have directly simulated the dynamical behavior of dust grains to determine their size distribution, and then calculated grain opacity in the accreting envelope. Their numerical simulations revealed that grain opacities in the envelope can be much lower than those in the protoplanetary disk. The reason is that small grains initially suspended in the outer envelope quickly grow large in size and then settle down into the deep envelope where temperature is high enough that grains evaporate. In the following section, the details of the gas opacity used in this study are described. Our numerical calculations and results for the critical core mass and the timescale of gas accretion are shown in Sections 3 and 4, respectively. We discuss a possibility of Jupiter formation with a small core in Section~5. ", "conclusions": "} There is a theoretical dilemma concerning Jupiter formation, as described in Introduction. Modelings of Jupiter's interior suggest that Jupiter has a small core of $<$~$10 M_\\oplus$ \\citep[e.g.,][]{sg04}, while many core-accretion models of Jupiter formation require a large core of $>$~$10 M_\\oplus$ to finish its formation by the time of disk dissipation \\citep[][]{jbp96,alibert04,alibert05,fortier07,fortier09}. The disk instability scenario has been revisited as an alternative scenario of their formation \\citep[e.g.][]{boss00,mayer02}. In this study, we have demonstrated that reduced opacities in the protoplanet's envelope have the potential to resolve this dilemma. From Fig.~\\ref{fig2}, one finds that $M_{\\rm core} = 0.75 M_\\oplus$ for $\\tau_g =$ 1~Myr in the metal-free case; $M_{\\rm core} = 1.7 M_\\oplus$ even in the alkali case. Given observed lifetimes of protoplanetary disks of several Myr, the fact above indicates the reduction of opacity allows Jupiter to have a small core that is consistent with interior modelings in principle. The feasibility of such minimum $M_\\mathrm{crit}$ depends on opacities in the protoplanet's envelope, while this does not change our conclusion that the minimum $M_\\mathrm{crit}$ obtained here provides the lowest limit to core masses of gas giants to which the core accretion model can apply. We need more extensive investigation of gas giant formation into which sedimentation and coagulation of grains in the accreting envelope are incorporated, although reduction of opacity was already pronounced \\citep{mp03,mp08}. Our results also shed light upon growth of solid cores. Many core accretion models assumed the presence of a single protoplanet \\citep[e.g.][]{jbp96}. However, a gas giant is in practice thought of as being formed in a system of multiple protoplanets embedded in a protoplanetary disk. Compared to cases of a single protoplanet, the final mass of a core should be small in the case of a multiple-protoplanet system. According to \\citet{ki98,ki00}, the isolation mass is a few $M_\\oplus$ around $5$AU. Even such a small core is enough for a gas giant to capture disk gas within several Myr, as demonstrated in this study. We need to perform comprehensive simulations on gas giant formation in multiple-protoplanet systems, which incorporate adequate models of planetary accretion such as fragmentation of planetesimals and $e$-damping due to gas drag \\citep[e.g.][]{duncan09}. These calculations will be our future work." }, "1003/1003.2031_arXiv.txt": { "abstract": "We report on optical---near-infrared photopolarimetric observations of a blazar 3C~454.3 over 200~d. The object experienced an optical outburst in July 2007. This outburst was followed by a short state fainter than $V=15.2$~mag lasting $\\sim 25$~d. The object, then, entered an active state during which we observed short flares having a timescale of 3--10~d. The object showed two types of features in the color--magnitude relationship. One is a ``bluer-when-brighter'' trend in the outburst state, and the other is a ``redder-when-brighter'' trend in the faint state. These two types of features suggest a contribution of a thermal emission to the observed flux, as suspected in previous studies. Our polarimetric observation detected two episodes of the rotation of the polarization vector. The first one was a counterclockwise rotation in the $QU$ plane during the outburst state. After this rotation event of the polarization vector, the object entered a rapidly fading stage. The second one was seen in a series of flares during the active state. Each flare had a specific position angle of polarization, and it apparently rotated clockwise from the first to the last flares. Thus, the object exhibited rotations of the polarization vector in opposite directions. We estimated a decay timescale of the short flares during the active state, and then calculated an upper limit of the strength of the magnetic field, $B$=0.2~G, assuming a typical beaming factor of blazars, $\\delta=20$. This upper limit of $B$ is smaller than those previously estimated from spectral analysis. ", "introduction": "Blazars are a subclass of active galactic nuclei (AGN), in which a relativistic jet is viewed at a small angle to the line of sight. They comprise two groups of objects, BL~Lac objects and Flat Spectrum Radio-loud Quasars (FSRQs). Emission lines which originates from the nuclear region (broad and narrow line regions) are observable in the optical range in FSRQs, and rarely in BL~Lac objects \\citep{Urry95}. The radio and optical fluxes from blazars are highly polarized because synchrotron radiation from jets is dominant in the radio---optical bands. Since the polarization can be a probe of the magnetic field in the jet, observations of temporal variations of polarization are important for the investigation of the structure of the jet. Rotations of polarization vectors have, in particular, attracted attentions (e.g. \\cite{Qian91}; \\cite{Sillanpaa93}). \\citet{Jones85} proposed that an apparent rotation is a result of random motion of the polarization vector. \\citet{Marscher08} have recently reported a smooth rotation of the polarization vector associated with an optical flare in BL~Lac. They proposed that the rotation cannot be explained by the random motion scenario, and indicates an emitting zone passing through a helical magnetic field. The timescale of the rotation event in BL~Lac is so short ($\\sim 5$~d) that a number of similar events might have been missed before. \\citet{Larionov08} reported that the polarization vector in 3C~279 had rotated smoothly over approximately two months. And another rotation event in PKS~1510$-$089 was reported by \\citet{Marscher10}, lasting 50$\\pm$10~d. Long-lasting and high time-resolved polarimetric observations are required to determine whether observed rotations indicate real structures of the magnetic field in jets. 3C~454.3 is classified as a FSRQ. \\citet{Raiteri07} and \\citet{Raiteri08} reported a UV excess over the synchrotron component in the spectral energy distribution (SED). The UV excess suggests a substantial contribution of the thermal emission from an accretion disk. Until $\\sim$2001 only a moderate variability in a range of $R\\;\\sim\\;15-17$~mag was observed in the optical regime \\citep{Villata06}. The object showed an unprecedentedly bright optical outburst in May 2005. \\citet{Fuhrmann06} reported that the object reached a maximum about $R\\sim12.0$~mag during this bright outburst. The object again experienced a major outburst in July---August 2007, and subsequently another short flares in November 2007---February 2008 \\citep{Raiteri08}. The object reached a maximum of $R=12.58$~mag during the outburst state in 2007. In the gamma-ray region, a flare was also detected by the AGILE satellite together with the 2007 optical outburst \\citep{Vercellone08}. In general, a blazar becomes bluer when it is brighter. This feature is so-called the ``bluer-when-brighter'' trend (e.g. \\cite{Racine70}). On the other hand, a ``redder-when-brighter'' trend has been observed in 3C~454.3 (\\cite{Miller81}; \\cite{Villata06}; \\cite{Raiteri07}). This trend showed a ``saturation'' at $R\\sim14$~mag, and the object possibly turned into a bluer-when-brighter trend in the bright states during the 2007 outburst \\citep{Raiteri08}. The redder-when-brighter trend in 3C~454.3 is probably caused by a substantial contribution of the thermal emission from the disk. Simultaneous optical and near-infrared (NIR) observations could detect the bluer-when-brighter trend in the bright state more clearly than the observations only within the optical regime. We performed photopolarimetric observations of 3C~454.3 simultaneously in the optical and NIR bands from July 2007 to February 2008. In this paper, we report the light curve, color and polarization variations of 3C~454.3 in the optical and NIR regions. Our polarimetric observations found intriguing rotations of the polarization vector. This paper is arranged as follows: In section 2, we present the observation method and analysis. In section 3, we report the result of the photometric and polarimetric observations. In section 4, first, we discuss the origin of a long term component under the short flares. Second, we discuss implications of the observed rotation events of the polarization vector, Finally, we estimate the strength of magnetic field using a decay timescale of flares. The conclusion is drawn in section~5. ", "conclusions": "We observed the outburst, faint and post-outburst active states of the blazar 3C~454.3 in 2007 in multicolor photopolarimetric mode. We found that 3C~454.3 experienced the rotations of the polarization vector in both clockwise and counterclockwise directions in the $QU$ plane during the outburst and post-outburst active states. The feature of the bluer-when-brighter trend was observed during the outburst state. On the other hand, the feature of the redder-when-brighter trend was observed during the faint state. It indicates that the relative contribution of the thermal component increased in the $V$ band during this state. \\\\ \\\\ This work was partly supported by a Grant-in-Aid from the Ministry of Education, Culture, Sports, Science and Technology of Japan (19740104). A part of the light curve presented in this paper is based on data taken and assembled by the WEBT collaboration and stored in the WEBT archive at the Osservatorio Astronomico di Torino-INAF (http://www.oato.inaf.it/blazars/webt/)." }, "1003/1003.0805_arXiv.txt": { "abstract": "Mid-infrared astronomy (operating at wavelengths ranging from 2 to 25 $\\mu$m) has progressed significantly in the last decades, thanks to the improvement of detector techniques and the growing diameter of telescopes. Space observatories benefit from the absence of atmospheric absorption, allowing to reach the very high sensitivities needed to perform 3D hyperspectral observations, but telescopes are limited in diameter ($< 1$ meter) and therefore provide observations at low angular resolution (typically a few seconds of arc). On the other hand, ground-based facilities suffer from strong atmospheric absorption but use large telescopes (above 8m diameter) to perform sub-arcsecond angular resolution imaging through selected windows in the mid-infrared range. In this Paper, we present a method based on Lee and Seung's Non-negative Matrix Factorization (NMF) to merge data from space and ground based mid-infrared (mid-IR) telescopes in order to combine the best sensitivity, spectral coverage and angular resolution. We prove the efficiency of this technique when applied to real mid-IR astronomical data. We suggest that this method can be applied to any combination of low and high spatial resolution positive hyperspectral datasets, as long as the spectral variety of the data allows decomposition into components using NMF. ", "introduction": "\\label{intro} In galaxies (including our own Galaxy, the Milky-Way), the ultraviolet (UV) and visible light (respectively wavelength ranges of 10-400 and 400-750 nm) emitted by stars is absorbed by dust particles having sizes ranging from a few nanometers to a few micrometers. The UV-visible energy they absorb is then re-emitted in the infrared (IR, 2-1000$\\mu$m \\cite{dra03}). Therefore, in the recent years, astronomers have focused their efforts in studying the Milky-Way and external galaxies to the IR, where the emitted light contains information both on the amount of absorbed energy originating from stars and on the composition of interstellar dust. Unfortunately, most of the IR light coming from space is absorbed by the atmosphere. This has motivated a number of IR space missions that provide the sensitivity required to perform hyperspectral observations. However, due to technical and cost constrains, space mission can only launch small diameter telescopes, which limits the spatial resolution of the obtained data. Conversely, the ground-based telescopes working in the IR usually have large apertures providing subarcsecond angular resolution, but low sensitivities, that only permit imaging through a few broad band filters. In the field of remote sensing, a set of methods referred to as \\emph{pansharpening} \\cite{alp07} have been developed in order to perform data fusion/improvement by combining hyperspectral datasets at different spatial resolutions. In this paper, we present such a method, developed to combine ground and space data, to benefit form the advantages of the two techniques. This method is based on the decomposition into \\emph{components} of hyperspectral data obtained by space telescopes, using Non-Negative Matrix Factorization (NMF), followed by non-negative least square fitting of ground-based data using these \\emph{components} (Sect.\\ref{for}). We apply this to real data obtained in the mid-IR range (5-15 $\\mu$m) in order to prove the efficiency of the proposed method (Sect.\\ref{app}). ", "conclusions": "" }, "1003/1003.5757_arXiv.txt": { "abstract": "During its second Antarctic flight, the CREAM (Cosmic Ray Energetics And Mass) balloon experiment collected data for 28 days, measuring the charge and the energy of cosmic rays (CR) with a redundant system of particle identification and an imaging thin ionization calorimeter. Preliminary direct measurements of the absolute intensities of individual CR nuclei are reported in the elemental range from carbon to iron at very high energy.\\\\\\\\ \\vspace{-0.6cm} ", "introduction": "The CREAM experiment was designed to measure the composition and energy spectra of cosmic rays approaching energies up to 10$^{15}$ eV. Since December 2004, four instruments were successfully flown on balloons over Antarctica where they collected several million CR events in the elemental range from hydrogen to iron, and with total particle energies reaching the 100 TeV scale and above. The final goal of the experiment is to provide a deeper understanding of the acceleration mechanism of cosmic rays and to test the validity of the astrophysical models describing their propagation in the Galaxy \\cite{1}.\\\\ In this paper, we present preliminary energy spectra of the even-charged, abundant nuclei from carbon to iron as measured by the instrument during its second flight (CREAM-II). ", "conclusions": "The CREAM-II instrument carried out measurements of high-Z cosmic-ray nuclei with an excellent charge resolution and a reliable determination of their energy. Energy spectra of the more abundant heavy nuclei are measured and found to agree well with other direct measurements." }, "1003/1003.4492_arXiv.txt": { "abstract": "The left-right twin Higgs model predicts a light stable scalar $\\hat{S}$, which is a candidate for WIMP dark matter. We study its scattering on nucleon and find that the cross section is below the CDMS II upper bound but can reach the SuperCDMS sensitivity. Then we study the Higgs phenomenology by paying special attention to the decay $h\\to \\hat{S}\\hat{S}$ which is strongly correlated with the dark matter scattering on nucleon. We find that such an invisible decay can be sizable, which can severely suppress the conventional decay modes like $h\\to VV (V=W,~Z)$ and $h\\to b\\bar{b}$. On the other hand, compared to the SM prediction, the rates of Higgs boson productions at the LHC via gluon-gluon fusion, weak boson fusion or in association with top quark pairs are all reduced significantly, e.g., the gluon-gluon fusion channel can be suppressed by about $30\\%$. ", "introduction": "The twin Higgs mechanism \\cite{twinhiggs,lrth} is proposed as an interesting solution to the hierarchy problem. The SM Higgs emerges as a pseudo-Goldstone boson once a global symmetry is spontaneously broken, which is similar to what happens in the little Higgs models \\cite{littlehiggs}. An additional discrete symmetry is imposed, which ensures the absence of one-loop quadratic divergence of Higgs mass. The resulting Higgs boson mass is naturally around the electroweak scale when the cut-off scale of the theory is around 5-10 TeV. The twin Higgs mechanism can be implemented in left-right models with the additional discrete symmetry being identified as the left-right symmetry \\cite{lrth}. In the left-right twin Higgs (LRTH) model, several physical Higgs bosons remain after the spontaneous symmetry breaking. Another additional discrete symmetry is introduced in the model under which the $SU(2)_L$ doublet $\\hat{h}$ is odd while all the other fields are even. The lightest particle $\\hat{S}$ in its neutral components is stable and thus can be a candidate for weakly interacting massive particle (WIMP) dark matter. The phenomenology of LRTH model has been studied by some authors \\cite{otherwork}. The density of cold dark matter in the universe has been determined precisely by WMAP \\cite{wmap}: \\beq \\Omega_{CDM}h^2=0.105^{+0.021}_{-0.030}. \\eeq The thermal production of WIMPs can naturally explain such a relic density. As a direct detection of WIMPs, the CDMS attempts to observe the recoil energy transferred to a target nucleus in an elastic collision with a WIMP. Very recently the CDMS collaboration has completed their analysis of the final data runs of the CDMS II experiment and reported two candidate events \\cite{cdms2}. Although these events cannot be interpreted as significant evidence for WIMP interacting with nucleons, the CDMS gives the most stringent upper limit on the WIMP-nucleon spin-independent cross section. For example, the cross section is constrained to be smaller than $3.8\\times10^{-44}$ cm$^2$ for a WIMP of 70 GeV at 90\\% confidence level \\cite{cdms2}. The implications of the new results from the CDMS II experiment have been discussed in many models \\cite{papercdms}. In this work we focus on the left-right twin Higgs model. We first examine the scattering of the dark matter candidate $\\hat{S}$ with nucleon and compare the rate with the CDMS II results. Then we study the Higgs phenomenology, paying special attention to the decay $h\\to \\hat{S}\\hat{S}$ which is strongly correlated with the dark matter scattering on nucleon. We will figure out the size of such an invisible decay rate and how severely to suppress the conventional decay modes like $h\\to VV (V=W,~Z)$ and $h\\to b\\bar{b}$. We also study the suppression for the rates of Higgs boson productions at the LHC via gluon-gluon fusion, weak boson fusion or in association with a pair of top quarks. Since the LHC will be able to discover the Higgs boson in the full mass range \\cite{lhchiggs}, our study will help to probe the left-right twin Higgs model. This work is organized as follows. In Sec. II, we briefly review the left-right twin Higgs model. In Sec. III, we examine the scattering of the dark matter candidate $\\hat{S}$ with nucleon and compare the rate with the CDMS II results. Also, the correlation of Higgs decays with the dark matter scattering on nucleon is studied. In Sec. IV, we calculate the main productions of the Higgs boson at the LHC. Finally, we give our conclusion in Sec. V. ", "conclusions": "In LRTH model, the scalar $\\hat{S}$ is a natural candidate for WIMP dark matter, and the Higgs boson mass is typically in the range of 160 - 180 GeV. Since the invisible decay $h\\to \\hat{S}\\hat{S}$ can affect other decay branching ratios, and also has a strong correlation with the scattering on nucleon, we in this work focused on the low mass region of $\\hat{S}$ so that the decay $h\\to \\hat{S}\\hat{S}$ can be open. We obtained the following observations: (i) The cross section of $\\hat{S}$ scattering on nucleon can naturally satisfy the CDMS II upper bound, and can be large enough to be accessible at SuperCDMS; (ii) The Higgs boson can have a sizable invisible decay $h\\to \\hat{S}\\hat{S}$, whose branching ratio can reach $32\\%$ and has a strong correlation with the cross section of $\\hat{S}$ scattering on nucleon. However, the branching ratios of other new decay modes $h\\to \\hat{A}\\hat{A}$ and $h\\to \\hat{h}_1\\hat{h}_1$ are small; (iii) The branching ratios of the conventional decay modes of the Higgs boson, $h\\to VV$ $(V=W,~Z)$ and $h\\to b\\bar{b}$, can be suppressed over $30\\%$ and $47.5\\%$, respectively; (iV) The Higgs production cross sections times the branching ratios of the conventional decays can be all sizably suppressed. So, it is possible to probe the LRTH model via the Higgs productions at the LHC." }, "1003/1003.0939_arXiv.txt": { "abstract": "To quantify how rare the bullet-cluster-like high-velocity merging systems are in the standard $\\Lambda$ Cold Dark Matter (CDM) cosmology, we use a large-volume ($27~h^{-3}~{\\rm Gpc}^3$) cosmological $N$-body MICE simulation to calculate the distribution of infall velocities of subclusters around massive main clusters. The infall-velocity distribution is given at $(1-3)R_{200}$ of the main cluster (where $R_{200}$ is similar to the virial radius), and thus it gives the distribution of realistic initial velocities of subclusters just before collision. These velocities can be compared with the initial velocities used by the non-cosmological hydrodynamical simulations of 1E0657-56 in the literature. The latest parameter search carried out by Mastropietro and Burkert have shown that the initial velocity of 3000~km~s$^{-1}$ at about $2R_{200}$ is required to explain the observed shock velocity, X-ray brightness ratio of the main and subcluster, X-ray morphology of the main cluster, and displacement of the X-ray peaks from the mass peaks. We show that such a high infall velocity at $2R_{200}$ is incompatible with the prediction of a $\\Lambda$CDM model: the probability of finding 3000~km~s$^{-1}$ in $(2-3)R_{200}$ is between $3.3\\times 10^{-11}$ and $3.6\\times 10^{-9}$. A lower velocity, 2000~km~s$^{-1}$ at $2R_{200}$, is also rare, and moreover, Mastropietro and Burkert have shown that such a lower initial velocity does not reproduce the X-ray brightness ratio of the main and subcluster or morphology of the main cluster. Therefore, we conclude that the existence of 1E0657-56 is incompatible with the prediction of a $\\Lambda$CDM model, unless a lower infall velocity solution for 1E0657-56 with $\\lesssim 1800~{\\rm km~s^{-1}}$ at $2R_{200}$ is found. ", "introduction": "The bow shock in the merging cluster 1E0657-57 (also known as the ``bullet cluster'') observed by {\\sl Chandra} indicates that the subcluster \\citep[found by][]{Barrena-etal02} moving through this massive ($10^{15}~ h^{-1}M_\\sun$) main cluster creates a shock, and the shock velocity is as high as $4700~ {\\rm km~s^{-1}}$ \\citep{Markevitch-etal02,Markevitch06}. A significant offset between the distribution of X-ray emission and the mass distribution has been observed \\citep{Clowe-etal04,Clowe-etal06}, also indicating a high-velocity merger with gas stripped by ram pressure. Several groups have carried out detailed investigations of the physical properties of 1E0657-57 using non-cosmological hydrodynamical simulations \\citep{Takizawa05,Takizawa06,Milos-etal07,SF07,MB08}. One of the key input parameters for all of these simulations is the {\\it initial velocity} of the subcluster, which is usually given at somewhere near the virial radius of the main cluster. An interesting question is whether the existence of such a high-velocity merging system is expected in a $\\Lambda$CDM universe. \\citet{HW06} were the first to calculate the likelihood of subcluster velocities using the Millennium Run simulation \\citep{Springel-etal05}. As the volume of the Millennium Run simulation is limited to $(0.5~h^{-1}~{\\rm Gpc})^3$, there are only 5 cluster-size halos with $M_{200}>10^{15}~h^{-1}~M_\\sun$, and 1 cluster with $M_{200}>2\\times 10^{15}~h^{-1}~M_\\sun$ at $z=0.28$ (close to the redshift of 1E0657-57, $z=0.296$). Therefore, \\citet{HW06} had to extrapolate their results for $M_{200}>10^{14}~h^{-1}~M_\\sun$ assuming that the likelihood of finding the bullet-cluster systems scales with $V_{\\rm sub}/V_{200}$, where $V_{\\rm sub}$ is the subcluster velocity in the rest frame of the main cluster, and $V_{200}=(GM_{200}/R_{200})^{1/2}$. Here, $R_{200}$ is the radius within which the mean mass density is 200 times the critical density of the universe, and $M_{200}$ is the mass enclosed within $R_{200}$. While \\citet{HW06} concluded that the existence of 1E0657-57 is consistent with the standard $\\Lambda$CDM cosmology, this conclusion was later challenged by \\citet{FR07} who showed that, once an updated mass of the main cluster of 1E0657-57 is taken into account, the probability of finding 1E0657-57 is as low as $10^{-7}$. This conclusion still relies on the extrapolation of the likelihood derived for $M_{200}>10^{14}~h^{-1}~M_\\sun$. As the probability of finding high-velocity merging systems decreases exponentially with velocities, an accurate determination of the {\\it subcluster velocity}, rather than the shock velocity, is crucial. \\citet{Milos-etal07} and \\citet{SF07} used hydrodynamical simulations to show that the subcluster velocity can be significantly lower than the shock velocity (which is $4700~{\\rm km~s^{-1}}$). \\citet{Milos-etal07} found that the subcluster velocity can be $4050~{\\rm km~s^{-1}}$, whereas \\citet{SF07} found that it can be as low as $2700~{\\rm km~s^{-1}}$. \\citet{MB08} showed that the subcluster velocity of $3100~{\\rm km~s^{-1}}$ best reproduces the X-ray data of 1E0657-57. These varying results are due in part to the varying assumptions about the initial velocity given to the subcluster at the beginning of their hydrodynamical simulations: \\citet{Milos-etal07} used zero relative velocity between the main cluster and subcluster at the initial separation of 4.6~Mpc (which is 2 times $R_{200}$ of the main cluster, 2.3~Mpc). The velocity is about 1600~km~s$^{-1}$ at a separation of 3.5~Mpc ($\\simeq 1.5R_{200}$)\\footnote{Milosavljevi\\'c (2010), private communication. All velocities quoted throughout this paper are calculated in the rest frame of the main cluster.}; \\citet{SF07} used the initial velocity of 2057~km~s$^{-1}$ when the separation was 3.37~Mpc ($\\simeq 1.5R_{200}$); and \\citet{MB08} explored various initial velocities such as 2057~km~s$^{-1}$ at the initial separation of 3.37~Mpc and 2000, 3000, and 5000~km~s$^{-1}$ at the initial separation of 5~Mpc ($\\simeq 2.2R_{200}$). \\citet{MB08} found that the simulation run with the initial velocity of 3000~km~s$^{-1}$ best reproduces the X-ray data. In this paper, we demonstrate that the initial velocities used by \\citet{Milos-etal07} and \\citet{SF07} are consistent with the prediction of a $\\Lambda$CDM model, but those of \\citet{MB08} at 5~Mpc are not. The simulations of \\citet{Milos-etal07} and \\citet{SF07} do not reproduce details of the X-ray and weak lensing data of 1E0657-57, and \\citet{MB08} argue that one needs the initial velocity of 3000~km~s$^{-1}$ to explain the data. If this is true, the existence of 1E0657-57 is incompatible with the prediction of a $\\Lambda$CDM model. ", "conclusions": "\\citet{MB08} showed that the subcluster initial velocity of 3000~km~s$^{-1}$ at the separation of 5~Mpc is required to explain the X-ray and weak lensing data of 1E0657-56 at $z=0.296$. They argued that a lower velocity, 2000~km~s$^{-1}$, seems excluded because it cannot reproduce the observed X-ray brightness ratio of the main and subcluster or the X-ray morphology of the main cluster. In this paper, we have shown that such a high velocity at 5~Mpc, which is about 2 times $R_{200}$ of the main cluster, is incompatible with the prediction of a $\\Lambda$CDM model. Using the results at $z=0$ and $M_{\\rm main}\\ge 0.7\\times 10^{15}~h^{-1}~M_\\sun$, $\\Lambda$CDM is excluded by more than 99.91\\% confidence level (none out of 1135 subclusters has $V_c\\ge 2000$~km~s$^{-1}$ in $2\\le r/R_{200}\\le 3$). For a lower minimum main cluster mass, $M_{\\rm main}\\ge 0.5\\times 10^{15}~h^{-1}~M_\\sun$, $\\Lambda$CDM is excluded by more than 99.95\\% confidence level (none out of 2189 subclusters has $V_c\\ge 2000$~km~s$^{-1}$ in $2\\le r/R_{200}\\le 3$). The results at $z=0.5$ are not yet fully conclusive due to the limited statistics: none out of 78 subclusters has $V_c\\ge 3000$~km~s$^{-1}$ in $2\\le r/R_{200}\\le 3$, while there is one subcluster with $V_c\\ge 2000$~km~s$^{-1}$ in $2\\le r/R_{200}\\le 3$. For $M_{\\rm main}\\ge 0.5\\times 10^{15}~h^{-1}~M_\\sun$, none out of 186 subclusters has $V_c\\ge 3000$~km~s$^{-1}$, while there is one subcluster with $V_c\\ge 2000$~km~s$^{-1}$. While these confidence levels are directly measured from the simulation, one can estimate the probability better by fitting the probability density, $p(\\log V_c)$, to a Gaussian distribution as \\begin{equation} p(\\log V_c) = \\frac{1}{\\sqrt{2\\pi\\sigma^2_{\\nu}}} {\\exp\\left[-\\frac{\\left(\\log V_c-\\nu\\right)^2}{2\\sigma^2_{\\nu}}\\right]}, \\label{eqn:fit} \\end{equation} where $V_c$ is in units of km~s$^{-1}$ and $\\nu$ and $\\sigma_{\\nu}$ are the two fitting parameters. The best-fit values of the two parameters for $z=0$ and $0.5$ are $(\\nu,\\sigma_\\nu)=(3.02,0.07)$ and $(3.13,0.06)$, respectively. The mean velocity at $z=0$ is smaller than that at $z=0.5$ by a factor of $10^{3.13-3.02}=1.29$. This may be understood as the effect of $\\Lambda$ slowing down the structure formation at $z<0.5$. Generally, one has to be careful about this approach, as we are probing the tail of the distribution, where the above fits may not be accurate. Using the above Gaussian fits, we find $P(>3000~{\\rm km~s^{-1}})=3.3\\times 10^{-11}$ and $3.6\\times 10^{-9}$ at $z=0$ and $0.5$, respectively. We also find $P(>2000~{\\rm km~s^{-1}})=2.9\\times 10^{-5}$ and $2.2\\times 10^{-3}$ at $z=0$ and $0.5$, respectively. These numbers pose a serious challenge to $\\Lambda$CDM, unless one finds a lower velocity solution for 1E0657-56. Here, a ``lower velocity'' may be somewhere between $V_c\\lesssim 1500$ and $1800$~km~s$^{-1}$ at $r\\sim 2R_{200}$, which give 1\\% probabilities at $z=0$ and $z=0.5$, respectively. The bullet cluster 1E0657-56 is not the only site of violent cluster mergers. For example, there are A520 \\citep{Markevitch-etal05} and MACS~J0025.4-1222 \\citep{Bradac-etal08}. Also, high-resolution mapping observations of the Sunyaev-Zel'dovich (SZ) effect have revealed a violent merger event in RX~J1347-1145 at $z=0.45$ \\citep{Komatsu-etal01,Kitayama-etal04,Mason-etal09}, which are confirmed by X-ray observations \\citep{Allen-etal02,Ota-etal08}. The shock velocity inferred from the SZ effect and the X-ray data of RX~J1347-1145 is 4600~km~s$^{-1}$ \\citep{Kitayama-etal04}, which is similar to the shock velocity observed in 1E0657-56 \\citep{Markevitch06}. The lack of structure in the redshift distribution of member galaxies of RX~J1347-1145 suggests that the geometry of the merger of this cluster is also closer to edge-on \\citep{Lu-etal10}. However, the lack of a bow shock in the {\\sl Chandra} image may suggest that it is not quite as edge-on as 1E0657-56. In any case, it seems plausible that there may be more clusters like 1E0657-56 in our universe. This too may present a challenge to $\\Lambda$CDM. Since the volume of the MICE simulation is close to the Hubble volume,\\footnote{For example, the comoving volume available from $z=0$ to $z=1$ over the full sky is $54~h^{-3}~{\\rm Gpc}^3$, which is only twice as large as the volume of the MICE simulation. The comoving volume out to $z=3$ is still $396~h^{-3}~{\\rm Gpc}^3$, which is nowhere near enough to overcome the probability of $10^{-9}$.} our results can be compared directly with observations, provided that detailed follow-up observations are available for us to calculate the shock velocity, gas distribution, and dark matter distribution. These three observations would then enable us to estimate the mass ratio and initial velocity of the collision which, in turn, can be compared to the probability distribution we have derived in this paper. Note also that the probabilities obtained in our work are the conditional ones. That is, the probability for which a fitting formula is provided is the probability of the velocity of bullet-systems that are nearly head-on, with 1:10 or more mass ratio, and with $M_{main}\\ge 0.7\\times10^{15}\\, h^{-1}M_{\\odot}$. If we computed the probability of finding high-velocity bullet systems among all clusters from the simulation, then the probability would be even smaller than those estimated above. Such a conditional probability is relevant to the observation, if we have sufficient amount of data for estimating the mass ratio and the initial velocity, as mentioned above. Note that we have precisely such data for 1E0657-56. An interesting question that we have not addressed in this paper is how many high-velocity bullet systems are expected for flux-limited galaxy cluster surveys, such as the South Pole Telescope and eROSITA (extended ROentgen Survey with an Imaging Telescope Array). To calculate, e.g., $dN_{\\rm bullet}/dz$, one needs the light-cone output of the MICE simulation. While we have not investigated this, we expect two major light-cone effects on the infall velocity distribution. First, the infall velocities at high $z$'s should be larger since the effect of $\\Lambda$ has yet to kick in at high $z$'s, which we have already demonstrated here by comparing the mean infall velocity at $z=0$ and $0.5$. Second, the massive bullet systems with mass greater than $10^{15}\\, h^{-1}M_{\\odot}$ are very rare at higher $z$'s. The first effect will make the high-velocity system more common, while the second effect will make the high-velocity system less common. In order to quantify the net effect, one needs the light-cone output. However, the light-cone effect alone would not be able to reconcile the existence of 1E0657-56 with the prediction of $\\Lambda$CDM." }, "1003/1003.5946_arXiv.txt": { "abstract": "We study the physics of electron acceleration at collisionless shocks that move through a plasma containing large-scale magnetic fluctuations. We numerically integrate the trajectories of a large number of electrons, which are treated as test particles moving in the time dependent electric and magnetic fields determined from 2-D hybrid simulations (kinetic ions, fluid electron). The large-scale magnetic fluctuations effect the electrons in a number of ways and lead to efficient and rapid energization at the shock front. Since the electrons mainly follow along magnetic lines of force, the large-scale braiding of field lines in space allows the fast-moving electrons to cross the shock front several times, leading to efficient acceleration. Ripples in the shock front occuring at various scales will also contribute to the acceleration by mirroring the electrons. Our calculation shows that this process favors electron acceleration at perpendicular shocks. The current study is also helpful in understanding the injection problem for electron acceleration by collisionless shocks. It is also shown that the spatial distribution of energetic electrons is similar to in-situ observations \\citep[e.g.,][]{Bale1999GRL,Simnett2005}. The process may be important to our understanding of energetic electrons in planetary bow shocks and interplanetary shocks, and explaining herringbone structures seen in some type II solar radio bursts. ", "introduction": "Collisionless shocks are widely believed to be the primary acceleration mechanism giving rise to the ubiquitous existence of energetic particles in space. The theory of diffusive shock acceleration (DSA) was proposed some $30$ years ago \\citep{Axford1978,Bell1978MNRAS,Blandford1978,Krymsky1977} and is currently believed to be the most important mechanism for a variety of astrophysical environments, for example, interplanetary shocks, the heliospheric termination shock, and shocks associated with supernova remnants. The theory predicts a universal power-law energy flux spectrum $dJ/dE\\propto E^{-1}$ for strong shocks with a density compression ratio of 4. The physical mechanism by which particles are accelerated from thermal energies to much higher energies where DSA is presumed to be applicable (the injection problem) has received some recent attentions but has not reached a common consensus explanation. Many acceleration theories, for example, shock drift acceleration (SDA) \\citep[see reviews,][]{Armstrong1985,Decker1988SSRv} and shock surfing acceleration \\citep{Sagdeev1966RvPP,Lee1996JGR,Zank1996JGR} have been proposed. The acceleration of low energy protons in the shocks containing large-scale pre-existing magnetic fluctuations is very efficient \\citep{Giacalone2005ApJa,Giacalone2005ApJb}, which suggests that there may not be an injection problem. It is generally thought that, pre-accelerated particles will interact resonantly with magnetic turbulence which results in isotropization and diffusion. Many previous works considering magnetic turbulence focused on the ions \\citep[e.g.,][]{Bell1978MNRAS,Giacalone1992GeoRL,Ng2003ApJ,Giacalone2004ApJ,Giacalone2005ApJa,Giacalone2005ApJb}. However, the acceleration of electrons is less well understood since for electrons whose gyroradii are very small, the cyclotron resonance condition is not easily satisfied thus they cannot interact resonantly with large-scale ambient turbulence on ion-scale. While the scattering provided by whistler waves \\citep{Shimada1999} is one possibility, \\citet{Jokipii2007ApJ} proposed an attractive solution to the injection problem that does not require pitch-angle scattering, i.e., conserving the first adiabatic invariant. The idea is that the low-rigidity particles, especially electrons, can move rapidly along meandering magnetic field lines and thus travel back and forth between shock front. The particles gain energy from the difference between upstream and downstream flow velocities. Energetic electrons are often observed to be associated with collisionless shocks. Accelerated electrons are thought to produce type II radio bursts in the solar corona and interplanetary space. \\citet{Anderson1979GeoRL} reported \\emph{ISEE} spacecraft measurements of upstream electrons ($>16$ keV) of the Earth's bow shock that originate from a thin region close to the point of tangency between interplanetary magnetic field lines and the shock surface. \\citet{Tsurutani1985JGR} showed observations of energetic electrons associated with interplanetary shocks showing \"spike-like\" flux enhancements for energies $\\gtrsim 2$ keV. The spike events were observed at quasi-perpendicular shocks with $\\theta_{Bn} \\gtrsim 70^\\circ$, where $\\theta_{Bn}$ is the angle between upstream magnetic field and shock normal. Some shock crossings had no enhancements of energetic electrons which were reported to be associated with low shock speeds and small $\\theta_{Bn}$. \\citet{Simnett2005} presented data which shows energetic electrons are accelerated close to shock front. They also showed some accelerated electrons can escape far upstream of quasi-perpendicular interplanetary shocks. The clear evidence of electron acceleration by DSA is rare, but a recent example was discussed by \\citet{Shimada1999} showing the importance of whistler waves. In order to explain the energization of electrons within the shock layer, \\citet{WuCS1984} and \\citet{Leroy1984} developed analytic models for electron acceleration from thermal energies by adiabatic reflection by a quasi-perpendicular shock. This is known as fast-Fermi acceleration. This theory describes a scatter-free electron acceleration process in a planar, time-steady shock. It obtains a qualitative agreement with observations at Earth's bow shock in terms of the loss-cone pitch-angle distribution and energy range of accelerated electrons. \\citet{Krauss-Varban1989a} used the combination of electron test particle simulation and 1-D hybrid simulation and verified Wu's basic conclusions. The main energy source of fast Fermi acceleration comes from the $-\\textbf{V} \\times \\textbf{B}/c$ electric field which is the same as SDA \\citep{Armstrong1985}. It can also be demonstrated that fast-Fermi acceleration and SDA are the same process in two different frames of reference \\citep{Krauss-Varban1989JGRb}. Thus one would expect electrons to drift in the direction perpendicular to the flow and magnetic field. For a single reflection, the fraction and energies of accelerated particles are limited \\citep[e.g.,][]{Ball2001PASA}. \\citet{Holman1983ApJ} proposed the basic outline for type II solar radio bursts in which energetic electrons are accelerated through SDA. It is expected that multiple reflections are required in order to explain herringbone structures in type II bursts, where the electrons are accelerated to a fraction of the speed of light. More recently, \\citet{Burgess2006} studied electron acceleration in 2-D quasi-perpendicular shocks using test-particle simulations and self-consistent hybrid simulations. He focused on the effect of the rippling of the shock front on particle acceleration in highly oblique shocks with $\\theta_{Bn} \\geq 80^\\circ$. The ripples, in this case, were produced by instabilities along the shock front. Burgess found that the acceleration of electron by SDA can be more efficient for a rippled shock. The shape of the resulting energy spectra has a flat plateau from the initial release energies to energies several times higher than this. Above the flat plateau the spectra drop off steeply as $\\theta_{Bn}$ become smaller. In this paper, we use test particle simulations combined with $2$-D hybrid simulations that include pre-existing large-scale magnetic field turbulence \\citep{Giacalone2005ApJb} to study the shock acceleration of electrons. In addition to the effect of the large-scale turbulence, the shock microphysics occuring on ion length and time scales is also included. Moreover, the shock front is rippled and distorted in response to the turbulence, which is also included in our model. In $\\S $ 2 we describe the numerical method we used to combine the fields from the hybrid simulation and test-particle simulation to obtain the electron distribution. $\\S $ 3 gives the main results of our simulation. In $\\S $ 4 we summarize the main conclusions and discuss the implication of our work. ", "conclusions": "We studied the acceleration of electrons at collisionless shocks by utilizing a combination of a 2-D hybrid simulation to obtain the shock structure and a test-particle simulation to determine the motion of electrons. The hybrid simulation provides realistic electric and magnetic fields within the transition layer of the shock that effect the motion of test-electrons, which is determined by solving the equation of motion. The interaction of the shock with pre-existing upstream fluctuations, and other nonlinear processes occuring in the hybrid simulation lead to a \"rippling\" of shock surface which also effects the transport of the electrons. We find that the electrons are efficiently accelerated by a nearly perpendicular shock. The turbulent magnetic field, leads to field-line meandering which allows the electrons to cross the shock front many times. The rippling of the shock front also contributes to the acceleration by mirroring electrons between the ripples. For the case that the averaged shock normal angle $<\\theta_{Bn}> = 90^\\circ$ and turbulence variance $\\sigma = 0.3$, the electrons can readily be accelerated to up to $200-300$ times the initial energy. The resulting spectrum is flat between about $0.1$ keV to $0.7$ keV. At higher energies, the spectrum falls off with energy like a power law with a spectral slope of about $-3$. This acceleration process is more efficient at perpendicular shocks. As $<\\theta_{Bn}>$ decreases from $90^\\circ$ , both the number fraction and highest achievable energy of accelerated particles decreases. Based on our calculations, we conclude that perpendicular shocks are the most important for the acceleration of electrons. The current study is helpful in understanding the injection problem for electron acceleration by collisionless shocks. It is also found that different value of magnetic turbulence variances strongly affects the maximum energy attainable. The case with larger turbulence variance has a flatter energy spectrum than the case of smaller turbulence variance, which suggests the enhanced motion of electrons normal to the shock front, due to enhanced field-line random walk, is of importance for the acceleration of electrons to high energies. In addition, we also found that the energetic electron density upstream and downstream of collisionless shocks show filamentary structures (Figure 2). This could help explain electron spike-like events observed upstream and downstream of terrestrial and interplanetary shocks \\citep{Anderson1979GeoRL,Tsurutani1985JGR,Simnett2005}. Observation by Voyager $1$ at the termination shock and in the heliosheath also show the evidence of electron spike-like enhancements at the shock front \\citep{Decker2005Sci}. The upstream spatial distribution of energetic electrons shows irregular features which depend on both the irregularity in the shock surface and the global topology of magnetic field lines. At first the electrons are accelerated and reflected at the shock front, and then they travel upstream along the magnetic field lines. The electrons could be taken far upstream by field line random walk. This result can possibly lead to an interpretation to the complex electron foreshock events recently observed to be associated with interplanetary shocks \\citep{Bale1999GRL,Pulupa2008ApJ}. \\citet{Bale1999GRL} and \\citet{Pulupa2008ApJ} proposed the complex upstream electron events are resulted from large-scale irregularities in shock surface. In this paper we have demonstrated that the upstream electron flux may be controlled by both an irregular shock surface and large-scale meandering magnetic field lines." }, "1003/1003.2647_arXiv.txt": { "abstract": "We present a joint analysis of the overlapping BLAST 250, 350, 500\\,\\micron, and Large APEX Bolometer Camera 870\\,\\micron\\ observations (from the LESS survey) of the Extended Chandra Deep Field South. Out to $z \\sim 3$, the BLAST filters sample near the peak wavelength of thermal far-infrared (FIR) emission from galaxies (rest-frame wavelengths $\\sim 60$--200\\,\\micron), primarily produced by dust heated through absorption in star-forming clouds. However, identifying counterparts to individual BLAST peaks is very challenging, given the large beams (FWHM 36--60\\,arcsec). In contrast, the ground-based 870\\,\\micron\\ observations have a significantly smaller 19\\,arcsec FWHM beam, and are sensitive to higher redshifts ($z\\sim1$--5, and potentially beyond) due to the more favourable negative $K$-correction. We use the LESS data, as well as deep {\\em Spitzer} and VLA imaging, to identify 118 individual sources that produce significant emission in the BLAST bands. We characterize the temperatures and FIR luminosities for a subset of 69 sources which have well-measured submm SEDs and redshift measurements out to $z\\sim3$. For flux-limited sub-samples in each BLAST band, and a dust emissivity index $\\beta=2.0$, we find a median temperature $T=30$\\,K (all bands) as well as median redshifts: $z=1.1$ (interquartile range 0.2--1.9) for $S_{250} > 40$\\,mJy; $z=1.3$ (interquartile range 0.6--2.1) for $S_{350} > 30$\\,mJy; and $z=1.6$ (interquartile range 1.3--2.3) for $S_{500} > 20$\\,mJy. Taking into account the selection effects for our survey (a bias toward detecting lower-temperature galaxies), we find no evidence for evolution in the local FIR-temperature correlation out to $z\\sim2.5$. Comparing with star-forming galaxy SED templates, about 8\\% of our sample appears to exhibit significant excesses in the radio and/or mid-IR, consistent with those sources harbouring an AGN. Since our statistical approach differs from most previous studies of submm galaxies, we describe the following techniques in two appendices: our `matched filter' for identifying sources in the presence of point-source confusion; and our approach for identifying counterparts using likelihood ratios. This study is a direct precursor to future joint far-infrared/submm surveys, for which we outline a potential identification and SED measurement strategy. ", "introduction": "\\label{sec:intro} Observations in the submillimetre (submm) wavelength band (defined here to be 200--1000\\,\\micron) are ideal for detecting light from massive star-forming galaxies out to cosmological distances. It has been known since the all-sky {\\em Infrared Astronomical Satellite (IRAS)} survey of the 1980's that such sources contain significant amounts of dust, so that the ultra-violet (UV) light of newly-formed stars is absorbed by the galaxies' interstellar medium (ISM) \\citep{sanders1996}. The dust is typically heated to tens of Kelvin, and most of the light is then thermally re-radiated at far-infrared (FIR) wavelengths ($\\sim$60--200\\,\\micron). In the submm, the thermal spectral energy distribution (SED) drops off steeply, so that there is a progressively stronger negative $K$-correction with increasing observing wavelength. The correction is so strong that near $\\sim$1\\,mm the observed flux density for a galaxy of fixed luminosity is approximately constant from $1 \\lsim z \\lsim 10$ \\citep{blain2002}. Even though much of the submm band is obscured to ground-based observations by atmospheric water vapour, a number of surveys over the last decade have exploited transparency in several spectral windows to successfully locate high-redshift ($z>1$) dusty star-forming galaxies solely through their submm emission (submillimetre galaxies, or SMGs). Their discovery was first made with the Submillimetre Common User Bolometer Array \\citep[SCUBA][]{holland1999} at 850\\,\\micron\\ \\citep[e.g.,][]{smail1997,hughes1998,barger1998,cowie2002, scott2002,borys2003,webb2003,coppin2006}. Several other instruments confirmed their existence in the slightly more transparent 1.1--1.2\\,mm band \\citep[e.g.,][]{greve2004,laurent2005,scott2008, perera2008,austermann2010}. These ground-based surveys at 850--1200\\,\\micron\\ typically cover $\\ll 1$\\,deg$^2$, and detect several tens of sources per field. The typical angular resolution of these surveys is in the range $\\sim$9--20 arcsec full-width at half-maximum (FWHM). It is worth noting that observations at 350 and 450\\,\\micron\\ have also been attempted from the ground \\citep[e.g.,][]{smail1997,hughes1998,fox2002,kovacs2006,khan2007,coppin2008}. However, this wavelength range is much more difficult, due to increased atmospheric opacity, so that these surveys have only detected a handful of sources. While the first generation surveys successfully demonstrated the existence of these ultra-luminous infrared galaxies (ULIRGs) at $z\\sim1.5$--4 \\citep[e.g][]{chapman2003b,aretxaga2003,chapman2005}, sample sizes have been modest (typically $\\ll 100$ sources in a given field). Most of what is known about SMGs is based on cross-identifications with sources in higher-resolution data, particularly in the radio \\citep[primarily 1.4\\,GHz Very Large Array maps, e.g.,][]{smail2000,ivison2007} and in the mid-IR \\citep[such as 24\\,\\micron\\ {\\em Spitzer} maps, e.g.,][]{ivison2004,pope2006}. While these counterpart identification strategies could be biased toward lower redshifts due to the {\\em positive} $K$-corrections in the radio/mid-IR, more observationally time-consuming mm-wavelength interferometric observations \\citep[e.g.,][]{lutz2001,dannerbauer2004,iono2006,younger2007} demonstrate reasonable correspondence with proposed radio/mid-IR counterparts for a handful of sources. With accurate positions it is then possible to identify optical counterparts, although they are usually extremely faint due to obscuration by the same dust that makes them bright in the submm--FIR, and the fact that stellar light from the most distant objects gets red-shifted out of the optical bands into the near-IR. Obviously, ground-based optical spectroscopy is even more challenging given the difficulty in imaging the counterparts. Recent observations by the 1.8-m Balloon-borne Large Aperture Submillimeter Telescope (BLAST) at 250, 350, 500\\,\\micron\\ \\citep[a pathfinder for {\\em Herschel}/SPIRE,][]{pascale2008} toward the Extended Chandra Deep Field South (ECDF-S) have provided the first confusion-limited submm maps at these wavelengths which cover areas larger than 1\\,deg$^2$. These bands were chosen to bracket the peak rest-frame FIR emission from the SMG population at $z\\sim1$--4. However, given the size of its primary mirror, the BLAST diffraction-limited angular resolution of 36--60\\,arcsec FWHM at 250--500\\,\\micron\\ has made associations between submm emission peaks and individual sources at other wavelengths considerably more challenging than with the existing ground-based surveys at longer wavelengths. For this reason many of the primary BLAST scientific results to date have been derived from the statistics of brightness fluctuations for entire maps, such as the number counts \\citep{devlin2009,patanchon2009}, contributions of known sources to the Cosmic Infrared Background \\citep[CIB][]{puget1996,fixsen1998} in the BLAST bands \\citep{devlin2009,marsden2009,pascale2009}, evolution in the FIR--radio correlation \\citep{ivison2009}, and the large-scale clustering of infrared-bright galaxies \\citep{viero2009}. We emphasize that {\\em none of these results depend on identifying individual submm sources in the BLAST maps}. More traditional analyses of BLAST sources identified through peaks in maps convolved with the point spread function (PSF) have also been attempted \\citep{dye2009,dunlop2010,ivison2009}. In general, it has been a struggle to determine whether these peaks are produced primarily by single galaxies, or blends of several faint sources, necessitating either: conservative cuts in the signal-to-noise ratio (\\snr) to consider only the very brightest sources; or a careful (though subjective) comparison of all the multi-wavelength data on a case-by-case basis to decide whether single or multiple objects are the likely source of the submm emission. In these earlier papers, Poisson chance alignment `$P$' probabilities \\citep{downes1986} have been used to rank potential counterparts to the BLAST peaks from external matching catalogues, showing that 10\\% spurious threshold probabilities must be adopted to obtain reasonable source statistics (unlike the more conservative 5\\% that is typical in the submm community). These methods yield limited results for these wide-area BLAST maps, despite the fact that the \\snr\\ of the individual peaks rival those of most previous ground-based observations. In this paper, driven by the apparent inadequacy of existing methods for studying individual sources in the low-resolution BLAST maps, we develop improved approaches for: (i) filtering confused maps to find emission peaks that are more likely to be produced by individual (or at least a small number of) sources; and (ii) identifying counterparts to these peaks in external matching catalogues using Likelihood Ratios (LR), a method which can incorporate more prior information than that assumed in the calculation of $P$. In addition to deep BLAST observations of the ECDF-S, we also make extensive use of the deepest wide-area submm map at $\\sim$1\\,mm to date: the Large APEX Bolometer Camera (LABOCA) ECDFS Submm Survey (LESS) at 870\\,\\micron\\ \\citep{weiss2009}, taken with the 12\\,m APEX telescope \\citep{gusten2006}. This first detailed comparison between BLAST and longer-wavelength ground-based submm data helps in two key ways. First, the LABOCA beam has a 19\\,arcsec FWHM (roughly half that of the 250\\,\\micron\\ BLAST beam), enabling us to ascertain {\\em directly} whether some of the BLAST peaks resolve into multiple submm counterparts. Second, like SCUBA, LABOCA \\citep{siringo2009} is more sensitive to $z>1$ sources than BLAST, and the most distant sources \\citep[$z>4$, e.g.,][]{coppin2009} are expected to be LABOCA-detected BLAST-dropouts. Therefore this study will offer superior constraints on the high-redshift submm galaxy population than earlier BLAST studies. Now that we have entered the era of {\\em Herschel} surveys, we also show that these techniques will be useful, despite the approximately twofold improvement in angular resolution offered by SPIRE compared to BLAST. We explore this issue using simulations of SPIRE maps using the smaller PSFs. While we find that the situation is certainly improved for SPIRE, confusion will continue to seriously hamper the interpretation of these new surveys. The analysis is organized as follows. In Section~\\ref{sec:submm} we summarize our treatment of the submm data using our new `matched filter' to identify individual peaks (full details are given in Appendix~\\ref{sec:matched}). We produce an external matching catalogue in Section~\\ref{sec:matchcat} combining 24\\,\\micron\\ mid-IR and 1.4\\,GHz radio priors to select sources from a deep {\\em Spitzer} IRAC near-IR catalogue. The LR identification technique is summarized in Section~\\ref{sec:lr_priors} (a full development, and calculation of priors are given in Appendix~\\ref{sec:lr}), and it is used to produce a list of potential matches to the submm peaks in Section~\\ref{sec:potential}. Also, for cases where matches in the catalogue could not be identified, we search for counterparts in the higher-resolution LESS 870\\,\\micron\\ peak catalogue. At this stage we have a collection of submm peaks, and a list of individual galaxies that we believe produce these peaks -- in many cases blends of several galaxies. To establish their submm flux densities we fit PSFs at all of their locations simultaneously, in each of the submm maps, in Section~\\ref{sec:photometry}. The effects of confusion, missing identifications, and clustering are explored in Sections~\\ref{sec:confusion}--\\ref{sec:clustering} using simulations. We derive redshifts for the proposed counterparts in Section~\\ref{sec:redshift}, and study the rest-frame properties of the sample in Section~\\ref{sec:restsed}, showing in particular how confusion may have biased some of the earlier BLAST results. Finally, in Section~\\ref{sec:future} we simulate SPIRE data to demonstrate the usefulness of our techniques for these new surveys. ", "conclusions": "\\label{sec:conclusions} We have performed a deep multi-wavelength study of individual peaks from the BLAST 250, 350, and 500\\,\\micron\\ survey within the Extended Chandra Deep Field South (ECDF-S). By comparing the BLAST data with LABOCA ECDFS Submm Survey (LESS) maps at 870\\,\\micron\\ we have been able to greatly improve the positional uncertainties and longer-wavelength (Rayleigh-Jeans tail) SEDs for our sample. Compared to the earlier BLAST studies of \\citet{dye2009}, \\citet{ivison2009} and \\citet{dunlop2010}, our methodology for identifying counterparts and measuring flux densities differ in several key respects: \\begin{itemize} \\item It is recognized that peaks in the submm maps (particularly the BLAST bands) are generally blends of several unrelated sources. It is therefore important to search for {\\em multiple} counterparts to each peak in a matching catalogue, whereas earlier BLAST catalogue-based studies tended to focus on {\\em single} counterparts. \\item We use a new `matched filter' that compensates for source confusion when searching for submm sources. For the BLAST bands, in which the contribution to the total noise by source confusion is roughly a factor of two larger than the instrumental noise, there is a significant improvement in \\snr\\ of approximately 15--20\\%. In contrast, there is only a minimal improvement of about 5\\% for the LESS data since the confusion and instrumental noise components are approximately the same. \\item Identifications in our combined radio/IR matching catalogue have been made using Likelihood Ratios that incorporate more prior information than traditional `$P$' statistics. With a threshold spurious ID rate set at 10\\%, we find 52, 50, and 31 matches to 64, 67, and 55 submm peak positions at 250, 350, and 500\\,\\micron, respectively. This is a significant improvement compared to the match rates obtained using 10\\% cuts on $P$ giving 45, 31, and 12 matches in the 24\\,\\micron\\ catalogue, and 51, 35, and 17 matches in 1.4\\,GHz catalogue. Most of the gain is at 500\\,\\micron\\ where the beam is so large that any additional information to help the process clearly makes a large difference. Combining the identifications made independently in each band, approximately 75\\% of the BLAST selected peaks have at least one potential counterpart. \\item We obtain submm photometry for the 118 unique identifications in the matching catalogue by performing a simultaneous fit of the submm PSFs at the precise ID locations in each of the four original (un-smoothed) submm maps. This procedure results in useful measurements for 76 sources; the remaining 42 sources are too close to one another to reliably disentangle the flux densities produced by each of the contributing objects. In order to make further progress in such cases, prior information for the SEDs will likely be needed \\citep[see, for example, the work of][]{roseboom2009}. \\item Our procedure for identifying counterparts and measuring their flux densities has compensated for much of the `flux-boosting' present in the original flux-limited submm peak lists. Comparing the number counts of our final catalogue in the BLAST bands with the $P(D)$ counts from \\citet{patanchon2009} we find good agreement above 40\\,mJy at 250\\,\\micron, 30\\,mJy at 350\\,\\micron\\, and with somewhat lower completeness reaching a maximum of $\\sim66$\\% above 20\\,mJy at 500\\,\\micron. \\end{itemize} We then proceeded to identify redshifts for the counterparts, and hence probe the rest-frame properties of our sample: \\begin{itemize} \\item Of the 76 sources with usable submm photometry, 69 counterparts have redshifts: 23 are optical spectroscopic redshifts; 35 are optical photometric redshifts; and 11 are IRAC photometric redshifts. The median of the entire distribution is $z=1.1$ with an interquartile range 0.3--1.9. Restricting ourselves to sub-samples above the flux-density limits mentioned previously shows a clear trend in increasing redshift with observed wavelength: a median $z=1.1$ with an interquartile range 0.2--1.9 at 250\\,\\micron; a median $z=1.3$ with an interquartile range 0.6--2.1 at 350\\,\\micron; and a median $z=1.6$ with an interquartile range 1.3--2.3 at 500\\,\\micron. In general, there is a higher-redshift tail for our sample than in the earlier study of \\citet{dye2009} which we believe is due to the fainter flux densities probed. \\item Fitting modified blackbody SEDs of the form $S_\\nu \\propto \\nu^{2.0}B_\\nu(T)$ to the 69 sources with useable submm photometry and redshifts, we establish rest-frame cold-dust temperatures. The total distribution has a median $T=29$\\,K and interquartile range 23--36\\,K. There is almost no variation in the temperatures of the different flux-limited sub-samples: a median $T=30$\\,K with an interquartile range 25--39\\,K at 250\\,\\micron; a median $T=30$\\,K with an interquartile range 23--35\\,K at 350\\,\\micron; and a median $T=30$\\,K with an interquartile range 22--38\\,K at 500\\,\\micron. These temperatures are systematically warmer by about 7\\,K than those inferred in \\citet{dye2009}. While this discrepancy may be partly caused by the relative depths of the samples, we believe it is mostly due to improved compensation for confusion (which reduces the estimated flux densities at longer BLAST wavelengths more than at shorter wavelengths), and the inclusion of higher-resolution 870\\,\\micron\\ data from LESS. \\item The primary improvement over earlier catalogue-based BLAST studies is the characterization of submm SEDs for the fainter, but higher-redshift ($z\\gsim1$) and more-luminous galaxies. Even though the earlier analyses typically considered single counterparts to the BLAST peaks, these counterparts {\\em are} probably significant submm emitters (although the measured flux densities were typically high, especially in the most confused 500\\,\\micron\\ band). Therefore the measured properties of the optical counterparts (approximate redshift distributions, galaxy types etc.), and even the FIR temperatures and luminosities of the brighter, lower-redshift galaxies are generally correct. \\item We have also fit star-forming galaxy SED templates from the library of \\citet{dale2001} (which span the radio--near-IR) to the submm and FIR photometry of our sample. Generally speaking there is good correspondence between these models and the photometry in the radio and near/mid-IR, confirming that most of the emission is probably powered by star-formation. However, about 8\\% of the sample exhibits a significant excess in either (or both) of these wavelength regimes, which could indicate AGN activity. We note that these features are restricted to sources with luminosities primarily in the range $11.5<\\log(L_\\mathrm{FIR})<12.8$, near the bright end of our sample. \\item We compare the distribution of luminosity and cold dust temperature in our sample with the local-Universe measurement of \\citet{chapin2009}. While the submm colours of our sources appear systematically cooler at all luminosities compared to the local distribution, we have determined that the distribution is consistent with selection effects in our survey. We therefore find no evidence for evolution in the temperature-luminosity correlation out to $z\\sim2.5$. \\item Finally, we investigate the utility of the methods described in this paper for the new generation of {\\em Herschel}/SPIRE surveys. Using a simple simulation we show that these new surveys are significantly confused, despite a factor of $\\sim$2 improvement in angular resolution over BLAST. We find that the matched-filter will yield a {\\em greater} improvement in the detection of point sources than for BLAST due to the significantly lower instrumental noise of SPIRE. However, even with these improvements, peaks in SPIRE maps will be confused, and we advocate the use of Likelihood Ratios to identify the expected multiple counterparts. \\end{itemize}" }, "1003/1003.0474_arXiv.txt": { "abstract": "Radiation pressure acting on gas and dust causes \\ion{H}{2} regions to have central densities that are lower than the density near the ionized boundary. \\ion{H}{2} regions in static equilibrium comprise a family of similarity solutions with 3 parameters: $\\beta$, $\\gamma$, and the product $Q_0\\nrms$; $\\beta$ characterizes the stellar spectrum, $\\gamma$ characterizes the dust/gas ratio, $Q_0$ is the stellar ionizing output (photons/s), and $\\nrms$ is the rms density within the ionized region. Adopting standard values for $\\beta$ and $\\gamma$, varying $Q_0\\nrms$ generates a one-parameter family of density profiles, ranging from nearly uniform-density (small $Q_0\\nrms$) to shell-like (large $Q_0\\nrms$). When $Q_0\\nrms\\gtsim 10^{52}\\cm^{-3}\\s^{-1}$, dusty \\ion{H}{2} regions have conspicuous central cavities, even if no stellar wind is present. For given $\\beta$, $\\gamma$ and $Q_0\\nrms$, a fourth quantity, which can be $Q_0$, determines the overall size and density of the \\ion{H}{2} region. Examples of density and emissivity profiles are given. We show how quantities of interest -- such as the peak-to-central emission measure ratio, the rms-to-mean density ratio, the edge-to-rms density ratio, and the fraction of the ionizing photons absorbed by the gas -- depend on $\\beta$, $\\gamma$, and $Q_0\\nrms$. For dusty \\ion{H}{2} regions, compression of the gas and dust into an ionized shell results in a substantial increase in the fraction of the stellar photons that actually ionize H (relative to a uniform density \\ion{H}{2} region with the same dust/gas ratio and density $n=\\nrms$). We discuss the extent to which radial drift of dust grains in \\ion{H}{2} regions can alter the dust-to-gas ratio. The applicability of these solutions to real \\ion{H}{2} regions is discussed. ", "introduction": "} \\citet{Stromgren_1939} idealized photoionized nebulae around hot stars as static, spherical regions with a uniform density of ionized gas out to a bounding radius $R$. The Str\\\"omgren sphere model continues to serve as the starting point for studies of \\ion{H}{2} regions around hot stars. However, a number of physical effects lead to departures from the simple Str\\\"omgren sphere model: dynamical expansion of the \\ion{H}{2} region if the pressure in the surrounding neutral medium cannot confine the ionized gas; deviations from sphericity due to nonuniform density; motion of the star relative to the gas; injection of energy and momentum by a stellar wind; absorption of H-ionizing photons by dust grains; and radiation pressure acting on gas and dust. Each of these effects has been the object of a number of investigations, beginning with the study of ionization fronts by \\citet{Kahn_1954}. \\citet{Savedoff+Greene_1955} appear to have been the first to discuss the expansion of a spherical \\ion{H}{2} region in an initially uniform neutral medium. \\citet{Mathews_1967,Mathews_1969} and \\citet{Gail+Sedlmayr_1979} calculated the dynamical expansion of an \\ion{H}{2} region produced by an O star in a medium that was initially neutral, including the effects of radiation pressure acting on the dust. \\citet{Mathews_1967,Mathews_1969} showed that radiation pressure on dust would produce low-density central cavities in \\ion{H}{2} regions. More recently, \\citet{Krumholz+Matzner_2009} reexamined the role of radiation pressure on the expansion dynamics of \\ion{H}{2} regions, concluding that radiation pressure is generally unimportant for \\ion{H}{2} regions ionized by a small number of stars, but is important for \\newtext{the expansion dynamics of} giant \\ion{H}{2} regions surrounding clusters containing many O-type stars. Their study concentrated on the forces acting on the dense shell of neutral gas and dust bounding the \\ion{H}{2} region, hence they did not consider the density structure within the ionized region. Dust absorbs $h\\nu > 13.6\\eV$ photons that would otherwise be able to ionize hydrogen, thereby reducing the extent of the ionized zone. \\citet{Petrosian+Silk+Field_1972} developed analytic approximations for dusty \\ion{H}{2} regions. They assumed the gas density to be uniform, with a constant dust-to-gas ratio, and found that dust could absorb a substantial fraction of the ionizing photons in dense \\ion{H}{2} regions. Petrosian et al.\\ did not consider the effects of radiation pressure. \\citet{Dopita+Groves+Sutherland+Kewley_2003, Dopita+Fischera+Crowley+etal_2006} constructed models of compact \\ion{H}{2} regions, including the effects of radiation pressure on dust, and presented models for different ionizing stars and bounding pressures. In these models, radiation pressure produces a density gradient within the ionized gas. The present paper provides a systematic discussion of the structure of dusty \\ion{H}{2} regions that are assumed to be in equilibrium with an external bounding pressure. The assumptions and governing equations are presented in \\S\\ref{sec:equilibrium}, where it is shown that dusty \\ion{H}{2} regions are essentially described by a 3-parameter family of similarity solutions. In \\S\\ref{sec:results} we show density profiles for selected cases, as well as surface brightness profiles. The characteristic ionization parameter $U_{1/2}$ and the fraction $(1-\\fion)$ of the ionizing photons absorbed by dust are calculated. Dust grain drift is examined in \\S\\ref{sec:dust drift}, where it is shown that it can alter the dust-to-gas ratio in the centers of high density \\ion{H}{2} regions. The results are discussed in \\S\\ref{sec:discussion}, and summarized in \\S\\ref{sec:summary}. ", "conclusions": "Discussion} \\subsection{Absorption of Ionizing Photons by Dust} \\begin{figure}[htb] \\begin{center} \\includegraphics[width=8cm,angle=0]% {f10.eps} \\caption{\\label{fig:inoue} \\capsize Photoionizing fraction $\\fion$ for 12 Galactic \\ion{H}{2} regions, as estimated by \\citet{Inoue_2002} from infrared and radio observations, vs.\\ $Q_{0,49}n_e T_4^{0.83}$ (see text). $\\fion$ cannot exceed 1, therefore the high value found for G298.22-0.34 give some indication of the uncertainties in estimation of $\\fion$. Solid lines: $\\fion$ for \\ion{H}{2} regions with radiation pressure for dust characterized by $\\gamma=5$, 10, and 20. Broken line: $\\fion$ for uniform \\ion{H}{2} regions with $\\sigma_d=10^{-21}\\cm^2\\,{\\rm H}^{-1}$.} \\end{center} \\end{figure} For a sample of 13 Galactic \\ion{H}{2} regions, \\citet{Inoue_2002} used infrared and radio continuum observations to obtain the values of $\\fion$ shown in Figure \\ref{fig:inoue}. The estimated values of $\\fion$ are much larger than would be expected for uniform \\ion{H}{2} regions with dust-to-gas ratios comparable to the values found in neutral clouds. \\citet{Inoue_2002} concluded that the central regions of these \\ion{H}{2} regions must be dust-free, noting that this was likely to be due to the combined effects of stellar winds and radiation pressure on dust. As seen in Fig.\\ \\ref{fig:inoue}, the values of $\\fion$ found by Inoue are entirely consistent with what is expected for static \\ion{H}{2} regions with radiation pressure for $5 \\ltsim \\gamma \\ltsim 20$ (corresponding to $0.5 \\ltsim \\sigma_{d,-21}\\ltsim 2$), with no need to appeal to stellar winds or grain destruction. \\subsection{The Density-Size Correlation for \\ion{H}{2} Regions} \\begin{figure}[htb] \\begin{center} \\includegraphics[width=8cm,angle=0]% {f11a.eps} \\includegraphics[width=8cm,angle=0]% {f11b.eps} \\caption{\\label{fig:ne vs D} \\capsize Density $\\nrms$ vs.\\ diameter $D$. (a) Models with $(\\beta,\\gamma)=$ (2,1), (2,5), (3,10), and (5,20). Results shown were calculated for $T_4=0.94$, $\\langle h\\nu\\rangle_i=18\\eV$. Solid lines show models with $p_{\\rm edge}$ fixed, and $Q_0$ varying from $10^{48}\\s^{-1}$ to $10^{54}\\s^{-1}$. Broken lines show models with $Q_0$ fixed, and $p_{\\rm edge}/k$ varying from $10^4\\cm^{-3}\\K$ to $10^{11}\\cm^{-3}\\K$. (b) Model grid for $\\beta=3$, $\\gamma=10$ together with observed values are shown for various samples of Galactic and extragalactic H\\,II regions. Cyan open triangles: \\citet{Kennicutt_1984}. Blue diamonds: \\citet{Churchwell+Goss_1999}. Green crosses: \\citet{Garay+Lizano_1999}. Cyan crosses: \\citet{Kim+Koo_2001}. Black open stars: \\citet{Martin-Hernandez+Vermeij+vanderHulst_2005}. Red solid triangles: radio sample from \\citet{Hunt+Hirashita_2009}. Red open circles: HST sample from \\citet{Hunt+Hirashita_2009}. } \\end{center} \\end{figure} \\ion{H}{2} regions come in many sizes, ranging from \\ion{H}{2} regions powered by a single O star, to giant \\ion{H}{2} regions ionized by a cluster of massive stars. The physical size of the \\ion{H}{2} region is obviously determined both by the total ionizing output $Q_0$ provided by the ionizing stars, and the r.m.s.\\ density $\\nrms$ of the ionized gas, which is regulated by the pressure $p_{\\rm edge}$ of the confining medium. With the balance between photoionization and recombination determining the size of an \\ion{H}{2} region, an anticorrelation between size $D$ and density $\\nrms$ is expected, and was observed as soon as large samples of \\ion{H}{2} regions became available \\citep[e.g.,][]{Habing+Israel_1979,Kennicutt_1984}. For dustless \\ion{H}{2} regions, one expects $\\nrms\\propto D^{-1.5}$ for fixed $Q_0$, but for various samples relations close to $\\nrms\\propto D^{-1}$ were reported \\citep[e.g.,][]{Garay+Rodriguez+Moran+Churchwell_1993, Garay+Lizano_1999, Kim+Koo_2001, Martin-Hernandez+Vermeij+vanderHulst_2005}. For ultracompact \\ion{H}{2} regions, \\citet{Kim+Koo_2001} attribute the $n_e\\propto D^{-1}$ trend to a ``champagne flow'' and the hierarchical structure of the dense gas in the star-forming region, but \\citet{Arthur+Kurtz+Franco+Albarran_2004} and \\citet{Dopita+Fischera+Crowley+etal_2006} argue that the $n_e\\propto D^{-1}$ trend is a result of both absorption by dust and radiation pressure acting on dust in static \\ion{H}{2} regions. \\citet{Hunt+Hirashita_2009} recently reexamined the size-density relationship. They interpreted the size-density relation for different observational samples in terms of models with different star formation rates [and hence different time evolution of the ionizing output $Q_0(t)$], and differences in the density of the neutral cloud into which the \\ion{H}{2} region expands. Their models did not include the effects of radiation pressure on dust; at any time the ionized gas in an \\ion{H}{2} region was taken to have uniform density, resulting in overestimation of the dust absorption. Figure \\ref{fig:ne vs D}a shows a grid of $\\nrms$ vs.\\ $D$ for the present models, for four combinations of $(\\beta,\\gamma)$. While differences between the models with different $(\\beta,\\gamma)$ can be seen, especially for high $Q_0$ and high $p_{\\rm edge}$, the overall trends are only weakly dependent on $\\beta$ and $\\gamma$, at least for $1\\ltsim\\gamma\\ltsim 20$. Figure \\ref{fig:ne vs D}b shows the model grid for $\\beta=3$ and $\\gamma=5$ together with observed values of $D$ and $\\nrms$ from a number of different studies. It appears that observed \\ion{H}{2} regions -- ranging from \\ion{H}{2} regions ionized by one or at most a few O stars ($Q_0<10^{50}\\s^{-1}$) to ``super star clusters'' powered by up to $10^3-10^5$ O stars ($Q_0=10^{52}-10^{54}\\s^{-1}$) can be accomodated by the present static equilibrium models with external pressures in the range $10^4 \\ltsim p/k \\ltsim 10^{10.3}\\cm^{-3}\\K$. Note that for diameters $D\\gtsim 10^2\\pc$, the assumption of static equilibrium is unlikely to be justified, because the sound-crossing time $(D/2)/15\\kms\\gtsim 3 \\Myr$ becomes longer than the lifetimes of high-mass stars. The fact that some \\ion{H}{2} region samples \\citep[e.g.,][]{Garay+Rodriguez+Moran+Churchwell_1993,Kim+Koo_2001} seem to obey a $\\nrms\\propto D^{-1}$ relationship appears to be an artifact of the sample selection. We see in Fig.\\ \\ref{fig:ne vs D}b that the overall sample of \\ion{H}{2} regions does not have a single $\\nrms$-vs.-$D$ relationship. But the observations appear to be generally consistent with the current models of dusty \\ion{H}{2} regions. \\subsection{Cavities in \\ion{H}{2} Regions: N49} Even without dust present, radiation pressure from photoelectric absorption by H and He can alter the density profile in a static \\ion{H}{2} region, lowering the central density and enhancing the density near the edge of the ionized region (see Fig.\\ \\ref{fig:nprofs gamma=0, 5, 10, 20}a). As seen in Figure \\ref{fig:Iprofs gamma=0, 5, 10, 20}a, for large values of $Q_0\\nrms$ the surface brightness profile can be noticeably flattened. If dust is assumed to be present, with properties typical of the dust in diffuse clouds, the equilibrium density profile changes dramatically, with a central cavity surrounded by a high-pressure shell of ionized gas pushed out by radiation pressure. In real \\ion{H}{2} regions, fast stellar winds will also act to inflate a low-density cavity, or ``bubble'', near the star; the observed density profile will be the combined result of the stellar wind bubble and the effects of radiation pressure. The GLIMPSE survey \\citep{Churchwell+Babler+Meade+etal_2009} has discovered and catalogued numerous interstellar ``bubbles''. An example is N49 \\citep{Watson+Povich+Churchwell+etal_2008}, with a ring of free-free continuum emission at 20~cm, surrounded by a ring of $8\\micron$ PAH emission. An O6.5V star is located near the center of the N49 ring. The image is nearly circularly symmetric, with only a modest asymmetry that could be due to motion of the star relative to the gas. The 20~cm image has a ring-peak-to-center intensity ratio $I({\\rm peak})/I({\\rm center})\\approx 2$. Is the density profile in N49 consistent with what is expected for radiation pressure acting on dust? From the 2.89~Jy flux from N49 at $\\lambda=20$~cm \\citep{Helfand+Becker+White+etal_2006} and distance $5.7\\pm0.6\\kpc$ \\citep{Churchwell+Povich+Allen+etal_2006}, the stellar source has $Q_{0,49}\\approx (0.78\\pm0.16)/\\fion$. If $\\fion\\approx 0.6$, then $Q_{0,49}\\approx (1.3\\pm0.3)$. The \\ion{H}{2} region, with radius $(0.018\\pm0.02)$~deg, has $\\nrms\\approx197\\pm63\\cm^{-3}$. Hence $Q_{0,49}\\nrms\\approx 260\\cm^{-3}$. If $\\sigma_{d,-21}=1$, then $\\tau_{d0}\\approx 1.3$. From Fig.\\ \\ref{fig:log fion vs taud0}a we confirm that $\\fion\\approx 0.6$ for $\\tau_{d0}\\approx 1.3$. Figure \\ref{fig:model summ}d shows that an \\ion{H}{2} region with $\\tau_{d0}=1.3$ is expected to have a central minimum in the emission measure, but with $I({\\rm peak})/I({\\rm center})\\approx 1.3$ for $\\beta=3,\\gamma=10$, whereas the observed $I({\\rm peak})/I({\\rm center})\\approx 2$. The central cavity in N49 is therefore significantly larger than would be expected based on radiation pressure alone. While the effects of radiation pressure are not negligible in N49, the observed cavity must be the result of the combined effects of radiation pressure and a dynamically-important stellar wind (which is of course not unexpected for an O6.5V star). \\subsection{Lyman-$\\alpha$} The original ionizing photon deposits a radial momentum $h\\nu_i/c$ at the point where it is absorbed by either a neutral atom or a dust grain. A fraction ($1-\\fion$) of the ionizing photons are absorbed by dust; this energy is reradiated isotropically, with no additional force exerted on the emitting material. Because the infrared optical depth within the \\ion{H}{2} region is small, the infrared emission escapes freely, with no dynamical effect within the \\ion{H}{2} region. A fraction $\\fion$ of the ionizing energy is absorbed by the gas. Subsequent radiative recombination and radiative cooling converts this energy to photons, but the isotropic emission process itself involves no net momentum transfer to the gas. We have seen above that the \\ion{H}{2} can have a center-to-edge dust optical depth $\\tau(R)\\approx 1.6$ for $\\tau_{d0}\\gtsim 5$, or $Q_{0,49}\\nrms\\gtsim10^{2}\\cm^{-3}$ (cf.\\ Fig.\\ \\ref{fig:model summ}c with $\\beta=3$, $\\gamma=10$). This optical depth applies to the $h\\nu>13.6\\eV$ ionizing radiation; the center-to-edge optical depth for the $h\\nu<3.4\\eV$ Balmer lines and collisionally-excited cooling lines emitted by the ionized gas will be significantly smaller, and much of this radiation will escape dust absorption or scattering within the \\ion{H}{2} region. That which is absorbed or scattered will exert a force on the dust at that point only to the extent that the diffuse radiation field is anisotropic. We conclude that momentum deposition from the Balmer lines and collisionally-excited cooling lines within the ionized zone will be small compared to the momentum deposited by stellar photons. Lyman-$\\alpha$ is a special case. At low densities ($n \\ll 10^3\\cm^{-3}$), $\\sim70\\%$ of Case B recombinations result in emission of a Ly-$\\alpha$ photon, increasing to $>95\\%$ for $n>10^5\\cm^{-3}$ as a result of collisionally-induced $2s\\rightarrow2p$ transitions \\citep{Brown+Mathews_1970}. After being emitted isotropically, the photon may scatter many times before either escaping from the \\ion{H}{2} or being absorbed by dust. Most of the scatterings take place near the point of emission, while the photon frequency is still close to line-center. On average, the net radial momentum transfer per emitted photon will likely be dominated by the last scattering event before the photon escapes from the \\ion{H}{2} region, or by the dust absorption event if it does not. At a given point in the nebula, the incident photons involved in these final events will be only moderately anisotropic. Since there is less than one Ly-$\\alpha$ photon created per case B recombination, the total radial momentum deposited by these final events will be a small fraction of the radial momentum of the original ionizing photons. \\citet{Henney+Arthur_1998} estimate that dust limits the Ly-$\\alpha$ radiation pressure to $\\sim$$6\\%$ of the gas pressure. We conclude that Ly-$\\alpha$ has only a minor effect on the density profile within the ionized zone. \\subsection{\\ion{H}{2} Region Expansion} \\ion{H}{2} regions arise when massive stars begin to emit ionizing radiation. The development of the \\ion{H}{2} region over time depends on the growth of the ionizing output from the central star, and the expansion of the initially-high pressure ionizing gas. Many authors \\citep[e.g.,][]{Kahn_1954,Spitzer_1978} have discussed the development of an \\ion{H}{2} region in gas that is initially neutral and uniform. If the ionizing output from the star turns on suddenly, the ionization front is initially ``strong R-type'', propagating supersonically without affecting the density of the gas, slowing until it becomes ``R-critical'', at which point it makes a transition to ``D-type'', with the ionization front now preceded by a shock wave producing a dense (expanding) shell of neutral gas bounding the ionized region. While the front is R-type, the gas density and pressure are essentially uniform within the ionized zone. When the front becomes D-type, a rarefaction wave propagates inward from the ionization front, but the gas pressure (if radiation pressure effects are not important) remains relatively uniform within the ionized region, because the motions in the ionized gas are subsonic. When radiation pressure effects are included, the instantaneous density profile interior to the ionization front is expected to be similar to the profile calculated for the static equilibria studied here. Let $V_i$ be the velocity of the ionization front relative to the star. When the ionization front is weak D-type, the velocity of the ionization front relative to the ionized gas just inside the ionization front is $\\sim0.5 V_i$ \\citep{Spitzer_1978}. Given the small dust drift velocities $v_{d,r}$ near the ionization front (i.e., $\\tau(r)\\rightarrow \\tau(R)$ in Fig.\\ \\ref{fig:vdrift gamma=10}), dust is unable to drift outward across the ionization front as long as the ionization front is propagating outward with a speed (relative to the ionized gas) $V_i\\gtsim 0.1\\kms$ Summary} \\begin{enumerate} \\item Dusty \\ion{H}{2} regions in static equilibrium consist of a three-parameter family of similarity solutions, parametrized by parameters $\\beta$, $\\gamma$, and a third parameter, which can be taken to be $Q_{0,49}\\nrms$ or $\\tau_{d0}$ (see eq.\\ \\ref{eq:taud0}). The $\\beta$ parameter (eq.\\ \\ref{eq:define beta}) characterizes the relative importance of $h\\nu<13.6\\eV$ photons, and $\\gamma$ (eq.\\ \\ref{eq:define gamma}) characterizes the dust opacity. A fourth parameter -- e.g., the value of $\\nrms$ or $Q_{0,49}$ -- determines the overall size and density of the \\ion{H}{2} region. \\item Radiation pressure acting on both gas and dust can strongly affect the structure of \\ion{H}{2} regions. For dust characteristic of the diffuse ISM of the Milky Way, static \\ion{H}{2} regions with $Q_{0,49}\\nrms\\ltsim 10^2\\cm^{-3}$ will have nearly uniform density, but when $Q_{0,49}\\nrms\\gg 10^2\\cm^{-3}$, radiation pressure acts to concentrate the gas in a spherical shell. \\item For given $\\beta$ and $\\gamma$, the importance of radiation pressure is determined mainly by the parameter $\\tau_{d0}$ (see eq.\\ \\ref{eq:taud0}). When $\\tau_{d0}\\gtsim 1$, radiation pressure will produce a noticeable central cavity. \\item If the dust-to-gas ratio is similar to the value in the Milky Way, then compression of the ionized gas into a shell limits the characteristic ionization parameter: $U_{1/2}\\ltsim 0.01$, even for $Q_0 \\nrms\\gg 1$ (see Fig.\\ \\ref{fig:U}). \\item For $\\tau_{d0}\\gtsim 1$, compression of the gas and dust into an ionized shell leads to a substantial {\\it increase} \\citep[compared to the estimate by][]{Petrosian+Silk+Field_1972} in the fraction $\\fion$ of $h\\nu>13.6\\eV$ photons that actually ionize H, relative to what would have been estimated for a uniform density \\ion{H}{2} region, as shown in Fig.\\ \\ref{fig:log fion vs taud0}. Eq.\\ (\\ref{eq:fion fit}) allows $\\fion$ to be estimated for given $Q_0\\nrms$, $\\beta$, and $\\gamma$. Galactic \\ion{H}{2} regions appear to have values of $\\fion$ consistent with the present results for \\ion{H}{2} regions with radiation pressure (see Fig.\\ \\ref{fig:inoue}). \\item Interstellar bubbles surrounding O stars are the result of the combined effects of radiation pressure and stellar winds. For the N49 bubble, as an example, the observed ring-like free-free emission profile is more strongly peaked than would be expected from radiation pressure alone, implying that a fast stellar wind must be present to help create the low-density central cavity. \\item For static \\ion{H}{2} regions, dust drift would be important on time scales $\\ltsim1\\Myr$ for $Q_{0,49}\\nrms\\gtsim 10^3\\cm^{-3}$. Real \\ion{H}{2} regions are not static, and the dust will not drift out of the ionized gas because the ionization front will generally be propagating (relative to the ionized gas just inside the ionization front) faster than the dust drift speed $\\ltsim 1\\kms$ (see Fig.\\ \\ref{fig:vdrift gamma=10}). \\end{enumerate}" }, "1003/1003.5121_arXiv.txt": { "abstract": "{HR\\,8799 is a $\\lambda$ Bootis, $\\gamma$ Doradus star hosting a planetary system and a debris disk with two rings. This makes this system a very interesting target for asteroseismic studies. In particular, this work is devoted to the determination of the internal metallicity of this star, linked with its $\\lambda$ Bootis nature, and its age, taking the advantage of its $\\gamma$ Doradus-type pulsations. To do so we have used the equilibrium code CESAM and the non-adiabatic pulsational code GraCo. We have applied the Frequency Ratio Method and the Time Dependent Convection theory to estimate the mode identification, the Brunt-Va\\\"is\\\"al\\\"a frequency integral and the mode instability, making a selection of the possible models fulfilling all observational constraints. Using the position of the star in the HR diagram, the solar metallicity models is discarded. This result contradicts one of the main assumptions of the most accepted hypothesis explaining the $\\lambda$ Bootis nature, the accretion/diffusion of gas from a star with solar metallicity. Therefore, in sight of these new results, a revision of this hypothesis is needed. The inclusion of accurate internal chemical mixing is necessary. The use of the asteroseismological constraints provides a very accurate determination of the physical characteristics of HR\\,8799: an age in the ranges [1123, 1625] and [26, 430] Myr, and a mass in the ranges [1.32, 1.33] and [1.44, 1.45] $M_{\\odot}$, respectively, depending on the visual angle $i$. The determination of this angle and more accurate multicolor photometric observations can definitively fix the mass, metallicity and age of this star. In particular, an accurate age estimation is needed for a correct understanding of the planetary system. In our study we have found that the age widely used for modelling the system is unlike.} ", "introduction": "Accurate ages stars hosting extrasolar planets are useful for, e.g., obtaining theoretical constraints on tidal interactions between Hot Jupiters and their hosts, dynamical modelling of exoplanetary systems, predicting exoplanetary luminosities, etc, since they constrain the age of the planets. The eventual aim of refining the age determination of the stars is to obtain more information on planetary system formation to discern between the different scenarios. Several methods can be applied to determine ages of stars, like the use of chromospheric activity proxies, lithium depletion, evolutionary track fitting, etc (Barrado y Navascu\\'es 1998; Barrado y Navascu\\'es et al. 1999). Nonetheless, for pulsating stars, asteroseismology is the best and most accurate method we know that provides mass, radii and ages with an accuracy order of magnitudes larger than obtained with other methods in the literature (Stello et al. 2009). The discovery by Marois et al. (2008), of the first planetary system by direct imaging around HR\\,8799 has made the age, mass and metallicity determination of this star a very important task. That was the starting point of a set of approximately ten studies about this system during 2009, none of them from an asteroseismic point of view, even considering that there are three detected pulsational frequencies (Zerbi et al. 1999) that can be used to better understand the host star. The main characteristics of this system are: \\begin{itemize} \\item The host star is a A5 Main Sequence star (Gray \\& Kaye 1999), $\\gamma$ Doradus pulsator (Zerbi et al. 1999), and it has $\\lambda$ Bootis-type surface chemical peculiarities (Sadakane 2006). \\item There are three sub-stellar objects orbiting the host star. They have been observed using direct imaging (Marois et al. 2008). \\item The system has a debris disk with two rings, one inner and another further out the three orbiting objects (Chen et al. 2009). \\item From a dynamical point of view, the stability of the system is a challenge, regarding the mass given for the objects and star, and the age assumed for the system (Reidemeister et al. 2009). \\end{itemize} These four characteristics make the planetary system HR\\,8799 a very interesting one, with many challenges to understand its structure, formation and evolution. In particular, the age of the star (and consequently of the system) is one of the key quantities for all the studies done around HR\\,8799, since it is directly related with the mass of the orbiting objects, and it imposes a goal for the dynamical stability of the system. In two recent works (Moya et al., 2010a, b), a comprehensive modelling of HR8799 using asteroseismology has been done. The main conclusions are the following: \\begin{itemize} \\item The use of asteroseismological observational data imposes large constraints to the acceptable models of stellar internal structure. \\item The actual uncertainty in the knowledge of the age of HR\\,8799 is larger than shown in the literature. Only less than 20\\% of the accepted model using all the observational constraints have an age in the range estimated by other works. \\item The internal subsolar metallicity obtained for the star provides some keys to understand the $\\lambda$ Bootis nature. \\end{itemize} These works have shown the enormous potential of the use of asteroseismology to understand this system and others. Nevertheless, these studies cannot be conclusive due to the poor asteroseismic observational available data. There is not very accurate multicolour photometric observations. Therefore, additional observations are imperative to provide better asteroseismological constraints. ", "conclusions": "The accurate determination of the mass, age and internal metallicity of a star hosting a planet or a planetary system is a necessary step to understand the formation and evolution of that planetary system. If the hosting star pulsates, the use of asteroseismology can provide these physical quantities with an accuracy hardly obtained with other techniques. The case of HR\\,8799 is a excellent example of this benefit provided by asteroseismology. In two recent works (Moya et al. 2010a, b), the first comprehensive asteroseismologic study of the planetary system host HR\\,8799, a $\\lambda$ Bootis star presenting $\\gamma$ Doradus pulsations has been carried out. This asteroseismic work is specially important for the determination of the internal abundances of this kind of stars, a previous step to understand the physical mechanism responsible for the surface chemical peculiarities of the $\\lambda$ Bootis group. On the other hand, an analysis of the age determination of the planetary system HR\\,8799 has also been done. The results found in the literature are not conclusive, and the only valid argument to estimate the age of the star is that using its radial velocity and proper motion, but it is an estatistical argument needed of additional estimations. The selection of models fulfilling the spectroscopic observations shows that there are no selected models with solar metallicity. This has an impact in the main assumption of the theory widely accepted to explain the $\\lambda$ Bootis nature, i.e. that these stars have solar metallicity, whereas the observed abundances are due to surface phenomena. On the other hand, the main complementary argument to the kinematic age determination is the position of the star in the HR diagram. We have shown that, due to the $\\lambda$ Bootis nature of HR\\,8799, that hides the internal metallicity of the star, this HR diagram position provides a range of ages [10,2337] Myr, a much wider than that estimated by other authors of [30,160] Myr (Marois et al. 2008). Only a small amount (18.1$\\%$) of models in our grid have ages in the range claimed in the literature. For inclination angles around $i=36^\\circ$, the models fulfilling all the observational constraints have masses in two separate ranges of M=[1.32, 1.33], [1.44, 1.45] $M_{\\odot}$, and the age of the system is constrained in two separate ranges: [1123, 1625] Myr and [26, 430] Myr respectively. A percentage of 16.7$\\%$ of the models are in the range given in the literature. This determination has an impact in the determination of the mass of the objects observed orbiting around HR\\,8799. A consequence of this study is the need for a precise determination of the inclination angle $i$, of the multicolour photometric amplitudes and phases of $f_2$, and some information of $m$ values through time-series if high resolution spectroscopy. These determinations would help to carry out a definitive selection of the models. In any case, the range of ages assigned to this star in the literature is unlikely to be the correct one. Only a stellar luminosity larger than that reported would allow young models with solar metallicity to fulfill all the observational constraints." }, "1003/1003.6064_arXiv.txt": { "abstract": "We analyze correlations between the first letter of the name of an author and the number of citations their papers receive. We look at simple mean counts, numbers of highly-cited papers, and normalized h-indices, by letter. To our surprise, we conclude that orthographically senior authors produce a better body of work than their colleagues, despite some evidence of discrimination against them. ", "introduction": "Citation counts and indices based on them are a key aspect of the sociology and practice of astrophysics, and of many other fields. Various metrics to associate the citation counts of an author with their value as a scientist have been developed, such as the h-index, h-b-index, and g-index\\cite{indices}, though all have been criticized as inappropriate in different contexts. Indeed, any metric based solely on citation counts is ultimately flawed, since it is unable to distinguish between citations of the form ``Work in \\cite{zunckel1} forms the basis of this discussion'' and ``We find the results in \\cite{zuntz1} to be utterly incorrect''. Some unscrupulous authors even inappropriately cite their own work to improve their indices. Despite their flaws, citation counts and indices certainly do have an important impact, both on science and scientists. Highly cited papers become part of the canon of standard works, and the authors of highly-cited papers are employed and feted\\footnote{An earlier version of this manuscript misspelled this word as \\emph{fetid}; we apologize for any confusion caused.}. Characterizations of patterns of citation can therefore be both metrics for wider socio-scientific trends and indicators of surprising aspects of paper generation, publication and promulgation systems. There have been previous studies of the correlation between number of citations and the length \\cite{Stanek1}, number of authors \\cite{Stanek2}, release timing \\cite{dietrich,haque}, field \\cite{leydesdorff} and publication method \\cite{gentil} of articles. Here we study another aspect of citation patterns: correlation between the number of citations received by an article and its orthography. Specifically, we correlate the number of citations a paper has with the first letter of its first author's surname. Selection by paper rather than by author does pose some problems, but it is significantly more tractable to collect such data given inconsistency in citation names. In Section \\ref{hypothesis and data section} we discuss the hypotheses we wish to test and the data used. In Section \\ref{analysis section} we analyze our data and discuss various citation measures. We conclude in Section \\ref{sec:discussion}. ", "conclusions": "Discussion} Based on the results generated, our third method of h-indices appears to be the most reliable; we are therefore reluctantly directed towards hypothesis H0, that orthographically high-ranking authors produce a better global body of work than their less alphabetically gifted colleagues. Two caveats apply to our selection of data. First, we are somewhat restrictive by limiting ourself to first authors only, particularly in light of the practice of alphabetizing authors on some papers, though this is more common for authors after the first. It is possible that any discrimination against or superiority of AWNNTEOTAs occurs at the point of choosing which author should be first. Second, we use the data themselves to estimate the raw number of authors with a given initial letter. This might skew results, if, for example, there are large numbers of authors with the surname ``Aaaronson'' who produced no papers at all during the referenced period; this would make our conclusions stronger. While future work on this topic could possibly improve on the analysis and interpretation methods we have used here, our conclusions can be made significantly more robust by improving the data directly \\cite{zwart}\\cite{zuntz2}\\cite{zuntz3}\\cite{zlosnik1}\\cite{zwolak}\\cite{zlosnik2}\\cite{zlosnik3}\\cite{zhang}\\cite{zunckel2}\\cite{zunckel3}\\cite{zurek}\\cite{zunckel4}. \\emph{Acknowledgments.} We would like to acknowledge the help of various colleagues in formulating this work, but they asked us not to because they didn't want their names associated with it. The exception was Laura Newburgh, who suggested our method for improving future data." }, "1003/1003.1644_arXiv.txt": { "abstract": "{ Hickson compact groups (HCGs) are among the densest galaxy environments of the local universe. To examine the effects of the environment on the infrared properties of these systems, we present an analysis of {\\em Spitzer} and{ \\em ISO } mid-infrared imaging, as well as of deep ground-based near-infrared imaging of 14 HCGs containing a total of 69 galaxies. Based on mid-infrared color diagnostics we identify the galaxies that appear to host an active nucleus, while using a suite of templates, and fit the complete infrared spectral energy distribution for each group member. We compare our estimates of galaxy mass, star formation rate, total infrared luminosities, and specific star formation rates (sSFR) for our HCG sample to samples of isolated galaxies and interacting pairs and find that overall there is no discernible difference among them. However, HCGs that can be considered as dynamically ``old'' host late-type galaxies with a slightly lower sSFR than the one found in dynamically ``young'' groups. This could be attributed to multiple past interactions among the galaxies in old groups, that have led to the build up of their stellar mass. It is also consistent with our prediction of the presence of diffuse cold dust in the intergalactic medium of 9 of the dynamically ``old'' groups.} ", "introduction": "Compact groups of galaxies were identified as systems of several galaxies that, owing to their small projected separations and signs of tidal distortion, appear to be real physical entities of gravitationally interacting systems \\citep{Hickson97}. Since the discovery in 1877 of Stephan's Quintet, the prototypical group, many others have been found by both visual and automated searches of the Palomar Observatory Sky Survey plates. Although the criteria for assigning group membership based only on imaging data have often been debated, it is now generally accepted that the most complete and better studied of these samples is the one compiled by \\citet{Hickson82}. His catalog of the so--called Hickson compact groups (hereafter HCGs) consists of 451 galaxies contained in 100 groups of four or more galaxies, that occupy compact configurations within relatively isolated regions where no excess of other surrounding galaxies can be seen \\citep[see][ for details on the criteria used]{Hickson82}. More recently, it has been revealed that the HCG catalog is incomplete because several compact groups are apparent associations of galaxies along the line of sight. However, compact group catalogs, produced with redshift information, to select galaxies that are physically close, are dense \\citep{Mamon09}. Over the past 20 years, various detailed analyses of HCGs have been performed using multi--wavelength imaging and spectroscopy. Based on optical morphology alone, it is obvious that the overwhelming majority of HCGs display an excess of elliptical galaxies ($\\sim$31\\% of the total) while their spiral fraction is just 43\\%, nearly a factor of two less than what is observed in the field \\citep{Hickson82}. \\citet{Rubin91} have found that two thirds of HCG spirals display peculiar rotation curves, while \\citet{Zepf93b} report that ellipticals in compact groups are likely to have irregular isophotes and exhibit lower internal velocity dispersions for their luminosities, so that they do not lie on the fundamental plane of isolated field ellipticals. Interestingly, the luminosity function derived for HCG members is deficient at the faint end compared to other samples. All these clues are consistent with an evolutionary pattern of tidal encounters and the accretion of small companions by the group members. Furthermore, \\citet{Mendes94} also show that 43\\% of all HCG galaxies display morphological features of interactions and mergers, such as bridges, tails, and other distortions. The HCGs appear to occupy a unique position in the range of galaxy environments found in the local universe. While their density enhancements are high, close to those seen in rich clusters, the overdensities appear to be more locally contained, with a much smaller population involved in the enhancement. Furthermore, the HCG velocity dispersions are $ \\sigma_V\\sim$250 \\,~km\\,s$^{-1}$, lower than what is seen in rich clusters but higher than that of typical loose groups. In addition, a clear correlation exists between the group velocity dispersion and the elliptical galaxy fraction, with the highest values of $\\sigma_V$ to be found in the most elliptical rich groups \\citep{Hickson88}. Despite this progress, several open questions remain. Are compact groups a transient phenomenon fading out after the merging of all their galaxies into a giant field elliptical? Were they more numerous in the early universe, and can they account for all, or most, present-day giant ellipticals? Are they dynamically closed systems, or can they replenish the intergalactic medium with reprocessed material in the form of diffuse tails and tidal dwarf galaxies? A necessary step in determining of the evolutionary state of HCGs is the complete census and analysis of the member galaxies' stellar population, gas content, and star formation properties. Even though only a few HCGs have been mapped in detail in neutral hydrogen, single-dish measurements reveal that they are generally deficient in HI with a median mass of M$_{\\rm HI}\\sim2.2\\times10^{10}\\,{\\rm M}_{\\odot}$, two times less than what is observed in loose groups \\citep{Williams87,Verdes01}. However, the molecular gas content of their individual galaxies, as traced by the CO emission, is similar to that of interacting pairs and starburst galaxies \\citep{Leon98}. Nearly 40\\% of the HCG members for which nuclear spectroscopy has been obtained display evidence of active galactic nuclei (AGN), and it has been argued that those contain copious amounts of dust \\citep{Shimada00}. Up to now, however, there is scant information on the far-IR colors and luminosities of the individual galaxies of the groups, because the available IRAS measurements have typically not resolved the groups, especially at 60 and 100\\,$\\mu$m, providing only a single value and color for the ensemble of their galaxies. Despite this limitation, it has been shown that compact groups have warm far-IR colors, similar to those of merging gas-rich galaxies \\citep{Zepf93a}. If one recalls that HCGs are in fact deficient in spirals and that the far-IR flux of field ellipticals is typically below the IRAS detection limit, one may speculate that vigorous star formation activity may actually take place in some locations within the groups. Could this activity be in circumnuclear regions of weak, enshrouded AGN because of the dynamical torques exerted in the gas from small nuclear bars funneling it to the center \\citep{Combes01}? Would it be possible instead that there is star formation due to shocked compressed gas outside the main bodies of galaxies as seen in Stephan's Quintet \\citep{Xu99, Appleton06}? A first analysis of the mid-infared properties of 12 HCG using Spitzer imagery has been presented by \\citet{Johnson07} and \\cite{Gallagher08}. These authors use near- and mid-IR color diagnostics to reveal a possible excess thermal emission due to an active nucleus (AGN). They also separate the groups into three types based on the ratio of their neutral hydrogen to dynamical mass, and find that most groups are either gas-rich or gas-poor, while very few groups are found in an intermediate state. As expected, gas-rich groups are the ones that are more actively forming stars. In addition, an absence of intermediate stage groups is apparent, which would be consistent with a rapid evolution of the group members, probably the result of dynamical interactions. In this paper we present the first results of our detailed analysis of the deep near-IR, mid-IR, and far-IR imagery of 14 HCGs and place them in the context of similar results for samples of isolated galaxies and early-stage interacting systems. In Section 2, we describe our data and reduction processes, along with the control samples used. In Section 3, we describe how we model the infrared spectral energy distribution (SED) of all galaxies in the groups. Our results on the physical properties of each galaxy such as star formation rate (SFR), stellar mass, dust content, and specific star formation rate (sSFR) are presented in Section 4. We discuss the implications of these findings in Section 5, and present our conclusions in Section 6. ", "conclusions": "In this paper we have presented our first analysis of the near- and mid-infrared SED of 69 galaxies contained in 14 Hickson compact groups, and found the following: - Nearly half of the galaxies in the groups, 14 out of the 32, which are optically classified as ellipticals, have mid-IR emission and colors consistent with those expected for late-type systems. We suggest that this stems from enhanced star formation as a result of recent gas accretion from companion galaxies in the groups. - Based on their integrated mid-IR color, 16 galaxies (23\\%) are AGN candidates. Detailed nuclear photometry reveals that 6 of these have nuclear mid-IR SEDs, which are inconsistent with a mid-IR spectrum dominated by star formation. - We find no evidence of the SFR and build up of stellar mass, which would substantially differentiate late-type galaxies in groups from galaxies in early-stage interacting pairs, or spiral galaxies in the field. This is a surprise given that the group environment has played an important role in the evolution of the galaxies, shown by the fact that most groups contain a large fraction of early-type systems. However, late-type galaxies in dynamically ``old'' HCGs might have a lower, though not statistically significant, sSFR than those in dynamically ``young'' groups, which could be attributed to multiple past interactions. - We investigated the contribution of each galaxy to the total far-IR emission of its group, and identify 9 groups where extended cold dust emission may be present." }, "1003/1003.4257_arXiv.txt": { "abstract": "We present Hubble Space Telescope observations in the far UV of the ultraluminous X-ray source in NGC 6946 associated with the optical nebula MF 16. Both a point-like source coincident with the X-ray source and the surrounding nebula are detected in the FUV. The point source has a flux of $5 \\times 10^{-16} \\rm \\, erg \\, s^{-1} \\, cm^{-2} \\, \\AA^{-1}$ and the nebula has a flux of $1.6 \\times 10^{-15} \\rm \\, erg \\, s^{-1} \\, cm^{-2} \\, \\AA^{-1}$, quoted at 1533 \\AA\\ and assuming an extinction of $A_V = 1.54$. Thus, MF 16 appears to host the first directly detected ultraluminous UV source (ULUV). The flux of the point-like source is consistent with a blackbody with $T \\approx 30,000$~K, possibly from a massive companion star, but this spectrum does not create sufficient ionizing radiation to produce the nebular He{\\sc ii} flux and a second, hotter emission component would be required. A multicolor disk blackbody spectrum truncated with an outer disk temperature of $\\sim 16,000$~K provides an adequate fit to the FUV, B, V, I, and He{\\sc ii} fluxes and can produce the needed ionizing radiation. Additional observations are required to determine the physical nature of the source. ", "introduction": "Ultraluminous X-ray sources (ULXs) are bright, irregularly variable, non-nuclear, X-ray sources with apparent luminosities exceeding the Eddington limit for a $20 M_{\\odot}$ compact object. ULXs could represent a new class of black holes with masses intermediate between stellar-mass and supermassive black holes \\citep{Colbert99,Makishima00,Kaaret01} or may be stellar-mass black holes with unusually high accretion rates and/or beamed emission. Over the past few years, optical and radio nebulae have been found which are spatially coincident and likely physically related to ULXs \\citep{Pakull02,Roberts03,Kaaret04,Lang07}. Study of nebulae associated with ULXs may enable us to determine the total radiation and particle fluxes from the ULXs if the nebulae are continuously powered by photoionization or particle flows. The nebula MF 16 in the nearby (5.1 Mpc) and nearly face-on spiral galaxy NGC 6946 has been described as an `ultraluminous supernova remnant complex' \\citep{Blair01}. The object is unusually luminous for a SNR in optical line emission, e.g.\\ the H$\\alpha$ luminosity is $2 \\times 10^{38} \\rm \\, erg \\, s^{-1}$, and in X-rays \\citep{Schlegel94}. \\citet{Roberts03} found that the X-ray source is variable and, thus, not diffuse X-ray emission associated with a SNR, but rather emission from an X-ray binary containing a compact object. MF 16 which had been previously thought of as a SNR is, instead, a nebula surrounding an X-ray binary. MF 16 shows many lines consistent with shock emission. However, several of the lines (e.g.\\ [O {\\sc iii}], H$\\alpha$, and [N {\\sc ii}]) contain a narrow component, generally associated with photoionization instead of shocks. \\citet{Blair01} concluded that unusually strong photoionization is needed, in addition to shocks, to explain the spectrum, but did not determine the source of ionizing photons. They also reported detection of He{\\sc ii} line emission, but did not discuss the origin of the line. \\citet{Abolmasov08} studied the optical emission from MF 16 with particular attention to the He{\\sc ii} $\\lambda 4686$ line. Their modeling shows that most of the power in the optical line emission comes from photoionization, although a component from shocks is still important. They find that an ultraviolet source more luminous than the observed ULX is needed to power the emission and that the UV source has a blackbody temperature of $10^{5.15 \\pm 0.10}$~K and a luminosity $1.2 \\times 10^{40} \\rm \\, erg \\, s^{-1}$. The luminosity is a lower bound on the source luminosity if the photoionized nebula is density bounded. The UV source proposed by \\citet{Abolmasov08} should be bright in the FUV. We obtained FUV observations using the solar blind channel (SBC) of the Advanced Camera for Surveys (ACS) on the Hubble Space Telescope (HST) to search for the predicted FUV emission. We describe the observations and analysis in \\S 2. We discuss the implications in \\S 3. \\begin{figure*} \\centerline{\\includegraphics[width=4.75in]{f1.eps}} \\caption{Images of MF 16. The panels show A: FUV (F140LP filter), B: B-band (F450W filter), C: V-band (F555W filter), and D: I-band (F814W filter). The small circles mark the counterparts to the X-ray source and have radii of $0.15\\arcsec$. The ellipses mark the extent of the nebula with a major diameter of $2.0\\arcsec$ and a minor diameter of $1.4\\arcsec$. The arrow in panel A points North and is $1\\arcsec$ long. \\label{images}} \\end{figure*} ", "conclusions": "FUV imaging of MF 16 reveals both a point-like source coincident with the X-ray source and a surrounding nebula with an angular extent of $2.0\\arcsec$ by $1.4\\arcsec$. The point source has a flux of $5 \\times 10^{-16} \\rm \\, erg \\, s^{-1} \\, cm^{-2} \\, \\AA^{-1}$ and the nebula has a flux of $1.6 \\times 10^{-15} \\rm \\, erg \\, s^{-1} \\, cm^{-2} \\, \\AA^{-1}$, both quoted at 1533 \\AA, corresponding to the pivot wavelength of the F140LP filter and assuming an extinction of $A_V = 1.54$. The monochromatic luminosity estimate is $\\lambda L_{\\lambda} = 1.0 \\times 10^{40} \\rm \\, erg \\, s^{-1}$, suggesting the presence of an ultraluminous UV source (ULUV). The measured flux is close to that predicted by \\citet{Abolmasov08} near 1000 \\AA. Figure~\\ref{mspec} shows the spectral energy distribution of the point-like source at the center of MF 16 using the data from the FUV (F140LP), B (F450W), V (F555W), and I (F814W) bands. The dereddened fluxes are from Table~\\ref{rates}. We note that the BVI data are quasi-simultaneous, obtained within 5 hours, while the FUV data were obtained 7 years later. Thus, the non-simultaneity of the observations should be of concern in fitting the data. However, \\citet{Fridriksson08} report that the rms variation in X-ray flux of this source in Chandra observations spread over several years is only 8\\%. Also, the I-band flux from the ACS/WFC measured in 2004 agrees with that measured with the WFPC2 in 2001 within 25\\%. Thus, fitting of these non-simultaneous data appears justified. A fit of a power-law, $F_{\\lambda} = A (\\lambda/6000 \\rm \\AA)^{\\alpha}$ to the BVI data dereddened assuming $A_V = 1.54$ gives best fit parameters $\\alpha = -4.1 \\pm 0.4$ and $A = (8.8 \\pm 0.8) \\times 10^{-18} \\rm \\, erg \\, s^{-1} \\, cm^{-2} \\, \\AA^{-1}$, see Fig.~\\ref{mspec}. Using $A_V = 1.8$ makes the curve somewhat steeper, $\\alpha = -4.4 \\pm 0.4$, while $A_V = 1.34$ leads to a slightly shallower slope, $\\alpha = -3.9 \\pm 0.4$. In all cases, a good fit is obtained with $\\chi^2 < 1$. The exponent of the the power-law is consistent with that expected for the Rayleigh-Jeans tail of a blackbody spectrum, $\\alpha = -4$. We fit all four data points with a single blackbody spectrum. This produces a good fit with a temperature $T =$~31000~K and a bolometric luminosity of $3.4\\times 10^{39} \\rm \\, erg \\, s^{-1}$. However, due to the relatively low temperature, the spectrum produces an insignificant flux of radiation capable of ionizing He. A second source producing radiation at shorter wavelengths is required to produce the He{\\sc ii} $\\lambda 4686$ luminosity seen from the nebula. Thus, we added a hotter blackbody with $T =$140,000~K and $L = 1.2 \\times 10^{40} \\rm \\, erg \\, s^{-1}$ as suggested by the emission line modeling \\citep{Abolmasov08}. The fit is acceptable, $\\chi^2/{\\rm DoF} = 3.8/2$, and the cooler blackbody temperature is reduced to $T =$26000~K and its luminosity to $2.2\\times 10^{39} \\rm \\, erg \\, s^{-1}$. These values are for $A_V = 1.54$. Varying the extinction in the range $A_V = 1.34-1.8$ and the hot blackbody parameters within the certainty range given by \\citet{Abolmasov08}, allows temperatures in the range 23,000--32,000~K and luminosities in the range $1.1-4.6 \\times 10^{39} \\rm \\, erg \\, s^{-1}$. These properties are consistent with a high mass star, specifically a late O or early B (O7 to B3) supergiant, that could be the companion star in the binary system. It would be of interest to obtain an optical and/or UV spectrum of the point source to search for spectral lines that would permit classification of the star, or rule out interpretation in terms of a stellar companion. The spectra of accretion disks around black holes are often described by a multicolor disk blackbody (MCD) spectrum where the local effective temperature is a decreasing function of radius, usually $T(r) \\propto r^{-p}$, between an inner disk radius, $R_{\\rm in}$, where the temperature is $T_{\\rm in}$ and an outer radius, $R_{\\rm out}$, with temperature $T_{\\rm out}$. At long wavelengths, $ hc/\\lambda \\ll kT_{\\rm out}$, the spectrum has the form of a Rayleigh-Jeans tail and at intermediate wavelengths, $kT_{\\rm out} \\ll hc/\\lambda \\ll kT_{\\rm in}$, the spectrum flattens. The observed FUV/optical spectrum of the point source is compatible with this form if the outer disk temperature lies between the FUV and B bands. The FUV/optical data can constrain MCD model spectra only in regards to the outer disk temperature and the flux in the FUV/optical range. A constraint on the short wavelength end of the spectrum can be obtained from the ionizing luminosity needed to produce the nebular He{\\sc ii} $\\lambda 4686$ emission. The solid lines in Fig.~\\ref{mspec} show MCD spectra consistent with the FUV/optical data and with a UV luminosity of $L(20 < \\lambda < \\rm 228 \\AA) \\approx 4 \\times 10^{39} \\rm \\, erg \\, s^{-1}$ as needed to produce the observed He{\\sc ii} line \\citep{Abolmasov08}. We consider two disk temperature profiles: a standard profile with $p = 0.75$ and a profile with $p = 0.5$ as found in spectral fits to a number of ULXs \\citep{Vierdayanti06}. Both fits are acceptable with $\\chi^2 < 3.4$. For $p = 0.75$, the bolometric luminosity is $L_{\\rm Bol} = 1.2 \\times 10^{40} \\rm \\, erg \\, s^{-1}$ and $T_{\\rm out} = 16000 \\pm 2000 $~K. However, the disk is unusually narrow, $R_{\\rm out}/R_{\\rm in}\\sim 40$. For $p = 0.5$, we find that $L_{\\rm Bol} = 8 \\times 10^{40} \\rm \\, erg \\, s^{-1}$, $T_{\\rm out} = 21000 \\pm 2000 $~K, and $R_{\\rm out}/R_{\\rm in}\\sim 6000$. The true situation may be more complex with $p \\sim 0.5$ only in the inner disk and $R_{\\rm out}/R_{\\rm in}$ intermediate between these two extremes. It is also be possible to fit the data using models where the outer disk is irradiated. However, the complexity of irradiation, and the number of parameters in such models, preclude obtaining useful constraints with the available data. In all cases, the outer disk radius is of order $10^{12} \\rm \\, cm$. In summary, two possible scenarios to explain the observational data are that the emission is the sum of a hot blackbody from the compact object and stellar emission from the companion star or that the emission is due to a multicolor disk with an outer disk radius $\\sim 10^{12} \\rm \\, cm$. The small number of broad band measurements currently available is inadequate to distinguish between these two cases. Additional observations are needed to determine the physical nature of the source." }, "1003/1003.2251_arXiv.txt": { "abstract": "Based on the drift-kinetic theory, we develop a model for particle acceleration and transport in solar flares. The model describes the evolution of the particle distribution function by means of a numerical simulation of the drift-kinetic Vlasov equation, which allows us to directly compare simulation results with observations within an actual parameter range of the solar corona. Using this model, we investigate the time evolution of the electron distribution in a flaring region. The simulation identifies two dominant mechanisms of electron acceleration. One is the betatron acceleration at the top of closed loops, which enhances the electron velocity perpendicular to the magnetic field line. The other is the inertia drift acceleration in open magnetic field lines, which produces antisunward electrons. The resulting velocity space distribution significantly deviates from an isotropic distribution. The former acceleration can be a generation mechanism of electrons that radiate loop-top nonthermal emissions, and the latter be of escaping electrons from the Sun that should be observed by {\\insitu} measurements in interplanetary space and resulting radio bursts through plasma instabilities. ", "introduction": "\\label{sec:introduction} { Many observations with such as hard X-rays (HXRs), $\\gamma$-rays, and microwaves have revealed that a solar flare is one of the strongest particle accelerator in our solar system. Electrons that are accelerated to several tens of keV to MeV radiate HXRs at footpoints of flare loops \\citep[e.g.,][]{1994PhDT.......335S,2003ApJ...595L.103K,2009ApJ...697..843M}, and sometimes above the top of a soft X-ray loop \\citep{1994Natur.371..495M,1995PASJ...47..677M}. They are also observed via microwaves in the gigahertz band. These emissions have been used to understand the properties of accelerated electrons \\citep[e.g.,][]{1983Natur.305..292N,1988ApJ...324.1118K,2000ApJ...545.1116S,2001ApJ...557..880K,2002ApJ...576..505W,2008ApJ...673..598M}. Recently, accelerated ions have been studied by the $\\gamma$-ray observations of the {\\it Reuven Ramaty High Energy Solar Spectroscopic Imager} \\citep[{\\it RHESSI};][]{2003ApJ...595L..69L,2006ApJ...644L..93H}.} In addition to these observations of accelerated particles at flare sites, escaping particles from the Sun into interplanetary space, the so-called solar energetic particles (SEPs), and resulting radio bursts are observed in association with flares \\citep[e.g.,][]{1982ApJ...253..949L,1985SoPh..100..537L,1998ApJ...503..435E,1999SSRv...90..413R,1999ApJ...519..864K,2007ApJ...663L.109K,2009ApJ...691..806K}. \\cite{2007ApJ...663L.109K,2009ApJ...691..806K} have reported detailed examinations of the relationship of electrons between at the flare site and in interplanetary space, by using {\\RHESSI} and {\\it WIND} observations. On the other hand, escaping electrons are often observed with no corresponding flare activity \\citep[e.g.,][]{1980ApJ...236L..97P,2003GeoRL..30m..30G,2007ApJ...657..567M,2010ApJ...708L..95E}. This indicates that small-scale particle accelerations frequently take place in the corona. { Based on the magnetic reconnection scenario \\citep{1995ApJ...451L..83S}, many authors have proposed models to explain the particle acceleration in solar flares. The particle acceleration in and near the reconnection region has been well discussed. For example, \\cite{1996ApJ...462..997L} has studied the DC electric field acceleration in reconnecting current sheets. \\cite{2001JGR...10625979H} have discussed the nonadiabatic acceleration in the vicinity of the reconnection region at which the gyro radius is comparable to the curvature radius of the magnetic field line. Fermi acceleration by the contraction of magnetic islands has been studied by \\cite{2006Natur.443..553D}.} { In addition to the acceleration near the reconnection region, stochastic acceleration mechanisms in the reconnection downstream region have been proposed by e.g., \\cite{1996ApJ...461..445M} (fast-mode magnetohydrodynamic waves) and \\cite{1998ApJ...495L..67T} (oblique shocks), although necessary high-frequency waves for the particle scattering is poorly known in the solar circumstance.} {The models greatly contribute to the understanding of how particles are accelerated to observed energies. For the complete understanding of many observed features, however, we should take into account the particle transport and dissipation processes as well as the acceleration in a realistic environment of the solar corona. This is because observed quantities do not necessarily reflect the particle distribution just at the acceleration site, but are the convolution of all of these processes \\citep{1998ApJ...502..455A}}. \\cite{1997ApJ...485..859S} have proposed a model called ``collapsing trap'', in which the particle acceleration, transport, and dissipation naturally follow from the shrinkage of magnetic loops driven by reconnection. These processes are almost adiabatic in the trap, because the particle gyro scale ($\\sim 10^{-9} \\; {\\rm s} \\; {\\rm and} \\sim 10^{0} \\; {\\rm cm}$ for electrons) is much smaller than the flare scale ($\\sim 10^{2} \\; {\\rm s} \\; {\\rm and} \\sim 10^{9} \\; {\\rm cm}$) due to very strong magnetic fields in the corona. The adiabatic acceleration may take place everywhere in the trap, which diminishes a serious problem that the number of accelerated particles required for observed HXR intensities is comparable to the number in a whole flare region \\citep{1974IAUS...57..105K,1997JGR...10214631M}. { The mechanism can be regarded as the subsequent acceleration process that occurs after the acceleration near the reconnection region, and as the co-existing process with the acceleration in the downstream region.} It is noteworthy to note that the model is very similar to the ``dipolarization'' associated with substorms in the terrestrial magnetosphere. In the collapsing trap, particles can adiabatically gain energy from convection electric fields ${\\vect E} = -{\\vect v} \\times {\\vect B}$ through $\\nabla B$ and inertia drift motions. \\cite{2004A&A...419.1159K} have studied the electron distribution function in the shrinking loop, with considering only the acceleration due to the $\\nabla B$ drift (the betatron acceleration). \\cite{2004ApJ...608..554A} has analytically considered the electron transport and the resulting impulsive HXR emissions in the loop. \\cite{2005ApJ...635..636G} have presented a rigorous treatment of the adiabatic motion of a single particle in the loop, by numerically solving guiding-center equations of motion. \\cite{2006ApJ...647.1472K} have performed a guiding-center test particle simulation in a cusp-shaped loop obtained from a magnetohydrodynamic (MHD) simulation. However, a complete modeling of the evolution of the adiabatic particle distribution in a flare region has not been studied yet. The correlation between the intensity of HXRs and the strength of convection electric fields has been reported by e.g., \\cite{2002ApJ...565.1335Q} and \\cite{2003ApJ...586..624A,2004ApJ...611..557A}. Recently, \\cite{2008ApJ...672L..69L} have observed the correlation between the hardness of the HXR spectrum and the strength of convection electric fields along flare ribbons. These observations indicate that the electric field greatly influences the distribution of electrons during their travel in the corona. The adiabatic model with actual coronal parameters can be tested through a direct comparison with these observations, because the model simultaneously describes particle acceleration, transport, and dissipation processes in a very wide area of the corona where particles are strongly magnetized. To understand the particle acceleration, transport, and dissipation processes in solar flares, we theoretically study the evolution of the particle distribution based on the numerical simulation. For this purpose, we newly develop a comprehensive model based on the drift-kinetic theory, in which the evolution of the particle distribution function is described by means of a numerical simulation of the drift-kinetic Vlasov equation. The simulation can be performed with actual coronal parameters, allowing us to directly compare it with observations. In {\\S}~\\ref{sec:models} we describe equations of particles and electromagnetic fields in our model. Simulation results are shown in {\\S}~\\ref{sec:calculation-results}. The simulation identifies two dominant mechanisms of electron acceleration; the betatron acceleration at the top of closed loops and the inertia drift acceleration in open magnetic field lines. We discuss the results in {\\S} \\ref{sec:discussion}, and then conclude the paper in {\\S} \\ref{sec:conclusion-1}. ", "conclusions": "\\label{sec:conclusion-1} Based on the drift-kinetic theory, we have developed a comprehensive model for particle acceleration and transport in solar flares. Using this model, we have simulated the time evolution of the electron distribution in a flaring region. There are two dominant mechanisms of electron acceleration. The betatron acceleration takes place at the top of closed loops. It can be a generation mechanism of electrons that radiate loop-top nonthermal emissions. The phase space distribution of accelerated particles in the loop strongly depends on the magnetic field configuration. The inertia drift acceleration also takes place in open magnetic field lines, producing escaping electrons from the Sun. The number of escaping electrons estimated from the simulation can account for the observed number of flare-associated escaping electrons. The inertia drift acceleration is caused by a motion of curved field lines, which is not necessarily driven by magnetic reconnection. In this sense, we propose the acceleration in this study as a mechanism for producing escaping electrons that are not always associated with flares." }, "1003/1003.2317_arXiv.txt": { "abstract": "A short overview is presented of current issues concerning the production and evolution of Li, Be and B in the Milky Way. In particular, the observed ``primary-like\" evolution of Be is re-assessed in the light of a novel idea: it is argued that Galactic Cosmic Rays are accelerated from the wind material of {\\it rotating} massive stars, hit by the forward shock of the subsequent supernova explosions. The pre-galactic levels of both Li isotopes remain controversial at present, making it difficult to predict their Galactic evolution. A quantitative estimate is provided of the contributions of various candidate sources to the solar abundance of Li. ", "introduction": "The idea that the light and fragile elements Li, Be and B are produced by the interaction of the energetic nuclei of galactic cosmic rays (CGR) with the nuclei of the interstellar medium (ISM) was introduced 40 years ago (Reeves et al. 1970, Meneguzzi et al. 1971, hereafter MAR). In those early works it was shown that, taking into account the relevant cross-sections and with plausible assumptions about the GCR properties - source composition, intensity and spectrum - one may reproduce reasonably well the abundances of those light elements observed in GCR and in meteorites (pre-solar). Among the required ingredients for such a calculation, the relevant spallation cross sections of CNO nuclei are accurately measured in the laboratory. The source composition and the equilibrium energy spectrum of GCR are inferred from a combination of observations and models of GCR propagation in the Milky Way (e.g. in the framework of the so-called ``leaky box\" model). Once the equilibrium spectra of GCR in the ISM are established, the calculation of the resulting abundances of LiBeB is straightforward, at least to first order\\footnote{The full calculation should include production by spallation of other primary and secondary nuclides, such as $^{13}$C; however, this has only second order effects.}. The production rate (s$^{-1}$) of the abundance $Y_L=N_L/N_H$ (by number) of LiBeB nuclei is given by \\begin{equation} \\frac{dY_L}{dt} \\ = \\ F^{GCR}_{p,a}\\sigma_{pa+CNO}Y^{ISM}_{CNO} \\ + F^{GCR}_{CNO}\\sigma_{pa+CNO}Y^{ISM}_{p,a} P_L\\ + F^{GCR}_{a}\\sigma_{a+a}Y^{ISM}_{a} P_L \\end{equation} where: $F$ (cm$^{-2}$ s$^{-1}$) is the average GCR flux of protons, alphas or CNO, $Y$ the abundances by number of those nuclei in the ISM, and $\\sigma$ (cm$^2$) is the average (over the equilibrium energy spectrum of GCR) cross-section for the corresponding spallation reactions producing LiBeB. The first term in the right hand member of this equation (fast protons and alphas hitting CNO nuclei of the ISM) is known as the ``direct\" term, the second one (fast CNO nuclei being fragmented on ISM protons and alphas) is the ``reverse\" term and the last one involves ``spallation-fusion\" reactions, concerning only the Li isotopes. $P_L$ is the probability that nuclide $L$ (produced at high energy) will be thermalized and remain in the ISM (see, e.g. Prantzos 2006). Obviously, the GCR flux term $F^{GCR}_{CNO} \\propto Y^{GCR}_{CNO}$ is proportional to the abundances of CNO nuclei in GCR, a fact of paramount importance for the evolution of Be and B (see next sections). Substituting appropriate values for GCR fluxes ($F^{GCR}_p\\sim$10 p cm$^{-2}$ s$^{-1}$ for protons and scaled values for other GCR nuclei), for the corresponding cross sections (averaged over the GCR equilibrium spectrum $\\sigma_{p,a+CNO\\longrightarrow Be}\\sim$10$^{-26}$ cm$^{-2}$) and for ISM abundances $Y_{CNO}\\sim$10$^{-3}$, and integrating for $\\Delta t \\sim$10$^{10}$ yr, one finds $Y_{Be}\\sim$2 10$^{-11}$, i.e. approximately the meteoritic Be value. Satisfactory results are also obtained for $^6$Li and $^{10}$B. Two problems were identified with the GCR production, compared to meteoritic composition: the \\lis ratio ($\\sim$2 in GCR, but $\\sim$12 in meteorites) and the \\bb ratio ($\\sim$2.5 in GCR, but $\\sim$4 in meteorites). It was then suggested in MAR that supplementary sources are needed for $^7$Li and $^{11}$B. Modern solutions to those problems involve {\\it stellar} production of $\\sim$60\\% of $^7$Li (in the hot envelopes of AGB stars and/or novae, see Sec. 7) and of $\\sim$40\\% of $^{11}$B (through $\\nu$-induced spallation of $^{12}$C in SN, see Sec. 5). In both cases, however, uncertainties in the yields are such that observations are used to constrain the yields of the candidate sources rather than to confirm the validity of the scenario. \\begin{figure}[t!] \\begin{center} \\includegraphics[width=0.99\\textwidth]{Prantzos_Be_OFe.eps} \\caption[]{ Observations of Be vs. Fe ({\\it left}) and vs. O ({\\it right}). In both panels, dotted lines indicate slopes of 1 (primary) and 2 (secondary). Be clearly behaves as a primary vs. Fe, whereas there is more scatter in the data vs. O. } \\label{eps1} \\end{center} \\end{figure} \\begin{figure}[t!] \\begin{center} \\includegraphics[angle=-90,width=0.7\\textwidth]{Prantzos_Energy.eps} \\caption[]{Energy input required from energetic particles accelerated by one CCSN in order to produce a given mass of Be, such as to have [Be/Fe]=0 (solar), assuming that a core collapse SN produces, on average, 0.1 M$_{\\odot}$ of Fe. {\\it Solid} curve corresponds to the case of a constant composition for GCR, {\\it dotted} curve corresponds to a time variable composition, following the one of the ISM. In the former case, the required energy is approximately equal to the energy imparted to energetic particles by supernovae, namely $\\sim$0.1 of their kinetic energy of $\\sim$1.5 10$^{51}$ ergs; in the latter case, the energy required to keep [Be/Fe]=0 becomes much larger than the total kinetic energy of a CCSN for metallicities [Fe/H]$\\leq$-1.6.} \\label{eps1} \\end{center} \\end{figure} ", "conclusions": "" }, "1003/1003.2910_arXiv.txt": { "abstract": "We present a multi-epoch and multi-frequency VLBI study of the compact radio source J0650+6001. In VLBI images the source is resolved into three components. The central component shows a flat spectrum, suggesting the presence of the core, while the two outer regions, with a steeper spectral index, display a highly asymmetric flux density. The time baseline of the observations considered to derive the source expansion covers about 15 years. During this time interval, the distance between the two outer components has increased by 0.28$\\pm$0.13 mas, that corresponds to an apparent separation velocity of $0.39c \\pm 0.18c$ and a kinematic age of 360$\\pm$170 years. On the other hand, a multi-epoch monitoring of the separation between the central and the southern components points out an apparent contraction of about $0.29 \\pm 0.02$ mas, corresponding to an apparent contraction velocity of $0.37c \\pm 0.02c$. Assuming that the radio structure is intrinsically symmetric, the high flux density ratio between the outer components can be explained in terms of Doppler beaming effects where the mildly relativistic jets are separating with an intrinsic velocity of $0.43c \\pm 0.04c$ at an angle between 12$^{\\circ}$ and 28$^{\\circ}$ to the line of sight. In this context, the apparent contraction may be interpreted as a knot in the jet that is moving towards the southern component with an intrinsic velocity of $0.66c \\pm 0.03c$, and its flux density is boosted by a Doppler factor of 2.0. ", "introduction": "The evolutionary stage of a powerful radio source originated by an active galactic nucleus (AGN) is related to its linear size. In this context, compact symmetric objects (CSO), which are powerful ($L_{\\rm 1.4\\,GHz} > 10^{24}$ W/Hz) and intrinsically small (i.e. linear size LS $<$ 1 kpc) radio sources, should represent a young stage in the individual radio source life. These objects usually have a roughly symmetric structure, with core, jets, and mini-lobes/hotspots, resembling a scaled-down version of the large classical radio galaxies which conversely have linear sizes of hundreds of kpc or even a few Mpc. The main characteristic displayed by young radio sources is the convex synchrotron radio spectrum that peaks at frequencies in the GHz regime. The spectral turnover is usually due to synchrotron self-absorption \\citep[SSA;][]{mo08,snellen00}, although an additional contribution from free-free absorption (FFA) has been found in the most compact objects \\citep[e.g.][]{dd09,mutoh02,kameno00}.\\\\ Conclusive evidence of the genuine {\\it youth} of this class of objects came from the determination of both kinematic and radiative ages.\\\\ The determination of the kinematic age is based on multi-epoch milliarcsecond-scale resolution observations spanning several years, and able to measure at which rate the hotspots are increasing their separation. Assuming that the separation velocity $v_{\\rm sep}$ has maintained constant, the kinematic age $t_{\\rm kin}$ can be estimated:\\\\ \\begin{displaymath} t_{\\rm kin} = {\\rm LS} \\cdot v_{\\rm sep}^{-1} \\, . \\end{displaymath} \\noindent From the analysis of a dozen of the most compact CSOs \\citep[LS $<$ 20 pc,][]{polatidis03}, it has been derived that the separation speed is generally in the range of 0.1$c$ and 0.4$c$, leading to kinematic ages of a few thousand years.\\\\ The radiative age can be estimated by multi-frequency observations able to constrain the radio spectral shape and thus to determine at which frequency the spectral break occurs. The break frequency $\\nu_{\\rm br}$ is strictly related to the radiative lifetime of the synchrotron emitting electrons $t_{\\rm syn}$. Once the magnetic field is known, for example assuming minimum total energy content corresponding to equipartition between particles and magnetic field ($H$) energies \\citep{pacho70}, and following some assumptions \\citep{mm03}, the radiative age can be easily estimated by means of:\\\\ \\begin{displaymath} t_{\\rm syn} = {\\rm const} \\cdot \\nu_{\\rm br}^{1/2} H^{-3/2} \\, . \\end{displaymath} \\noindent Several studies aimed at estimating the radiative age in compact radio objects indicate ages between 10$^{3}$ - 10$^{4}$ years \\citep{mm99,mm03}, in excellent agreement with the kinematic ages.\\\\ Given all the assumptions mentioned earlier, these methods do not provide the accurate source age, that can be improved either by increasing the time interval spanned by the observations or through a better sampling of the frequency coverage used to derive the radio spectrum.\\\\ Among all the CSOs studied in this framework, the radio source J0650+6001 represents a peculiar case. This source is identified with a quasar at redshift $z = 0.455$, and its optical spectrum is characterized by very weak broad lines and prominent narrow lines \\citep{stickel93}. It has a non-aligned triple radio structure with a total angular size of 7 mas ($\\sim$40 pc), and its total radio spectrum turns over around 5.5 GHz \\citep{mo07}. From a previous study no evidence of proper motion has been found \\citep{mo08}, although the resolution at 5 GHz (i.e. the frequency at which several observations spanning almost a decade have been performed) was not adequate to reliably estimate small changes in the component positions. However, another analysis of the position of the source components based on global VLBI observations \\citep{akujor96}, indicated an apparent contraction between the central and the southern components. \\\\ In this paper we report on the results on multi-epoch VLBI and 5-GHz VLA data of the compact symmetric object J0650+6001, and we present an interpretation to explain the radio properties of this source.\\\\ Throughout this paper, we assume the following cosmology: $H_{0} = 71\\, {\\rm km\\, s^{-1}\\, Mpc^{-1}}$, $\\Omega_{\\rm M} = 0.27$, and $\\Omega_{\\Lambda} = 0.73$, in a flat Universe. At the redshift of the target 1$^{''} = 5.773$ kpc. The spectral index is defined as $S {\\rm (\\nu)} \\propto \\nu^{- \\alpha}$.\\\\ \\begin{table} \\caption{Log of the archival VLBA observations analyzed in this paper.} \\begin{center} \\begin{tabular}{ccccc} \\hline Freq.&Obs. date&Beam&$uv_{\\rm max}$&Antennas\\\\ GHz& &mas$^{2}$&M$\\lambda$& \\\\ \\hline &&&&\\\\ 5.0&22 Nov 1999&2.46$\\times$1.94&125&VLBA -Sc -Nl\\\\ 5.0&17 Dec 2004&1.65$\\times$0.88&172&VLBA + Eff\\\\ 8.4&17 Dec 2004&0.90$\\times$0.54&290&VLBA + Eff\\\\ &&&&\\\\ \\hline \\end{tabular} \\label{vlba} \\end{center} \\end{table} \\begin{table} \\caption{Multi-epoch VLA flux density at 5 GHz of the source J0650+6001. Column 1: flux density in mJy; Col. 2: observing date; Col. 3: Reference: 1: \\citet{ulvestad81}; 2: \\citet{perley82}; 3: \\citet{odea90}; 4: this work; 5: \\citet{dd00}, 6: \\citet{tinti05}; 7: \\citet{mo07}; 8: \\citet{mo08}.} \\begin{center} \\begin{tabular}{ccc} \\hline Flux&Obs. date&Ref.\\\\ \\hline &&\\\\ 906&2 Feb 1979&1\\\\ 830& 18 Nov 1980&2\\\\ 916&30 Dec 1984&3\\\\ 940&21 Jul 1991&4\\\\ 969&23 Feb 1993&4\\\\ 1058&24 Sep 1995&4\\\\ 1236&14 Nov 1998&5\\\\ 1136&3 Jul 2002&6\\\\ 1150&14 Sep 2003&7\\\\ 1106&28 Jan 2004&7\\\\ 1180&19 Nov 2006&8\\\\ &&\\\\ \\hline \\end{tabular} \\label{oss_vla} \\end{center} \\end{table} ", "conclusions": "We have presented multi-epoch VLBI and VLA observations of the radio source J0650+6001. The source shows a peculiar triple structure, where the outer components are misaligned with respect to the central flat-spectrum region, and show a high flux density asymmetry. The comparison between multi-epoch high spatial resolution observations indicated that the farthest components are increasing their distance with an apparent separation velocity $v_{\\rm s,a} = 0.39c \\pm 0.18c$. Assuming that this velocity is representative of the mean separation speed, we obtain a source kinematic age of 360$\\pm$170 years. On the other hand, the distance between the central component and the southern lobe has decreased with a mean apparent contraction speed $v_{\\rm c,a}= 0.37c \\pm 0.02c$. To explain these peculiar characteristics, we discussed an interpretation based on relativistic beaming effects. From this analysis we drew a picture in which the source is oriented with an angle between 12$^{\\circ}$ and 28$^{\\circ}$ to the line of sight with the southern component being the approaching side of the source. The intrinsic separation velocity is $v_{\\rm s,i} = 0.43c \\pm 0.04c$. The apparent contraction observed between the central and the southern components may be explained in terms of a knot in the jet that dominates the radio emission of the central component, that is moving towards the southern component with a mildly relativistic intrinsic velocity $v_{\\rm c,i} = 0.66c \\pm 0.03c$, and with a Doppler factor $D = 2$, which causes a moderate boost of its flux density. The monitoring of the flux density variability and the study of the component separations on a longer time scale will be crucial to better determine the nature of this source.\\\\" }, "1003/1003.4531_arXiv.txt": { "abstract": "It is shown that a first-order cosmological perturbation theory for the open, flat and closed Friedmann-Lema\\^\\i tre-Robertson-Walker universes admits one, and only one, gauge-invariant variable which describes the perturbation to the energy density and which becomes equal to the usual Newtonian energy density in the non-relativistic limit. The same holds true for the perturbation to the particle number density. Using these two new variables, a new manifestly gauge-invariant cosmological perturbation theory has been developed. Density perturbations evolve diabatically. Perturbations in the total energy density are gravitationally coupled to perturbations in the particle number density, irrespective of the nature of the particles. There is, in first-order, no back-reaction of perturbations to the global expansion of the universe. Small-scale perturbations in the radiation-dominated era oscillate with an increasing amplitude, whereas in older, less precise treatments, oscillating perturbations are found with a decreasing amplitude. This is a completely new and, obviously, important result, since it makes it possible to explain and understand the formation of massive stars after decoupling of matter and radiation. ", "introduction": "\\label{sec:introduction} Since measurements of the fundamental parameters of our universe are today very precise, cosmology is nowadays a mature branch of astrophysics. Despite advances in observational as well as theoretical cosmology, there is, as yet, no manifestly covariant cosmological perturbation theory. It is the purpose of this article to fill up this gap. We redo the calculations of Lifshitz~\\cite{lifshitz1946} and Lifshitz and Khalatnikov~\\cite{c15}, and we use new results of the literature, which were not known at the time when Lifshitz and Khalatnikov developed their theory. Combining these results with new insights obtained by us and described in detail in our report~\\cite{miedema-2008}, we are able to develop a perturbation theory with the properties that both the evolution equations and their solutions are invariant under general infinitesimal coordinate transformations $x^\\mu\\rightarrow x^\\mu-\\xi^\\mu(x^\\nu)$. We refer to such a theory as a \\emph{manifestly covariant gauge-invariant} perturbation theory. \\subsection{Former Insights} \\label{sec:ins-lit} Firstly, it is mandatory~\\cite{c13,mfb1992,Ellis1,Ellis2,ellis-1998,Mukhanov-2005} to use gauge-invariant variables to construct a perturbation theory. Secondly, although Lifshitz and Khalatnikov were aware of the fact that the system of perturbation equations can be divided into three independent systems of equations, namely for gravitational waves, vortices, and scalar perturbations, they did not use explicitly the decompositions (\\ref{eq:decomp}) and (\\ref{eq:div-curl-u}) found by York~\\cite{York1974} and Stewart~\\cite{Stewart}. This decomposition makes the computations much more tractable. Finally, an important result~\\cite{York1974,Stewart} is that the perturbed metric tensor for scalar perturbations can be written in terms of two potentials (\\ref{decomp-hij-par}). This facilitates the derivation of the Newtonian results (\\ref{eq:poisson}) and (\\ref{eq:newt-ngi}) in the non-relativistic limit. \\subsection{Our New Insights} \\label{sec:our-ins} In addition to the results found in the literature, we use the new insights obtained by us~\\cite{miedema-2008}. Firstly, we have found, using the combined First and Second Law of Thermodynamics, that, in general, density perturbations do not evolve adiabatically, as we will demonstrate in Section~\\ref{sec:diabatic}\\@. Therefore, it is essential that in a cosmological perturbation theory one uses an equation of state for the pressure $p=p(n,\\varepsilon)$ rather than $p=p(\\varepsilon)$, where $n$ is the particle number density and $\\varepsilon$ is the energy density of the universe. \\emph{This enables us to show that a small negative internal pressure fluctuation may have a favorable effect on the growth rate of a density perturbation}. Secondly, we have found that for scalar perturbations, the perturbed momentum constraint equations (\\ref{basis-2}) can be rewritten as an evolution equation (\\ref{FRW6}) for the local perturbation $\\mbox{$^3\\!R_{\\een\\parallel}$}$ (\\ref{driekrom}) to the global spatial curvature $\\mbox{$^3\\!R_\\nul$}$ (\\ref{eq:glob-curve}). Using this new insight we can rewrite the evolution equations for scalar perturbations as simple extensions (\\ref{subeq:pertub-flrw}) of the background equations (\\ref{subeq:einstein-flrw}). The sets of equations (\\ref{subeq:einstein-flrw}) and (\\ref{subeq:pertub-flrw}) are written with respect to the same system of reference. From these two sets of equations it follows that only three independent scalars, namely the energy density $\\varepsilon$, the particle number density $n$, and the expansion scalar $\\theta$, play a role in a perturbation theory. This reduces the number of possible gauge-invariant variables considerably. Finally, in a first-order perturbation theory the gravitational field is already weak. Therefore, we only need to take the limit that all particle velocities are negligible with respect to the speed of light, $v/c\\rightarrow0$, the so-called \\emph{non-relativistic limit}, see (\\ref{eq:non-rel-lim}), in order to show that our treatment reduces to the Newtonian theory of gravity. In fact, we show that there exist two, and only two, gauge-invariant variables defined by (\\ref{e-n-gi}) with the property that in the non-relativistic limit the well-known Newtonian results (\\ref{eq:poisson}) and (\\ref{eq:newt-ngi}) show up. Consequently, $\\varepsilon^\\gi_\\een$ and $n^\\gi_\\een$ are the real, physical, energy density perturbation and the particle number density perturbation. ", "conclusions": "" }, "1003/1003.4194_arXiv.txt": { "abstract": "We use Total Solar Irradiance (TSI) measurements from the VIRGO (Variability of solar IRradiance and Gravity Oscillations) instrument on board SOHO to obtain preliminary estimates of the mean total radiative energy emitted by X-class solar flares. The basic tool is that of summed-epoch analysis, which has also enabled us to detect and partially characterize systematic errors present in the basic data. We describe these errors, which significantly degrade the photometry at high frequencies. We find the ratio of GOES 1-8\\AA~luminosity to total bolometric luminosity to be of order~0.01. ", "introduction": "The detection of solar flares against the glare of the photosphere has only previously been possible via imaging, whereas many kinds of \\textit{stellar} flares can be readily detected via broad-band photometry of integrated starlight. This is often because of a cooler, fainter photospheric background, as in the classical dMe flare stars. The complete energy distribution of a solar flare across many wavebands (especially, for example the vacuum UV) has also remained poorly understood \\citep[e.g.][]{emslie}. Another problem for characterizing the bolometric luminosity of a solar flare lies in the photospheric fluctuations, mainly a broad continuum of power due to granulation \\citep[e.g.][]{hudson} and having an rms amplitude of roughly 50~ppm in the 2-8~mHz band. The Total Solar Irradiance (TSI) instruments flown in space since the late 1970s typically sample this band. Flare light curves at various wavelengths suggest that this band contains much of the flare power, although variations well above 1~Hz do occur \\citep[e.g.][]{dennis}. Only recently has it become possible to detect solar flares bolometrically \\citep{woods}, and only very convincingly for a single event, the X17~flare of 2003 October~28. The solar background noise level corresponds to 50~$\\mu$mag in astronomical terms, and it is remarkable that this fluctuation obscures most solar flares, whereas stellar flares often have increases exceeding one full magnitude. Kretzschmar (2008) \\nocite{kretzsch} has introduced the technique of superposed epoch analysis to detect multiple flares in the mean, and we adopt this scheme for the analysis presented here. We also study TSI observations from the VIRGO instrument on board SOHO, but the same techniques should work well for other instruments. Our analysis has revealed sources of unwanted variance in the TSI signal, which we intend to correct in a future analysis. ", "conclusions": "We have reported initial results of summed-epoch analysis for the total irradiance of solar flares. This analysis suggests a strong correlation with the impulsive phase of a flare, identified here by the time derivative of the soft X-ray flux. We use the analysis to obtain an estimate of $L_X/L_{bol}$, which can then be compared with similar estimates for different kinds of stellar flares, for which parameters such as gravity, rotation, or composition can be studied. Our analysis, that of Kretzschmar (2008), and the synthesis of Emslie et al. (2005) all agree that this ratio is of order~0.01, using the GOES 1-8\\AA~passband. We have also found a source of unwanted noise in the observations, which we plan to correct in a future more sensitive analysis of these data. \\bigskip\\noindent {\\bf Acknowledgements:} SOHO is a project of international cooperation between ESA and NASA. HSH thanks NASA for support under grant NNX07AH74G." }, "1003/1003.3257_arXiv.txt": { "abstract": "The GeV and TeV emission from M82 and NGC 253 observed by Fermi, HESS, and VERITAS constrains the physics of cosmic rays (CRs) in these dense starbursts. We argue that the $\\gamma$-rays are predominantly hadronic in origin, as expected by previous studies. The measured fluxes imply that pionic losses are efficient for CR protons in both galaxies: we show that a fraction $F_{\\rm cal} \\approx 0.2 - 0.4$ of the energy injected in high energy primary CR protons is lost to inelastic proton-proton collisions (pion production) before escape, producing $\\gamma$-rays, neutrinos, and secondary electrons and positrons. We discuss the factor $\\sim 2$ uncertainties in this estimate, including supernova rate and leptonic contributions to the GeV-TeV emission. We argue that $\\gamma$-ray data on ULIRGs like Arp 220 can test whether M82 and NGC 253 are truly calorimetric, and we present upper limits on Arp 220 from the \\emph{Fermi} data. We show that the observed ratio of the GeV to GHz fluxes of the starbursts suggests that non-synchrotron cooling processes are important for cooling the CR electron/positron population. We briefly reconsider previous predictions in light of the $\\gamma$-ray detections, including the starburst contribution to the $\\gamma$-ray background and CR energy densities. Finally, as a guide for future studies, we list the brightest star-forming galaxies on the sky and present updated predictions for their $\\gamma$-ray and neutrino fluxes. ", "introduction": "M82 and NGC 253 are nearby ($D \\approx 2.5 - 4.0\\ \\Mpc$), prototypical starburst galaxies, each having an intense star-forming region of about $200\\ \\pc$ radius in the center of a more quiescent disk galaxy. The starbursts are expected to have high supernova (SN) rates of about $0.03 - 0.3\\ \\yr^{-1}$. SN remnants are believed to accelerate primary cosmic ray (CRs) protons and electrons. The high SN rates in starbursts imply high CR emissivities. The presence of CR electrons and positrons in these starbursts is inferred from the nonthermal synchrotron radio emission they produce \\citep[e.g.,][]{Klein88,Volk89}. However, most of the CR energy is believed to be in the form of CR protons. When high energy CR protons collide with interstellar medium (ISM) nucleons, they create pions, which decay into secondary electrons and positrons, $\\gamma$-rays, and neutrinos. With their dense ISMs ($\\mean{n} \\approx 100 - 500\\ \\cm^{-3}$) and possible high CR energy densities \\citep[as evinced by the bright radio emission][]{Volk89,Akyuz91,Persic10}, M82 and NGC 253 are predicted to be bright $\\gamma$-ray sources (e.g., \\citealt{Akyuz91,Sreekumar94,Volk96,Paglione96,Romero03,Domingo05,Thompson07} [TQW]; \\citealt{Persic08,deCeaDelPozo09,Rephaeli09,Lacki10} [LTQ]). As prototypical starbursts, if M82 and NGC 253 are seen in $\\gamma$-rays, starbursts in general may be sources of $\\gamma$-rays \\citep{Pohl94,Torres04a}, with important implications for the diffuse $\\gamma$-ray and neutrino backgrounds (e.g., \\citealt{Pavlidou02}; \\citealt{Loeb06} [LW06]; TQW). However, the $\\gamma$-ray luminosity of starbursts depends not only on the injection rate of CRs, but also on the efficiency of converting CR proton energy into pionic $\\gamma$-rays, neutrinos, and secondary electrons and positrons. In turn, this efficiency depends on the ratio of the timescale for pion production to the escape timescale. The hypothesis that CR protons in starbursts lose all of their energy to pionic collisions before escaping is called ``proton calorimetry'' \\citep[c.f.][]{Pohl94}.\\footnote{Here, we consider only CR protons with kinetic energy above the threshold for pion production.} If proton calorimetry is strongly violated, then M82 and NGC 253 and, by extension, other starbursts could in fact be weak $\\gamma$-ray sources. Although $\\gamma$-ray emission from M82 and NGC 253 has been sought for several years with no success (at GeV, \\citealt{Cillis05}; and at TeV, \\citealt{Aharonian05,Itoh07}), the launch of \\emph{Fermi} and the advent of powerful VHE $\\gamma$-ray telescopes has led to recent detections of both starbursts at GeV energies (with \\emph{Fermi}; \\citealt{Abdo10a}) and in VHE $\\gamma$-rays (M82 with VERITAS, \\citealt{Acciari09}; NGC 253 with HESS, \\citealt{Acero09}). These GeV and TeV detections constrain the cosmic ray (CR) population in these dense star-forming environments. In this paper, we discuss the implications of the $\\gamma$-ray detections of M82 and NGC 253. The ratio of the $\\gamma$-ray luminosities to the bolometric luminosities informs the question of whether or not these systems are proton calorimeters (TQW). The emission also has implications for the energy density of CRs in starbursts \\citep[e.g.,][]{Akyuz91}. Finally, combined with the radio emission, the energy losses of CR electrons and positrons are constrained \\citep[c.f.,][]{Paglione96,Domingo05,Persic08,deCeaDelPozo09,Rephaeli09}. Pionic $\\gamma$-rays must be accompanied by secondary positrons and electrons; the ratio of the power in these expected electrons and positrons to the observed radio emission informs us of the energy losses of the CR electrons and positrons. In particular, we derive in \\S~\\ref{sec:FRCImplications} the expected synchrotron luminosity from the pionic luminosity if synchrotron cooling is the dominant loss process. In \\S\\ref{sec:Detections} we describe the detections of M82 and NGC 253 at GeV and TeV energies. We then interpret the detections as $\\gamma$-rays from diffuse CR protons in \\S\\ref{sec:Interpretation}. Our interpretation includes comparison of the $\\gamma$-ray luminosities of M82 and NGC 253 with their CR luminosities and their IR luminosities (\\S\\ref{sec:CalorimetryFraction}), and a discussion of the uncertainties in these estimates (\\S\\ref{sec:Uncertainties}). We find that a fraction 0.4 and 0.2 of luminosity in $\\ge \\GeV$ CR protons is lost to pion production in M82 and NGC 253, respectively. We discuss the implications for our estimates mean for proton calorimetry in M82 and NGC 253 are at GeV energies (\\S\\ref{sec:Calorimetry}). Other possible sources for the observed $\\gamma$-rays are considered in \\S\\ref{sec:Sources}. The implications of the detections of M82 and NGC 253 for the detection of other star-forming galaxies, the starburst contribution to the diffuse extragalactic $\\gamma$-ray and neutrino backgrounds, the dynamical importance of CRs in starbursts, and for the physics of the FIR-radio correlation are described in \\S\\ref{sec:Implications}. We summarize our results in \\S\\ref{sec:Conclusion}. ", "conclusions": "\\label{sec:Conclusion} M82 and NGC 253 have now been detected in GeV and TeV $\\gamma$-rays, with fluxes roughly comparable to previous detailed predictions. We have shown that the observed $\\gamma$-ray fluxes imply that a fraction $F_{\\rm cal} \\approx 0.2 - 0.4$ of the energy injected into high energy CR protons is lost to inelastic collisions (pion production) with protons in the ISM (for $\\eta^{\\prime} = 0.1$). However, $F_{\\rm cal}$ in the range of 0.1 - 1 can be accommodated with different SNe rates and acceleration efficiencies (see the uncertainties in \\S~\\ref{sec:Uncertainties} and \\S\\ref{sec:Sources}). We find a significantly higher $F_{\\rm cal}$ for NGC 253 than \\citet{Acero09} because NGC 253 has more GeV emission than they expected. The uncertainty in $F_{\\rm cal}$ will decrease significantly with more observations by {\\it Fermi}, HESS, and VERITAS. A future test of proton calorimetry in M82 and NGC 253 would be a $\\gamma$-ray detection of a ULIRG like Arp 220 \\citep[c.f.][]{Torres04}. Arp 220 is more likely to be a proton calorimeter than M82 and NGC 253, with its extremely high average gas density. If M82 and NGC 253 are not proton calorimeters but Arp 220 is, the ratio of Arp 220's pionic luminosity to its stellar luminosity will be greater than M82 and NGC 253 -- it will be brighter in $\\gamma$-rays than expected (see Figure~\\ref{fig:GeVtoTIR}, Tables \\ref{table:flux_normal} \\& \\ref{table:flux_starburst}). Unfortunately, Arp 220's flux is expected to be challenging to detect with \\emph{Fermi}, although upper limits alone may be constraining (as in the Appendix). Stacking searches of ULIRGs may also prove useful. Pionic $\\gamma$-ray emission implies secondary $e^{\\pm}$ production in these starbursts \\citep[c.f.][]{Rengarajan05}; from the GHz to GeV ratio, we found evidence of significant non-synchrotron losses. This suggests that bremsstrahlung and ionization are important energy loss mechanisms for CR electrons and positrons \\citep[c.f.][]{Murphy09}. This would support the idea presented in \\citet{Thompson06} that these losses flatten the GHz radio spectrum of starbursts (\\S~\\ref{sec:FRCImplications}). It would also support the ``high-$\\Sigma_g$ conspiracy'' suggested by LTQ to explain the linearity of the FIR-radio correlation for starbursts, whereby bremsstrahlung, ionization, and IC losses suppress the synchrotron radio emission of CR electrons in starbursts, but proton calorimetry leads to secondary electrons and positrons that boost the radio emission. Whatever the underlying physics of $\\gamma$-ray production in M82 and NGC 253 is, the high fluxes of these starbursts suggest that other starbursts should also be $\\gamma$-ray sources. We compile our predictions in Tables \\ref{table:flux_normal} \\& \\ref{table:flux_starburst}. Considering that much of the star formation in the universe at high-$z$ is in luminous infrared galaxies \\citep[e.g.,][]{Elbaz99,Chary01,PerezGonzalez05,Magnelli09}, starbursts might make up a significant fraction ($\\sim 1/2$) of the entire $\\gamma$-ray background (e.g., \\citealt{Pavlidou02}, TQW, \\citealt{Bhattacharya09}; \\S\\ref{sec:Backgrounds}). If the hadronic interpretation of the $\\gamma$-ray flux holds, the neutrino background should also be large (LW06). Finally, the conclusion that M82 has $F_{\\rm cal} \\approx 0.4$ and NGC 253 has $F_{\\rm cal} \\approx 0.2$ implies that the pion cooling timescale is nearly equal to the wind escape timescale, $\\sim 2 \\times 10^5$\\,yr for these systems. This, in turn, suggests that the CR protons on average interact with ISM near the mean density. If this is correct, then the CR pressure is significantly below the pressure needed to support each starburst gravitationally, and CRs are not on average dynamically important deep within the starbursts (\\S~\\ref{sec:Feedback})." }, "1003/1003.1252_arXiv.txt": { "abstract": "{} {We investigate Galactic bulge planetary nebulae without emission-line central stars for which peculiar infrared spectra have been obtained with the Spitzer Space Telescope, including the simultaneous signs of oxygen and carbon based dust. Three separate sub-groups can be defined characterized by the different chemical composition of the dust and the presence of crystalline and amorphous silicates.} {We use literature data to analyze the different nebular properties and deduce both the evolutionary status and the origin of these three groups. In particular, we check whether there are signs of evolutionary links between dual-dust chemistry planetary nebulae without detected emission-line central stars and those with emission-line stars.} {Our primary finding is that the classification based on the dust properties is reflected in the more general properties of these planetary nebulae. However, some observed properties are difficult to relate to the common view of planetary nebulae. In particular, it is challenging to interpret the peculiar gas chemical composition of many analyzed objects in the standard picture of the evolution of planetary nebulae progenitors. We confirm that the dual-dust chemistry phenomenon is not limited to planetary nebulae with emission-line central stars.} {Our results clearly indicate that there is no unique road to the formation of planetary nebulae even in a homogeneous environment such as the Galactic bulge. The evolution of a single asymptotic giant branch star may lead to the formation of different types of planetary nebulae. In addition, the evolution in a close binary system should sometimes also be considered.} ", "introduction": "After the completion of hydrogen and helium burning in their cores, low- to intermediate-mass stars (0.8 $\\leq$ M $\\leq$ 8 \\Msun) evolve towards the asymptotic giant branch \\citep[AGB; e.g.,][]{Herwig2005} and then pass through the planetary nebula (PN, plural PNe) phase before ending their lives as white dwarfs. At the tip of the AGB phase, these stars experience strong mass loss that efficiently enriches the surrounding interstellar medium with huge amounts of gas and dust. Stars leave the AGB when the strong mass loss stops and then the future central star (CS, plural CSs) rapidly evolves towards hotter effective temperatures in the Hertzsprung-Russell diagram. Thus, when the ionization of the ejected gas takes place, a new PN is formed. However, in most cases the total amount of ionized gas is very small compared to the total mass previously ejected. An important fraction of this material remains neutral in the form of dust grains, molecules, or atoms, which can be easily detected in the infrared domain. Thanks to the analysis of the features observed by the Infrared Space Observatory (ISO) in the spectra of PNe, it was possible to confirm the presence of large amounts of dust grains around PNe as well as the dominant dust chemistry (C-rich versus O-rich). Features at 3.3, 6.2, \"7.7\", 8.6, and 11.3\\um\\ attributed to polycyclic aromatic hydrocarbons (PAHs) are common in C-rich PNe \\citep[e.g.,][and references therein]{Garcia-Lario1999} while strong features attributed to crystalline silicates (e.g., those centered on 23.5, 27.5 and 33.8 \\um) are usually found in O-rich PNe \\citep[e.g.,][]{Molster2002}. A few Galactic disk PNe exhibited a remarkable dual-dust (C-rich and O-rich) chemistry showing both PAH and crystalline silicate features in ISO spectra \\citep{Waters1998a, Waters1998b, Cohen1999, Cohen2002}. The fact that this was an infrequent phenomenon may be due to the instruments used, which in many cases may have been unable to detect crystalline silicates. For example, the Spitzer Space Telescope \\citep[Spitzer,][]{Werner2004}, detected crystalline silicates in 10 post-AGB sources \\citep{Cerrigone2009} while after completion of the ISO mission only 2 such sources were known \\citep{Szczerba2003}. The mixed-chemistry PNe discovered by ISO pertain to the class of objects with C-rich Wolf-Rayet type nuclei (the so-called [WR] PNe), which usually show a lack of hydrogen in their atmospheres. These atmospheres are instead mostly composed of helium, carbon, and oxygen and the CSs show intense mass-loss \\citep[e.g.,][]{Crowther2008}. The evolution of an AGB star with a stellar or substellar companion that undergoes the common-envelope phase is another possible way of creating a PN. Some authors argue that a companion object is often mandatory for a planetary nebula to be created \\citep[see in][and references therein]{DeMarco2009}. We noted that some hypotheses compiled to explain the simultaneous presence of carbon and oxygen dust also require a binary system. The Galactic bulge is characterized by an old population of mostly low-mass stars (\\citet{Zoccali2003}, but see also \\citet{Uttenthaler2007} and references therein). It is also known that the abundances of PNe in the Galactic bulge (GBPNe) differ from those located in the Galactic disk as they have higher metallicities and lower C/O ratios \\citep[e.g.,][]{WangLiu2007}. The differences in metallicity seems to play a dominant role in the chemical evolution of low- to intermediate-mass stars (e.g., \\citet{Garcia-Hernandez2007}; \\citet{Stanghellini2007}; \\citet{Chiappini2009}). Thus, studding the Galactic bulge enables us to investigate the stellar evolution of low- and intermediate-mass stars in higher metallicity environments and at the same time an insight into the chemical evolution and formation of our Milky Way. \\citet{Gutenkunst2008} analyzed Spitzer spectra acquired using the Infrared Spectrograph \\citep[IRS,][]{Houck2004} of 11 PNe towards the Galactic bulge and inferred dual-dust chemistry in 6 of them. They suggested that the high percentage of dual-dust chemistry sources is unsurprising because the fraction of [WR] PNe is significantly higher in the bulge than in the Galactic disk. However, as we checked, only one of their dual-dust chemistry sources have the wind characteristics of [WR] type CS and the higher proportion of genuine [WR] PNe in the Galactic bulge is not confirmed \\citep{Gorny2009}. \\citet{P-C2009} found that dual-dust chemistry is truly widespread among GBPNe. They analyzed a larger sample of 26 GBPNe observed with Spitzer/IRS among which 21 exhibit dual-dust chemistry. \\citet{P-C2009} observations shown that the simultaneous presence of oxygen and carbon-rich dust features in the infrared spectra of [WR] PNe is not restricted to objects with late/cool [WC] class stars. In addition, dual-dust chemistry was seen in all observed PNe with WEL stars ({\\em weak emission-line} stars, \\citealp{Tylenda1993}) as well as members of recently discovered VL group (low ionization PNe around stars with {\\em very late} [WC\\,11]-like spectra, \\citealp{Gorny2009}). Surprisingly, \\citet{P-C2009} found dual-dust chemistry also in some PNe without detected emission-line CSs. Another interesting property of the PNe observed by \\cite{P-C2009} was the amorphous silicate emission at 10\\,$\\mu$m, which was detected in four dual-dust chemistry GBPNe and in most of the O-rich PNe that they observed. Note that before Spitzer there was known only one such PN, namely SwSt\\,1 \\citep[e.g.,][]{Szczerba2001}, belonging to the [WR]~PNe. In contrast, the 10\\um\\ feature objects found by \\cite{P-C2009}, neither belong to this group nor exhibit stellar emission lines. In this work, we analyze the multiple properties of PNe without emission-line CSs that are found to have peculiar infrared spectra acquired with Spitzer/IRS. The paper is organized as follows. In Sect.~2 and 3 we describe our working sample and the main properties of their infrared spectra, respectively. We analyze the nebular properties and evolutionary status of these PNe in Sect.~4. Finally, in Sect.~5 we discuss the results obtained. Our concluding remarks are given in Sect.~6. \\begin{table} \\caption{ List of analyzed PNe. The Galactic bulge (b) or disk (d) association of the object is marked in Col.~3. The reference to original observers is given in Col.~4. Objects are divided into three groups according to their infrared properties as marked in Cols.~5--7 (see description in Sect.~3).} \\begin{tabular}{ l @{\\hspace{0.40cm}} l @{\\hspace{0.40cm}} c @{\\hspace{0.40cm}} c @{\\hspace{0.60cm}} c @{\\hspace{0.20cm}} c @{\\hspace{0.20cm}} c} \\hline \\hline PN G & name & pop. & Ref. & \\multicolumn{3}{c}{IR spectra} \\\\ & & & & PAH & SiO cr. & SiO am. \\\\ \\hline \\DCcr: & & & & & & \\\\ 000.1+04.3 & H 1-16 & b & P & + & + & - \\\\ 007.2+01.8 & Hb 6 & d & P & + & + & - \\\\ 354.2+04.3 & M 2-10 & b & G & + & + & - \\\\ 358.7-05.2 & H 1-50 & b & P & + & + & - \\\\ \\vspace{0.15cm} 358.9+03.2 & H 1-20 & b & G & + & + & - \\\\ \\DCacr: & & & & & & \\\\ 354.5+03.3 & Th 3-4 & b & P & + & + & + \\\\ 356.9+04.4 & M 3-38 & b & P & + & + & + \\\\ 358.2+04.2 & M 3-8 & b & P & + & + & + \\\\ \\vspace{0.15cm} 359.7-02.6 & H 1-40 & b & P & + & + & + \\\\ \\OCacr: & & & & & & \\\\ 002.2-02.7 & M 2-23 & b & P & - & + & + \\\\ 008.2+06.8 & He 2-260 & d & P & - & + & + \\\\ 355.6-02.7 & H 1-32 & b & P & - & + & + \\\\ 355.7-03.5 & H 1-35 & b & P & - & + & + \\\\ \\hline \\hline \\end{tabular} Ref: G -- Gutenkunst et~al. (2008); P -- Perea-Calder{\\'o}n et~al. (2009) \\end{table} ", "conclusions": "We have investigated PNe without emission-line central stars located towards the Galactic bulge that have peculiar infrared spectra acquired by Spitzer/IRS. Among these objects, we have found three separate groups divided according to their composition of dust grains: \\begin{itemize} \\item{\\DCcr\\ -- dual-dust chemistry PNe with simultaneous presence of both carbon-based dust (PAHs) and oxygen-based dust (crystalline silicates);} \\item{\\DCacr\\ -- dual-dust chemistry PNe with simultaneous existence of PAHs and crystalline silicates as well as amorphous silicates;} \\item{\\OCacr\\ -- PNe characterized by oxygen dust chemistry with only oxygen-rich grains in both crystalline and amorphous forms.} \\end{itemize} \\noindent We have analyzed a wide range of different properties of these PNe. Our main results are: \\begin{itemize} \\item{We confirm that dual-dust chemistry is a common phenomenon of PNe in the Galactic bulge and can occur in objects not related to emission-line central stars.} \\item{The Properties of \\DCcr\\ PNe do not distinguish them clearly from the majority of other PNe in the Galactic bulge. They have intermediate or higher-mass central stars. Their infrared spectra closely resemble those of PNe with emission-line nuclei. Some \\DCcr\\ may be evolutionary related to the latter objects or may have undiscovered emission-line central stars.} \\item{\\DCacr\\ objects belong to the densest PNe in the Galactic bulge. Their derived {\\em{m}}$_d$/{\\em{m}}$_g$ mass ratios and infrared excesses IRE are higher than average. There is a possibility of extensive internal extinction. \\DCacr\\ PNe have intermediate and higher-mass central stars. The chemical composition of nebular gas is peculiar as oxygen seems underabundant relative to hydrogen and nitrogen but not to other elements (except possibly argon). This composition of \\DCacr\\ PNe cannot be explained in the standard picture of AGB star chemical evolution.} \\item{The \\OCacr\\ PNe are the only analyzed PNe not showing dual-dust chemistry. They have hottest dust temperature T$_d$ and highest infrared excess IRE. \\OCacr\\ have also very small diameters and are among the densest PNe in the Galactic bulge. However, their expansion velocities are smaller than average and therefore their evolutionary status indicates that \\OCacr\\ can have intermediate-mass central stars. In contrast, the surrounding nebulae show low metallicity with an underabundance of nitrogen, neon, and argon. The domination of oxygen-based dust indicates in addition a low abundance of carbon. We argue that their properties are in qualitative agreement with scenarios of PNe formation not from single AGB stars but from binary systems.} \\end{itemize}" }, "1003/1003.3243_arXiv.txt": { "abstract": "{We present the GalMer database, a library of galaxy merger simulations, made available to users through tools compatible with the Virtual Observatory (VO) standards adapted specially for this theoretical database. To investigate the physics of galaxy formation through hierarchical merging, it is necessary to simulate galaxy interactions varying a large number of parameters: morphological types, mass ratios, orbital configurations, etc. On one side, these simulations have to be run in a cosmological context, able to provide a large number of galaxy pairs, with boundary conditions given by the large-scale simulations, on the other side the resolution has to be high enough at galaxy scales, to provide realistic physics. The GalMer database is a library of thousands simulations of galaxy mergers at moderate spatial resolution and it is a compromise between the diversity of initial conditions and the details of underlying physics. We provide all coordinates and data of simulated particles in FITS binary tables. The main advantages of the database are VO access interfaces and value-added services which allow users to compare the results of the simulations directly to observations: stellar population modelling, dust extinction, spectra, images, visualisation using dedicated VO tools. The GalMer value-added services can be used as virtual telescope producing broadband images, 1D spectra, 3D spectral datacubes, thus making our database oriented towards the usage by observers. We present several examples of the GalMer database scientific usage obtained from the analysis of simulations and modelling their stellar population properties, including: (1) studies of the star formation efficiency in interactions; (2) creation of old counter-rotating components; (3) reshaping metallicity profiles in elliptical galaxies; (4) orbital to internal angular momentum transfer; (5) reproducing observed colour bimodality of galaxies.} ", "introduction": "In the framework of the present cosmological paradigm, mergers and interactions are among the most important mechanisms governing galaxy formation and evolution. \\citet{SB51} proposed that collisions of late-type disc galaxies should produce early-type ones. A natural consequence of this phenomenon is the morphology--density relation discovered three decades later \\citep{Dressler80}: dense regions of the Universe, where galaxy collisions are supposed to be more frequent, contain larger fractions of early-type galaxies than sparsely populated areas. \\citet{TT72} were the first to suggest that interactions `tend to bring \\emph{deep} into a galaxy a fairly \\emph{sudden} supply of fresh fuel in the form of interstellar material$\\dots$' The gas in the bar formed during the interaction loses its angular momentum due to the torques \\citep{CDG90,BH96}, falls onto the galaxy centre possibly inducing strong bursts of star formation \\citep{MH94a,MH96} and creating young compact central stellar component \\citep{MH94b} often observed in present-day early-type galaxies \\citep{Silchenko06,Kuntschner+06}. Intense star formation episodes accompanied by supernova explosions, enrich the interstellar medium (ISM) with heavy elements, consumed into the stars hence increasing the observed metal abundances of the stellar population. Large scale cosmological $N$-body simulations \\citep[e.g.][]{SFW06,OPT08} often lack the spatial resolution to trace in detail star formation and morphological transformation at galaxy scale. Therefore, usually they are complemented by semi-analytical prescriptions qualitatively accounting for phenomena strongly affecting galaxy evolution, such as star formation \\citep[e.g.][]{Blaizot+04,Somerville+08}. However, the parameters of the semi-analytical models have to be chosen based on more detailed simulations of galaxy interactions. High resolution galaxy simulations \\citep[e.g.][]{BDE08} cannot be performed in large statistical numbers. A compromise has then to be done between statistics and resolution. This becomes one of the main motivations for studying large numbers of galaxy interactions by means of dedicated intermediate-resolution numerical simulations. Merger-induced star formation, as well as morphological transformation, strongly depends on the mass ratio of the interacting galaxies. Generally, the intensity of starburst decreases as the merger mass ratio increases \\citep[e.g.][]{Cox+08}. Equal mass mergers of disc galaxies (mass ratios below 4:1) usually result in early-type elliptical-like remnants \\citep{Toomre77,NB03} while minor mergers (ratios above 10:1) do not destroy the progenitor's disc preserving its exponential mass distribution although making it thicker and dynamically hotter \\citep{QHF93,WMH96,VW99,BJC05}. A sequence of repeated minor mergers can form elliptical galaxies, with global morphological and kinematical properties similar to that observed in real ellipticals \\citep{BJC07}. Orbital parameters of the interaction and orientation of galaxies also strongly affect the process of merger, e.g. star formation efficiency on retrograde orbits is generally higher than for direct encounters \\citep{dMCMS07}. One needs to explore a large multi-dimensional parameter space (different initial morphologies related to gas content, orbital configurations, mass ratios, etc.) by running thousands of simulations in order to fully understand the astrophysical consequences of galaxy interactions for the modern picture of galaxy evolution. The GalMer project, developed in the framework of the French national HORIZON\\footnote{http://www.projet-horizon.fr/} collaboration, has the ambitious goal of providing access for the astronomical community to the results of massive intermediate-resolution numerical simulations of galaxy interactions in pairs, covering as much as possible the parameter space of the initial conditions and, thus allowing to study statistically the star formation enhancements, structural and dynamical properties of merger remnants. An important aspect is to integrate the data services into the framework of the International Virtual Observatory in order to take the full advantage of already developed technologies and data visualisation and processing tools. This paper is a technical presentation of: (1) the GalMer TreeSPH simulations providing all essential details on initial conditions for galaxies of different masses and prescriptions used to model the processes of star formation including supernova feedback and metallicity evolution; (2) the GalMer database, the first VO resource containing results of TreeSPH simulations; (3) the GalMer value-added services aimed at facilitating the comparison between simulations and observations such as modelling the spectrophotometric galaxy properties using an evolutionary synthesis code. The paper is organised as follows: in Section~2 we describe initial conditions of the numerical simulations and orbital parameters of galaxy interactions; Section~3 contains information on the numerical method, prescriptions for star formation and metallicity evolution; Section~4 provides the description of the GalMer database structure, its access interface, and mechanisms of data visualisation; Section~5 presents value-added services of the GalMer database; in Section~6 we define some astrophysical applications which can be tackled with our simulations; Section~7 contains a brief summary. ", "conclusions": "We present the GalMer database providing online access to the results of TreeSPH simulations. The structure of the database conforms to recent International Virtual Observatory standards. We describe the interactive data access web-interface, advanced mechanisms for data visualisation and manipulation using connection to the Virtual Observatory tools dedicated for dealing with tabular, imaging and spectral data. After having stored the snapshots in FITS binary table format, giving access to all coordinates and data on individual simulation particles, we provide a set of value-added tools using the results of TreeSPH simulations. These include: (1) generator of projected maps of various physical quantities traced by the simulations; (2) engine to perform the modelling of spectrophotometric properties of interacting and merging galaxies based on the \\soft{PEGASE.2/PEGASE.HR} evolutionary synthesis code. The latter tool can be used as a virtual telescope and synthetic images, spectra, and datacubes generated using it are directly comparable to observational data. We provide examples of several use-cases for the GalMer database using our value-added tools. The database will be updated by including new simulations as they have been completed." }, "1003/1003.3419_arXiv.txt": { "abstract": "{% A considerable fraction of the central stars of planetary nebul\\ae\\ (CSPNe) are hydrogen-deficient. As a rule, these CSPNe exhibit a chemical composition of helium, carbon, and oxygen with the majority showing Wolf-Rayet-like emission line spectra. These stars are classified as CSPNe of a spectral type [WC]. We perform a spectral analysis of CSPN PB\\,8 with the Potsdam Wolf-Rayet (PoWR) models for expanding atmospheres. The source PB\\,8 displays wind-broadened emission lines from strong mass loss. Most strikingly, we find that its surface composition is hydrogen-deficient, but not carbon-rich. With mass fractions of 55\\% helium, 40\\% hydrogen, 1.3\\% carbon, 2\\% nitrogen, and 1.3\\% oxygen, it differs greatly from the 30--50\\% of carbon which are typically seen in [WC]-type central stars. The atmospheric mixture in PB\\,8 has an analogy in the WN/WC transition type among the massive Wolf-Rayet stars. Therefore we suggest to introduce a new spectral type [WN/WC] for CSPNe, with PB\\,8 as its first member. The central star of PB\\,8 has a relatively low temperature of $T_*=52\\,\\text{kK}$, as expected for central stars in their early evolutionary stages. Its surrounding nebula is less than 3000 years old, i.e.\\ relatively young. Existing calculations for the post-AGB evolution can produce hydrogen-deficient stars of the [WC] type, but do not predict the composition found in PB\\,8. We discuss various scenarios that might explain the origin of this unique object. } ", "introduction": "% A planetary nebula (PN) surrounds a central star which is hot enough ($T_\\ast > 25\\,000\\,\\text{K}$) to ionize its circumstellar matter. According to the well-established scenario {\\citep[e.g.][]{paczynski1970,schoenberner1989}}, the central star of the planetary nebula (CSPN) ejects its nebula while suffering thermal pulses at the tip of the asymptotic giant branch (AGB). In the subsequent PN phase it evolves rapidly towards the white dwarf cooling sequence. Most of the CSPNe show a hydrogen-rich surface composition. Among the Galactic central stars, 5-6\\,\\% are hydrogen-deficient and show emission lines in their spectra \\citep{tylack1993,acknein2003}. Moreover, half of these hydrogen-deficient stars have spectra similar to those of massive Wolf-Rayet stars of the carbon sequence and are therefore classified as spectral type [WC]. They have a strong stellar wind composed of helium, carbon, and oxygen. Typical carbon surface-abundances have been found to lie between 30\\% and 50\\% by mass \\citep[see the reviews by][]{koesterke2001,crow2008}. The central star (CS) of the planetary nebula PB\\,8 (PN\\,G292.4+04.1) was first classified by \\cite{mendez1991} as a hydrogen-rich Of-WR(H) star due to the H$\\gamma$ P Cygni profile and the appearance of an unusually strong He{\\,\\sc ii} 4686 emission line. In contrast, \\cite{acknein2003} classified this star as a [WC5-6] type star. Below we analyze optical, IUE, and FUSE spectra of the central star PB\\,8 by means of the Potsdam Wolf-Rayet (PoWR) model atmosphere code. The observations are introduced in Sect.\\,\\ref{sect:pb8-obs}. Spectral modeling is briefly explained in Sect.~\\ref{sect:pb8-methods}. In Sect.~\\ref{sect:pb8-analysis} we describe the spectral analysis, and the results are discussed in the final section (Sect.\\,\\ref{sect:pb8-discussion}). ", "conclusions": "\\label{sect:pb8-discussion} \\subsection{PN and Central star status} \\label{sect:pb8-PNstatus} The nebula PB\\,8 appears as a roughly spherical nebula, nearly round in the composite image of H$\\alpha$, [N{\\,\\sc ii}], and [O{\\,\\sc iii}] (Fig.~\\ref{fig:pb8-linear}, left panel \\footnote{from {\\tt http://www.astro.washington.edu/balick/PNIC/}}), although the shell shows some knotty structure. In particular, there is a bright structure extending from the center to the northern side of the shell. The long-slit spectrogram also reveals a good symmetry in the radial velocities (Fig.~\\ref{fig:pb8-linear}, right panel). Given the unique chemical abundances of the central star PB\\,8, one must consider the possibility that this object is actually a massive star with a ring nebula. However, the low nebular expansion velocity discussed in Sect.~\\ref{sect:pb8-nebula} is rather characteristic for PNe. \\cite{medina2006} found expansion velocities for PNe with Wolf-Rayet nuclei in the range of $8-44\\,\\text{km}\\,\\text{s}^{-1}$ from direct observations. Expansion velocities for ring nebul\\ae\\ around massive stars are systematically higher, $16-110\\,\\text{km}\\,\\text{s}^{-1}$ \\citep{chu1999}. Moreover, the electron density in PB\\,8 measured by \\cite{gar2009}, ${n_\\text{e} = 2550\\pm 550\\,\\text{cm}^{-3}}$ is typical for young planetary nebul\\ae, but several times higher than found in ring nebul\\ae\\ \\citep{mathis1992}. Furthermore, if the central star of PB\\,8 were a massive star, this would imply a luminosity of at least $\\log (L/L_\\odot) = 5.3$, which shifts the distance to {$\\approx 24.2 \\,\\rm kpc$. With a Galactic latitude of $4^\\circ$ this corresponds to a height of $1.7{\\,\\rm kpc}$} above the fundamental plane of the Galaxy. This is much more than the scale height of the thin disk and therefore an unlikely location for a massive star. \\subsection{Re-classification of the central star of PB\\,8} The central star of PB\\,8 has been classified as spectral type [WC5-6] by \\cite{acknein2003}. Yet we showed above that the central star of PB\\,8 is not a member of the [WC] sequence; its spectrum shows strong lines of nitrogen, reflecting that its chemical composition rather resembles that of a WN star. Nevertheless, carbon is slightly enhanced, in contrast to the typical WN composition where carbon is strongly depleted due to the CNO cycle equilibrium. Among massive WR stars are a few objects with a similar composition as our program star, which are usually considered to be caught in the transition phase between the WN and WC stage. These are classified as spectral type WN/WC or WNC. Therefore in analogy to these massive stars we suggest to classify the central star of PB\\,8 as [WN/WC]. The detailed subtype of PB\\,8 is WN6 when applying the classification scheme established by \\cite{smith1968} for massive WN stars. With the scheme in \\cite{vanderhucht1981} for massive WC stars, the WC7 subtype seems to be appropriate. In combination with these two schemes, we determine the detailed subtype classification as [WN6/WC7]. The [WC5-6] classification of PB\\,8 by \\cite{acknein2003} was partly based on the identification of spectral features with stellar C\\,{\\sc ii}, but we cannot confirm any stellar C\\,{\\sc ii} line from our high-resolution data. \\cite{tylack1993} alternatively defined the class of ``weak emission line stars'' (WELS) for those spectra that show much fainter and narrower emission lines than massive WC stars. \\cite{gesicki2006} assign this WELS classification to PB\\,8. However, the nature and homogeneity of the WELS class seems to be still unclear. There are two other known WR-type central stars with non-carbon-rich winds. One is LMC-N\\,66 in the LMC, which is only sometimes of the Wolf-Rayet type. It has an irregular nebula and seems to be a close binary \\citep[see discussion in][]{pena2004}. The other example, the central star of PMR\\,5 discovered by \\cite{morgpark2003}, is probably a Galactic [WN] star. Its spectrum shows only helium and nitrogen lines, while any carbon lines are missing. In case of PB\\,8, carbon and oxygen lines are visible. \\cite{morgpark2003} discuss the PN status of PMR\\,5 on the basis of the nebular expansion velocity and electron density. They conclude that PMR\\,5 is a normal PN. \\subsection{Evolutionary status} The surface composition of PB\\,8 appears unique among all CSPNe that have been analyzed so far. Only two other CSPNe (PMR\\,5 and the enigmatic variable LMC-N\\,66) are known to show a WN-type composition, which is dominated by helium with a significant amount of nitrogen. Two more CSPNe are known to be helium-rich, but without strong winds \\citep[LoTr\\,4 and K\\,\\mbox{1-27},][]{rauch1998}. Our program star PB\\,8 is unique in showing a significant amount of carbon, while carbon is usually depleted in WN-type compositions. Note that there is a He-sdO star without PN, KS\\,292, that shows a similar composition as PB\\,8, including the enhanced carbon abundance \\citep{rauch1991}. This poses the question of how to explain the evolutionary origin of PB\\,8. The formation of hydrogen-deficient post-AGB stars can be explained by a final thermal pulse which leads to the ingestion of the hydrogen envelope \\citep{herwig1999,herwig2001,werner2006,althaus2005}. This final thermal pulse may occur either at the tip of the AGB (AGB final thermal pulse, AFTP) or later, when the AGB has been left (late thermal pulse, LTP, or very late thermal pulse, VLTP). These models lead to a carbon-rich surface composition (carbon abundance larger than $20\\,\\%$ by mass), which is what is needed to explain the observed abundance patterns of [WC]-type central stars. However, the predicted nitrogen abundance is very small, except for the VLTP case where $X_\\text{N} \\geq 1\\%$ has been predicted \\citep{althaus2005,werner2006}. As a tentative explanation for PB\\,8, we propose that the final thermal pulse has only been ``weak'', so that only a small amount of carbon has been dredged up to the surface. The bulk of matter at the surface is then enriched by helium from the former intershell region. This material also contains nitrogen according to the equilibrium from the CNO cycle. In addition, a part of the hydrogen-rich envelope must have survived the last pulse and become mixed into the present outer layers. While it is not clear whether such a ``weak'' last thermal pulse can happen on the AGB, it might occur in an extremely late VLTP when the star is already too cool to undergo a full He-shell flash (F.\\,Herwig, private communication). Further constraints for the evolutionary origin of PB\\,8 may be derived from the planetary nebula. In Sect.\\,\\ref{sect:pb8-nebula} we showed that the present nebula is younger than 3000\\,years. There is no visible remnant of an older PN. Moreover, the nebula abundance ratios $\\text{He}\\,/\\,\\text{H}=0.123$ and $\\text{N}\\,/\\,\\text{O}=0.28$ by number \\citep{gar2009} show that PB\\,8 is not a helium-enriched Peimbert's Type\\,I PN. For the latter, $\\text{He}\\,/\\,\\text{H}>0.125$ or $\\text{N}\\,/\\,\\text{O}>0.5$ is expected \\citep{peimbert1987}. A VLTP origin of the nebula is therefore implausible. The low $\\text{N}\\,/\\,\\text{O}$ ratio also indicates the absence of hot bottom burning (HBB), which is predicted for more massive AGB stars. From a comparison with stellar evolutionary tracks, \\cite{kaler1989} deduce $\\text{N}\\,/\\,\\text{O} > 0.8$ as a sharp limit for N-enriched PNe, which are supposed to indicate HBB in AGB-stars with $M_\\text{core} > 0.8\\,\\text{M}_\\odot$. Thus the following alternative scenarios might explain our results: \\begin{enumerate} \\item The CSPN of PB\\,8 has a low mass and evolves slowly. For instance, a $0.6\\,\\text{M}_\\odot$ post-AGB star on the way to an LTP has a crossing time of $4000\\,\\text{years}$ from $10^4\\,\\text{K}$ to its maximum effective temperature \\citep{bloecker2003}. Then, either \\begin{enumerate} \\item the present nebula was ejected by a born-again AGB-star after occurrence of a ``weak'' VLTP. A possible older PN from the first AGB phase has already dissolved. As mentioned above, this scenario does not fit well to PB\\,8, as the PN is not enriched in helium; or, \\item the CSPN suffered an ``anomalous AFTP'', resulting in the observed surface abundances. The nebula was formed only during this AGB phase of the star. In this scenario it is difficult to explain the enhanced nitrogen abundance of the star, as nitrogen enrichment is neither predicted for the AFTP nor can it originate from HBB. \\end{enumerate} \\item Alternatively, the CS of PB\\,8 may have a relatively high mass and therefore may have evolved very fast. The crossing time for e.g.\\ a $0.94\\,\\text{M}_\\odot$ post-AGB star is only $50\\,\\text{years}$ \\citep{bloecker2003}. A VLTP has already occurred, but most of the nebula observed now still originates from the first AGB period, not from the born-again AGB star after the VLTP. Albeit possible, this scenario has a low probability because the empirical mass distribution of central stars has a sharp maximum at $0.6\\,\\text{M}_\\odot$ and declines substantially towards higher values \\citep{tylenda2003}; furthermore there are no hints of HBB, which would be indicative for a more massive CSPN. \\end{enumerate} However, one has to keep in mind that the appropriate stellar evolution models are still calculated only in 1D. Especially the convective mixing during H-ingestion flashes of the TPs were treated in diffusion approximation, whereas recent hydrodynamical studies, e.g.\\ by \\cite{woodward2008}, emphasize that convective mixing is rather an advection process, making 2D or 3D calculations necessary. Alternatively to single star evolution, one may consider binarity with a common envelope phase as the origin of the hydrogen deficiency. However, PB\\,8 shows no evidence of binarity. The nebula does not look bipolar. Also, \\cite{mendez1989} found no indication of radial velocity variations between three spectra taken within one year, but only changes of the P-Cygni line profiles, which must be attributed to variability of the stellar wind. Summarizing, the evolutionary origin of PB\\,8 cannot be explained by any existing model for a post-AGB star which lost its hydrogen envelope in a final thermal pulse. However, one can imagine scenarios of a weak or anomalous thermal pulse, occurring on the AGB or later, which may explain the unique chemical composition of this star and its young nebula. The chemical composition found in the expanding atmosphere of the central star of PB\\,8 differs from any known central star abundance. However, it resembles the rare transition class of WN/WC subtypes of massive Wolf-Rayet stars. Therefore we suggest to open a new class of [WN/WC]-type central stars with PB\\,8 as its first member." }, "1003/1003.2629_arXiv.txt": { "abstract": "We present observations at 250, 350, and 500 \\micron~of the nearby galaxy cluster Abell 3112 ($z$ = 0.075) carried out with BLAST, the Balloon-borne Large Aperture Submillimeter Telescope. Five cluster members are individually detected as bright submillimetre sources. Their far-infrared SEDs and optical colours identify them as normal star-forming galaxies of high mass, with globally evolved stellar populations. They all have $(B-R)$ colours of 1.38 $\\pm$ 0.08, transitional between the blue, active population and the red, evolved galaxies that dominate the cluster core. We stack to determine the mean submillimetre emission from all cluster members, which is determined to be 16.6$\\pm$2.5, 6.1$\\pm$1.9, and 1.5$\\pm$1.3 mJy at 250, 350, and 500 \\micron, respectively. Stacking analyses of the submillimetre emission of cluster members reveal trends in the mean far-infrared luminosity with respect to cluster-centric radius and \\kmag-band magnitude. We find that a large fraction of submillimetre emission comes from the boundary of the inner, virialized region of the cluster, at cluster-centric distances around \\rfive. Stacking also shows that the bulk of the submillimetre emission arises in intermediate-mass galaxies ($L < L^{*}$), with \\kmag~magnitude $\\sim$1 mag fainter than the giant ellipticals. The results and constraints obtained in this work will provide a useful reference for the forthcoming surveys to be conducted on galaxy clusters by {\\it Herschel}. ", "introduction": "\\label{intro} The evolution of cluster galaxies and their star-formation rates have been studied using several different approaches in the last few decades. Optical surveys have shown clear correlations between galaxy colours and local galaxy density or cluster-centric radius (\\citealt{Dressler80, Dressler97, Kodama01}). Spectroscopic observations have consistently identified trends in the star-formation activity of cluster galaxies, both as a function of cluster-centric distance (e.g.~\\citealt*{Verdugo08}; \\citealt{Braglia09} and references therein) and at different redshifts~\\citep{Poggianti06, Poggianti09}. Variation of the star-formation rate (SFR) and correlation with the local environment has also been investigated at different mass scales, from groups \\citep{Wilman08} to large superclusters \\citep{Porter07}, and also in relation to local large-scale filamentary structures (\\citealt{Braglia07}; \\citealt{Porter08}). Observations with {\\it IRAS} and {\\it ISO} provided a way to investigate the nature of dust and to correlate the SFRs of cluster galaxies with their dust content, albeit mostly covering the spectral regions dominated by warm ($> 40$~K) dust. While part of these studies was aimed at detecting diffuse emission from warm intracluster dust (e.g. \\citealt{Stickel98}; \\citealt{Stickel02}; \\citealt{Montier05}; \\citealt{Giard08}), several results were also obtained with observations of individual cluster members in several clusters. \\citet{Edge01} used combined {\\it IRAS}, IRAM and JCMT observations to detect CO line emission from molecular gas in the central galaxies of a sample of 16 cooling core galaxies. Tuffs et al. (2002) and Popescu et al. (2002a; 2002b) observed a large sample of galaxies in Virgo, finding a dependence in the dust content of galaxies with Hubble type. \\citet{Pierini03} found that dust luminosity and mass depend on galaxy geometry and shape as well as stellar mass. Several recent $24\\,$\\micron~observations with {\\it Spitzer}-MIPS have detected dusty star-forming galaxies in intermediate- to high-$z$ clusters. \\citet{Geach06} find an increase of the total SFR in clusters with increasing redshift from {\\it Spitzer} observations, although with large scatter. \\citet{Bai07} investigate the IR properties and the mid-IR luminosity function of cluster galaxies in a higher redshift cluster at $z = $ 0.83, confirming the presence of evolution in the star-formation rate of cluster galaxies. \\citet{Fadda08} identify consistent overdensities of $24\\,$\\micron~sources along two filaments between the clusters Abell 1770 and Abell 1763 ($z = $ 0.23) with respect to the surrounding field. Similar to \\citet{Bai07}, \\citet{Haines09a} confirm an excess of $24\\,$\\micron~sources in the cluster Abell 1758 at $z = $ 0.28. \\citet{Tran09} also compare the $24\\,$\\micron~luminous members of a cluster, a supergroup and the field, concluding that the mid-IR inferred SFR is higher in the intermediate environment of the groups than in the field, while it is globally lower in the cluster. Local dependence of the density of 24\\,\\micron~sources in clusters is investigated in the LoCuSS survey by \\citet{Haines09b}, who find a global decrease of star-forming systems with decreasing cluster-centric radius. Recently, \\citet{Wardlow10} have used the AzTEC camera to observe a field centred on the cluster MS0451.6--0305 at $z = 0.54$, identifying two luminous infrared galaxies (LIRG) with a combined SFR of 100 \\sfryr. They suggest that, if these are indeed cluster members, they can be examples of a population of galaxies undergoing transformation to the red sequence through interaction with the cluster environment. All the studies presented have investigated the star-formation activity of cluster galaxies either in the mid-IR or at millimetre wavelengths. However, a complete characterization of the output of star-formation requires coverage of the 200--800\\,\\micron~spectral region, where the peak of the far-IR emission is expected to lie. The Balloon-borne Large Aperture Submillimeter Telescope (BLAST: \\citealt{Pascale08}; \\citealt{Devlin09}) is a pathfinder experiment to {\\it Herschel}/SPIRE, and has provided the first maps of selected areas of the sky at 250, 350, and 500\\,\\micron. These wavelengths were mainly chosen to constrain the peak of the FIR emission from galaxies at redshifts $z \\gtrsim 1$. Several studies carried out by the BLAST collaboration on extragalactic fields, either on individual sources (\\citealt{Dye09}; \\citealt{Dunlop09}; \\citealt{Ivison10}), using stacking (\\citealt{Devlin09}; \\citealt{Marsden09}; \\citealt{Pascale09}; \\citealt{Ivison10}), or other statistical analyses (\\citealt{Patanchon09}; \\citealt{Viero09}), have been performed on blank-field maps. A few other studies (\\citealt{Rex09}; \\citealt{Wiebe09}) have been conducted on known targets. In particular, Rex et al. have provided the first sub-mm maps of the `Bullet' cluster (\\citealt{Tucker98}), investigating the nature of a bright sub-mm source identified as a counterpart of a lensed high-$z$ star-forming galaxy. However, no direct investigation of sub-mm emission from cluster members has been conducted so far. We present here 250, 350, and 500~\\micron~observations of a field centred on the nearby cluster Abell 3112 ($z = $ 0.075; A3112 hereafter) carried out by BLAST, and the results of a combined analysis of the optical and sub-mm properties of a spectroscopic sample of its cluster members. These results demonstrate that observation of cluster galaxies at sub-mm wavelengths can provide insight into the star-formation activity in clusters and help understanding galaxy evolution within these overdense environments. This paper is organized as follows. Section \\ref{obsdata} introduces the BLAST observations of A3112 and the ancillary optical and near-IR data used for our study. Section \\ref{results} shows the results from stacking analyses of cluster member catalogues and the properties of sub-mm bright cluster members. These results are discussed in Section \\ref{discuss} and summarized in Section \\ref{summ}. Throughout the paper, we use a standard $\\Lambda$CDM cosmology, where $\\Omega_\\rmn{M} = 0.3$, $\\Omega_{\\Lambda} = 0.7$ and $h \\equiv H_0/100~\\rmn{km}~\\rmn{s}^{-1}~\\rmn{Mpc}^{-1} = 0.7$. ", "conclusions": "\\label{summ} We have been able to provide a description of the star-formation activity of member galaxies in a cluster at sub-mm wavelengths, as observed by the BLAST experiment. A combination of stacking analyses based on cluster-centric distance, galaxy magnitude and colour allowed us to identify correlations between optical properties of cluster galaxies and their FIR emission. Studying the sub-mm SEDs of cluster members identifies BLAST-detected cluster members mostly as normal star-forming galaxies. Their optical colours and \\kmag~magnitude identify them as massive early-type galaxies in an advanced stage of evolution. The BCG is found to have low dust content, its mid-IR emission possibly due to the central AGN. Together with an expected increase of the SFR in the cluster outskirts, we find a significant increase in the 250\\,\\micron~flux density at distances around~\\rfive, i.e. further inside the cluster with respect to the infall region. The increase peaks to values 3--4 times as large as in the inner regions of the cluster. This confirms previous studies in the mid-IR that show large fractions of star-forming galaxies closer to the dense cores of galaxy clusters. The combined results of our analysis show that the cluster members identified at sub-mm wavelengths can be part of a population of evolved systems on the verge of transition from the population of blue active galaxies to the quenched systems (ellipticals and S0s) dominating the cluster cores, and suggests that environmental effects at distances of order of \\rfive~play a role regulating star-formation activity during this transition. Deeper and more complete studies of galaxy clusters at far-IR and sub-mm wavelengths with {\\it Herschel} and SCUBA-2 should provide more complete coverage of the physical processes at work in cluster galaxies. In particular, combining the higher sensitivity and resolution of SPIRE and the spectral coverage offered by parallel observations with PACS and SPIRE will allow for the detection and characterization of cluster galaxies down to smaller luminosities and masses, providing a more complete description of the star-formation activity in galaxy clusters." }, "1003/1003.0595_arXiv.txt": { "abstract": "% A new estimation of the orbital period of YY Her on the base of our and published observations is presented. Phased light curves in $(RI)_{\\rm C}$ bands show evidently ellipsoidal effect connected with the tidal distortion of the giant surface. ", "introduction": " ", "conclusions": "" }, "1003/1003.2559_arXiv.txt": { "abstract": "The clustering of X-ray selected AGN appears to be a valuable tool for extracting cosmological information. Using the recent high-precision angular clustering results of $\\sim 30000$ XMM-{\\it Newton} soft (0.5-2\\,keV) X-ray sources (Ebrero et al.), which have a median redshift of $z\\sim 1$, and assuming a flat geometry, a constant in comoving coordinates AGN clustering evolution and the AGN bias evolution model of Basilakos et al., we manage to break the $\\Omega_m-\\sigma_8$ degeneracy. The resulting cosmological constraints are: $\\Omega_m=0.27^{+0.03}_{-0.05}$, w$=-0.90^{+0.10}_{-0.16}$ and $\\sigma_{8}=0.74^{+0.14}_{-0.12}$, while the dark matter host halo mass, in which the X-ray selected AGN are presumed to reside, is $M=2.50^{+0.50}_{-1.50}\\times 10^{13}h^{-1}M_{\\odot}$. For the constant $\\Lambda$ model (w=$-1$) we find $\\Omega_m=0.24 \\pm 0.06$ and $\\sigma_{8}=0.83^{+0.11}_{-0.16}$, in good agreement with recent studies based on cluster abundances, weak lensing and the CMB, but in disagreement with the recent bulk flow analysis. {\\bf Keywords:} cosmology: cosmological parameters, large scale structure of the universe ", "introduction": "A large variety of cosmologically relevant data, based on the combination of galaxy clustering, the supernova Ia's Hubble relation, the cosmic microwave background (CMB) fluctuations and weak-lensing strongly support a flat universe, containing cold dark matter (CDM) and ``dark energy'' which is necessary to explain the observed accelerated cosmic expansion (eg., % Komatsu et al. 2010; Hicken et al. 2009; Fu et al. 2008 and references therein). The nature of the mechanism that is responsible for the late-time acceleration of the Hubble expansion is a fundamental problem in modern theoretical physics and cosmology. Due to the absence of a physically well-motivated fundamental theory, various proposals have been suggested in the literature, among which a cosmological constant, a time varying vacuum quintessence, $k-$essence, vector fields, phantom, tachyons, Chaplygin gas, etc (eg., % Weinberg 1989; Peebles \\& Ratra 2003; % Boehmer \\& Harko 2007; Padmanabhan 2008 and references therein). Note that the simplest pattern of dark energy corresponds to a scalar field having a self-interaction potential with the associated field energy density decreasing with a slower rate than the matter energy density. In such case the dark energy component is described by an equation of state $p_{Q}={\\rm w}\\rho_{Q}$ with w$<-1/3$ (dubbed ``quintessence'', eg. Peebles \\& Ratra 2003 and references therein). The traditional cosmological constant ($\\Lambda$) model corresponds to w$=-1$. The viability of the different dark-energy models in reproducing the current excellent cosmological data and the requirements of galaxy formation is a subject of intense work (eg. Basilakos, Plionis \\& Sol\\'a 2009 and references therein). Another important cosmological parameter is the normalization of the cold dark matter power spectrum in the form of the rms density fluctuations in spheres of radius 8$h^{-1}$ Mpc, the so-called $\\sigma_8$. There is a degenerate relation between $\\sigma_8$ and $\\Omega_m$ (eg. Eke, Cole \\& Frenk 1996; Wang \\& Steinhardt 1998; Henry et al. 2009; Rozo et al. 2009 and references therein) and it is important to improve current constraints in order to break such degeneracies. Furthermore, there are also apparent inconsistencies between the values of $\\sigma_8$ provided by different observational methods, among which the most deviant and problematic for the {\\em concordance} cosmology, is provided by the recent bulk flow analysis of Watkins, Feldman \\& Hudson (2009). In this paper we extend our previous work (Basilakos \\& Plionis 2009; hereafter BP09), using the angular clustering of the largest sample of high-$z$ X-ray selected active galactic nuclei (Ebrero et al. 2009a), in an attempt to break the $\\sigma_{8}-\\Omega_m$ degeneracy within spatially flat cosmological models. ", "conclusions": "We have used the recent angular clustering measurements of high-$z$ X-ray selected AGN, identified as soft (0.5-2 keV) XMM point sources (Ebrero et al. 2009a), in order to break the degeneracy between the rms mass fluctuations $\\sigma_{8}$ and $\\Omega_m$. Applying a standard likelihood procedure, assuming a constant in comoving coordinates AGN clustering evolution, the bias evolution model of Basilakos et al. (2008) and a spatially flat geometry, we put relatively stringent constraints on the main cosmological parameters, given by: $\\Omega_m=0.27^{+0.03}_{-0.05}$, w$=-0.90^{+0.10}_{-0.16}$ and $\\sigma_{8}=0.74^{+0.14}_{-0.12}$. We also find that the dark matter host halo mass, in which the X-ray selected AGN are assumed to reside, is $M=2.50^{+0.50}_{-1.50}\\times 10^{13}h^{-1}M_{\\odot}$. Finally, if we marginalize over the previous host halo mass and w$=-1$ ($\\Lambda$ cosmology), we find $\\Omega_m=0.24\\pm 0.06$ and $\\sigma_{8}=0.83^{+0.11}_{-0.16}$." }, "1003/1003.3463_arXiv.txt": { "abstract": "We present high-sensitivity 2$'$$\\times$4$'$ maps of the $J$=2$\\rightarrow$1 rotational lines of SiO, CO, $^{13}$CO and C$^{18}$O, observed toward the filamentary Infrared Dark Cloud (IRDC) G035.39-00.33. Single-pointing spectra of the SiO $J$=2$\\rightarrow$1 and $J$=3$\\rightarrow$2 lines toward several regions in the filament, are also reported. The SiO images reveal that SiO is widespread along the IRDC (size $\\geq$2$\\,$pc), showing two different components: one bright and compact arising from three condensations (N, E and S), and the other weak and extended along the filament. While the first component shows broad lines (linewidths of $\\sim$4-7$\\,$km$\\,$s$^{-1}$) in both SiO $J$=2$\\rightarrow$1 and SiO $J$=3$\\rightarrow$2, the second one is only detected in SiO $J$=2$\\rightarrow$1 and has narrow lines ($\\sim$0.8$\\,$km$\\,$s$^{-1}$). The maps of CO and its isotopologues show that low-density filaments are intersecting the IRDC and appear to merge toward the densest portion of the cloud. This resembles the molecular structures predicted by flow-driven, shock-induced and magnetically-regulated cloud formation models. As in outflows associated with low-mass star formation, the excitation temperatures and fractional abundances of SiO toward N, E and S, increase with velocity from $\\sim$6 to 40$\\,$K, and from $\\sim$10$^{-10}$ to $\\geq$10$^{-8}$ respectively, over a velocity range of $\\sim$7$\\,$km$\\,$s$^{-1}$. Since 8$\\,$$\\mu$m sources, 24$\\,$$\\mu$m sources and/or extended 4.5$\\,$$\\mu$m emission are detected in N, E and S, broad SiO is likely produced in outflows associated with high-mass protostars. The excitation temperatures and fractional abundances of the narrow SiO lines, however, are very low ($\\sim$9$\\,$K and $\\sim$10$^{-11}$, respectively), and consistent with the processing of interstellar grains by the passage of a shock with $v_s$$\\sim$12$\\,$km$\\,$s$^{-1}$. This emission could be generated i) by a large-scale shock, perhaps remnant of the IRDC formation process; ii) by decelerated or recently processed gas in large-scale outflows driven by 8$\\,$$\\mu$m and 24$\\,$$\\mu$m sources; or iii) by an undetected and widespread population of lower mass protostars. High-angular resolution observations are needed to disentangle between these three scenarios. ", "introduction": "Infrared Dark Clouds (IRDCs) are high-extinction regions viewed against the diffuse mid-IR Galactic background \\citep[][]{per96,egan98}. These clouds are cold \\citep[$T<$25$\\,$K;][]{pill07} and exhibit a range of densities from $n$(H)$\\geq$10$^3$$\\,$cm$^{-3}$ to $\\geq$10$^4$-10$^5$$\\,$cm$^{-3}$ in their clumps and cores \\citep[][]{tey02,but09}. Since these structures have masses and mass surface densities similar to regions that are known to be forming massive protostars and star clusters, they may represent the initial conditions for massive star and star cluster formation \\citep[Rathborne, Jackson \\& Simon 2006;][]{zha09,rag09}. It is well-known that silicon monoxide (SiO) is an excellent tracer of molecular gas processed by shocks. While SiO is heavily depleted onto dust grains in the quiescent gas of dark clouds such as L183 \\citep[upper limits of the SiO fractional abundance of $\\leq$10$^{-12}$; Ziurys, Friberg \\& Irvine 1989;][]{req07}, this molecule is enhanced by large factors (in some cases by $>$10$^6$) toward molecular outflows \\citep[][]{mar92}. This is due to the injection of molecular material into the gas phase by the processing of dust grains \\citep[e.g.][]{cas97,sch97,jim08,gui07,gui09}. The typical SiO emission measured in molecular outflows shows broad line profiles with linewidths of some tens of km$\\,$s$^{-1}$ \\citep[][]{mar92}. Narrower SiO lines have also been detected toward low-mass star forming regions such as NGC1333 and L1448-mm \\citep[][]{lef98,jim04}. Although the nature of this emission is not clear yet, \\citet{lef98} have proposed that these lines could trace shocked material deflected and decelerated by the interaction with pre-existing clumps. Alternatively, narrow SiO could arise from gas recently processed by the magnetic precursor of young magneto-hydrodynamic (MHD) shocks \\citep{jim04}. In the case of IRDCs, \\citet{mott07} and \\citet{beu07} have recently carried out two large surveys of SiO emission toward the Cygnus X molecular cloud complex and toward a sample of IRDCs, respectively. Their single-pointing observations show that the detection rate of SiO toward IR-quiet massive cores close to sites of on-going star formation, is relatively high. This is expected since SiO is tightly associated and restricted to shocked gas in outflows. However, no widespread narrow SiO emission, as seen toward NGC1333, has been reported in IRDCs so far. We present the detection of widespread SiO $J$=2$\\rightarrow$1 emission (size of $\\geq$2$\\,$pc) toward the IRDC G035.39-00.33 \\citep[Cloud H in][]{but09}. From its filamentary morphology, this IRDC is believed to be at the early stages of its evolution as predicted by dynamical models of giant molecular cloud formation \\citep[][Heitsch, Stone \\& Hartmann 2009]{van07,hen08}. The observed large-scale SiO feature is probably a composition of broad and compact emission, linked to outflows associated with high-mass star formation; and extended narrow SiO emission. The observed narrow SiO lines could be explained i) by a large-scale shock, remnant of the IRDC formation process; ii) by decelerated or recently processed gas in the precursor of MHD shocks in large-scale outflows, probably driven by the 8$\\,$$\\mu$m and 24$\\,$$\\mu$m sources observed in the IRDC; or iii) by an undetected and widespread population of lower-mass protostars. ", "conclusions": "" }, "1003/1003.1716_arXiv.txt": { "abstract": "We examine correlations between the masses, sizes, and star formation histories for a large sample of low-redshift early-type galaxies, using a simple suite of dynamical and stellar populations models. We confirm an anti-correlation between size and stellar age, and survey for trends with the central content of dark matter (DM). An average relation between central DM density and galaxy size of $\\langle\\rho_{\\rm DM}\\rangle \\propto \\Re^{-2}$ provides the first clear indication of cuspy DM haloes in these galaxies---akin to standard $\\Lambda$CDM haloes that have undergone adiabatic contraction. The DM density scales with galaxy mass as expected, deviating from suggestions of a universal halo profile for dwarf and late-type galaxies. We introduce a new fundamental constraint on galaxy formation by finding that the central DM fraction decreases with stellar age. This result is only partially explained by the size-age dependencies, and the residual trend is in the opposite direction to basic DM halo expectations. Therefore we suggest that there may be a connection between age and halo contraction, and that galaxies forming earlier had stronger baryonic feedback which expanded their haloes, or else lumpier baryonic accretion that avoided halo contraction. An alternative explanation is a lighter initial mass function for older stellar populations. ", "introduction": "\\label{sec:intro} The formational history of early-type galaxies (ETGs: ellipticals and lenticulars) remains an outstanding question. While these dynamically hot systems may be basically understood as end-products of galaxy mergers, the details of these mergers and their cosmological context are unclear. High-redshift ($z$) observations have made initially surprising discoveries that many ETGs were already present at early times with mature stellar populations, and that these galaxies were much more compact than those in the present day. (e.g. \\citealt{2004Natur.430..181G,2005ApJ...626..680D,2006ApJ...650...18T}). The evolution in ETG sizes is still controversial in both observation and interpretation (e.g. \\citealt{2009Natur.460..717V,2010ApJ...709.1018V,2010MNRAS.401..933M,2009arXiv0907.2392V}). However, the most likely scenario is for a combination of effects where individual galaxies grow in size by accretion of smaller, gas-poor galaxies (an ``inside-out'' picture of galaxy formation), and where younger ETGs are formed with larger sizes because of their decreased cold gas content and the lower background densities of dark matter (DM; e.g. \\citealt{2009MNRAS.392..718S,2009ApJ...697.1290B,2009ApJ...698.1232V,2009ApJ...699L.178N,2010MNRAS.401.1099H,2009arXiv0912.0012S}). The role of DM is considered fundamental to the formation of galaxies, and DM halo properties have been extensively studied in cases such as gas-rich spirals and nearby dwarfs which have suitable observational tracers (e.g. \\citealt{2007ApJ...659..149M,2009ApJ...704.1274W,2009arXiv0911.1998K}). Studying DM in the general population of ETGs is in many ways more difficult, with ongoing surveys of the available large-radius halo tracers attempting to remedy our ignorance in this area (e.g. \\citealt{2009AJ....137.4956R,2009MNRAS.394.1249C,2009MNRAS.398...91P}; \\citealt[hereafter T+09]{2009ApJ...691..770T}; \\citealt{2010arXiv1002.3142W}). DM can alternatively be studied with less precision but in more extensive ETG samples by considering the well-studied central regions---inside the ``effective radius'' (\\Re) enclosing half the stellar light, where DM is generally thought to be a minor yet potentially detectable contributor to the mass. Here one of the classic approaches is to analyze the ``fundamental plane'' (FP) relating ETG sizes, luminosities, and central velocity dispersions (\\sigc). The FP shows a ``tilt'' or systematic deviation from simple expectations based on galaxies with constant dynamical mass-to-light ratios \\ML{}, probably implying systematic differences in the stellar populations or in DM content. After many years of debate, there is still not a consensus on what is driving the FP tilt (e.g. \\citealt{2004ApJ...600L..39T,2005ApJ...623L...5F}; \\citealt[C+06 hereafter]{Cappellari06}; \\citealt{2006MNRAS.369.1081B,2006MNRAS.370.1445D,2007ApJ...665L.105B,2008ApJ...684..248B,2008ApJ...678L..97J}; \\citealt[hereafter paper I]{2009MNRAS.396.1132T}; \\citealt{2009ApJ...702.1275A,2009PhDT.........9G,2010MNRAS.402L..67G,2010MNRAS.tmp..135H,2009arXiv0912.4558L}). Some work suggests stellar populations variations or non-homologies in the luminosity profiles as the tilt drivers. However, it is a fairly generic expectation from the standard cosmological framework for galaxy formation that the central DM content of ETGs will systematically increase with luminosity---a point we discussed in paper I and develop further in this paper. If the FP tilt is {\\it not} caused in large part by DM, there could be problems implied for galaxy formation theory. Here one could pursue two different philosophies: to empirically and robustly determine the reasons for the FP tilt without recourse to theory (e.g. making no assumptions about the underlying DM profiles); or to adopt the theoretical framework as broadly correct and consider the detailed implications for ETG composition and formation. Following the second approach, it is now time to begin moving on from phenomenological questions about the FP tilt, and to establish more direct connections between DM in ETGs and their formational histories. In particular, the central DM content could prove crucial to solving the size-evolution puzzles mentioned above, and more fundamentally to understanding the assembly of ETGs. To this end, we now extend the analysis of paper I, which combined models of stellar dynamics and population synthesis to infer {\\it total} and {\\it stellar} masses in a large sample in nearby galaxies, and thereby to analyze the FP tilt. There have been previous suggestions that age and star formation timescales are important fourth parameters in the FP \\citep{2002MNRAS.330..547T,2009MNRAS.397...75G,2009PhDT.........9G}. We will now explore this possibility systematically, using the results from paper I to consider additional correlations involving DM content and star formation histories (SFHs). It should be noted at the outset that the data and the analysis techniques that we use for deriving mass and SFH parameters may not be the most state-of-the-art. However, our aim is to pioneer a framework for interpreting any data of this kind in a broad cosmological context (where in particular, emerging high-$z$ data-sets provide only crude observational constraints on ETGs), and we are so far able to tentatively identify some basic and intriguing trends. The paper is organized as follows. In Section~\\ref{sec:observ} we present our basic observational results. We go on to analyze some implications of the trends of DM with mass in Section~\\ref{sec:dmimp} and with age in Section~\\ref{sec:lcdm2}. In Section~\\ref{sec:concl} we summarize our conclusions. Two Appendices include analyses of statistical and systematic uncertainties. ", "conclusions": "\\label{sec:concl} We have continued an analysis begun in paper I of a large data-set of nearby early-type galaxies (ETGs), combining dynamics and stellar populations to constrain the central DM content. After having identified variations in DM as the main cause of the tilt of the fundamental plane, we have moved on to consider various scaling relations of the DM haloes, and connections to the star formation histories of the galaxies. Our basic observational findings are that the central DM fraction \\fDM\\ within an effective radius \\Re\\ has a strong anti-correlation with stellar age, and that the galaxy sizes also have an age anti-correlation. We have constructed composite profiles of DM density with radius, finding that they are on average cuspy, with inferred density exponents of $\\sim -1.6$ near \\Re. These profiles are steeper than literature findings for spiral galaxies, and the central DM densities of the early-types are denser overall, suggesting that gas-rich mergers would need to produce a net halo contraction. To further interpret the data, we have generated a series of $\\Lambda$CDM toy models, including variable contributions from adiabatic contraction (AC). The results from comparisons of models to data are: \\begin{itemize} \\item Models with AC fit well overall with a Kroupa IMF, while models without AC prefer a Salpeter IMF. \\item The size-age trends can explain {\\it part} of the \\fDM-age trends: older galaxies show less evidence for DM because their more compact stellar centres probe less volume of the DM halo. \\item The remaining \\fDM-age trends are not easily explained by variations in halo mass or concentration, and suggest differences in baryonic effects on the DM, in the sense that younger galaxies have undergone AC while older galaxies have not. \\item An alternative scenario is for the IMF to be less massive for older stellar populations. \\end{itemize} There is ample scope for future insights and improvements. We plan to further investigate the galaxies' star formation histories in the context of theoretical mass assembly histories. Environmental trends can be investigated, as these are expected to be important (e.g. \\citealt{Thomas+05,2006ApJ...652...71W,2010arXiv1002.0835R,2010arXiv1002.0847N,2010arXiv1003.1119L}). Forthcoming high-quality, homogeneous, multiwavelength large surveys of low-redshift ETGs should also be able to refute, confirm or extend the trends presented here \\citep{2009PhDT.........9G,2009arXiv0908.1904C}. Finally, the gold standard for probing galaxy mass profiles is extended kinematics data along with detailed dynamical modelling---both to provide more leverage on the DM independently of the stellar mass, and to sift individual galaxies for the presence of a DM core or cusp (e.g. \\citealt{2007MNRAS.382..657T,2009MNRAS.395...76D,2009MNRAS.398..561W,2009PhDT.........7F}). Even if some of our current conclusions turn out to be completely wrong, we hope to have introduced a useful framework for interpreting mass results for large data sets of ETGs over cosmic time. The DM constraints are part of a spectrum of clues that can ultimately be combined to pin down the modes of ETG formation (e.g. \\citealt{2008MNRAS.384...94C}). Other avenues with considerable promise include central rotation \\citep{2009MNRAS.397.1202J}, extra light \\citep{2010MNRAS.402..985H} and orbital structure \\citep{2008ApJ...685..897B}, as well as halo rotation \\citep{2010arXiv1001.0799H} and metallicity gradients \\citep{2009MNRAS.398..561W,2009MNRAS.400.2135F}, and globular cluster constraints \\citep{2007AJ....134.1403R,2009ApJ...691...83S,2010MNRAS.401L..58B}." }, "1003/1003.3239_arXiv.txt": { "abstract": "We study filamentary structure in the galaxy distribution at $z \\sim 0.8$ using data from the Deep Extragalactic Evolutionary Probe 2 (DEEP2) Redshift Survey and its evolution to $z \\sim 0.1$ using data from the Sloan Digital Sky Survey (SDSS). We trace individual filaments for both surveys using the Smoothed Hessian Major Axis Filament Finder, an algorithm which employs the Hessian matrix of the galaxy density field to trace the filamentary structures in the distribution of galaxies. We extract 33 subsamples from the SDSS data with a geometry similar to that of DEEP2. We find that the filament length distribution has not significantly changed since $z \\sim 0.8$, as predicted in a previous study using a $\\Lambda$CDM cosmological N-body simulation. However, the filament width distribution, which is sensitive to the non-linear growth of structure, broadens and shifts to smaller widths for smoothing length scales of $5-10$ \\mpch from $z \\sim 0.8$ to $z \\sim 0.1$, in accord with N-body simulations. ", "introduction": "The observed large-scale distribution of galaxies shows dense linear features: filaments of galaxies which surround huge voids that appear largely empty, while rich clusters are found at their intersection \\citep{dav82,del86,bon96,got05}. These structures are widely believed to have evolved through gravitational instability from small density fluctuations in the early universe. The evolution of large scale structure with cosmic time can probe the complex physics that governs the creation of galaxies in their host dark matter potential wells. A number of studies \\citep[e.g.][]{coh96,con96,gia98,bro03,phl03,coi04b,ouc04,lef05, men06} have focused on the redshift dependence of the galaxy two point correlation function as a critical test of both cosmological and galaxy evolution models. The comoving correlation length of galaxies is observed to be almost constant with redshift, which is interpreted as a consequence of the increasing bias of galaxies with redshift. The two-point correlation function is a complete statistical measure of galaxy clustering only in the linear regime. Statistics of galaxy filaments, such as their lengths and widths, can be used as another useful tool to measure the large scale structure and test both cosmology and galaxy formation models. Filaments, with a typical length of $50-70$ \\mpch \\citep{bha04}, have been seen in every wide-field redshift survey, from the Great Wall of the CFA2 \\citep{del86,gel89} to the very long filaments found \\citep{got05} in the Sloan Digital Sky Survey (SDSS) \\citep{yor00}. Features qualitatively similar to the observed filamentary structures are also seen in numerical cosmological simulations. Various techniques have been proposed to identify and characterize filaments in observational and simulated samples \\citep[e.g.][] {moo83,eri04,sto05,lac05,nov06,ara07,sou08a,sou08b,sou09,for09}. \\citet[][hereafter Paper I and II, repectively]{bon09a,bon09b} use the eigenvectors of the Hessian matrix of the smoothed galaxy density field to identify and quantify filamentary structures. The filament length and width distributions of the observed local galaxy distribution from SDSS are consistent with those from N-body simulations at $z \\sim 0$ adopting a standard cosmology. The time evolution of the filament network was studied in Paper II using cosmological N-body simulations; they found that the backbone of the filamentary structure is in place at $z=3$. These simulations show that non-linear growth of structure has little impact on the length of filamentary structures, but a great deal on the width. The dark matter filament width distribution evolves from $z\\sim3$ to $z\\sim0$, broadening and peaking at smaller widths as the universe expands. Although many recent papers have studied the filamentary structures of local galaxy surveys \\citep{sto07,sou08a,sou08b,gay09}, no study has been done on the evolution of filaments in redshift survey data. This is mainly due to the small volumes and the resulting severe cosmic variance of existing high redshift surveys. However, thanks to the successful completion of the Deep Extragalactic Evolutionary Probe 2 (DEEP2) Galaxy Redshift Survey \\citep{dav03}, we can study the galaxy distribution at $z \\sim 1$ over a large comoving volume ($5 \\times 10^6$ $h^{-3} \\rm Mpc^3$) over four widely separated fields. In this study, we present measurements of filament statistics both for the galaxy distribution at $z \\sim 0.8$ using the DEEP2 Galaxy Redshift Survey \\citep{dav03}, which is an R-band-selected survey with a sampling density comparable to local surveys, and for the local galaxy distribution from the SDSS. We identify filaments in the galaxy distribution in DEEP2 using the Hessian matrix method of Paper I and II, and draw subsamples from the SDSS redshift survey with the DEEP2 geometry and sampling, to make a direct comparison between the two. This paper is organized as follows: \\S~\\ref{Method} summarizes the methods we use to find filamentary structures and measure their properties, referring the reader to Paper I and Paper II for more details. In \\S~\\ref{sec:data}, we provide details of the data samples used here. \\S~\\ref{Result} presents our results and \\S~\\ref{sec:summary} discusses their meaning and implications. We assume a flat concordance $\\Lambda$CDM cosmology with $\\Omega_m =0.3$, $\\Omega_{\\Lambda} = 1-\\Omega_m = 0.7$ and $H_0 = 100$ $h$ km s$^{-1}$ Mpc$^{-1}$ throughout this paper. ", "conclusions": "\\label{sec:summary} We study the time evolution of the filament network in the galaxy distribution by comparing the filamentary structure at $z \\sim 0.8$ from the Deep Extragalactic Evolutionary Probe 2 (DEEP2) Redshift Survey and those at $z \\sim 0.1$ from the Sloan Digital Sky Survey (SDSS). We trace individual filaments for both surveys using SHMAFF, an algorithm which employs the Hessian matrix to trace the filamentary structures in the distribution of structure. We define three subsamples from DEEP2 and 33 subsamples from SDSS, with the same sampling and geometry, namely a box of $320 \\times 40 \\times 14$ $\\rm (h^{-1} Mpc)^3$. We smooth the galaxy distribution with length scales of $l=5$ and $10$ $h^{-1}$ Mpc, and trace individual filaments along the axis of structure, and mark the end of filaments when the axis orientation changes more rapidly than a preset threshold of $C=30$ and $40$$^{\\circ} l^{-1}$ respectively. We found 97 filaments in DEEP2 subsamples and 957 filaments in SDSS subsamples with smoothing length $l=5$ $h^{-1}$ Mpc, and 24 filaments for DEEP2 and 230 for SDSS for $l=10$$h^{-1}$ Mpc. Thus the number of filaments per unit volume is unchanged from high to low redshift. We find that filament length distribution has not changed significantly since $z \\sim 1$, however, the filament width distribution, which is sensitive to non-linear growth of structure, broadens and shifts to smaller widths for smoothing length scales of $5$ and $10$ \\mpch from $z \\sim 0.8$ to $z \\sim 0.1$. The evolution in the length and width distributions is consistent with predictions from a $\\Lambda$CDM cosmological N-body simulation. As found in Paper II, non-linear growth of structure has a great impact on the width of filamentary structures. We restricted our study to two-dimensional analysis due to the the geometry of the DEEP2 survey. In order to better show the filamentary evolution, however, the filamentary structures should be studied in larger volumes and analyzed in three dimensions. The next generation of galaxy surveys, which will target the early universe with larger volume and depth can open up the possibility of detailed study of filamentary structures and their evolution. These surveys include the Advanced Dark Energy Physics Telescope (ADEPT), a space-based spectroscopic survey that promises to determine the location of 100 million galaxies at $1 < z < 2$, the BigBOSS \\citep{sch09}, a proposed ground-based wide field spectroscopic survey at $0.2\\dMign$, which are strong emitters of $\\nu$ and $\\bar\\nu$. We found that the net energy deposition rate due to this process, $\\dE$, is well approximated by a simple formula (see Fig.~4), \\begin{eqnarray} \\label{eq:Edot}\\nonumber \\dE & \\approx & 1.1\\times 10^{52}\\,\\xms^{-4.8}\\, \\left(\\frac{M}{3M_\\odot}\\right)^{-3/2} \\\\ & \\times & \\left\\{ \\begin{array}{ll} 0 & \\dM<\\dMign \\\\ \\dot{m}^{9/4} \\;\\; & \\dMign<\\dM<\\dMtr \\\\ \\dot{m}_{\\rm trap}^{9/4}\\;\\; & \\dM>\\dMtr \\\\ \\end{array} \\right\\} \\,{\\rm erg~s}^{-1}, \\end{eqnarray} where $\\dot{m}=\\dM/M_{\\odot}$~s$^{-1}$, $\\xms=\\rms(a)/r_{g}$, $r_g=2GM/c^2$, and $\\dMtr$, $\\dMign$ are given in equation~(\\ref{eq:pw}). The dependence of $\\dE$ on the black hole spin is huge: $\\xms^{-4.8}$ varies by a factor of 200 for $0r_{\\rm c}(a)$, where $r_{\\rm c}$ is a critical $a$-dependent radius at which the radial SSF vanishes. The dominant conservative effect of the SSF in Schwarzschild spacetime is known to be of 3rd post-Newtonian (PN) order (with a logarithmic running). Our numerical results suggest that the leading-order PN correction due to the black hole's spin arises from spin-orbit coupling at 3PN, which dominates the overall SSF effect at large $r_0$. In PN language, the change-of-sign of the radial SSF is attributed to an interplay between the spin-orbit term ($\\propto -ar_0^{-4.5}$) and the ``Schwarzschild'' term ($\\propto r_0^{-5}\\log r_0$). ", "introduction": "The gravitational two-body problem is extremely difficult to tackle in a general-relativistic context, due to the intrinsic nonlinearities of the theory. However, when one of the two components is much more massive than the other the problem simplifies and can sometimes be attacked via black hole perturbation theory. Nature provides us with such extreme mass-ratio systems in the form of compact objects inspiraling into massive black holes in galactic nuclei. Such systems are key targets for the planned space-based gravitational wave detector LISA (Laser Interferometer Space Antenna) \\cite{LISA}. Detection of the gravitational waves and accurate extraction of the physical parameters requires precise theoretical templates of the waveforms, which, in turn, necessitate knowledge of the radiative evolution of the system. The underlying theoretical problem, in its most fundamental form, is that of a pointlike particle orbiting a black hole of a much larger mass. The interaction of the particle with its own gravitational field gives rise to a gravitational self-force (GSF), which is responsible in particular for the radiative inspiral. How to calculate this GSF has been the subject of extensive study over the last decade \\cite{Barack-review}. The fundamental formalism for calculations of the GSF in curved spacetime was first laid down by Mino, Sasaki, and Tanaka \\cite{Mino-Sasaki-Tanaka} and independently by Quinn and Wald \\cite{Quinn-Wald}, with important later supplements by Detweiler and Whiting \\cite{Detweiler-Whiting}, Gralla and Wald \\cite{Gralla-Wald}, Pound \\cite{Pound} and Harte \\cite{Harte} (See Poisson for a review \\cite{Poisson-review}). The resulting equations of motion are known as the MiSaTaQuWa equations. The analogous self-forced equation of motion for the electromagnetic case was derived by DeWitt and Brehme long ago \\cite{DeWitt-Brehme} (with corrections by Hobbs \\cite{Hobbs}) and reproduced more recently using other methods in \\cite{Quinn-Wald, Gralla-Harte-Wald}. Quinn obtained the equivalent results for the scalar field self-force (SSF) \\cite{Quinn}. The MiSaTaQuWa equations of motion are hard to implement directly and so they were later recast into forms more amenable to practical calculation. One of the standard methods is the mode-sum scheme first introduced in Ref.\\ \\cite{mode-sum-orig}. Using this method, self force calculations have been performed for a range of problems. These include calculations of the SSF for radial infall \\cite{Barack-Burko}, circular \\cite{Burko-circular,Diaz-Rivera} and eccentric \\cite{Haas} orbits; the electromagnetic self-force for eccentric orbits \\cite{Haas-em-eccentric}, and the GSF for radial infall \\cite{Barack-Lousto}, circular \\cite{Barack-Sago-circular, Detweiler-circular}, and eccentric orbits \\cite{Barack-Sago-eccentric}. More recently, researchers have been exploring alternative calculation methods which are based on direct regularization of the self interaction in 2+1 and 3+1 dimensions \\cite{Barack-Golbourn-Sago, Lousto-Nakano, Vega}. Common to all calculations presented so far is the fact that they specialize to the simpler (but less astrophysically relevant) case where the central object is a non-rotating, Schwarzschild black hole. In this paper we open a new front in self force calculations by considering extreme mass-ratio systems where the central black hole is rotating. The motivation for this is clear: Although little is known about the spin distribution of astrophysical massive black holes (but see, e.g., \\cite{Lousto-Nakano-Zlochower-Campanelli,Blum-etal}), there is no reason to think that massive holes in nature are non-rotating. Hence, a useful model of a LISA-relevant inspiral must incorporate a Kerr black hole as a central object. Indeed, as this work demonstrates, the spin of the central hole may have a very pronounced effect on the value of the self force and hence on the inspiral dynamics. Computing the GSF for generic inspiral orbits in Kerr is an extremely challenging task, and this work only represents a first step toward this ultimate goal. The recent advance in calculations of the GSF in Schwarzschild \\cite{Barack-Sago-ISCO-shift} was achieved after nearly a decade of development, in which the necessary computational techniques had been devised mainly by using the SSF as a simple test bed. In preparing to tackle the Kerr problem, we once again resort here to the simplicity of the scalar field toy model. Furthermore, as a primer, we specialize to (geodesic) orbits which are both circular and equatorial. This setup already captures much of the complexity of the Kerr problem (and, indeed, offers an opportunity to explore some qualitatively new physics), while providing a more manageable environment for development. Our calculation represents a first application of the standard mode-sum scheme for orbits in Kerr. As such, it provides a first test of the regularization parameter values derived in Ref.\\ \\cite{Barack-Ori} (we shall review the notion of regularization parameters in Sec.\\ III below). We opt here to work in the frequency domain, with the obvious advantage that we then only need to deal with {\\em ordinary} differential equations (ODEs). We decompose the scalar field equation in a basis of spheroidal harmonics (which are frequency-dependent), and solve the resulting ODEs numerically, with suitable boundary conditions. Since the mode-sum scheme requires as input the {\\em spherical}-harmonic modes of the scalar field gradient, we then need to re-expand the spheroidal-harmonic solutions into spherical-harmonic components. A major technical hurdle intrinsic to this procedure is that the discontinuity of the spherical-harmonic components across the particle's orbit hampers the convergence of the frequency series there, due to the Gibbs phenomenon. This problem was analyzed in depth in Ref.\\ \\cite{Barack-Ori-Sago}, and a simple and elegant solution was proposed, which entirely circumvents the problem. With this recent development, the frequency-domain approach becomes an attractive option for SSF studies, in our view. (We remark that the above Gibbs phenomenon issue does not manifest itself in the case of circular orbits considered in our current work.) In this work we calculate the dissipative and conservative components of the SSF for a variety of orbital radii and black hole spins. Our results for the dissipative component are found to agree well with the numerical results of Gralla \\etal \\cite{Gralla} (computed from asymptotic fluxes), as well as with the analytic results of Gal'tsov \\cite{Galtsov} at large orbital radii. As a further important test of our code we verify that the work done by the dissipative component of the SSF precisely balances the flux of energy in the scalar waves radiated out to infinity and through the event horizon, as extracted from our numerical solutions. For the conservative component our code recovers the results of Diaz-Rivera \\etal \\cite{Diaz-Rivera} in the Schwarzschild case. This conservative piece is calculated here for the first time for a nonzero Kerr spin parameter, revealing several interesting new features. Our main results for the conservative SSF are displayed in figure \\ref{fig:r-a-plane}. The remainder of this paper is structured as follows. In Sec.\\ \\ref{section:setup-and-formalism} we review the relevant features of circular equatorial geodesics of the Kerr geometry, and describe the setup of our problem. In Sec.\\ \\ref{section:mode-sum} we discuss the application of the mode-sum scheme for orbits in Kerr, attempted here for the first time. Section \\ref{section:numerical-implementation} describes our numerical method, and in Sec.\\ \\ref{section:analysis} we provide various validation tests of our code and present our results. Lastly in Sec.\\ \\ref{sec:conclusion} we summarize our results and consider future work. Throughout this work we use Boyer-Lindquist coordinates $(t,r,\\theta, \\phi)$, with metric signature $(-+++)$ and geometrized units such that the gravitational constant and the speed of light are equal to unity. ", "conclusions": "\\label{sec:conclusion} In this work we presented a first calculation of the SSF experienced by a particle orbiting a Kerr black hole, specializing to circular and equatorial geodesic orbits. This represented a first application of the mode-sum method in Kerr, and as a by-product we confirmed the analytic values of the regularization parameters $A_{\\alpha}$, $B_{\\alpha}$ and $C_{\\alpha}$, as calculated in \\cite{Barack-Ori}, for the above class of orbits. Our numerical calculation relied on a standard frequency-domain decomposition of the scalar field equation in terms of spheroidal harmonics; the spherical-harmonic contributions required within the regularization procedure were obtained by projecting the spheroidal-harmonic contributions onto a basis of spherical harmonics. We tested the performance of our code in various ways. The contribution to the SSF from the high-$l$ modes was found to possess the expected behavior, falling off exponentially for the time component and as $\\sim l^{-2}$ for the radial component. We confirmed numerically that the work done by the time component of the SSF precisely balances the energy in scalar waves radiated out to infinity and down through the event horizon. The energy flux calculated from our code also agreed closely with the previous numerical results by Gralla \\etal \\cite{Gralla} as well as with Galt'sov's analytic formula \\cite{Galtsov} in the large-$r_0$ regime. The radial, conservative component of the SSF was calculated here for the first time. Our code produces good agreement with the previous results of Diaz-Rivera \\etal \\cite{Diaz-Rivera} in the Schwarzschild case. For non-zero spin, we observed a qualitatively new behavior: The SSF on prograde orbits with radius larger than a certain $a$-dependent radius $r_c$ turns from repulsive (as in the Schwarzschild case) to attractive. While we have no genuine physical intuition to explain the direction of the radial SSF (not even in the Schwarzschild case), we observed, at a formal level, that the above change-of-sign may be attributed to a competition between a repulsive ``Schwarzschild'' term and an attractive spin-orbit coupling term. This observation came from fitting our numerical SSF data to an analytic PN model at large $r_0$. We thus derived a numerical approximation for the leading-order, 3PN spin-correction term. It would be interesting to test our result against an analytic PN computation of the radial SSF, once the PN result becomes available. To further make contact with PN theory it would be necessary to extract higher-order terms in the PN series, and for this it may be necessary to improve the accuracy of our code at large orbital radii. The main limiting factor, and by far the dominant source of error in our calculation, is the large contribution to the SSF from the long uncomputed tail of the $l$-mode series. The relative contribution of this tail increases with $r_0$; in our analysis the uncomputed tail contribution for $r_0=100M$ is more than twice that of the computed modes! The problem can be mitigated in future work by pushing our numerical calculation to higher $\\ell$, or---better still---by obtaining analytic expressions for some of the higher-order terms in the $1/l$ mode-sum, thereby accelerating the convergence of the mode-sum. (This latter technique was applied successfully by Detweiler \\etal in the Schwarzschild case \\cite{Detweiler-Messaritaki-Whiting}.) As mentioned in the introduction, in general a frequency-domain application of the mode-sum method is made difficult by the bad convergence of the frequency mode sum along the particle's orbit (``Gibbs phenomenon''). The problem is unnoticed for circular orbits, since in this case the scalar field is a smooth function of time along the orbit. However, the issue will need to be addressed in contemplating the extension of our code to more generic orbits. The recently introduced method of ``extended homogeneous solutions'' \\cite{Barack-Ori-Sago} proposes a simple method to overcome the above difficulty and we envisage incorporating this method in a future extended version of our code. We have already started work to generalize the code to eccentric orbits (which, as a first step, we keep equatorial). Extension to the gravitational problem is more challenging. The main obstacle is the lack of a formal framework for analyzing Lorenz-gauge metric perturbations in the frequency-domain in Kerr. A potential avenue of approach would be to work with coupled tensorial spherical-harmonics, although this may pose a significant technical challenge. Another possibility would be to develop a suitable tensorial spheroidal-harmonic basis for decomposition in Kerr, akin to the tensorial spherical harmonics that can be used in the Schwarzschild case." }, "1003/1003.4442_arXiv.txt": { "abstract": "We point out that there are regions in the MSSM parameter space which successfully provide a dark matter (DM) annihilation explanation for observed $e^+$ excess (e.g. PAMELA), while still remaining in agreement with all other data sets. Such regions (e.g. the uplifted Higgs region) can realize an enhanced neutralino DM annihilation dominantly into leptons via a Breit-Wigner resonance through the CP-odd Higgs channel. Such regions can give the proper thermal relic DM abundance, and the DM annihilation products are compatible with current antiproton and gamma ray observations. This scenario can succeed without introducing any additional degrees of freedom beyond those already in the MSSM. ", "introduction": "There has been growing interest in the interpretation of an observed positron excess in light of the recent data from PAMELA \\cite{adri1,adri2,pamelapbar} and Fermi-LAT \\cite{fermi1,port1,atw}, while satisfying antiproton and gamma ray constraints. Among the possible sources for the positron excess are astrophysical sources, such as pulsars \\cite{dan09,yuk,profu,maly,maly2,dela10} and supernovae \\cite{shav,sn1,fuji,subir1,subir}, and dark matter \\cite{bal3,gordy1,bal5,cmssm3,dan4,cir,cire5,gordy3,chen5,yasu,ishi,ibarra2,nojiri2,gog,meade,moro2,gao1,eds,mea,gra,paolo5,joe,boem3,chen3,dib,mssm}. The dark matter (DM) possibility is of great interest from the particle theory viewpoint, and we seek in this paper the cosmic and gamma ray signatures of DM annihilation in a supersymmetric model within the MSSM (Minimal Supersymmetric Standard Model) parameter space. Most of the conventionally explored MSSM parameter space cannot successfully explain the observed positron data. There is typically a significant branching ratio of DM annihilation into gauge bosons and Higgs as well as quarks and hence too much hadronic production of antiprotons, in excess of what is observed. The MSSM, however, has more than a hundred parameters and it would be worth seeking such a possibility to realize the current cosmic ray observations in the framework of the MSSM without introducing any additional degrees of freedom. Another persistent problem in the dark matter annihilation scenarios to explain the observed positron excess is the requirement of large annihilation cross sections in the halo far bigger (typically a boost factor of a factor 100 or more) than the canonical thermally averaged cross section value at freeze-out for the weakscale dark matter inferred from the observed dark matter relic density. Several resolutions to this apparent discrepancy in the annihilation cross sections have been proposed, such as unconventional cosmological histories \\cite{paolo4,pall,cate} which can affect the dark matter freeze-out temperature; Sommerfeld enhancement \\cite{hisa1,cire,nima1} which requires new light particles to allow new long range interactions between dark matter particles; Breit-Wigner enhancement \\cite{posp,dan,ibe,guo} which requires a particle whose mass is close to the twice of the dark matter mass; and substructure clumps \\cite{boost2, die2,kuh,alb,brun,mark1} which could provide a partial contribution to the enhancement (say, by a factor of a few). The annihilation enhancements make the the current annihilation cross section higher in galaxies today than it was at the time of freeze-out, so that it can explain the positron excess (which requires a high cross section today) while still satisfying the correct relic density (which requires a lower cross section at freeze-out). The main results of the paper are to point out the existence of parameter regions in the MSSM which can potentially realize a DM annihilation scenario that explains the observed positron excess. We show that there can be dominant leptonic final states and a large boost factor to obtain the positron excess and thermal relic abundance without introducing any additional degrees of freedom beyond those already in the MSSM; in addition these scenarios can be compatible with the current antiproton and gamma ray data. As a concrete example, we consider the following scenario: The requirement of a dominant leptonic final state (in this case taus) can be satisfied in the uplifted Higgs regions \\cite{bog} within the MSSM as described in Section \\ref{mot}. In addition, we obtain the required boost factor from the pseudo-scalar Higgs s-channel resonance in the MSSM which can induce the Breit-Wigner enhancement \\cite{posp,dan,ibe,guo} via $\\chi \\chi \\rightarrow A \\rightarrow \\tau^+ \\tau^-$ for $m_A \\sim 2 m_\\chi$ ($\\chi$ denotes the neutralino dark matter and $A$ denotes the CP-odd Higgs). We also note that we have here found a situation where annihilation to taus satisfies all existing observations. The reason is that we treat the FERMI data as astrophysical background (fit by a simple power law) while we take the PAMELA excess to be due to DM annihilation. Our interpretation is in contrast to some of the previous literature, where the requirement was made of explaining the excesses in both PAMELA and FERMI as due to DM annihilation. In this paper we take the dark matter to have a cored isothermal density profile \\cite{bah,core} in the Galactic Halo unless stated otherwise. This relatively flat core density distribution helps to ameliorate the severe gamma ray constraints \\cite{cirel,pap,kev2} (from observations e.g. by FERMI and HESS) in contrast to other profiles such as the NFW profile \\cite{nfw}. Our choice of profile should be adequate for the purpose of illustrating the potential significance of the previously unexplored MSSM parameter space discussed here. Sec. \\ref{mot} reviews the uplifted supersymmetric Higgs regions in the MSSM parameter space and its unique properties well motivated for the current cosmic ray observations. Sec. \\ref{cosray} then discusses the cosmic and gamma rays signals expected for such regions, followed by the conclusion/discussion in Sec. \\ref{conc}. ", "conclusions": "\\label{conc} We have illustrated that dark matter annihilation scenarios in the MSSM can be viable candidates to explain all current cosmic ray observations without necessarily introducing any additional degrees of freedom. In particular, we have studied pseudoscalar Higgs s-channel resonance in the uplifted Higgs region to obtain boosted leptophilic annihilation cross sections which can explain PAMELA data while not in conflict with any other data sets. We here mention several directions for possible future improvement of our analysis. The positron fraction constraints mainly come from the highest energy bins of the data with large error bars and it should improve by the forthcoming data from PAMELA and AMS-02 with better controls of the systematic/statistical errors. A proper treatment of the charge dependent solar modulation effects (which we did not consider) could account for discrepancies in the low energy data ($E\\lesssim 10$GeV) among different experiments. In addition to the positron fraction which also suffer from the electron background estimation uncertainties, the absolute flux of positrons such as those from AMS-02 would give more definite probe of the underlying physics. Our studies assumed the same value of the boost factor for all species of cosmic rays. However, this is not necessarily the case. In particular, the antiprotons we observe originate from a large region of the halo in contrast to the positrons which are produced locally (say within a few kpc). Hence, in principle, the antiproton flux could be less boosted than the positron flux, for instance if the positrons are boosted partially by the local clumpiness \\cite{boost2, die2,kuh,alb,brun,mark1}. Such relaxation of the antiproton constraints could be important for the annihilation scenarios within the framework of the MSSM, most of which are excluded due to antiproton overproduction. We presented our discussion primarily assuming an isothermal profile for the galactic halo. In fact a more cuspy profile towards the center, such as an NFW profile, was shown to be severely constrained \\cite{cirel,pap,kev2}. Currently there is a great deal of uncertainty regarding the density distribution in the inner regions of galaxies, though there exist observations of flattened cores \\cite{salucci} in spiral galaxies. More precise observations as well as detailed simulations including the proper treatment of the gas physics will be required for more realistic modeling of the halo profiles. Further studies of $\\gamma$-ray constraints beyond the analysis of this paper would be warranted. We refer the reader to, for instance, Ref. \\cite{pap,cirel,kev2,regis} for gamma ray constraints from additional regions of the sky and halo profiles. We focused on gamma rays from the hadronic processes which are characteristic of DM annihilation scenarios in which $\\tau$ channel dominates; the consideration of additional effects such as the inverse Compton scattering including those to the extragalactic gamma rays \\cite{stepha1,danbbb,kawa} could give additional constraints depending on the parameter ranges of interest. More stringent constraints could however come from the particle physics rather than from the astrophysics once we have a concrete particle physics model. For instance, for the uplifted Higgs scenario, the tuning of the heavy Higgs decay width $\\Gamma_A/m_A \\lesssim {\\cal O}(10^{-3}$) implies $y_{\\tau}\\lesssim {\\cal O}(10^{-1})$, which needs to be checked with flavor physics constraints. We also could in principle consider a large $y_{\\mu}$ in an analogous manner to a large $y_{\\tau}$ by adjusting the slepton, squark masses and the phase of the gluon mass \\cite{bog}. If $\\mu$ final states could be significant by such (possibly fine-tuned) adjustments, then the $\\gamma$-ray constraints could be relaxed relative to the $\\tau$ dominant scenarios considered in this paper \\cite{pap,cirel,mea,eds} even though here again the particle physics constraints (e.g. the flavor changing neutral currents) would need to be carefully checked. In the vast MSSM parameter space, it is worth searching for other regions previously missed that could explain all the cosmic ray data. The uplifted Higgs region itself deserves further study. Even though our cosmic ray analysis is based on the properties of the uplifted Higgs regions, we kept our analysis fairly general so that a similar study could be applicable to any other parameter regions with similar properties. The unprecedented wealth of data expected from upcoming astrophysical and terrestrial experiments (e.g. AMS-02 and LHC) could well provide further motivation towards a fuller exploration of the MSSM parameter space in the coming years." }, "1003/1003.3865_arXiv.txt": { "abstract": "Although the gravitational wave kick velocity in the orbital plane of coalescing black holes has been understood for some time, apparently conflicting formulae have been proposed for the dominant out-of-plane kick, each a good fit to different data sets. This is important to resolve because it is only the out-of-plane kicks that can reach more than 500~\\kms and can thus eject merged remnants from galaxies. Using a different ansatz for the out-of-plane kick, we show that we can fit almost all existing data to better than 5\\%. This is good enough for any astrophysical calculation, and shows that the previous apparent conflict was only because the two data sets explored different aspects of the kick parameter space. ", "introduction": "When two black holes spiral together and merge, the gravitational radiation they emit is usually asymmetric and thus the remnant black hole acquires a linear velocity relative to the original center of mass. The speed can reach over 3,000~\\kms \\citep{Dain:2008ck}, which would eject the remnant from any galaxy in the universe (see Figure~2 of \\citet{Merritt:2004xa}). As discussed in \\citet{Merritt:2004xa}, the kick magnitude and distribution are important for discussions of hierarchical merging, supermassive black hole formation, galactic nuclear dynamics, and the degree to which black holes influence galaxy formation. The community has converged on the formula for the kick speed in the original orbital plane \\citep{Baker:2007gi,Campanelli:2007cg,Gonzalez:2006md}, but apparently conflicting dependences for the out-of-plane kick (which dominates the total kick for most configurations) have been proposed. \\cite{Lousto:2008dn} suggested that for a binary with component masses $m_1$ and $m_2\\geq m_1$, the kick scales as $\\eta^2$, where $\\eta\\equiv m_1m_2/(m_1+m_2)^2$ is the \\emph{symmetric mass ratio}; however the data in \\cite{Baker:2008md} were fit much better with an $\\eta^3$ dependence. The conflict is only apparent, however, because there were no runs in common between the two data sets. This suggests that an analysis might be performed with a new ansatz that can fit all of the existing data. Here we perform such a fit, and demonstrate that there is a single formula for the out-of-plane kick that fits almost all existing data to better than 5\\% accuracy. The new ansatz is similar to one recently suggested by \\cite{Lousto:2009mf}, for which, however, a fit was not attempted. Its form is based straightforwardly on the post-Newtonian (PN) approximation and includes both the aforementioned $\\eta^2$ and $\\eta^3$ terms, as well as a slightly more complicated spin-angle dependence. In \\S~2 we describe our new runs and we describe the new ansatz in \\S~3. In \\S~4 we list all 95 runs we have fit, from different numerical relativity groups, and our fitting procedure and best-fit parameters. In \\S~5 we discuss the implications of our results, and indicate the fraction of kicks above 500~\\kms and 1,000~\\kms for representative spin and mass ratio distributions, comparing it with previous results. The goodness of our fit to the entire usable data set of out-of-plane kicks suggests that the full three-dimensional kick is now modelled well enough that it will not limit the accuracy of any astrophysical calculation. ", "conclusions": "One of the most important outputs of kick calculations and fits is the probability distribution of kicks given assumptions about the mass ratio, spin magnitudes, and spin directions. This distribution is critical to studies of hierarchical merging in the early universe (e.g., \\citet{Volonteri:2007et}) as well as to the gas within galaxies \\citep{Devecchi:2008qy} and an evaluation of the prospects for growth of intermediate-mass black holes in globular clusters \\citep{HolleyBockelmann:2007eh}. In Table~2 we show the results of our work (from fit \\#1), compared with the proposed fit formula of \\cite{Campanelli:2007ew}. It is clear that our work gives distributions very close to those of \\cite{Campanelli:2007ew}, with perhaps slightly smaller kicks because of the $\\eta^3$ term we include. \\begin{deluxetable}{lrrr} \\tabletypesize{\\small} \\tablecolumns{4} \\tablewidth{0pt} \\tablecaption{Fraction of kick speeds above a given threshold, compared with the results of \\cite{Campanelli:2007ew} (CLZM). In all cases we assume an isotropic distribution of spin orientations.\\label{tab:kickcomp}} \\tablehead{\\colhead{Mass ratio and spin} & \\colhead{Speed threshold} & \\colhead{CLZM} &\\colhead{This work}} \\startdata $1/10 \\leq q \\leq 1$, $\\alpha_1=\\alpha_2=0.9$ & $v>500 \\kms$ & 0.364$\\pm$0.0048 & 0.342526$\\pm$0.00019\\\\ & $v>1000 \\kms$ & 0.127$\\pm$0.0034 & 0.120974$\\pm$0.00011 \\\\ $1/4 \\leq q \\leq 1$, $\\alpha_1=\\alpha_2=0.9$ & $v>500 \\kms$ & $ 0.699 \\pm 0.0045$ & 0.697818$\\pm$0.00026\\\\ & $v>1000 \\kms$ & 0.364$\\pm$0.0046 & 0.353393$\\pm$0.00019 \\\\ $1/4\\leq q \\leq 1$, $0\\leq \\alpha_1,\\alpha_2\\leq 1$ & $v>500 \\kms$ & 0.428$\\pm$0.0045 & 0.415915$\\pm$0.00020\\\\ & $v>1000 \\kms$ & 0.142$\\pm$0.0034 & 0.134615$\\pm$0.00012\\\\ \\enddata \\end{deluxetable} In summary, we have demonstrated that a modified formula fits all available out-of-plane kicks extremely well. The wide range of mass ratios, spin magnitudes, and angles explores all the major aspects of parameter space for the out-of-plane kicks, and thus we do not expect new results to deviate significantly from our formula. The excellence of these fits suggests that the kick distribution is known to an accuracy that is sufficient for any astrophysical purpose." }, "1003/1003.4990_arXiv.txt": { "abstract": "We study the clustering properties of $z\\sim 5.7$ \\lya emitters (LAEs) in a cosmological reionization simulation with a full \\lya radiative transfer calculation. \\lya radiative transfer substantially modifies the intrinsic \\lya emission properties, compared to observed ones, depending on the density and velocity structure environment around the \\lya-emitting galaxy. This environment-dependent \\lya selection introduces new features in LAE clustering, suppressing (enhancing) the line-of-sight (transverse) density fluctuations and giving rise to scale-dependent galaxy bias. In real space, the contours of the three-dimensional two-point correlation function of LAEs appear to be prominently elongated along the line of sight on large scales, an effect that is opposite to and much stronger than the linear redshift-space distortion effect. The projected two-point correlation function is greatly enhanced in amplitude by a factor of up to a few, compared to the case without the environment-dependent selection effect. The new features in LAE clustering can be understood with a simple, physically motivated model, where \\lya selection depends on matter density, velocity, and their gradients. We discuss the implications and consequences of the effects on galaxy clustering from \\lya selection in interpreting clustering measurements and in constraining cosmology and reionization from LAEs. ", "introduction": "\\lya emitters (hereafter LAEs) are galaxies with strong \\lya emission lines. Owing to the strong \\lya line feature, LAEs can be efficiently detected through narrowband imaging or with integral-field-units (IFU) spectroscopy, which makes them natural targets for searches of high-redshift galaxies. Large samples of high-redshift LAEs are expected with ongoing and forthcoming LAE surveys. In this paper, we investigate the clustering of $z\\sim 5.7$ LAEs in a cosmological reionization simulation, focusing on the effects of \\lya radiative transfer on their clustering properties. Large samples of LAEs would provide exciting opportunities to probe the high redshift universe. The resonance nature of \\lya line makes LAEs a sensitive probe of the high-redshift intergalactic medium (IGM), especially across the reionization epoch (e.g., \\citealt{Miralda98a,Miralda98,Haiman99, Santos04,Haiman05,Wyithe07,Malhotra04,Kashikawa06,Furlanetto06,Dijkstra07a, Dijkstra07b,McQuinn07,Mesinger08,Iliev08,Dayal08,Dayal09,Dayal10}). In particular, the clustering of LAEs at redshift $z>6$ can potentially put tight constraints on the ionization status of the IGM \\citep[e.g.,][]{Furlanetto06,McQuinn07,Iliev08}. Large samples of LAEs should enable accurate measurements of their clustering. As with galaxy clustering in general, the clustering of LAEs encodes useful information of galaxy formation and evolution. We expect to learn about the relation between these young galaxies and dark matter halos, and to obtain insights on the early stage of structure formation. Galaxy clustering also encodes important cosmological information. The fluctuation power spectrum of galaxies is related to that of matter, usually differing by a constant, multiplicative galaxy bias factor on large scales. Because nonlinearity is weaker at higher redshift, the galaxy power spectrum may be used down to smaller spatial scales for constraining cosmology, tightening constraints on parameters like the neutrino mass and providing tests of inflation \\citep[e.g.,][]{Takada06}. The efficient detection of LAEs at high redshift makes them attractive candidates in this endeavor. Large-volume surveys of LAEs, such as the Hobby--Eberly Telescope Dark Energy Experiment (\\citealt{Hill08}), might also enable the detection of the baryon acoustic oscillations feature \\citep[e.g.,][]{Eisenstein05} in the LAE power spectrum. Baryon acoustic oscillations can be used to measure the expansion history of the universe, contributing to constraints on the evolution of dark energy and the curvature of the universe. At present, measurements of LAEs clustering are very limited by small survey volumes and small samples of LAEs. \\citet{Ouchi03} present the angular two-point correlation function (hereafter, 2PCF) of 87 $z=4.86$ LAEs in a 0.15deg$^2$ narrowband survey in the Subaru Deep Field. The correlation length is estimated to be $(3.5\\pm 0.3) \\hMpc$ (and $(6.2\\pm 0.5)\\hMpc$ if a maximum contamination correction is applied), larger than that of $z\\sim 4$ Lyman break galaxies. With a set of LAEs found in a larger area ($\\sim$0.3 deg$^2$) around the same field, \\citet{Shimasaku04} also report strong clustering for 41 $z=4.86$ LAEs, but they find almost no clustering for 51 $z=4.79$ LAEs. The strong clustering of $z=4.86$ LAEs is difficult to reproduce with a simple model that relates LAEs to dark matter halos \\citep{Hamana04}. With 151 $z=4.5$ LAEs in a 0.36deg$^2$ field of the narrowband Large Area Lyman Alpha (LALA) survey, \\citet{Kovac07} estimate a correlation length of $3.2\\pm 0.4\\hMpc$ ($4.6\\pm 0.6\\hMpc$ with contamination correction). The clustering can be reproduced if these LAEs reside in halos more massive than (1--2)$\\times 10^{11}\\hMsun$. \\citet{Gawiser07} measure the angular 2PCF of 162 $z=3.1$ LAEs discovered in a 0.28deg$^2$ field of the MUltiwavelength Survey by Yale-Chile (MUSYC). They find a moderate clustering with a correlation length of $2.5_{-0.7}^{+0.6}\\hMpc$, corresponding to that of halos with a minimum mass of $\\sim$2.8$\\times 10^{10}\\hMsun$ (and a median mass of $5.6\\times 10^{10}\\hMsun$). The 261 $z=2.1$ LAEs in the same MUSYC field have a correlation length of $3.2\\pm 0.6\\hMpc$ \\citep{Guaita10}, corresponding to a median halo mass of $1.8\\times 10^{11}\\hMsun$. From the clustering of $z=3$--7 LAEs in the 1deg$^2$ Subaru/{\\it XMM-Newton} Deep Survey (SXDS), \\citet{Ouchi10} infer that the average host halo mass is $10^{10}$--$10^{11}M_\\odot$. Understanding the LAE clustering results needs to take into account the differences in LAE samples and redshifts, as well as sample variance caused by small survey volumes. The development of a physically based theoretical model of LAEs and their clustering is also needed. In models of LAEs, the \\lya flux of an LAE is usually computed from the emitted ionizing photons in the galaxy residing in a dark matter halo, assuming Case B recombination \\citep{Osterbrock89}. The theoretical modeling of LAE clustering depends on how LAEs and dark matter halos are connected, and how the observed \\lya luminosity is determined from the intrinsic \\lya luminosity. Generally speaking, there are two scenarios considered in LAE models, the duty cycle scenario and the \\lya escape fraction scenario. In the duty cycle scenario, LAEs are short-lived and at any given time only a fraction of all galaxies are active as LAEs. In the \\lya escape fraction scenario, it is assumed that only a fraction of \\lya photons can escape from the source, and therefore the observed \\lya luminosity is a fraction of the intrinsic one. Either scenario can make the predicted \\lya luminosity function (LF) match the observation. For a given number density of LAEs, the masses of host halos in the duty cycle scenario would be on average lower than those in the escape fraction scenario. As a consequence, the clustering of LAEs would be different in the two scenarios, with a stronger clustering in the escape fraction scenario. \\citet{Nagamine10} predict LAE clustering based on cosmological smoothed particle hydrodynamic (SPH) simulation and consider both scenarios. They find that LAE clustering measurements from observations are in favor of their duty cycle scenario. \\citet{Tilvi09} present an LAE model in which \\lya luminosity or star formation rate (SFR) is related to the halo mass accretion rate, rather than halo mass, and the model naturally gives rise to the duty cycle of LAEs. Their model predict correlation lengths of LAEs in agreement with the observations. \\citet{Orsi08} combine a semi-analytic model of galaxy formation with a large $N$-body simulation to predict the clustering of LAEs. They adopt the scenario of \\lya escape fraction and assume the escape fraction to be constant (2\\%) and independent of galaxy properties. By accounting for the large sample variance, the model is found to reproduce the angular clustering measurements from current surveys of LAEs. \\citet{McQuinn07} develop a model of LAEs using reionization simulations with cosmological volume (see also \\citealt{Iliev08}) and discuss the effect of reionization on LAE clustering. Their model computes the escape fraction based on a simplified treatment of \\lya radiative transfer (hereafter, RT) that consists of multiplying the intrinsic line profile by $\\exp(-\\tau_\\nu)$, where $\\tau_\\nu$ is the optical depth at frequency $\\nu$ along the line of sight. The $\\exp(-\\tau_\\nu)$ model is reasonably accurate in the case of studying the absorption feature of \\lya photons passing through some neutral regions along the line of sight, especially for absorption caused by the damping wing or small optical depth. For studying the observed \\lya emission, the model is not accurate. Even if one limits the $\\exp(-\\tau_\\nu)$ model to study the transfer outside of a radius much larger than halo size by assuming a \\lya line profile at that radius, it is not clear what radius to use, what line profile to assume, and what angular distribution of \\lya emission at that radius to adopt. The model neglects the spatial and frequency diffusion of \\lya photons, which is important because of the scattering, not the absorption nature of the transfer of \\lya photons. \\citet{Zheng10} (hereafter Paper I) present a simple physical model of LAEs, where \\lya RT is the primary physical process transforming intrinsic \\lya emission properties to observed ones. For the first time, a full RT calculation of \\lya photons \\citep{Zheng02} in gas halos around LAEs is performed in a self-consistent fashion with the radiation-hydrodynamic reionization simulations \\citep{Trac08}. While in this model the number of \\lya photons in the IGM is correctly conserved, only a fraction of them can be observed, those included in the central part of the extended \\lya emission with high enough surface brightness. The model predicts a broad distribution of apparent (observed) \\lya luminosity at fixed intrinsic \\lya luminosity or ultraviolet (UV) luminosity, a consequence of a variable intergalactic environment of LAEs and the environment-dependent RT of \\lya photons. Therefore, the model predicts an effective \\lya escape fraction that is not constant, but has a broad distribution and is correlated with the environment. This simple physical model is able to explain an array of observed properties of $z\\sim$5.7 LAEs in \\citet{Ouchi08}, including \\lya spectra, morphology, and apparent \\lya LF. The broad distribution of apparent \\lya luminosity at fixed UV luminosity provides a natural explanation for the observed UV LF, especially the turnover toward the low-luminosity end. The model also reproduces the observed distribution of \\lya equivalent width (EW) and explains the deficit of UV bright, high-EW sources. In this paper, we investigate the clustering of LAEs within the model presented in Paper I. As we will show, the environment-dependent \\lya RT introduces new and significant effects in the clustering of galaxies selected by \\lya emission, a real physical effect that has not been properly taken into account in previous studies. We first present the environment dependence of the \\lya selection and the dependence on halo mass in Section~\\ref{sec:environ}. In Section~\\ref{sec:clustering}, we present the results of LAE clustering from our model, in terms of the 2PCFs. Following an intuitive interpretation of the features seen in LAE clustering, we provide a simple physical model to further aid our understanding of LAE clustering. In Section~\\ref{sec:hod}, we show the environment dependence of the halo occupation distribution (HOD) of LAEs. In Section~5, we summarize our main findings and discuss the implications. In the appendices, we provide an extended simple physical model of LAE clustering, present the power spectrum of LAEs, make comparisons to the LAE clustering in the $\\exp(-\\tau_\\nu)$ model, and present tests on factors that may mask the new clustering effects. Throughout the paper, we adopt the same cosmological model as in the reionization simulation \\citep{Trac08} used in our RT calculation. It is a spatially flat $\\Lambda$CDM cosmological model with Gaussian initial density fluctuations, and the cosmological parameters are consistent with the {\\it Wilkinson Microwave Anisotropy Probe} 5 year data \\citep{Dunkley09}: $\\Omega_m=0.28$, $\\Omega_\\Lambda=0.72$, $\\Omega_b=0.046$, $h=0.70$, $n_s=0.96$, and $\\sigma_8=0.82$. Our \\lya RT calculation is based on the $z=5.7$ output of the simulation, which has a box size of $100\\hMpc$ on a side. In our calculation, a 768$^3$ grid is used to represent the neutral hydrogen density, temperature, and peculiar velocity fields in the simulation box. The Hubble flow is added to the velocity field. LAEs are assumed to reside in dark matter halos with positions and velocities from the halo catalog. To reduce source blending in the \\lya image and spectra, \\lya photons are collected with a finer spatial resolution, a $6144^2$ grid for the image of the whole box, corresponding to 16.3$\\hkpc$ (comoving) or 0.58\\arcsec per pixel. The spectral resolution and range are 0.1\\AA (25$\\kms$) and 24\\AA in rest frame, respectively. We divide the whole simulation box into three layers so that the depth of each layer approximates that from the width of the narrowband filter used in searching for $z\\sim 5.7$ LAEs \\citep{Ouchi08}. The calculation result for each layer is saved in an IFU-like datacube of dimension 6144$\\times$6144$\\times$240. We refer the readers to Paper I for more details about the characteristics of the simulation and calculation. \\begin{figure*} \\plotone{f1.ps} \\caption[]{ \\label{fig:denvelgrad_mass} Dependence of \\lya flux suppression of LAEs on density and peculiar velocity, as a function of halo mass. The suppression is characterized by the ratio of the apparent (observed) and intrinsic \\lya luminosity $L_{\\rm apparent}/L_{\\rm intrinsic}$. (a) Dependence on the smoothed overdensity field at the source position. The overdensity field is smoothed with a 3D top-hat filter of radius 2$\\hMpc$ (comoving). (b) Dependence on the density gradient along the $Z$-direction. The derivative is with respect to comoving coordinate. (c) Dependence on the host halo velocity. (d) Dependence on the linear peculiar velocity gradient along the $Z$-direction. The linear peculiar velocity is obtained from the smoothed overdensity field based on the continuity equation (see the text for details). The velocity gradient is put in units of the Hubble parameter. Different colors are for LAE host halos of different masses, as labeled in panel~($d$). The median of the ratio is plotted as a solid curve. The two dotted curves delineate the upper and the lower quartiles, and for clarity we only plot those for the lowest mass range. Note that the line-of-sight direction (from the observer to sources) is along the $-Z$-direction, which matters for interpreting the results in panels~($b$) and ($c$). See Section~\\ref{sec:environ} for further discussion. } \\end{figure*} \\begin{figure*} \\epsscale{1.1} \\plottwo{f2a.ps}{f2b.ps} \\caption[]{ \\label{fig:jointdep} Joint dependence of \\lya flux suppression of LAEs on environments in the full \\lya RT model (left) and the $\\exp(-\\tau_\\nu)$ model (right). In each panel, the gray scale shows the distribution of halos ($\\sim10^{10}\\hMsun$) in the plane of density and line-of-sight velocity gradient, darker for higher probability density. Contours indicate the median \\lya flux suppression, defined as the ratio of observed to intrinsic \\lya luminosity, for LAEs residing in these halos. Thicker contours correspond to higher ratios, and adjacent contours differ by 0.1 dex in contour levels. Note the difference in the trend of flux ratios at fixed velocity gradient in the two models. See the text for details. } \\epsscale{1.0} \\end{figure*} \\begin{figure} \\epsscale{1.2} \\plotone{f3.ps} \\caption[]{ \\label{fig:schematic} Schematic illustration of RT effect on the observed flux of \\lya emission. In each panel, \\lya photons are emitted from the central source and scattered by neutral gas surrounding it. The gas is shown as the shaded region, and the gray scale indicates the \\lya optical depth along each radial direction with darker for higher optical depth. The flux of escaped \\lya photons along each direction is illustrated by the arrows with length proportional to flux. Although the optical depths along the line-of-sight direction (from the observer to the source) are the same in all four panels, we expect to have different observed flux because of different distribution of optical depth in other directions. } \\epsscale{1.0} \\end{figure} \\begin{figure*} \\epsscale{1.15} \\plotone{f4.ps} \\epsscale{1.0} \\caption[]{ \\label{fig:slice} Spatial distribution of LAEs. The slice is $5\\hMpc$ thick. The gray scale map shows the matter density distribution (smoothed at 2$\\hMpc$ scales; darker for higher density) to delineate the large-scale filamentary structures. Points represent halos above $5\\times 10^9 \\hMsun$. Left: distribution of LAEs as a function of the ratio of the apparent (observed) and intrinsic \\lya luminosity $L_{\\rm apparent}/L_{\\rm intrinsic}$. Red (blue) points are for sources with \\lya luminosity weakly (strongly) suppressed, with the apparent to intrinsic luminosity ratio higher than 14\\% (lower than 0.6\\%), which approximately corresponds to the top (bottom) 10\\% of the luminosity ratio distribution. Right: distribution of LAEs (red points) above a threshold in apparent (observed) luminosity. The luminosity threshold corresponds to a LAE number density of $3.8\\times 10^{-3}\\denhMpc$. Note that the line-of-sight direction (from the observer to sources) is along the $-Z$-direction. That is, the distant observer observes the sources from the top of the panels, which matters for interpreting the relation between \\lya flux suppression and the local environment (see the text). } \\end{figure*} ", "conclusions": "\\label{sec:discussion} We investigate the clustering of LAEs, galaxies that are selected by their \\lya emission, within a physical model that fully accounts for the \\lya RT. Our model of LAEs combines radiation-hydrodynamic cosmological reionization simulations with Monte Carlo RT for \\lya photons. It is a simple model, which assumes that RT is the single factor in transforming intrinsic \\lya emission properties to observed ones. As previously shown in Paper I, the simple model is able to explain an array of observational properties of $z=5.7$ LAEs. The model has strong predictive power, and we predict the clustering properties of LAE clustering in this paper. \\lya RT depends on the circum-galactic and inter-galactic environments of \\lya emitting galaxies. \\lya photons emitted from a source at a halo center see a complex density and velocity structure and an anisotropic optical depth distribution. \\lya photons tend to find the easiest way out, and the observed \\lya emission is usually anisotropic. {\\it We emphasize that the observed \\lya flux at a given line of sight is not purely determined by the line-of-sight optical depth, but the line-of-sight optical depth relative to those in all other directions (see Section~\\ref{sec:environ}).} Because of the resonance nature of the \\lya line and the large scattering cross-section, the observed \\lya emission is highly sensitive to the local environments, including the matter density, line-of-sight velocity, their line-of-sight gradients, and the velocity gradients in the transverse directions (not all of these environmental variables are independent). It is this strong coupling between observed \\lya emission and environments that gives rise to new effects in LAE clustering that need to be taken into account in interpreting the clustering of observed LAE samples. The overall effects in clustering caused by the \\lya selection are anisotropic clustering and scale-dependent galaxy bias. For density fluctuations along the line of sight, \\lya selection leads to a higher probability in detecting LAE sources in underdense regions than in overdense regions, so it suppresses the line-of-sight density fluctuations. For density fluctuations perpendicular to the line of sight, the anisotropic distribution of \\lya emission makes sources in overdense regions preferentially selected, hence the transverse density fluctuations are enhanced. Roughly speaking, filamentary or planar structures tend to be preferentially selected when they are parallel to the line of sight. The suppression of line-of-sight fluctuations and the enhancement of transverse fluctuations create anisotropy in LAE clustering. The iso-contour curves in the 3D 2PCFs, which reflect LAE pair counts as a function of line-of-sight and transverse separations, show a distinct elongation pattern along the line of sight. The elongation appears on all scales where we have reliable measurements of the 2PCFs, ranging from sub-Mpc to over 10 Mpc. The anisotropic pattern is opposite to the linear redshift-space distortion effect (Kaiser effect), which makes contours squashed along the line of sight. We emphasize that the anisotropy caused by \\lya selection originates in real space. At $z=5.7$, the cases we consider in this paper, the Ly$\\alpha$-selection-induced distortion in LAE clustering is much stronger than the linear redshift-space distortion. Therefore, even in redshift space the elongation pattern along the line of sight in the 3D 2PCFs is well preserved. We note that the elongation also differs from the FoG effect seen in galaxy clustering, which is caused by random motions of galaxies in virialized structures. While the FoG effect only shows up in redshift space and on small scales (e.g., $\\lesssim$ 1Mpc), the \\lya selection effect can appear in both real and redshift spaces and on much larger scales. The \\lya RT induced features are not ``Fingers of God'' but ``Arms of God''. The anisotropic clustering induced by \\lya selection is a completely new phenomenon in galaxy clustering. Other than the usual redshift-space distortion (FoG and Kaiser effect), there are other forms of anisotropic clustering discussed in the literature. \\citet{Padilla05} study the cross-correlation between voids and galaxies. In the redshift-space two-point cross-correlation function of voids and galaxies, an elongation pattern along the line of sight is found on large scales (up to a few times the radius of voids in the sample). Unlike the \\lya selection effect, which is a real-space effect, the phenomenon in the void-galaxy cross-correlation is a redshift-space effect. It simply reflects that galaxies tend to stream out of void regions. The closest analogy to the \\lya selection effect we study is the effect of orientation-dependent galaxy selection investigated by \\citet{Hirata09}, in the case that galaxies are aligned by large scale tidal fields. Both effects can be attributed to an environment-dependent surface brightness selection and act in both real and redshift spaces. However, the \\lya RT selection effect is much stronger. While the \\citeauthor{Hirata09} effect may change the clustering at a level of a few percent, the \\lya selection can change the clustering amplitude by a factor of a few, owing to the high sensitivity of \\lya RT to local environments of galaxies. In addition, unlike the surface brightness distribution of stars in galaxies, which largely has a parity symmetry (i.e., similar surface brightness if viewed at opposite directions), the anisotropic distribution of \\lya surface brightness does not necessarily have any symmetry. The 3D anisotropic clustering of LAEs, if measured, will be a strong test to the RT model of LAEs. To achieve this goal, a large spectroscopic sample of LAEs will be needed. With current narrow-band surveys of LAEs, however, the clustering measurements are limited to the angular 2PCFs, which resemble the projected 2PCFs. Unlike the redshift-space distortion effect, which is largely eliminated in the projected 2PCFs, we find that the \\lya selection effect is imprinted in the projected 2PCFs. Projection keeps the transverse fluctuations, which are always enhanced by the \\lya selection. The projected 2PCF of LAEs has a higher amplitude than that of the control LAE sample with environment effect removed. Our model makes the following distinctive prediction (see Figure~\\ref{fig:biasfactor}): {\\it the amplitude of the LAE 2PCF has a very weak dependence on the observed \\lya luminosity.} The prediction breaks down for faint LAEs if a large dispersion (1 dex) between intrinsic \\lya luminosity and halo mass is introduced, which may result from stochastic star formation, but it remains valid for luminous LAEs (see Appendix~\\ref{sec:effect_lineprof}). This is testable with large narrow-band surveys. The relation between LAEs and underlying halos has interesting properties. For LAE samples defined by a threshold in observed \\lya luminosity, faint LAEs appear to be more strongly clustered than mass threshold samples of dark matter halos having the same number density. The trend is reversed for very bright LAEs (with number density lower than $2\\times 10^{-4}\\denhMpc$), where halos are more strongly clustered than LAEs of the same number density. If an LAE sample with low luminosity threshold were simply related to a mass threshold sample of halos and one used the observed LAE clustering amplitude to infer the halo number density, one would reach the conclusion that there are multiple LAEs per halo. This is not clearly seen from currently available clustering measurements yet, probably because of the small sample sizes (hence large variance) and the existence of contamination from low-redshift objects. The Ly$\\alpha$-selection-induced anisotropic clustering can be largely understood by accounting for the dependences of the selection on density and line-of-sight velocity gradient. Besides these dependences, \\lya selection is also related to line-of-sight peculiar velocity and line-of-sight density gradient, which can give rise to scale-dependent galaxy bias, changing the shape of the fluctuation power spectrum with respect to matter (Appendices~\\ref{sec:extendedtoy} and \\ref{sec:powerspectrum}). We develop a simple physical model that incorporates all the identified environment variables in the selection to aid the understanding and interpretation of LAE clustering (Appendix~\\ref{sec:extendedtoy}). Although the model is able to provide qualitative explanations to features in the LAE clustering, it is an over-simplified model, which assumes linear dependence on environment variables. A more sophisticated model will have to rely on a detailed study of the dependence of \\lya RT on environments and an accurate description of the statistical relation between observed \\lya emission and environmental variables. As a powerful tool, the HOD framework has been successfully applied to interpret galaxy clustering data from many galaxy surveys. For LAEs, the usual assumption in the HOD framework that galaxy properties are only a function of halo mass and are independent of large-scale environments breaks down, as a result of the selection caused by environment dependent \\lya RT. As shown in Section~4, the HOD of LAEs is strongly dependent on environments. \\citet{Hamana04} perform (environment-independent) HOD modeling of the angular 2PCF of $z=4.86$ LAEs in one Subaru Deep Field and find that the observed strong clustering cannot be reproduced by the model. Given the large sample variance in the small field and crude error estimate of the clustering measurement, however, it is not clear whether the failure of the model in this case is caused by the neglecting of the \\lya selection effect. Our study shows that to correctly model LAE clustering within the HOD framework, one has to extend the framework to incorporate the strong environment dependence of \\lya selection. As discussed in Paper I, the uncertainties in our current model lie in the intrinsic \\lya luminosity and spectra, which deserve detailed investigation. By ``intrinsic'', we mean the properties of \\lya photons after escaping the ISM. The intrinsic \\lya luminosity is related to the SFR and the initial mass function of stars. If the model (luminosity-threshold) LAEs are matched to the observed LAEs in terms of the number density, the uncertainty in the intrinsic \\lya luminosity is largely removed in studying LAE clustering. On the other hand, the intrinsic spectrum of \\lya emission, especially the intrinsic line width, is an important factor in determining the strength of the \\lya selection effect. A larger intrinsic line width would lead to a weaker dependence of \\lya RT on environments (Paper I), which in turn would make the effects of \\lya RT on LAE clustering weaker. For example, we would see a less elongated pattern in the 3D 2PCF of LAEs with a larger intrinsic line width. The intrinsic line shape may also be modified by galactic winds, which can also affects the \\lya RT (e.g., \\citealt{Kunth98,Atek08}) and the effect on clustering. If isotropic galactic winds shift the initial \\lya line by a few hundred $\\kms$, the coupling of observed \\lya emission to the environment can be weakened. However, galactic winds usually display a collimated bipolar pattern, and photons escaping in directions other than the bipolar direction may not achieve a large shift. High-resolution simulation of individual galaxies with galactic wind included is useful for further studying the effect of wind on the observational properties of LAEs. See the tests and discussions in Appendix~\\ref{sec:effect_lineprof}. Consequently, if galactic wind can have a strong effect, the strength of Ly$\\alpha$-selection-induced clustering effects would allow us to potentially constrain the intrinsic \\lya line width and shape of LAEs, which would otherwise be difficult to discern. For low-mass halos ($\\sim 10^{10}\\hMsun$) in the simulation, the grid for hydrodynamic calculation marginally resolves the virial radius. The \\lya RT for sources in these halos may be limited by the resolution. However, the infall region (inside the turnaround radius), which is about 5.6 times larger and is important in shaping the \\lya RT, is well resolved. Therefore, the \\lya RT for sources in these low-mass halos is not expected to suffer the resolution effect significantly. \\lya RT is a physical process that must exist around LAEs. The new effects on LAE clustering, including the enhancement in the projected 2PCFs, are strong, which means that they cannot be easily masked by other effects. If observation showed a null detection of the effects, e.g., finding no enhancement in the angular 2PCFs, it would have important implications in our study of LAEs. A substantial scatter in the intrinsic \\lya emission properties at fixed halo mass remains as a possibility to reduce or even mask the \\lya RT selection effect. In our model, the intrinsic \\lya emission properties, especially the \\lya luminosity, are tightly correlated with halo mass. The \\lya luminosity is based on the SFR averaged over 10Myr time scale and the star formation prescription does not lead to significant scatter in the SFR at fixed halo mass and redshift \\citep{Trac07}. To have a large scatter in the instantaneous SFR, one needs to introduce a broad distribution of the star formation efficiency or make the star formation stochastic in halos of fixed mass. Dust in the ISM could further enlarge the scatter in the \\lya luminosity. If in the end the intrinsic \\lya luminosity distribution at fixed halo mass were much broader than that from \\lya RT effect (Paper I), the \\lya RT selection could be masked. In Appendix~\\ref{sec:effect_lineprof}, we present simple tests on the effect of the dispersion in the relation between intrinsic \\lya luminosity and halo mass. The dispersion can change the clustering amplitude of LAEs of fixed number density by including LAEs residing in lower mass halos. However, the anisotropic clustering pattern persists even if a large dispersion (1 dex in luminosity) is introduced. A broad or stochastic \\lya luminosity distribution would give rise to an effective duty cycle such that only a tiny fraction of galaxies are in the \\lya emitting phase at a given time, which means that at a given number density LAEs reside in halos of much lower mass, compared to the case of a narrow intrinsic \\lya luminosity distribution to start with. \\citet{Nagamine10} argue that a duty cycle scenario provides a reasonable explanation to existing LAE clustering measurements. To fully address the magnitude of the distribution and stochasticity of star formation efficiency, the relevant processes have to be incorporated in the reionization simulation, which is limited by our understanding of baryon physics. Potential consequences, when combined with \\lya RT selection, on LAE clustering deserves detail investigations. Strong tests to the model and constraints on different processes are expected to come from the \\lya LF, UV LF, and clustering of LAEs. The completely new effects on LAE clustering due to \\lya RT add another layer of complexity in modeling their clustering. This means that the existence of \\lya selection effect would complicate the inference of cosmological parameters from LAE clustering in large LAE surveys. To map from galaxy clustering to dark matter clustering, which directly encodes the cosmological information, the effects that are commonly considered in current galaxy clustering analysis include the redshift-space distortion, the nonlinear evolution of structure, and the scale-dependent bias induced by halo biasing. For \\lya selected galaxies, the new effects that add to the above list are the anisotropic clustering (opposite to and stronger than the redshift-space distortion) and the scale-dependent bias induced by environment-dependent \\lya RT. For our model with $z=5.7$ LAEs, the Ly$\\alpha$-selection-induced anisotropy is overwhelming. At lower redshifts $z=2-3$ we expect the \\lya RT induced selection effect to continue to operate in the regions surrounding galaxies, as at $z=5.7$ shown here. However, at present, it is not clear what the strength of the \\lya selection effect will be and up to what scales the clustering of LAEs is affected at these redshifts. Several competing factors prevents us from having a simple guess. On one hand, lower matter density and higher UV background at $z=2-3$ (hence lower neutral hydrogen fraction in the scattering regions), compared to $z=5.7$, permit easier escape of \\lya photons through the circum-galactic and inter-galactic gas. On the other hand, a lower Hubble velocity may provide a countering factor. We reserve a more detailed investigation of clustering of LAEs at $z=2-3$ for a future work. The strong environmental dependence of the \\lya selection, on the other hand, provides a sensitive way to probe the late stage of cosmological reionization. With simple treatments of \\lya RT, it has been shown that \\lya emission from LAEs can be used to probe the mean hydrogen neutral fraction as a function of redshift (e.g., from the \\lya LF; \\citealt{Malhotra04, Haiman05}). Both \\citet{Furlanetto06} and \\citet{McQuinn07} show that the clustering of LAEs is enhanced by the patchy reionization. Our current work is limited to $z=5.7$, when the reionization is completed, and we plan to perform full RT modeling of LAEs at higher redshifts in a subsequent paper. Because of the fast evolution of neutral hydrogen density as reionization proceeds, we would see a rapid change in the distribution of the observed-to-intrinsic \\lya luminosity ratio, which would lead to a rapid change in the clustering amplitude of LAEs relative to the underlying population \\citep{Furlanetto06}. The isolated \\ion{H}{2} bubbles introduce a characteristic scale, which would result in a scale dependence in the bias factor \\citep{Furlanetto06}. We also expect that isolated \\ion{H}{2} bubbles and large-scale fluctuations of photoionization rate may introduce additional effects in the 3D anisotropic clustering of LAEs. As a whole, we expect that different features in the clustering of LAEs, if detected, would provide a wealth of information about reionization." }, "1003/1003.1051_arXiv.txt": { "abstract": "{Cold fronts have been observed in a large number of galaxy clusters. Understanding their nature and origin is of primary importance for the investigation of the internal dynamics of clusters.} {To gain insight on the nature of these features, we carry out a statistical investigation of their occurrence in a sample of galaxy clusters observed with XMM-Newton and we correlate their presence with different cluster properties. } {We have selected a sample of 45 clusters starting from the B55 flux limited sample by Edge et al. (1990) and performed a systematic search of cold fronts.} {We find that a large fraction of clusters host at least one cold front. Cold fronts are easily detected in all systems that are manifestly undergoing a merger event in the plane of the sky while the presence of such features in the remaining clusters is related to the presence of a steep entropy gradient, in agreement with theoretical expectations. Assuming that cold fronts in cool core clusters are triggered by minor merger events, we estimate a minimum of 1/3 merging events per halo per Gyr.} {} ", "introduction": "The unprecedented angular resolution of the X-ray telescope {\\it Chandra} led to the discovery of several new phenomena within various astrophysical systems. One of these are cold fronts detected in galaxy clusters. Initially observed in merging clusters, the prototypes are found in A2142 \\citep{Maxim:2000}, A3667 \\citep{Vikhlinin1:2001, Vikhlinin2:2001, Vikhlinin:2002} and 1E0657-56 \\citep{Maxim:2002}. All these systems feature very sharp discontinuities in their X-ray images where the drop of the surface brightness (and correspondingly of the gas density) is accompanied by a jump in the gas temperature, with the denser region colder than the more rarefied region, unlike shock fronts. For this reason, these features have been dubbed ``cold fronts'' \\citep{Vikhlinin1:2001}. The density and the temperature discontinuities have similar amplitude so that pressure is approximatively continuous across the front. Cold fronts have been initially interpreted as the edge of the cool core of a merging substructure which has survived the merger and is rapidly moving through the ambient gas \\citep{Maxim:2000}. Cold fronts have successively been detected in the cores of some relaxed clusters (e.g. A1795: \\citealp{Maxim:2001}; RX J1720.1+2638: \\citealp{Mazzotta:2001}; A496: \\citealp{Dupke:2003}; 2A 0335+096: \\citealp{Mazzotta:2003}) and to date a large number of relaxed systems are known to host one. Since the presence of cold fronts in cool cores provides evidence of gas motions and possibly of departures from hydrostatic equilibrium, understanding the nature of such a widespread phenomenon is mandatory to characterize the dynamics of galaxy clusters. High resolution hydrodynamical simulations are, at present, the main technique to investigate the mechanisms generating cold fronts. Indeed, cold front features could already be detected in simulations published prior to the launch of {\\it Chandra} \\citep{Roettiger:1997, Roettiger:1998}. After cold fronts discovery, several hydrodynamical simulations have been developed to model the effect of the ram-pressure stripping in a merger event and the formation of the cold front feature in merging clusters \\citep{Heinz:2003, Nagai:2003,Mathis:2005}. Several simulations have also been employed to understand the origin of cold fronts in relaxed non-merging clusters \\citep[e.g.][]{Churazov:2003, TH:2005, AM06}. The emerging picture (\\citealp{AM06}; see also \\citealp{MM_review:2007} for a review) is that cold fronts arise during major merging events through ram-pressure stripping mechanisms which induce the discontinuity among the merging dense subcluster and the less dense surrounding ICM. In relaxed clusters, the cold fronts features are induced by minor merger events which produce a disturbance on the gas in the core, displace it from the center of the potential well and decouple it from the underlying dark matter through ram-pressure. Subsequently, a sloshing mechanism sets in, generating cold fronts. The necessary condition for triggering this mechanism is the presence of a steep entropy profile for the central gas which is generally fulfilled at the center of relaxed cool core clusters. Cold fronts are at present observed in a large number of galaxy clusters. \\citet{Maxim:2003} analyzed a sample of 37 relaxed clusters observed with Chandra showing that cold fronts are present in the majority of the cores of relaxed clusters. Recently, \\citet{Owers:2009} characterized a sample of nine cold fronts with quantitative measurements of the thermodynamic discontinuities across the edges and associated the presence of a cold front with evidence of merger activity.\\\\ While many objects have been studied in detail to understand the nature of cold fronts, we still lack a systematic investigation of the characteristics of these phenomena and of their host clusters through a large sample. The aim of this paper is to perform a systematic search of cold fronts in a representative sample and to investigate the properties of their parent clusters. Such a study is necessary to inspect the nature and origin of cold fronts and eventually to test the reliability of the picture emerging from the simulations. The sample is selected starting from the B55 flux limited sample by \\citet{Edge:B55}. We use for our analysis {\\it XMM-Newton} data. In spite of its limited spatial resolution with respect to {\\it Chandra}, {\\it XMM-Newton} has the positive attribute of having a large field of view, allowing a significant coverage of most of the clusters. In most cases, the clusters are inside the EPIC field of view up to a radius $\\simgt 0.3 r_{180}$, allowing the characterization of the main thermodynamical properties well beyond the core regions. Additionally, the {\\it XMM-Newton} large collecting area allows a good statistics for a large number of objects. Among the several physical properties characterizing the intracluster medium, we focus our attention on entropy. Entropy plays a key role in describing the thermodynamical state of the ICM, its distribution is a signature of the thermodynamical history of the cluster and it is also intimately related to the non-gravitational processes which may have occurred \\citep{Voit_entropy:2002, Voit_entropy:2005}. Moreover, as previously stressed, simulations highlight how the steep gas entropy profile is a necessary condition for the onset of the sloshing mechanism and therefore for the presence of cold fronts in cool core clusters \\citep{AM06}. The structure of the paper is the following. In \\S\\ \\ref{sec:CF_sample} we describe the sample of clusters that we have analyzed and in \\S\\ \\ref{sec:data_red} we provide details about the data reduction. Then we describe (see \\S\\ \\ref{sec:search_CF}) the algorithm used for the systematic search of cold fronts in the cluster sample. We present our results about the occurrence and the origin of cold fronts in \\S\\ \\ref{sec:occurr} and we discuss them in \\S\\ \\ref{sec:disc}. We summarize our findings in \\S\\ \\ref{sec:summary}. We adopt a $\\Lambda$CDM cosmology with $\\Omega_{\\rm{m}} = 0.3$, $\\Omega_\\Lambda = 0.7$, and $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$. ", "conclusions": "\\label{sec:disc} The origin of cold fronts in clusters manifestly undergoing a merger event can be related to the motion of a dense cold cloud of gas within the atmosphere of another subcluster. Conversely, the presence (or the absence) of these features in the subsample of non-merging clusters and clusters where the merger is not close to the plane of the sky is not clearly understood. Fig. \\ref{fig:entropy_shade} provides some hints to help understand what determines the presence of cold fronts in these systems. The general picture emerging here is that the entropy profile discriminates among the two classes (CF and NCF) of clusters. While at large radii the (scaled) entropy profiles of these clusters are very similar, in the innermost regions ($r \\simlt 0.08r_{180}$) their behaviors differ. Our finding of a steep entropy gradient in CF cluster is in agreement with theoretical expectations. Indeed, simulations by \\citet{AM06} show that cold fronts can rise and develop in the cores of clusters if the entropy sharply decreases towards the center (as typically occurs in the center of cool core clusters). According to these simulations, cold fronts develop as a consequence of minor merger events; during its passage near the center of a cluster, a merging subclump induces some disturbance on the low entropy gas of the core and displaces it from the center of the potential well. If the entropy profile is steep, the cool gas starts sinking towards the minimum of the gravitational potential, a sloshing mechanism sets in and cold fronts arise. If the entropy profile is not steep, the entropy contrast is insufficient for the cool gas to flow back and for the sloshing mechanism to set in. In agreement with this picture, we find that cold fronts form only in regions where the entropy profile sharply decreases. In Fig. \\ref{fig:entropy_shade}, we plot (blue filled circles) the cold fronts positions measured for the non-merging clusters of our sample. Excluding the outermost cold front of A496 (this cluster hosts three cold fronts) which lies at a distance of $\\sim 0.08 r_{180}$ from the peak, all the cold fronts we detect lie at distances smaller than $\\sim 0.05 r_{180}$, where entropy profiles steepen, and greater than $\\sim 0.01r_{180}$ where, in many systems, {\\it Chandra} detects a flattening \\citep{Donahue_entropy:2006}. \\begin{figure} \\centering \\includegraphics[angle=0,width=9.5 truecm] {mean_prof_entropy_wr_fit_065_pc.ps} \\caption{Mean scaled entropy profiles as in Fig. \\ref{fig:entropy_shade}. Red dot-dashed curve is the profile of A3558. The blue filled square is the A3558 cold front distance from the cluster center. } \\label{fig:entropy_outliers}% \\end{figure} The large majority of clusters of the final subsample obey the general rule that cold fronts are hosted by systems with a steep entropy profile in their centers. However, as already pointed out, A3558 is a peculiar case: although its entropy profile is similar to the NCF clusters profiles, it hosts a cold front. Some comments are needed to understand why this outlier does not follow the general behavior. The cold front for A3558 (blue square in Fig. \\ref{fig:entropy_outliers}) is located at larger distance ($ r \\sim 0.05 r_{180}$) with respect to all the other cold fronts we detect, in a region where a weak entropy gradient, not as sharp as in the other CF clusters profiles, is present. This cluster lies at the center of the Shapley supercluster and its special behavior of this cluster is likely related to the unique environment in which it is embedded. To understand the reason why cold fronts can arise in such a system, we refer once more to \\citet{AM06} simulations. When the entropy profile of the cluster does not sharply decrease in the center, the central cold gas is easily pushed away from the dark matter peak, at the merging subclump passage. Accordingly, the cold front emerges at a large distance from the core. However, there is no entropy contrast to trigger the sloshing mechanism and this cold front will not develop further. Cold fronts rising in these systems are short-lived and therefore rare phenomena (see Fig.\\, 12 in \\citealt{AM06}). A3558 is embedded in a very unrelaxed environment where merging events are frequent, and therefore the probability to form (and to observe) such fronts is higher. Alternatively, this cold front might be a merger cold front that we failed to recognize due to the fact that A3558 cannot be classified easily as a merging cluster. As discussed in \\citet{Rossetti_A3558:2007}, it presents some features similar to those of cool core clusters and other properties that are more common in merging clusters. One of the main findings of our paper is that we detect at least one cold front in all steep entropy gradient clusters in the final subsample. \\citet{AM06} show that once the sloshing mechanism sets in cold fronts can be recognized in all the projection planes, even if they are more prominent on the merger plane (see Fig. 19 of their paper). However, the limited resolution of our instruments allow us to recognize only the most apparent brightness discontinuity. Indeed, we have performed some simulations of cold fronts projection with the {\\it XMM-Newton} PSF and we found that cold fronts can only be observed if they lie within some 30\\deg\\ of the plane of the sky. This means that our 100\\% detection rate implies that most steep entropy clusters must host more than one ``prominent'' cold front. This abundance of cold fronts suggests that, whatever the triggering mechanism might be, it must have a high occurrence rate. Since the prominent cold fronts that we can detect are located on the merger plane, the detection of one or more cold fronts in all our steep entropy systems seems to indicate that a sizeable fraction of them are currently experiencing more than one minor merger. Assuming that, crudely speaking, cold fronts are visible for a timescale of about 3 Gyr (this is the case for the dark matter + gas simulation in \\citealt{AM06}, while for the dark matter only this timescale is longer), our cold front detection rate translates into a minimum merger frequency of 1/3 merger event per halo per Gyr. If we further assume a minimum mass ratio of 1/10 we can compare our rate with rates expected from cosmological simulations. Using Fig.8 in a recent paper by \\citet{Fakhouri:2008} we find a merger rate of $\\sim 0.2$ merger per halo per Gyr for mass ratio larger that 1/10. This is somewhat smaller than the minimum rate implied by our observed cold front rates however, given the numerous simplifications we have applied in our calculation, we deem it to be in acceptable agreement. Gas sloshing may provide an important contribution to the cooling-heating problem in cool core clusters \\citep{ZuHone:2009}. The sloshing gas typically moves at sub-(or trans-)sonic velocities carrying a kinetic energy comparable to the thermal energy but the dissipation of this kinetic energy to thermal energy is too slow compared to cooling \\citep{Maxim:2001}. However the sloshing mechanism also brings the outer high entropy gas into the core, mixing it with the cooling gas and resulting in a heat inflow which can prevent the formation of a ``cooling flow'' for periods of time 1-3 Gyr \\citep{ZuHone:2009}. If subcluster encounters are frequent enough, as it is suggested by our high detection rate, the sloshing mechanism can efficiently offset cooling. Intriguingly the sloshing mechanism operates preferentially in steep entropy profile clusters, i.e. precisely those which require heating to offset the cooling. With the coming into operation of the first space-borne micro-calorimeter, quite likely the one onboard the ASTRO-H mission \\citep{Taka_NEXT:2008}, it will be possible to investigate gas motions in the direction of the line of sight, i.e. orthogonally with respect to that of the plane of the sky sampled with cold fronts. The combination of the two informations will afford a reliable estimate of the motions of the ICM in clusters core and estimate their role in offsetting cooling. We have performed a systematic search of cold fronts using {\\it XMM-Newton} data for a sample of 45 objects extracted from the B55 flux limited sample \\citep{Edge:B55}. The main results of our work are the following: \\begin{itemize} \\item{Excluding three unclassified cases, we find that 19 clusters out of 42 host at least one cold front.} \\item{We do not detect any cold front in systems having redshift greater than about 0.075. This is most likely related to {\\it XMM-Newton} resolution limit. By cutting our sample at $z = 0.075$, we restrict our sample to 32 objects with a cold front occurrence of 59\\% .} \\item{Cold fronts are easily detected in systems that are manifestly undergoing a merger event in (or close to) the plane of the sky.} \\item{Out of the 23 clusters of the remaining subsample (systems undergoing a merger event which is not lying in the plane of the sky and non-merging clusters) 10 objects exhibit a cold front. For this final subsample, the entropy profile of systems hosting cold fronts is found to be steeper than that of clusters without them. The difference is observed at radii smaller than about $0.08 r_{180}$ where all our cold fronts are found.} \\item{Our findings are in agreement with simulation based predictions. As shown by \\citet{AM06} an entropy gradient is a necessary ingredient to trigger gas sloshing. } \\item{Since projection effects highly limit the capability of detecting cold fronts, the finding that all the clusters with a steep entropy profile host a cold front implies that most clusters with a steep entropy profile must have more than one cold front. } \\item{Under the assumption that cold fronts in cool core clusters are triggered by minor mergers, we estimate a minimum of 1/3 events per halo per Gyr, which is somewhat larger than that expected from cosmological simulations \\citep{Fakhouri:2008}.} \\item{Gas sloshing may provide an important contribution to the cooling-heating problem in cool core clusters. A robust assessment of the gas motions associated to the sloshing phenomenom will become possible with the coming into operation of the first space borne microcalorimeter.} \\end{itemize}" }, "1003/1003.4732_arXiv.txt": { "abstract": "We investigate the rotational emission from dust grains that rotate around non-principal axes. We argue that in many phases of the interstellar medium, the smallest grains, which dominate spinning dust emission, are likely to have their nutation state (orientation of principal axes relative to the angular momentum vector) randomized during each thermal spike. We recompute the excitation and damping rates associated with rotational emission from the grain permanent dipole, grain-plasma interactions, infrared photon emission, and collisions. The resulting spinning dust spectra generally show a shift toward higher emissivities and peak frequencies relative to previous calculations. ", "introduction": "One of the difficulties in measuring the anisotropies in the cosmic microwave background (CMB) is that the interstellar medium (ISM) also emits microwave radiation through several mechanisms. This ``foreground'' radiation must be modeled and subtracted in order to measure the cosmological parameters accurately using the CMB. The standard theory of ISM microwave emission contains three major emission mechanisms \\citep[e.g. ][]{2000ApJ...530..133T, 2003ApJS..148...97B, 2009AIPC.1141..265F}: synchrotron radiation from relativistic electrons spiralling in the Galactic magnetic field; free-free radiation from ionized gas; and thermal emission from dust grains. These are typically traced by external templates: low-frequency radio maps for the synchrotron \\citep{1982A&AS...47....1H}, H$\\alpha$ for the free-free \\citep{2003ApJS..146..407F}, and far-infrared continuum for the dust \\citep{1999ApJ...524..867F}. \\citet{Kogut1996a, Kogut1996b} reported a spatial correlation between Galactic microwave emission at 31.5, 53 and 90 GHz and the thermal infrared continuum from dust. They interpreted the microwave emission as dust-correlated free-free radiation, on top of the Rayleigh-Jeans tail of thermal emission from dust. Their observations were confirmed by \\citet{deOliveiraCosta1997}, who measured the microwave intensity of the Galaxy at 30 and 40 GHz. \\citet{1997ApJ...486L..23L} claimed the presence of an \"anomalous\" component of Galactic microwave emission, which they observed as a signal at 14.5 and 32 GHz strongly correlated with the diffuse 100 $\\mu$ m intensity. It was far too bright to be thermal dust and had a flat spectrum across these bands, and low-frequency radio and H$\\alpha$ observations predicted far too little synchrotron or free-free emission to explain the signal. \\citet{1997ApJ...486L..23L} proposed that the signal originated from hot gas at $T\\ge10^6\\,$K, which could produce free-free radiation but little H$\\alpha$; however \\citet{1998ApJ...494L..19D} showed that this gas would cool rapidly and that keeping it hot was energetically unfeasible. Several alternative explanations have been proposed. Spinning dust emission is due to the rotation of small dust grains with permanent electric dipole moments. The basic mechanism has been known for decades \\citep{1957ApJ...126..480E, 1970Natur.227..473H, 1992A&A...253..498, 1994ApJ...427..155}, and was suggested as an explanation for the anomalous emission by \\citet{1998ApJ...508..157D} (hereafter DL98b). Magnetic dust emission is due to thermal fluctuations of the magnetic dipole moments of grains including ferromagnetic or ferrimagnetic materials \\citep{1999ApJ...512..740D}. Hard synchrotron radiation would be a new synchrotron component from young (recently-accelerated) high energy electrons, proposed to be strongly correlated with the far-infrared emission from dust due to their common association with recent star formation \\citep{2003ApJS..148...97B}. Both the spinning and magnetic dust hypotheses predict an emission spectrum that peaks in the microwave (the former due to the rotation rates of the smallest grains, and the latter due to the gyrofrequency in ferromagnetic materials). The hard synchrotron hypothesis is now disfavoured due to the low polarization of the anomalous component observed by the {\\slshape Wilkinson Microwave Anisotropy Probe} ({\\slshape WMAP}; \\citealt{2007ApJS..170..335P}), its strong morphological correlation with dust maps \\citep{2004ApJ...614..186F, 2006MNRAS.370.1125D}, and evidence that the anomalous emission has a rising spectrum at low frequencies ($<20$ GHz; \\citealt{1999ApJ...527L...9D, 2004ApJ...617..350F, 2005ApJ...624L..89W}). A key test to distinguish the various models for anomalous emission is to construct predicted emission spectra and compare them to observations. DL98b computed spinning dust spectra for a variety of interstellar environments, accounting for the main processes that affect grain rotation: collisions, grain-plasma interactions, infrared emission, and radiation-reaction torque on the grain electric dipole moment. Model spinning dust spectra have been used extensively to test (and in some cases disfavour or rule out) the spinning dust hypothesis for the anomalous emission seen in the diffuse high-Galactic latitude ISM \\citep{2003ApJS..148...97B, 2004ApJ...614..186F, 2008ApJ...680.1235D, 2009ApJS..180..265G}, in the Galactic Plane \\citep[e.g.][]{2004ApJ...617..350F}, and in dense regions such as molecular clouds \\citep{2004ApJ...614..186F, 2005ApJ...624L..89W, 2006ApJ...639..951C, Casassus2008} and H{\\sc\\,ii} regions \\citep{2007MNRAS.379..297D, 2009ApJ...690.1585D}, supernova remnants \\citep{2007MNRAS.377L..69S}, planetary nebulae \\citep{2007MNRAS.382.1607C}, and an external galaxy \\citep[NGC6946;][]{2010ApJ...709L.108M}. \\citet{2009ApJ...699.1374D} have even used the anomalous emission seen by {\\slshape WMAP} in the warm ionized medium (WIM; traced by H$\\alpha$) to test dust models; they observe a factor of $\\sim 3$ {\\em lower} anomalous emission than predicted, which they tentatively interpret as due to depletion of the smallest dust grains (the polycyclic aromatic hydrocarbons, or PAHs) in the WIM. Recently, the grain rotation problem has been revisited by two theoretical groups. \\citet[][hereafter AHD09]{2009MNRAS.395.1055A} constructed a more detailed model of grain rotation, following the angular velocities of grains using a Fokker-Planck equation and re-evaluating the rotational excitation and damping rates using updated grain properties and a more sophisticated model for the grain-plasma interactions. They also released a public code, {\\sc SpDust}, to compute spinning dust spectra for any input physical conditions and grain properties. \\citet{2010A&A...509A..12Y} presented a quantum-mechanical treatment of several of these processes and computed the resulting emission spectra. The existing theoretical treatments of spinning dust, however, still contain a number of simplifying assumptions. One of the major uncertainties is the grain size distribution and typical dipole moment, however this uncertainty can be turned into a virtue by using it to constrain dust models \\citep[e.g.][]{2009ApJ...699.1374D}. Additionally, there are uncertainties in the physics of grain rotation, such as the validity of the Fokker-Planck approximation or the assumed properties such as the evaporation temperature of departing adsorbed atoms. Some of these pieces of physics are not readily amenable to improvement by theoretical calculations, but others are. The purpose of this paper is to revisit the assumption by DL98b and AHD09 that grains rotate around the axis of largest moment of inertia due to internal dissipation processes. We argue in particular that PAHs in the diffuse and high UV flux phases are likely to be in a random nutation state. This is not a trivial detail: a dust grain rotating around a non-principal axis emits at multiple frequencies, including frequencies well in excess of the instantaneous grain angular velocity. The fact that electric dipole emission depends on the second derivative of the dipole moment $\\ddot\\bmu$ rather than just $\\bmu$ enhances the importance of these higher frequencies.\\footnote{\\citet{2010A&A...509A..12Y} allowed for an arbitrary nutation state, but imposed the assumption that the grain dipole moment be exactly parallel to the axis of greatest moment of inertia, which eliminates three of the four frequencies of emission from an axisymmetric grain. They also did not re-consider the collisional and plasma excitation and drag coefficients.} We show in Section~\\ref{ss:em} that for disc-like grains, at fixed angular momentum incorporating a random nutation state increases the spinning dust emissivity by roughly an order of magnitude. Of course, having a random nutation state also modifies the processes that change grain angular momenta. We investigate each of the major processes and find that the typical grain angular momentum is reduced, but still find a factor of 1.6 increase in the peak spinning dust emissivity $j_\\nu$ and a factor of 1.3 increase in the peak frequency for WIM conditions.\\footnote{For ease of comparison with previous results, our WIM conditions are those of DL98b: density $n_{\\rm H}=0.1\\,$cm$^{-3}$, gas temperature $T=8000\\,$K, H ionization fraction $n($H$^+)/n_{\\rm H}=0.99$, and radiation field normalization $\\chi=1$.} This paper is organized as follows. Section~\\ref{sec:grain} reviews the key parameters of the grain models. Section~\\ref{sec:rotation} describes the expected rotational state of grains and the formalism used in this paper (and in the updated {\\sc SpDust}) for describing the grain angular momentum distribution. Section~\\ref{sec:ed} considers the electric dipole emission from grains rotating in a random nutation state. Subsequent sections consider spin-up and spin-down processes for the grains, taking account of nutation: Section~\\ref{sec:plasma} considers grain-plasma interactions; Section~\\ref{sec:ir} considers infrared photon emission; and Section~\\ref{sec:coll} considers collisions. Predicted spinning dust spectra are shown in Section~\\ref{sec:results}, where we also explore the sensitivity to some of our assumptions. We conclude in Section~\\ref{sec:disc}. The physical processes affecting grains in non-uniform rotation are very complex, and this paper contains some unavoidably long calculations. The reader interested primarily in the results may skip directly from the end of Section~\\ref{ss:em} to the beginning of Section~\\ref{sec:results}. We note that \\citet{Hoang10} have recently completed a related analysis in which axisymmetric dust grains are followed through a 2-dimensional space of angular velocities $(\\omega_\\parallel,\\omega_\\perp)$. Our analyses agree on the basic conclusion that allowing grains to rotate around a non-principal axis results in an increase in the spinning dust emissivity and an increase in the peak frequency. ", "conclusions": "\\label{sec:disc} The purpose of this work was to revisit the assumption of DL98b and AHD09 that PAHs rotate about their axis of main inertia. The motivation in doing so is that thermal spikes following the absorption of UV photons randomize the orientation of the grain with respect to the angular momentum axis. These absorption events happen frequently enough (i.e. on timescale shorter than the timescale for significant changes in the total angular momentum) that we expect such a randomization to be effective in most environments. Thus we expect the results from this work (``case 2'') to be a better approximation to diffuse or high-radiation environments (CNM, WNM, WIM, PDR, and RN) than those from AHD09, which assumed rapid dissipation of the nutational energy ($\\theta=0$ or ``case 1''). However, the new release of {\\sc SpDust} allows the user to choose either case; for example, one may wish to explore the range of cases in dark cloud environments where thermal spikes are infrequent, or what happens if an as-yet-unidentified dissipational process is active and restores $\\theta=0$. In this work, we showed that, for a given angular momentum, the power radiated by a grain in case 2 is $\\sim 10$ times higher than that radiated by a grain in case 1. This is because in case 2, the grain emits at higher frequencies, including above the one corresponding to the instantaneous angular velocity, as it is not rotating around the axis of greatest inertia. We evaluated the rotational excitation and damping rates in case 2 as a function of grain size and environment conditions, and the resulting angular momentum distribution. We showed that in a given environment, grains in case 2 have a lower rms angular momentum than those in case 1, by a factor of $\\sim 0.7$. This is due to larger damping rates, in particular radiation-reaction damping, in case 2. \\begin{figure} \\begin{center} \\includegraphics[width = 3.2in]{Omega_rms_mu_ip.eps} \\includegraphics[width = 3.2in]{Power_peak_mu_ip.eps} \\caption{Top panel: normalized rms angular momentum $\\langle \\Omega^2\\rangle^{1/2}$ as a function of the ratio of in-plane to total dipole moment. Bottom panel: estimate of the peak frequency $\\int \\nu (\\rmd P/\\rmd \\nu) \\rmd \\nu \\big{/} P$, as a function of this ratio. Both are for a dust grain of radius $a = 5\\,$\\AA\\ and dipole moment per atom $\\beta = 0.38$ debye, in WIM conditions.} \\label{fig:rms_peak} \\end{center} \\end{figure} \\begin{figure} \\begin{center} \\includegraphics[width = 3.2in]{Power_mu_ip.eps} \\caption{Total power emitted by a dust grain of radius $a = 5\\,$\\AA\\ and dipole moment per atom $\\beta = 0.38$ debye, in WIM conditions, as a function of the ratio of in-plane to total dipole moment. } \\label{fig:pow_muip} \\end{center} \\end{figure} The combination of these results leads to a spinning dust spectrum peaking at slightly higher frequencies in case 2, and a total power approximately twice as large as that emitted in case 1. Finally, we showed that the spectrum in case 2 is only weakly sensitive to the precise value of the $\\mu_{\\rm ip}^2:\\mu_{\\rm op}^2$ ratio. \\citet{2009ApJ...699.1374D} found a tension between theoretical results and microwave observations of the WIM: the theory was a factor of $\\sim 3$ larger than the observations, and the peak frequency of the spinning dust and its amplitude could not be simultaneously reconciled by changing $\\beta$ (the normalization of the dipole moment). By increasing the theoretical emissivity and moving its peak to higher frequencies, our results may worsen this tension. This seems likely to strengthen the empirical case for depletion of the PAH population in the WIM phase, however there are other conceivable explanations for this discrepancy. The random walk model for the dipole moment may not apply well to the smallest grains (e.g. one could imagine that some of the small PAHs have symmetries that guarantee $\\bmu=0$ exactly), or one could imagine extra low-frequency internal degrees of freedom which allow the grain to relax to a state where it rotates around the axis of greatest moment of inertia. A detailed exploration of the parameter space \\citep[as was done by][]{2009ApJ...699.1374D} is beyond the scope of this paper. As a final note, we present some of the remaining issues in the treatment of the rotational physics of the smallest dust grains: \\begin{itemize} \\item {\\em Triaxiality}: Many PAHs have triaxial moment of inertia tensors (e.g. ovalene C$_{32}$H$_{14}$, circumpyrene C$_{42}$H$_{16}$, and their derivatives). This case was not treated in the present paper due to its much greater complexity: since the dipole moment then depends on elliptic functions of the angle conjugate to the nutation action, a countably infinite number of frequencies are emitted. Aside from this aspect, however, the underlying formalism in this paper would be applicable: the nutation action (rather than $h K = 2\\pi L\\cos\\theta$) would be conserved in free rotation and we would average over this action instead of $\\cos\\theta$. The analysis would also break into two cases depending on whether the grain lies on the short-axis or long-axis side of the separatrix. \\item {\\em Impulsive torques}: Some of the sources of torque, such as ion impacts, impart large but infrequent changes in angular momentum. This could in principle lead to ``rotational spikes'' analogous to the well-known thermal spikes in the grains' internal energy, and would not be treated correctly by the Fokker-Planck equation (which is a diffusive approximation).\\footnote{This issue is treated in \\citet{Hoang10}; they find that the principal effect on the spinning dust spectrum is the existence of a ``tail'' to high frequencies resulting from transient spin-up of the grains.} \\item {\\em Ancillary data}: We have not fully quantified the uncertainties in the ancillary data, such as evaporation temperatures, the emissivity in the lowest-frequency vibrational modes, and the grain charging model (photoelectric and electron/ion impact). However, our hope in making the {\\sc SpDust} code publicly available is to provide users the flexibility to explore deviations from default or fiducial parameters. \\end{itemize}" }, "1003/1003.0327_arXiv.txt": { "abstract": "We hereby study the stability of a massless probe orbiting around an oblate central body (planet or planetary satellite) perturbed by a third body, assumed to lie in the equatorial plane (Sun or Jupiter for example) using an Hamiltonian formalism. We are able to determine, in the parameters space, the location of the frozen orbits, namely orbits whose orbital elements remain constant on average, to characterize their stability/unstability and to compute the periods of the equilibria. The proposed theory is general enough, to be applied to a wide range of probes around planet or natural planetary satellites. The BepiColombo mission is used to motivate our analysis and to provide specific numerical data to check our analytical results. Finally, we also bring to the light that the coefficient $J_2$ is able to protect against the increasing of the eccentricity due to the Kozai-Lidov effect. ", "introduction": "BepiColombo (MPO and MMO orbiters) is a joint European and Japanese space agencies space mission aimed at studying the planet Mercury. The MPO (Mercury Planetary Orbiter) will be brought into a polar elliptical orbit around Mercury with an inclination of $88\u201390\\Deg$, an eccentricity of $0.1632$ and a semi-major axis of $3\\,394$ km. The MMO (Mercury Magnetospheric Orbiter) will also be brought into a polar elliptical orbit with % an eccentricity of $0.6679$ and a semi-major axis of $8\\,552$ km. Actually polar orbits are very interesting for scientific missions to planetary satellites (with near polar low-altitude) or to planet (with high-eccentric high-altitude). The orbital dynamics of such space probes is governed by the oblateness ($J_2$ effect) of the central body around which the space probe is orbiting and the gravity field from the third body. A well-known effect of the third-body perturbation is the change in the stability of circular orbits related to orbit inclination. This effect is a natural consequence of the Kozai-Lidov resonance (\\citealt{Kozai1962,Lidov1963}). The final fate of such a satellite is the collision with the central body. Therefore the control of the orbital eccentricity leads to the control of the satellite lifetime. \\cite{Scheeres2001} studied near-circular % orbits in a model that included both the third body's gravity and $J_2$. In addition \\cite{SanJuan2006} studied orbit dynamics about oblate planetary satellites using a rigorous averaging method. \\cite{Paskowitz2006} added the effect of the coefficient $J_3$. These authors mainly focused their attention to an orbiter around planetary satellites especially for Europa orbiter. So they did not take into account the eccentricity of the third body and they detailed the near-circular orbits. Our purpose is to build a simplified Hamiltonian model, as simple as possible, which will reproduce the motion of probes orbiting an oblate central body also taking into account the third body effect. Especially we are looking for the conditions that give rise to frozen orbits. Frozen orbits are orbits that have orbital elements constant on average. These particular orbits are able to keep constant the eccentricity. Therefore in a neighbourhood of these orbits there is a stability area where even a limited control could be used to avoid the crash onto the central body. Beside the oblateness of the central body and the gravity effect of the third body, our averaged model takes into account also the eccentricity of the orbit of the third body. Moreover let us observe that our results are given in closed form with respect to eccentricity and inclination of the probe, namely we do not perform any power series expansion; therefore, our theory applies for arbitrary eccentricities and inclinations of the space probe, and is not limited to almost-circular orbits. We can thus conclude that the theory is general enough to be applied to a wide range of probes around a planet or around a natural planetary satellite and, can be formulated and presented in a general way that allows extension of the results to other cases. The Mercury orbiter mission (BepiColombo) is used to motivate our analysis and to provide specific numerical data to check our analytical results. We are able to provide the location of frozen orbits and study their stability as a function of the involved parameters, using implicit equations and graphics. Finally we give the analytical expressions of the periods at the stable equilibria. The analytical results are verified and confirmed using dedicated numerical simulations of the whole model. To conclude, we discuss the effect of the protection of $J_2$ on the increase of the eccentricity due to Kozai-Lidov effect and the apparition of an asymmetry caused by the addition of the coefficient $J_3$. ", "conclusions": "The orbit dynamics of a space probe orbiting a planet or a natural planetary satellite has been investigated. The proposed model includes the effects of $J_2$ for the central body and the perturbation of the third body. We have developed a doubly averaged Hamiltonian and studied the location of the stable and unstable frozen orbits. Our analytical approach allows us to compute also the periods of the free librations at the equilibria. The analytical results have been checked and validated numerically by performing numerical integrations of the complete systems. Our theory is able to explain the behavior of our preliminary numerical investigations where the variation of the amplitude of the eccentricity is null and the presence of a separatrix has been found by numerical investigation. The theory is general enough to be applied to a wide range of probes around any planet or any natural planetary satellite, provided that they respect the hypotheses used to obtain our Hamiltonian model. We have shown the protection mechanism of the coefficient $J_2$ on the increasing of the eccentricity due to Kozai-Lidov effect. This mechanism is therefore able to find a larger number of frozen orbits than for the only Kozai-Lidov problem. We have also explained the asymmetry of the frozen equilibria caused by the addition of the coefficient $J_3$. We have also brought to the light a local deformation of the Kozai-Lidov equilibria that appears at high eccentricity, high inclination and large value of $\\gamma$. It would be interesting to take this theory into account to choose the intial semi-major axis and eccentricity of an orbiter for future missions around planets or planetray satellites." }, "1003/1003.5672_arXiv.txt": { "abstract": "The outer regions of galactic disks have received increased attention since ultraviolet observations with GALEX demonstrated that nearly 30\\% of galaxies have UV emission beyond their optical extents, indicating star formation activity. These galaxies have been termed extended UV (XUV) disks. Here, we address whether these observations contradict the gas surface density threshold for star formation inferred from \\Ha \\ radial profiles of galaxies. We run smoothed particle hydrodynamics simulations of isolated disk galaxies with fiducial star formation prescriptions and show that over-densities owing to the presence of spiral structure can induce star formation in extended gas disks. For direct comparison with observations, we use the 3-D radiative transfer code \\textsc{Sunrise} to create simulated FUV and \\Ks -band images. We find that galaxies classified as Type I XUV disks are a natural consequence of spiral patterns, but we are unable to reproduce Type II XUV disks. We also compare our results to studies of the Kennicutt-Schmidt relation in outer disks. ", "introduction": "Star formation rate profiles derived from \\Ha \\ emission indicate that galaxies follow a ``Kennicutt-Schmidt Law'' where the star formation rate and gas surface densities averaged in azimuthal annuli are related by $\\Sigma_{\\rm{SFR}} \\propto \\Sigma_{\\rm{gas}}^{1.5}$ \\citep{Kennicutt-1998, Kennicutt-1989}. However, sudden drops in \\Ha \\ intensity create strong departures at a gas surface density of 3-5 \\sdunits. This is generally interpreted as a threshold density for star formation \\citep{Kennicutt-1989, Martin-Kennicutt-2001}. Several explanations for this threshold have been proposed, including a transition between dynamically unstable and stable regions of the galaxy \\citep[e.g.][]{Toomre-1964} or a phase transition of the gas \\citep[e.g.][]{Elmegreen-Parravano-1994, Schaye-2004, Krumholz-et-al-2009}. Outer disks have consequently been considered inhospitable environments for star formation despite the existence of \\Ha \\ knots at large radii \\citep{Kennicutt-1989,Martin-Kennicutt-2001, Ferguson-et-al-1998}. Interest in outer disks was rekindled by observations with the Galaxy Evolution Explorer (GALEX: \\citet{Martin-et-al-2005}) which revealed that $\\sim 30$\\% of disk galaxies have UV emitting sources beyond their optical disks \\citep{Thilker-et-al-2007, Gildepaz-et-al-2005, Thilker-et-al-2005, Zaritsky-Christlein-2007}. In the case of M83 and NGC 4625, UV knots have been identified as low mass stellar complexes and, if visible in the \\Ha, are generally ionized by a single star \\citep{Gildepaz-et-al-2007}. These knots are dynamically cold and rotating, indicating that outer disk complexes are extensions of the inner disk \\citep{Christlein-Zaritsky-2008}. \\citet[][hereafter T07]{Thilker-et-al-2007} classified extended UV (XUV) disks into two types. In 2/3 of XUV disks, the UV emission is structured, lying in filamentary, spiral patterns (Type I XUV disks). In the other 1/3, a large zone in the outer regions of the galaxy is dominated by UV emission (Type II XUV disks). The prevalence of star formation in outer disks raises new questions about the nature of star formation in diffuse environments. In particular, we address whether a threshold density to star formation is consistent with the observed emission. \\citet[][hereafter B07]{Boissier-et-al-2007} selected a sample of galaxies of large angular size, including XUV and non-XUV disks, and studied their extinction corrected UV and gas profiles. While the uncertainties are large, B07 do not find a departure from the Kennicutt-Schmidt relation at low gas density. They suggest that the high mass end of the stellar initial mass function (IMF) is not sampled in every stellar cluster at the low star formation rates observed and \\Ha \\ emitting stars are not formed in all star forming complexes. They argue that this is responsible for the difference in the UV and \\Ha \\ profiles and the Kennicutt-Schmidt law is followed to low surface density. \\citet{Pflamm-Kroupa-2008} propose instead that the Kennicutt-Schmidt law is obeyed below the traditional threshold density but that high mass stars capable of emitting in the \\Ha \\ simply do not form at low gas densities. This is supported by a deficiency of H$\\alpha$ emission in dwarf galaxies that cannot be explained by IMF sampling effects \\citep{Lee-et-al-2009}. \\citet{Roskar-et-al-2008b} note that stars migrate radially from their formation sites in the disk, which can modify the stellar populations of outer disks. However, it is not clear that this process is efficient enough to relocate UV emitting stars to 2-3 \\D25 \\citep{Roskar-et-al-2008a, Roskar-et-al-2008b}. Alternatively, at some sites the gas density may exceed a star formation threshold locally, allowing stars to form beyond the radius where the azimuthally averaged gas density is at or below a threshold density \\citep{Kennicutt-1989, Martin-Kennicutt-2001, Schaye-2004, Elmegreen-Hunter-2006, Gildepaz-et-al-2007}. Outer disk stellar complexes are often coincident with local H\\,I over-densities \\citep[][T07]{Ferguson-et-al-1998}, supporting this interpretation. XUV disks are gas rich for their star formation rates, indicating that a high fraction of their gas is not undergoing star formation (T07). Studies of UV selected sources in the XUV disk of M\\,83 indicate that star formation occurs only above a critical gas density defined by the Toomre criterion for dynamical instability \\citep{Dong-et-al-2008} and, where it occurs, follows a Kennicutt-Schmidt law. This theory has been explored by \\citet{Elmegreen-Hunter-2006}. In their model, gas clumping triggered by spiral density waves, radial variations in the interstellar medium (ISM) turbulence, and gas phase transitions lead to localized regions of active star formation in extended gas disks. They implement a Kennicutt-Schmidt law and a radially varying threshold density based on the Toomre criterion \\citep{Elmegreen-Hunter-2006}. These processes enable star formation to extend into the outer disk. In \\citet{Bush-et-al-2008}, we used smoothed particle hydrodynamic (SPH) simulations of an isolated disk galaxy with an extended gas disk of initially constant surface density to explore the morphology resulting from such in situ outer disk star formation. We showed that spiral structure can propagate from the inner disk to the extended gas disk, producing regions of locally enhanced gas density and triggering star formation in filaments similar to those observed in Type I XUV disks. However, this study was restricted to one galaxy model, limiting our ability to draw general conclusions. In this paper, we expand on the work of \\citet{Bush-et-al-2008} by varying the structure of the simulated galaxy, in particular the gas profile and the disk to halo mass ratio, to explore the amount and morphology of outer disk star formation that result naturally without altering traditional star formation laws. In addition, we produce simulated images to directly compare UV and \\Ks \\ emitting populations to observations. This allows us to classify our models as normal, Type I or II XUV disks. We also compare our results to studies of the Kennicutt-Schmidt law, radially and locally. While we test whether the UV and \\Ks \\ emission of XUV disks can be reproduced using traditional star formation laws, we do not create H$\\alpha$ images and therefore do not address the discrepancy between H$\\alpha$ and UV radial profiles. In \\S 2 we describe our models, in \\S 3 we describe the evolution and uncertainties in one model in detail, in \\S 4 we describe the full set of models and how they compare to observations and finally in \\S 5 and \\S 6 we discuss our results. Our models succeed in producing the UV morphologies of Type I XUV disks. However, they do not reproduce the UV bright outskirts of Type II XUV disks. ", "conclusions": "\\label{sec:conc} We have explored how outer disk star formation in extended gas disks resulting from a Kennicutt-Schmidt star formation prescription and volume density threshold compares to observations of XUV disks. We used a set of models with outer disk structures motivated by H\\,I surveys. We find that models with slowly varying high surface density outer gas disks would be classified as Type I XUV disks. Possible variations in the star formation rate normalization or spiral structure amplitude mean that gas distributions that result in only a few knots of outer disk UV emission in our models may not always have UV emitting sources in their observed outer disks. While we cannot rule out other explanations for Type I XUV disks, since the Type I XUV emission in our models results naturally from empirical star formation laws and the observed structure of galaxy disks, we argue that the theory outlined here is a viable explanation. The star formation in the outer disk of our models is not prevalent enough to erase the tell-tale truncation of the profile caused by the star formation volume density threshold. Since B07 see no truncation in the UV with gas radial profile, this could indicate that we need to revise our star formation prescription. However, the star formation rates in our models are consistent with star formation rates in the Type I XUV disk M\\,83 \\citet{Dong-et-al-2008}. More detailed observational studies of the relation of gas density to star formation rate in outer disks are needed to determine whether the observed fields of M\\,83 are a special case or whether disks selected for XUV emission are subject to different processes than the B07 sample. While XUV emission does not depend enough on the detail of gas density profiles to distinguish easily between them, it is a good indicator of the presence of gas at large radii in galaxies. Our results imply that nearly 1/3 of galaxies have large amounts of gas at 2-3 times their optical extents. In addition, forming Type II XUV disks with traditional star formation laws must require some process, such as rapid gas accretion, for creating a large young stellar population without an accompanying older population. These results may be outlining a picture where accreting gas in the outer parts of disks is common at late times, as suggested by the inside out disk formation scenario demonstrated by simulations of structure formation \\citep[e.g.][]{Mo-Mao-White-1998, Brook-et-al-2006}, and supported by the possibility that cold clouds developing in cosmological simulations may accrete onto disks at late times \\citep{Keres-Hernquist-2009}. The next step in exploring the observed signatures of in situ star formation would be to run very high resolution simulations that could be compared to observational work such as that of \\citet{Herbert-Fort-et-al-2009}. These authors have computed the cross correlation of young clusters in NGC 3184 and found a significant inter-cluster signal for separations of less than 1 kpc, with cluster masses of $\\sim 10^{3}-10^{4}$ \\Msun. With higher resolution simulations, it is possible that measurements like these could differentiate between in situ star formation and stellar migration scenarios. Moreover, simulations with star formation prescriptions that are linked directly to the molecular gas may help reconcile in situ star formation and the results of B07. \\citet{Robertson-et-al-2008} present models indicating that directly relating the star formation rate to the molecular gas may eliminate truncations in the star formation rate profile. Continued observational studies and sophisticated simulations are needed to explore this complex regime in star formation." }, "1003/1003.2787_arXiv.txt": { "abstract": "We build on the formulation developed in Sridhar \\& Singh~({\\emph{JFM}, {\\bf 664}, 265, 2010}), and present a theory of the \\emph{shear dynamo problem} for small magnetic and fluid Reynolds numbers, but for arbitrary values of the shear parameter. Specializing to the case of a mean magnetic field that is slowly varying in time, explicit expressions for the transport coefficients, $\\alpha_{il}$ and $\\eta_{iml}$, are derived. We prove that, when the velocity field is non helical, the transport coefficient $\\alpha_{il}$ vanishes. We then consider forced, stochastic dynamics for the incompressible velocity field at low Reynolds number. An exact, explicit solution for the velocity field is derived, and the velocity spectrum tensor is calculated in terms of the Galilean--invariant forcing statistics. We consider forcing statistics that is non helical, isotropic and delta--correlated--in--time, and specialize to the case when the mean--field is a function only of the spatial coordinate $X_3$ and time $\\tau\\,$; this reduction is necessary for comparison with the numerical experiments of Brandenburg, R{\\\"a}dler, Rheinhardt \\& K\\\"apyl\\\"a (\\emph{ApJ}, {\\bf 676}, 740, 2008). Explicit expressions are derived for all four components of the magnetic diffusivity tensor, $\\eta_{ij}(\\tau)\\,$. These are used to prove that the shear--current effect cannot be responsible for dynamo action at small $\\re$ and $\\rem$, but for all values of the shear parameter. ", "introduction": "Astrophysical systems like planets, galaxies and clusters of galaxies possess magnetic fields which exhibit definite spatial ordering, in addition to a random component. The ordered (or ``large--scale'') components are thought to originate from turbulent dynamo action in the electrically conducting fluids in these objects. The standard model of such a process involves amplification of seed magnetic fields due to turbulent flows which lack mirror--symmetry (equivalently, which possess helicity) \\cite{M78,KR80,BS05}. The turbulent flows generally possess large--scale shear, which is expected to have significant effects on transport properties \\cite{LK09}; however, it is not clear whether the turbulent flows are always helical. Recent work has explored the possiblity that non--helical turbulence in conjunction with background shear may give rise to large--scale dynamo action \\cite{BRRK08,Yousef08a,Yousef08b,KKB,HP09,RK03}. The evidence for this comes mainly from direct numerical simulations \\cite{BRRK08,Yousef08a, Yousef08b}, but it by no means clear what physics drives such a dynamo. One possibility that has received some attention is the shear--current effect \\cite{RK03}, where an extra component of the mean electromotive force (EMF) is thought to result in the generation of the cross--shear component of the mean magnetic field from the component parallel to the shear flow. However, there is no agreement yet whether the sign of such a coupling is favorable to the operation of a dynamo; some analytic calculations \\cite{RS06,RK06} and numerical experiments \\cite{BRRK08} find that the sign of the shear--current term is unfavorable for dynamo action. A quasilinear kinematic theory of dynamo action in a linear shear flow of an incompressible fluid which has random velocity fluctuations was presented in \\cite{SS09}, who used the ``second order correlation approximation'' (SOCA) in the limit of zero resistivity. Unlike earlier analytic work which treated shear as a small perturbation, this theory did not place any restriction on the strength of the shear. They arrived at an integro--differential equation for the evolution of the mean magnetic field and argued that the shear-current assisted dynamo is essentially absent. The theory was extended to take account of non zero resistivity in \\cite{SS10}; this is again nonperturbative in the shear strength, uses SOCA, and is rigorously valid in the limit of small magnetic Reynolds number ($\\rem$) but with no restriction on the fluid Reynolds number ($\\re$). The kinematic approach to the \\emph{shear dynamo problem} taken in \\cite{SS09, SS10} uses in an essential manner the shearing coordinate transformation and the Galilean invariance (which is a fundamental symmetry of the problem) of the velocity fluctuations. The present work extends \\cite{SS10} by giving definite form to the statistics of the velocity field; specifically, the velocity field is assumed to obey the forced Navier--Stokes equation, in the absence of Lorentz forces. In section~II we begin with a brief review of the salient results of \\cite{SS10}. The expression for the Galilean--invariant mean EMF is then worked out for the case of a mean magnetic field that is slowly varying in time. Thus the mean--field induction equation, which is an integro--differential equaton in the formulation of \\cite{SS10} now simplifies to a partial differential equation. This reduction is an essential first step to the later comparision with the numerical experiments of \\cite{BRRK08}. Explicit expressions for the transport coefficients, $\\alpha_{il}$ and $\\eta_{iml}$, are derived in terms of the two--point velocity correlators. We then recall some results from \\cite{SS10}, which express the velocity correlators in terms of the velocity spectrum tensor. This tensorial quantity is real when the velocity field is non helical; we are able to prove that, in this case, the transport coefficient $\\alpha_{il}$ vanishes. Section~III develops the dynamics of the velocity field at low $\\re$, using the Navier--Stokes equation with stochastic external forcing. An explicit solution for the velocity field is presented and the velocity spectrum tensor is calculated in terms of the Galilean--invariant forcing statistics. For non helical forcing, the velocity field is also non helical and the transport coefficient $\\alpha_{il}$ vanishes, as noted above. We then specialize to the case when the forcing is not only non helical, but isotropic and delta--correlated--in--time as well. In section~IV we specialize to the case when the mean--field is a function only of the spatial coordinate $X_3$ and time $\\tau\\,$; this reduction is necessary for comparision with the numerical experiments of \\cite{BRRK08}. Explicit expressions are derived for all four components of the magnetic diffusivity tensor, $\\eta_{ij}(\\tau)\\,$, in terms of the velocity power spectrum; the late--time saturation values, $\\eta^\\infty_{ij}\\,$, have direct bearing on the growth (or otherwise) of the mean magnetic field. Comparisons with earlier work---in particular \\cite{BRRK08}---are made, and the implications for the shear--current effect are discussed. We then conclude in section~V. ", "conclusions": "Building on the formulation of \\cite{SS10}, we have developed a theory of the \\emph{shear dynamo problem} for small magnetic and fluid Reynolds numbers, but for arbitrary values of the shear parameter. Our primary goal is to derive precise analytic results which can serve as benchmarks for comparisons with numerical simulations. A related goal is to resolve the controversy surrounding the nature of the shear--current effect, without treating the shear as a small parameter. We began with the expression for the Galilean--invariant mean EMF derived in \\cite{SS10}, and specialized to the case of a mean magnetic field that is slowly varying in time. This resulted in the simplification of the mean--field induction equation, from an integro--differential equation to a partial differential equation. This reduction is the first step to the later comparison with the numerical experiments of \\cite{BRRK08}. Explicit expressions for the transport coefficients, $\\alpha_{il}$ and $\\eta_{iml}$, were derived in terms of the two--point velocity correlators which, using results from \\cite{SS10}, were then expressed in terms of the velocity spectrum tensor. Then we proved that, when the velocity field is non helical, the transport coefficient $\\alpha_{il}$ vanishes; just like everything else in this paper, this result is non perturbative in the shear parameter. We then considered forced, stochastic dynamics for the incompressible velocity field at low Reynolds number. An exact, explicit solution for the velocity field was derived, and the velocity spectrum tensor was calculated in terms of the Galilean--invariant forcing statistics. For non helical forcing, the velocity field is also non helical and the transport coefficient $\\alpha_{il}$ vanishes, as noted above. We then specialized to the case when the forcing is not only non helical, but isotropic and delta--correlated--in--time as well. We considered the case when the mean--field was a function only of the spatial coordinate $X_3$ and time $\\tau\\,$; the purpose of this simplification was to facilitate comparison with the numerical experiments of \\cite{BRRK08}. Explicit expressions were derived for all four components, $\\eta_{11}(\\tau)$, $\\eta_{22}(\\tau)$ $\\eta_{12}(\\tau)$ and $\\eta_{21}(\\tau)$, of the magnetic diffusivity tensor, in terms of the velocity spectrum tensor. Important properties of this fundamental object are as follows: \\begin{enumerate} \\item All the components of $\\eta_{ij}$ are zero at $\\tau=0$, and saturate at finite values at late times, which we denote by $\\eta^\\infty_{ij}\\,$. \\item The off--diagonal components, $\\eta_{12}$ and $\\eta_{21}$, vanish when the microscopic resistivity vanishes. \\item The sign of $\\eta_{12}^{\\infty}$ is sensitive to the values of the control parameters. This may help reconcile, to some extent, the fact that different signs for $\\eta_{12}^{\\infty}$ are reported in \\cite{RK06} and \\cite{BRRK08}. \\end{enumerate} We derived the condition --- the inequality~(\\ref{dyncond}) --- required for the growth of the mean magnetic field: the sum of three terms must exceed unity. It was demonstrated that two of the terms are very small in magnitude, and hence dynamo action was controlled by the behaviour of one term. i.e. the mean magnetic field would grow if $\\left(\\eta^{\\infty}_{21}S/\\eta_T^2K_3^2\\right)$ exceeds unity. This is possible for small enough $K_3^2$, so long as $\\left(\\eta^{\\infty}_{21}S\\right)$ is positive. However, we see from Figures~(1--3) that $\\eta^{\\infty}_{21}$ is always positive, implying that the product $\\left(\\eta^{\\infty}_{21}S\\right)$ is always negative. Thus the mean--magnetic field always decays, a conclusion which is in agreement with those of \\cite{BRRK08}, \\cite{RS06} and \\cite{RK06}. We then related the above conclusions to the shear--current effect, and demonstrated that the shear--current effect cannot be responsible for dynamo action, at least for small $\\re$ and $\\rem$, but for all values of the shear parameter. In \\cite{BRRK08}, it is suggested that the dynamo action observed in their numerical experiments might be due to a fluctuating $\\alpha$--effect; addressing this issue is being the scope of our present calculations." }, "1003/1003.2052_arXiv.txt": { "abstract": "Core helium burning is the dominant source of energy of extreme horizontal branch stars, as the hydrogen envelope is too small to contribute to the nuclear energy output. The evolution of each mass in the HR diagram occurs along vertical tracks that, when the core helium is consumed, evolve to higher \\teff\\ and then to the white dwarf stage. The larger is the mass, the smaller is the \\teff\\ of the models, so that the zero age horizontal branch (ZAHB) is ``horizontal\". In this paper we show that, if the helium mass fraction (Y) of the envelope is larger than Y$\\sim$0.5, the shape of the tracks changes completely: the hydrogen burning becomes efficient again also for very small envelope masses, thanks to the higher molecular weight and to the higher temperatures of the hydrogen shell. The larger is Y, the smaller is the envelope mass that provides strong H--shell burning. These tracks have a curled shape, are located at a \\teff\\ following the approximate relation \\teff=8090+ 32900$\\times$Y, and become more luminous for larger envelope masses. Consequently, the ZAHB of the very high helium models is ``vertical\" in the HR diagram. Synthetic models based on these tracks nicely reproduce the location and shape of the ``blue hook\" in the globular cluster \\ocen, best fit by a very high \\teff\\ (bluer) sequence with Y=0.80 and a cooler (redder) one with Y=0.65. Although these precise values of Y may depend on the color--\\teff\\ conversions, we know that the helium content of the progenitors of the blue hook stars can not be larger than Y$\\sim$0.38--0.40, if they are descendants of the cluster blue main sequence. Consequently, this interpretation implies that all these objects must in fact be progeny of the blue main sequence, but they have all suffered further deep mixing, that has largely and uniformly increased their surface helium abundance, during the red giant branch evolution. A late helium flash can not be the cause of this deep mixing, as the models we propose have hydrogen rich envelopes much more massive than those required for a late flash. We discuss different models of deep mixing proposed in the literature, and conclude that our interpretation of the blue hook can not be ruled out, but requires a much deeper investigation before it can be accepted. ", "introduction": "\\label{sec:intro} The vast problematics involving the globular cluster $\\omega$ Cen goes from the presence of multiple evolutionary sequences all along the main sequence \\citep{bedin2004}, the subgiant and the giant branches \\citep{lee1999nature,pancino2000,rey2004,sollima2005}, to the presence of a wide spread in metal content (e.g., Norris \\& Da Costa 1995; Suntzeff \\& Kraft 1996; Smith et al. 2000; Johnson et al. 2009) and a likely spread in helium content \\citep{norris2004,piotto2005}. Even if present in other clusters, these features are not all found together as in this most massive GC in the Galaxy. Another striking feature in the colour-magnitude (CM) diagram of $\\omega$ Cen is the large number of very faint and very blue stars at the end of the horizontal branch (HB), beyond the temperature and magnitude limits attainable by currently known HB structures. Such objects are commonly called \"blue hook\" stars and are found in other clusters ---M54 \\citep{rosenberg2004}, NGC 2808 \\citep{brown2001}, NGC 2419 \\citep{ripepi2007}, perhaps NGC 6388 \\citep{busso2007}--- but in $\\omega$ Cen their distribution in the CM diagram is different from what observed elsewhere (Figure \\ref{f2}, left panel). Most of the blue hook objects are arranged along two ``parallel sequences'' inclined towards high temperatures \\citep{anderson2009}. Spectroscopic analysis of the blue hook stars has been done by \\cite{moehler2004} for NGC~2808 resulting in \\teff$\\sim$35--36$\\times 10^3$K for the hottest objects, with a couple of stars probably already evolving towards the white dwarf stage. In \\ocen, they are in the range 32--36$\\times 10^3$K, according to \\cite{moehler2007}, but many hotter stars are also present. More recently, \\cite{villanova2009} have re--examined several blue hook stars in \\ocen\\ and find a clustering in a range of \\teff$\\sim 35-38 \\times 10^3$K. Being the spectroscopic tools hampered by uncertainties much larger than the HST photometry, we can infer from the vertical shape of the extreme blue hook that all its stars cluster at very similar \\teff. We do know that these extreme HB stars should be mainly burning helium in their cores, and evolve along a ``vertical\" track that can simulate the shape of the blue hook, but without covering its whole luminosity extension. In fact, in order to explain the blue hook with tracks of this kind, \\cite{cassisi2009} adopt a superposition of stars having a normal helium content (that start their evolution at a larger luminosity) plus stars with enhanced helium (starting their evolution at lower luminosity), born from the high helium blue MS stars, plus stars that suffered a late helium core flash, and are then also carbon enhanced \\citep{moehler2007}. A reason to justify the hypothesis that the blue hook is the superposition of the progeny of very different evolutionary parents comes from the analysis of helium abundance, since blue hook stars show a great variety of results: they are both helium--normal and helium rich objects, the latter ones close to the range Y=0.38--0.40 believed to apply to the blue MS \\citep[see, e.g. figure 2 in][]{moehler2007}. The hottest (\\teff$>$40000K) stars are generally extremely helium rich, or even very helium poor. In our opinion, however, it is not necessary to connect these different helium abundances to different evolutionary paths. The atmospheric analysis is very difficult, and the resulting surface helium abundance may be affected by errors difficult to be fully quantified. In addition, it is well possible that the main reasons why the surface abundances of these stars are largely different from each other, are surface phenomena acting on timescales much shorter than the evolutionary time in HB, like helium sedimentation in the thin atmospheric layer, or residual mass loss. The presence of a well defined locus in the HR diagram, together with what looks like a broader (cooler) sequence at redder colors (see Figure 2, left panel) appears more like some kind of evolutionary boundary that limits the stars' evolution, than as the casual result of many different evolutionary paths. The HST data set shown in Figure 2 is the same used and fully described in \\cite{cassisi2009}, namely a mosaic of 3$\\times$3 fields obtained with the ACS/WFC (GO--9442, PI A. Cool) through the F435W and F625W filters. More than 350 stars populate the thin vertical part of the blue hook, and $\\sim$170 are spread on the right of this sequence. \\begin{figure*} \\includegraphics[width=8cm]{f1a.eps} \\includegraphics[width=8cm]{f1b.eps} \\caption{On the left we plot the evolutionary tracks of different total masses for surface helium content Y=0.80, 0.70, 0.60, 0.50, 0.45 from left to right. Total masses for Y=0.8, 0.7 and 0.6 are 0.46, 0.47, 0.48, 0.49, 0.50 and 0.51\\msun. For Y=0.45 we have added M=0.55\\msun. The core mass is fixed at \\Mc=0.455\\msun. For comparison, we add the ZAHB and the corresponding (dash--dotted) tracks, having Y=0.32, total masses 0.472, 0.475, 0.48, 0.49, 0.50, 0.51, 0.55 and 0.60\\msun\\ from left to right, and \\Mc=0.469\\msun: these tracks do not show the curling shape at any mass. The transition from pure--core helium burning tracks to tracks in which also hydrogen burning is efficient is at $\\sim$0.53\\msun. On the right side we plot (top panel) the \\% of luminosity provided by the core helium burning for a mass of 0.51\\msun, with Y=0.24, 0.32, 0.5, 0.60, 0.70 and 0.8. Decreasing Y, the hydrogen shell burning contributes a larger fraction of luminosity for a longer time. If Y$<$0.5, however, the hydrogen shell does not ignite any longer for this mass. In the bottom panel, we see that, for Y$\\geq$0.5, the models spend most of their lifetime in the ``curling\" region (see left panel and text).} \\label{f1}% \\end{figure*} \\begin{figure} \\includegraphics[width=8cm]{f2.eps} \\caption{We plot the ratio {\\sl q} of the core mass to total mass as a function of time for the Y=0.80 tracks of 0.46, 0.47, 0.48, 0.50, 0.51, 0.52, 0.53 and 0.55\\Msun\\ (from top to bottom). When the H--burning begins to be efficient, {\\sl q} rapidly increases due to the efficiency of the H--shell, giving the same values of {\\sl q} for all these masses. Thanks to the same {\\sl q}, the tracks remain at the same \\teff, at increasing luminosity, due to the larger total mass. The lowest track plotted shows the {\\sl q} for M=0.55\\Msun\\ and Y=0.45: the lower shell efficiency does not allow the fast increase of the core mass. } \\label{f1bis}% \\end{figure} In Section 2 we present new models of helium core burning stars of metallicity adequate to describe \\ocen's stars, characterized by extremely high values of helium abundance (from Y=0.45 to Y=0.8) around the burning helium core. While, in Section 3, we will give reasonable but unproven justifications for adopting these extravagant compositions, these models are the first that can reproduce the vertical shape of the blue hook stars in \\ocen, as the ZAHB of these models is ``vertical\", and not ``horizontal\" as it happens for less helium rich compositions. Simulations of the blue hook star population allow us to speculate on the evolutionary path leading to this kind of stellar structures. ", "conclusions": "" }, "1003/1003.0866_arXiv.txt": { "abstract": "The Vulpecula OB association, \\vul, is a region of active star formation located in the Galactic plane at 2.3$\\;$kpc from the Sun. Previous studies suggest that sequential star formation is propagating along this 100$\\;$pc long molecular complex. In this paper, we use \\emph{Spitzer} MIPSGAL and GLIMPSE data to reconstruct the star formation history of \\vul, and search for signatures of past triggering events. We make a census of Young Stellar Objects (YSO) in \\vul~based on IR color and magnitude criteria, and we rely on the properties and nature of these YSOs to trace recent episodes of massive star formation. We find 856 YSO candidates, and show that the evolutionary stage of the YSO population in \\vul~is rather homogeneous - ruling out the scenario of propagating star formation. We estimate the current star formation efficiency to be $\\sim8$\\%. We also report the discovery of a dozen pillar-like structures, which are confirmed to be sites of small scale triggered star formation. ", "introduction": "The physical mechanisms describing stellar birth are fairly well understood for low- and intermediate-mass stars but still under debate for their high-mass analogues \\citep{mckee07}. On larger scales, observations show that star forming mechanisms are of relatively poor efficiency, as only a fraction of the gas reservoir in the Universe is turning into stars. Typical Star Formation Efficiencies (SFE) are of the order 3-6\\% in the Galaxy \\citep{evans09}, and 5\\% or less in other galaxies \\citep{rownd99}. Still, in extreme environments like in the starburst galaxy Arp~220, the SFE can reach 50\\% \\citep{anan00} suggesting that the star forming mechanism at work in such objects could be of a different nature. For instance the feedback into the interstellar medium (ISM) from short-lived massive stars seems to influence the yield of star formation in their local environment. As they evolve off the main sequence, high-mass stars produce a copious amount of energy while still embedded in their native cocoon; such disruption of a molecular cloud leads to gravitational instabilities and possibly to the onset of a new episode of star formation. \\citet{hosokawa06} showed that under certain conditions runaway triggering can take place around massive OB stars. In other cases however, turbulence and magnetic fields can have a negative feedback on the local ISM and lead to the suppression of the star forming activity \\citep{price09,stone70}. \\citet{elmegreen98} provides an exhaustive review of the theoretical framework to study the dynamical triggers of star formation. In this context, we use the data from the \\emph{Spitzer} Legacy surveys MIPSGAL \\citep{carey09} and GLIMPSE \\citep{benjamin03} to study the star forming activity currently taking place in the Vulpecula OB association (hereafter \\vul). \\vul~hosts nearly one hundred OB~stars and three bright \\hii regions known as Sharpless objects 86, 87 and 88 \\citep{sharpless59}. According to \\citet{ehlerova01} and \\citet{turner86}, the star forming activity occurring in \\vul~might have been triggered by a common external source, and star formation might be propagating from one \\hii region to another through the expansion of a supernova shock front. The aim of this paper is to study the triggered star formation on scales as large as $\\sim$100~pc in \\vul, which could help us understand star formation mechanisms within other Giant Molecular Clouds in the Galaxy and beyond. We use the method developed by \\citet{gutermuth08} to obtain a census of Young Stellar Objects (hereafter YSOs) in \\vul~based on infrared color and magnitude criteria, and we rely on the properties and nature of these YSOs to trace episodes of star formation. The MIPSGAL and GLIMPSE sensitivity makes our search for YSOs in \\vul~biased towards massive objects that reveal the most recent star forming activity. In Section~\\ref{sec:obs_cat} we present the dataset used in the paper. In Section~\\ref{sec:vulOB1} we give a comprehensive description of \\vul~encompassing the three \\hii regions as well as the dozen pillar-like structures we discovered in this region. Section~\\ref{sec:YSO} describes the identification process of Young Stellar Objects based on their IR-excess emission. We present and discuss our results in Section~\\ref{sec:result}, and we give our conclusions in Section~\\ref{sec:conclu}. ", "conclusions": "\\label{sec:conclu} We have presented a thorough description of the Vulpecula OB association. We have complemented the existing observations of isolated objects in \\vul~with a global view of the whole star forming complex from an infrared perspective. We exploited \\emph{Spitzer} legacy surveys MIPSGAL and GLIMPSE data to identify 856 young stellar objects with IR-excess emission. We rely on the nature and properties of these objects to highlight the recent activity of star formation in the complex, and we look for evidences of triggered star formation. We find two populations of YSO candidates: one population of distributed objects that likely contains IR-bright evolved stars and some genuine YSOs born in isolation or ejected from their parental cloud; and another population of clustered YSO candidates whose spatial distribution correlates very well with the mid-IR morphology of the complex. YSO candidates surface density peaks locally around the three Sharpless objects, the Hyper Compact \\hii regions, and other embedded star clusters like Cr404. The vigorous star forming activity around these energetic sources is consistent with scenarios of triggered star formation mechanisms. Still we cannot ascertain that these stars were born as a direct consequence of their extreme environment, nor that they would have formed in the absence of what we consider the triggering agents. Nevertheless our analysis allowed to rule out the scenario of \\citet{ehlerova01} according to which star formation was propagating from Sh2-88 to Sh2-86. We rather find that the evolutionary stage of the YSO population across \\vul~is homogeneous, and thus consistent with the scenario of \\citet{turner86}. We have further reported the discovery of a dozen pillar-like structures in \\vul, and we comment on their morphology from the near-IR to the radio regime. We were not able to identify the energetic source(s) responsible for the molding of the pillars, but we argue that these objects are indeed associated with the OB~association. Our finding of YSO candidates at the tip of most of the pillars is consistent with mechanisms of triggered star formation on small scales." }, "1003/1003.2863_arXiv.txt": { "abstract": "We investigate the evolution of 5 granular-scale magnetic flux cancellations just outside the moat region of a sunspot by using accurate spectropolarimetric measurements and G-band images with the Solar Optical Telescope aboard \\textit{Hinode}. The opposite polarity magnetic elements approach a junction of the intergranular lanes and then they collide with each other there. The intergranular junction has strong red shifts, darker intensities than the regular intergranular lanes, and surface converging flows. This clearly confirms that the converging and downward convective motions are essential for the approaching process of the opposite-polarity magnetic elements. However, motion of the approaching magnetic elements does not always match with their surrounding surface flow patterns in our observations. This suggests that, in addition to the surface flows, subsurface downward convective motions and subsurface magnetic connectivities are important for understanding the approach and collision of the opposite polarity elements observed in the photosphere. We find that the horizontal magnetic field appears between the canceling opposite polarity elements in only one event. The horizontal fields are observed along the intergranular lanes with Doppler red shifts. This cancellation is most probably a result of the submergence (retraction) of low-lying photospheric magnetic flux. In the other 4 events, the horizontal field is not observed between the opposite polarity elements at any time when they approach and cancel each other. These approaching magnetic elements are more concentrated rather than gradually diffused, and they have nearly vertical fields even while they are in contact each other. We thus infer that the actual flux cancellation is highly time dependent events at scales less than a pixel of \\textit{Hinode} SOT (about 200 km) near the solar surface. ", "introduction": "The magnetic fields in the photosphere are highly heterogeneous, as is well known \\citep[e.g.][]{Parker1979, Zwaan1987, Stenflo1989}. These fields exist in a dynamical state in a spectrum of sizes, ranging from the kilo-gauss fibrils with diameters of a few hundred kilo-meters, through the mesoscale pores, to sunspots of various sizes. In the turbulent convection zone immediately below the quiet-Sun photosphere, the magnetic field is dominated by the fluid. Therefore, it is expected that the field distribution is organized into the convective cells. This field distribution is necessarily complex with flux connectivities of various scales, across individual and multiple cells. Magnetic concentrations that have vertical and strong (kilo-gauss) field are observed along the boundaries of the convective cells. A most prominent feature in the quiet Sun is the network magnetic field that partially outlines supergranular cells. It is believed that the network magnetic field is formed by the advection of internetwork fields via the supergranular flows. Magnetic elements in internetwork areas tend to move toward the nearest network concentration at a speed of about 0.2 km s$^{-1}$ \\citep{deWijn2008}, and their rms velocity is about 1.5 km s$^{-1}$ in internetwork areas \\citep{Nisenson2003, deWijn2008}. Velocities of 0.2 km s$^{-1}$ and 1.5 km s$^{-1}$ are similar to a typical speed of supergranular flows and granular flows, respectively. Down at the scales of a few hundred kilo-meters, equal to a few photospheric density scale heights, magnetic elements of both polarities form to often merge and collide with each other in the photospheric fluid. The net magnetic flux increases or decreases, depending on whether like-polarity or opposite-polarity magnetic elements have merged. In this paper, we report an interesting set of ``magnetic flux cancellation'' events of this kind observed at the high spatial resolution of the Solar Optical Telescope \\citep[SOT;][]{Tsuneta2008} aboard \\textit{Hinode} \\citep{Kosugi2007}. The magnetic flux cancellation is a descriptive term to indicate a mutual flux loss due to the apparent collision of the opposite-polarity magnetic elements \\citep{Martin1985}. This flux cancellation is essential to the process of replacement of old magnetic flux with newly emerging flux in the quiet Sun on a timescale of a few days \\citep{Schrijver1997, Hagenaar2001}, and also to the process of removal of sunspot magnetic flux from the photosphere \\citep{Kubo2008}. Various possible processes have been proposed to explain the observed flux cancellation \\citep[e.g.][]{Zwaan1987, Ryutova2003}, involving submergence (retract) of $\\Omega$-shaped loops or emergence of U-shaped loops across the photosphere. In both cases, the canceling opposite-polarity magnetic elements correspond to the two intersections of such loops with the photospheric layer. The opposite-polarity magnetic elements disappear when the top of a submerging $\\Omega$-loop has passed through the photospheric layer \\citep[see Fig.2 in][]{Zwaan1987}. Alternatively, these elements disappear when the bottom of a rising U-loop has passed clear through the photospheric layer. When magnetic field lines have emerged into the chromosphere and corona, they can hardly submerge back below the photosphere because of magnetic buoyancy. Magnetic reconnection taking place within several scale heights above the solar surface is probably needed to create low-lying $\\Omega$-loops whose magnetic tension force can then overcome the magnetic buoyancy force \\citep{Parker1975}. In the photospheric magnetic reconnection cases, reconnection should be most efficient around the temperature minimum region: about 600 km above the solar surface \\citep{Litvinenko1999, Takeuchi2001}. In contrast, both magnetic tension and buoyancy forces are directed upward in the case of a U-loop rising through the photosphere. Magnetic reconnection is therefore not crucial for the emerging U-loops. An important observable signature for understanding flux cancellation is the motion of the horizontal magnetic field connecting the canceling magnetic elements. Horizontal magnetic fields have been observed between the opposite-polarity magnetic elements during the cancellations of moving magnetic features around a sunspot \\citep{Chae2004}. Similar horizontal fields have also been observed in events of cancellations of pores and sunspots \\citep{Kubo2007}. A flux cancellation without increase of the horizontal field has also been reported for the moving magnetic features \\citep{Bellot2005}. Knowledge of the full vector field permits one to determine whether the field geometry has $\\Omega$-loop topology or U-loop topology, but in the case of cancellation of small, isolated flux elements, such a determination is often compromised by usual difficulty of resolving the 180$\\degr$ azimuth ambiguity. Regarding the motions at such a cancellation site, \\citet{Harvey1999} show that the magnetic flux disappears in the chromosphere before it does in the photosphere for at least about half of the cancellation events. They suggest that magnetic flux is submerging in most, if not all, of the cancellation sites. On the other hand, both Doppler red shift \\citep{Chae2004} and Doppler blue shift \\citep{Yurchyshyn2001} are reported in the cancellation sites. The center-to-limb variations of the Doppler velocities at the polarity inversion lines in the cancellation sites suggest that the observed velocities with the spatial resolution of about 1$\\arcsec$ mainly show the gas flows along the horizontal fields rather than actual emerging or submerging motions of the field lines \\citep{Kubo2007}. Observationally, the physical process of the magnetic flux cancellation is still not well understood. Recent observations with the high spatial resolution by \\textit{Hinode} SOT show that downward motions (red-shifts) are continuously observed during the flux cancellation process \\citep{Iida2010}. From the Stokes-V signals far from the disk center, they also suggest that the canceling opposite polarity elements tend to have an $\\Omega$-shaped configuration rather than a U-shaped configuration. In this paper, we investigate the detailed evolution of canceling magnetic elements and their surrounding convective motions near the disk center by using vector magnetic fields and velocity fields observed with the SOT. We are particularly interested in the evolution that precedes the cancellation in order to investigate why the opposite-polarity magnetic elements approach and collide with each other. For a full understanding of the phenomenon, this dynamical development is as important as the flux removal process at the cancellation site. ", "conclusions": "We have presented an observational study of flux cancellation events on the photosphere that are sufficiently resolved by the \\textit{Hinode} SOT to show that they are characterized by the absence of a horizontal field during the cancellation process. These events are interesting because in the usual idea of the submergence of a low-lying $\\Omega$-loop or the buoyant rise of a U-loop, the appearance of a horizontal field is the observational signature of the loop top (or bottom) passing across the photosphere. Such flux cancellations appear to be more common than the cancellation with the appearance of the horizontal fields. Although the nature of the flux removal process at the cancellation site is an open question for cancellations without the appearance of the horizontal fields, our study shows that it takes place in local areas at scales less than the $\\sim$200km resolution of the SOT and close to the solar surface. The distinction between the cancellations with and without the appearance of the horizontal fields might arise from whether the canceling opposite polarity elements have emerged into the photosphere as a pair or not. However, we have investigated only 5 events that have the size less than a few arcseconds just outside the moat region of a sunspot. We need to investigate what really determines the appearance of the horizontal fields between the canceling magnetic elements by using more events. In particular, the origin of the canceling magnetic elements and their surrounding magnetic field configuration (formation of pre-existing horizontal fields) may be important. We have confirmed the converging and downward convective flows are essential for the approaching and canceling process of the opposite-polarity magnetic elements. The flux cancellations are observed at the intergranular junction characterized by the strong converging and downward flows. However, the approaching opposite-polarity elements seem not to always follow the surrounding surface flow patterns, at least in our observations. We are accustomed to thinking of the connectivities of the surface magnetic fields that we can observe in the solar atmosphere. Our observational study suggests that it is important to also consider the connectivities of the surface fields that occur below the photosphere. Lanes of downdrafts of cool fluids entraining magnetic fluxes that thread across these lanes below the photosphere must be a common occurrence. Such a process is a simple explanation for a pair of opposite-polarity elements to appear to seek each other on the photosphere. Information of subsurface convective flows would be extremely helpful for better understanding of the magnetic flux cancellation." }, "1003/1003.3442_arXiv.txt": { "abstract": "Observations indicate that roughly 60\\% of the baryons may exist in a Warm-Hot Intergalactic Medium (\\WHIM) at low redshifts. Following up on previous results showing that gas is released through galaxy mergers, we use a semi-analytic technique to estimate the fraction of gas mass lost from haloes solely due to mergers. We find that up to $\\sim$ 25\\% of the gas in a halo can unbind over the course of galaxy assembly. This process does not act preferentially on smaller mass haloes; bigger haloes {\\emph {always}} release larger amounts of gas in a given volume of the Universe. However, if we include multi-phase gas accretion onto haloes, we find that only a few percent is unbound. We conclude that either non-gravitational processes may be in play to heat up the gas in the galaxies prior to unbinding by mergers or most of the baryons in the \\WHIM have never fallen into virialised dark matter haloes. We present a budget for stocking the \\WHIM compiled from recent work. ", "introduction": "The baryon budget shows significant evolution from $z \\sim 3$, and results in an apparent baryon deficit today \\citep{FHP98,FP04}. At high redshift, most of the baryonic mass is in the Ly-$\\alpha$ forest \\citep{FHP98,FP04}, while at low redshifts over half of the baryons are as yet undetected. The consensus is that the majority of the `missing' baryons are actually in regions of low overdensity, $\\delta\\rho/\\rho \\sim 10-100 $ \\citep[e.g][]{CO99,DHKW99,DC01,KRCS05,CO06,DM06,DO07} with temperatures in the range $10^5-10^7$ K -- commonly referred to as the Warm-Hot Intergalactic Medium (\\WHIM). The immediate question is: how is this \\WHIM produced~? Some form of mass and energy injection is essential to create this hot reservoir of gas; this form of feedback must both regulate the gas in galaxies and the metal content of the Intergalactic Medium (\\IGM). There has been much numerical work to incorporate various feedback mechanisms in an attempt to solve this puzzle \\citep[e.g.,][and references therein]{CO99,NS01,DC01,KRCS05,CO06,DO07}. Cosmological simulations seem to suggest that gravitational collapse during galaxy formation can produce and maintain the majority of the \\WHIM at $10^5 - 10^7$ K \\citep{CO99,DHKW99,CO06,CDM01,DO07}. Supernova feedback provides another avenue to generate the \\WHIM; for star bursts of 100 $M_\\odot$ per year, as much as $20\\%$ of the hot gas in a Milky Way mass galaxy can be unbound \\citep{STWS06,KSW07}. However, SNe feedback may be a self regulating process, in that a starburst also heats the remaining gas and may damp the star formation rate, which in turn would reduce the fraction of unbound gas \\citep[e.g.][]{STWS08}. Combining these effects, it is commonly thought that galaxies with host halo mass $\\gtrsim 10^{11} $ M$_{\\odot}$ lose $\\lesssim 10\\%$ of their gas through SN feedback, while low mass haloes may be entirely depleted of gas by this mechanism\\citep{YK97,MF99,E00,STWS06}. A third possibility is that the radiation from an accreting supermassive black hole could power large-scale winds to blow mass out of the galaxy \\citep[see][]{SO04,MQT05,HH05,HHM05,HH06,CS06,SSMH07}. For a fixed amount of energy, all the non-gravitational feedback mechanisms are more effective in low mass galaxies due to their shallower potential. However, observations suggest that low-mass galaxies are in general more gas-rich and are less likely to have suffered a gas blow-out \\citep{K04,GBMW06}. In \\citet[][hereafter, SH09]{SH09}, we show that hot gas is driven into the \\IGM by galaxy mergers. The amount of hot halo gas lost depends strongly on the energy of the merger; it is possible for low mass galaxies to retain their gas in this scenario during low-speed or distant encounters. However, SH09 only estimated the mass lost during a single merger. When all the mergers in the Universe are considered, this could heat and drive a significant portion of the total baryon budget into the \\WHIM. In principle, this process could join \\AGN and star formation feedback as a way to populate the \\WHIM, and we find that this method operates preferentially in {\\em more} massive haloes. To estimate the total fraction of gas released by mergers, we construct a series of analytic halo merger trees using a publicly available\\footnotemark semi-analytic Extended Press-Schechter (\\EPS) code \\citep{PCH08}. \\footnotetext{\\url{http://star-www.dur.ac.uk/~cole/merger_trees/}} In Section~\\ref{section:eps} we describe the theory of halo merger trees, in Section~\\ref{section:method} we outline the experiments designed to track the gas ejected via galaxy mergers, in Section~\\ref{section:results} we present the results for the halo gas ejected by this process and Section~\\ref{section:discussion} contains the discussion. ", "conclusions": "\\label{section:discussion} In this paper we show that a significant portion of the \\WHIM can be generated by gas ejected from galaxies during mergers. Our semi-analytic prescription shows that up to $\\sim$ 25\\% of the gas (assuming universal gas fraction) in haloes of mass $10^8 - 10^{13}\\, \\Msun$ can be ejected by mergers. Given an observed gas mass at $z=0$, it is possible to infer the typical gas mass that was unbound from assembling that halo (column 4, Table~\\ref{table:results}) For comparison with SN feedback, a quiescent Milky-Way type halo with star formation rate of 1-10 $\\Msun{\\rm yr^{-1}}$ would unbind $\\leq 2\\%$ of the gas content \\citep{STWS08}. We also find that multiphase gas accretion drastically reduces the amount of unbound gas from mergers, down to a few percent of the gas mass. In contrast with previous numerical work involving non-gravitational feedback, where the effects of mass loss are severe in smaller haloes, this merger mechanism unbinds gas preferentially from massive haloes. There is no selective unbinding of gas from dwarf galaxies, in line with observational evidence suggesting that dwarf galaxies are more gas-rich and therefore may not have suffered a gas blow-out \\citep{K04,GBMW06}. This form of {\\it gravitational feedback} may even play a larger role in regulating the stellar mass function: \\citet{K09a} show that simulated galaxies exhibit a discrepancy with the observed stellar mass function \\citep{B03} for both high and the low mass galaxies. They conjecture that the key to solving this discrepancy is through a feedback mechanism that works for halos $\\gtrsim 10^{12}\\, \\Msun$ -- akin to our scenario. If the merger-ejection process is very efficient, then the current day haloes may be very gas-poor. It is conceivable that the current stellar mass {\\em and} the gas fraction of a galaxy constrains the mean stellar mass and gas content of the galaxies of the past. Overall, we find that for our most reasonable scenario, $\\sim 15\\%$ of the \\WHIM can be generated through galaxy mergers. If previous work on large-scale gravitational shocks proves correct \\citep{CO99, DC01}, $\\sim 66\\%$ of the \\WHIM can be attributed to gas that may have never fallen into a halo. In addition, recent studies have shown that roughly 20\\% can be produced via non-gravitational feedback, such as SNe and AGN \\citep[e.g.,][]{CO06}. Therefore, with these three mechanisms to populate the \\WHIM, it may well be true that the baryon budget is balanced." }, "1003/1003.4262_arXiv.txt": { "abstract": "The Birmingham Solar-Oscillations Network (BiSON) has collected helioseismic data over three solar cycles. We use these data to determine how the internal properties of the Sun during this minimum differ from the previous two minima. The cycle 24 data show oscillatory differences with respect to the other two sets, indicating relatively localized changes in the solar interior. Analysis of MDI data from Cycle 23 and Cycle 24 also show significant signs of differences. ", "introduction": "\\label{sec:intro} The minimum before Solar Cycle 24 has been much quieter than many before, with the lowest sustained 10.7cm radio flux since observation of this proxy began in 1947. Other differences have been observed too, as has been discussed in detail in other articles of this volume. Most of the discussions concern observations at or above the solar surface. In this paper we try to determine whether or not there were changes in the solar interior between this minimum and others before it using helioseismic data. The Birmingham Solar-Oscillations Network (BiSON) has been collecting helioseismic data for over three solar cycles (Broomhall et al. 2009). No other helioseismic data set spans such a long time interval. We analyze data for one-year periods starting 1986/04/01, 1995/11/01 and 2008/05/01 to determine whether there are detectable differences between the thermal structure of the Sun during these epochs. These epochs corresponds to the minima of cycles 22, 23 and 24 respectively. ", "conclusions": "Using data from the Birmingham Solar Oscillation Network and the Michelson Doppler Imager on board {\\it SoHO}, we find evidence that suggests the structure of the Sun may have been different during the the minimum of Cycle 24 compared with the minimum of Cycle 23. There is some evidence that it was different compared to the minimum of Cycle 22 too, however, Cycle 22 data need to be re-analyzed before we can make firm conclusions. The evidence points to a lowered sound-speed during Cycle 24 in layers deeper than 0.98R$_\\odot$." }, "1003/1003.4748_arXiv.txt": { "abstract": "With present and upcoming SUSY searches both directly, indirectly and at accelerators, the need for accurate calculations is large. We will here go through some of the tools available both from a dark matter point of view and at accelerators. For natural reasons, we will focus on public tools, even though there are some rather sophisticated private tools as well. ", "introduction": "The general theory behind relic density calculations of dark matter particles is given in Chapter 7 \\cite{Chap:Gelmini}. Here we will focus on supersymmetric dark matter (neutralinos) and various tools available for calculating the relic density. There are currently three publicly available codes for calculating the relic density of neutralinos: {\\tt DarkSUSY}, {\\tt micrOMEGAs} and {\\tt IsaRED}\\footnote{After this chapter was completed, a fourth code, {\\tt SuperIso Relic} \\cite{Arbey:2009gu} has appeared as well.}. All three of these codes are capable of reading (and sometimes writing) SUSY Les Houches Accord (SLHA) files~\\cite{Skands:2003cj,Allanach:2008qq} which allows for an easy interface between these codes and other tools to be described in Section~\\ref{section-tools-collider}. We will in the following refer to these codes and how they calculate the relic density. We will use the notation of Chapter 7 \\cite{Chap:Gelmini} and only write down the equations needed to facilitate our discussions here. \\subsection{The Boltzmann equation} In most supersymmetric models of interest for dark matter phenomenology, the lightest neutralino, $\\chi$ is the lightest supersymmetric particle and our dark matter candidate. As such, we want to calculate the relic density of neutralinos in the Universe today as accurately as possible, which means that we need to solve the Boltzmann equation. \\begin{equation} \\frac{dn_\\chi}{dt}=-3Hn_\\chi-\\langle \\sigma_{\\rm ann} v \\rangle \\left( n_\\chi^2 - n_{\\chi,\\,\\rm eq}^2 \\right) \\label{eq:boltzmann} \\end{equation} where $n_\\chi$ is the number density of neutralinos, $H$ is the Hubble parameter, $\\langle \\sigma_{\\rm ann} v \\rangle$ is the thermally averaged annihilation cross section and $n_{\\chi,\\,\\rm eq}$ is the equilibrium number density of neutralinos. This equation needs to be solved over time (or temperature) properly calculating the thermal average at each time step. When the neutralinos no longer can follow the chemical equilibrium density $n_{\\chi,\\,\\rm eq}$, they are said to freeze-out. There are several complications in solving Eq.~(\\ref{eq:boltzmann}); for example, we may have resonances and thresholds in our annihilation cross section. The solution to this is to calculate the annihilation cross section in general relativistic form, for arbitrary relative velocities, $v$. Another complication is that other supersymmetric particles of similar mass will be present during freeze-out of the neutralinos. To solve this we need to take into account the so-called coannihilations between all the SUSY particles that are almost degenerate in mass with the neutralino (in practice, it is often enough to consider coannihilations between all SUSY particles up to about 50\\% heavier than the neutralino). Following the discussion in Chapter~7~\\cite{Chap:Gelmini}, we can solve for the total number density of SUSY particles, \\begin{equation} n= \\sum_{i=1}^N n_{i} \\end{equation} instead of only the neutralino number density. It is also advantageous to rephrase the Boltzmann equation in terms of the abundance $Y=n/s$ and use $x=m_\\chi/T$ as independent variable instead of time or temperature. When coannihilations are included, the Boltzmann equation (\\ref{eq:boltzmann}) can then be written as \\begin{equation} \\frac{dY}{dx}=-\\left(\\frac{45}{\\pi M_P^2}\\right)^{-1/2}\\frac{g_*^{1/2}m_\\chi}{x^2}\\left<\\sigma_{\\rm eff} v\\right> (Y^2-Y_{eq}^2)\\,. \\label{eq:Boltzco} \\end{equation} where $\\langle \\sigma_{\\rm eff} v \\rangle$ is given by \\cite[Eq.~(7.18)]{Chap:Gelmini} \\subsection{Solving the Boltzmann equation} To solve the Boltzmann equation (\\ref{eq:Boltzco}) we need to calculate the thermally averaged annihilation cross section $\\langle \\sigma_{\\rm eff} v \\rangle$ for each given time (temperature). This is typically quite CPU-intensive, and we therefore need to use some tricks. In {\\tt DarkSUSY} ~\\cite{Gondolo:2004sc}, the solution is speeded up by tabulating $W_{\\rm eff}$ in \\cite[Eq.~(7.19)]{Chap:Gelmini} , but using the momentum of the $\\chi$, $p_{\\rm eff}$ as independent variable instead of $s$. This tabulation takes extra care of thresholds and resonances making sure that they are tabulated properly. This tabulated $W_{\\rm eff}$ is then used to calculate the thermal average $\\langle \\sigma_{\\rm eff} v \\rangle$ for each time (temperature), using \\cite[Eq.~(7.18)]{Chap:Gelmini} . The advantage with this method is that $W_{\\rm eff}$ does not depend on temperature and instead the temperature dependence of $\\langle \\sigma_{\\rm eff} v \\rangle$ is completely taken care of by the other factors in \\cite[Eq.~(7.18)]{Chap:Gelmini}. Numerically, one needs to take special care of the modified Bessel functions $K_1$ and $K_2$ which both contain exponentials that need to be handled separately to avoid numerical underflows. The Boltzmann Equation (\\ref{eq:Boltzco}) is then solved with a special implicit method with adaptive stepsize control, which is needed because the equation is stiff and develops numerical instabilities unless an implicit method is used. The details of the DarkSUSY method are as follows (see {\\tt DarkSUSY} manual~\\cite{Gondolo:2004sc}). The derivative $dY/dx$ in Eq.~(\\ref{eq:Boltzco}) is replaced with a finite difference $(Y_{n+1}\\!-\\!Y_n)/h$, where $Y_n=Y(x_n)$ and $x_{n+1}=x_n+h$. Then $Y_{n+1}$ is computed in two ways: first, the right hand side of Eq.~(\\ref{eq:Boltzco}) is approximated with $-\\tfrac{1}{2}[\\lambda_{n}(Y_{n}^2-Y_{eq,n}^2)+\\lambda_{n+1}(Y_{n+1}^2-Y_{eq,n+1}^2)]$, where $\\lambda(x)=\\left(45/\\pi M_P^2\\right)^{-1/2} (g_*^{1/2}m/x^2)\\left<\\sigma_{\\rm eff} v\\right>$, and an analytic solution is used for the resulting second-degree algebraic equation in $Y_{n+1}$; second, the right hand side of Eq.~(\\ref{eq:Boltzco}) is approximated with $-\\lambda_{n+1}(Y_{n+1}^2-Y_{eq,n+1}^2)$ and an analytic solution of the algebraic equation for $Y_{n+1}$ is used. The stepsize $h$ is reduced or increased to maintain the difference between the two approximate values of $Y_{n+1}$ within a specified error. Overall, the solution of the Boltzmann equation and the tabulation of $W_{\\rm eff}$ solves for the relic density to within about 1\\%. If needed, higher accurary can also be chosen as an option. In {\\tt micrOMEGAs}~\\cite{Belanger:2006is}, $\\langle \\sigma_{\\rm ann} v \\rangle$ at a given temperature $T$ is arrived at by performing a direct integration and does not therefore rely on a tabulation of the matrix elements squared. Two modes are provided to perform the integration. In the {\\it accurate} mode the program evaluates all integrals by means of an adaptive Simpson integration routine. It automatically detects all singularities of the integrands and checks the precision of the calculation increasing the number of points until an accuracy of $10^{-3}$ is reached. In the default mode ({\\it fast} mode) the accuracy is not checked but a set of points is sampled according to the behaviour of the integrand: poles, thresholds and Boltzmann suppression at high momentum. The first integral over scattering angles is performed by means of a $5$-point Gauss formula. The accuracy of this mode is generally about $1\\%$. The user can also test the precision of the approximation based on expanding the cross section in terms of its $s$ and $p$-wave components. In the Boltzmann Equation, we need to know $g_{\\rm eff}(T)$ and $h_{\\rm eff}(T)$ that enter $g_*^{1/2}$ through \\cite[Eq.~(7.9)]{Chap:Gelmini}. In {\\tt DarkSUSY}, the default is to use the estimates in Ref.~\\cite{Hindmarsh:2005ix}, but other options are also available. Typically, different estimates of $g_{\\rm eff}(T)$ and $h_{\\rm eff}(T)$ translate into relic densities different by a few percent. {\\tt IsaRED} on the other hand does not solve the Boltzmann equation numerically, instead it finds the freeze-out temperature (the temperature when the annihilation rate equals the expansion rate of the Universe) and calculates the relic density from that (including remnant annihilations at later times). For the thermal averaged annihilation cross section, it uses the same relativistic treatment as {\\tt DarkSUSY} and {\\tt micrOMEGAs}. \\subsection{Coannihilation criteria} \\begin{figure} \\centerline{\\includegraphics[width=0.75\\textwidth]{k1rateex.eps}} \\caption{Total differential annihilation rate per unit volume $dA/dp_{\\rm eff}$ for the same model as in \\cite[Fig.~7.3]{Chap:Gelmini}, evaluated at a temperature $T=m_\\chi/20$, typical of freeze-out. Notice the Boltzmann suppression at high $p_{\\rm eff}$.} \\label{fig:k1effrate} \\end{figure} In principle, one should include all SUSY particle coannihilations when calculating the relic density. However, the heavier they are, the less abundant they will be and can thus be neglected. This is important to speed up the calculation, as we would otherwise spend most of our CPU cycles on calculating non-important coannihilation cross sections. One can estimate which particles to include, by investigating \\cite[Eq.~(7.18)]{Chap:Gelmini}. The modified Bessel function $K_1$ contains an exponential, the so called Boltzmann suppression, that will suppress all heavier particles. In Fig.~\\ref{fig:k1effrate}, we show $dA/dp_{\\rm eff}$ for the same model as in \\cite[Fig.~7.3]{Chap:Gelmini}. $dA/dp_{\\rm eff}$ is essentially (apart from normalization) the integrand in the numerator in \\cite[Eq.~(7.18)]{Chap:Gelmini}. Comparing the two figures, we clearly see the Boltzmann suppression of larger $p_{\\rm eff}$, i.e.\\ heavier coannihilating particles. We can quantify this by comparing the Boltzmann suppression factor for two coannihilating particles with masses $m_i$ and $m_j$ with the corresponding factor for the LSP, the $\\chi$. The suppression factor for the coannihilating particles compared to the $\\chi$ is roughly (neglecting the $p_{\\rm eff}^2$ in \\cite[Eq.~(7.18)]{Chap:Gelmini}) given by \\begin{equation} B = \\frac{K_1((m_i+m_j)/T)}{K_1(2m_\\chi/T)} \\simeq e^{-x \\frac{m_i+m_j-2m_\\chi}{m_\\chi}} \\quad ; \\quad x=\\frac{m_\\chi}{T} \\end{equation} At freeze-out we typically have $x\\simeq 20$, which gives a suppression factor of $B\\simeq 10^{-6}$ for coannihilation particles about 50\\% heavier than the LSP, $\\chi$. In {\\tt DarkSUSY}, one can set the maximum mass fraction of coannihilation particles, $f=m_i/m_\\chi$ that will be included in the calculation, whereas in {\\tt micrOMEGAs} one sets the minimum $B$ to allow instead. The defaults in {\\tt DarkSUSY} ($f=1.5$) and {\\tt micrOMEGAs} ($B_{\\rm min} = 10^{-6}$) are roughly equivalent. One should remember, though, that the value to choose depends on the particle physics model. For example, for chargino coannihilations, the $\\chi^+ \\chi^-$ coannihilation cross section can be orders of magnitudes larger than the $\\chi \\chi$ annihilation cross sections and one should choose $f$ or $B_{\\rm min}$ so that one does not accidentally neglect coannihilations that are important. For the MSSM, the default values of {\\tt DarkSUSY} and {\\tt micrOMEGAs} are typically sufficient for all interesting cases. {\\tt IsaRED} instead includes a preset collection of particles that are of relevance for the mSUGRA setup. \\subsection{Annihilation cross section} At the heart of the relic density calculation are the annihilation and coannihilation cross sections. In the MSSM there are over 2800 sub-processes (not counting charged-conjugate final states) that can in principle contribute in the relic density calculation. It appears at first sight to be a daunting task to provide such a general code. In {\\tt DarkSUSY}, all annihilation and coannihilation cross sections for the MSSM\\footnote{Gluino coannihilations are currently not included.} are calculated at tree level by hand with the help of symbolic programs like {\\tt Reduce}, {\\tt Form} or {\\tt Mathematica}. The calculations are performed with general expressions for the vertices in the Feynman rules and the results are converted to {\\tt Fortran} code. The vertices are then calculated numerically for any given MSSM model. The analytically calculated cross sections are differential in the angle of the outgoing particles, and the integration over the outgoing angle is performed numerically. In {\\tt micrOMEGAs}, on the other hand, any annihilation and coannihilation cross sections are calculated automatically and generated on the fly. This is possible thanks to an interface to {\\tt CalcHEP}~\\cite{Pukhov:2004ca}, which is an automatic matrix element/cross sections generator. This automation is carried one step further in that {\\tt CalcHEP} itself reads its MSSM model file (Feynman rules) from {\\tt LanHEP}~\\cite{Semenov:1998eb}, which outputs the complete set of Feynman rules from a simple coding of the Lagrangian, see section~\\ref{section-tools-collider}. In the first call to {\\tt micrOMEGAs} only those subprocesses needed for the given set of the MSSM parameters are generated. The corresponding ``shared\" library is stored on the user disk space and is accessible for all subsequent calls, thus each process is generated and compiled only once. This library is then filled with more and more processes whenever the user needs new processes for different MSSM scenarios. This avoids having to distribute a huge code with all the possible 2800 processes. Both methods have advantages and disadvantages. In the {\\tt DarkSUSY} setup, no recalculation of the (analytical) annihilation cross sections is needed, which can speed things up. Also, the analytically calculated annihilation cross sections can be optimized to be faster. On the other hand, the {\\tt micrOMEGAs} setup makes it easier to adapt the code to non-MSSM cases. In both codes though, the actual Boltzmann equation solver is very general and works for any kind of WIMP dark matter, not only SUSY dark matter. In {\\tt IsaRED}, {\\tt CompHEP}~\\cite{Boos:2004kh} is used to calculate the annihilation and coannihilation cross sections for a subset of SUSY particles of relevance mostly for mSUGRA (the two lightest neutralinos, the lightest chargino, the left-handed eigenstates of sleptons and squarks, and gluinos). The expressions for the annihilation cross sections in {\\tt IsaRED} are not calculated on the fly, but are instead precalculated and included with the code. ", "conclusions": "" }, "1003/1003.3732_arXiv.txt": { "abstract": "{Infrared dark clouds (IRDCs) likely represent very early stages of high-mass star/star cluster formation.} {The aim is to determine the physical properties and spatial distribution of dense clumps in the IRDC MSXDC G304.74+01.32 (G304.74), and bring these characteristics into relation to theories concerning the origin of IRDCs and their fragmentation into clumps and star-forming cores.} {G304.74 was mapped in the 870 $\\mu$m dust continuum with the LABOCA bolometer on APEX. The 870 $\\mu$m map was compared with the 1.2 mm continuum map of the cloud by B\\'eltran et al. (2006). Archival MSX and IRAS infrared data were used to study the nature and properties of the submillimetre clumps within the cloud. The H$_2$ column densities were estimated using both the 870 $\\mu$m dust emission and the MSX 8 $\\mu$m extinction data. The obtained values were compared with near-infrared extinction which could be estimated along a few lines of sight. We compared the clump masses and their spatial distribution in G304.74 with those in several other recently studied IRDCs.} {Twelve clumps were identified from the 870 $\\mu$m dust continuum map. Three of them are associated with the MSX and IRAS point sources. Moreover, one of the clumps (SMM 6) is associated with two MSX 8 $\\mu$m point-like sources. Thus, there are 8 clumps within G304.74 which are not associated with mid-infrared (MIR) emission. The H$_2$ column densities derived from the dust continuum and extinction data are similar. The comparison suggests that the dust temperature may be elevated (20-30 K) near the southern end of the cloud, whereas the starless clumps in the centre and in the north are cool ($T_{\\rm d}\\sim15$ K). There is a high likelihood that the clump mass distributions in G304.74 and in several other IRDCs represent the samples of the same parent distribution. In most cases the spatial distributions of clumps in IRDCs do not deviate significantly from random distributions.} {G304.74 contains several massive clumps that are not associated with MIR emission. On statistical grounds it is likely that some of them are or harbour high-mass starless cores (HMSCs). The fact that the clump mass distributions (resembling the high-mass stellar IMF), and in some cases also the random-like spatial distributions, seem to be comparable between different IRDCs, is consistent with the idea that the origin of IRDCs, and their further sub-fragmentation down to scales of clumps is caused by supersonic turbulence in accordance with results from giant molecular clouds.} ", "introduction": "\\subsection{Infrared dark clouds} The so-called infrared dark clouds (IRDCs) were discovered by the Infrared Space Observatory (ISO; \\cite{perault1996}) and the Midcourse Space Experiment (MSX; \\cite{egan1998}); the clouds were detected as dark absorption objects against the bright mid-infrared (MIR) Galactic background radiation. Based on the MSX 8 $\\mu$m data of the Galactic plane from $l=0-360\\degr$ and $\\left| b \\right|\\leq5\\degr$, Simon et al. (2006a) identified almost 11\\,000 IRDC candidates. Recently, Peretto \\& Fuller (2009) used Spitzer satellite data to extract about 9000 new IRDC candidates. Molecular line and dust continuum studies of IRDCs have shown that they are cold ($T<25$ K), dense ($n({\\rm H_2})\\gtrsim10^5$ cm$^{-3}$, $N({\\rm H_2})\\gtrsim10^{22}$ cm$^{-2}$), and massive ($\\sim10^2-10^5$ M$_{\\sun}$) structures with sizes of $\\sim1-15$ pc (e.g., \\cite{carey1998}; \\cite{simon2006b}; \\cite{rathborne2006}, hereafter RJS06; \\cite{du2008}; \\cite{vasyunina2009}). Most IRDCs are filamentary (e.g., \\cite{peretto2009}), and contain density enhancements, or clumps\\footnote{We prefer to use the term ``clump'' according to e.g., Ragan et al. (2009, hereafter RBG09), rather than the term ``core'' which was used by e.g., RJS06. Clumps have masses and sizes (radii) of $\\sim10-10^3$ M$_{\\sun}$ and $\\sim0.1-1$ pc, respectively (e.g., \\cite{bergin2007}).}, that are visible in (sub)millimetre dust continuum maps (e.g., \\cite{carey2000}; \\cite{garay2004}; \\cite{ormel2005}; RJS06). The cold clumps (i.e., clumps unassociated with MSX 8 $\\mu$m emission) identified by RJS06 have typical sizes and masses of $\\sim0.5$ pc and $\\sim120$ M$_{\\sun}$, respectively. Because IRDCs have clumpy structures, they are likely to be in an early stage of fragmentation (e.g., RBG09 and references therein). The radial galactocentric distribution of IRDCs peaks at $R_{\\rm GC}=5$ kpc in the 1st Galactic quadrant, and at $R_{\\rm GC}=6$ kpc in the 4th quadrant, which correspond to the location of the Scutum-Centaurus spiral arm (see \\cite{simon2006b} and \\cite{jackson2008}). This, together with the fact that IRDCs have sizes and masses similar to those of warm massive cluster-forming regions (e.g., \\cite{lada2003}; \\cite{motte2003}), has led to the suggestion that IRDCs represent the very early stages of high-mass star/star cluster formation. Indeed, several studies have found signs of ongoing star formation within IRDCs. These include CH$_3$OH and H$_2$O masers (\\cite{beuther2002b}; \\cite{pillai2006a}; \\cite{wang2006}; \\cite{ellingsen2006}), outflow signatures (\\cite{beuther2005}; \\cite{beuthersrid2007}; \\cite{sakai2008}), and associated infrared sources (e.g., \\cite{rathborne2005}; \\cite{beutherstein2007}; \\cite{chambers2009}; RBG09). Cold, dense clumps are suggested to host/represent precursors of {\\sl hot} molecular cores, i.e., high-mass starless cores (HMSCs; e.g., \\cite{sridharan2005}; \\cite{beuther2007}). A few hot cores have been found within IRDCs (\\cite{rathborne2007}, 2008), but it is important to increase the sample in order to study the sequence of events. It should be noted that the formation of hot cores may not be universal in IRDCs - some of them seem to form only low- to intermediate-mass stars (e.g., \\cite{vanderwiel2008}). \\subsection{Infrared dark cloud MSXDC G304.74+1.32} The IRDC studied in this paper was designated as MSXDC G304.74+1.32 (hereafter, G304.74) by Simon et al. (2006a). G304.74 was observed by Beltr{\\'a}n et al. (2006) with the SIMBA bolometer array on SEST at 1.2 mm. Beltr{\\'a}n et al. (2006) identified eight millimetre clumps within the cloud, out of which four were found to be associated with MSX point sources. The distance to the cloud, 2.4 kpc, is a kinematic distance estimated from the velocity of the CS line observed by Fontani et al. (2005). The galactocentric distance is $\\sim7.4$ kpc. G304.74 was chosen for the present study because its relatively close distance allows for reasonably good spatial resolution in order to study its substructure. Moreover, a relatively high number of clumps (8) were already identified from the cloud (see above), and thus it was considered an appropriate object for studying IRDC fragmentation. \\vspace{0.7cm} In this paper, we present the results of our 870 $\\mu$m dust continuum mapping of G304.74. The paper pays special attention to the clumpy structure of the cloud, and the clump mass distribution, and thereby addresses the fragmentation of IRDCs. The observations and data-reduction procedures are described in Sect. 2. The observational results are presented in Sect. 3. The MSX 8 $\\mu$m optical thicknesses toward the submm peaks are derived in Sect. 4. In Sect. 5, we describe the methods used to derive the physical properties of the observed clumps. In Sect. 6, we discuss the results of our study, and in Sect. 7, we summarise our main conclusions. ", "conclusions": "We have mapped IRDC G304.74 in the 870 $\\mu$m dust continuum emission with the APEX telescope. The submm dust continuum observations have been used together with 1.2 mm data from Beltr{\\'a}n et al. (2006), and archival MSX and IRAS infrared data to derive the physical characteristics of the clumps within the cloud. Besides the dust continuum we used dust extinction data from MSX and 2MASS to derive the H$_2$ column densities, and the mass distribution in the cloud. The results obtained via different methods are in reasonable agreement with each other taking into account the uncertain nature of some of the dust properties and the relation between the H$_2$ column density and extinction. However, the agreement can be improved by assuming an elevated temperature in four clumps near the southwestern end of the cloud, and the possibility of a temperature gradient from about 15 K in the north and centre to about 20-30 K in the south cannot be ruled out. The filamentary cloud G304.74 contains 12 submm clumps. Star formation has already started in the cloud as three of the clumps are associated with both MSX and IRAS point sources. The SEDs of the two IRAS sources indicate bolometric luminosities in the range $\\sim1.5-2\\times10^3$ L$_{\\sun}$. These are likely to be intermediate- or high-mass protostars. In addition, one of the clumps (SMM 6) is associatied with two MSX 8 $\\mu$m point sources. The remaining eight submm clumps are MIR dark. The masses of these clumps ($\\sim40-200$ M$_{\\sun}$) are sufficiently large to enable high-mass star formation, or alternatively, they can represent the cold precursors of proto-clusters. Thus, some of the candidate starless clumps in G304.74 could represent/harbour the so-called high-mass starless cores (HMSCs, e.g., \\cite{beuther2007}). Further studies of these high-mass starless clump/core candidates are important in order to constrain the initial conditions of high-mass star and star cluster formation. The clump masses in G304.74 were compared with the clump mass spectra from more extensive surveys of IRDCs. We found that IRDC clump masses from the present work, and those from Rathborne et al. (2006) and Ragan et al. (2009) probably represent subsamples of the same parent distribution. Also, average distances between a clump and its nearest neighbour in different IRDCs are comparable (within a factor of $\\sim6$), suggesting that the fragmentation length-scale does not vary much from cloud to cloud. Moreover, in most IRDCs, clumps seem to be randomly distributed within the cloud area. These characteristics, and the fact the star formation in IRDCs predominantly occurs in the cluster mode, agree with models where fragmentation is driven by large-scale turbulence. It is not clear, yet, how effectual turbulence is for the fragmentation of IRDC clumps into dense cores. High spatial resolution studies, like the ones presented in Zhang et al. (2009) and Beuther \\& Henning (2009) have recently started to throw light on scales where gravity is likely to dominate." }, "1003/1003.0677_arXiv.txt": { "abstract": "The observation of massive black hole binaries (MBHBs) with Pulsar Timing Arrays (PTAs) is one of the goals of gravitational wave astronomy in the coming years. Massive ($\\simgt 10^8\\,\\msun$) and low-redshift ($\\simlt 1.5$) sources are expected to be individually resolved by up-coming PTAs, and our ability to use them as astrophysical probes will depend on the accuracy with which their parameters can be measured. In this paper we estimate the precision of such measurements using the Fisher-information-matrix formalism. For this initial study we restrict to ``monochromatic\" sources, \\emph{i.e.} binaries whose frequency evolution is negligible during the expected $\\approx 10$ yr observation time, which represent the bulk of the observable population based on current astrophysical predictions. In this approximation, the system is described by seven parameters and we determine their expected statistical errors as a function of the number of pulsars in the array, the array sky coverage, and the signal-to-noise ratio (SNR) of the signal. At \\emph{fixed} SNR (regardless of the number of pulsars in the PTA), the gravitational wave astronomy capability of a PTA is achieved with $\\approx 20$ pulsars; adding more pulsars (up to 1000) to the array reduces the source error-box in the sky $\\Delta\\Omega$ by a factor $\\approx 5$ and has negligible consequences on the statistical errors on the other parameters, because the correlations amongst parameters are already removed to a large extend. If one folds in the increase of coherent SNR proportional to the square root of the number of pulsars, $\\Delta\\Omega$ improves as $1/\\mathrm{SNR}^2$ and the other parameters as $1/\\mathrm{SNR}$. For a fiducial PTA of 100 pulsars uniformly distributed in the sky and a coherent SNR = 10, we find $\\Delta\\Omega\\approx 40\\,\\mathrm{deg}^2$, a fractional error on the signal amplitude of $\\approx 30\\%$ (which constraints only very poorly the chirp mass - luminosity distance combination ${\\cal M}^{5/3}/D_L$), and the source inclination and polarization angles are recovered at the $\\approx 0.3$ rad level. The ongoing Parkes PTA is particularly sensitive to systems located in the southern hemisphere, where at SNR = 10 the source position can be determined with $\\Delta \\Omega \\approx 10\\,\\mathrm{deg}^2$, but has poorer (by an order or magnitude) performance for sources in the northern hemisphere. ", "introduction": "\\label{s:intro} Pulsar Timing Arrays (PTAs), such as the Parkes PTA (PPTA) ~\\cite{man08}, the European PTA (EPTA)~\\cite{jan08}, Nanograv~\\cite{NANOGrav}, the International Pulsar Timing Array (IPTA) project~\\cite{HobbsEtAl:2009}, and in the future the Square Kilometre Array (SKA)~\\cite{laz09} provide a unique means to study the population of massive black hole (MBH) binary systems with masses above $\\sim 10^7\\,M_\\odot$ by monitoring stable radio pulsars: in fact, gravitational waves (GWs) generated by MBH binaries (MBHBs) affect the propagation of electromagnetic signals and leave a distinct signature on the time of arrival of the radio pulses~\\cite{EstabrookWahlquist:1975,Sazhin:1978,Detweiler:1979,HellingsDowns:1983}. MBH formation and evolution scenarios~\\cite{vhm03, kz06, mal07, yoo07} predict the existence of a large number of MBHBs. Whereas the high redshift, low(er) mass systems will be targeted by the planned Laser Interferometer Space Antenna ({\\it LISA}~\\cite{bender98})~\\cite{enelt,uaiti,ses04,ses05,ses07}, massive and lower redshift ($z\\simlt 2$) binaries radiating in the (gravitational) frequency range $\\sim 10^{-9}\\,\\mathrm{Hz} - 10^{-6} \\,\\mathrm{Hz}$ will be directly accessible to PTAs. These systems imprint a typical signature on the time-of-arrival of radio-pulses at a level of $\\approx 1-100$ ns~\\cite{papII}, which is comparable with the timing stability of several pulsars~\\cite{Hobbs:2009}, with more expected to be discovered and monitored in the future. PTAs therefore provide a direct observational window onto the MBH binary population, and can contribute to address a number of astrophysical open issues, such as the shape of the bright end of the MBH mass function, the nature of the MBH-bulge relation at high masses, and the dynamical evolution at sub-parsec scales of the most massive binaries in the Universe (particularly relevant to the so-called ``final parsec problem''~\\cite{mm03}). Gravitational radiation from the cosmic population of MBHBs produces two classes of signals in PTA data: (i) a stochastic GW background generated by the incoherent superposition of radiation from the whole MBHB population~\\cite{rr95, phi01, jaffe, jen05, jen06, papI} and (ii) individually resolvable, deterministic signals produced by single sources that are sufficiently massive and/or close so that the gravitational signal stands above the root-mean-square (rms) level of the background ~\\cite{papII}. In~\\cite{papII} (SVV, hereafter) we explored a comprehensive range of MBH population models and found that, assuming a simple order-of-magnitude criterion to estimate whether sources are resolvable above the background level $\\approx 1$-to-10 individual MBHBs could be observed by future PTAs surveys. The observation of GWs from individual systems would open a new avenue for a direct census of the properties of MBHBs, offering invaluable new information about galaxy formation scenarios. The observation of systems at this stage along their merger path would also provide key insights into the understanding of the interaction between MBHBs and the stellar/gaseous environment~\\cite{KocsisSesana}, and how these interactions affect the black hole-bulge correlations during the merger process. If an electro-magnetic counterpart of a MBHB identified with PTAs was to be found, such a system could offer a unique laboratory for both accretion physics (on small scales) and the interplay between black holes and their host galaxies (on large scales). The prospects of achieving these scientific goals raise the question of what astrophysical information could be extracted from PTA data and the need to quantify the typical statistical errors that will affect the measurements, their dependence on the total number and spatial distribution of pulsars in the array (which affects the surveys observational strategies), and the consequences for multi-band observations. In this paper we estimate the statistical errors that affect the measurements of the source parameters focusing on MBHBs with no spins, in circular orbits, that are sufficiently far from coalescence so that gravitational radiation can be approximated as producing a signal with negligible frequency drift during the course of the observation time, $T \\approx 10$ yr (\"monochromatic\" signal). This is the class of signals that in SVV we estimated to produce the bulk of the observational sample. The extension to eccentric binaries and systems with observable frequency derivative is deferred to a future work. GWs from monochromatic circular binaries constituted by non-spinning MBHs are described by seven independent parameters. We compute the expected statistical errors on the source parameters by evaluating the variance-covariance matrix -- the inverse of the Fisher information matrix -- of the observable parameters. The diagonal elements of such a matrix provide a robust lower limit to the statistical uncertainties (the so-called Cramer-Rao bound~\\cite{JaynesBretthorst:2003, Cramer:1946}), which in the limit of high signal-to-noise ratio (SNR) tend to the actual statistical errors. Depending on the actual structure of the signal likelihood function and the SNR this could underestimate the actual errors, see \\emph{e.g.}~\\cite{NicholsonVecchio:1998,BalasubramanianDhurandhar:1998,Vallisneri:2008} for a discussion in the context of GW observations. Nonetheless, this analysis serves as an important benchmark and can then be refined by carrying out actual analyses on mock data sets and by estimating the full set of (marginalised) posterior density functions of the parameters. The main results of the paper can be summarised as follows: \\begin{itemize} \\item At least three (not co-aligned) pulsars in the PTA are necessary to fully resolve the source parameters; \\item The statistical errors on the source parameters, at \\emph{fixed} SNR, decrease as the number of pulsars in the array increases. The typical accuracy greatly improves by adding pulsars up to $\\approx 20$; for larger arrays, the actual gain become progressively smaller because the pulsars ``fill the sky\" and the effectiveness of further triangulation saturates. In particular, for a fiducial case of an array of 100 pulsars randomly and uniformly distributed in the sky with optimal coherent SNR = 10 -- which may be appropriate for the SKA -- we find a typical GW source error box in the sky $\\approx 40$ deg$^2$ and a fractional amplitude error of $\\approx$ 30\\%. The inclination and polarization angles can be determined within an error of $\\sim 0.3$ rad, and the (constant) frequency is determined to sub-frequency resolution bin accuracy. These results are independent on the source gravitational-wave frequency. \\item When an anisotropic distribution of pulsars is considered, the typical source sky location accuracy improves linearly with the array sky coverage. The statistical errors on all the other parameters are essentially insensitive to the PTA sky coverage, as long as it covers more than $\\sim 1$ srad. \\item The ongoing Parkes PTA aims at monitoring 20 pulsars with a 100 ns timing noise; the targeted pulsars are mainly located in the southern sky. A GW source in that part of the sky could be localized down to a precision of $\\lesssim 10$ deg$^2$ at SNR$=10$, whereas in the northern hemisphere, the lack of monitored pulsars limits the error box to $\\simgt 200$ deg$^2$. The median of the Parkes PTA angular resolution is $\\approx 130\\,(\\mathrm{SNR}/10)^{-2}$ deg$^2$. \\end{itemize} The paper is organised as follows. In Section II we describe the GW signal relevant to PTA and we introduce the quantities that come into play in the parameter estimation problem. A review of the Fisher information matrix technique and its application to the PTA case are provided in Section III. Section IV is devoted to the detailed presentation of the results, and in Section V we summarize the main findings of this study and point to future work. Unless otherwise specified, throughout the paper we use geometric units $G=c=1$. ", "conclusions": "In this paper we have studied the expected uncertainties in the measurements of the parameters of a massive black hole binary systems by means of gravitational wave observations with Pulsar Timing Arrays. We have investigated how the results vary as a function of the signal-to-noise ratio, the number of pulsars in the array and their location in the sky with respect to a gravitational wave source. Our analysis is focused on MBHBs in circular orbit with negligible frequency evolution during the observation time (\"monochromatic sources\"), which we have shown to represent the majority of the observable sample, for sensible models of sub--parsec MBHB eccentricity evolution. The statistical errors are evaluated by computing the variance-covariance matrix of the observable parameters, assuming a coherent analysis of the Earth-terms only produced by the timing residuals of the pulsars in the array (see Section II B). For a fiducial case of an array of 100 pulsars randomly distributed in the sky, assuming a coherent total SNR = 10, we find a typical error box in the sky $\\Delta \\Omega \\approx 40$ deg$^2$ and a fractional amplitude error of $\\approx 0.3$. The latter places only very weak constraints on the chirp mass-distance combination ${\\cal M}^{5/3}/D_L$. At fixed SNR, the typical parameter accuracy is a very steep function of the number of pulsars in the PTA up to $\\approx 20$. For PTAs containing more pulsars, the actual gain becomes progressively smaller because the pulsars ``fill the sky\" and the effectiveness of further triangulation weakens. We also explored the impact of having an anisotropic distribution of pulsars finding that the typical source sky location accuracy improves linearly with the array sky coverage. For the specific case of the Parkes PTA where all the pulsars are located in the southern sky, the sensitivity and sky localisation are significantly better (by an order of magnitude) in the southern hemisphere, where the error-box is $\\lesssim 10 \\,\\mathrm{deg}^2$ for a total coherent SNR = 10. In the northern hemisphere, the lack of monitored pulsars prevent a source location to be in an uncertainty region $\\lesssim 200\\,\\mathrm{deg}^2$. The monitoring of a handful of pulsars in the northern hemisphere would significantly increase both the SNR and the parameter recovery of GW sources, and the International PTA~\\cite{HobbsEtAl:2009} will provide such a capability in the short term future. The main focus of our analysis is on the sky localisation because sufficiently small error-boxes in the sky may allow the identification of an electro-magnetic counterpart to a GW source. Even for error-boxes of the order of tens-to-hundreds of square degrees (much larger than \\emph{e.g.} the typical {\\it LISA} error-boxes~\\cite{v04,k07,lh08}), the typical sources are expected to be massive (${\\cal M} \\simgt 10^{8}\\msun$) and at low redshift ($z\\simlt 1.5$), and therefore the number of associated massive galaxies in the error-box should be limited to a few hundreds. Signs of a recent merger, like the presence of tidal tails or irregularities in the galaxy luminosity profile, may help in the identification of potential counterparts. Furthermore, if nuclear activity is present, \\emph{e.g.} in form of some accretion mechanism, the number of candidate counterparts would shrink to an handful, and periodic variability \\cite{hkm07} could help in associating the correct galaxy host. We are currently investigating the astrophysical scenarios and possible observational signatures, and we plan to come back to this important point in the future. The advantage of a counterpart is obvious: the redshift measurement would allow us, by assuming the standard concordance cosmology, to measure the luminosity distance to the GW source, which in turn would break the degeneracy in the amplitude of the timing residuals $R \\propto {\\cal M}^{5/3}/(D_L f^{1/3})$ between the chirp mass and the distance, providing therefore a direct measure of ${\\cal M}$. The study presented in this paper deals with monochromatic signals. However, the detection of MBHBs which exhibit a measurable frequency drift would give significant payoffs, as it would allow to break the degeneracy between distance and chirp mass, and enable the direct measurement of both parameters. Such systems may be observable with the Square-Kilometre-Array. In the future, it is therefore important to extend the present analysis to these more general signals. However, as the frequency derivative has only modest correlations with the sky position parameters, we expect that the results for the determination of the error-box in the sky discussed in this paper will still hold. A further extension to the work is to consider MBHBs characterised by non-negligible eccentricity, which is currently in progress. Another extension to our present study is to consider both the Earth- and pulsar-terms in the analysis of the data and the investigation of the possible benefits of such scheme, assuming that the pulsar distance is not known to sufficient accuracy. This also raises the issue of possible observation campaigns that could yield an accurate (to better than 1 pc) determination of the pulsar distances used in PTAs. In this case the use of the pulsar-term in the analysis would not require the introduction of (many more) unknown parameters and would have the great benefit of breaking the degeneracy between chirp mass and distance. The final world of caution goes to the interpretation of the results that we have presented in the paper. The approach based on the computation of the Fisher Information matrix is powerful and straightforward, and is justified at this stage to understand the broad capabilities of PTAs and to explore the impact on astronomy of different observational strategies. However, the statistical errors that we compute are strictly \\emph{lower limits} to the actual errors obtained in a real analysis; the fact that at least until SKA comes on line, a detection of a MBHB will be at a moderate-to-low SNR should induce caution in the way in which the results presented here are interpreted. Moreover, in our current investigation, we have not dealt with a number of important effects that in real life play a significant role, such as different calibrations of different data sets, the change of systematic factors that affect the noise, possible non-Gaussianity and non-stationarity of the noise, etc. These (and other) important issues for the study of MBHBs with PTAs should be addressed more thoroughly in the future by performing actual mock analyses and developing suitable analysis algorithms. \\appendix" }, "1003/1003.5052_arXiv.txt": { "abstract": "{The correlation between distant Gamma-Ray Bursts (GRBs) and foreground galaxy clusters is re-examined by using the well localized (with an accuracy down to a few arcseconds) \\it Swift/\\rm XRT GRBs. The galaxy clusters are compiled from both X-ray selected \\it ROSAT \\rm brightest cluster sample (BCS) and BCS extension by requiring $\\delta \\geq0\\degr$ and $b\\geq20\\degr$. The \\it Swift/\\rm XRT GRBs fulfilling the above selection criteria are cross-correlated with the clusters. Both Nearest-Neighbor Analysis and angular two-point cross-correlation function show that there is no enough evidence supporting the correlation between the GRBs and foreground clusters. We suggest that the non-correlation is probably related to the GRB number-flux relation slope. ", "introduction": "% \\label{sect:intro} It is now no doubt that Gamma-Ray Bursts (GRBs) take place at cosmological distance. Thanks to the prompt localizations and deep follow-up observations, the record of the highest redshift of GRBs has been progressively broken in past a few years, especially after the launch of the \\it Swift \\rm satellite (Gehrels et al. 2004). At present, GRB\\,090423 detected by the \\it Swift \\rm satellite is the most distant GRB with a redshift of $\\sim8.1$ (Salvaterra et al. 2009). So far, there are about 50 \\it Swift \\rm detected GRBs with measured redshifts. A majority of these GRBs lie beyond $z=1$ with a redshift distribution that peaks at $z\\sim1-2$. The cosmological origin and high luminosities offer an opportunity to use GRBs as tracers to study: star formation history of the Universe (e.g., Savaglio et al. 2009; Jakobsson et al. 2005); properties and evolution of intergalactic medium and high redshift galaxies (e.g., Prochaska et al. 2006; Prochter et al. 2006; Vergani et al. 2009), similar as done with the high redshift quasars; and the nearby mass distribution through the weak lensing of GRBs caused by the local large-scale structure (Williams \\& Frey 2003 and references therein). Because GRBs are point-like sources, their weak lensing could only be detected trough the angular correlation between sources and corresponding lenses\\footnote{Another way looking for candidates of GRB lensing effect is based on the time delay of two bursts from the same sky region (e.g., Veres et al. 2009)}. A number of authors previously examined whether subsets of GRBs are correlated with subsets of foreground galaxy clusters. The results obtained by these authors are, however, contradictious. Kolatt \\& Piran (1996) claimed the 136 GRBs selected from the Burst and Transient Source Experiment (BATSE) 3B catalog are correlated with the Abell cluster (Abell et al. 1989) within an angular separation of 4\\degr at a significance level 95\\%. Marani et al. (1997) obtained a stronger correlation by using the BATSE GRBs with more accurate positions. By contrast, Hurley et al. (1999) did not find any evidence for the correlation between GRBs and galaxy clusters by extending the GRB sample to the BATSE 4B/Third Interplanetary Network catalog. In addition, Williams \\& Frey (2003) reported an anti-correlation between the BATSE GRBs and Abell clusters. A caveat in these previous studies is the large BATSE error box that usually ranges from a fraction of a degree to a few degrees. The error box in some cases is as high as $\\sim30\\degr$. The poor localization has been greatly improved after the launch of the \\it Swift \\rm satellite. The spacecraft of the \\it Swift \\rm satellite can quickly slew to the GRB position given by the BAT instrument within 100 seconds. Due to the high sensitivity, the XRT onboard the \\it Swift \\rm satellite has the capability that measure X-ray afterglow position with an accuracy better than 5\\arcsec within 100 seconds for about 90 percent BAT triggers (Burrows et al. 2005). Here, we re-examine the correlation between GRBs and foreground galaxy clusters by using the \\it Swift\\rm/XRT sample. As mentioned before, the highly accurate position provided by the XRT instrument allows us to regard these GRBs as point sources. The correlation is studied by the nearest-neighbor distance method and angular two-point cross-correlation function. Both methods indicate that there is no significant correlation between the GRBs and foreground galaxy clusters. .... The paper is organized as follows. \\S 2 describes the sample selection. The analysis and results are presented in \\S 3. A short discussion and a conclusion are provided in \\S4. ", "conclusions": "\\label{sect:con} With much improved burst locations down to a few arcseconds, the \\it Swift\\rm/XRT GRB sample allows us to re-examine the correlation between GRBs and foreground galaxy clusters. By using the NNA and angular two-point cross-correlation function, our analysis indicates that there is no enough evidence supporting the correlation between the foreground X-ray selected \\it ROSAT \\rm brightest clusters and \\it Swift\\rm/XRT GRBs. Although our study is not the first one examining the relation between GRBs and foreground galaxy clusters, it is indeed the first study using the GRBs position information with accuracy down to a few arcseconds. The contradiction of the results obtained before 2004 (see more details and the references listed in Section 1) is likely due to the poor localization of the BATSE GRB sample. It is note that the error boxes of the BATSE GRBs range from a fraction of a degree to as large as $\\sim30\\degr$. Generally speaking, the motivation of searching the correlation between distant GRBs and nearby massive structure is to study the local large structure in terms of the weak lensing effect of the GRBs. Assuming the weak lensing effect indeed occurs for the GRBs, how could we understand the non-correlation revealed in the current study? In fact, the association due to weak lensing effect not only depends on the mass of the lenses, but also depends on the number-flux relation of the background objects. Weak lensing increases the brightness of faint objects, while expands the area behind the lenses. Combining the competing two factors, the correlation function could be related with the number-flux relation as (e.g., Williams \\& Irwin 1998; Myers et al. 2003) \\begin{equation} \\omega(\\theta)=\\mu^{2.5\\beta-1}-1 \\end{equation} where the number-flux relation is approximately described as a powerlaw: $N(>\\log f)\\propto 10^{-2.5\\beta\\log f}$ (Boyle et al. 1988), $\\mu$ is the magnification factor. This equation indicates that $\\omega(\\theta)=0$ when $\\beta=0.4$. The value of $\\beta$ could be constrained from the \\it Swift \\rm/XRT observations. To estimate $\\beta$, we use the sample recently compiled by Zheng et al. (2009). Zheng et al. (2009) collected the information of amount of GRBs observed by \\it Swift \\rm\\ satellite before 2008 from the literature and on-line databases. The number-flux relation of these GRBs is displayed in Figure 5 for the 3\\,keV X-ray flux density at 11 hour. The data are evenly binned into 9 bins in the logarithmic space. The over-plotted error bars correspond to the 1$\\sigma$ Poisson noise in each bin. The solid line plots the best fitted smoothed powerlaw model with an expression: \\begin{equation} \\frac{dN}{d\\log f}=\\frac{N_0}{10^{a(\\log f-\\log f_0)}+10^{b(\\log f-\\log f_0)}} \\end{equation} where $N_0=218$ is the total number of the GRBs. Comparing the equation with Eq (5) yields $\\beta=0.4b$ at the bright end. A weighted least-squared fitting yields the following parameters: $a=-0.0003\\pm0.003$, $b=1.05\\pm0.03$, and $\\log f_0=1.55\\pm0.02$. Although the sample seems to be clearly incomplete at the faint end, $\\beta$ is inferred to be $0.42\\pm0.01$ from the well sampled bright end. In fact, about half of the GRBs listed in the GRB sample are bright in X-ray with $\\log(f_x/10^{-3} \\mu\\mathrm{Jy})>1.5$. This exercise suggests that the zero correlation function of the GRBs is likely due to their number-flux relation, although the weak lensing effect may indeed occur for these GRBs. \\begin{figure}[h!!!] \\centering \\includegraphics[width=9.0cm, angle=0]{ms395figure5.eps} \\begin{minipage}[]{85mm} \\caption{ The number-flux relation for the GRBs compiled by Zheng et al. (2009). The data are evenly binned into 9 bins in the logarithmic space. The errorbars over-plotted are 1$\\sigma$ Poisson noise. The solid line is the best fitted smoothed powerlaw model (see Eq. 6). } \\end{minipage} \\label{Fig1} \\end{figure}" }, "1003/1003.0455_arXiv.txt": { "abstract": "We report on an on-going test campaign of more than 5000 Schottky CdTe detectors ($4\\times4\\times1$ mm$^3$), over a sample of twelve thousands, provided by Acrorad Co., Ltd (Japan). 6400 of these detectors will be used to build the detection plane of the ECLAIRs camera on the Chinese-French gamma-ray burst mission SVOM. These tests aim at eliminating, from the flight model, detectors showing high leakage current and unstable behaviors. They are mandatory to fulfill the prime requirement of ECLAIRs to detect gamma-ray burst photons down to 4 keV. For better functional performance and stability, the detectors will be operated at $-20^\\circ$C under a reverse bias of 600 V. Under these conditions, we found that 78\\% of the detectors already tested could be considered for the flight model. They show a low initial leakage current (with a mean value around 20 pA), and remain stable below 100 pA within two hours. Using a standard analog electronics chain, we measured a mean energy resolution of 1.8 keV at 59.6 keV using an $^{241}$Am source. Because the Schottky detectors are well known to be unstable due to the bias-induced polarization effect, the high voltage power supplies on ECLAIRs will have to be switched off at regular time intervals. We investigated the polarization effect first at room temperature and low bias voltage for faster analysis. We found that the spectroscopic degradation in quantum efficiency, gain and energy resolution, starts as soon as the bias is turned on: first slowly and then dramatically after a time $t_p$ which depends on the temperature and the voltage value. Preliminary tests under in-flight conditions ($-20^\\circ$C, -600 V) showed that the detectors should remain stable over a timescale larger than a day. As a by product of our test campaign, we measured the mean activation energy of 170 Schottky CdTe detectors. We found evidence for two distinct populations of detectors: the main one centered at 0.64 eV, interpreted as due to cadmium vacancies in the crystal, and the second population centered at 0.54 eV, correlated with a lower apparent resistivity. ", "introduction": "The SVOM (\\textit{Space-based multi-band astronomical Variable Objects Monitor}) gamma-ray burst mission will carry the ECLAIRs 2D coded-mask telescope, operating in the 4-250 keV range \\cite{2009AIPC.1133...25G,2009AIPC.1133...76R}. CdTe semi-conductor detectors are well suited for detecting X- and gamma-rays in this energy range, thanks to the high detection efficiency of CdTe, due to a high atomic number ($Z_{\\rm Cd}$=48, $Z_{\\rm Te}$=52) and a high density ($\\rho_{\\rm CdTe}$ = 5.85 g.cm$^{-3}$) \\cite{2001ITNS...48..950T}. The 4 keV low-energy threshold is a driver for the mission, as it shall improve significantly the sensitivity of ECLAIRs to high redshift gamma-ray bursts \\cite{2009NIMPA.603..365G}. For this purpose, the front-end electronics integrated in the ASIC IDeF-X \\cite{2006NIMPA.567..140G,2009ITNS...56.2351G} for Cd(Zn)Te detectors has been chosen. It shows a low intrinsic noise (Equivalent Noise Charge (ENC) floor of 33 e$^{-}$ rms at 6 $\\mu$s peaking time), low consumption ($\\sim 3$ mW/channel) and is radiation hard (Single Event Latchup Linear Energy Transfer threshold of 56 MeV.cm$^{2}$.mg$^{-1}$). However, it is suited only to low noise detectors and low capacitive materials. Therefore, we must choose detectors with the lowest leakage current, in order to minimize the parallel noise. Given the theoretical expression of the ENC at the output of the shapers of the ASIC \\cite{2006NIMPA.567..140G}, while using measured values of the noise coefficients $\\alpha_{\\rm d}$ and $\\alpha_{\\rm 1/f}$ from IDeF-X ECLAIRs (table II in \\cite{2009ITNS...56.2351G}), we can simulate the evolution of the ENC as a function of the detector leakage current ($I_{\\rm leak}$). These values depend on both the peaking time chosen (between 2 and 6 $\\mu$s typically), and the total capacitance at the input of the charge sensitive amplifier of the ASIC. The later may vary between 2 pF and 10.5 pF \\cite{2009ITNS...56.2351G}. It includes the detector capacitance together with the parasitic capacitance, mainly due to interconnections between the detector and the electronics (ceramic boards in our case). Assuming a 6 sigma low-energy threshold, and considering the worst case (i.e. 10.5 pF input capacitance) and 6 $\\mu$s peaking time, detecting photons of 4 keV requires \\ileak\\ to be lower than 150 pA during the in-flight operation of the camera (Figure \\ref{remoue:fig1}). \\begin{figure}[!h] \\begin{center} \\centerline{\\includegraphics[width=.45\\textwidth]{remoue_fig1.ps}} \\caption{The maximum leakage current accepted to ensure a 4 keV low energy threshold (defined at 6 sigma of the centre of the noise peak) as a function of the input capacitance (detector + parasitic). Two values are compared for the peaking time: 2 $\\mu$s and 6 $\\mu$s. We used the theoretical ENC expression and measured noise parameters of the ASIC IDeF-X ECLAIRs \\cite{2006NIMPA.567..140G,2009ITNS...56.2351G}. The 150 pA limit is shown for indication (dashed line).} \\label{remoue:fig1} \\end{center} \\end{figure} Such a low value of \\ileak\\ could be achieved with ohmic CdTe detectors operated at low temperature, as were previously used for the ISGRI hard X-ray imager aboard INTEGRAL \\cite{2003A&A...411L.141L}, but only under the condition of a low bias voltage (because of the linear increase of \\ileak\\ with the bias in this case). But lower bias voltage means less charge collected, because of the low mobility and short lifetime of holes in CdTe \\cite{2001ITNS...48..950T}. Therefore, Schottky CdTe diodes are favoured. Indeed, the Schottky blocking junction enables to apply higher bias voltage without increasing too much the leakage current, which in turn improves the charge collection efficiency, reducing the distortion of the X-ray spectrum. The detection plane of ECLAIRs will then be an assembly of 6400 Schottky CdTe monolithic pixels with a $4\\times4$ mm$^{2}$ area and 1 mm thickness provided by Acrorad Co., Ltd (Japan) \\cite{2009AIPC.1133...76R}. The anode is made of Indium and the cathode of Platinum. The crystal is grown by the so-called Traveling Heater Method (THM), and is doped with Cl \\cite{1999NIMPA.436..120F}. This is the first time such planar monolithic In/CdTe/Pt detectors are considered for hard X-ray imaging in space, as ISGRI aboard INTEGRAL used Pt/CdTe/Pt detectors \\cite{2003A&A...411L.141L} and the BAT telescope on \\textit{Swift} used CdZnTe pixels \\cite{2005SSRv..120..143B}. More recently, pixellated devices are considered as for the Caliste 64 micro-camera \\cite{2009NIMPA.610..154M} for the Symbol-X mission, or the Si/CdTe Compton camera \\cite{2007NIMPA.579..871W} on board Astro-H. Unfortunately, Schottky-type CdTe detectors suffer from the bias-induced polarization phenomenon, which manifests itself by a decrease in quantum efficiency, a degradation of the spectral resolution and a shift of the photo-peak towards lower energies \\cite{1974ITNS...21..322M,2006JaJAP..45.8842T,2009ApPhL..94j2113C,2007ApPhL..90f3504O}. A way to minimize and delay considerably this effect is to operate the detectors at low temperature and high bias voltage \\cite{2006JaJAP..45.8842T,1998ITNS...45..428M}. Thus ECLAIRs will keep an in-flight nominal temperature of $-20^\\circ$C $\\pm$ $2^\\circ$C over the whole detection plane, with the detectors reverse biased at 600 V. Moreover, this effect being reversible, a way to maintain stable performance of the detectors over time is to turn off the high voltage at regular time intervals. In the current baseline configuration, the high voltage is expected to be switched off once a day for ECLAIRs, when the spacecraft will pass in the deepest part of the South Atlantic Anomaly. Another type of polarization induced by high photon fluence irradiation has been reported for Cd(Zn)Te detectors, typically above $10^{5}$ photons/mm$^{2}$ integrated over about 100 ms \\cite{1988NIMPA.263..457V,1999MedPh..26...38J,2008ITNS...55..572S}. It shall not be an issue for ECLAIRs, as an on-axis exceptionally bright gamma-ray burst should generate less than 1 photon/mm$^{2}$/100 ms. Thirty two pixels are to be mounted on a monolithic Al$_2$O$_3$ ceramic plate and hybridized with the ASIC IDeF-X ECLAIRs, to form the so-called XRDPIX elementary module \\cite{2009AIPC.1133...76R,remoue09}. In order to minimize the parallel noise of the ASIC, we decided to select the most suitable CdTe detectors to be used for the flight before mounting them on each XRDPIX. In fact, each channel of the ASIC has its own adjustable discrimination threshold \\cite{2009ITNS...56.2351G}. But, because we aim for homogeneous modules, we considered during the selection process that the performance will be dictated by the worst detector on a given XRDPIX. In total, more than twelve thousands will be tested. Their leakage currents at two operating temperatures ($-20^\\circ$C and $+25^\\circ$C) will be measured and their spectroscopic performance assessed. In this paper, we first report on the results of a uniform testing of more than 5000 detectors, which led to the selection criteria to be applied to the flight detectors. Albeit similarly large numbers of detectors currently flowing on INTEGRAL and \\textit{Swift} have been tested \\cite{1999NIMPA.428..216L,2005NIMPA.541..372S}, this is the first time Schottky CdTe detectors are tested so extensively. We then describe the results of complementary tests of a few selected samples of detectors showing either acceptable performance or atypical behaviors during the automatic tests; the latter could only be identified because of the large sample of detectors considered. Those tests include the measurement of the polarization effect, as well as the time evolution of the dark current, and finally the measurement of the detector performance as a function of the operating temperature, in particular to define a precise value of their activation energy. Those tests are required to anticipate the detector behavior in space and optimize the in-flight performance of the ECLAIRs camera, under the general constraints of the SVOM mission (low-earth orbit, frequent South Atlantic Anomaly passages) and science requirements (detection of GRB photons down to 4 keV). \\label{intro} \\begin{figure} \\begin{center} \\centerline{\\includegraphics[width=.49\\textwidth]{remoue_fig2.ps}} \\caption{The leakage current measured at $-20^\\circ$C against the leakage current at $+25^\\circ$C, in pA, for 5000 detectors (at -600 V). The cluster of plus signs identifies the main population of detectors, containing about 90\\% of the total sample. A linear fit is shown for indication (a=$1.6\\times10^{-2}$, b=$-3.5$ pA). Detectors deviating from the main correlation line are considered suspicious; some of them have been further investigated (see Figure \\ref{remoue:fig4}). The cluster of cross marks at the bottom left of the plot (less than 2.5\\% of the total number of detectors tested) can be explained by the aging of the electronic chain.} \\label{remoue:fig2} \\end{center} \\end{figure} \\begin{figure}[!h] \\begin{center} \\centerline{\\includegraphics[width=.49\\textwidth]{remoue_fig3.ps}} \\caption{The histogram of leakage currents measured at $-20^\\circ$C (-600 V) for the main population (corresponding to the red cluster in Figure \\ref{remoue:fig2}). For indication, the histogram was fitted by a Landau function, which peaks at 12 pA. The distribution at $+25^\\circ$C shows the same pattern, with a peak value at 1.1 nA.} \\label{remoue:fig3} \\end{center} \\end{figure} ", "conclusions": "After extensive testing of an homogeneous set of more than 5000 Schottky CdTe detectors, foreseen for the detection plane of the X-/Gamma-ray telescope on board the SVOM mission, we can draw the following conclusions, relevant to their in-flight operation and to the requirement for ECLAIRs to detect X-ray photons down to 4 keV. \\begin{itemize} \\item First, we defined selection criteria for the detectors to be mounted on the XRDPIX flight modules. Under ECLAIRs in-flight operating conditions ($-20^\\circ$C and a 600 V reverse bias), they show a low initial leakage current (with a mean value around 20 pA), which remains below 100 pA within 2 hours. They are located less than three sigma away from the main correlation line describing the relation between \\ileak\\ at $+25^\\circ$C and \\ileak\\ at $-20^\\circ$C. These criteria are based on the worse possible hybridization scheme of the detectors with the ASIC IDeF-X, and therefore they are very conservative. Using a standard analog electronics chain, we measured a mean energy resolution of 1.8 keV at 59.6 keV, with the FWHM-value being dominated by the chain ENC ($\\sim 135$ electrons rms), i.e. these detectors should intrinsically better resolve the X-ray lines than can be measured here. Yet, we found about 10\\% detectors with bad spectra, to be removed from the selection. Therefore, taking into account all these criteria, we found that about 78\\% of flight pre-selected detectors at the moment. \\item The absorption due to the Pt layer on top of the 1 mm thick detector determines its efficiency below 20 keV, with about 45-55\\% of photons lost at 4 keV. Because the Pt layer was found to be variable, between 200 and 300 nm, precise quantum efficiency measurements will be carried out on a sample of XRDPIX units using the SOLEX facility. A $^{55}$Fe radioactive source, showing a characteristic X-ray line at 5.9 keV, will also be used during the XRDPIX performance assessments. \\item The bias-induced polarization effect has been shown to occur on a very short timescale at high temperatures and low bias voltages, justifying the needs for operating ECLAIRs at $-20^\\circ$C and -600 V. We investigated the conventional and modified charge accumulation models described in \\cite{2006JaJAP..45.8842T}, using our leakage current measurements at room temperature. Our data were found to be more in agreement with the conventional model. We also derived the time $t_w$ when the depletion depth begins to decrease, and we found that, even if the $t_w$-values give a rough estimation of the $t_p$-values measured in spectroscopy, they are systematically higher. \\item The time stability of the detector performance is essential to preserve the science capability of the instrument ECLAIRs. To do so, we need to define the in-flight operation conditions taking into account the intrinsically unstable nature of the Schottky CdTe detectors. It will be necessary to switch off the high voltage on a regular timescale in order to limit the effect of the polarization phenomenon as well as the constant increase in the leakage current with time. The measurements performed on pre-selected detectors, when biased at -600 V and cooled down to $-20^\\circ$C, showed that their spectroscopic performance (quantum efficiency, gain and energy resolution) should be stable on a timescale larger than a day. The long duration leakage current measurements, performed on a 100 hour timescale, showed that in most cases the $I_{\\rm leak}$-values should stay below the maximum 150 pA allowed to maintain a 4 keV low-energy threshold. Yet, precise determination of the time to polarization $t_p$ of these detectors at $-20^\\circ$C and -600 V is needed in order to determine the nominal duration of operation in orbit, which could be longer than a day following our preliminary tests, as well as the timescale for full recovery of the detector performance when switching off the high voltage. This will be investigated at the XRDPIX level. \\item As a by product of our test campaign, we measured the mean activation energy of 170 Schottky CdTe detectors, representative of all the detectors tested. We found evidence for two distinct populations: the main one centered at 0.64 eV and the second one centered at 0.54 eV. Those two values have been previously reported, albeit over a limited sample of detectors. The 0.64 eV energy activation is often reported \\cite{2006JaJAP..45.8842T,2009ApPhL..94j2113C,2003ITNS...50..229A} and is likely to be due to cadmium vacancies, while the less frequent 0.54 eV level might be present under specific concentration of Cl dopant \\cite{2003ITNS...50..229A,{laasch1995}}. We showed that the second population centered at 0.54 eV is correlated with: i) an apparent resistivity lower than that computed for the main population (typically below 110 G$\\Omega$.m at $-20^\\circ$C); ii) a sample of non-selected detectors showing higher $I_{\\rm leak}$-values at $-20^\\circ$C. \\end{itemize}" }, "1003/1003.5793_arXiv.txt": { "abstract": "We studied several representative circumstellar disks surrounding the Herbig Ae star MWC~480 and the T Tauri stars LkCa~15 and DM~Tau at (sub-)millimeter wavelengths in lines of CCH\\footnote{Based on observations carried out with the IRAM Plateau de Bure Interferometer. IRAM is supported by INSU/CNRS (France), MPG (Germany) and IGN (Spain).}. Our aim is to characterize photochemistry in the heavily UV-irradiated MWC~480 disk and compare the results to the disks around cooler T Tauri stars. We detected and mapped CCH in these disks with the IRAM Plateau de Bure Interferometer in the C- and D-configurations in the (1-0) and (2-1) transitions. Using an iterative minimization technique, the CCH column densities and excitation conditions are constrained. Very low excitation temperatures are derived for the T Tauri stars. These values are compared with the results of advanced chemical modeling, which is based on a steady-state flared disk structure with a vertical temperature gradient, and a gas-grain chemical network with surface reactions. Both model and observations suggest that CCH is a sensitive tracer of the X-ray and UV irradiation. The predicted radial dependency and source to source variations of CCH column densities qualitatively agree with the observed values, but the predicted column densities are too low by a factor of several. The chemical model fails to reproduce high concentrations of CCH in very cold disk midplane as derived from the observed low excitation condition for both the (1-0) and (2-1) transitions. ", "introduction": "Our present understanding of the chemical composition and evolution of protoplanetary disks is far from being complete. Recent years have revealed increased interest in such studies as a necessary prerequisite to characterize the physical and chemical conditions for planet formation. Apart from CO and its isotopologues, and occasionally HCO$^+$, CN, HCN, and CS, the molecular content of protoplanetary disks characterized by millimeter line observations remains poorly known \\citep[see reviews by][]{DGH07,Bergin_ea07}. Recently, infrared spectroscopy provided evidence for the presence of C$_2$H$_2$, HCN, H$_2$O, OH, and CO$_2$ in the inner regions of protoplanetary disks and indicated the importance of photochemistry \\citep[e.g.,][]{Lahuis_ea06,Carr_Najita08,Salyk_ea08,Pascucci_ea09}. Molecular line data are limited in their sensitivity and spatial resolution, which implies that the spatial distribution of molecular abundances in disks remains poorly determined \\citep[e.g.,][]{Pietu_ea05,Semenov_ea05,Qi_ea06,Qi_ea08,Panic_ea09}. Thus a detailed comparison with existing chemical models is difficult and often based on global data (e.g. integrated line profiles), while for testing the validity of sophisticated models one has to include the thermal disk structure and the time-dependent chemical evolution together with line data and channel maps. As part of the Heidelberg-Bordeaux ``Chemistry In Disks'' (CID) project we are investigating the molecular content and physical properties of a sample of well-studied T~Tauri and Herbig~Ae disks of various age, followed by comprehensive physico-chemical modeling. In the first CID paper, we have presented the results of a deep search for N$_2$H$^+$ and HCO$^+$ toward two T Tauri stars (DM~Tau, LkCa~15) and one Herbig Ae star (MWC~480), see \\citet[CID1 paper hereafter]{Dutrey_ea07}. The N$_2$H$^+$ emission has been detected in LkCa 15 and DM Tau, with the N$_2$H$^+$ to HCO$^+$ ratio of a few percent, similar to that of cold dense cores, and the disk ionization degree of $\\sim 10^{-8}$, as predicted by chemical models. In the second CID paper by \\citet[]{Schreyer_ea08}, we have found that the modeled and observed column densities of several key species relative to $^{13}$CO are lower for the Herbig AB Aur disk than the values measured in DM Tau, while the absolute amount of CO gas is similar in both disks. This has been interpreted as an indication of a poor molecular content of the Herbig A0e system compared to the disk around the M1e star DM Tau due to intense UV irradiation of the disk by the A0 star. In this third CID paper, we report on the observations of CCH in three objects: DM~Tau, LkCa~15 and MWC~480. The CCH molecule has been selected because it should be sensitive to the UV radiation field and because its chemistry is relatively well-studied, with many accurately determined relevant reaction rates \\citep[e.g.,][]{Woodall_ea07}. Using the $\\chi^2$-minimization technique in the $uv$-plane as described in \\citet{GD98}, we derive the column density of CCH in the upper layers of the outer disks. We compare these values with the CCH column densities computed with realistic two-dimensional steady-state disk models and a gas-grain chemistry with surface reactions. ", "conclusions": "\\label{summary} We observed protoplanetary disks around the Herbig Ae star MWC~480 and the T Tauri stars LkCa~15 and DM~Tau in the (1-0) and (2-1) lines of CCH. We detected and mapped these disks in the CCH lines, using the IRAM Plateau de Bure Interferometer in the C- and D-configurations. Using the iterative minimization technique, the column densities of ethynyl are derived, which are lower in MWC~480 ($N({\\rm CCH})\\la 10^{13}$~cm$^{-2}$) than in DM Tau and LkCa~15 ($N({\\rm CCH})\\sim 3\\,10^{13}$~cm$^{-2}$). The derived excitation temperature of CCH in DM Tau is very low, $T_{\\rm ex}\\sim 6$~K, which is along with previous detections of cold gas in this object (CO, HCO$^+$, CH, and HCN). The observed CCH column densities are compared with the results of advanced chemical modeling, which is based on a steady-state flared disk structure with a vertical temperature gradient, and a gas-grain chemical network with surface reactions. The CCH abundances in all disk models show two layers of high concentration. Overall, the disk around the Herbig Ae star MWC~480 has less CCH, which we explain by the strong UV radiation field of this object, and the lack of X-ray irradiation from the star. The modeled CCH column densities are in qualitative agreement with the observed values at an evolutionary time of a few million years. Altogether, our measurements and modeling support the observational trend that disks around Herbig Ae stars tend to be more molecule deficient compared to those surrounding T Tauri stars due to stronger UV irradiation and lower X-ray luminosities (affecting inner disk chemistry). Our model reproduces the radial slope of the CCH column densities. Yet it fails to explain the low excitation temperature of the observed CCH lines as the modeled CCH molecular layers are located in disk regions dense enough to thermalize both transitions. The chemical simulations show that only $\\la 20\\%$ of the CCH column densities are in the region of the DM Tau disk that is colder than 20~K. Other mechanisms (stronger photodesorption, dynamic processes, grain growth) are required to enhance the abundance of CCH in the cold disk midplanes." }, "1003/1003.0663_arXiv.txt": { "abstract": "{Our universe may have formed via bubble nucleation in an eternally-inflating background. Furthermore, the background may have a compact dimension---the modulus of which tunnels out of a metastable minimum during bubble nucleation---which subsequently grows to become one of our three large spatial dimensions. Then the reduced symmetry of the background is equivalent to anisotropic initial conditions in our bubble universe. We compute the inflationary spectrum in such a scenario and, as a first step toward understanding the effects of anisotropy, project it onto spherical harmonics. The resulting spectrum exhibits anomalous multipole correlations, their relative amplitude set by the present curvature parameter, which appear to extend to arbitrarily large multipole moments. This raises the possibility of future detection, if slow-roll inflation does not last too long within our bubble. A full understanding of the observational signal must account for the effects of background anisotropy on photon free streaming, and is left to future work.} \\begin{document} ", "introduction": "\\label{sec:introduction} Inflation is generically eternal \\cite{V83,L86}. That is, for many scalar field potentials the physical volume of inflating spacetime is divergent, with inflation ending only in localized ``pockets'' within the inflating background. This is the case, for instance, when inflation is driven by the positive vacuum energy of some metastable ``parent'' vacuum, in which the vacuum phase of our ``daughter'' universe arises due to tunneling through a potential barrier. The tunneling process is described by an instanton that interpolates between parent and daughter vacua, and appears as bubble nucleation in the inflating spacetime \\cite{CDL,LW}. Indeed this view of cosmology is supported by the present understanding of string theory, which argues for the existence of an enormous landscape of such metastable vacua \\cite{BP,S03,D,DK}. In the standard picture, both the parent and daughter vacua have three large (expanding) spatial dimensions. Then the symmetries of the de Sitter parent vacuum suggest the daughter bubble should possess a homogeneous and isotropic geometry---in particular it should possess an O(3) rotational symmetry on any homogeneous foliation \\cite{CDL}. Yet the number of large spatial dimensions may vary from vacuum to vacuum \\cite{G,B-PS-PV,CJR,B-PS-PV-2}, as is expected in string theory. In particular, one of the large spatial dimensions of the daughter vacuum may be compact in the parent vacuum. The size of this dimension can be characterized by a ``volume'' modulus that, during bubble nucleation, tunnels out of a metastable minimum, and subsequently grows to very large values. While the resulting bubble is still assumed to reflect the symmetries of the parent vacuum, the presence of the compact dimension breaks O(3) rotational invariance. Indeed, we expect the bubble geometry to be toroidal, with O(2) rotational symmetry in the two large spatial dimensions, uniformly wrapping around the compact space. This would appear as anisotropic initial conditions to an observer in the daughter vacuum. While the formerly-compact spatial dimension remains globally closed in the daughter bubble, this will not necessarily be evident to a local observer. In order for the daughter vacuum to correspond to our universe, its local evolution should approach an approximately O(3) rotationally-symmetric Friedmann-Robinson-Walker (FRW) cosmology, and the circumference of each large spatial dimension should become and/or remain much larger than the Hubble radius. In fact these conditions are easy to satisfy---the latter is akin to solving the horizon problem of classical big bang cosmology, and is accomplished by a sufficiently long period of slow-roll inflation (after bubble nucleation) along each large spatial dimension \\cite{Guth81}. (A period of inflation after bubble nucleation is required even in the standard O(3)-symmetric situation, in order to redshift away the large initial spatial curvature of the bubble.) Meanwhile, it is well known that during such inflation an initially homogeneous but anisotropic universe rapidly approaches local isotropy \\cite{BT,W}. Although the initial anisotropy rapidly redshifts away, background anisotropy present at the onset of inflation will generate statistical anisotropy in quantum fluctuations as they expand beyond the Hubble radius, and this in turn can modify the spectrum of primordial density perturbations \\cite{GCP,ACW,PPU,WKS10}. During inflation the affected modes are pushed to physical scales far beyond the local horizon; however if the duration of inflation is appropriate they will have re-entered in time to form the largest observable scales in the cosmic microwave background (CMB). Indeed, seemingly anomalous correlations have already been detected among the low-multipole CMB anisotropies \\cite{O-CTZH,SSHC,EHBGL,LM,JBEGH}, which might indicate deviations from statistical isotropy in the inflationary spectrum. (It should be noted that the significance of these ``anomalies'' is difficult to assess, and their source(s) could be non-cosmological.) In this context a number of models of anisotropy during inflation have been proposed; see e.g.~\\cite{BBK05,DDR,A-P,BM08,DGW,GMV,YS08,WKS09,DM10}. However unlike other approaches, anisotropic bubble nucleation provides concrete theoretical constraints on the form of initial conditions, and serves as a natural extension of the standard inflationary scenario. Furthermore, it provides an opportunity to confirm aspects of the landscape/multiverse hypothesis. When the background is homogeneous and isotropic, inflationary perturbations generally decouple into scalar, vector, and tensor modes (for a review of cosmological perturbation theory see e.g.~\\cite{MFB}). The same cannot be said for fluctuations about anisotropic backgrounds, which complicates the corresponding analysis. For this reason we assume that metric perturbations are suppressed, until the background geometry of the bubble has become essentially isotropic, after which standard cosmological perturbation theory can be used. As a concrete model one can imagine the spectrum of isocurvature fluctuations in a sub-dominant scalar field, which are much later converted into adiabatic density perturbations (as in the curvaton mechanism \\cite{curvaton}). Note that this simplification comes at little cost: we are interested not in the (model-dependent) amplitude or tilt of the spectrum, but in its statistical anisotropy, and this should not depend strongly on the back-reaction of the scalar on the metric. Even with this and some other simplifications, the analysis is rather complicated. Although we obtain an analytic expression for the inflationary power spectrum in terms of an appropriate set of anisotropic mode functions, we must resort to numerical evaluation to project this spectrum onto spherical harmonics. Still, we find that certain patterns are evident: whereas the standard (isotropic) picture gives a multipole correlator $C_{\\ell\\ell'mm'}=\\<\\hat{a}_{\\ell m}^{\\phantom{\\dagger}}\\, \\hat{a}^\\dagger_{\\ell'm'}\\>$ that is diagonal in both $\\ell$ and $\\ell'$ and in $m$ and $m'$, and is independent of $m$, our scenario introduces off-diagonal components in $\\ell$ and $\\ell'$ (when $\\ell-\\ell'=\\pm 2$), and introduces $m$-dependence into $C_{\\ell\\ell'mm'}$ (it is still diagonal in $m$ and $m'$). (These results are not unlike those of \\cite{ACW}, which studied a Bianchi type I anisotropic cosmology.) Our approximations limit us to a region in parameter space where the corrections to $C_{\\ell\\ell'mm'}$ are suppressed relative to the leading order terms by roughly the present-day curvature parameter $\\Omega^0_{\\rm curv}$. While this greatly constrains the size of these effects, they appear to extend to arbitrarily large $\\ell$, giving hope for statistically significant future detection \\cite{PK}. Note that because the statistically anisotropic contributions to $C_{\\ell\\ell'mm'}$ are suppressed by $\\Omega^0_{\\rm curv}$. We have focused on statistical anisotropies in inflationary perturbations, however in this scenario the background spatial curvature of the bubble is itself anisotropic. In particular, the spatial geometry is flat along one direction and open in the two-dimensional planes orthogonal to that direction. While our computation of the inflationary spectrum accounts for this spatial curvature anisotropy, our projection onto spherical harmonics does not. Indeed, anisotropic spatial curvature induces anisotropic expansion, which affects the free streaming of photons and thus deforms the surface of last scattering, along with our perception of angular scales on it \\cite{DD}. This affects the appearance of the inflationary spectrum, inducing corrections to the observed multipole correlator $C_{\\ell\\ell'mm'}$. A full understanding of the observable signatures of this model involves combining both of these effects; this is left to future work \\cite{B-PS2}. The remainder of this paper is organized as follows (an effort has been made to make the major sections self-contained). We study the dynamics of anisotropic bubble nucleation, within the context of a toy model of modulus stabilization, in Section \\ref{sec:bubble}. The primary goal of this section is to obtain the instanton boundary conditions that determine the initial conditions for the subsequent bubble evolution. However, because our compactification of one extra dimension with positive vacuum energy is (to our knowledge) novel, we present the model in some detail. In Section \\ref{sec:background} we describe the salient features of the post-nucleation, background evolution of the bubble, focusing on the (pre-)inflationary geometry (including obtaining a simple analytic approximation of the metric). In Section \\ref{sec:perts} we compute the spectrum of inflationary perturbations in a massless scalar field. To better understand the observational signatures of this spectrum, we here also perform a basic analysis of its projection onto spherical harmonics. Some issues pertaining to the plausibility of observing this scenario are discussed in Section~\\ref{sec:likelihood}. Meanwhile a final summary, including a discussion of various avenues for future work, is provided in Section~\\ref{sec:discussion}. \\vspace{11pt} \\noindent Preliminary accounts of this work were presented in \\cite{MPS}. ", "conclusions": "\\label{sec:discussion} The present understanding of string theory and inflationary cosmology points to a picture of spacetime containing countless bubbles endlessly nucleating within an eternally-inflating background. In the context of the string landscape, the complete set of bubbles contains a wide range of low-energy physics, including different numbers of compact dimensions (in fact it is the enormous variety of compactifications that inspires the diversity of the landscape in the first place). It is therefore a priori possible that our bubble, containing three large (expanding) spatial dimensions, nucleated within a vacuum containing only two such dimensions. For instance, the eternally-inflating vacuum in which our bubble nucleates could contain a compact dimension, the size of which is governed by a metastable modulus that tunnels through a potential barrier upon bubble nucleation. The tunneling instanton and initial bubble geometry then respect reduced symmetry from the (3+1)-dimensional perspective, due to the presence of the additional compact dimension. Yet, as expected, a round of slow-roll inflation within the bubble is sufficient to redshift away the initial anisotropy and curvature, creating an O(3)-symmetric FRW cosmology consistent with the observed universe. Nevertheless, if inflation within the bubble does not last too long, effects of the initial anisotropy may be observable. We here focus on one such effect: the generation of statistical anisotropy among large-scale inflationary perturbations. We find that, when projected onto a two-sphere approximating the surface of last scattering, the inflationary spectrum generates a multipole correlator $C_{\\ell\\ell'mm'}$ that features (in an appropriate coordinate system) off-diagonal elements in $\\ell$ and $\\ell'$ (when $\\ell-\\ell'=\\pm 2$), as well as dependence on the multipole moment $m$ (it is still diagonal in $m$ and $m'$). These effects are suppressed relative to the statistically-isotropic components of $C_{\\ell\\ell'mm'}$ by the present-day curvature parameter $\\Omega^0_{\\rm curv}$, but appear to extend to arbitrarily large $\\ell$. There are a number of remaining issues to be explored. Most importantly, as a first approach to the problem we have ignored the effects of spatial curvature and expansion anisotropy on the free streaming of photons from the surface of last scattering to the point of present detection. In fact, anisotropic spatial curvature sources anisotropic expansion, which in turn deforms the surface of last scattering away from the surface on which we project the inflationary spectrum, in addition to perturbing the trajectories of geodesics as they radiate away from the point of observation. A full understanding of the observable signatures of anisotropic bubble nucleation requires combining both of these effects. Also, we have ignored metric perturbations, focusing on the spectrum of a subdominant scalar field and assuming its isocurvature perturbations translate directly into adiabatic density perturbations. While we do not expect this to have a large effect on the spectrum of statistical anisotropies (compared to for instance the standard scenario where the primordial perturbations are sourced by the inflaton itself), it does not allow us to study the tensor perturbations generated during inflation. Because scalar and tensor metric perturbations in general do not decouple in an anisotropic background, there are possibly interesting correlations between these signals. Ignoring metric perturbations also hides any interesting effects that might come from fluctuations in the bubble wall itself. It would also be interesting to explore the nature and degree of non-Gaussianity implied by the existence of statistical anisotropy among inflationary perturbations. Finally, another potential signature of multiverse cosmology is observable bubble collisions, see e.g.~\\cite{JJBP,GGV,AJS,worldscollide,FKNS}. It would be interesting to understand whether the reduced symmetry at early times of anisotropic bubble nucleation affects the spectrum of bubble collisions on typical observer's sky, or if there is any special signature of collisions with bubbles containing a reduced number of large spatial dimensions \\cite{MPS10}." }, "1003/1003.0039_arXiv.txt": { "abstract": "We present observations of two strongly lensed $z\\sim5$ Lyman-$\\alpha$ Emitting (LAE) galaxies that were discovered in the Sloan Giant Arcs Survey (SGAS). We identify the two sources as SGAS J091541+382655, at $z=5.200$, and SGAS J134331+415455 at $z=4.994$. We measure their AB magnitudes at $(i,z)=(23.34\\pm0.09,23.29\\pm0.13$) mags and $(i,z)=(23.78\\pm0.18,24.24^{+0.18}_{-0.16}$) mags, and the rest-frame equivalent widths of the Lyman-$\\alpha$ emission at $25.3\\pm4.1$\\AA~and $135.6\\pm20.3$\\AA~for SGAS J091541+382655 and SGAS J134331+415455, respectively. Each source is strongly lensed by a massive galaxy cluster in the foreground, and the magnifications due to gravitational lensing are recovered from strong lens modeling of the foreground lensing potentials. We use the magnification to calculate the intrinsic, unlensed Lyman-$\\alpha$ and UV continuum luminosities for both sources, as well as the implied star formation rates (SFR). We find SGAS J091541+382655 and SGAS J134341+415455 to be galaxies with (L$_{Ly-\\alpha}$, L$_{UV})\\leq(0.6$L$_{Ly-\\alpha}^{*}, 2$L$_{UV}^{*}$) and (L$_{Ly-\\alpha}$, L$_{UV})=(0.5$L$_{Ly-\\alpha}^{*}, 0.9$L$_{UV}^{*}$), respectively. Comparison of the spectral energy distributions (SEDs) of both sources against stellar population models produces estimates of the mass in young stars in each galaxy: we report an upper limit of M$_{stars} \\leq 7.9^{+3.7}_{-2.5} \\times 10^{7}$ M$_{\\sun} h_{0.7}^{-1}$ for SGAS J091531+382655, and a range of viable masses for SGAS J134331+415455 of $2\\times10^{8}$ M$_{\\sun} h_{0.7}^{-1} <$ M$_{stars} < 6\\times10^{9}$ M$_{\\sun} h_{0.7}^{-1}$. ", "introduction": "Understanding the evolution of galaxies -- especially the first generation of galaxies -- remains one of the most important topics in astrophysics and cosmology. Populations of galaxies in the distant universe are identified over a wide range of wavelengths, including Distant Red Galaxies (DRGs), Ultra Luminous Infra-Red Galaxies (ULIRGs) and Sub-Millimieter Galaxies (SMGs). Many efforts to study the properties of high redshift galaxies at optical wavelengths focus on two distinct classes selected by their rest-frame UV properties: 1) Lyman-Break Galaxies (LBGs) and 2) Lyman-$\\alpha$ Emitters (LAEs). LBGs are selected via deep wide-band photometry, identified by the `Lyman limit' continuum break that appears at 912\\AA~ in the rest frame \\citep{steid1996a,steid1996b,lowenthal1997} -- though for sources at higher redshift this spectral break moves redward, approaching 1216\\AA~ in the rest frame due to the Lyman-$\\alpha$ forest \\citep{steid1987,rauch1998} absorption by intergalactic neutral hydrogen. LAEs are selected by either narrow-band imaging \\citep{cowie1998,rhoads2000,rhoads2003,ajiki2004,gawiser2006,yamada2005} or blind spectroscopy \\citep{kurk2004,sawicki2008} tuned to detect Lyman-$\\alpha$ line emission redshifted into near-ultraviolet, optical, or near-infrared wavelengths. Over the past decade large samples of LBGs and LAEs have driven studies of star-forming galaxies at $z\\gtrsim2.5$. Surveys for LBGs and LAEs are efficient for collecting statistical samples of high-redshift galaxies, but at $z\\gtrsim3$ they produce objects which are generally too faint to have their galactic continuum emission studied spectroscopically. The standard approach for studying the properties of these galaxy samples relies on stacking the photometric signal from many objects and fitting the observed mean SED against a variety of stellar population synthesis models in order to constrain parameters such as the ages and masses of the underlying stellar populations, as well as the amount of dust extinction \\citep{shapley2003,chary2005,pirzkal2007,lai2007,fink2008,fink2009a,nilsson2009,yabe2009}. In principal, stellar population synthesis modeling can also provide information about dust properties (i.e. the shape of the dust law) and metallicity, but even the stacked SED signal at $z\\gtrsim3$ is insufficient to constrain these additional parameters with much confidence. Broadly speaking, galaxies selected as LBGs are believed to sample more massive star-forming galaxies with an underlying older stellar population, and possibly higher dust content, while LAE selected galaxies tend to be lower mass galaxies with low metallicities and very little dust \\citep{giaval2002,venem2005,gawiser2007}. Hubble Space Telescope imaging studies of high redshift LAE galaxies imply that these sources are compact, and likely either disk-like or irregular in structure \\citep{pirzkal2007,tanigu2009}. Recently \\citet{fink2009c} modeled individual SEDs of 14 bright $z\\sim4.5$ LAEs from the Chandra Deep Field South and found a broad range in stellar population age, stellar mass, and dust extinction, which suggests that stacking SED analyses of high redshift galaxies may not be the best approach. The main hurdle involved in studying any high redshift source is the general lack of signal. Distant galaxies are faint and therefore difficult to detect, and those which are identified are rarely -- if ever -- amenable to detailed follow-up. Furthermore, those sources which are bright enough to be studied individually are drawn from the extreme bright tail of the luminosity function of high redshift galaxies, and are therefore not necessarily representative of the bulk of the populations. In this paper we present two serendipitously discovered, strongly lensed high redshift galaxies: \\object{SGAS J091541+382655}, spectroscopically confirmed at $z=5.200\\pm0.001$, with $r_{AB}=24.68\\pm0.25$, $i_{AB}=22.92\\pm0.09$ and $z_{AB}=22.75\\pm0.13$ mags, and \\object{SGAS J134330+415455} spectroscopically confirmed at $z=4.994\\pm0.001$, with $r_{AB}\\geq25.47$ mags, $i_{AB}=23.36\\pm0.18$ mags and $z_{AB}=23.70^{+0.18}_{-0.16}$ mags. Both objects have $riz$ colors and magnitudes that satisfy selection criteria for $r-$band dropouts in $z\\sim5$ dropout surveys, as well as Lyman-$\\alpha$ equvialent widths (see Section 2) sufficiently large to be selected in surveys for Lyman-$\\alpha$ excess. At $z\\gtrsim5$ these objects are the two brightest LAEs in the literature to date, and both sources are projected on the sky within ~$30\\arcsec$ of the cores of confirmed strong lensing galaxy clusters. This means that the sources -- once corrected for the lensing magnification -- are intrinsically much fainter than the observed flux implies, and therefore provide a rare opportunity to study individual LAE properties at the fainter end of the luminosity function. There is a small but growing number of magnified galaxies at high redshift that are excellent candidates for high resolution spectroscopic follow-up \\citep{koester2010,wuyts2010}; some of these galaxies have been observed in detail at optical and near-infrared wavelengths, including cB58 \\citep{pettini2000}, \"the 8 O'clock Arc\" \\citep{allam2007,fink2009b}, and \"the Cosmic Eye\" \\citep{smail2007,siana2009,quider2010}. The two galaxies discussed in this paper are the first $z\\sim5$ galaxies that present similar opportunities for detailed study via follow-up spectroscopy. Where necessary we calculate cosmological distances assuming a flat cosmology with $H_{0}=70$ km s$^{-1}$ Mpc$^{-1}$, and matter density $\\Omega_{M}=0.3$. All magnitudes are AB. ", "conclusions": "We have identified two lensed Lyman-$\\alpha$ emitting galaxies at $z\\sim5$ near the cores of strong lensing selected galaxy clusters. These sources are among the brightest galaxies identified at such high redshift, but their intrinsic luminosities are much lower than the observed flux due to magnification by the gravitational potential of foreground galaxy clusters. We use the available data to investigate the underlying stellar populations for these galaxies and find that the light -- continuum and line emission -- for SGAS 091541+382655 likely originates from a population of young stars with low dust content. Both sources are in the process of undergoing active star formation. Our analysis of these two LAEs corroborates our current understanding of the nature of Lyman-$\\alpha$ emitting galaxies at high redshift, and the large magnification of these sources due to gravitational lensing makes them excellent candidates for studying the individual properties of galaxies on the faint end of the L$_{Ly-\\alpha}$ and L$_{UV}$ luminosity functions at $z\\sim5$. We encourage efforts to followup these sources aggressively on 8-10m class telescopes in order to better study the properties of the underlying stellar populations via continuum light. These sources are also excellent targets for space-based observations, both with current $HST$ instruments and with $JWST$ in the future." }, "1003/1003.0808_arXiv.txt": { "abstract": "Some of radiopulsars have anomalous braking index values $n = \\Omega \\ddot{\\Omega} / \\dot{\\Omega}^{2} \\sim \\pm ( 10^{3} \\div 10^{4} ) $. It is shown that such anomalous values may be related to nondipolar magnetic field. Precession of a neutron star leads to rotation (in the reference frame of the star) of its angular velocity $\\vec{\\Omega}$ around the direction of neutron star magnetic dipole moment $\\vec{m}$ with an angular velocity $\\vec{\\Omega}_{p}$. This process may cause the altering of electric current flowing through the inner gap and consequently the current losses on the time scale of precession period $T_{p} = 2\\pi / \\Omega_{p}$. It occurs because the electric current in the inner gap is determined by Goldreich-Julian charge density $\\rho_{GJ} = -\\frac{ \\vec{\\Omega} \\cdot \\vec{B} }{ 2\\pi c }$, that depends on the angle between direction of small scale magnetic field and the angular velocity $\\vec{\\Omega}$. ", "introduction": "Radio pulsars have been discovered more than 40 years ago \\cite{Hewish1968} and at present many thousand papers are devoted to these objects. Despite of large progress in the understanding of processes in pulsar magnetospheres still many important questions are unclear. One of them is the value of pulsar braking indices. If a pulsar were just a simple magnet with dipolar magnetic momentum $\\vec{m}$ rotating with angular velocity $\\vec{\\Omega}$, then the braking index $n = \\ddot{\\Omega}\\Omega / \\dot{\\Omega}^{2}$ would be equal to $3$. The taking into account of evolution of angle $\\chi$ between vectors $\\vec{m}$ and $\\vec{\\Omega}$ yields $n = 3 + 2cot\\chi$ \\cite{DavisGoldstein1970,Beskin_book}. Observations of some pulsars are well within theoretical predictions. For example, the braking index of pulsar Crab is equal to 2.5 and pulsar Vela has $n\\approx 1.4$ \\cite{Beskin_book}. So it seems that more sophisticated magnetospheric models will be able to remove this discrepancy, see for example \\cite{Melatos1997,Timokhin2005,Timokhin2006,Contopoulos2006,Contopoulos2007, Timokhin2007_Force_free,Timokhin2007_Differentially}. However, many isolated pulsars have very large positive or, sometimes, negative braking indices upto $|n| \\sim 10^{4}$. Sometimes such large values may be due to unobserved glitches or timing noise \\cite{Alpar2006}. For example, in \\cite{Alpar2006} it is shown that all negative braking indices may be associated with unresolved glitches which has occurred between intervals of observations of the pulsars. But in some cases refined measurements of braking indices show that at least some pulsars have large and sometimes negative braking indices $n \\sim \\pm (10-10^{2})$ \\cite{Galloway1999}. For example, the pulsars B0656+14 and B1915+13 have braking indices $n \\approx 14.1 \\pm 1.4$ and $n \\approx 36.08 \\pm 0.48$, correspondingly, while pulsar B2000+32 and B1719-37 have very large negative braking indices $n \\approx -226 \\pm 4.5$ and $n \\approx -183 \\pm 10$, correspondingly \\cite{Galloway1999}. There are some explanations of such braking index values. For example, the values like $n \\sim \\pm 10$ may be related to nonstandard mechanisms of pulsar braking, like neutron star slowing down due to neutrino emission \\cite{Peng1982} or due to interaction with circumpulsar disk \\cite{Menou2001,Malov2004,Li2006}, see \\cite{Malov2001} for review of possible mechanisms and its comparison with observations. The large braking indices may be explained by the rapid changing of magnetic field. Such changes may be caused by Hall-drift instabilities \\cite{Geppert2002_Hall_drift,Geppert2002_Non_linear,Geppert2007}. Also the positive braking indices of old pulsars may be explained by the relaxation of angular velocities of neutron star crust and superfluid between two glitches \\cite{Alpar2006}. The large braking indices of some pulsars may be related to Tkachenko waves \\cite{Popov2008}. In this paper we present a some model based on works \\cite{Biryukov2006} and \\cite{Contopoulos2007_Note_on_cyclic}, see also \\cite{Biryukov2007,Link2006_Strong,Geppert2007}, where it has been shown that large braking indices may be explained by the existence of some internal cyclic process in a pulsar. ", "conclusions": "In this paper we present a some explanation of large values of braking indices of pulsars. The proposed model is based on four main assumptions: \\begin{enumerate} \\item The neutron stars have a small scale magnetic field. The strength of this field must be enough large to curve pulsar tube but enough small to allow free emission of electrons from star surface. \\item There are inner gaps in pulsar tubes and a electric current flowing across inner gaps depends on angle between small scale magnetic field and angular velocity $\\vec{\\Omega}$ of the star. \\item Braking torque depends on current that flows across inner gaps. \\item The neutron stars precess with periods like $10^{3}-10^{4}$ years. \\end{enumerate} In this paper we also neglect the contribution of outer gaps. If a some part of electric current flows through outer gaps then this part is determined only by outer gap electrodynamics and, consequently, does not depend from small scale magnetic field. It decreases the variation of current losses over the precession period and, consequently, leads to the decreasing of braking index. In the case of young pulsars like Crab and Vela, which have small braking indices, we suppose that current losses is almost fully determined by currents flowing through outer gaps and the contribution of inner gaps current is negligible. In this paper it is also assumed that magnetic dipole braking exists and does not depends on electric current flowing across inner gaps. This assumption is widely used, c.f. \\cite{Beskin2006_Be_born,Xu2007,GurevichIstomin2007,IstominShabanova2007}, but, as mentioned in \\cite{Beskin_book}, it must be treated with caution. It is shown that in the case of force free magnetosphere and absence of electric current flowing along pulsar tube the orthogonal pulsar $\\chi = \\pi / 2$ does not slow down at all \\cite{Beskin1983,Beskin1984,Shibata1999}. It gives the reason to suppose that magnetic dipole braking does not exist or, at least, must be depend on the electric current \\cite{Beskin_book}. In such case presented model can provide the only qualitative explanation of the existence of large braking indices. And the quantitative estimations, of course, will strongly depend on relation between magnetic dipole braking and current losses torques. The presence of a long period precession is the weakest point of the model. At present the precession is discovered only at a few isolated neutron stars. And these pulsars have the precession periods like $T_{p} \\sim 1-10$ years \\cite{Link2007}. The presence of pined superfluid in neutron star crust substantially increases precession speed \\cite{Shaham1977}. The precession periods $T_{p}$ larger than $(10^{2}-10^{4}) P$ may exist only when vortices of superfluid can not been pinned anywhere in star, are able to move freely and do not coexist with magnetic flux tubes \\cite{Link2006_Incompatibility,Link2007}. The small number of isolated pulsars with observed precession force us to exclude the triaxial precession and to assume that neutron star is axisymmetrical and symmetry axis coincide with magnetic dipole moment $\\vec{m}$. In this case in frame reference related to \"rigid stars\" vector $\\vec{m}$ rotate with constant angular velocity $\\vec{\\omega}$. And because of $\\vec{\\omega} \\approx \\vec{\\Omega}$ its trajectory coincides with the case of unprecessing star, see Appendix A. Also, in order to prevent the observation of such precession it is necessary to assume that pulsar tube structure does not precess too. It particularly means that pulsar tube crossection is circular or depends on only vectors $\\vec{\\Omega}$ and $\\vec{m}$ and does not depend from small scale magnetic field. Also, it means that distributions of energy of primary electrons and pair multiplicity over pulsar tube crossection are close to axisymmetrical." }, "1003/1003.0725_arXiv.txt": { "abstract": "{ Perfectly matched layers are a very efficient way to absorb waves on the outer edges of media. We present a stable convolutional unsplit perfectly matched formulation designed for the linearized stratified Euler equations. The technique as applied to the Magneto-hydrodynamic (MHD) equations requires the use of a sponge, which, despite placing the perfectly matched status in question, is still highly efficient at absorbing outgoing waves. We study solutions of the equations in the backdrop of models of linearized wave propagation in the Sun. We test the numerical stability of the schemes by integrating the equations over a large number of wave periods.} \\begin{document} ", "introduction": "\\label{intro} The choice of appropriate boundary conditions in simulations remains a significant challenge in computational physics. Amidst the vast number of options that exist, a widely invoked construct is that of the absorbing boundary, one that is expected to relieve the computational domain of outgoing waves or other structures without affecting the solution in the region of interest \\citep[for a review, see e.g.,][]{colonius}. A number of different methods exist to accomplish this, such as an attenuating `sponge' \\citep[e.g.,][]{lui,colonius}, characteristics-based boundary conditions \\citep[e.g][]{thompson}, an artificially imposed supersonic outward directed advective flow \\cite[e.g.,][]{lui}, perfectly matched layers \\citep[PMLs;][]{Berenger1994}, etc. Unfortunately, a number of these methods is afflicted with disadvantages, a primary cause of concern being that of absorption efficiency. For example, the characteristics-based boundary conditions do not perform very well when the incident waves at the boundary are significantly inclined - and a fair degree of reflection is observed. The sponges are better at absorbing the waves and are relatively easy to implement; however, the reflectivity can still be substantial. The criterion of low reflectivity is perhaps best satisfied by the PML formulation, developed first by \\citet{Berenger1994} as an absorbing layer for the Maxwell equations. The central idea of this technique lies in performing an analytic continuation of the wave vector into the complex plane. Wave vectors perpendicular to the boundary are forced to assume a decaying form in the complex plane, thereby dramatically reducing the amplitudes of the waves in the boundary PML region. In the absence of discretization errors, the PML as set out by \\citet{Berenger1994} is highly absorbent. After numerical discretization however, there are weak associated reflections. An important issue in the classical split formulation of \\citet{Berenger1994} is that the absorption efficiency decreases rapidly at grazing incidence \\citep[e.g.,][]{CoMo98a,WiRa00}. There have been numerous advances in constructing stable analytic continuations of wave vectors in the boundary region. In particular, the stable convolutional PMLs \\citep[C-PML; e.g.,][]{roden2000,FeVi05,Komatitsch2007,DrGi07b}, contain a Butterworth filter inside the PML, thereby dramatically improving the absorption efficiency at grazing incidence. This formalism was adopted to the study of waves in anisotropic geophysical media by \\citet{Komatitsch2007}; the numerical stability of said technique was studied by \\citet{Komatitsch2007,Meza2008}. Extensions to the poro- and viscoelastic cases have been introduced by \\citet{MaKoEz08} and \\citet{MaKo09}, respectively. Our goal in this contribution is to develop a C-PML for the 3D linearized ideal MHD equations in stratified media. Astrophysical media such as stellar interiors and atmospheres may be strongly stratified and highly magnetized. Simulating wave propagation in such environments is of interest because in understanding the effects of stratification and magnetic fields on the oscillations, we may be able to better interpret observations and constrain certain properties of the object in question. In particular, we focus on the Sun, a star that has been well studied and for which high-quality observations exist. The interior of the Sun is opaque to electromagnetic waves - consequently, only photons that arrive from very close to the surface (photosphere) of the Sun are visible. Helioseismology is the inference of the internal structure and dynamics of the Sun by observing the surface manifestation of its acoustic pulsations \\citep[e.g.,][]{dalsgaard02,gizon05,gizon09}. Armed with accurate interaction theories of waves and measurements of the acoustic field at the surface, one can attempt to determine the sub-surface structure and dynamics of various solar features such as sunspots \\citep[see e.g.,][]{gizon_etal_2009}, large-scale meridional flow, interior convective length scales, etc. In order to construct these interaction theories, it is important to simulate and study small amplitude (linear) wave propagation in a solar-like medium - filled with scatterers like sunspots or mean flows. The time scales of wave propagation are typically much smaller than the rate at which the scatterer itself evolves; thus one may invoke the assumption of time stationarity of the background medium. Numerical simulations of waves in solar-like media have been performed by numerous groups \\citep[e.g.,][]{werne04,hanasoge1,shelyag06,khomenko06,cameron07,parchevsky_method,Hanasoge_couvidat08, hanasoge_mag,cameron08}. Two groups in particular \\citep{khomenko06, parchevsky_method} currently utilize the classical split PML formulation in order to solve the wave equations in stratified media. There exist drawbacks in these formulations. For example, \\citet{khomenko06} have noted that there are long term instabilities associated with their method. The technique discussed in \\citet{parchevsky_method} involves the addition of a small arbitrary constant (see Eq. [21] and related discussion of their article) that could possibly be acting as a sponge; it is unclear if their method is perfectly matched after the introduction of this constant. The modal power spectrum that \\citet{parchevsky_method} display in Figure 9 of their article shows substantial reflected power from their lower boundary, atypical of perfectly matched formulations. The plan of this article is as follows: in Section 2, we recall the linearized ideal MHD equations and discuss the C-PML method as was previously developed for the seismic wave equations. The C-PML is then extended to the stratified Euler/MHD equations in Section 3 and results from numerical tests are discussed. We summarize and conclude in Section 4. ", "conclusions": "We have developed a Convolutional Perfectly Matched Layer (C-PML) for the stratified linearized Euler equations and a highly efficient absorption layer for the ideal MHD equations. This boundary formulation is quite useful for the calculation of wave propagation in astrophysical plasmas, where stratification and magnetic fields abound. The absorption layer for MHD waves is a slightly altered version of the convolutional formulation of \\citep[e.g.,][]{roden2000}, requiring an extra sponge-like term in order to stabilize the system. While the boundary methods as discussed here are perfectly matched for the stratified Euler equations, the inclusion of the sponge term casts some doubt as to whether the same could be said of the MHD equations. Simulations with the absorption layers for the stratified MHD and Euler equations were performed and found to be stable over a time length of 300 wave periods. With increasingly sophisticated boundary formulations, such as that of \\citet{Hu2008}, who developed the PML for turbulent flows, we may in the near future be able to extend these techniques to fully non-linear MHD turbulence/dynamo calculations." }, "1003/1003.1665_arXiv.txt": { "abstract": "{The physical conditions of the gas forming the narrow line regions (NLR) in active galactic nuclei (AGN) have been extensively studied in the optical band. Recently, detailed X-ray studies have shown how the emission in the 0.1-2 keV band detected in Seyfert 2 galaxies is associated to gas lying close to or associated with the NLR.} {We take advantage of the spectacular extension ($\\sim$15'') of the NLR in the type II Seyfert galaxy NGC~5252 and of the complementary characteristics of $XMM$--$Newton$ and $Chandra$ to investigate the physical conditions of the gas in this galaxy. } {The X-ray data from $XMM$--$Newton$ are used to define the spectral properties of the ionising nuclear source. The $Chandra$ data are used to trace the spatial characteristics of the soft X-ray emission. This information is then compared to the optical HST characteristics of the NLR in NGC~5252.} {The X-ray spectrum of the nucleus of NGC~5252 is intrinsically flat ($\\Gamma$$\\sim$1.4-1.5), and absorbed by neutral gas with a column density N$_{H}$$\\sim$10$^{22}$ cm$^{-2}$. Below $\\sim$1 keV a soft excess is detected. The high-resolution spectrum obtained with the $XMM$--$Newton$ RGS shows the presence, in the 0.2-1.5 keV range, of emission lines which strongly indicate that the soft X-ray component is essentially due to ionised gas. Moreover, the soft X-ray emission is spatially resolved around the nucleus and well overlaps the images obtained in narrow optical bands centered around the [OIII] emission line at 5007$\\AA$. The [OIII]/soft-X flux ratios along the ionisation cones is basically constant. This indicates that the electron density does not significantly deviates from the r$^{-2}$ law (constant ionisation parameter) moving outward from the nucleus.} { This result combined with previous optical studies suggest two plausible but different scenarios in the reconstruction of the last $\\sim$30000 years history of the central AGN. The most promising one is that the source is indeed a ``quasar relic'' with steady and inefficient energy release from the accretion of matter onto the central super-massive black-hole. This scenario is suggested also by the flat nuclear X-ray spectrum that suggests an advection dominate accretion flow (ADAF) like emission mechanism.} ", "introduction": "Obscuration of the nuclear emission in type~II AGN allows the study of soft X-ray spectral components, which are normally outshone by the direct component in type~I unobscured objects. It has been well known since the early day of X-ray spectroscopy that excess emission above the extrapolation of the absorbed nuclear radiation is present in almost all bright Seyfert~2s (Turner et al. 1997). This excess appears smooth when measured with instruments with moderate energy resolutions such as CCD. However, high-resolution (grating) measurements with {\\it Chandra} and XMM-Newton revealed that this excess is generally due to a blending of strong recombination lines from He- and H-like transitions of elements from Carbon to Nitrogen (Sako et al. 2000, Sambruna et al. 2001, Kinkhabwala et al. 2002, Armentrout et al. 2007). X-ray spectral diagnostics (Kinkhabwala et al. 2002, Guainazzi \\& Bianchi 2006) and a close morphological coincidence between the soft X-rays and the [OIII] in Extended Narrow Line Regions (ENLR; Bianchi et al. 2006, Bianchi et al. 2010) strongly indicate that the gas is photoionised by the AGN, with an important role played by resonant scattering. In this context, NGC~5252 represents an extraordinary laboratory to study the feedback between the AGN output and circumnuclear gas on kpc scale, thanks to its spectacular ionisation cones (Tadhunter \\& Tsvetanov 1989). NGC 5252 is classified as Seyfert 1.9 (\\cite{ost}) S0 (\\cite{dev}) nearby (z=0.023,) galaxy (N$_{H, Gal}$=2.14$\\times$10$^{+20}$ cm$^{-2}$, Dickey \\& Lockman, 1990). Small radio jets (r$\\sim$4'') have been detected and found to be aligned with the ionisation cones (\\cite{wilson94}). Nonetheless, the host galaxy luminosity (M$_{R}$$\\sim$-22. \\cite{capetti}), mass (M$_{bulge}$$\\sim$2.4$\\times$10$^{11}$M$_{\\sun}$, \\cite{marc}) and the mass of the central super-massive black-hole (M$_{BH}$$\\sim$10$^{9}$M$_{\\sun}$, \\cite{capetti}) are more typical of quasar than Seyfert galaxies. These pieces of evidence led \\cite{capetti} to speculate that NGC 5252 is most probably to be considered a QSO relic. This view is in agreement with \"downsizing\" scenarios about the evolution of super-massive black-hole (SMBH) in cosmic times. Accordingly with these scenarios, most massive SMBHs formed and evolved earlier than lower mass ones. Ionisation cones are one of the strongest argument in favour of the Seyfert unification scenarios (Antonucci 1993). For this reason, NGC~5252 is also an important laboratory to test AGN geometrical models. From a diferent point of view, AGN activity has been recognized, since a while as a key component of the SMBH host galaxy co-evolution and AGN feedback is likely to self-regulate or be responsible of the observed properties (Menci et al. 2004). The very existence of ionization cones witness that feedback/winds are or were active and thus these sources are ideal laboratories for feedback. X-ray measurements allows to directly link the properties of the gas emitting optical lines with the intrinsic AGN power, which in type~II AGN can be truly measured only at energies larger than the soft photoelectric cut-off due to the AGN obscuring matter. Furthermore, the morphological coincidence between X-rays and optical emission in ENLR (Bianchi et al. 2006) points to a fundamental physical link between the two wavebands. They need to be studied simultaneously in order to derive the correct energy budget in the ionisation cones. Prompted by these motivations, we have performed deep X-ray observations of NGC~5252 at the highest spatial and spectral resolution currently available with {\\it Chandra} and XMM-Newton. The results of these observations are the subject of this paper. ", "conclusions": "The soft X-ray emission of NGC~5252 is clearly extended and ACIS images demonstrate that the spectacular ionisazion cones observed in [OIII] have counterparts in the 0.1-1 keV band. The cumulative soft X-ray spectrum observed by $XMM$--$Newton$ is described by a soft power-law ($\\Gamma$$\\sim$3). The ACIS images suggest that this is probably due to a blend of emission lines that mimics such steep power-law as demonstrated in other type II Seyferts like NGC 1068, Circinus galaxy and Mrk 3 (\\cite{kin02}, \\cite{brink02}, \\cite{ogle03}, \\cite{sam01}, \\cite{sako00}, \\cite{b05}, \\cite{pp05}). This scenario is supported also by the detection in the RGS high resolution spectrum of three emission lines, of CV, OVII and OVIII, probably due to photoionised gas. This is consistent also by previous optical studies that excluded collisional ionisation along the cones of NGC~5252 (Tsvetanov et al. 1996). Moreover, the presence of {\\it in situ} ionisation sources due to shocks formed by large scale outflows interacting with the interstellar matter has been excluded (\\cite{morse98}). Thus the source of ionising photons is most probably the nucleus. Under this assumption, we can use the imaging of the arcs to study the physical condition of the gas along the ionisation cones. In particular, the constant of [OIII]/(0.5-2 keV) flux ratio along the ionisation cones within the inner\\footnote{The outer arcs and filaments (\\cite{tad}) are most probably too weak to be detected in X-rays. Considering the extension of the outer [OIII] arcs, the minimum detectable flux between 0.1-1 keV is F$_{0.1-1 keV}$$\\sim$5$\\times$10$^{-15}$erg s$^{-1}$ cm$^{-2}$ while, assuming a constant [OIII]/soft X-ray ratio, the expected flux should be $\\sim$10 times lower. } $\\sim$1.5 kpc suggests a r$^{-2}$ law for the ion density. Optical spectroscopic studies (Acosta-Pulido et al. 1996) suggest that the radial dependency of the ionization parameter\\footnote{U=$\\frac{L}{4 \\pi \\rho r^{2}}$, where L is the source's luminosity, $\\rho$ is the density of the ionized gas, $r$ is the distance between the source of ionizing photons and the ionized matter.} U follows a different law in the south-east ($U \\propto r^{-0.4}$) with respect to the north-east ($U \\propto r^0$) cone. The authors speculate that the intrinsic behavior should be the one shown in the former, while the radial-independence of the ionization parameter in the latter may be due to a ``conspiracy'' introduced by the existence of two counterotating disks of gas (\\cite{morse98}): one is coplanar to the stellar one, and another is inclined by $\\sim$40$^{\\circ}$. Morse et al. (1998) speculated that the prominence of the southeast [OIII] cone in the nuclear regions is due to the fact that this component is seen directly, while the northeast [OIII] cone is seen through the gas of the other disk. If so, the absorption due to this component could alter the line ratios presented by \\cite{ap96} and thus the correct behavior of U should be the one derived from the south-east cone. U$\\propto$r$^{-0.4}$ implies that the luminosity of the nucleus increased by a factor $\\Delta$L$\\sim$3-6 in the last $\\Delta$t$\\sim$5000 years. These numbers become $\\Delta$L$\\sim$10-30 and $\\Delta$t$\\sim$30000 years if we further assume that the U and the ion density laws are still valid up to 10 kpc from the nucleus, i.e. where the optical cones are still detectable in [OIII] but not in X-rays. On the contrary, having U constant and $\\rho \\propto$ r$^{-2}$ would imply that L has remained constant during the last 5000 (30000) years. This is consistent with the ``quasar-relic'' scenario proposed by \\cite{capetti}. These authors suggested that the nucleus of NGC~5252 is indeed the ``relic'' of a nucleus that already experienced the activity phase in the past and that now persists in an almost quiescent phase. This is suggested by the high mass of the SMBH (M$_{BH}$=10$^{9}$M$_{\\sun}$, Capetti et al. 2005) that indicates that the nucleus has already accreted in the past, the low Eddington ratio (L$\\sim$10$^{-3}$L$_{Edd}$, assuming the bolometric correction from Marconi et al. (2004), L$_{hard-x}$$\\sim$(1/22)$\\times$L$_{bol}$), and the early type (S0) morphology of the AGN host galaxy. In literature it is also reported that the optical emission line ratios in the inner 30\" are typical of LINERS (\\cite{gon}), thus suggesting that a low efficiency engine is acting at the nucleus of the source. It is worth noting that also the detection of two counterotating disks suggests that NGC~5252 is a \"quasar-relic\". These disks are tracers of a major merging event that occurred, most probably, more than 10$^8$ years ago, since the stellar disk of NGC~5252 is undisturbed. If the merging event triggered a phase of AGN activity (see Jogge 2006, and references therein for a discussion on this topic), we can expect that it lasted few/some $\\sim$10$^{7}$ years (\\cite{mar}; \\cite{ste}; \\cite{jac}; \\cite{gon2}) after which the source has persisted in a quiescent state. Finally, it is worth noting that the spectrum of the nucleus hosted by NGC~5252 is confirmed to be quite flat (Cappi et al. 1996). As shown, if modeled with a simple absorbed power-law its photon index points to a very hard spectrum ($\\Gamma$$\\sim$1). The low EW ($\\sim$50 eV) of the neutral (E$_{FeK\\alpha}$$\\sim$6.4 keV) iron line is consistent with what expected if the FeK$\\alpha$ line is produced via transmission in the observed column (N$_{H}$$\\sim$2$\\times$10$^{22}$ cm$^{-2}$, Makishima 1986) thus excluding a reflection dominated spectrum. To reconcile, at least marginally, the hardness of the NGC~5252 nuclear spectrum, with previous results for Seyfert galaxies ($\\Gamma$$\\sim$1.5-2.5; Turner \\& Pounds, 1989; Nandra \\& Pounds, 1994; Smith \\& Done, 1996; Dadina 2008), we must invoke complex absorption models involving partial covering of the source and/or the presence of ionised absorbers along the line of sight. In this case the spectral index becomes $\\Gamma$$\\sim$1.4-1.5. It is interesting to note that the flat photon index may be a further clue suggesting that the X-rays may be produced in an ADAF, Narayan \\& Yi, 1994) as expected in a ``quasar-relic''." }, "1003/1003.2270_arXiv.txt": { "abstract": "We present analysis of the three-dimensional shape of intracluster gas in clusters formed in cosmological simulations of the $\\Lambda$CDM cosmology and compare it to the shape of dark matter (DM) distribution and the shape of the overall isopotential surfaces. We find that in simulations with radiative cooling, star formation and stellar feedback (CSF) intracluster gas outside the cluster core ($r\\gtrsim 0.1r_{500}$) is more spherical compared to non-radiative (NR) simulations, while in the core the gas in the CSF runs is more triaxial and has a distinctly oblate shape. The latter reflects the ongoing cooling of gas, which settles into a thick oblate ellipsoid as it loses thermal energy. The shape of the gas in the inner regions of clusters can therefore be a useful diagnostic of gas cooling. We find that gas traces the shape of the underlying potential rather well outside the core, as expected in hydrostatic equilibrium. At smaller radii, however, the gas and potential shapes differ significantly. In the CSF runs, the difference reflects the fact that gas is partly rotationally supported. Interestingly, we find that in non-radiative simulations the difference between gas and potential shape at small radii is due to random gas motions, which make the gas distribution more spherical than the equipotential surfaces. Finally, we use mock {\\sl Chandra} X-ray maps to show that the differences in shapes observed in three-dimensional distribution of gas are discernible in the ellipticity of X-ray isophotes. Contrasting the ellipticities measured in simulated clusters against observations can therefore constrain the amount of cooling in the intracluster medium and the presence of random gas motions in cluster cores. ", "introduction": "\\label{section:intro} In the prevailing, hierarchical cold dark matter (CDM) paradigm of cosmological structure formation, galaxy- and cluster-sized CDM halos are formed via accretion and merging with smaller halos. The CDM paradigm predicts that DM halos are generally triaxial and are elongated along the direction of their most recent major mergers. The triaxiality of DM halos has been demonstrated in a number of studies using numerical simulations \\citep{frenk_etal88,dubinski_carlberg91, warren_etal92, thomas_etal98,jing_suto02,suwa_etal03,hopkins_etal05,kasun_evrard05, allgood_etal06, bett_etal07, gottloeber_yepes07,paz_etal08} and arises due to anisotropic accretion and merging along filamentary structures. The degree of triaxiality strongly correlates with the halo formation time \\citep[e.g.,][]{allgood_etal06,ho_etal06,wray_etal06}, which implies that at a given epoch more massive halos are more triaxial. For the same reason, triaxiality is sensitive to the linear structure growth function and is higher in cosmological models in which halos form more recently \\citep[e.g.,][]{maccio_etal08}. Although shapes of DM halos have been studied extensively in dissipationless $N$-body cosmological simulations, DM shape is difficult to probe observationally, though some handle on shape is provided by lensing studies \\citep[e.g.,][]{hoekstra_etal04,parker_etal07,rozo_etal07,evans_bridle09,hawken_bridle09}. Moreover, it is well known that including baryons in simulations modifies the shapes of DM halos, especially in the case of significant gas dissipation during galaxy formation \\citep[][see \\citeauthor{debattista_etal08} \\citeyear{debattista_etal08} and \\citeauthor{valluri_etal10} \\citeyear{valluri_etal10} for discussion of the physical nature of this effect]{katz_gunn91,evrard_etal94,dubinski_94,kazantzidis_etal04, springel_etal04, hayashi_etal07,tissera_etal10}. It is therefore of paramount importance to examine predictions for halo shapes using cosmological simulations that include gas dynamics and dissipative processes accompanying galaxy formation. Further, the shape of the gas itself can be examined in such simulations and compared to the shape of the underlying potential, sourced predominantly by DM. Gas is expected to follow isopotential surfaces in hydrostatic equilibrium, so simulations may test whether cluster gas is in equilibrium on average and whether it can be used as a reliable tracer of the shape of the underlying potential \\citep{buote_tsai95,lee_suto03,flores_etal07,kawahara10}. Probing the shape of the gravitational potential via gas, as was first suggested by \\citet{binney_strimpel78} and observationally tested by \\citet{fabricant_etal84} and \\citet{buote_canizares96}, can open interesting avenues for using the shapes of DM halos around observed galaxy clusters to both test the CDM paradigm and constrain the amount of halo gas that dissipated and was converted into stars during halo formation. This is particularly relevant for galaxy clusters, where high-quality X-ray imaging data now exists for large samples of clusters. \\citet{kawahara10} recently analyzed axis ratios of X-ray clusters from the {\\em XMM-Newton} catalogue and found relatively good agreement with the CDM predictions of \\citet{jing_suto02} based on dissipationless simulations of a large cluster sample, confirming findings of \\citet{buote_tsai95} and \\citet{flores_etal07} based on a single simulated system.\\footnote{A similar test of the shape determined using the observed projected galaxy distribution of clusters have been presented by \\citet{plionis_etal06}.} \\citet{buote_tsai95} were also the first to show that for the simulated cluster they studied the shape of the X-ray isophotes reflected the shape of the underlying three-dimensional gas distribution and potential. \\citet{fang_etal09} compared the ellipticities of X-ray surface brightness isophotes of clusters simulated with and without radiative cooling and star formation. Focusing on a single cluster from the sample we analyze in this paper, they showed that the shape of gas can be quite flattened when gas cools significantly and settles into rotating thick disk. They also showed that this flattening is detectable in the shape of X-ray isophotes. \\citet{fang_etal09} also argued that the flattened shape of the gas distribution in the simulated cluster implies that gas does not trace potential in the inner regions due to rotational support. Their results therefore demonstrate that shapes of X-ray isophotes in cluster cores are a useful diagnostic of amount of cooling and gas motions in cluster cores. \\citet{fang_etal09} have also compared isophote shapes for synthetic {\\it Chandra} observations of a sample of clusters simulated with cooling to observations and concluded that ellipticity profiles in runs with radiative cooling do not match observations. They attributed the discrepancy to ongoing, significant cooling in the cores of simulated clusters that is absent from the cores of real clusters \\citep[e.g.,][]{peterson_fabian06}. In this paper we present analysis of the three-dimensional shapes of intracluster gas, DM, and underlying gravitational potential using high-resolution cosmological simulations of galaxy clusters formed in the $\\Lambda$CDM cosmology. Our work extends the work of \\citet{fang_etal09} by presenting more detailed analysis of three-dimensional shape profiles of gas, DM, and gravitational potential for the full sample of clusters. In addition, we focus on the effects of cumulative cooling and dissipation during the entire cluster evolution on the shape of potential and gas distribution at intermediate radii ($r>0.2r_{500}$), where dissipation makes potential more spherical \\citep[e.g.,][]{kazantzidis_etal04}. This effect was not investigated in \\citet{fang_etal09}. We show explicitly that dissipation leads to more spherical shapes for ICM gas outside cluster cores ($0.1\\lesssim r/r_{500} \\lesssim 1$)\\footnote{Here and throughout this paper $r_{500}$ denotes the cluster-centric radius enclosing a mean overdensity of $500\\rho_c(z)$, where $\\rho_c(z)$ is the critical density of the universe at the redshift of analysis.}, reflecting the corresponding effect on the DM distribution. We also find that the shape of gas matches the shape of the gravitational potential at these radii in general, but deviates from it at smaller radii and at $r \\gtrsim r_{500}$ where assumption of the hydrostatic equilibrium breaks down. At smaller radii ($r\\lesssim 0.1r_{500}$) gas distributions in simulations with cooling become oblate, reflecting the partially rotation-supported thick disks into which gas settles as it loses its thermal energy by cooling, a result that is qualitatively consistent with \\citet{fang_etal09}. We predict gas shapes as may be determined observationally by estimating ellipticities from mock X-ray maps of the same clusters and show that one can constrain cluster gas physics by comparing ellipticity profiles of simulations with and without dissipation to those of observations. This paper is organized as follows. In Section~\\ref{sec:simulations} we describe our cluster simulations. In Section~\\ref{sec:3d_results} we describe the method of estimating axis ratios and the results for the three dimensional shapes of clusters. We give our estimates for ellipticities of mock X-ray maps as an observable prediction in Section~\\ref{sec:mock_xray}. We provide a summary and discussion of our results in Section~\\ref{sec:discussions}. ", "conclusions": "\\label{sec:discussions} We have investigated the axis ratios of DM mass, hot gas, and gravitational potential using 16 clusters simulated from the same initial conditions, but with different baryonic physics (with and without radiative cooling and star formation). Our results can be summarized as follows. \\begin{itemize} \\item We show that gas distribution in simulated clusters has a rather spherical shape at large cluster-centric radii with the average axis ratios of $\\sim 0.8$ at $r\\gtrsim 0.5r_{500}$. This implies that the standard assumption of spherical symmetry in analyses of ICM using X-ray and Sunyaev-Zeldovich observations should be quite accurate. \\item We show that baryonic dissipation makes gas distribution more spherical at $0.1\\lesssim r/r_{500}\\lesssim 1$, where the axis ratios are larger by $\\sim 0.2$ on average in the cooling (CSF) runs compared to the non-radiative (NR) runs. At small radii $r/r_{500} \\lesssim 0.1$, the short-to-long axis ratios $c/a$ are lower in the CSF runs compared to the NR runs by as much as $\\sim -0.5$, but the intermediate axis ratios $b/a$ are reduced to a much smaller degree. The CSF gas distributions are oblate in their centers reflecting the presence of cool gas supported by rotation. \\item We present predictions for X-ray ellipticity profiles of intracluster gas based on mock {\\em Chandra} X-ray maps of our simulated clusters. We show that the effect of cooling on the ellipticity of gas is consistent with the three-dimensional results. Specifically, the ellipticities at larger radii ($0.1\\lesssim r/r_{500c}\\lesssim 1$) decrease for the CSF clusters compared to NR clusters reflecting the rounder gravitational potential of the CSF clusters. At $r \\lesssim 0.1r_{500}$, ellipticities in the CSF clusters increase with decreasing radius. \\item NR clusters exhibit the opposite trends. NR clusters are more triaxial than the CSF clusters at $r\\gtrsim 0.1r_{500}$, but become considerably rounder at $r \\lesssim 0.1r_{500}$. The latter trend is due to random gas motions, which are present at all radii but become considerable compared to the thermal energy of the gas (as reflected by increasing $\\sigma_{\\rm gas}/v_{\\rm circ}$ ratio). \\item Our results indicate that observed ellipticity profiles of X-ray clusters can be used to constrain both the amount of cooling in the last several billion years of cluster evolution and presence of significant random gas motions. \\end{itemize} There are several issues to bear in mind when interpreting our results. First is the issue of overcooling in galaxy formation simulations. Our CSF simulations suffer from overcooling in the cluster cores ($r \\lesssim 0.1r_{500}$), such that the effect of halo contraction in response to the formation of the central galaxy is likely overestimated. The dissipational effect on gas shapes that we present can therefore be considered as an upper limit. At the same time, some amount of cooling must have occurred because galaxies do exist in clusters and the observed ICM gas fractions are significantly below the expected values \\citep[e.g.,][]{kravtsov_etal09}. We therefore expect the ellipticity of the intracluster gas in real clusters to be between our results for CSF and NR runs. Although our cluster sample is not drawn to sample the mass function of clusters, the mass-dependence of ellipticities is very weak (Figure~\\ref{fig:sm_z0}), at least in the range of masses we probe ($3.5\\times10^{13}\\,h^{-1}M_{\\odot}r_{500}$. The difference in shape between gas and potential can be attributed to deviation from hydrostatic equilibrium as shown in Section~\\ref{subsec:gas_vs_pot}. In the intermediate radial range $0.1\\leq r/r_{500}\\leq 1$ , however, we have shown that the gas shape generally coincides with that of the potential, and therefore the shape of the potential can be inferred from the shape of gas. If there is an independent way of determining the shape of gravitational potential, e.g. by gravitational lensing, one may be able to constrain the amount of gas motions by comparing the shape of the gravitational potential and the shape of gas. However, even without independent information about the potential, our results indicate that ellipticities derived from X-ray images alone can constrain the amount of cooling and the presence of random gas motions. This is because these effects result in a rapid change of ellipticity with decreasing radius at $r<0.1r_{500}$ but of an opposite sign. We will present comparisons of gas ellipticities in simulations and observations from {\\em Chandra} and {\\em ROSAT} in a separate companion paper ({E.} {T.} Lau et al. 2011, in preparation)." }, "1003/1003.5516_arXiv.txt": { "abstract": "In this paper we take the reported measurements of black hole spin for black hole X-ray binaries, and compare them against measurements of jet power and speed across all accretion states in these systems. We find no evidence for any correlation between the properties of the jets and the reported spin measurements. These constraints are strongest in the hard X-ray state, which is associated with a continuous powerful jet. We are led to conclude that one or more of the following is correct: (i) the calculated jet power and speed measurements are wrong, (ii) the reported spin measurements are wrong, (iii) there is no strong dependence of the jet properties on black hole spin. In addition to this lack of observational evidence for a relation between black hole spin and jet properties in stellar mass black holes, we highlight the fact that there appear to be at least three different ways in which the jet power and/or radiative efficiency from a black hole X-ray binary may vary, two of which are certainly independent of spin because they occur in the same source on relatively short timescales, and the third which does not correlate with any reported measurements of black hole spin. We briefly discuss how these findings may impact upon interpretations of populations of active galactic nuclei in the context of black hole spin and merger history. ", "introduction": "\\begin{figure} \\centerline{\\epsfig{file=histomontage.eps, width=8cm, angle=0}} \\caption{A compilation of reported spin measurements for black holes. The first panel shows measurements for black hole X-ray binaries (BH XRBs) based on disc continuum fitting; the three purple measurements are the three different spins reported for the system GRS 1915+105. The second panel shows measurements for BH XRBs based on disc reflection (including iron line) fitting. The third panel shows measurements for AGN, all of which are based upon reflection fits. The fourth and final panel shows the sum total of the three previous histograms.} \\label{spincomp} \\end{figure} Black holes remain one of the most bizarre and intriguing aspects of astrophysics. In general relativity a black hole is entirely described by only three parameters, mass, spin and charge. Mass is the easiest of these parameters to measure (most accurately by observing the orbits of other bodies around the black hole, such as a binary companion, nearby stars, masers etc), and charge is generally supposed to be unimportant, with astrophysical source electrically neutral on average on macroscopic scales. Black hole spin is not just a curiosity; a spinning black hole has a smaller event horizon than a non-rotating hole, and consequently a deeper gravitational well outside of the horizon, potentially increasing the efficiency of accretion. In addition, the rotational energy of spinning black holes may be enormous ($\\sim 30$\\% $Mc^2$ for a maximally spinning black hole), and could potentially be tapped as an energy source (Penrose 1969; Christodolou 1970). This concept was placed into the framework of accretion by Blandford \\& Znajek (1977) who investigated the extraction of black hole spin by a magnetic field supported by an accretion disc, and concluded that energy and angular momentum could be extracted from the black hole in this way. This concept was extended by MacDonald \\& Thorne (1982), and more recently discussed by Livio, Ogilvie \\& Pringle (1999) who concluded that the likely extraction of rotational energy of the black hole had been overestimated. McKinney (2005) however arrived at the opposite conclusion, deriving a very strong dependence of jet power on black hole spin (see also De Villiers et al. 2005 for parallel work on the influence of spin on jets from numerical simulations). The most frequent discussion of black hole spin is in the context of the apparent radio loud:radio quiet `dichotomy' in active galactic nuclei (e.g. Sramek \\& Weedman 1980; Stocke et al. 1992; Miller, Rawlings \\& Saunders 1993; Xu, Livio \\& Baum 1999), which may have an origin in the powering of AGN jets by black hole spin (e.g. Rees et al. 1982; Wilson \\& Colbert 1995; Sikora, Stawarz \\& Lasota 2007), and may tell us about the merger history of galaxies (e.g. Volonteri, Sikora \\& Lasota 2007). In recent years it has become clear that many aspects of black hole accretion and jet formation are directly comparable between AGN and lower-mass (typically $\\sim 10M_{\\odot}$) black holes in X-ray binary systems (XRBs). This is to be expected, given the very simple scalings with mass for black holes in general relativity, although there is likely to be a larger diversity of environments in AGN. Scalings between mass, radio luminosity and X-ray luminosity are reported in Merloni, Heinz \\& di Matteo (2003) and Falcke, K\\\"ording \\& Markoff (2004; see also Koerding, Falcke \\& Corbel 2006); scaling of fast variability properties with mass and accretion rate are reported in McHardy et al. (2006) and K\\\"ording et al. (2007); more qualitative similarities between XRB and AGN accretion are noted in K\\\"ording, Jester \\& Fender (2006) and also discussed in Marscher et al. (2002) and Chatterjee et al. (2009). The temporal evolution of XRB jets, relatively rapid compared to AGN, has allowed many estimates of the power (e.g. Fender 2001; Gallo et al. 2005; K\\\"ording, Fender \\& Migliari 2006) and speed (e.g. Mirabel \\& Rodriguez 1994; Miller-Jones, Fender \\& Nakar 2006) of the jets and their connection to accretion `state' as characterized by the X-ray emission (e.g. Fender et al. 1999; Fender, Belloni \\& Gallo 2004 [hereafter FBG04]; Corbel et al. 2004). Importantly these studies have shown that the jet power of a black hole XRB, as well as the radiative efficiency of the accretion flow, can change dramatically in the same source at the same overall radiative luminosity on timescales far shorter than those associated with significantly changing mass or angular momentum. In very brief summary, in black hole XRBs the coupling of radio emission (and hence jets) to X-ray state and luminosity is as follows: at Eddington ratios (in terms of X-ray luminosity) below about 0.01, sources seem to be exclusively in the `hard' X-ray state in which the X-ray emission is dominated by a component extending to $\\sim 100$ keV, widely (but not universally) accepted to arise via thermal Comptonisation of seed photons by a hot flow / corona. In this state there is strong aperiodic variability and a steady, powerful, flat-spectrum jet. The luminosities of the two components scales roughly as $L_{\\rm radio} \\propto L_{X}^{0.6-0.7}$. At higher Eddington ratios, reached generally by transient outbursting systems, sources can switch into `softer' states in which the X-ray spectrum is dominated by a cooler ($\\sim 1$ keV) component with a near-blackbody spectrum, generally interpreted as the inner accretion disc. In this state the radio emission is either dramatically suppressed by a factor $\\geq 50$ or evolves to a fading, optically thin, state, both scenarions suggesting the `quenching' of the core jet (possibly with some remnant extended emission). In transitions from hard to soft states major radio flares, often resolved as discrete, powerful, ejection events, are commonly observed. Sources generally fade in the soft state until they are once again at a few \\% Eddington (in $L_{X}$), and then make a transition back to the hard state in which mode they fade further. The initial hard $\\rightarrow$ soft state transition is usually at a higher luminosity than the soft $\\rightarrow$ hard return branch, i.e. hysteresis when spectral hardness is compared to luminosity. Note that the same source has been observed to make both hard $\\rightarrow$ soft and soft $\\rightarrow$ hard transitions at different luminosities in different outbursts; note further that some sources e.g. Cyg X-1 never drop below the 1\\% Eddington threshold and remain 'persistent and variable'. For comprehensive reviews on these phenomena, see FBG04; Remillard \\& McClintock 2006; Done, Kubota \\& Gierlinski 2007; Fender, Homan \\& Belloni 2009; Belloni 2009. The most comprehensive compilation of X-ray data on black hole binaries is presented in Dunn et al. (2010). Note that X-ray binary systems with comparable properties but hosting a neutron star instead of a black hole (candidate) also show jets, but with a lower ratio of $L_{\\rm radio}$ to $L_{X}$ (Fender \\& Kuulkers 2001; Migliari \\& Fender 2006). \\begin{table*} \\begin{tabular}{ccccc} \\hline Source & Mass & \\multicolumn{2}{c}{Spin estimate} & Refs \\\\ & $(M_{\\odot}$ & Disc & Reflection & \\\\ \\hline M33 X-7 & $15.6 \\pm 1.5$ & $0.77 \\pm 0.05$& & 1,6,7,17 \\\\ LMC X-1 & $10.9 \\pm 1.4$ & $0.90^{+0.04}_{-0.09}$ & & 1,7,18 \\\\ LMC X-3 & $11.6 \\pm 2.1$ & $<0.8$ & & 4,7, 19 \\\\ & & $-0.03$ & & 13 \\\\ GS 2000+25 & $7.2 \\pm 1.7$ & 0.03 & & 1,13 \\\\ GS 1124-68 & $6.0 \\pm 1.5$ &-0.04 & & 1,13 \\\\ 4U 1543-47 & $9.4 \\pm 1.0$ & 0.7--0.85 & $0.3 \\pm 0.1$ & 1,2,3,7,8 \\\\ GRO J1655-40 & $6.30 \\pm 0.27$ & 0.65--0.8 & $0.98 \\pm 0.01$ & 1,2,3,7,9\\\\ & & 0.93 & & 13 \\\\ GRS 1915+105 & $14 \\pm 4$ & 0.98--1.0 & & 1,2,5,7 \\\\ & & 0--0.15 & & 10 \\\\ & & $\\sim 0.7$ & & 11 \\\\ & & 0.998 & & 13 \\\\ XTE J1550-564 & 9.7--11.6 & $<0.8$ & $0.76 \\pm 0.01$ & 1,4,7\\\\ XTE J1650-500 & $5 \\pm 2$ & & $0.79 \\pm 0.01$ & 1,7 \\\\ GX 339-4 & $\\geq 6$ & & $0.94 \\pm 0.02$ & 1,7 \\\\ SAX J1711.6-3808 & & & $0.6^{+0.2}_{-0.4}$ & 7 \\\\ XTE J1908+094 & & & $0.75 \\pm 0.09$ & 7 \\\\ Cygnus X-1 & $10 \\pm 5$ & & $0.05 \\pm 0.01$ & 1,7 \\\\ 4U 1957+11 & 3--16 & 0.8--1.0 & & 1,12 \\\\ A 0620-00 & $6.6 \\pm 0.3$ & $0.12^{+0.18}_{-0.20}$ & & 21 \\\\ \\hline MCG 6-30-15 & $(4.5 \\pm 2) \\times 10^6$ & & $0.989^+{0.009}_{-0.002}$ & 14 \\\\ SWIFT J2127.4+5654 & $\\sim 10^7$ & & $0.6 \\pm 0.2$ & 15 \\\\ Fairall 9 & $(2.6 \\pm 0.6) \\times 10^8$ & & $0.60 \\pm 0.07$ & 16 \\\\ 1H 0707-495 & $\\sim 10^7$ & & $\\geq 0.98$ & 20 \\\\ \\hline \\end{tabular} \\caption{A compilation of published spin (and mass) measurements for black holes in both X-ray binary systems and AGN, based on disc and reflection/line measurements. All of these measurements, except those of Zhang et al. (1997; see text for discussion) and the two upper limits, are presented in Fig 1. Ref 1 = Remillard \\& McClintock (2006) and McClintock \\& Remillard (2009), Ref 2 = McClintock, Narayan \\& Shafee (2007), Ref 3 = Shafee et al. (2006), Ref 4 = Davis, Done \\& Blaes (2006), Ref 5 = McClintock et al. (2006), Ref 6 = Liu et al. (2008), Ref 7 = Miller et al. (2009), Ref 8 = Gallo, Fender \\& Pooley (2003), Ref 9 = fender, homan \\& belloni and references therein, Ref 10 = Kato (2004), Ref 11 = Middleton et al. (2006), Ref 12 = Nowak et al. (2008), Ref 13 = Zhang et al. (1997), Ref 14 = Brenneman \\& Reynolds (2006), Ref 15 = Miniutti et al. (2009), Ref 16 = Schmoll et al. (2009, Ref 17 = Orosz et al. (2007), Ref 18 = Orosz et al. (2009), Ref 19 = Val-Baker, Norton \\& Negueruela (2007), Ref 20 = Fabian et al. (2009), Ref 21 = Gou et al. 2010 and references therein.} \\label{spintable} \\end{table*} \\begin{figure*} \\centerline{\\epsfig{file=lrlx-full-key.eps, width=17cm, angle=0}} \\caption{The radio:X-ray plane for low/hard state black hole X-ray binaries. All currently available data are plotted, illustrating both the overall correlation over more than $10^8$ in X-ray luminosity, and also the increasing number of 'radio quiet' sources being found at relatively high X-ray luminosities. The first nine sources in the key, indicated in bold, have reported spin measurements (see Table 1). For those sources we have fitted a function with same slope as the ensemble (+0.6) but with variable normalization. In turn, we have used this normalisation as a measure of the relative jet power of the source, and compare it later to the reported spin measurements.} \\label{radio} \\end{figure*} \\begin{figure*} \\centerline{\\epsfig{file=lirlx-full-key.eps, width=17cm, angle=0}} \\caption{As Fig\\ref{radio} but for near-IR data. See Russell et al. (2006), Homan et al. (2005); Migliari et al. (2007); Russell et al. (2007b); for observational details. Note that there appear to be two tracks for the same source, XTE J1550-564 (solid blue inverted triangles), corresponding to the rise and decay phases of an outburst. This is an example of changing jet efficiency in the same source, same luminosity, same state, which is clearly not a spin effect. The two tracks for XTE J1550-564 are fitted separately in our analysis.} \\label{IR} \\end{figure*} In parallel with these advances in the study of black hole jet power, speed and relation to accretion state, there has been a rapid recent growth in the number of estimates of spin of black holes in XRB systems (the spin is generally discussed in terms of the dimensionless spin parameter $a_* = cJ / GM^2$ which has a range from 0 [non-rotating, or `Schwarzschild' black hole] to 1 [maximally rotating, or `Extreme Kerr' black hole]). Two approaches have been taken based on detailed fitting of X-ray spectra. In the first approach, the accretion disc continuum is modelled; in the second approach the `reflection' component, including the iron line around 6.4 keV, is also modelled. One of the earliest attempts to measure spin from accretion disc continua was made by Zhang, Cui \\& Chen (1997), who reported that the two superluminal jet sources, GRS 1915+105 and GRO J1655-40 had high ($a_* > 0.9$) spin whereas three other X-ray transients had much lower ($|a_*| < 0.05$), and possibly in some cases retrograde (compared to the inner accretion disc) spins. More recent disc-modelling results have been presented by, amongst others, Shafee et al. (2006), McClintock et al. (2006), Davis, Done \\& Blaes (2006), Middleton et al. (2006), Nowak et al. (2008), Steiner et al. (2009), Gou et al. (2010). Recent results from modelling of the reflection component are compiled in Miller et al. (2009). General points to take from the presented results are reports of very high spins for some black holes (e.g. 0.98--1.00 for GRS 1915+105 from disc measurements in McClintock et al. 2006; $0.98 \\pm 0.01$ for GRO J1655-40 from reflection components in Miller et al. 2009), some discrepancies (see discussion in Miller et al. 2009 and our table 1), and low spin measurements for both Cygnus X-1 ($0.05 \\pm 0.01$; Miller et al. 2009) and A 0620-00 ($0.12^{+0.18}_{-0.20}$; Gou et al. 2010). Several criticisms of the spin-fitting methods have appeared in the literature (e.g. Kolehmainen \\& Done 2010; Done \\& Diaz-Trigo 2010). In the context of disc-fitting, we further note that Fragile (2010) has reported that fits to a system where the black hole spin and inner accretion disc axes are misaligned by only 15$^{\\circ}$ are enough to render essentially useless inferred measurements of spin via this method. \\begin{table} \\begin{tabular}{ccc} \\hline Source & Distance (kpc) & Refs \\\\ \\hline GRS 1915+105 & 11.0 & 1 \\\\ GX 339-4 & 8.0 & 2 \\\\ 4U 1543-47 & 7.5 & 2 \\\\ XTE J1550-564 & 5.3 & 3 \\\\ XTE J1650-500 & 2.6 & 4 \\\\ GRO J1655-40 & 3.2 & 2 \\\\ Cygnus X-1 & 2.1 & 1 \\\\ Swift J1753.5-0127 & 8.0 & 4\\\\ GRO J0422+32 & 2.5 & 2 \\\\ 1E1740.7-2942 & 8.5 & 1 \\\\ A 0620-00 & 1.2 & 2 \\\\ GRS 1758-258 & 8.5 & 1 \\\\ GS 1354-64 & 25.0 & 5\\\\ XTE J1118+480 & 1.7 & 2 \\\\ XTE J1720-318 & 6.5 & 6 \\\\ V404 Cyg & 2.4 & 7\\\\ H1743-322 & 7.5 & 8\\\\ \\hline \\end{tabular} \\caption{Source distances adopted in this paper. Ref 1 = Gallo, Fender \\& Pooley (2003) and references therein; Ref 2 = Russell et al. (2006) and references therein; Ref 3 = Hannikanen et al. (2009); Ref 3 = Homan et al. (2006); Ref 4 = Zurita et al. (2008); Ref 5 = Casares et al. (2009); Ref 6 = Chaty \\& Bessolaz (2006); Ref 7 = Miller-Jones et al. (2009); Jonker et al. (2010).} \\end{table} In this paper we take these reported measurements of black hole spin and compare them against different methods of estimating the power and, in some cases, speed of the jet observed in such systems. From these comparisons we will draw conclusions about evidence for the dependence of jet power, or speed, on spin, in accreting black holes. Note that in this paper we are not considering estimates of black hole spin based upon other methods, such as frequencies of quasi-periodic oscillations. ", "conclusions": "So, is there any case for the reported black hole spins being correlated with jet power or jet velocity in black hole X-ray binaries ? Almost certainly no. Our view on the evidence is summarised in Table \\ref{summ}. This leads us to conclude that either: \\begin{enumerate} \\item{One or more of the methods used for estimating jet power or velocity are in error} \\item{One or more of the methods used for estimating black hole spin are in error} \\item{Jet power and/or velocity are {\\em not} related to black hole spin} \\end{enumerate} In addition to this lack of observational evidence for a relation between black hole spin and jet power or speed, we have highlighted the fact that there appear to be at least three different ways in which the jet power and/or radiative efficiency -- both of which in the context of AGN are used as estimators of spin -- from a black hole X-ray binary may vary. Two of which are certainly independent of spin because they occur in the same source on relatively short timescales, and the third which does not correlate with any reported measurements of black hole spin. This paper is not setting out to argue that black hole spin does not, in some cases, affect the power or speed of jets formed by that black hole. However, current estimates of all three parameters (spin, jet power, jet speed) of black hole X-ray binaries show no evidence for a strong relation between them. Furthermore, it is suggested that as well as pursuing the spin--jet connection, researchers working on AGN populations should consider more carefully the fact that observations of black hole binaries suggest there may be parameters other than spin which determine the radio loudness of a system." }, "1003/1003.0105_arXiv.txt": { "abstract": "We present a constructive numerical example of fast magnetic reconnection in a three dimensional periodic box. Reconnection is initiated by a strong, localized perturbation to the field lines. The solution is intrinsically three dimensional, and its gross properties do not depend on the details of the simulations. $\\sim 50\\%$ of the magnetic energy is released in an event which lasts about one Alfven time, but only after a delay during which the field lines evolve into a critical configuration. We present a physical picture of the process. The reconnection regions are dynamical and mutually interacting. In the comoving frame of these regions, reconnection occurs through an X-like point, analogous to Petschek reconnection. The dynamics appear to be driven by global flows, not local processes. ", "introduction": "Most of the matter in the universe exists in the plasma state and plasma also plays an important role in gas dynamics in astrophysics. When magnetic field are present, the dynamics of an electrically conducting plasma is sensitive to magnetic forces; as a result, magnetohydrodynamics(MHD) is used to understand the dynamical evolution of astrophysical fluids. The ideal limit of MHD poses a new class of problems in dissipative processes. In ideal hydrodynamics, irreversible processes, such as shock waves and vorticity reconnection, occur at dynamical speeds, independent of microscopic viscosity parameters. Weak solutions describe these irreversible discontinuous solutions of the Euler equations. While smooth flows conserve entropy and vorticity, the infinitesimal discontinuity surfaces generate entropy and reconnect vorticity. This can also be understood as a limiting case starting with finite viscosity, where these surfaces have a finite width. With magnetic fields, a more dramatic problem emerges. If two opposing field lines sit nearby, a state of higher entropy can be reached by reconnecting the field lines, and converting their magnetic energy into fluid entropy. In the presence of resistivity, this process occurs on a resistive time scale for some relevant scale. This exaggerates the problem somewhat. Extensive theoretical research on magnetic reconnection(\\cite{2000mrp..book.....B}, \\cite{2000mare.book.....P}) has shown that scales intermediate between the size of a system and resistive scales can be important. Nevertheless, in many astrophysical settings, simple models for reconnection give time scales that are very long, and reconnection is observed or inferred to occur on much shorter time scales, e.g. for solar flares, more than $10^{10}$ times faster than the theory\\cite{1991JGR....96.9399D}. This has led to the suggestion that magnetic reconnection in the limit of vanishing resistivity might also go to a weak (discontinuous) solution, occuring at a finite speed which is insensitive to the value of the resistivity. The problem is best illustrated by the Sweet-Parker configuration (\\cite{1958IAUS....6..123S}, \\cite{1957JGR....62..509P}), where opposing magnetic fields interact in a thin current sheet, the reconnection layer. This unmagnetized layer becomes a barrier to further reconnection. In a finite reconnection region, fluid can escape the reconnection region at alfvenic speeds. Because the reconnection region is thin, the reconnection speed is reduced from the alfven speed by a factor of the ratio of the current sheet width to the transverse system size. In the Sweet-Park model this factor is the inverse of the square root of the Lundquist number ($V_A L/\\eta$). The predicted sheet widths are typically extremely thin. Petschek proposed a fast magnetic reconnection solution (\\cite{1964NASSP..50..425P}) based on the idea that magnetic reconnection happens in a much smaller diffusive region, called the X-point, instead of a thin sheet. The global structure is determined by the log of the Lundquist number, and stationary shocks allow the fluid to convert magnetic energy to entropy. However, Biskamp's simulations (\\cite{1986mrt..conf...19B}) showed that Petschek's solution is unstable when Ohmic resistivity becomes very small. In their two dimensional incompressible resistive MHD simulations, they injected and ejected plasma and magnetic flux across the boundary. They also changed the boundary condition during the simulation to eliminate the boundary current layer. However, considering the current sheet formed in their simulation, the computation domain may not be big enough. After reproducing different scaling simulations results(\\cite{1986mrt..conf...19B}, \\cite{1986JGR....91.6807L}), Priest and Forbes \\cite{1992JGR....9716757P} pointed out that it is the boundary conditions that determine what happens (including Biskamp's unstable Petscheck's simulation) and that sufficiently free boundary conditions can make fast reconnection happen. However, there is no self-consistent simulation of fast reconnection reported, except with artificially enhanced local resistivity\\cite{1989JGR....94.8805S}. To reconcile the observed fast reconnection with its absence in simulations leads to two possible resolutions: 1) ideal MHD are not the correct equations, and long range collisionless effects are required, or 2) assumptions about the reconnection regions are too restrictive. This includes the 2-dimensionality and the boundary conditions. In exploring of the first possibility, it was found that when integrating with the Hall term in the MHD equations, or using a kinetic description(\\cite{2001JGR...106.3715B}), it was possible to find fast reconnection. However, this still didn't offer any help to the collisional system, which still has fast magnetic reconnection no matter whether Hall term is present or not; and also the increase of local resistivity is not generic in astrophysical environments, which mostly has highly conducting fluids. For the second possibility, we note that Lazarian \\& Vishniac (LV99) \\cite{1999ApJ...517..700L} proposed a model of fast magnetic reconnection with low amplitude turbulence. Subsequent simulation results \\cite{2009arXiv0903.2052K} support this model. They found that the reconnection rate depends on the amplitude of the fluctuations and the injection scale, and that Ohmic resistivity and anomalous resistivity do not affect the reconnection rate. The result that only the characteristics of turbulence determine the reconnection speed provides a good fit for reconnection in astrophysical systems. LV99 offered a solution to fast magnetic reconnection in collisional systems with turbulence. In this paper, we consider a different problem, whether we could still have fast reconnection without turbulence. We present an example of fast magnetic reconnection in ideal three dimensional MHD simulation in the absence of turbulence. Here we explore a different aspect: 3-D effects and boundary conditions. Traditionally, simulations have searched for stationary 2-D solutions, or scaling solutions. In the case of fast reconnection, the geometry changes on an alfvenic time, so these assumptions might not be applicable. Specifically, we bypass the choice of boundary condition by using a periodic box. The primary constructive fast reconnection solution, the Petscheck solution, has some peculiar aspects. The global geometry of the flow, and the reconnection speed, depend on the details of a microscopic X-point. This X-point actually interacts infinitesimal matter and energy, so it seems rather surprising that this tiny volume could affect the global flow. Instead, one might worry about the global flow of the system, which dominates the energy. We will see that this is particularly important in our simulations. ", "conclusions": "To summarize, we have found a global flow pattern which reinforces X-point reconnection, and the resulting fast reconnection in turn drives the global flow pattern. The basic picture is two dimensional. We did find that a pure 2-D simulation does not show this fast reconnection. This is easy to understand, since the reconnected field loops are loaded with matter, and would require resistivity to dissipate. In 3-D, these loops are twists which are unstable to a range of instabilities, allowing the field loops to collapse. So three basic ingredients are needed: 1. A global flow which keeps the field lines outside the X-point at a large opening angle to allow the reconnected fluid to escape, and avoid the Sweet-Parker time scale. 2. The reconnection energy drives this global flow 3. A three dimensional instability allows closed (reconnected) field lines to collapse, releasing all the energy stored in the field. The problem described here has two geometric dimensionless parameters: the 2 axis ratios of the periodic box. In addition, there are a number of numerical parameters. We have varied them to study their effects. Extending the box in the Y direction (separation between reconnection regions) shuts off this instability, which might be expected: there are no global flows possible if the two interaction regions are too far separated. We found the threshold to be $Y < 1.2Z$. In the other direction, there appears to be no limit to make $Y << Z$. Increasing the size of the Z dimension does not diminish this instability. There is also a dependence on X (extend along field symmetry axis). Shortening it to one grid cell protects the topology of field loops, and reconnection is not observed in 2-D simulations. We changed different initial condition to see whether the fast reconnection is sensitive to how the initial setup is. After changing the angle of the opposite magnetic field(from beyond 90 degree to 180 degree), the strength of the rotational perturbation, and axis of the rotational perturbation, we found that the fast reconnection still appeared. The boundary condition is kept periodic and we found that the evolution of fluid dynamics of different initial conditions are similar. It can be seen that the fast reconnection happens at the two interfaces of the straight magnetic field at the same time, with a magnetic twist moving towards it on each side. They are not head-on collision on the magnetic field, but a little separated in transverse direction. This special geometry helps the magnetic reconnection happen fast, since each magnetic twist pushes the field line, it also affect the velocity field at the other side and it helps to increase the outflow speed. If we look back to Sweet-Parker's solution(\\cite{1958IAUS....6..123S}, \\cite{1957JGR....62..509P}), the main problem is that the current sheet is so thin, that even if one accelerates the outflow to Alfven speed, the mass of outflow is still small, which slows down the speed of the reconnection. Petschek's configuration\\cite{1964NASSP..50..425P} can resolve this problem with a small reconnection region and finite opening angle for the outflow. In our simulation the speed of the outflow is further increased by the feedback between the two reconnection regions. The solar flare reconnection time scale is about Alfven time scale\\cite{1991JGR....96.9399D}, which is the order of seconds to minutes. If there is only magnetic diffusivity($\\eta$) present, the diffusive time is $\\tau_{D}=L^2/\\eta$, with $L$ is the characteristic length. Taking the values from \\cite{1991JGR....96.9399D}, $L=1000km$ and $\\eta$ is $10^{-3}{m^2}s^{-1}$, $\\tau_{D}$ is $10^{15}s$. Sweet-Parker's thin current sheet proposed a reconnection time as $\\tau_{SP}=L/(V_{Ai}/R_{mi}^{1/2})$, with $R_{mi}=L\\upsilon_{Ai}/\\eta$. This makes the reconnection time about $10^5$ Alfven times. Petschek's configuration has a reconnection time as $\\tau_{P}=L/(\\alpha\\upsilon_{A})$, with $\\alpha$ is between 0.01 and 0.1 and Alfven speed $\\sim 100km/s$, and this makes the time scale as $100-1000s$. Our fast reconnection time has the order of Alfven time scale, and Alfven time $\\tau_{A}=L/\\upsilon_{A}$, which is the same order as observed time scales of $20-60s$ \\cite{1991JGR....96.9399D}. Furthermore, comparing to LV99, no turbulence is needed or added in our simulations. Our fast magnetic reconnection time scale is qualitatively similar to the energy release time scale for solar flares." }, "1003/1003.2336_arXiv.txt": { "abstract": "We present timing analyses of eight X-ray light curves and one optical/UV light curve of the nova V4743\\,Sgr (2002) taken by \\chandra\\ and \\xmm\\ on days after outburst: 50 (early hard emission phase), 180, 196, 302, 371, 526 (super soft source, SSS, phase), and 742 and 1286 (quiescent emission phase). We have studied the multifrequency nature and time evolution of the dominant peak at $\\sim 0.75$\\,mHz using the standard Lomb-Scargle method and a 2-D sine fitting method. We found a double structure of the peak and its overtone for days 180 and 196. The two frequencies were closer together on day 196, suggesting that the difference between the two peaks is gradually decreasing. For the later observations, only a single frequency can be detected, which is likely due to the exposure times being shorter than the beat period between the two peaks, especially if they are moving closer together. The observations on days 742 and 1286 are long enough to detect two frequencies with the difference found for day 196, but we confidently find only a single frequency. We found significant changes in the oscillation frequency and amplitude. We have derived blackbody temperatures from the SSS spectra, and the evolution of changes in frequency and blackbody temperature suggests that the 0.75-mHz peak was modulated by pulsations. Later, after nuclear burning had ceased, the signal stabilised at a single frequency, although the X-ray frequency differs from the optical/UV frequency obtained consistently from the OM onboard \\xmm\\ and from ground-based observations. We believe that the late frequency is the white dwarf rotation and that the ratio of spin/orbit period strongly supports that the system is an intermediate polar. ", "introduction": "\\label{intro} Cataclysmic variables are interacting binary systems, consisting of a white dwarf primary and a late main sequence star, where accretion takes place from the cool companion to the white dwarf. If the accretion rate is in the range between $(1-4)\\times 10^{-7}$\\,M$_\\odot$\\,yr$^{-1}$, steady nuclear burning can establish, thus the hydrogen content of the accreted material is fused to helium at the accretion rate upon arrival on the white dwarf surface \\cite{vandenheuvel1992}. If the accretion rate is $<10^{-7}$\\,M$_\\odot$\\,yr$^{-1}$ then the hydrogen-rich material settles on the white dwarf surface and ignites explosively in a thermonuclear runaway, after a critical amount of mass of $10^{-6}-10^{-4}$\\,M$_\\odot$ (depending on white dwarf mass) has been reached. Such outbursts are commonly known as a Classical Nova outbursts. Within hours after outburst, the white dwarf is engulfed in an envelope of optically thick material that is driven away from the white dwarf by radiation pressure. Early in the evolution, the nova is bright in optical, but the peak of the spectral energy distribution shifts to higher energies as the mass ejection rate (and thus the opacity) decreases, exposing successively hotter layers (see, e.g., \\citealt{bode2008}). A few weeks after outburst, the nova becomes bright in X-rays, and the observed X-ray spectra resemble those of the class of Super Soft Binary X-ray Sources (SSSs). The nova V4743\\,Sgr was discovered in September 2002 by \\cite{haseda2002}. The first measured envelope ejection velocities exceeded 1200\\,km\\,s$^{-1}$ \\cite{kato2002}. \\cite{nielbock2003} studied the nova at $\\lambda$1.2\\,mm with the SEST instrument and concluded that the dominant emission source is the heated dust rather than free-free emission. The spectrum of the Nova was classified as Fe II type by \\cite{morgan2003}. \\cite{petz2005} found significant absorption from Fe and N from atmosphere modelling to a SSS X-ray spectrum taken with \\chandra\\ 180 days after outburst, that had been presented earlier by \\cite{ness2003}. Ness et al. found large-amplitude fast variability with a period of $\\sim 22$ minutes with clear overtones of this signal in the periodogram. During this observation a strong decline in X-ray brightness was observed with simultaneous spectral change from continuum spectrum to emission lines. The count rate dropped from 44 counts per second to 0.6 within $\\sim 6$\\,ks and stayed low for the rest of the observation (Ness 2010, in preparation). \\cite{kang2006} presented CCD unfiltered optical photometry analyses from observations taken in 2003 and 2005. They detected a period of 6.7\\,h which was interpreted as the orbital period of the underlying binary system. Six observations taken in short succession between days 1003 and 1011 days after outburst showed fast variability with a period of $\\sim 24$ minutes. The authors attributed this signal to the beat period between the orbital and spin period of the white dwarf, where the 22 minute signal present in X-ray (\\citealt{ness2003}) was assumed to be the spin period of the central white dwarf. Sophisticated X-ray period studies were presented by \\cite{leibowitz2006}. They reanalysed the \\chandra\\ data taken on day 180, together with an \\xmm\\ observation taken on day 196. The large-amplitude variations were also found in the \\xmm\\ data set. They detected at least 6 frequencies on day 180 and at least 12 on day 196. Most of the variability was explained by a combination of oscillations at a set of discrete frequencies, and the main feature with a period of $\\sim 22$ minutes has a double peak in the periodogram. At least 5 signals were constant between the two observations. \\cite{leibowitz2006} suggested that the main feature in the periodogram at $\\sim 22$ minutes ($\\sim 0.75$\\,mHz) is related to the white dwarf spin and that the other observed frequencies are produced by non-radial white dwarf pulsations. In this paper we present timing analyses of all \\chandra\\ and \\xmm\\ observations, including those already analysed, for consistency. The main goal is to characterise the main feature at $\\sim 0.75$\\,mHz and its possible multiple structure as well as following the evolution of the oscillations. We address the question whether these variations are produced by the white dwarf rotation or by its pulsations. In Sect.~\\ref{obs} we present all analysed X-ray data sets, and in Sect.~\\ref{period_anal_tech} we describe our period analysis methods. The results are described in Sect.~\\ref{results}. This section is structured into the presentation of two different methods described in Sect.~\\ref{period_anal_tech}, the evolution of the oscillation amplitude of the main signal, studies on beat periods, and other signals within a larger frequency radius around the main signal as well as overtones. Our results are discussed in Sect.~\\ref{discussion}, and a summary with conclusions is given in Sect.~\\ref{summary}. ", "conclusions": "\\label{discussion} \\begin{figure} \\includegraphics[width=84mm]{pevol.eps} \\caption{Graphical illustration of the evolution of the measured X-ray frequencies. The top panel shows the main frequency around 0.75\\,mHz, and in the bottom panel we show the first overtone. The bullets, connected by the solid line are the results from the 2-D method. The results from the Lomb-Scargle method are included with the dashed lines for comparison.} \\label{pevol} \\end{figure} With a data set covering more than three years of evolution, we see large changes in brightness between observations. On day 180, a sharp decline was observed by \\cite{ness2003}, and brightness variations thus occur on long and short time scales. The origin of the observed emission is not the same in all observations. On day 50, the X-ray emission likely originated from shocks within the ejecta, while on days 180-526, the spectrum was dominated by SSS emission that comes from the photosphere around the white dwarf. The post-SSS phase in novae is usually dominated by optically thin X-ray emission originating from the nebular ejecta that are radiatively cooling like in V382\\,Vel (see \\citealt{ness2005}). V4743\\,Sgr has a quiescent X-ray emission source (\\citealt{ness2007a}) that could resemble those typically observed in intermediate polars, where X-ray emission originates from an accretion shock close to the white dwarf. In addition, \\cite{kang2006} detected a similar frequency of $0.7047\\pm 0.0009$\\,mHz in optical observations which is likely associated with magnetically controlled accretion columns. The authors interpretted this frequency as the beating between the spin frequency of the white dwarf and the orbital modulation. Furthermore our optical/UV detection of a frequency of $0.699 \\pm 0.015$\\,mHz during day 742 is consistent with the \\cite{kang2006} optical finding within the errors. It thus stands to reason that the origin of the main X-ray frequency is of a fundamental nature as, e.g., the spin period of the white dwarf. However, systematic inspection of the frequency evolution yields some inconsistencies with the interpretation of pure rotational modulation. We confirm the multifrequency substructure found by \\cite{leibowitz2006} for days 180 and 196. However, for day 180, we need the higher sensitivity of our 2-D sine fitting method, as the Lomb-Scargle method yields only a single peak. Similarly, Leibowitz et al. had to use a least square fitting approach in order to detect the double nature of the main peak for day 180. Since in addition, the formal measurement uncertainties from the Lomb-Scargle method are higher than for our 2-D sine fitting method, we concentrate only on the results from the 2-D sine fitting method for the discussion. \\subsection{Evolution of the main frequency} The evolution of the main frequency is illustrated in Fig.~\\ref{pevol}. In the early observations, it is split in two frequencies, as the two signals on days 180 and 196 are different at $<99.7$\\% confidence. Since they have moved closer together from day 180 to day 196, it is possible that this was a monotonic trend that continued until they merged into a single frequency. Since the exposure times on days 301, 372, and 526 were much shorter than the beat period that corresponds to the difference between the two signals on day 180, this possibility can not be tested. However, as illustrated in Fig.~\\ref{beat}, the observations on days 742 and 1286 were long enough to detect two frequencies if they were as far apart as on day 180. Unless the relative amplitudes of the two signals have changed significantly, we can confidently conclude for at least these two observations that only a single signal was present and that the double nature of the 0.75-mHz signal has disappeared. These results are also supported by the overtone studies. The detected values satisfy the characteristics of a fundamental and a first harmonic. After day 196, the frequency appears to have decreased and then increased again by day 742 (see Fig.~\\ref{pevol}). While this is formally not statistically significant in the 3-$\\sigma$ level, other nearby signals have undergone such kind of change on a statistically significant level (see Table~\\ref{tab:othersigs}). Furthermore, the decrease of the relative amplitude as illustrated in Fig.~\\ref{aevol} is highly significant. In addition to being split in two nearby frequencies in the early evolution, the main frequency is thus not a stable signal. In spite of the likely different origin of the optical emission observed by \\cite{kang2006}, they found a similar period to our main frequency, yielding $0.7047\\,\\pm\\,0.0009$\\,mHz, which is consistent with our optical/UV measurements obtained with \\xmm\\ OM on day 742. While their frequency is remarkably close to our main X-ray frequency, the statistical uncertainty ranges indicate a significant difference from all our measurements (see Table~\\ref{det_signals}), including the late observations up to 1286 days after outburst. Nevertheless, the similarity of the frequency is striking and suggests some fundamentally common underlying origin. However, if the spin period of the white dwarf were to be responsible for the modulations in X-ray and optical, then the frequencies in both bands should be identical. The fact that the optical and the X-ray frequencies differ, plus the obvious changes in frequency and amplitude, demonstrate that there is more than simply rotational modulation from the spin of the white dwarf. \\subsection{Pulsations versus Spin Modulation} \\label{disc:pulsations} \\begin{figure} \\includegraphics[width=84mm]{pulse.eps} \\caption{Graphical illustration of the connection between changes in X-ray frequency and colour temperature. Each data point represents a predicted temperature $T_i$ for each observation taken at time day$_i$ after outburst (x-axis), according to the formula given in the top, where $T_j$ is the blackbody temperature (see Sect.~\\ref{spectra}) for day$_j$ as indicated by the plot symbol (see legend), and $f_j/f_i$ is the ratio of frequencies from our measurements. The dashed grey line represents the observed colour temperatures (see Fig.~\\ref{bevol}) for comparison with the predictions (the symbols that are connected by this line are the cases $i=j$, thus the measured temperatures).} \\label{pulse} \\end{figure} While \\cite{kang2006} rule out pulsations, the new evidence from our analysis requires us to reopen the case. Kang et al. argued that the interpretation of this period as white dwarf pulsations would yield a significant slow-down with time, owing to cooling of the temperature of the white dwarf (\\citealt{somers1999}). From their estimates of changes in effective temperature, derived from observed luminosity changes, they calculated that the period reported by \\cite{ness2003} for day 180 should be 38 minutes around day 1000, and they rule out pulsations because no such period is in their data. While it is true that we are also not observing a 38 minute period in the X-ray data around day 1000, we do see significant frequency changes in Fig.~\\ref{pevol}. \\cite{kang2006} suggested using equation 10 in \\cite{kjeldsen1995} for the conversion of temperature changes to expected pulsation frequency changes. In Sect. 4.2 of \\cite{kjeldsen1995}, a scaling relation between stellar parameters and pulsation frequencies is derived. It is argued that the acoustic cut off frequency, $\\nu_{\\rm max}$, that defines a typical time scale of the atmosphere, scales with the maximum envelope frequency of stellar pulsations. Linear adiabatic theory is applied to derive $\\nu_{\\rm max}$ and thus the relation given in their equation 10. The ratio of two pulsation frequencies in two different atmospheres scales with the square root of the inverse ratio of the respective effective temperatures (assuming the same mass and radius). With this method, our observed changes in frequency can be converted to expected changes in photospheric temperature, if pulsations were assumed. In Table~\\ref{tevol} we list squared ratios of frequencies from our 2-D sine fitting method taken from Table~\\ref{det_signals}. We indicate the two respective observation dates with numbers (1) and (2) for which the ratios were taken. For days 180 and 196 we indicate the high and low frequencies by superscript '+' and '-' behind the day of observation. Since the 3-$\\sigma$ uncertainties given in Table~\\ref{tevol} do not include the possibility that two unresolved frequencies are present, we used the larger uncertainty range resulting from half of the inverse exposure times for days 302, 371, and 526. The inverse exposure time sets a strict limit on the detectability of two separate frequencies (see Sect.~\\ref{sect:beat} and Fig.~\\ref{beat}), and if the signal was split in two components on days 302, 371, and 526, then a value somewhere in between would result (depending on the individual amplitudes), while the statistical uncertainty would not include the upper or lower frequency. As an estimate of the evolution of the effective temperature, we use colour temperatures derived in Sect.~\\ref{spectra}. As pointed out earlier, the blackbody temperature is not equivalent to the photospheric temperature, and also the evolution might not follow the same trend. However, we consider it a strong possibility that changes in the blackbody temperature indicate equivalent changes in temperature and assume this as a working hypothesis. The following conclusions depend on whether or not this assumption is valid. In the last three columns of Table~\\ref{tevol} we quote the colour temperatures derived in Sect.~\\ref{spectra} for the respective days and the inverse ratios of these values for direct comparison with the squared frequency ratios in the third column. The extent of agreement between these ratios is an indicator for the validity of the underlying assumption of pulsations. Those ratios which agree within the given uncertainty ranges of the frequency measurement are indicated by underlined numbers. For illustration we show in Fig.~\\ref{pulse} the evolution of the observed colour temperature in comparison with temperatures computed from \\begin{equation} T_i=T_{j}*(f_j/f_i)^2\\,, \\end{equation} where $T_i$ and $f_i$ are predicted temperature and observed frequency for day$_i$, respectively, and $T_j$ and $f_j$ are the observed temperature and frequency for day$_j$. The best agreement between observed and predicted effective temperature is seen for the observations on days 180 and 302, and between days 196 and 371. All predictions derived from day 526 are far off the mark and are not included in the plot. This can be seen from the predictions from other observations for day 526, which are too high, indicating that the frequency changes do not predict the decline in temperature. The agreement between predicted and observed temperature between days 196 and 302 and between days 180 and 196 is poorer. In light of the underlying assumptions made by \\cite{kjeldsen1995}, a main sequence (MS) stars with an effective temperature of order 5500\\,K, plus the considerable uncertainty in the effective temperature, the agreements are surprisingly good, strongly suggesting that pulsations play an important role. The breakdown of the relationship towards the end of the SSS phase can be understood as the turn off of the central energy source, leading to a highly non-equilibrium situation, and pulsations may not propagate the same way as in equilibrium. We thus argue that the possibility of pulsations can at least not easily be ruled out. It must be noted that for a frequency of $\\sim 0.75$\\,mHz and a mass of order 0.5-1.4\\,M$_\\odot$, equation 10 of \\cite{kjeldsen1995} yields effective radii of order 2-7\\,R$_\\odot$, which appears rather large for a hot WD and is about an order of magnitude larger than the radii derived in Sect.~\\ref{spectra} from blackbody fits, which have been discussed to be unrealistically large. If this apparent discrepancy could be resolved by a larger scaling factor, then the relative changes of effective temperature and oscillation changes would still hold for a hot white dwarf. Without computing new models for hot white dwarfs, the validity of the relation must be taken with care. However, if we assume this relation to be valid, then we have good reason to conclude that pulsations are occurring between days 180 and 371.\\\\ A fact that complicates the interpretation is the similarity between the frequencies measured during the SSS phase and during the last phase, long after the nova has turned off. This similarity suggests that a fundamentally similar origin is accountable, which brings us back to the spin period of the white dwarf. While pulsations seem to be a valid explanation for the SSS phase, the spin period of the white dwarf should also play a major role. We have no explanation but hesitate to believe in a coincidence. Perhaps magnetic fields that are present in an intermediate polar might stir up the ejecta during the SSS phase and stimulate pulsations. However, this is pure speculation. \\cite{leibowitz2006} interpreted the 0.75\\,mHz structure as due to the rotation of the white dwarf and other close signals as the nonradial pulsations of the central star. This explanation would be consistent with a similar interpretation by \\cite{drake2003} and \\cite{rohrbach2009} for the nova V1494\\,Aql. A dominant short-term oscillation with a period of 2498.9\\,s with other close signals was detected. The authors interpreted the variability as the pulsations of the central white dwarf. These central acretors after the nova explosion resemble planetary nebula nuclei where pulsations were observed in the range of $\\sim 1000$\\,s to $\\sim 5000$\\,s (\\citealt{ciardullo1996}). For example, the white dwarf in the planetary nebula NGC\\,246 had a principal frequency of 0.67\\,mHz in September 1989 and 0.54\\,mHz in June 1990 (\\citealt{ciardullo1996}). This gives a change of 0.13\\,mHz within 9 months, i.e., $\\sim 0.01$\\,mHz per month. This is remarkably similar to our finding of changes in each of the two frequencies detected on days 180 and 196 (Table~\\ref{det_signals}), which change by the amount of $\\sim 0.01$\\,mHz within two weeks. Such frequency changes are thus not unprecedented in pulsating white dwarfs. \\cite{sastri1973} studied instabilities in hydrogen burning nova shells and found pulsations as a possible phenomenon. However, while the multiperiodic and unstable behaviour is consistent with our results, the predicted typical periods are much smaller. For example, the unstable 35\\,s X-ray oscillations in RS\\,Oph (\\citealt{beardmore2008}) shortly after the start of the super soft phase would fall into this range. Another mechanism with periods in accordance to our finding could be pulsations in isolated hot GW\\,Vir white dwarfs. The origin of these pulsations is the compression-induced opacity and ionisation increase in the partial ionisation zones of carbon and oxygen (\\citealt{starrfield1984}). Such nonradial pulsations are typical for hydrogen-deficient white dwarfs with carbon-oxygen envelopes. While in accreting white dwarfs like V4743\\,Sgr the white dwarf is not hydrogen poor on its surface, \\cite{dreizler1996} found some GW\\,Vir stars that have pulsations and hydrogen on their surface. The X-ray spectra of V4743\\,Sgr contain relatively deep carbon and oxygen absorption lines (\\citealt{ness2003}), which is an indication for V4743\\,Sgr being a CO-type nova as opposed to an ONe-type nova. \\cite{drake2003} proposed the same interpretation in the case of the nova V1494\\,Aql, but without any spectral indication of the white dwarf to be the required CO type. \\begin{table} \\caption{Comparison of squared X-ray frequency ratios (frequencies from 2-D method from Table~\\ref{det_signals}) with inverse ratios of colour temperature ($T_{\\rm col}$ in $10^3$\\,K, derived in Sect.~\\ref{spectra}. For days 180 and 196, the low and high frequencies are marked with - and + superscripts, respectively. The errors include systematic uncertainties from the exposure times (see text for details). The underlined numbers indicate cases for which the relation $(f_1/f_2)^2\\simeq T_{\\rm col}^{(2)}/T_{\\rm col}^{(1)}$ holds. Note that mismatches involve only days 196 and 526, times of non-equilibrium situations (see Sect.~\\ref{disc:pulsations}).} \\begin{center} \\begin{tabular}{lccccc} \\hline \\hline Day(1) & Day(2) & $(f_1/f_2)^2$ & $T_{\\rm col}^{(1)}$ & $T_{\\rm col}^{(2)}$ & $T_{\\rm col}^{(2)}/T_{\\rm col}^{(1)}$\\\\ \\hline 180$^-$ & 196$^-$ & \\underline{0.96}$\\,\\pm\\,0.01$ & 427 & 406 & \\underline{0.95}\\\\ 180$^+$ & 196$^+$ & $1.03\\,\\pm\\,0.01$ & 427 & 406 & 0.95\\\\ 196$^-$ & 302 & $0.97\\,\\pm\\,0.11$ & 406 & 451 & 1.11\\\\ 196$^+$ & 302 & \\underline{1.05}$\\,\\pm\\,0.11$ & 406 & 451 & \\underline{1.11}\\\\ 302 & 371 & \\underline{1.06}$\\,\\pm\\,0.16$ & 451 & 443 & \\underline{0.98}\\\\ 371 & 526 & $0.99\\,\\pm\\,0.17$ & 443 & 336 & 0.76\\\\ 180$^+$ & 302 & \\underline{1.08}$\\,\\pm\\,0.11$ & 427 & 451 & \\underline{1.06}\\\\ 180$^+$ & 371 & \\underline{1.15}$\\,\\pm\\,0.12$ & 427 & 443 & \\underline{1.04}\\\\ 180$^+$ & 526 & $1.14\\,\\pm\\,0.13$ & 427 & 336 & 0.79\\\\ 302 & 526 & $1.05\\,\\pm\\,0.17$ & 451 & 336 & 0.74\\\\ \\hline \\end{tabular} \\end{center} \\label{tevol} \\end{table} \\subsection{Connection between spin and orbital period} \\cite{leibowitz2006} and \\cite{kang2006} argued that the two nearby frequencies can arise from the white dwarf spin period on the one hand and the beat period between the spin and the orbital period of the binary system on the other hand. The orbital period was reliably determined by \\cite{kang2006} as $6.718\\,\\pm\\,0.005$ hours from both observing campaigns. This can be compared to the beat periods between various pairs of observed frequencies. In Table~\\ref{periods} we list the detected periods in minutes, the difference between their frequencies, and the derived beat periods. We calculated the beat periods for the two signals in the X-ray light curves taken on days 180 and 196 (top two rows), for all possible combinations between these two signals and the optical period of 23.651\\,min (derived from the $f_1=0.7047$\\,mHz frequency found by \\cite{kang2006}; following four rows), and the combinations of periods observed in the late X-ray light curves and the optical period (bottom two rows). The beat periods between the two periods detected on days 180 and 196 are clearly inconsistent with the orbital period. Also, the beat periods between these frequencies and the optical frequency are not consistent with the orbital period. Only the beat periods between the frequencies derived from the late X-ray observations and the optical frequency are in agreement with the orbital period. The latter agreement suggests that either the optical or the late X-ray period is the spin period, while the respective other one is the beat between orbital and spin period. The most probable situation is that the X-ray signal is the spin of the white dwarf. The fact that none of the earlier frequencies are connected to the orbit-spin relation, lends additional support to our interpretation that the oscillations during the SSS phase are more than only the spin period of the white dwarf. \\begin{table} \\caption{Beat periods $P_{\\rm beat}$ (last column in unit hours), derived from various pairs of observed periods '$P_1$' and '$P_2$' (in unit minutes) and their corresponding frequencies, $f_1$ and $f_2$. In the top two rows, the beat periods for the two close frequencies that we found in the X-ray light curves on days 180 and 196 (see Table~\\ref{det_signals}) are given. Next, the beat period between each X-ray period and optical period of 23.651 minutes, observed between days 1003 and 1011 by Kang et al. (2006), are given. The last two rows show the beat period between the late X-ray periods and the optical period, which compare well to the orbital period of 6.72 hours.} \\begin{center} \\begin{tabular}{lccrr} \\hline \\hline Day & $P_1$ [min] & $P_2$ [min] & $f_2$ - $f_1$ [mHz] & $P_{\\rm beat}$ [h]\\\\ \\hline 180 & $23.180$ & $21.533$ & $0.055\\,\\pm\\,0.005$ & $5.1\\,\\pm\\,0.5$\\\\ 196 & $22.707$ & $21.872$ & $0.028\\,\\pm\\,0.006$ & $9.9\\,\\pm\\,2.2$\\\\ \\hline 180 & $23.180$ & $23.652^a$ & $0.069\\,\\pm\\,0.003$ & $4.0\\,\\pm\\,0.2$\\\\ 180 & $21.533$ & $23.652^a$ & $0.014\\,\\pm\\,0.004$ & $19.8\\,\\pm\\,6.2$\\\\ 196 & $22.707$ & $23.652^a$ & $0.057\\,\\pm\\,0.004$ & $4.9\\,\\pm\\,0.3$\\\\ 196 & $21.872$ & $23.652^a$ & $0.029\\,\\pm\\,0.005$ & $9.6\\,\\pm\\,1.7$\\\\ \\hline 742 & $22.252$ & $23.652^a$ & $0.044\\,\\pm\\,0.011$ & $6.3\\,\\pm\\,1.7$\\\\ 1286 & $22.282$ & $23.652^a$ & $0.043\\,\\pm\\,0.011$ & $6.4\\,\\pm\\,1.8$\\\\ \\hline \\end{tabular} \\end{center} $^a$Detected optical period between days 1003 and 1011 ($f_2=0.7047 \\pm 0.0009$\\,mHz) \\label{periods} \\end{table} \\subsection{The coincidence between pulsation and spin modulation} In Sect.~\\ref{disc:pulsations} we have discussed that during the SSS phase, the main period could be characterised by pulsations, while the frequency detected in the later observations is more likely the spin period of the white dwarf. In that case we wonder why the frequencies are so similar, despite of their different nature (see, e.g., Fig.~\\ref{LSpower_all_run}). We consider it unlikely that the spin period of the white dwarf has changed, yet small but significant changes in the observed oscillation frequency are undeniable. We can only speculate about how to explain the presence of a persistent signal that is variable at the same time. Perhaps, the spin period of the white dwarf could be modulated by additional processes that could depend on the physical conditions of the emitting plasma. One possibility would be that the magnetic field axis is not aligned with the rotation axis, and the spinning of the non-aligned magnetic dipole stirs up the ejecta surrounding the white dwarf. In this way pulsations could be stimulated by the spin period of the white dwarf. The flux is then modulated by these pulsations and the time evolution would depend on the properties of the ejecta. We also note that the expansion of the shell during the SSS phase that rotates at a slower velocity in the outer layers, owing to conservation of angular momentum, could slow down the oscillation frequency. We caution, however, that these ideas are unsubstantiated without theoretical support and are as such highly speculative. \\subsection{The decline on day 180} In Figs.~\\ref{aevol} and \\ref{pevol} it can be seen that the largest changes in frequency and amplitude as well as the shrinking difference between the two components of the main peak appeared shortly after the sudden brightness decline on day 180 that was accompanied by spectral changes from soft optically thick to hard optically thin (\\citealt{ness2003}). Since this decline has only been seen in the first SSS observation, thus during the early SSS phase, it may be part of an early variability phase similar to that first seen by Swift in RS Oph (e.g., \\citealt{beardmore2008}). Such an early variability phase is now routinely being observed, e.g., in V458\\,Vul (see figure 6 in \\citealt{ness2009}). Since these novae did not show any persistent oscillations as in V4743\\,Sgr, our findings are a unique contribution to the discussion about the origin of the early variations in the SSS phase. Additional evidence comes from high-resolution spectroscopy. Ness (2010, in preparation) reports the rapid disappearance of some of the nebular emission lines that were seen during the dark phase (\\citealt{ness2003}). This indicates that the surrounding nebular emission was recombining after the bright central X-ray source had disappeared. One possibility for the complete disappearance of the X-ray source could be an expansion of the photospheric radius, followed by photospheric adjustment, thus shifting the peak of the SED back into the ultraviolet. This would be consistent with the spectral changes parameterised in Fig.~\\ref{bevol}. From day 180 to day 196, a significant drop in the blackbody temperature and hardness is seen which could indicate a lower photospheric temperature on day 196. In that case, the atmospheric structure could change significantly, which could lead to changes in pulsation frequency and amplitude, as the conditions for the propagation of pulsations through the atmosphere would change. Our observations of changes in the oscillations are important evidence for a physical explanation of the variations rather than geometrical causes such as an eclipse." }, "1003/1003.4510_arXiv.txt": { "abstract": "{The primordial nature of the Spite plateau is at odds with the WMAP satellite measurements, implying a primordial Li production at least three times higher than observed. It has also been suggested that A(Li) might exhibit a positive correlation with metallicity below [Fe/H]$\\sim$-2.5. Previous samples studied comprised few stars below [Fe/H]=$-$3.} { We present VLT-UVES Li abundances of 28 Halo dwarf stars between [Fe/H]=$-$2.5 and $-$3.5, ten of which have [Fe/H]$<-$3.} {We determined stellar parameters and abundances using four different \\teff\\ scales. The direct infrared flux method was applied to infrared photometry. \\halpha\\ wings were fitted with two synthetic grids computed by means of 1D LTE atmosphere models, assuming two different self-broadening theories. A grid of \\halpha\\ profiles was finally computed by means of 3D hydrodynamical atmosphere models. The \\ion{Li}{i} doublet at 670.8 nm has been used to measure A(Li) by means of 3D hydrodynamical NLTE spectral syntheses. An analytical fit of A(Li)$_{\\rm 3D, NLTE}$ as a function of equivalent width, \\teff, log g, and [Fe/H] has been derived and is made available.} {We confirm {{previous claims}} that A(Li) does not exhibit a plateau below [Fe/H]=$-$3. We detect a strong positive correlation with [Fe/H] that is {{\\em insensitive}} to the choice of \\teff\\ estimator. From a linear fit, we infer a steep slope of about 0.30 dex in A(Li) per dex in [Fe/H], which has a significance of 2-3 $\\sigma$. The slopes derived using the four \\teff\\ estimators are consistent to within 1$\\sigma$. A significant slope is also detected in the A(Li)--\\teff\\ plane, driven mainly by the coolest stars in the sample (\\teff$<$6250), which appear to be Li-poor. However, when we remove these stars the slope detected in the A(Li) -- [Fe/H] plane is not altered significantly. When the full sample is considered, the scatter in A(Li) increases by a factor of 2 towards lower metallicities, while the plateau appears very thin above [Fe/H]=$-$2.8. At this metallicity, the plateau lies at $\\rm\\left\\langle{ A(Li)_{3D,NLTE}}\\right\\rangle=2.199\\pm0.086$.} {The meltdown of the Spite plateau below [Fe/H]$\\sim -3$ is established, but its cause is unclear. If the primordial A(Li) were that derived from standard BBN, it appears difficult to envision a single depletion phenomenon producing a thin, metallicity independent plateau above [Fe/H]=$-$2.8, and a highly scattered, metallicity dependent distribution below. That no star below [Fe/H]=$-$3 lies above the plateau suggests that they formed at plateau level and experienced subsequent depletion.} ", "introduction": "\\citet{spite82Natur,spite82} first noted that metal-poor ($-2.4$$\\le$[Fe/H]$\\le$$-1.4$), warm (5700 K $\\le$ \\teff $\\le$ 6250 K), dwarf stars exhibit a remarkably constant Li abundance, irrespective of metallicity and effective temperature, and interpreted this {\\it plateau} in Li abundance (hereafter the {\\it Spite plateau}) as being representative of the abundance of Li synthesized during the primordial hot and dense phase of the Universe (Big Bang, \\citealt{wfh}; see \\citealt{iocco08} for a review). Determining the lithium abundance in unevolved metal-poor stars has since developed into an active research topic, because of its potential role as a cosmological diagnostic. In the standard Big Bang nucleosynthesis (SBBN) scenario, \\sli\\ is formed immediately after the Big Bang, together with \\hyd, \\deu, $^{3}$He, and $^{4}$He. \\deu\\ is formed first, and is subsequently required as a seed to form any heavier element (the so-called ``deuterium bottleneck''). The abundance of all the subsequent BBN products thus depend on the equilibrium \\deu\\ abundance, which is determined by the \\deu\\ photodissociation reaction \\deu($\\gamma$,\\hyd)\\hyd. As a result, all the abundances of BBN products ultimately depend on the primordial baryon/photon ratio $\\eta_{B}\\equiv n_B / n_\\gamma$ \\citep{steigman01}, and can in principle be employed to constrain this fundamental cosmological parameter. Following the cosmic microwave background anisotropy measurements of WMAP \\citep[e.g.][]{dunkley09}, \\etab~ can be inferred from the value of the baryonic density, $\\Omega_B$, i. e. determining the primordial abundance of BBN products is no longer the only means by which it is estimated. On the other hand, the comparison between the two estimates remains of paramount importance as a test of the reliability of the BBN theory, of our present understanding of the subsequent chemical evolution of the elements involved, and, in the case of Li, of our understanding of stellar atmospheres. Among the available BBN products, \\deu\\ and \\sli\\ are the most reliable \\etab~ indicators. Being \\deu\\ never produced in stars, its abundance in a low-metallicity environment can be assumed to be quite close to the cosmological value. In addition, its sensitivity to \\etab\\ is monotonic and quite strong \\citep[ (\\deu/\\hyd)$\\propto \\eta_B^{-1.6}$, ][]{steigman09}. On the other hand, \\deu\\ can be effectively measured only in high-redshift, low-metallicity damped Lyman $\\alpha$ (DLAs) or Lyman limit systems, for which the observations are so challenging that only seven such high quality measurements exist to date, which were all obtained after 10m-class telescopes became available \\citep[][]{pettini08}. The \\etab~ value inferred from them is in good agreement with that derived from WMAP \\citep[][]{steigman09}. In contrast, \\sli\\ can be measured with relative ease in the photospheres of warm, unevolved stars. The observations are typically restricted to dwarfs, at least when one is interested in determining the primordial Li abundance, because the fragile Li nucleus is destroyed by the $^7$Li(p,$\\alpha$)$^4$He reaction as soon as the temperature reaches 2.6 million K. This implies that giants should not be considered, since their deep convective zones mix the surface material with layers that exceed this temperature, and almost all Li is rapidly destroyed. The ease with which \\sli\\ is destroyed has always constituted a challenge to existing models of convection and diffusion in stellar atmospheres, which predict a depletion of at least a factor of four relative to the primordial abundance \\citep{michaud84}. While one could infer that some depletion might have occurred, it appeared impossible to obtain a {\\em constant} depletion over such a wide range of effective temperatures. The simplest solution was to assume that no depletion was indeed taking place. This is in marked contrast to the solar case, where the photospheric Li abundance is about two dex lower than the meteoritic value. The original interpretation of the Spite plateau has been challenged in many ways in the years since its discovery. Surely the most compelling challenge was the independent measurement of \\etab~ by the WMAP satellite, placing the expected primordial Li abundance at A(Li)\\footnote{A(Li)=$\\log [N({\\rm Li}) / N({\\rm H})]+12$}$_P$=2.65$^{+0.05}_{-0.06}$ \\citep{steigman07}, or even higher, A(Li)$_P$=2.72$\\pm0.05$ when updated rates are taken into account for the $^3$He($\\alpha,\\gamma$)$^7$Li reaction \\citep[][]{cyburt08}. The highest estimate of the Spite plateau does not exceed A(Li)=2.4, a more typical value being A(Li)$\\sim$ 2.2. The discrepancy can in principle be eliminated in two ways, by either rejecting the standard BBN scenario \\citep[for a review see][]{iocco08}, or by assuming that some degree of Li depletion has occurred. Two main mechanisms could again be invoked. Li could be subject to depletion {\\em before} the currently observed stars are formed \\citep{piau06}, by means of the reprocessing of the primordial gas in a first generation of massive, hot stars. This phenomenon does not appear to be able to explain the entire WMAP ~/~ Spite plateau gap, but, removing up to 0.3 dex of the discrepancy could considerably reduce the problem. The maximum possible depletion is nevertheless dependent on the initial mass function and lifetime of Pop. III stars, as well as on the effectiveness of the mixing of their ejecta in the interstellar medium, which are all poorly known. Alternatively, Li can be depleted {\\em within} the stars we currently observe, as a consequence of phenomena within the envelope, such as diffusion, gravity waves, rotational mixing, or any combination of these. As stated above, the negligible scatter, and apparent lack of slope in the Spite plateau are observational constraints that models of Li depletion have failed to reproduce. This could apparently be achieved by combining diffusion with some form of turbulence at the bottom of the atmospheric convective zone \\citep{richard05,korn06,korn07,piau08,lind09b}. Unfortunately, the effect of turbulence is introduced basically as a free parameter, and its tuning is made quite difficult by the subtlety of the effects expected on elements other than Li \\citep[see sect. 6.4 in][]{bonifacio07}. Claims have been made \\citep[e.g.][]{asplund06} that the lighter $^6$Li isotope has been detected in the atmospheres of dwarf stars displaying Spite plateau $^7$Li abundances. These measurements are very difficult and sensitive to subtle details of the analysis \\citep[][]{cayrel07}. If the detection of $^6$Li in EMP dwarf stars were to be confirmed, it would severely undermine any claim of a substantial atmospheric depletion of $^7$Li during the star's lifetime, since the $^6$Li is even more easily destroyed than $^7$Li. One additional problem is constituted by repeated claims that the Spite plateau might display a tilt towards lower Li abundances at lower metallicities, on the order of 0.1-0.2 dex in A(Li) per dex in [Fe/H] \\citep{ryan96,ryan99,boesgaard05,asplund06}, although other studies failed to confirm this \\citep[e.g.][]{bonifacio1997}. Roughly below [Fe/H]=$-$2.5, more and more stars appear to exhibit Li abundances below the plateau level, while the scatter increases. The extreme case is possibly represented by the lithium abundance upper limit of the hyper-iron-poor subgiant \\object{HE 1327--2326} \\citep[][and references therein]{frebel08}, which should have A(Li)$\\le$0.7 (from 1D analysis). The interpretation of this result is not straightforward. Even rejecting the interpretation \\citep{venn08} that this star might be a chemically-peculiar {\\em evolved} object, the unusual photospherical composition of this star has not yet found a satisfactory explanation. Were the composition of \\object{HE 1327--2326} to be indeed primordial, its lack of Li would support the \\citet{piau06} suggestion of a pollution by material cycled through massive Pop. III stars. Adopting the \\citet{piau06} hypothesis, one could then envision a scenario in which partial pollution by this astrated material induces varying degrees of Li ``depletion'' in EMP stars according to how much this reprocessed gas is available locally at the location and time of each stars' formation. A linear fit to EMP stellar Li abundances would then naturally lead to an expected trend in A(Li) with [Fe/H], whose slope would appear steeper the more the sample is limited to low metallicities. An alternative explanation would be to postulate that a Li ``over-depletion'' mechanism operates in the photospheres of the most metal-poor stars, a mechanism that would not act uniformly in every star of a given metallicity (possibly depending on \\teff\\ or rotation speed or both), but would be more efficient at lower [Fe/H]. These stars would then begin with a Li abundance corresponding to the Spite plateau, but most of them would then develop some degree of Li depletion. This explanation would, at the same time, explain the apparent slope at low metallicities and the increase in the scatter. It would also explain why, even at very low metallicity, one still finds some stars lying on the Spite plateau. A striking example of this is the EMP double-lined binary system \\object{CS 22876--032} \\citep{gonzalez08}, in which, at [Fe/H]=$-$3.6, the primary lies on the Spite plateau, while the secondary has a Li abundance {\\em lower} by about 0.4 dex. \\begin{table} \\caption{Observations log for the 11 new targets. \\label{obslog}} \\begin{center} {\\scriptsize \\begin{tabular}{llllll} \\hline {\\bf Star} & {\\bf Obs. date} & {\\bf MJD$^a$} & {\\bf Exp. time} & {\\bf V$_{\\mathrm rad}^b$} & {\\bf S/N$^c$} \\\\ & {\\bf (UT) } & {\\bf (UT) } & {\\bf sec. } & {\\bf km/s } & \\\\ \\hline \\object{BS 17572--100} & 21 Feb 2006 & 53787.05628641 & 3190\t\t\t\t& 189 & 191 \\\\ \\\\ \\object{CS 22188--033} & 7 Jul 2006 & 53923.31274319 & 3035 & 14 & 107 \\\\ \\\\ \\object{CS 22882--027} & 6 Jul 2006 & 53922.31576391 & 4100 & 182 & 78 \\\\ & & 53922.36486180 & 4100 & 182 & \\\\ \\\\ \\object{CS 22950--173} & 23 Apr 2006 & 53848.33688352 & 3600 & 69 & 92 \\\\ \\\\ \\object{CS 29491--084} & 17 May 2006 & 53872.31418739 & 3600 & -8 & 104 \\\\ \\\\ \\object{CS 29514--007} & 7 Jul 2006 & 53923.35103355 & 3600 & 41 & 91 \\\\ \\\\ \\object{CS 29516--028} & 19 May 2006 & 53874.32992778 & 3600 & -179 & 62 \\\\ & & 53874.37328577 & 3600 & -179 & \\\\ \\\\ \\object{CS 30302--145} & 23 Apr 2006 & 53848.28216715 & 2160 & 195 & 70 \\\\ & & 53848.30823918 & 2160 & 195 & \\\\ \\\\ \\object{CS 30344--070} & 17 May 2006 & 53872.35925533 & 2160 & -141 & 82 \\\\ & 27 May 2006 & 53882.31201609 & 2160 & -140 & \\\\ \\\\ \\object{HE 0148--2611} & 12 Jul 2006 & 53928.37446515 & 3600 & -227 & 72 \\\\ \\\\ \\object{HE 1413--1954} & 21 Feb 2006 & 53787.27197427 & 3600 & -101 & 52 \\\\ \\hline \\multicolumn{6}{l}{\\scriptsize $a$ Modified Julian date of observation start: MJD=JD-2400000.5}\\\\ \\multicolumn{6}{l}{\\scriptsize $b$ Rounded to the nearest km/s, barycentric correction applied}\\\\ \\multicolumn{6}{l}{\\scriptsize $c$ Near \\ion{Li}{i} 670.8 nm doublet, if more than one spectrum has been used, S/N}\\\\ \\multicolumn{6}{l}{\\scriptsize ~~~is measured in the coadded spectrum.}\\\\ \\end{tabular}} \\end{center} \\end{table} ", "conclusions": "We have presented the largest sample to date of Li abundances for EMP halo dwarf stars (27 abundances and one upper limit), including the largest sample to date below [Fe/H] = $-$3 (10 abundances). Lithium abundance determination is highly sensitive to biases in the effective temperature scale, and we have tried to account for this using four different temperature estimators. In an additional effort to accurately represent the stellar atmospheres of the sample stars, 3D, time-dependent, hydrodynamical atmosphere models have been used to determinine our preferred \\halpha-based temperature scale, and a detailed 3D NLTE spectrosynthesis has been applied to the determination of lithium abundance. Both these techniques have been employed here for the first time, to our knowledge, in the analysis of EMP stars. This has also allowed us to develop a useful fitting formula allowing one to derive A(Li)$_{\\rm 3D, NLTE}$ directly as a function of EW, \\teff, \\glog, and [Fe/H] for EMP turn-off and early subgiant stars (see Appendix \\ref{fitappendix}). The first obvious conclusion of this work is that we have confirmed what was merely suggested by the analysis of \\citet{bonifacio07}, and previous works, that at the lowest metallicity there is sizable dispersion in the Li abundances and that there is a trend of decreasing Li abundance with decreasing metallicity. We have also shown that these two conclusions do not depend on the adopted temperature scale, as suggested by \\citet{molaro08}. The results hold, qualitatively, using both IRFM temperatures and H$\\alpha$ temperatures, regardless of the broadening theory adopted and irrespective of the use of either 1D or 3D model atmospheres. Quantitatively, the results differ in the mean level of the Li abundance, while the slopes in the A(Li) versus [Fe/H] relations agree within errors. None of the temperature scales investigated produces a ``flat'' Spite plateau over the full range in [Fe/H] (see Table \\ref{fitstable}). Our results are in substantial agreement with those of \\citet{aoki09}. While these authors do not detect a slope with either effective temperature or metallicity, this happens simply because of the small extent of their sample in both these parameters. On the other hand, they do point out that their sample has a lower Li abundance than that observed at higher metallicities. \\begin{figure} \\centering \\includegraphics[width=9cm]{13282_fi15.eps} \\caption{A unified view of A(Li) vs. [Fe/H] from some studies for which a common temperature scale can be assumed. Blue circles, \\citet{asplund06} data, red triangles, \\citet{aoki09} data, magenta squares, \\object{CS 22876--032} from \\citet{gonzalez08}, filled symbol primary star, open symbol secondary star. Black diamonds, this work, BA temperature scale. Dot-dashed gray line, best linear fit to \\citet{asplund06} data, continuous dark gray line, best fit to our data. Typical error bars for our data are displayed.} \\label{alltogethernow} \\end{figure} The picture outlined by the aforementioned results acquires more significance, once we place it in a broader context among the latest studies regarding lithium in EMP stars. Figure \\ref{alltogethernow} compares our results with those of three investigations employing compatible temperature scales. In this figure, open blue circles represent stars from the \\citet{asplund06} sample, red triangles from the \\citet{aoki09} data, and the two magenta squares the two components of the double-lined binary system \\object{CS 22876--032} \\citep[][the filled square corresponds to the primary star]{gonzalez08}. Our data are represented as black diamonds (the results of the BA scale are shown, for compatibility with the temperature scales used in the other three works)\\footnote{\\citet{gonzalez08} derived \\teff\\ from photometry and isochrones, but a cross-check with \\halpha\\ profiles computed in 1D with \\citet{barklem00} broadening confirmed the result.}. The best linear fit to our data is shown as a dark gray solid line, while the best fit to \\citet{asplund06} data (A(Li)=2.409+ 0.103[Fe/H]) is shown by a dot-dashed gray line. The \\citet{asplund06} Li abundances are increased here by 0.04 dex to account for the known offset already mentioned in Sect. \\ref{compare_others}, and their metallicty is decreased by 0.2 dex to correspond to the metallicity-scale offset detected by \\citet{bonifacio07}. It is now even more evident that the Spite plateau does not exist anymore at the lowest metallicity, and is replaced by an increased spread of abundances, apparently covering a roughly triangular region ending quite sharply at the plateau level. This region appears here to be populated in a remarkably even manner; at any probed metallicity some star remains at, or very close to, the Spite plateau level, but many do not. The rather different slopes of the best-fit relations in \\citet{asplund06} and in this work appear to be the obvious consequence of fitting two subsamples covering different metallicity regimes. This could provide also an explanation for the numerous claims, starting from \\citet{ryan99}, of a thin, but tilted Spite plateau. From this view, the difference was produced simply because the tail of these samples had been falling in the low-metallicity ``overdepletion zone'' as we have been able to discern more clearly. We are not aware of any theoretical explanation of this behavior. After the measurements of the fluctuations of the CMB made it clear that there is a ``cosmological lithium problem'', i.e., the Li predicted by SBBN and the measured baryonic density is too high with respect to the Spite plateau (by about 0.6\\, dex for our sample), there have been many theoretical attempts to provide Li-depletion mechanisms that would reduce the primordial Li to the Spite plateau value in a uniform way. Our observations now place anadditional constraint on these models -- below a metallicity of about [Fe/H] = $-$2.5, they should cause a dispersion in Li abundances and an overall lowering of A(Li). If Li depletion from the WMAP-prescribed level were to happen in the stellar envelopes of very metal-poor stars, the mechanism would have to be remarkably {\\em metallicity insensitive} to account for the thin, flat plateau observed between [Fe/H]=$-$2.5 and $-$1. And yet, the same phenomenon must become sharply {\\em metallicity sensitive} around and below [Fe/H]=$-$2.5, i.e., precisely where metallicity effects on the atmospheric structure are expected to become vanishing small. We are tempted to imagine that two different mechanisms may need to be invoked to explain the production of the Spite plateau for stars with [Fe/H] $> -2.5$, and of the low-metallicity dispersion for stars with [Fe/H] $< -2.5$. One could envision such a two-step process as follows: \\begin{enumerate} \\item Metal-poor halo stars are always formed at the Spite plateau level, regardless of their metallicity. Whether the plateau represents the cosmological Li abundance or is the result of some primordial uniform depletion taking place {\\em before} the star formation phase is immaterial in this context. \\item A second phenomenon, possibly related to atmospheric diffusion, becomes active around [Fe/H]=$-$2.5 and below, depleting Li further in the atmosphere of EMP stars. This phenomenon, aside from the metallicity sensitivity, would exhibit different star-to-star efficiency, being possibly dependent on additional parameters, such as stellar rotation or \\teff. Its efficiency must in any case be higher for more metal-poor stars. \\end{enumerate} In this scenario, the ``primordial'' plateau would be preserved above [Fe/H]$\\sim$ $-$2.5, but below that metallicity, a systematic ``leakage'' of stars towards lower A(Li) would take place, more effectively for more metal-poor stars, but naturally scattered due to the sensitivity to parameters other than [Fe/H]. This scheme would have a number of advantages. First of all, it would naturally explain our observations, ``mimicking'' a slope in A(Li) versus [Fe/H], but with increased scatter at low [Fe/H]. It would also explain why, while the scatter in A(Li) increases at low metallicities, not a single star in this metallicity regime has been found to lie above the Spite plateau level. It would then be consistent with a small number of stars remaining close to the plateau at any metallicity \\citep[e.g., \\object{CS 22876--032} A,][ filled magenta square in Fig. \\ref{alltogethernow}]{gonzalez08}; in these objects, the depletion process would be somehow inhibited. Finally, attributing the extra depletion to atmospheric diffusion / settling would not require a physical ``conspiracy'' capable of producing exactly the same depletion level regardless of metallicity, stellar rotation, gravity, or effective temperature, as is often invoked when diffusion is used to explain the Spite plateau. The nature of what we refer to above as the ``second phenomenon'', the one responsible for the departures from the Spite plateau below [Fe/H] = $-2.5$, is perhaps the most intriguing. Above, we have proposed some kind of photospheric settling mechanism, but one could as well envision a chemical evolution scenario, on the basis of some gas pre-processing with Li depletion ({\\it \\`a la}} \\citealt{piau06}) -- while it may not be able to account for the entire WMAP-Spite plateau discrepancy, this mechanism could easily account for the mild (0.2-0.4 dex) departure from the plateau observed at lower metallicities. Moreover, this mechanism would naturally produce a spread of abundances as a consequence of the local level of gas pre-processing. There are hints that the recently discovered ultra-faint dwarf galaxies (uFdg) might have been the source of the bulk of the EMP stars now found in the halo of the Milky Way \\citep[][and references therein]{tolstoy09}. If this were indeed the case, a sizeable fraction of our sample could have formed in uFdg systems, possibly more so for the most metal-poor objects. It has been suggested \\citep{komiya09} that the paucity of stars below [Fe/H] = $-3.5$ may be due to the onset of self-pollution in the primordial mini-halos when they started to merge to form larger structures, and ultimately the halo. One could then envision that gas reprocessing could have fairly significantly altered the Li abundances in the heavily dark-matter dominated cores of these sub-halos, but become progressively negligible when they merged, and their ISM mixed more and more completely with gas of pristine Li abundance. Sub-plateau stars may then originate from star formation that occurred in the sub-halo cores, or while mixing with pristine gas that progressively diluted the effect of the reprocessing. In any case, the main drawback of this two-phenomena scenario is to leave the WMAP-Spite plateau discrepancy unexplained. It also implies that any depletion from WMAP-based primordial A(Li) should have taken place before the currently observed stars formed. Non-standard primordial nucleosynthesis thus remains entirely viable, until the contrary is proven. Looking retrospectively at the history of the abundances of light elements and SBBN, one concludes that if the Spite plateau is the result of significant uniform Li depletion among metal-poor stars, one faces a formidable case of cosmic conspiracy. It is necessary to admit that the Li has been depleted quite precisely to a level consistent with SBBN production and also consistent with the abundances of the other measurable light elements. Had Li been depleted to the level of A(Li)=1.8, that is, below the minimum allowed by SBBN, the Spite plateau might not have been interpreted as being related to primordial Li, right from the beginning. A caveat is in order when looking at Fig. \\ref{alltogethernow}: the \\citet{asplund06} sample is biased, having been purposedly selected not to include objects significantly deviating from the Spite plateau. In fact, stars with varying degrees of Li depletion have long been known to exist \\citep[e.g.,][]{charbonnel05}. On the other hand, {\\em warm} (i.e., \\teff$>$ 6000 K) dwarf stars deviating from the Spite Plateau appear to be quite rare at [Fe/H]$>$-3. Again, our sample is limited towards higher metallicities, which prevents us from properly quantifying this statement. A re-evaluation of the Li abundances above [Fe/H]=-2.5 is in order, using an analysis technique homogeneous with the present one, and applied to a sample for which a well-known selection technique has been applied. It is plain to see that, had previous authors encountered Li abundance distributions as the one observed by us and \\citet{aoki09}, the concept of a lithium abundance plateau would have been short lived. Instead, there is a clear {\\em perception} that mildly Li depleted stars do exist among MP stars, but are a rare occurrence. We cannot easily identify any reason why all the previous studies should have been biased in a different way with respect to the present one, but until a coherent work, based on a well defined selection criterion, is extended to higher metallicities, all we are left with is just this perception. Two other lines of future investigation are suggested by the present results: ({\\em i}) increase the sample of {\\em hot} (\\teff $\\ge6250$\\,K)) EMP stars ([Fe/H]$\\le -3.0$), and ({\\em ii}) increase the number of cool (\\teff $<6250$\\,K) stars over the entire metallicity range. The present sample has seven {\\em hot} EMP stars, five of which have A(Li) below the Spite plateau. Even if ``standard'' ZAMS Li depletion was setting in at higher than usual temperatures in EMP stars, these objects are unlikely to be affected by it. Statistics is still weak in this metallicity regime, and that the primary component of CS~22876--032 lies at the level of the Spite plateau, at a metallicity of [Fe/H] = $-$3.6, clearly calls for a search for similar objects. At the same time, the cool end of our sample is presently too poorly populated to draw any definitive conclusion about whether these stars are experiencing convective depletion, which again, calls for an enlargement of the cool stars sample at low ([Fe/H]$<$-2.5) metallicities. To help understand the phenomena involved (since there could be several at work) that bring about the observed Li abundance pattern, another important issue is to assign an accurate evolutionary status to each star, that is, to confidently assess which stars are dwarfs, main-sequence turnoff, or subgiants. Our surface gravities are not sufficiently accurate for this purpose. The GAIA satellite will provide accurate parallaxes for all of the presently studied stars, thus allowing one to more clearly elucidate this problem." }, "1003/1003.1759_arXiv.txt": { "abstract": "As an initial investigation into the % long-term evolution of protostellar disks, we explore the conditions required to explain the large outbursts of disk accretion seen in some young stellar objects. We use one-dimensional time-dependent disk models with a phenomenological treatment % of the magnetorotational instability (MRI) and gravitational torques to follow disk evolution over long timescales. Comparison with our previous two-dimensional disk model calculations (Zhu et al. 2009b, Z2009b) indicates that the neglect of radial effects and two-dimensional disk structure in the one-dimensional case makes only modest differences in the results; this allows us to use the simpler models to explore parameter space efficiently. We find that the % mass infall rates typically estimated for low-mass protostars generally result in AU-scale disk accretion outbursts,% as predicted by our previous analysis (Zhu \\etal 2009a, Z2009a). We also confirm quasi-steady accretion behavior for high mass infall rates if the values of $\\alpha$-parameter for the magnetorotational instability is small, while at this high accretion rate convection from the thermal instability may lead to some variations. We further constrain the combinations of the $\\alpha$-parameter and the MRI critical temperature, which can reproduce observed outburst behavior. Our results suggest that dust sublimation may be connected with full activation of the MRI. This is consistent with the idea that small dust captures ions and electrons to suppress the MRI. In a later paper we will explore both long-term outburst and disk% evolution with this model, allowing for infall from protostellar envelopes with differing angular momenta. ", "introduction": "In the standard model of low-mass star formation, a molecular cloud core collapses to a protostar over timescales of $\\sim 10^{5}$~yr (e.g., Shu, Adams, \\& Lizano 1987), consistent with observations (Kenyon \\etal 1990; Enoch \\etal 2008). However, steady accretion of this mass onto central stars with a plausible mass-radius relation% results in accretion luminosities that are larger than those observed in low-mass protostars (Kenyon \\etal 1990, 1994; Enoch \\etal 2009). One solution to this ``luminosity problem'' is that most infalling matter first falls to the circumstellar disk and then is accreted to the star during short-lived outbursts; in this model protostars are usually observed in quiescence. The FU Orionis objects provide direct evidence for this type of behavior, with maximum accretion rates of $10^{-4} \\msunyr$ over periods of decades to centuries (Hartmann \\& Kenyon 1996), which also directly suggests $\\sim$10$^{-2}$ M$_{\\odot}$ in the disk at $\\sim$ 1 AU. A number of theories have been proposed to explain FU Orionis outbursts, including thermal instability in the inner disk (Bell $\\&$ Lin 1994), binary interactions \\citep{Bonnel1992}, and gravitational clumping at several AU (Vorobyov $\\&$ Basu 2005, 2006). In a recent paper (Z2009a), we explored the possibility that outbursts might result because of a mismatch between the mass fluxes that can be transported by the magnetorotational instability (MRI) in the inner disk, and the gravitational instability (GI) in the outer disk. Using steady thin disk theory, we argued that outbursts are to be expected when disks are driven by mass addition lower than 10$^{-4}\\msunyr$, as initially found by Armitage, Livio, \\& Pringle (2001). We then developed a two-dimensional model of FU Orionis disks, which verified that outbursts of accretion similar to those observed could be produced using reasonable parameters for MRI transport when thermal ionization dominates (Z2009 b). The computationally-intensive nature of two-dimensional (let alone three-dimensional) simulations of outbursts makes it difficult to conduct studies of the effects of differing parameters on disk evolution over significant timescales. We have therefore developed one-dimensional disk models to follow disk evolution. While such models have limitations, they can serve as a starting point to investigate the landscape of possible disk evolutionary pathways. The diversity of disk properties among stars of nearly the same age and mass (e.g., Hartmann 2009) is plausibly% the result of differing initial conditions, and models of the type we explore here can begin to address this possibility. In \\S 2 we describe our 1-D, two-zone model, and compare its outburst properties with our 2-D models in \\S 3. In \\S 4 we present the results of a parameter study designed to show when outbursts occur and how the outburst strength and frequency depend on the adopted parameters. We present our discussion and conclusions in \\S 5 and \\S 6. The present work, in which we assume constant mass addition to the outer disk, serves as a starting point for our subsequent investigation of long-term disk evolution with mass addition due to infall from a rotating protostellar envelope in a following paper. ", "conclusions": "In this paper, we have evolved a one-dimensional layered disk model including both MRI and GI to study the unsteady disk accretion of protostars. The 1-D models reproduce the general properties of 2-D (axisymmetric) outbursts reasonably well, such as the outbursting mass accretion rate, duration% and the accreted mass during one outburst. Because the 1-D model is faster, it% enables us to study outbursts in an extended parameter space. Our results confirm that the disk can accrete steadily with high infall rates ($\\dot{M_{in}}$$\\sim$10$^{-4}$$\\msunyr$ if $\\alpha_{M}$=0.01; \\cite{armitage01}). This steady accretion may still have short timescale variations, however, due to the thermal instability in the inner disk, as suggested by our earlier, 2-D simulations. We also confirm the prediction by Z2009a that protostars are likely to accrete unsteadily/in outbursts for infall rates $\\lesssim 10^{-5} \\msunyr$ with $\\alpha_M = 0.01$ and $\\lesssim 10^{-4} \\msunyr$ for $\\alpha_M = 0.1$. Outbursts are triggered at $r \\sim 1-10$~AU for protostellar infall rates $\\sim 10^{-5} - 10^{-6} \\msunyr$. The outbursts are stronger and shorter with larger $\\alpha_{M}$ or T$_{M}$. The total mass accreted during one outburst mainly depends on T$_{M}$. While the outbursts are slightly shorter for more massive central stars, the outburst $\\dot{M}$ is nearly independent of central star mass. The active layer surface density only affects the mass accretion rate in% the low state; it has little effect on the outburst. By comparing with % the mass accretion rate and duration of% observed FU Orionis events, we can constrain a combination of $\\alpha_{M}$ and T$_{M}$. If $\\alpha_{M}$ is low (0.01), T$_{M}$ needs to be high (1800 K, higher than the dust sublimation temperature); if $\\alpha_M$ is high (0.1) then $T_M$ needs to be low ($\\lesssim$1400 K). Our results show that 1-D, two zone models can capture the basic features of disk evolution, given our assumptions about the action of the MRI and GI. In a later paper we will address disk evolution over a much longer timescale, explicitly taking into account mass infall from a rotating protostellar cloud." }, "1003/1003.1790_arXiv.txt": { "abstract": "We extend the study of the core of the \\fekalfa emission line at $\\sim 6.4$~keV in Seyfert galaxies reported in Yaqoob \\& Padmanabhan (2004) using a larger sample observed by the \\chandra High Energy Grating (HEG). The sample consists of 82 observations of 36 unique sources with $z<0.3$. Whilst heavily obscured active galactic nuclei (AGNs) are excluded from the sample, these data offer some of the highest precision measurements of the peak energy of the \\fekalfa line, and the highest spectral resolution measurements of the width of the core of the line in unobscured and moderately obscured ($N_{H}<10^{23} \\ \\rm cm^{-2}$) Seyfert galaxies to date. From an empirical and uniform analysis, we present measurements of the \\fekalfa line centroid energy, flux, equivalent width (EW), and intrinsic width (FWHM). The \\fekalfa line is detected in 33 sources, and its centroid energy is constrained in 32 sources. In 27 sources the statistical quality of the data is good enough to yield measurements of the FWHM. We find that the distribution in the line centroid energy is strongly peaked around the value for neutral Fe, with over 80\\% of the observations giving values in the range 6.38--6.43~keV. Including statistical errors, 30 out of 32 sources ($\\sim 94\\%$) have a line centroid energy in the range 6.35--6.47~keV. The mean equivalent width, amongst the observations in which a non-zero lower limit could be measured, was $53 \\pm 3$~eV. The mean FWHM from the subsample of 27 sources was $2060 \\pm 230 \\ \\rm km \\ s^{-1}$. { The mean EW and FWHM are somewhat higher when multiple observations for a given source are averaged.} From a comparison with the H$\\beta$ optical emission-line widths (or, for one source, Br$\\alpha$), we find that there is no universal location of the \\fekalfa line-emitting region relative to the optical BLR. In general, a given source may have contributions to the \\fekalfa line flux from parsec-scale distances from the putative black hole, down to matter a factor $\\sim 2$ closer to the black hole than the BLR. We confirm the presence of the X-ray Baldwin effect, an anti-correlation between the \\fekalfa line EW and X-ray continuum luminosity. The HEG data have enabled isolation of this effect to the narrow core of the \\fekalfa line. ", "introduction": "\\label{intro} The narrow core (FWHM~$<10,000 \\ \\rm km \\ s^{-1}$) of the \\fekalfa fluorescent emission, peaking at $\\sim 6.4$~keV is a common and dominant feature of the X-ray spectrum of active galactic nuclei (AGNs) that have a 2--10~keV X-ray luminosity less than $\\sim 10^{45} \\rm \\ erg \\ s^{-1}$ (e.g. Sulentic \\etal 1998; Lubi\\'{n}ski \\& Zdziarski 2001; Weaver, Gelbord, \\& Yaqoob 2001; Perola \\etal 2002; Yaqoob \\& Padmanabhan 2004 (hereafter YP04); Levenson \\etal 2006; Winter \\etal 2009). The luminosity in the core of the \\fekalfa emission may be comparable to any additional, relativistically-broadened \\fekalfa line emission that may be present, and indeed, in many cases may be the {\\it only} component of the \\fekalfa line (e.g. see Guainazzi, Bianchi, \\& \\dovname 2006; Nandra \\etal 2007; Miller 2007; Turner \\& Miller 2009; Bianchi \\etal 2009). Measurement of the properties of the core of the \\fekalfa line in AGN is important for two principal reasons. One is to constrain the physical properties of the large-scale structure in the central engine. The peak energy of the \\fekalfa line constrains the ionization state of the line-emitting matter, and the width of the line gives kinematic information that can be used to estimate the size and location of the X-ray reprocessor. The equivalent width (EW) of the \\fekalfa line is a function of geometry, column density, covering factor, element abundances, and orientation of the line-emitter. Another reason why spectroscopy of the \\fekalfa line core is important is that it is necessary to model the narrow component of the line in order to deconvolve any relativistically-broadened emission-line component that may be present. The \\chandra high energy grating (HEG; see Markert \\etal 1995) is still unsurpassed in spectral resolution in the Fe~K band, which at 6.4~keV is $\\sim 39$~eV, or $\\sim 1860 \\ \\rm km \\ s^{-1}$ FWHM. This is a factor of $\\sim 4$ better than the spectral resolution of X-ray CCD detectors aboard \\xmm and \\suzakup. Although broad \\fekalfa emission lines are better studied with CCD spectrometers (due to their higher throughput), the \\chandra HEG is well-suited for studying the narrow core of the \\fekalfa line. One can then utilize the HEG measurements to deconvolve narrow and broad \\fekalfa line components in lower spectral resolution data. In YP04 the results of a uniform analysis of the properties of the \\fekalfa emission-line core were presented, based on \\chandra HEG data of a modest sample of fifteen AGN. There are now a larger number of \\chandra HEG observations for which the data are available, and in the present paper we extend the study of YP04 to include 82 observations of 36 unique AGN. The paper is organized as follows. In \\S\\ref{data} we describe the observations and data. In \\S\\ref{hegspec} we describe the methodology and basic spectral-fitting results. In \\S\\ref{properties} we discuss the implications of the results for the properties of the core of narrow \\fekalfa emission line in the HEG AGN sample. In \\S\\ref{baldwineffect} we investigate whether the narrow core of the \\fekalfa line as isolated by the HEG, supports the so-called X-ray Baldwin effect (an anti-correlation between the line EW and X-ray luminosity and between the EW and a proxy for the accretion rate). In \\S\\ref{summary} we summarize our results and findings. ", "conclusions": "\\label{summary} We have presented an empirical and uniform analysis of the narrow core of the \\fekalfa emission line in a sample of 82 observations of 36 AGNs with low to moderately low X-ray absorption ($N_{H}<10^{23} \\ \\rm cm^{-2}$), using \\chandra HEG data. The \\fekalfa line was detected in 33 sources, and its centroid energy was measured in 32 sources (68 observations). The distribution in the centroid energy is strongly peaked around $\\sim 6.4$~keV, with over 80\\% of the measurements lying in the range 6.38--6.43~keV. Including the statistical errors and utilizing the best measurements for each source, the line centroid energy lies entirely in the range 6.35--6.47~keV for 30 out of 32 sources. Thus we confirm, for the largest sample of AGN observed with such a high spectral resolution (FWHM~$\\sim 1860 \\ \\rm km \\ s^{-1}$ at 6.4~keV), the ubiquity of the narrow core of \\fekalfa line, and its preferred origin in cool, neutral or only mildly-ionized matter. The equivalent width (EW) of the core of the \\fekalfa line was constrained in 70 out of 82 observations, with only upper limits obtained from the remaining 12 spectra. The weighted mean EW was $53 \\pm 3$~eV, and $\\sim 70\\%$ of the individual measurements had a 68\\% confidence upper limit on the EW of less than 100~eV. { Similar results were obtained when considering the EW distribution by source, although the weighted mean was somewhat higher from measurements that allowed the intrinsic line width to be free ($70\\pm4$~eV).} The EW distribution can be produced by both Compton-thin and Compton-thick matter distributions and a more detailed analysis with a physical model is required to distinguish between the two scenarios. We also presented measurements of the flux of the core of the \\fekalfa line and found that for sources that had multiple observations, there was no case in which the line flux varied between observations, within the statistical errors. The intrinsic width of the core of the \\fekalfa line was measured for 27 sources (53 observations) and we obtained a weighted mean value of FWHM~$=2060 \\pm 230 \\ \\rm km \\ s^{-1}$ { (or $=2200 \\pm 220 \\ \\rm km \\ s^{-1}$ when considering measurements by source, not by observation)}. Of the 27 sources, 12 yielded 99\\% confidence, two-parameter contours of line flux versus FWHM that were good enough to investigate the relation between the width of the \\fekalfa line and the width of the H$\\beta$ line (or Br$\\alpha$ for one of the sources). We found that the ratio of the X-ray to optical line width varies from source to source. The 99\\% confidence, two-parameter upper limit lies in the range $\\sim 0.5-4$ for the 12 sources. This means that contributions to the flux of the core of the \\fekalfa line are allowed down to a factor $\\sim 0.7-2$ times the radius of the optical BLR. The upper limit on the size of the X-ray line emitter is not constrained because line flux contributions from large, parsec-scale distances could be unresolved by the HEG. We note that our results are suggestive of the fact that the location of the X-ray line-emitting region relative to the BLR may actually be different in different sources. These conclusions are subject to the caveat that derivation of the true velocity width of the \\fekalfa line core requires a proper physical model, such as that of MY09, that includes a possible Compton shoulder. This will be the subject of future work. However, we note that such an analysis can only reduce the derived velocity widths of the \\fekalfa lines. Finally, having isolated the narrow core of the \\fekalfa line with the best available spectral resolution we confirm the anti-correlation (albeit with a large scatter) between the line EW and X-ray luminosity, $L_{x}$ (the X-ray Baldwin effect), and between the line EW and $L_{x}/L_{\\rm Edd}$. We thank the referee for his/her useful comments. Partial support for this work was provided by NASA through \\chandra Award AR8-9012X, issued by the Chandra X-ray Observatory Center, which is operated by the Smithsonian Astrophysical Observatory for and on behalf of the NASA under contract NAS8-39073. X.W.S. and J.X.W. acknowledge support from Chinese National Science Foundation (Grant No. 10825312, 10773010), and Knowledge Innovation Program of CAS (Grant No. KJCX2-YW-T05). This research made use of the HEASARC online data archive services, supported by NASA/GSFC. This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with NASA. The authors are grateful to the \\chandra instrument and operations teams for making these observations possible." }, "1003/1003.4660_arXiv.txt": { "abstract": "Loop quantum cosmology provides an efficient framework to study the evolution of the Universe beyond the classical Big Bang paradigm. Because of holonomy corrections, the singularity is replaced by a ``bounce.\" The dynamics of the background is investigated into the details, as a function of the parameters of the model. In particular, the conditions required for inflation to occur are carefully considered and are shown to be generically met. The propagation of gravitational waves is then investigated in this framework. By both numerical and analytical approaches, the primordial tensor power spectrum is computed for a wide range of parameters. Several interesting features could be observationally probed. ", "introduction": "Loop quantum gravity (LQG) is a nonperturbative and background-independent quantization of general relativity. Based on a canonical approach, it uses Ashtekar variables, namely SU(2) valued connections and conjugate densitized triads. The quantization is obtained through holonomies of the connections and fluxes of the densitized triads (see, {\\it e.g.}, \\cite{rovelli1} for an introduction). Basically, loop quantum cosmology (LQC) is the symmetry reduced version of LQG (although it is fair to underline that the relations with the full theory are still to be investigated into the details). While predictions of LQC are very close to those of the old quantum geometrodynamics theory in the low curvature regime, there is a dramatic difference once the density approaches the Planck scale: the big bang is replaced by a big bounce due to huge repulsive quantum geometrical effects (see, {\\it e.g.}, \\cite{lqc_review} for a review). Among the successes of LQC, one can cite: the excellent agreement between the trajectories obtained in the full quantum theory and the classical Friedman dynamics as far as the density in much below the Planck scale, the resolution of past and future singularities, the ``stability\" of states which remain sharply peaked even after many cycles (in the k=1 case) and the fact that initial conditions for inflation are somehow naturally met. The latter point is especially appealing as the inflationary scenario is currently the favored paradigm to describe the first stages of the evolution of the Universe (see, {\\it e.g.}, \\cite{linde} for a recent review). Although still debated, it has received many experimental confirmations, including from the WMAP 7-Years results \\cite{wmap}, and solves most cosmological paradoxes. It is rather remarkable that, as will be explained in this paper, the canonical quantization of general relativity naturally leads to inflation without any fine tuning. Inflation would have been unavoidably predicted by LQC, independently of its usefulness in the cosmological paradigm.\\\\ Two main quantum corrections are expected from the Hamiltonian of LQG when dealing with a semiclassical approach, as will be the case in this study mostly devoted to potentially observable effects. The first one comes from the fact that loop quantization is based on holonomies, {\\it i.e.} exponentials of the connection rather than direct connection components. The second one arises for inverse powers of the densitized triad, which when quantized become an operator with zero in its discrete spectrum thus lacking a direct inverse. As the status of \"inverse volume\" corrections is not clear due to the fiducial volume cell dependence, this work focuses on the holonomy term only and derives, for the first time in a fully consistent way, the entire dynamics up to the explicit computation of the tensor power spectrum. The background evolution is first studied and a specific attention is paid to the investigation of the inflationary stage following the bounce. Then, analytical formulas are given for the primordial tensor spectrum for either a pure de Sitter or a slow-roll inflation. Finally, numerical results are given for many values of the parameters of the model. ", "conclusions": "This study establishes the full background dynamics in bouncing models with holonomy corrections. Although this has already been claimed before, we confirm that due to the sudden change of sign of the Hubble parameter, inflation is nearly unavoidable. In this paper, we have considered a particular model of inflation where the content of the universe is dominated by a massive scalar field. We have investigated the details by both analytical and numerical studies the primordial power spectrum of gravitational waves. It exhibits several characteristic features, namely a $\\mathcal{P}_{\\text{T}}\\propto k^2$ IR power suppression, oscillations, and a bump at $k_*$. In the UV regime, the standard inflationary spectrum $\\mathcal{P}_{\\text{T}}\\propto k^{-2\\epsilon}$ is recovered. The primordial tensor power spectrum transforms into $B$-type CMB polarization. The performed investigations therefore open the window for observational tests of the model, in particular through the amplification which occurs while approaching $k \\rightarrow k_*$. The observed structures correspond to the UV region in the spectrum. If the present scale $\\lambda_* \\sim 1/k_*$ is not much larger than the size of horizon, then the effects of the bounce should be, in principle, observable. In particular, one should expect amplification, rather than suppression of the $B$-type polarization spectrum at the low multipoles. The suppression for $k 20$ (UH = the cluster's cores). The fraction of early-type galaxies increases with the log of the over-density. A well defined \"red sequence\" composed of early-type galaxies exists in all environments at high luminosity, but it lacks of low luminosity (dwarf) galaxies in the lowest density environment. Conversely low luminosity isolated galaxies are predominantly of late-type. In other words the low luminosity end of the distribution is dominated by red dE galaxies in clusters and groups and by dwarf blue amorphous systems in the lowest density regions. At $z$=0 we find evidence for strong evolution induced by the environment (Nurture). Transformations take place mostly at low luminosity when star forming dwarf galaxies inhabiting low density environments migrate into amorphous passive dwarf ellipticals in their infall into denser regions. The mechanism involves suppression of the star formation due to gas stripping, without significant mass growth, as proposed by Boselli et al. (2008a). This process is more efficient and fast in ambients of increasing density. In the highest density environments (around clusters) the truncation of the star formation happens fast enough (few 100 Myr) to produce the signature of post-star-burst in galaxy spectra. PSB galaxies, that are in fact found significantly clustered around the largest dynamical units, represent the remnants of star forming isolated galaxies that had their star formation violently suppressed during their infall in clusters in the last 0.5-1.5 Gyrs, and the progenitors of future dEs. } ", "introduction": "In 1986 Valerie de Lapparent, Margaret Geller and John Huchra announced their \"Slice of the Universe\", a strip of 1100 galaxy optical redshifts from the CFA redshift survey, going through the Coma cluster (de Lapparent et al. 1986). A spectacular connected filament of galaxies popped out from their data, later nicknamed the \"Great Wall\" by Ramella et al. (1992, who more than doubled the number of redshift measurements) running through the spring sky at an approximate distance of 100 Mpc (adopting the modern value of $Ho$=73 km s$^{-1}$ Mpc$^{-1}$, as in the remaining of this paper), and spanning a similar length. Its narrowness in the redshift space, bracketing large voids, was shocking for the epoch, for an universe supposedly \"homogeneous and isotropic\" on that scale.\\\\ Hundreds individual galaxies in the region $11.5^h99\\%$ by Stoughton et al. (2002).} (with effective surface brightness $\\mu_e<24.5 \\rm ~mag~ arcsec^{-2}$).\\\\ The present \"near-field cosmology\" project makes use of the photometry and the spectroscopy provided by SDSS much in the way astronomers use telescopes in \"service observing\" mode, but counting on an almost unlimited amount of observing time. By saying so we wish to emphasize that, opposite to most SDSS statistical analyzes that are based on such large number of objects that cannot be inspected individually, the number of galaxies ($\\sim$ 4000) in the present study is large enough to ensure statistical significance, but not too large to prevent individual inspection of all images and spectra. The Coma supercluster constitutes the ideal bridge between the near-field and the high-$z$ cosmology.\\\\ Several modern studies based on hundreds of thousands galaxies from the SDSS in the local universe (usually $0.05 -19$) galaxies, as summarized in Fig. \\ref{fig14}. \\item Although the migration takes place also in loose structures it is more efficient in denser environments, such as clusters of galaxies (\"Nurture\"). \\item Quenching of the star formation is induced in clusters by fast gas-ablation mechanism, namely ram-pressure. \\item Post-Star-Burst galaxies appear the tracers of the migration taking place in dense environments. \\end{enumerate} \\begin{figure} \\centering \\includegraphics[width=8.0cm]{fig14.ps} \\caption{The $g-i$ color versus $i$-band absolute magnitude relation for late-type galaxies (blue contours) and for early-type galaxies (red contours). The arrow marks the effect of Nurture at $z$=0 in dense environments. Low-luminosity blue galaxies have their star formation truncated as a result of gas stripping. \\label{fig14}} \\end{figure}" }, "1003/1003.4984_arXiv.txt": { "abstract": "We describe a accurate and fast pixel-based statistical method to interpolate fields of arbitrary spin on the sphere. We call this method Fast and Lean Interpolation on the Sphere (\\flints{}). The method predicts the optimal interpolated values based on the theory of isotropic Gaussian random fields and provides an accurate error estimate at no additional cost. We use this method to compute lensed Cosmic Microwave Background (CMB) maps precisely and quickly, achieving a relative precision of $0.02\\%$ at a \\healpix{} resolution of $N_\\text{side}=$ 4,096, for a bandlimit of $\\ell_\\text{max}=$ 4,096 in the same time it takes to simulate the original, unlensed CMB map. The method is suitable for efficient, distributed memory parallelization. The power spectra of our lensed maps are accurate to better than 0.5\\% at $\\ell=$ 3,000 for the temperature, the $E$ and $B$ mode of the polarization. As expected theoretically, we demonstrate that, on realistic cases, this method is between two to three orders of magnitude more precise than other known interpolation methods for the same computational cost. ", "introduction": "Gravitational lensing of the Cosmic Microwave Background (CMB) is a unique probe of the distribution of mass in the entire visible Universe. The first signatures of this CMB lensing signal have already been seen in cross-correlation with large scale structure templates \\citep{SZD07,HHNS08}. The effect of lensing on the CMB power spectra will provide powerful additional constraints on the physics of the dark sector \\citep{AS03,SHK06}. Including reconstruction of the lensing potential in cosmological parameter analysis removes further degeneracies. Once the lensing deflection is mapped with high precision it will likely allow detecting the absolute mass scale of neutrinos \\citep{KKS03,LPPP06} and provide tight constraints on the presence of dark energy at redshifts $z\\sim2$. Lensing creates polarization $B$ modes by rotating the stronger $E$ modes \\citep{ZS97}. These lensing $B$ modes are not affected by cosmic variance and can therefore significantly improve reconstructions of the lensing potential \\citep{SCR06}. The modeling of lensing $B$ modes is further motivated since they are a foreground for the search for the $B$-mode signal of inflationary graviational waves \\citep{KS02,KCK02}. \\planck{} will yield the first all-sky temperature and polarization maps of the CMB with sufficient signal to noise to detect this signal. Ground based experiments such as QUBIC \\citep{QUBIC}, SPTpol \\citep{SPTpol} or QUIET \\citep{QUIET} contain detailed information about the fine-scale structure imprinted by lensing while still covering significant portions of the sky. For the extraction of lensing and polarization science from these data sets, fast and precise methods to simulated lensed all-sky CMB maps are indispensable. In the Born approximation, which specifies that the lensing effects may be modeled by the impact of a single lens plane on the CMB, the lensing corresponds to looking up at the value of the temperature or the polarization at displaced position on the sky. Thus, making lensed maps is effectively a resampling of the CMB on a different set of positions. Different technical solutions were proposed. Currently three main techniques are known. The first technique consists in doing a brute force resummation of the spherical harmonics at a new set of positions. Acceleration methods using symmetry properties of the sampling positions can not be used. While prohibitively expensive for practical use, this method is exact and we therefore employ it on a subset of pixels as a precision benchmark for other methods. The second technique is polynomial interpolation on the sphere to compute the value of the temperature and polarization field at the displaced positions. The most simple of the interpolation technique is a bilinear interpolation of fields on the sphere (hereafter called the naive technique). Another algorithm, a Lagrange polynomial interpolation on the Equi-cylindrical Projection, was first proposed by \\cite{HPSSB04} and then later used by \\cite{DB08} for a full-sky simulation of the CMB lensing on a light cone of a cosmological simulation. A variant of this method, involving a bicubic interpolation scheme, is implemented in the publicly available code \\lenspix\\footnote{\\url{http://cosmologist.info/lenspix}} described in \\cite{Lewis05} and \\cite{HL08} (hereafter the {\\sc ECP} algorithm). A third technique has been recently developed by \\cite{Basak08}. This algorithm consists in recasting the spherical harmonic coefficient of the unlensed field into the Fourier basis in the $(\\theta,\\phi)$ variables. Then, we may use a Non-equispaced Fast Fourier Transform (NFFT) to compute the field at the displaced positions. This method achieves high precision and is significantly faster than the first technique though still too slow to allow for the production of a large set of lensed maps, which is required for the statistical analysis of observed CMB data. We propose a fourth technique that relies on the statistical properties of the considered fields to be lensed. It is based on the idea of interpolating the original field but using the known spectral information to compute the correct weighing coefficients for the interpolation. This method is related to a Wiener filter \\citep{Wiener49} but not limited to pure Gaussian random field. We call this method Fast and Lean Interpolation on the Sphere (\\flints{}).\\footnote{The reference implementation is written in C++/OpenMP and is available at \\url{http://www.iap.fr/users/lavaux/flints.php}.} In this work, we analyze the result obtained based on a \\healpix{} pixelization. We note that our framework may be used on any pixelization of the sphere, including the Equi-cylindrical projection already used in previous works. The implementation that we propose takes advantage of the geometrical properties of \\healpix{} for a number of memory and computational optimizations. We note that our interpolation method is of general interest beyond its application to lensing. In fact, since it is based on Wiener filtering it is guaranteed to give the best possible mean squared error of any method for a field with the same power spectrum and for fixed interpolation stencil. The structure of this paper is the following. In Section~\\ref{sec:interpolating}, we describe the general interpolation method. Then, in Section~\\ref{sec:applens}, we discuss its performance in generating lensed CMB maps. In Section~\\ref{sec:conclusion}, we conclude. ", "conclusions": "\\label{sec:conclusion} For a given interpolation stencil, we describe the optimal interpolation technique for isotropic band-limited fields of arbitrary spin, sampled on the sphere. Taking advantage of the symmetry properties of the \\healpix{} pixelization, the method is fast and memory-efficient. To test this approach we implement a supersampling filter for \\healpix{} temperature and polarization maps. A Monte Carlo study confirms both the predicted precision and our estimates of memory and CPU time scaling. We demonstrate this interpolation method to be powerful tool to simulate lensed CMB temperature and polarization maps from unlensed maps. Our Monte Carlo comparison to exact reference maps computed by \\lenspix{} and to predicted lensed power spectra by \\camb{} demonstrate that we achieve an accuracy which exceeds the requirements of the \\planck{} data \\citep{planckbook} while reducing the required computational time by an order of magnitude compared to using, e.g., the Torus method. In addition, the method allows very easy parallelization as the procedure is strictly local in pixel space. We compared the performance and the precision of our method to the Equi-cylindrical projection interpolation method. The two methods have similar speed. {\\flints{}} is more precise and has no catastrophic errors. We conclude that this method is a very promising technique in terms of speed, precision and scalability for the simulation of high resolution maps of the lensed CMB temperature and polarization anisotropies. \\flints{} enables us to produce lensed maps as cheaply as making a spherical harmonic transform, and makes us capable of producing thousands of simulations of the lensed sky within an acceptable computational time. This advance may allow us to run a full likelihood analysis of the lensing potential in observed CMB data with current computer technology, which is not possible with other known methods of computing lensed maps." } }